{
    "0801/0801.2870_arXiv.txt": {
        "abstract": "{The SKA will be a unique instrument with which to study the evolution of the gas content of galaxies. A proposed deep ($\\sim$ 8 Msec) `pencil-beam' survey is simulated using recently updated specifications for SKA sensitivity and survey speed. Almost $10^7$ galaxies could be detected in the redshifted 21cm line, most at redshifts in excess of two.  This will enable confident statements to be made about the evolution of the cosmic HI density and the HI mass function to $z=3$, corresponding to a lookback time of 11 Gyr. However, galaxies or groups of galaxies with masses the same as the most HI-massive galaxies at $z=0$ will be detectable at redshifts of 6, if they exist. The ideal instrument for studying HI evolution would have an instantaneous sensitivity at least a factor of two higher than current specifications in the critical frequency range 200-500 MHz, or $A/T>2\\times 10^4$ m$^2$ K$^{-1}$.  The capabilities of the SKA will be highly complementary to ALMA which will be able to study the evolution of the molecular gas component over the same redshift range.}  \\FullConference{From planets to dark energy: the modern radio universe\\\\ October 1-5 2007\\\\ University of Manchester, Manchester, UK}  \\begin{document} ",
        "introduction": "Understanding the evolution of galaxies is one of the goals of modern cosmology and one of the five key science goals of the SKA \\cite{cr04}. A key ingredient in galaxy evolution is the gas supplied through various accretion, merger and feedback processes that occur during the assembly of galaxies. This gas, which is mainly hydrogen, passes through a neutral atomic phase and later condenses into massive molecular clouds and stars. During the atomic phase, it can readily be traced with the 21cm hyperfine spin-flip transition and the Gunn-Peterson effect.  At high redshifts, $z>1.6$, neutral hydrogen is currently traced by ground-based observations of Gunn-Peterson absorption lines against bright background QSOs. Such observations demonstrate that the bulk of the neutral hydrogen is in the Damped Ly-$\\alpha$ systems with column densities exceeding $2\\times10^{20}$ cm$^{-2}$. However, interpretation of these observations tends hampered by insufficient lines of sight and serious uncertainties associated with dust obscuration, gravitational lensing and intrinsic source size. These uncertainties lead to contradictory results from measurements associated with bright QSOs, faint QSOs and Gamma Ray Bursts \\cite{Jas05} \\cite{Pro06} \\cite{Por07}. However, with the SKA, galaxies will be detected at similar redshifts in 21cm line emission, which will lead to a clearer understanding of the distribution of gas in the Universe, and the manner in which the gas content of galaxies evolves with time.  In order to measure the gas content of galaxies at the highest redshift, a deep pencil beam survey with the SKA is proposed. Currently proposed SKA specifications \\cite{ska07} are adopted, and used to generate artificial galaxy catalogues which are used to simulate the accuracy with which one simple parameter can be recovered, namely the cosmic HI density -- the comoving volume density of neutral hydrogen. ",
        "conclusions": "This simulation has demonstrated the large numbers of high-redshift galaxies that can be detected by the SKA in a significant, but feasible, HI survey of galaxies, and has demonstrated the high accuracy with which the cosmic HI density can be measured. Whilst galaxy numbers are low at redshifts below 1.8, this is largely due to the small field of view available to single pixel technology. Adoption of widefield detection technology has the potential to greatly increase numbers. Although the subsequent reduction of errors associated with shot noise and cosmic variance is important for many science goals, greater gains in the field of galaxy evolution study are likely to be made with better instantaneous sensitivity at lower frequencies. This will increase the ability to detect changes in the HI mass function and reduce the density extrapolation required to account for low-mass galaxies at redshifts approaching 3. An appropriate goal of $A/T>2\\times 10^4$ m$^2$ K$^{-1}$ is therefore suggested for frequencies below 500 MHz."
    },
    "0801/0801.4517_arXiv.txt": {
        "abstract": "Models with dark energy decaying into dark matter have been proposed in Cosmology to solve the coincidence problem. We study the effect of such coupling on the cosmic microwave background temperature anisotropies. The interaction changes the rate of evolution of the metric potentials, the growth rate of matter density perturbations and modifies the Integrated Sachs-Wolfe component of cosmic microwave background temperature anisotropies, enhancing the effect. Cross-correlation of galaxy catalogs with CMB maps  provides a model independent test to constrain the interaction. We particularize  our analysis for a specific interacting model and show that galaxy catalogs with median redshifts $z_m=0.1-0.9$ can rule out models with an interaction parameter strength of $c^2\\simeq 0.1$ better than 99.95\\% confidence level. Values of $c^2\\le 0.01$ are compatible with the data and may account for the possible discrepancy between the fraction of dark energy derived from WMAP 3yr data and the fraction obtained from the ISW effect. Measuring the fraction of dark energy by these two methods could provide evidence of an interaction. ",
        "introduction": "Measurements of luminosity distances using supernovae type Ia (SNIa) \\cite{riess}, of the cosmic microwave background (CMB) temperature anisotropies with the WMAP satellite \\cite{wmap3}, large scale structure \\cite{lss,cole}, the integrated Sachs--Wolfe effect \\cite{fosalba-gazta2003,boughn-crittenden}, and weak lensing \\cite{weakl}, indicate that the Universe has entered a phase of accelerated expansion -see \\cite{reviews} for recent reviews. This acceleration is explained in terms of an unknown and nearly unclustered matter component of negative pressure, dubbed ``dark energy\" (energy density, $\\rho_{x}$), that currently contributes with about $75\\%$ of the total density. The remaining $25\\%$ is shared between cold dark matter ($\\rho_{c} \\sim 21\\%$), baryons ($\\rho_{b} \\sim 4\\%$), and negligible amounts of relativistic particles (photons and neutrinos). The pressureless non-relativistic matter component redshifts faster with expansion than the dark energy, giving rise to the so-called `coincidence problem' that seriously affects many models of late acceleration, particularly the $\\Lambda$CDM model: ``Why are matter and dark energy densities of the same order precisely today?\" \\cite{steinhardt}. To address this problem within general relativity one must either accept an evolving dark energy field or adopt an incredibly tiny cosmological constant and admit that the ``coincidence\" is just that, a coincidence that might be alleviated by turning to anthropic ideas \\cite{anthropic}. (It should be mentioned, however, the existence of proposals in which a vacuum energy density of about the right order of magnitude arises from the Casimir effect at cosmic scales -see \\cite{emili} and references therein).  One way to address the coincidence problem, within general relativity, is to assume an interaction (coupling) between the dark energy component and cold dark matter such that the ratio $r \\equiv \\rho_{c}/\\rho_{x}$ evolves from a constant but unstable value at early times (in the radiation and matter dominated epochs) to a lower, constant and stable attractor at late times well in the accelerated expansion era \\cite{iqm0,iqm}. Before the late accelerated expansion was observed, Wetterich introduced interacting quintessence models to reduce the theoretical high value of the cosmological constant \\cite{wetterich}. Later, these kind of models were rediscovered, sometimes (but not always) in connection with the coincidence problem -see, e.g. \\cite{in-connection}. Other solutions (known as $f(R)$ models) require the Einstein-Hilbert action to be modified \\cite{f(R)}.  Interacting quintessence models are testable since they predict differences in the  expansion rates of the Universe, the growth of matter density perturbations, the Cosmic Microwave Background (CMB) temperature anisotropies and in other observables. In this paper we shall demonstrate that, as the Integrated Sachs-Wolfe effect measures directly the growth rate of matter density perturbations, can be used to detect variations on the evolution of the large scale structure with respect to the prediction of the concordance $\\Lambda$CDM model. As a toy model,  we shall consider a spatially flat Friedmann--Robertson--Walker universe filled with radiation, baryons, dark matter and dark energy such that the last two components interact with each other through a coupling term, $Q = 3 {\\cal H} c^{2} (\\rho_{c} +\\rho_{x})$. Thus, the energy balance equations for dark matter and dark energy take the form \\\\ \\begin{equation} \\dot{\\rho}_{c}+3{\\cal H} \\rho_{c} = 3 {\\cal H} c^{2} (\\rho_{c} +\\rho_{x})\\,  ,\\qquad \\dot{\\rho}_{x}+3 {\\cal H} (1+w_{x})\\rho_x= -3{\\cal H} c^{2} (\\rho_{c} +\\rho_{x}) \\, , \\label{balance} \\end{equation} \\\\ where $w_{x} = p_{x}/\\rho_{x} < -1/3$ is the equation of state parameter of the dark energy component and $c^{2}$ is a constant parameter that measures the strength of the interaction (it should not be confused with the speed of light that we set equal to unity). Derivatives are taken with respect to the conformal time, and ${\\cal  H}= \\dot a/a$, with $a$ the scale factor of the Friedmann-Robertson-Walker metric. To satisfy the severe constraints imposed by local gravity experiments \\cite{peebles_rmph, hagiwara}, baryons and radiation couple to the other two energy components only through gravity. Our ansatz for $Q$ guarantees that the ratio between energy densities, $r$, tends to a fixed, attractor, value at late times. It yields  a constant but unstable ratio at early times. The details of the calculation can be found in \\cite{iqm,olivares1,olivares2,olivares3}. This result holds irrespective of whether the dark energy is a quintessence (i.e., $ -1 <w_{x} < -1/3$), phantom ($w_{x} < -1$), $k$-essence or tachyon scalar field \\cite{dualk}. It is worth to remark that the above ansatz implies  an effective power law potential for the quintessence dark energy field at early times and an effective exponential potential at late times \\cite{iqm0}.  Beyond the physical motivation for the existence of a coupling between dark matter and dark energy, their interaction parameter $c^{2}$ needs to be constrained by observations. In \\cite{olivares2} we showed that the slope of the matter power spectrum at small scales decreases and the scale of matter radiation equality is shifted to larger scales with the interaction.  CMB temperature anisotropies are equally affected if the interaction is present at decoupling. The luminosity distance of SNIa depends on the expansion rate of the Universe and can be used to test the existence of a DM/DE coupling. Comparison of model predictions with data showed that only WMAP 3yr data and the Sloan Digital Sky Survey (SDSS) data provided constraints on the value of $c^2$ \\cite{olivares2,olivares3}. In particular, the SNIa data proved to be rather insensitive to the coupling. In this paper we show that the Integrated Sachs-Wolfe (ISW) component of the CMB temperature anisotropies is much more sensitive than the SNIa data at redhisfts $z\\sim 1$. It measures the growth factor of linear matter density perturbations, that is strongly affected by the interaction and can effectively set constraints on Interacting Quintessence Models (IQM). Other experiments that also constrain the growth rate are discussed in \\cite{amendola}. In Section II we briefly summarize the formalism behind the measurements of the ISW effect by cross-correlating CMB data with galaxy catalogs. In Section III we describe the observational prospects and in Section IV we present our conclusions. ",
        "conclusions": "The ISW effect may provide an excellent tool to detect, in a model independent way, the coupling between dark energy and dark matter at low redshifts, i.e., $z\\le 1$. The interaction damps the growing mode of the Newtonian potential faster compared to models with no interaction and  enhances the ISW effect; then, the ISW effect is sensitive to the growth rate of matter density perturbations, but model predictions need to be refined to model more realistically the selection function of galaxy catalogs. We have shown that the correlation of WMAP with templates constructed from different galaxy catalogs can rule out an interaction parameter of $c^2=0.1$ at better than the 99.95\\% confidence level. For this type of models, this upper limit is stronger than those obtained previously using the CMB radiation power spectrum \\cite{olivares1,olivares3}, but at a different redshift. The ISW effect shows more statistical power to discriminate among IQM models than the luminosity distance test obtained using SNIa.  The recent results in \\cite{cabre,cabre2007} suggest a discrepancy between the fraction of dark energy necessary to explain the amplitude of the ISW effect (measured from cross-correlating CMB data and galaxy catalogs) larger than allowed by the WMAP 3yr data. We have shown that the tension between both data sets could be explained by introducing a moderate coupling between dark energy and dark matter. This is in line with recent claims that the Layzer-Irvine equation \\cite{layzer} when applied to  galaxy clusters reveals the existence of a DM/DE coupling \\cite{clusters}. If the aforesaid discrepancy is confirmed by deeper and wider galaxy samples, it could turn out an indirect evidence for the existence of a DM/DE interaction."
    },
    "0801/0801.3271_arXiv.txt": {
        "abstract": "Gravitational collapse of dark matter, merger of dark matter haloes and tidal disruption of satellites are among processes which lead to the formation of fine and dense dark matter shells, also known as dark matter caustics. The putative weakly interacting species which may form the dark matter are expected to strongly annihilate in these dense regions of the Milky Way halo and generate in particular antiprotons and positrons. We derive the flux of these rare antimatter particles at the Earth and show that it depends significantly on the cut-off radius of the dark matter distribution at the galactic centre. Boost factors of $\\sim 30$ are found with respect to a smooth NFW profile for high-energy antiprotons and low-energy positrons if this cut-off radius is taken to be 300 pc -- a somewhat extreme value though. This yields a detectable antiproton signal around hundreds of Gev in models where the annihilation cross section today is enhanced by non--perturbative effects as in the generic case of a heavy Wino. However, dark matter caustics cannot provide a better explanation for the HEAT excess reported above $\\sim$ 10 GeV than a smooth NFW or isothermal cored distribution.  \\noindent {\\small\\sf preprint no: LAPTH-1235/08} ",
        "introduction": "\\label{sec:introduction}  The halo of our galaxy is believed to have formed from the gravitational collapse of dark matter (DM) and a continuous merger with other haloes. The formation of haloes, described by the Jeans-Vlasov-Poisson equation, proceeds with the formation of dark matter {\\it shells}, also referred to as dark matter {\\it caustics} \\citep{sikivieipser,sikivietkachev,tremaine,alardcolombi,natarajan}. As the density peaks collapse under self-gravity, at the surfaces of (formally) infinite density where dark matter streams meet, caustics emerge. The infalling satellites are also distrupted by a mechanism similar to the formation of the {\\it primordial} caustics. As a satellite moves inside the potential well of its host halo, a tidal tail forms, and around the apapsis of the orbit high-density shells or caustics emerge [see {\\it e.g.} \\citet{hayashi}]. Since each infalling satellite also has his own caustics, the hierarchical formation of structures yields a hierarchy of caustics. Observational examples for the caustics are the shell galaxies, where shells of stars form during the merger of galaxies, by a mechanism very similar to the formation of dark matter caustics \\citep{malincarter,carterallen,hernquistquinn1,hernquistquinn2}.  Once formed, they cannot be destroyed~: caustics are permanent structures in real and in phase space. However, their over-density with respect to the background density can diminish with time due to continuous increase in the number of streams. Therefore, it is likely that we are surrounded by a large number of primordial shells and fossil shells of disrupted satellites. The density of the shells can exceed that of the diffuse background by an amount which depends on the age of the galaxy and the number of times the satellite has wrapped around our galaxy. Due to their vicinity and over-density, these shells can be important for dark matter search experiments.  The nature of dark matter remains a mystery. Supersymmetry and extra-dimension extensions of the standard electroweak model provide a natural candidate in the form of a weakly interacting and massive particle (hereafter WIMP). These species should fill up the galactic halo. If dark matter consists of WIMPS then they are expected to strongly annihilate in the dense regions of our halo ({\\it i.e.} in the caustics) and generate in particular gamma-rays and charged cosmic rays such as positrons and antiprotons. Indeed, due to their very short diffusion length, positrons can be excellent tracers of nearby caustics. Although the implication of substructures, or surviving satellites, in boosting the cosmic-ray signal has been studied extensively [see {\\it e.g.} \\citet{Lavalle:2006vb,Lavalle_Maurin}], the signature of dark matter caustics on cosmic-rays has been rarely considered.  Here, we propose to study the boost in positron and antiproton signals due to annihilation in dark matter caustics. The ideal way to pursue this study would be to use a very high-resolution simulation. However, even the latest simulations have not yet achieved a high enough resolution to identify the caustics in a typical galaxy, although impressive progresses are being made \\citep{vogelsberger}. Analytic models for the formation of caustics are mainly based on the self-similar secondary infall \\citep{fillmoregoldreich,bertschinger}. We use a generalized version of the secondary infall model which takes into account the finite velocity dispersion of DM \\citep{mohayaeeshandarin}. Although secondary infall model has serious limitations because of the assumption of spherical symmetry and smooth accretion and because it ignores the hierarchical scenario of structure formation, it has so far provided a paradigm for the study of dark matter haloes [{\\it e.g.} see \\citet{ascasibar}].  In Section~2, we review and generalize the secondary infall model for dark matter with finite velocity dispersion and give an analytic fit for the density profile which includes the caustics for a typical galaxy like our own. The cosmic ray signature of dark matter caustics is studied in Section~3 for large antiproton and positron horizon radii. In Section~4, the possibility of detecting a nearby caustic through its short range positron signal is investigated. We finally conclude in Section~5.  \\begin{figure*} \\centering{ \\includegraphics[width=\\columnwidth]{trajectory-pruned} {\\hspace{0.2cm}} \\includegraphics[width=\\columnwidth]{phase-space-pruned}} \\caption{ Particle trajectories in real and phase spaces are shown for two values of the parameter $\\epsilon$ where $\\lambda$ is the non-dimensional radius and $\\xi$ is the non-dimensional time. The curves for $\\epsilon=1$ are featured in order to highlight the phase transition that occurs at $\\epsilon=2/3$. The physical origin for that transition is that for $\\epsilon<2/3$, at late times, the mass within a given radius $r$ is dominated by the contribution from particles from outside of $r$. In contrast for $\\epsilon>2/3$, the mass is dominated by particles within $r$ (see \\citep{fillmoregoldreich} for further details). } \\label{fig:trajectory} \\end{figure*} \\begin{figure*} \\centering{ \\includegraphics[width=\\columnwidth]{density-MW-pruned} {\\hspace{0.2cm}} \\includegraphics[width=\\columnwidth]{density-numeric-analytic-pruned}} \\caption{ The left panel features the total numerically-obtained density for $\\epsilon=0.2$ [solid black curve, using expression~(\\ref{eq:total density})]. This solution is compared in the right panel to our analytic fit [dashed red curve, using expression~(\\ref{eq:rho_total})]. } \\label{fig:total density} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusions}  We have studied whether DM caustics in the halo of the Milky Way can amplify the flux of cosmic ray antiprotons and positrons which is received at the Earth. We have used the secondary infall model for our halo which naturally includes the caustics and assumed that the galactic DM is made of weakly interacting massive species. We have then taken into account the smearing of the caustics due to a present-day velocity dispersion of these particles of $0.03$~cm~s$^{-1}$.  The cosmic ray antiprotons and positrons that are detected at the Earth originate from a region whose typical size is much larger than the shell thickness or even the shell separation. The associated horizon probes a large portion of the Milky Way and the coarse-grained average density~(\\ref{DM_shell_QD}) provides an adequate description. In the solar neighbourhood, the coarse-grained caustic density is the same as the smooth NFW or isothermal cored distributions usually assumed in the literature. The difference lies at the centre of the Milky Way. The lower-envelope density diverges there with an index of $\\sim 2.11$, hence a DM profile steeper than even in the NFW case. The coarse-grained shell density is assumed to be constant inside a sphere of radius $r_{\\rm cut-off}$ whose value is unknown. The smaller this cut-off radius, the more abundant DM at the galactic centre and consequently the stronger its signal should it reach the Earth.  We have then computed the antiproton flux which the coarse-grained shell density yields at the Earth. The reach of the antiproton sphere depends on the cosmic ray propagation model but is always of order a few kiloparsecs. The MIN configuration is associated to a rather small antiproton range. The associated signal is not different from what is derived assuming NFW or isothermal cored distributions. The MAX propagation model is characterized on the contrary by efficient diffusion taking over a moderate galactic convection. The antiproton sphere probes the inner and denser regions of the Milky Way. We find that the antiproton signal is enhanced by a factor of $\\sim 30$ should the conventional smooth NFW profile be replaced by the coarse-grained shell density~(\\ref{DM_shell_QD}) for which a cut-off radius of 300 pc has been assumed. The MED set of propagation parameters corresponds to the best fit to the B/C data and features the intermediate situation. It leads to the exciting possibility that the antiproton signal is only boosted above 10 GeV in the presence of caustics. Depending on the WIMP annihilation cross section, the antiproton flux could be severely distorted at high energy as shown in Fig.~\\ref{fig:flux_WINO} where no artificial boost factor is required. Thus a promising window opens up around a few hundreds of Gev where future antiproton measurements are eagerly awaited.  We are less optimistic for the positron signal. Depending on the cosmic ray propagation model, the positron flux at the Earth may be enhanced in the presence of shells with respect to a smooth NFW or isothermal cored DM profile. However, this situation arises only at low energy where the observations are already well explained by the sole secondary background component. Thus DM caustics cannot provide an explanation for the HEAT excess reported above $\\sim$ 10 GeV and consequently produced in the solar neighbourhood. The solution so far invoked is based on a WIMP with a hard positron annihilation spectrum like a Kaluza-Klein particle \\citep{Cheng:2002ej} or a neutralino with a dominant $W^{+}W^{-}$ channel \\citep{Hooper:2004bq,Delahaye:2007fr}. Boost factors of $\\sim$ 10 with respect to a smooth NFW DM halo are nevertheless necessary. Such a value is marginally possible \\citep{Lavalle_Maurin} in scenarios inspired by the $\\Lambda$-CDM N-body numerical simulations which point towards the existence of numerous and dense DM clumps inside which WIMP annihilation can be enhanced. We showed that the coarse-grained shell density~(\\ref{DM_shell_QD}) does not provide an alternative to explaining the HEAT excess. However, further data with higher precision and also a better understanding of the secondary positron background are needed to firmly exclude caustics as a possible explanation of the HEAT measurements. It remains very unlikely nevertheless that the present spectral form could be reproduced by caustics, at least by its coarse-grained distribution~(\\ref{DM_shell_QD}).  We finally explored the possibility that the fine-grained distribution~(\\ref{fine_grained}) could yield a strong positron flux should the Earth be embedded inside the densest part of a caustic, a region where the DM density reaches its peak value $\\rho_{k,{\\sf max}}$. At high energy, the positron sphere shrinks. Averaging the fine-grained shell density~(\\ref{eq:density-caustic}) by its coarse-grained approximation~(\\ref{DM_shell_QD}) is no longer possible. Positrons that are received at the Earth have short diffusion lengths $\\lD$ and hence can be excellent tracers of nearby caustics. We investigated here the case of a positron line. Positrons that are detected exactly at the line energy $E = E_{S} \\equiv m_{\\chi}$ have vanishing diffusion length and if the Earth sits exactly inside the shell, an enhancement of the positron signal by a factor of $\\sim 13,000$ is naively expected. However, energy is measured with a limited accuracy and the energy bin of the line has a non-vanishing width. Averaging $\\lD$ over the line bin leads to a diffusion length which is still far larger than the typical caustic thickness or even separation. Our study showed no difference between the results of Section~\\ref{sec:cr_signal} and those derived with the fine-grained distribution~(\\ref{fine_grained}) once energy is properly averaged. We hypothesize that an extremely high energy resolution, presently unavailable, would allow to detect the nearby caustics.  A word of caution is necessary at this stage though. The analysis of Section~\\ref{sec:draperies} is based on the assumption that cosmic rays diffuse on the inhomogeneities of the galactic magnetic field. The mean free path $\\lambda_{\\rm free}$ of their random walk may be derived from the space diffusion coefficient $K$ through the canonical relation \\be K(E) \\equiv {\\displaystyle \\frac{1}{3}} \\, \\lambda_{\\rm free} \\, \\beta \\;\\; . \\ee In this diffusion scheme and with the MED set of parameters, positrons with energy $\\epsilon$ cover on average a distance \\be \\lambda_{\\rm free} = 1.1 \\times 10^{-4} \\; {\\rm kpc} \\; \\epsilon^{0.7} \\label{lfree_close_line} \\ee before their next scattering on Alfv\\'en waves. Our treatment of cosmic ray propagation is definitely supported by the fact that $\\lambda_{\\rm free}$ is much smaller than the horizon size set by $\\lD$ and equation~~(\\ref{lD_close_line}). However, if the nearest DM caustic is very close to the Earth and lies at a distance which does not exceed $\\lambda_{\\rm free}$ or if the Earth is embedded inside the densest part of a shell, the diffusion hypothesis breaks down. Positrons will essentially drift along the lines of the magnetic field without encountering on their way any obstacle. Depending on the relative orientation of the local magnetic field with respect to the Earth and its nearest caustic, we could be possibly exposed to an intense high-energy positron flux whose evaluation is clearly beyond the scope of this article.   {\\small Acknowledgment~: We thank Pierre Brun and Julien Lavalle for having provided us with a few typical positron and antiproton spectra arising from WIMP annihilation. RM thanks Sergei Shandarin for contributions, Niayesh Afshordi, Ed Bertschinger and Mike Kuhlen for discussions and LAPTH Annecy for hospitality and French programms PNC \\& ANR (OTARIE) for travel grants. Special thanks are due to Mark Vogelsberger for a careful reading of the manuscript and many useful comments and corrections. }"
    },
    "0801/0801.1184.txt": {
        "abstract": "% context heading (optional) {} % aims heading (mandatory) {The aim of this work is to investigate the physical, structural and evolutionary properties of old, passive galaxies at $z>1.4$ and to place new constraints on massive galaxy formation and evolution.} % methods heading (mandatory) {We combine ultradeep optical spectroscopy from the GMASS project (\\emph{Galaxy Mass Assembly ultradeep Spectroscopic Survey}) with GOODS multi-band (optical to mid--infrared) photometry and HST imaging to study a sample of spectroscopically identified passive galaxies at $1.39<z<1.99$ selected from \\emph{Spitzer Space Selescope} imaging at 4.5$\\mu$m.} % results heading (mandatory) {A stacked spectrum with an equivalent integration time of $\\sim$ 500 hours was obtained and compared with libraries of synthetic stellar population spectra. The stacked spectrum is publicly released. The spectral and photometric SED properties indicate very weak or absent star formation, moderately old stellar ages of $\\approx$1 Gyr (for solar metallicity) and stellar masses in the range of $10^{10-11}$ M$_{\\odot}$, thus implying that the major star formation and assembly processes for these galaxies occurred at $z>2$. No X-ray emission was found neither from individual galaxies nor from a stacking analysis of the sample. Only one galaxy shows a marginal detection at 24$\\mu$m. These galaxies have morphologies that are predominantly compact and spheroidal. However, their sizes ($R_e \\lesssim$ 1 kpc) are much smaller than those of spheroids in the present--day Universe. Their stellar mass surface densities are consequently higher by $\\approx$1 dex if compared to spheroids at $z\\approx0$ with the same mass. Their rest-frame $B$-band surface brightness scales with the effective radius, but the offset with respect to the surface brightness of the local Kormendy relation is too large to be explained by simple passive evolution. At $z\\approx1$, a larger fraction of passive galaxies follows the $z\\approx0$ size -- mass relation. Superdense relics with $R_e \\approx$ 1 kpc are extremely rare at $z\\approx0$ with respect to $z>1$, and absent if $R_e<1$ kpc. Because of the similar sizes and mass densities, we suggest that the superdense passive galaxies at $1<z<2$ are the remnants of the powerful starbursts occurring in submillimeter--selected galaxies at $z>2$. The results are compared with theoretical models and the main implications discussed in the framework of massive galaxy formation and evolution. } % conclusions heading (optional), leave it empty if necessary {} ",
        "introduction": "Deep surveys provide the observational constraints needed to understand galaxy formation and evolution. In particular, many studies have been focused on massive galaxies (i.e. stellar mass ${\\cal M} > 10^{11}$ M$_{\\odot}$) as cosmological probes of the history of galaxy mass assembly. However, despite the remarkable success in finding and studying massive galaxies over a wide range of cosmic time, the global picture is far from being clear. Thanks to their simple and homogeneous properties (morphology, colors, passively evolving stellar populations, scaling relations) and being the most massive galaxies in the present-day Universe, early-type galaxies (ETGs) are crucial to investigate the cosmic history of massive galaxies (e.g. \\cite{alvio} and references therein).  At $z<1$, the most recent results seem now to agree in indicating that the majority of massive ETGs (${\\cal M} > 10^{11}$ M$_{\\odot}$) were already in place at z$\\approx$ 0.7-0.8, with a number density consistent with the one at z=0, whereas the evolution is more pronounced for the lower mass ETGs (e.g. \\cite{fon04,yamada,bundy06,cdr06,borch,scarlata,brown,bundy07}). This mass--dependent evolutionary scenario, known as \"downsizing\" (\\cite{cowie,gavazzi}), was proposed to explain the galaxy star formation histories, i.e. with massive galaxies forming their stars earlier and faster than the low mass ones. Recent results suggest that the downsizing concept should extend also to the stellar mass assembly evolution itself, i.e. with massive galaxies assemblying their mass earlier (e.g. \\cite{cdr06,bundy06,franceschini, bundy07,perez}), thus providing new and stringent constraints for the current models of galaxy formation (e.g. \\cite{delucia}).  The above results trace the evolution of the number density, luminosity and mass, but do not explain what is the mechanism with which ETGs assemble their mass and shape their morphology. Dissipationless ETG--ETG major merging (also called \"dry\" merging) has been advocated as an important mechanism to build up the masses of ETGs at $0<z<1$ (\\cite{bell06,vd05}). However, this scenario seems difficult to be reconciled with the properties of the small-scale clustering of ETGs at $0.16<z<0.36$ (\\cite{masjedi,masjedi2}), and (for massive ETGs) with the weak evolution of the high--mass end of the stellar mass function (e.g. \\cite{cdr06,bundy06,bundy07,pozz07}). An alternative process proposed recently to explain the assembly of ETGs is based on multiple, frequent minor mergers at moderate redshifts (\\cite{bournaud}).  Several studies suggest that the critical redshift range where the strongest evolution and assembly took place is at $1 \\lesssim z \\lesssim2$ (e.g. \\cite{fon04,glaze,abraham07,arnouts}). However, the picture at these redshifts is even more controversial than at $z<1$, partly because of the observational difficulty to identify spectroscopically and study large samples of ETGs at $z>1$. The few ETGs identified spectroscopically so far up to $z\\approx 2.5$ are very red ($R-K>5-6$), dominated by passively evolving old stars with ages of 1-4 Gyr, have stellar masses typically ${\\cal M} \\gtrsim 10^{11} M_{\\odot}$, and are strongly clustered with $r_0 \\approx$ 8--10 Mpc (\\cite{cim04,glaze,mcc04,daddi05a,sar05,kong,kriek}). Daddi et al. (2005) were the first to realize that a large fraction of these ETGs have smaller sizes ($R_e \\lesssim 1$ kpc) (see also \\cite{cassata05}) and higher mass internal densities than present--day ETGs. This result was soon confirmed by other observations (\\cite{trujillo06,zirm07,longhetti07,toft07, trujillo07}). However, it is still unclear how to explain such size -- mass -- density properties in the context of ETG evolution.  The existence of a substantial population of old, massive, passively evolving ETGs up to $z\\approx 2$ was not predicted in galaxy formation models available in 2004-2005, and opened the question on how it was possible to assemble such systems when the Universe was so young. A promising mechanism which can provide a better agreement with the observations seems to be the \"quenching\" of the  star formation at high redshifts with AGN \"feedback\" (e.g. \\cite{gra,menci06}).  The stellar ages and masses of the passive ETGs at $z\\approx1-2$ require precursors characterized by strong ($>100$ M$_{\\odot}$/yr) and short-lived (0.1-0.3 Gyr) starbursts occurring at $z > 2-3$. In addition, such precursors should also have a large clustering correlation length $r_0$ comparable to that of passive galaxies at lower redshifts ($z<2$) (e.g. \\cite{daddi00,firth02,kong,farrah06}), and compatible with that expected in the $\\Lambda$CDM models for galaxies located in massive dark matter halos and strongly biased environments. Examples of precursor candidates have been found amongst starburst galaxies selected at $z>2$ with a variety of techniques (e.g. $BzK$, \\cite{daddi04}; submm/mm, \\cite{chapman}; ``Distant Red Galaxies'', \\cite{franx}; optically-selected systems with high luminosity, \\cite{sha}; IRAC Extremely Red Objects, IEROs, \\cite{yan}; HyperEROs, \\cite{totani}, and ULIRGs at $z\\sim 1-3$ selected with {\\it Spitzer Space Telescope} \\cite{berta07}). Deep integral-field near-IR spectroscopy is being used to perform detailed studies of these precursor candidates in order to understand what are the main mechanisms capable to assemble massive galaxies with short timescales (e.g. \\cite{swin,nfs06,wright,law}). To date, the most detailed case is represented by a $BzK$--selected star-forming galaxy at $z=2.38$ which shows a massive rotating disk with high velocity dispersion which may become unstable and lead to the rapid formation of a massive spheroid (\\cite{genzel}).  In this paper, we exploit the combination of GMASS ultradeep spectroscopy, HST imaging and optical -- to -- mid-infrared photometry to study the global properties of a new sample of passive galaxies at $1.3<z<2.0$ in order to place new constraints on massive galaxy formation and evolution. This paper refers to other papers where more details can be found on the GMASS project (Kurk et al. 2007a in preparation), the large scale structure at $z=1.61$ (Kurk et al. 2007b in preparation), the photometric Spectral Energy Distribution (SED) fitting analysis (Pozzetti et al. 2007 in preparation) and the morphological analysis (Cassata et al. 2007 in preparation). We adopt $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm m}= 0.3$ and $\\Omega_{\\Lambda}=0.7$, give magnitudes in AB photometric system, and assume a Chabrier (2003) Initial Mass Function (IMF). ",
        "conclusions": "We presented a study of a 4.5$\\mu$m--selected sample of passive galaxies spectroscopically identified at $z\\approx$1.4--2 with the GMASS survey (Kurk et al. 2007a). This work benefits from ultradeep VLT+FORS2 optical spectroscopy complemented with multi-band photometry (0.4--8$\\mu$m) and HST imaging to investigate the physical, structural and evolutionary properties of these galaxies. The main results can be summarized as follows.  $\\bullet$ The GMASS passive galaxies have spectra and SEDs dominated by old stars and very weak or absent star formation. A comparison of the stacked rest-frame UV spectrum (equivalent to an integration time of $\\sim$500 hours) with three different libraries of stellar population model spectra indicates an age of $\\approx$0.7--2.8 Gyr for a metallicity range of 1.5--0.2 $Z_{\\odot}$. Extending the model fitting at longer wavelengths using near-infrared and IRAC photometry helps to reduce the age--metallicity degeneracy and indicates ages of $\\approx$1--1.6 Gyr, $Z=Z_{\\odot}$, $e$-folding timescales $\\tau \\sim$0.1--0.3 Gyr, where $SFR(t) \\propto exp(-t/ \\tau)$, and very low dust extinction. A comparison of the two libraries of model spectra which include TP-AGB stars shows that for a fixed age of 1 Gyr, the Maraston (2005) templates show a better agreement with the photometry than those of Charlot \\& Bruzual (2007) available to date.  $\\bullet$ Neither individual galaxies nor their stacked images are detected in the X-rays with \\textit{Chandra} data. This implies that a luminous AGN source ($L_X>10^{42}$ erg/s) is absent or is very heavily obscured. However, the possibility of dust obscuration seems unlikely because none of the galaxies has been detected at 24$\\mu$m, with only one marginal exception of one galaxy at $z=1.61$.  $\\bullet$ The stellar masses, estimated through the photometric SED fitting, are in the range of $10^{10-11}$ M$_{\\odot}$ and the specific star formation rates are very low ($\\lesssim 3 \\times 10^{-2}$ Gyr$^{-1}$). The stellar masses estimated with model spectra including TP-AGB stars are systematically lower by 0.1--0.2 dex than those estimated with models which do not include this phase of stellar evolution.  $\\bullet$ The HST+ACS morphological and surface brightness profile analysis indicate that the majority of the spectroscopically--selected passive galaxies have spheroidal morphologies consistent with being analogous to present-day ETGs. However, their sizes are smaller by a factor of $\\approx$2-3 than at $z\\approx0$, and imply that the stellar mass surface and volume internal densities are up to $\\approx$10 and $\\approx$ 30 times larger respectively. If literature data are added to the GMASS sample, we find that only a few passive systems at $1.2<z<2.5$ lie within the $z\\approx0$ size -- mass relation, whereas the majority has systematically much smaller sizes.  $\\bullet$ The Kormendy relation at $z\\approx1.5$ shows a large offset (2--3 mag) in effective surface brightness with respect to the local relation. This is difficult to explain with simple luminosity evolution models and requires that these high-$z$ compact and dense galaxies increase their size in order to reach the local Kormendy relation.  $\\bullet$ Samples of ETGs at lower redshifts ($0.7<z<1.2$) show that a larger fraction of passive galaxies follow the $z\\sim0$ size -- mass relation with respect to $z>1.3$. The ETGs at $z\\approx1$ which have the largest offsets with respect to the $z\\sim0$ size -- mass relation are the ones having the highest internal velocity dispersion, as expected from the ETG scaling relations. We find a hint that, for a fixed redshift $z\\approx$1, ETGs located within massive clusters are more preferentially located within the $z\\sim0$ size -- mass relation than ETGs located in lower density environments.  $\\bullet$ Superdense massive ETGs with $R_e \\approx$ 1 kpc are extremely rare at $z\\approx0$ with respect to $z>1$, and absent if $R_e<1$ kpc. However, it might be possible that compact and dense remnants are \"hidden\" inside present-day ETGs if the size of ETGs grew through mechanisms such as the \"smooth envelope accretion\".  $\\bullet$ Submillimeter--selected galaxies are the only systems at $z\\gtrsim$2 with sizes and mass surface densities (in gas) similar to those of the passive galaxies at $z\\approx$1--2. This suggests that a strong evolutionary link is present between these two galaxy populations.  $\\bullet$ It is currently unclear how the possible link between SMGs and compact passive galaxies fits within a more general framework which takes into account also the other galaxy populations so far identified at $1<z<3$. A plausible scenario could be outlined as follows and summarized in Fig. 23. {\\bf (1)} Massive star-forming galaxies selected in the optical/near-IR are gas--rich disky systems with long--lived star formation (e.g. $\\approx$0.5--1 Gyr, \\cite{daddi05b}). {\\bf (2)} These systems can become unstable (e.g. \\cite{genzel}) or participate in major merger events with other gas-rich systems. {\\bf (3)} In both cases a major starburst event is triggered, and this phase could correspond to the SMG stage characterized by short-lived ($\\approx$0.1 Gyr) vigorous starburst (\\cite{tacconi,tacconi2}). {\\bf (4)} The concomitant AGN provides enough feedback to ``quench'' the star formation in massive systems (Daddi et al. 2007a), and {\\bf (5)} compact, superdense, passively evolving remnants are formed, {\\bf (6)} and evolve subsequently by increasing gradually their sizes with mechanisms like major dry merging and/or envelope accretion more or less rapidly depending on their mass and environment. {\\bf (7)} The majority of most massive ETGs reach the assembly completion around $z\\approx$0.7, while lower mass ETGs continue to assemble down to lower redshifts (downsizing).  \\begin{figure} \\begin{center} \\includegraphics[width=\\columnwidth]{8739f23.ps} \\end{center} \\caption{A possible scenario for the formation and evolution of a massive spheroidal galaxy (see text).} \\end{figure}  $\\bullet$ Larger samples in the range of $1<z<3$ are needed to place stringent constraints on the mechanism(s) with which the sizes of high--$z$ passive galaxies grow as a function of cosmic time (e.g. dissipationless merging, envelope accretion, ...). High--resolution spectroscopy in the near--infrared with the next generation of telescopes (e.g. E-ELT, JWST) will also be crucial to study the internal dynamics of these systems. Further studies at $1<z<3$ will shed light on the global processes which lead to formation and evolution of massive galaxies."
    },
    "0801/0801.2382_arXiv.txt": {
        "abstract": "We present a new XMM-Newton spectrum of the Seyfert 2 nucleus of IC 2560, which hosts H$_{2}$O maser emission from an inclined Keplerian accretion disk. The X-ray spectrum shows soft excess due to multi-temperature ionized plasma, a hard continuum and strong emission features, from Mg, Si, S, Ca, Fe and Ni, mainly due to fluorescence. It is consistent with reflection of the continuum from a mostly neutral medium and obscuration due to a high column density, $>$ 10$^{24}$ cm$^{-2}$. The amplitude of the reflected component may exceed 10\\% of the central unobscured luminosity. This is higher than the reflected fraction, of a few percent, observed in other Seyfert 2 sources like NGC 4945. We observe an emission line at 6.7 keV, possibly due to FeXXV, undetected in previous Chandra observations. The absorption column density associated with this line is less than 10$^{23}$ cm$^{-2}$, lower than the obscuration of the central source. We hypothesize that this highly ionized Fe line emission originates in warm gas, also responsible for a scattered component of continuum emission that may dominate the spectrum between 1 and 3 keV. We compare X-ray and maser emission characteristics of IC 2560 and other AGN that exhibit water maser emission originating in disk structures around central engines. The temperature for the region of the disk associated with maser action is consistent with the expected 400-1000K range. The clumpiness of disk structures (inferred from the maser distribution) may depend on the unobscured luminosities of the central engines. ",
        "introduction": "IC2560 is a nearby spiral ($cz\\sim 2925$\\,km\\,s$^{-1}$) that is classified as a Seyfert 2 \\citep{fai86} and known to host 22\\,GHz H$_2$O maser emission originating in an accretion disk around the central engine \\citep{ishi01}.  Here we present analysis of a deep (archival) XMM-Newton EPIC observation of this active galactic nucleus (AGN).   Previous observations of IC 2560 in the X-ray band using ASCA GIS \\citep{ishi01} and Chandra ACIS-S \\citep{iwa02,mad06} both pointed to a reflection dominated spectrum and high obscuring column density.  IC 2560 is a relatively weak X-ray source, with an observed 2-10 keV flux of 3.3$\\times$10$^{-13}$ erg s$^{-1}$ cm$^{-2}$ that originates predominantly from a central source \\citep{mad06}. Interpretation of early GIS and ACIS-S observations was constrained by relatively limited sensitivity and energy resolution, in particular above 6 keV, where the bulk of the emission due to nuclear activity escapes into our line of sight. The EPIC instrument provides $\\sim$5$\\times$ larger collecting area at 6 keV and $\\sim$20\\% finer energy resolution (160 eV) than ACIS-S (chip 3). Although dominated by the central point source, flux detected with EPIC (1.1 kpc pixel$^{-1}$ at 1 keV) will include contribution from extended and off-nuclear emission components. As such, we note that the nuclear structure in IC 2560 appears to be relatively complex, as may be anticipated for a type-2 object. Hubble Space Telescope ACS/HRC observations at 3300 \\AA\\ detect patchy emission with multiple peaks within the central $\\sim$ 200 pc \\citep{mun07}, presumably due to dust. Relatedly, population synthesis models fit to optical spectra, obtained with the ESO 1.5m (3470-5450 \\AA) for the central $\\sim$ 300 pc, are suggestive of active star formation during the last 1.5-10 Myr \\citep{fer04}. Ground-based spectroscopic study has also detected asymmetry in the [OIII] line profile (a blue excess) that may be indicative of outflow at $\\sim$ 100 km s$^{-1}$ projected on the sky plane in a 300$\\times$400 pc central patch \\citep{schu03}. Because the inclination of the galaxy is 63$^\\circ$, and the inclination of the central engine (as inferred from  maser emission) may be even closer to edge-on, the actual flow speed may be larger.  The observed H$_2$O maser spectrum exhibits emission close to the systemic velocity as well as blue and red-shifted complexes, offset, approximately symmetrically, by $\\pm$ 200 - 420 km s$^{-1}$ \\citep{bra03}. This pattern is an indicator of emission from a rotating edge-on disk structure (e.g., NGC 4258; \\cite{nak93}). \\cite{ishi01} reported line-of-sight accelerations for the ``systemic emission'' and mapped its emission distribution using Very Long Baseline Interferometry (VLBI). Based on the assumption that the emission traced a thin, rotating, circular disk that is observed edge-on, they used the velocities of high-velocity emission to estimate disk radius and dynamical mass. Follow-up VLBI observation by Greenhill \\etal (2008, in preparation) enabled mapping of both systemic and high-velocity emission, thus tracing the disk structure, demonstrating Keplerian rotation of material around the central engine in a relatively thin distribution, and enabling more certain modeling of the accretion disk geometry and estimation of the central mass. The observed molecular gas lies at radii 0.08 - 0.27 pc and the inferred enclosed mass is 2.9$\\times$10$^6$M$_\\odot$ with $\\sim$ 20\\% accuracy. The thinness of disk (h/r $\\ll$1) as well as presence of H$_2$O emission suggest disk material that is relatively cool ($\\la$10$^3$ K).  Due to generally high extinction toward Seyfert-2 central engines, X-ray and radio observations together are especially useful in the study of the physical processes that occur at radii $\\la$1 pc. Measurement of X-ray spectra enable estimation of intrinsic luminosity, obscuring column density, temperature and abundance. Time monitoring facilitates estimation of size scales for high-energy phenomenon. VLBI observations of H$_2$O maser emission enable mapping of accretion disk geometry and orientation to the line of sight, and estimation of orbital radii for molecular material, central engine mass and Eddington luminosity (\\cite{gre07} and references therein). The presence of maser emission implies the existence of a reservoir of relatively cool material and at the same time, a heating mechanism to maintain population inversion. The necessary energy may be  imparted by viscous heating within the disk (e.g., \\citep{des98}) or external irradiation \\citep{neu94,wat94}. However, in either case, survival of molecular gas demands a shielding column from UV and X-ray emission generated by, e.g., the central engine and coronae. This column may arise in outflowing material, disk surface layers, or disk material at small radii, thus providing additional quantitative constraints on overall structure around the central engine and energetics. In section 2, we discuss the data reduction of the EPIC data, while section 3 focuses on analysis of the source spectrum. Estimation of the intrinsic emission of the central engine (i.e., corrected for absorption along line of sight) is discussed in section 4. X-ray properties of IC 2560 and other extragalactic H$_2$O maser sources believed to originate in accretion disks are considered in section 5; there we put IC 2560 into context and we discuss possible trends that may indicate connection between characteristics estimated from VLBI observations of masers and X-ray observations of their host AGN. Throughout the paper we adopt H$_{0}$=70 km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "A deep 70 ks XMM-Newton spectrum for IC 2560 shows evidence for a warm scattered continuum component as well as number of emission lines not detected in earlier Chandra observations. These include K$\\alpha$ emission from Mg (1.25 keV), Ca (3.6 keV) and Ni (7.47 keV) as well as Fe K$\\beta$ (7.05 keV) and K$\\alpha$ line emission from Fe XXV (6.7 keV). We do not detect line emission from species at intermediate ionization states, between the nearly neutral material that gives rise to the K$\\alpha$ emission, and the warm scattering medium with FeXXV. The direct X-ray emission from the central source appears to be completely obscured up to 10 keV; the inferred column density is in excess of a few times 10$^{24}$ cm$^{-2}$. This is in agreement with a 25\\% upper limit on the magnitude of the Compton shoulder. The reflecting medium is mostly neutral, optically thick and could also provide the obscuring column toward the central engine.  The 6.7 keV line is associated with the optically thin, mostly ionized, scattering medium,  with significantly lower density. The dominance of the resonance line at 6.7 keV indicates the density has to be lower than 10$^{23}$ cm$^{-2}$, while the estimated fraction of intrinsic emission scattered along line-of-sight indicates a lower limit of 10$^{22}$ cm$^{-2}$. The scattering medium can plausibly exist on the scale of a few parsec, but the exact geometry is unknown. Limits on density of the scatterer indicate a covering fraction between 0.1 and 1. Using a small sample of H$_{2}$O disk maser hosts, selected using physically motivated criteria, we find that the unobscured luminosity of the central source may shape molecular disk structure in this sample. The inner radius as well as degree of substructure in the disk may be related to the unobscured X-ray emission, though a larger sample is required to confirm this trend."
    },
    "0801/0801.1662_arXiv.txt": {
        "abstract": "We construct a little Higgs model with the most minimal extension of the standard model gauge group by an extra $U(1)$ gauge symmetry. For specific charge assignments of scalars, an approximate $U(3)$ global symmetry appears in the cutoff-squared scalar mass terms generated from gauge bosons at one-loop level. Hence, the Higgs boson, identified as a pseudo-Goldstone boson of the broken global symmetry, has its mass radiatively protected up to scales of 5-10~TeV. In this model, a $Z_2$ symmetry, ensuring the two $U(1)$ gauge groups to be identical, also makes the extra massive neutral gauge boson stable and a viable dark matter candidate with a promising prospect of direct detection. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.3667_arXiv.txt": {
        "abstract": "Through the modelling of the Spectral Energy Distribution of blazars we can infer the physical parameters required to originate the flux we see. Then we can estimate the power of blazar jets in the form of matter and fields. These estimate are rather robust for all classes of blazars, although they are in part dependent of the chosen model (i.e. leptonic rather than adronic). The indication is that, in almost all cases, the carried Poynting flux is not dominant, while protons should carry most of the power. In emission line blazars the jet has a comparable, and often larger, power that the luminosity of the accretion disk. This is even more true for line--less BL Lacs. If the jet is structured at the sub--pc scale, with a fast spine surrounded by a slower layer, then one component sees the radiation of the other boosted, and this interplay enhances the Inverse Compton flux of both. Since the layer emission is less beamed, it can be seen also at large viewing angles, making radio--galaxies very interesting GLAST candidates. Such structures need not be stable components, and can form and disappear rapidly. Ultrafast TeV variability is challenging all existing models, suggesting that at least parts of the jets are moving with large bulk Lorentz factors and at extremely small viewing angles. However, these fast ``bullets\" are not necessarily challenging our main ideas about the energetics and the composition of the bulk of the jet. ",
        "introduction": "EGRET, onboard {\\it CGRO}, and the ground based Cherenkov telescopes showed that blazars are the most powerful high energy extragalactic emitters, and allowed to know the bolometric luminosity of these objects. Now, just after the launch of {\\it AGILE} and just before {\\it GLAST}, we are preparing for (and already tasting) the possibility to have simultaneous data in the optical, X--rays and the GeV bands (and possibly the TeV one). This will be possible mainly by {\\it Swift}: its rapid slew and flexible scheduling will ensure good quality optical and X--ray data while the $\\gamma$--ray observations are still ongoing. What was an exception in the EGRET era, will be routine. This is therefore a risky and at the same time healthy time to put forward new ideas concerning blazars, that {\\it GLAST} can falsify. In this contribution, I will try to summarize how the power of jets can be derived, and what inferences can we draw from that. I will discuss the ultra--fast TeV variability recently observed in PKS 2155--304, arguing that, contrary to previous claims, it is unlikely that the jet of this source is magnetically dominated. Furthermore, I will point out that the fact that the TeV emitter BL Lac objects are also the least powerful blazars opens up the possibility to slow down their jets by the Compton rocket effect.   \\begin{figure} \\vskip -0.2 true cm \\hskip -0.5 true cm \\centerline{\\psfig{figure=ghisellini_f1.ps,width=11truecm}} \\vskip -0.5 true cm \\caption{SEDs of 3C 454.3 at different epochs. Top panel: the SED in 2000, corresponding to the {\\it Beppo}SAX observations discussed in \\cite{tav02}. % The other points are not simultaneous (see \\cite{pian06} % and references therein). Mid panel: the SED during the huge optical flare in 2005, as described in \\cite{pian06, giommi06}.  % Bottom panel: the SED on July 26, 2007, as observed by {\\it AGILE}\\cite{agile07}, and {\\it Swift}. The optical flux in the $R$ band comes from the Tuorla observatory. The optical and X--ray fluxes are corrected for Galactic extinction ($A_V=0.355$). The solid and dashed lines correspond to our modelling. The dotted line is the contribution from the accretion disk (assumed to be a simple black--body). From \\cite{gg07}. % } \\label{454} \\end{figure} ",
        "conclusions": ""
    },
    "0801/0801.4876_arXiv.txt": {
        "abstract": "Multi-frequency (4.6, 5, 5.5, 8, 8.8, 13, 15, 22 \\& 43 GHz) polarization observations of 6 ``blazars'' were obtained on the American Very Long Baseline Array (VLBA) over a 24-hr period on 2 July 2006. Observing at several frequencies, separated by short and long intervals, enabled reliable determination of the distribution of Faraday Rotation on a range of scales. In all cases the magnitude of the RM increases in the higher frequency observations, implying that the electron density and/or magnetic field strength is increasing as we get closer to the central engine. After correcting for Faraday rotation, the polarization orientation in the jet is either parallel or perpendicular to the jet direction. A transverse Rotation Measure (RM) gradient was detected in the jet of 0954+658, providing evidence for the presence of a helical magnetic field surrounding the jet. For three of the sources (0954+658, 1418+546, 2200+420), the sign of the RM in the core region changes in different frequency-intervals, indicating that the line-of-sight component of the magnetic field is changing with distance from the base of the jet. We suggest an explanation for this in terms of bends in a relativistic jet surrounded by a helical magnetic field; where there is no clear evidence for pc-scale bends, the same effect can be explained by an accelerating/decelerating jet. ",
        "introduction": "AGN jets emit synchrotron radiation that can be detected at all radio frequencies. This radiation is often highly linearly polarized, which provides important information on the degree of order and orientation of the magnetic field in these jets. The type of AGN studied in this experiment are known as ``blazars'', which have jets pointed close to our line of sight (LoS) and often exhibit strong variability in total flux and linear polarization over a broad range of frequencies from $\\gamma$-ray to radio.  We observed 6 sources with the American Very Long Baseline Array (VLBA) at 8 frequencies from 4.6 GHz to 43 GHz over a 24-hr period on 2 July 2006. This radio interferometer provides milliarcsecond resolution which corresponds to the parsec scale structure of these jets. Even though these AGN are intrinsically two-sided, pc-scale observations typically show a one sided ``core-jet'' structure due to Doppler boosting of the radiation from the relativistic jet pointed towards us. The jet moving away from us is usually too faint to be detected.  The radiation detected from the jet is generally optically thin while the core region usually displays a flat spectrum attributed to synchrotron self-absorption. In the optically thin regime the polarization orientation is perpendicular to the magnetic field, while in the optically thick/flat spectrum region the polarization is parallel to the magnetic field.  Recent MHD simulations have provided an almost complete explanation of how these jets are launched, accelerated and collimated close to the black hole, see \\cite{MeierJapan} and references therein. The global magnetic field structure expected on these scales is helical. It is possible that this structure changes when the flow gets disrupted by shocks after a few hundred Schwarzschild radii, but observational evidence of helical fields on scales larger than this \\cite{Asada2002, GabuzdaMurray2004, Mahmud2008} suggests that remnants of the earlier magnetic field structure remain or that a current driven helical kink instability is generated \\cite{NakamuraJapan, Carey2008}. ",
        "conclusions": "Several phenomenon, for example, the polarization structure \\cite{Lyutikov2005}, RM gradients \\cite{Asada2002} and circular polarization generation \\cite{GabCP2007}, can be understood by considering a helical magnetic field geometry for pc-scale AGN jets. In this paper, we consider how regions with different RM signs can also be explained within a helical magnetic field model, as places where the jet is observed at angles greater than or less than $1/\\Gamma$ due to bends in the jet or due to an accelerating/decelerating straight jet. (A longitudinal jet magnetic field with a change in the angle to the LoS could also cause a RM sign change, but this does not correspond to the observed magnetic field in 0954+658 or 2200+420.) It's important to note that VLBI resolution is usually not sufficient to completely resolve the true optically thick core, therefore, the VLBI ``core'' consists of emission from the true core and some of the optically thin inner jet. So if bends occur on scales smaller than the observed VLBI ``core'', core RMs with different signs can be derived from observations at different frequencies (ie. probing different scales of the inner jet).  In our future work, we will attempt to reconstruct the 3-D path of the jet through space using the combined information from the observed distributions of the total intensity, linear polarization, spectral index and rotation measure."
    },
    "0801/0801.1986_arXiv.txt": {
        "abstract": "The year 2007 has furnished us with outstanding results about the origin of the most energetic cosmic rays: a flux suppression as expected from the GZK-effect has been observed in the data of the HiRes and Auger experiments and correlations between the positions of nearby AGN and the arrival directions of trans-GZK events have been observed by the Pierre Auger Observatory.  The latter finding marks the beginning of ultra high-energy cosmic ray astronomy and is considered a major breakthrough starting to shed first light onto the sources of the most extreme particles in nature.  This report summarizes those observations and includes other major advances of the field, mostly presented at the 30$^{\\rm th}$ International Cosmic Ray Conference held in M\\'erida, Mexico, in July 2007. With increasing statistics becoming available from current and even terminated experiments, systematic differences amongst different experiments and techniques can be studied in detail which is hoped to improve our understanding of experimental techniques and their limitations. ",
        "introduction": "Understanding the origin of the highest energy cosmic rays is one of the most pressing questions of astroparticle physics.  Cosmic rays with energies exceeding $10^{20}$ eV have been observed for more than 40 years (see e.g.\\ \\cite{Nagano-Watson}) but due to their low flux only some ten events of such high energies could be detected up to recently.  There are no generally accepted source candidates known to be able to produce particles of such extreme energies.  Moreover, there should be a steeping in the energy spectrum near $10^{20}$ eV due to the interaction of cosmic rays with the microwave background radiation (CMB).  This Greisen-Zatsepin-Kuzmin (GZK) effect \\cite{GZK} severely limits the horizon from which particles in excess of $\\sim 6\\cdot10^{19}$ eV can be observed.  For example, the sources of protons observed with $E\\ge 10^{20}$ eV need to be within a distance of less than 50 Mpc \\cite{Harari-06}.  The non-observation of the GZK-effect in the data of the AGASA experiment \\cite{Takeda-03} has motivated an enormous number of theoretical and phenomenological models trying to explain the absence of the GZK-effect and has stimulated the field as a whole.  Only this year, with the final analysis of the HiRes-data \\cite{Abbasi-07a} and the advent of high-statistics and high quality hybrid data from the Pierre Auger Observatory (PAO) \\cite{Yamamoto-07}, the situation has changed considerably: a suppression such as expected from the GZK-effect is now observed with high statistical significance.  The very recent breaking news about the observation of directional correlations of the most energetic Pierre Auger events with the positions of nearby AGN \\cite{Abraham-07} complements the observation of the GZK effect very nicely and provides evidence for an astrophysical origin of the most energetic cosmic rays.  Another key observable allowing to discriminate different models about the origin of high-energy cosmic rays is given by the mass composition of cosmic rays. Unfortunately, such measurements are much more difficult due to their strong dependence on hadronic interaction models.  Only primary photons can be discriminated safely from protons and nuclei and recent upper limits to their flux largely rule out top-down models, originally invented to explain the apparent absence of the GZK-effect in AGASA data.  In this article, prepared for the TAUP conference in Sendai (Japan), we describe the status of each of these topics, as reported during the recent International Cosmic Ray Conference held in M\\'erida, Mexico, in July 2007 (ICRC2007) and in publications becoming available since. ",
        "conclusions": "Remarkable progress has been made in cosmic ray physics at the highest energies, particularly by the start-up of the (still incomplete) Pierre Auger Observatory.  The event statistics above $10^{19}$\\,eV available by now allows detailed comparisons between experiments and indicates relative shifts of their energy scales by $\\pm 25$\\,\\%.  Given the experimental and theoretical difficulties in measuring and simulating extensive air showers at these extreme energies, this may be considered a great success. On the other hand, knowing about overall mismatches of the energy scales between experiments may tell us something.  Clearly, in case of fluorescence detectors better measurements of the spectral and absolute fluorescence yields and their dependence on atmospheric parameters are needed and will hopefully become available in the very near future \\cite{fluorescence}.  This should furnish all fluorescence experiments with a common set of data.  Differences in the calibration between surface detectors and fluorescence telescopes, best probed by hybrid experiments like Auger and TA, may then be used to test the modelling of EAS. The muon component at ground, known to be very sensitive to hadronic interactions at high energies \\cite{Drescher-04}, could in this way serve to improve hadronic interaction models in an energy range not accessible at man-made accelerators.  In fact, several studies (e.g.\\ \\cite{Engel-07}) indicate a deficit of muons by 30\\,\\% or more in interaction models like QGSJET.  The energy scale is of great importance also for the AGN correlation discussed in the previous section. As shown in \\cite{AGN-long-07}, the correlation sets in abruptly at an (Auger) threshold energy of about 57 EeV. Already a downshift in energy by 17\\,\\% (the mismatch between Auger and HiRes) would weaken the signal by more than 3 orders of magnitude to make it basically disappear.  Thus, verification of the correlation signal by HiRes or AGASA would need to be done for a threshold energy (on their scale) of 67 EeV and 85 EeV, respectively.  In this energy range, HiRes observes a spectral slope of $\\gamma=5.1\\pm0.7$ \\cite{Bergman-07b}, i.e.\\ the number of events available for a correlation analysis would, according to Table \\ref{tab:expts}, drop to about 12 (HR-I) and 4 (HR-II) when taking the rise of the apertures into account.  This would amount to about half the statistics of Auger, well in agreement with the quoted exposures. Unfortunately, the angular resolution of monocular reconstruction is by far too poor for such a test.  Only stereo data could provide the required angular resolution. However, in this case the expected statistics of about 5 events above threshold (based upon the numbers and exposures given above) appears too small for any verification.  In fact, the distance parameter of the correlation of 71 Mpc may indicate a mismatch of the energy scale: For protons above 57 EeV the GZK horizon would be 200 Mpc \\cite{Harari-06} but already for 20\\,\\% higher energy it would shrink by more than a factor of two to become consistent to the correlation parameter.  Another puzzling feature is the observed small deflection of particles which suggests dominantly protons as primaries.  Note that 90\\,\\% of the events (20/22) off the galactic plane are correlated to within $\\sim 3^\\circ$ which AGN positions which is very unlikely for heavy nuclei.  On the other hand, the elongation curves in Fig.\\,\\ref{fig:Xmax} suggests an admixture of heavy nuclei by more than 10\\,\\%.  This may be related again to imperfections of the hadronic interaction models used for comparison in Fig.\\,\\ref{fig:Xmax}.  All of this tells us that the near future will be highly exciting: The question of the energy scales will soon be settled and more detailed comparisons between experiments will become possible. The shape of the energy spectrum in the GZK region will tell us about the source evolution, the composition in the ankle region will answer the question about the G-EG transition, observations of cosmogenic photons and neutrinos are in reach and in case of neutrinos will probe the GZK effect over larger volumes, the correlations will be done with better statistics, with improved search techniques and with more appropriate source catalogues and source selection parameters to tell us about source densities, and the true sources of EHECRs.  Very important to note is that different pieces of information start to mesh and are being accessed from different observational techniques and can be cross-checked: {\\em The big picture is being painted!}  Given the scientific importance of this, it would be a mistake to have only one observatory - even when operated as a hybrid detector - taking data.  The TA project and its extensions will be very important particularly in the sub-GZK range but, unfortunately, will be too small to collect sufficient statistics at the highest energies.  Auger-North will be imperative here and needs immediate vigorous support.  The next generation experiment JEM EUSO to be mounted at the Exposed Facility of Japanese Experiment Module JEM EF will potentially reach much larger exposures but still faces many experimental challenges to be addressed.  \\vspace*{-5mm} \\subsection*{Acknowledgement} I would like to thank the organizers of the excellent TAUP meeting in Sendai for the invitation to give a review talk at this exciting time. Also, its a pleasure to thank many of my colleagues for stimulating discussions. The German Ministry for Research and Education (BMBF) and the Deutsche Forschungsgemeinschaft (DFG) are gratefully acknowledged for financial support."
    },
    "0801/0801.3818_arXiv.txt": {
        "abstract": "Simultaneous multiwavelength studies of X-ray binaries have been remarkably successful and resulted in improved physical constraints, a new understanding of the dependence of mass accretion rate on X-ray state, as well as insights on the time-dependent relationship between disk structure and mass-transfer rate.  I will give some examples of the tremendous gains we have obtained in our understanding of XRBs by using multiwavelength observations.  I will end with an appeal that while Spitzer cryogens are still available a special effort be put forth to obtaining coordinated observations with emphasis on the mid-infrared: Whereas the optical and near-IR originate as superpositions of the secondary star and of accretion processes, the mid-IR crucially detects jet synchrotron emission from NSs that is virtually immeasurable at other wavelengths. A further benefit of Spitzer observations is that mid-infrared wavelengths can easily penetrate regions that are heavily obscured.  Many X-ray binaries lie in the Galactic plane and as such are often heavily obscured in the optical by interstellar extinction. The infrared component of the SED, vital to the study of jets and dust, can be provided {\\it only} by Spitzer; in the X-rays we currently have an unprecedented six satellites available and in the optical and radio dozens of ground-based facilities to complement the Spitzer observations. ",
        "introduction": "X-ray binaries (XRBs) owe their prominence to one of the most efficient energy release mechanisms known: accretion onto a compact object. The energy released by accretion can be spread over essentially the entire electromagnetic spectrum. Since each part of the spectrum provides distinct and often time-variable information (see, e.g., Fig. 1), attempts to understand these systems by concentrating on information from a limited wavelength range may lead to contradictory and misleading conclusions (see, e.g., Figs. 2-4). It is important to study them simultaneously over a broad range of the electromagnetic spectrum; spectral energy distributions (SEDS) are particularly useful because they clearly reveal multiple emission components and reflect the physics and geometry of the emitting regions.  Multiwavelength studies have been remarkably successful and resulted in improved physical constraints on the systems, a new understanding of the dependence of mass accretion rate on X-ray state, as well as insights on the time-dependent relationship between disk structure and mass-transfer rate, (e.g., Vrtilek~\\etal~1990,1991; Hasinger~\\etal~1990; McClintock~\\etal~2001; Fuchs~\\etal~2003; Homan~\\etal~2005)  One of the fundamental unsolved problems of accretion physics is how the outgoing power is distributed between electromagnetic radiation and mechanical power (relativistic jets), and what determines a switch between those two output channels. Relativistic jets are amongst the most energetic phenomena in the Universe. In our Galaxy, jets have been detected in young stellar objects, massive binaries, symbiotic stars, BH and NS XRBs, super-soft X-ray sources, and planetary nebulae. They are also believed to be responsible for $\\gamma$-ray bursts (Fender~\\etal~2004). Despite the tremendous range of scale, jets in stellar mass systems share physical properties with those in super-massive BHS. The most significant property shared is that {\\it all systems that produce jets also have accretion disks} (Livio 2002).  This implies that jet energy is obtained through accretion power and that they may play a major role in the transport of angular momentum of the inflalling gas.  As spatial separation is generally impossible owing to limitation in resolution, {\\it multiwavelength studies are essential to separate the stellar, wind, disk, and jet components. The broad band spectrum is also necessary to test theoretical models for jet formation and to constrain fundamental parameters of jet physics}.  \\begin{figure} \\centering { \\scalebox{0.22}{\\includegraphics{vrtilek_s_fig1.ps}} } { \\scalebox{0.30}{\\includegraphics{vrtilek_s_fig2.ps}} } \\caption{ LEFT: Broadband energy spectra of the black-hole binary GRO~J1655-40 illustrating variability within one source. Each color represents simultaneous observations. (Migliari~\\etal~2007; see also Tomsick contribution to these proceedings). RIGHT: Broadband energy spectra of a selection of black hole and neutron star binaries showing variability between sources (ATCA and RXTE values are from a survey conducted by the Disk/Jet Consortium in June of 2006; values for Sco X-1 are from Wachter~\\etal~2006). } \\label{fig1sub} \\vspace{-3mm} \\end{figure}    \\begin{figure} \\centering { \\scalebox{0.3}{\\includegraphics{vrtilek_s_fig3.ps}} } { \\scalebox{0.25}{\\includegraphics{vrtilek_s_fig4.ps}} } \\caption{ LEFT: RXTE ASM countrates (courtesy of the RXTE ASM team) of the black hole candidate Cygnus X-1 with typical countrates for ``Low\", ``Intermediate\", and ``High\" states indicated. RIGHT: Broadband energy spectrum of Cygnus X-1 (Gierlin'ski~\\etal~1993). It is clear that the designations ``High\" and ``Low\" states are based on a narrow energy range (2-10 keV) and are not valid when the full X-ray band is available. } \\label{fig2sub} \\vspace{-3mm} \\end{figure}  \\begin{figure} \\centering { \\scalebox{0.23}{\\includegraphics{vrtilek_s_fig5.ps}} } { \\scalebox{0.32}{\\includegraphics{vrtilek_s_fig6.ps}} } \\caption{ LEFT: RXTE ASM countrates (courtesy of the RXTE ASM team) of the pulsing XRB Her X-1. Arrows mark the dates of the SEDs plotted on the right. RIGHT: Simultaneous SEDs of Her X-1 clearly show that the optical and ultraviolet flux (left side) remain largely unaffected at times when the extreme ultraviolet to hard X-ray flux (right side) takes a dramatic plunge.  An indication that X-ray flux is not a dependable measure of mass accretion rate. Vrtilek~\\etal~2001). } \\label{fig2sub} \\vspace{-3mm} \\end{figure}  \\begin{figure} \\centering { \\scalebox{0.18}{\\includegraphics{vrtilek_s_fig7.ps}} } \\hspace{1cm} { \\scalebox{0.35}{\\includegraphics{vrtilek_s_fig8.ps}} }  \\caption{ LEFT: Hardness vs. intensity data from RXTE observations of the Galactic Z-source Sco X-1 (Barnard, Kolb, \\& Osborne 2003). RIGHT: UV continuum vs line flux in 1224-1986$\\AA$ band of Sco X-1 verifying that mass accretion rate increases from HB to NB to FB and does not correlate with X-ray flux. (Vrtilek~\\etal~1991).} \\label{cygscocyg} \\end{figure} ",
        "conclusions": ""
    },
    "0801/0801.1506_arXiv.txt": {
        "abstract": "The scaling relations between rotation velocity, size and luminosity form a benchmark  test for any theory of  disk galaxy formation.  We confront  recent theoretical models  of disk  formation to  a recent large  compilation  of  such   scaling  relations.   We  stress  the importance  of  achieving  a  fair  comparison  between  models  and observations. ",
        "introduction": "Understanding the origin  and nature of galaxy scaling  relations is a fundamental quest  of any successful theory of  galaxy formation.  The success  of a  particular  theory will  be  judged by  its ability  to predict  the  slope, scatter,  and  zero-point  of  any robust  galaxy scaling relation at any  particular wavelength.  The scaling relations between rotation velocity, $V$, size, $R$, and luminosity, $L$, are of special interest as they are linked via the virial theorem.  Courteau \\etal  (2007; hereafter C07)  compiled a sample of  1303 disk galaxies  for which accurate  rotational velocities  and near-infrared sizes and luminosities are available.   We refer the reader to C07 for details about this compilation.  The $VLR$ scaling relations for these galaxies are  shown in  Fig.~1.  The solid  black lines in  each panel show orthogonal linear fits to the combined data set, while the dashed lines show the  $2\\sigma$ scatter in these relations.   The points are color-coded according to central surface brightness.  This reveals the fundamental independence  of surface brightness of  the $VL$ relation. As discussed  in C07,  the $VL$ relation  is also independent  of disk size  and  residuals  from  the  size-luminosity  relation.   This  is unexpected, as the simplest models of galaxies embedded in dark matter halos  predict a  strong  surface brightness  dependence, unless  disk galaxies  are dominated by  dark matter  within their  optical regions (Courteau \\& Rix 1999; Dutton \\etal 2007). ",
        "conclusions": ""
    },
    "0801/0801.4924_arXiv.txt": {
        "abstract": "We review recent 3D cosmological hydrodynamic simulations of primordial star formation from cosmological initial conditions (Pop III.1) and from initial conditions that have been altered by radiative feedback from stellar sources (Pop III.2). We concentrate on simulations that resolve the formation of the gravitationally unstable cloud cores in mini-halos over the mass range $10^5 < M/\\Msun < 10^7 $ and follow their evolution to densities of at least $10^{10} \\cmm3$ and length scales of $<10^{-2}$ pc such that accretion rates can be estimated. The advent of ensembles of such simulations exploring a variety of conditions permits us to assess the robustness of the standard model for Pop III.1 star formation and investigate scatter in their formation redshifts and accretion rates, thereby providing much needed information about the Pop III IMF. The simulations confirm the prediction that Pop III.1 stars were massive ($\\sim 100 \\Msun$), and form in isolation in primordial mini-halos. Simulations of Pop III.2 star forming in relic HII regions suggest somewhat lower masses ($\\sim 30 \\Msun$) which may help explain the chemical abundances of extremely metal poor stars. We note that no 3D simulation at present has achieved stellar density let alone followed the entire accretion history of the star in any scenario, and thus the IMF of Pop III stars remains poorly determined theoretically. ",
        "introduction": "Dark matter mini-halos in the mass range $10^5-10^6$ solar mass virializing at high redshift are believed to collect enough primordial gas within them to host the formation of the first generation of stars (Population III). Mediated by molecular hydrogen cooling, gas in the centers of these halos condenses and becomes unstable to gravitational fragmentation with a typical mass scale of a few hundred solar masses. Pioneering high resolution 3D hydrodynamic cosmological simulations by Abel et al. (2002) and Bromm et al. (2002), and more recently by Yoshida et al. (2006) and O'Shea \\& Norman (2007a) have shown that gravitational fragmentation produces only one massive fragment per halo. If this gas does not fragment further as it approaches stellar density, Pop III stars with this typical mass would be produced. The fates and feedback effects of such stars has many interesting consequences for the early structure formation as discussed elsewhere in this volume (see review by Ciardi).  While this picture is now firmly established, based as it is on rather simple physics elucidated by the simulations, it is important to realize the remaining uncertainties. First, as of 2003 when the first reviews of primordial star formation were being written (Barkana \\& Loeb 2001, Bromm \\& Larson 2004, Glover 2005), a rather small number of such simulations had been done, raising the question of how ubiquitous this mode of star formation was in the early universe. Second, despite the large range of scales covered by the simulations, they were necessarily terminated well before stellar density was reached due to missing physics. The possibility of sub-fragmentation could not be ruled out, with consequent uncertainty on the mass scale of the first stars. Third, as emphasized by Barkana \\& Loeb (2004), for technical reasons the simulations were carried out in quite small cosmological volumes (< 1 Mpc comoving), the result being that the redshift of formation of the first stars found ($z\\sim 20$) was underestimated. Larger volumes would contain rarer peaks in the density field, and these would presumably form Pop III stars earlier (White \\& Springel 2000). Could the higher densities and temperatures of this earlier epoch alter the mass scale of the truly first stars? And finally, there was a host of complicated feedback effects--radiative, kinetic, and chemical--which were not included in these first simulations which could ``mess up\" the simple picture by altering the initial conditions and the cooling properties of the collapsing cloud cores.  The most concerning of these feedback effects was the build-up of an FUV background by emissions from the first stars that could photo-dissociate the hydrogen molecules which mediate Pop III star formation in the first place (Haiman, Rees \\& Loeb 1997, Ciardi, Ferrara \\& Abel 2000, Haiman, Abel \\& Rees 2001). The first simulations examining this ``negative feedback effect\" (Machacek, Bryan \\& Abel 2001, hereafter MBA; and Yoshida et al 2003, hereafter Y03) concluded that Pop III star formation would not be suppressed, but merely delayed. The concept of a critical halo mass capable of forming a Pop III star was introduced by MBA. They found that the critical halo mass is an increasing function of the FUV background mean intensity in the Lyman-Werner bands (11.2-13.6 eV) suggesting that Pop III star formation would become self-limiting. Y03 found that gas cooling becomes inefficient above $J_{21}=0.01$ due to the photo-destruction of \\h2. Here $J_{21}$ is the mean intensity of the FUV background in units of $10^{-21}$ ergs/cm$^2$/sec/ster/Hz. Using this information Y03 constructed a semi-analytic model of the cosmic history of Pop III star formation. They found a rapidly rising population of Pop III stars over the redshift interval 35 > z > 20 ``the rise of Pop III\", which begins to become self-regulated at $z \\sim 25$ when the FUV background attains $J_{21}=0.01$. They did not consider the lower redshift evolution as this was out of the range of their simulations, but rather speculated that Pop II star formation dominated the cosmic star formation history below $z \\sim 20$.  It is important to point out that neither the MBA simulations nor the Y03 simulations were evolved below redshifts of $z \\sim 20$ or had the spatial resolution to follow the chemo-thermal-hydrodynamics of cloud collapse on scales below $\\sim 0.1$ pc, making ``the fall of Pop III\" by radiative feedbacks alone quite uncertain and understudied. Indeed, Y03 found the Pop III global SFR was still rising at the end of their simulations, implying that we don't even know when the Pop III epoch peaked.  Ricotti, Gnedin \\& Shull (2002) addressed this question using hydrodynamic cosmological simulations that included the feedback of both photo-dissociating and photo-ionizing radiation from primeval galaxies. Although their simulations were of even lower resolution than MBA and Y03 and parameterized star formation with ad hoc prescriptions, they found a cosmic star formation history quite similar to that predicted by Y03 despite small box size effects. They too could not simulate sufficiently large boxes to obtain converged results in the interesting redshift regime 20 > z > 6, but rather found that rare luminous objects strongly affected the average SFR within the simulation volume. It is fair to say that we have a much better idea when the first Pop III star formed in the universe than when the last one formed. Indeed, some models have been presented which predict Pop III stars forming as low as redshift 3 (Schneider et al. 2006).  This review updates progress on direct numerical simulations of Pop III star formation by several groups since the reviews of Bromm and Larson (2004), Glover (2005), and Ciardi and Ferrara (2005). New high resolution simulations have been carried out by Yoshida et al. (2006), O'Shea \\& Norman (2007a) and Gao et al. (2007) which test the robustness of the basic picture described above by pushing the collapse to higher central densities, although not yet stellar densities, and by considering ensembles of simulations. The effects of a FUV background on Pop III star formation has recently been revisited by O'Shea \\& Norman (2007b) and by Wise and Abel (2007). It is found that \\h2 cooling remains important in halos with masses approaching $10^8 \\Msun$ and FUV backgrounds as strong as $J_{21}=1$. This suggests that the Pop III epoch may be more extended than previously thought and therefore occurring in rather different environments than originally simulated.  Pop III star formation has been studied in new and different environments as well. Pop III star formation in warm dark matter models has been studied by O'Shea \\& Norman (2006) and by Gao \\& Theuns (2007). Finally, Pop III star formation in gas that has been pre-processed by earlier generation of Pop III stars has been simulated by O'Shea et al. (2005), Mesinger, Bryan \\& Haiman (2006), and by Yoshida et al. (2007a,b). The surprising result is that despite the diversity of the environments studied, it is found that Pop III.2 stars form in basically the same way as Pop III.1 stars and that they are massive, although they appear to be somewhat less massive than Pop III.1 based on lower cloud core temperatures and accretion rates. However, the new simulations show that environment and formation redshift/history can have a substantial effect on protostellar accretions rates, suggesting some spread in the primordial IMF.  No 3D simulation has yet been done that follows the entire accretion history of the star, and therefore final masses are still uncertain. However, approximate integrations have been done assuming spherical symmetry which suggest that the entire protostellar envelope can be accreted (Omukai \\& Palla 2003, Yoshida et al. 2006, 2007a,b). Verifying this result with fully 3D simulations represents a grand challenge for the field. ",
        "conclusions": ""
    },
    "0801/0801.0670_arXiv.txt": {
        "abstract": "{} {We conducted a search for brown dwarfs (BDs) and very low mass (VLM) stars in the 625~Myr-old Hyades cluster in order to derive the cluster's mass function across the stellar-substellar boundary.} {We performed a deep (I=23, z=22.5) photometric survey over 16~deg$^2$ around the cluster center, followed up with K-band photometry to measure the proper motion of candidate members, and optical and near-IR spectroscopy of probable BD and VLM members.} {We report the discovery of the first 2 brown dwarfs in the Hyades cluster. The 2 objects have a spectral type early-T and their optical and near-IR photometry as well as their proper motion are consistent with them being cluster members. According to models, their mass is 50 Jupiter masses at an age of 625~Myr. We also report the discovery of 3 new very low mass stellar members of the cluster, and confirm the membership of 16 others. We combine these results with a list of previously known cluster members to build the present-day mass function (PDMF) of the Hyades cluster from 50 Jupiter masses to 3~M$_\\odot$. We find the Hyades PDMF to be strongly deficient in very low mass objects and brown dwarfs compared to the IMF of younger open clusters such as the Pleiades. We interpret this deficiency as the result of dynamical evolution over the past few 100~Myr, i.e., the preferential evaporation of low mass cluster members due to weak gravitational encounters.} {We thus estimate that the Hyades cluster currently hosts about 10-15 brown dwarfs, while its initial substellar population may have amounted up to 150-200 members.} ",
        "introduction": "The determination of the initial mass function (IMF), i.e., the mass frequency distribution of objects formed in molecular clouds, yields strong constraints to star formation theories (Padoan \\& Nordlund 2002; Bate \\& Bonnell 2005; Larson 2005; Jappsen et al. 2005). In the last 10 years, the IMF has been derived from the most massive stars down to the substellar domain in a variety of environments : star forming regions (e.g. Luhman et al. 2003), young open clusters (e.g. Bouvier et al. 1998) and in the field (e.g. Reid et al. 1999). As brown dwarfs continuously cool down as they evolve (Chabrier et al. 2000), they are brighter when younger. Thus, the IMF could be derived down to a mass of a few Jupiter masses in star forming regions (Lucas et al. 2006; Caballero et al. 2007) and down to about 30 Jupiter masses in rich young open clusters (Moraux et al. 2003, 2007; Barrado et al. 2004; de Wit et al. 2006). While a large number of field brown dwarfs have also been discovered within a few tens of parsecs from the Sun, the derivation of the field IMF below 0.1~M$_\\odot$ still remains uncertain due to the lack of knowledge of the ages of substellar objects (Chabrier 2002; Allen et al. 2005; Cruz et al. 2007).  In a previous series of paper, we derived the lower mass function of a number of young open clusters down to about 30~M$_{Jup}$ (see Moraux, Bouvier \\& Clarke 2005 and Bouvier, Moraux \\& Stauffer 2005 for short reviews) and discussed the implications of an apparently universal cluster mass function for star formation scenarios (Moraux et al. 2007). Here, we report an extended, deep optical survey of the Hyades cluster aimed at detecting substellar objects and deriving the cluster's mass function across the stellar-substellar boundary. The motivations for this survey are twofold. Firstly, all previous searches for substellar objects in the Hyades have failed to report any positive brown dwarf detections (Reid \\& Hawley 1999; Gizis, Reid \\& Monet 1999; Dobbie et al. 2002). The lowest mass members reported so far includes LH~0418+13, with a spectral type M8.5 and an estimated mass of 0.083~M$_\\odot$ (Reid \\& Hawley 1999), the $\\simeq$0.081~M$_\\odot$ companions of a couple of 0.1~M$_\\odot$ Hyades probable members, LP~415-20 and LP~475-855 (Siegler et al. 2003), and an unresolved very low mass companion (0.070-0.095~M$_\\odot$) in the spectroscopic binary RHy~403 (Reid \\& Mahoney 2000). Yet, the discovery of brown dwarfs at an age of 625~Myr (Perryman et al. 1998) would provide a unique benchmark to calibrate the substellar evolution models. Secondly, the Hyades cluster is dynamically evolved (Adams et al. 2002). A significant fraction of its initial low mass population is therefore expected to have drifted away beyond the cluster boundaries (e.g. Reid 1993). Deriving the lower mass function of such an evolved cluster would provide a direct measurement of the rate at which low mass cluster members evaporate and populate the field. In turn, this measurement would allow us to test the validity of N-body simulations of the dynamical evolution of young clusters (e.g. Kroupa 1995; Portegies Zwart et al. 2001).  The Hyades cluster (Melotte 25, $\\alpha_{2000}$=04$^h$26$^m$54$^s$, $\\delta_{2000}$=+15$\\degr$52$\\arcmin$; $l$=180.05$\\degr$, $b$=-22.40$\\degr$) is the closest rich open cluster to the Sun. Perryman et al. (1998) derived its main structural and kinematical properties based on Hipparcos measurements~: a distance of 46.3$\\pm$0.27~pc, an age of 625$\\pm$50~Myr, a metallicity [Fe/H] of 0.14$\\pm$0.05, a present-day total mass of about 400~M$_\\odot$, a tidal radius of 10.3~pc, a core radius of 2.5-3.0~pc and negligible extinction on the line of sight. The large proper motion of the cluster ($\\mu\\simeq$100~mas yr$^{-1}$) can be easily measured from imaging surveys over a timeframe of only a few years, which helps in assessing cluster's membership.  In Section~2, we describe the optical survey we performed over the central 16~deg$^2$ of the Hyades cluster, as well as follow-up K-band photometry and both optical and infrared spectroscopy. The proper motion of optically-selected candidate members is derived from optical and near-IR images obtained 2 to 3 years apart. In Section~3, we describe our selection of candidate members combining optical and inrafred photometry, as well as proper motion measurements and follow up spectroscopy. We thus report 21 probable members, of which 5 are new and 2 are the first Hyades brown dwarfs, with an estimated mass of $\\simeq$50~M$_{Jup}$. In Section~4, we discuss the spectral properties of the newly found low mass Hyades members, and derive a spectral type of T1 and T2 for the 2 brown dwarf candidates. We proceed in deriving the present-day mass function of the Hyades cluster from 0.050 to 3.0~M$_\\odot$, which we find to be strongly deficient in very low mass stars and brown dwarfs compared to the mass function of the younger Pleiades cluster. We discuss this result in the light of N-body simulations of the dynamical evolution of young open clusters. ",
        "conclusions": "From a deep, wide-field survey of the central region of the Hyades cluster, we identified new very low-mass cluster members, including 2 T-dwarfs with a mass of $\\simeq$ 50 Jupiter masses. The comparison of the Hyades lower mass function with that of younger clusters indicates that the cluster is strongly depleted in very low mass objects at an age of 625~Myr. We thus estimate a total of about 15 brown dwarfs in the present-day Hyades cluster compared to about 500 stars. This depletion appears to result from the preferential evaporation of the lowest mass cluster members over a timescale of several 100~Myr, as predicted by N-body models of the dynamical evolution of young clusters. A fraction of the substellar escapers should populate the solar neighborhood, and may be detected as members of the Hyades moving group from large scale surveys."
    },
    "0801/0801.2169_arXiv.txt": {
        "abstract": "We study the extent and covering fraction of cool baryons around galaxies of different luminosity and mass, based on a survey of Mg\\,II $\\lambda\\lambda\\,2796, 2803$ absorption features near known galaxies. The initial sample consists of 13 galaxy and absorber pairs and 10 galaxies that do not produce Mg\\,II absorption lines to within sensitive upper limits.  The redshifts of the galaxy and absorber pairs range from $z = 0.2067$ to 0.892 with a median of $z = 0.3818$. We find that galaxies at larger impact parameters produce on average weaker Mg\\,II absorbers.  This anti-correlation is substantially improved when accounting for the intrinsic luminosities of individual galaxies.  In addition, there exists a distinct boundary at $\\rho=R_{\\rm gas}$, beyond which no Mg\\,II absorbers are found.  A maximum likelihood analysis shows that the observations are best described by an isothermal density profile and a scaling relation $R_{\\rm gas}=91\\times (L_B/L_{B_*})^{(0.35\\pm 0.05)}\\ h^{-1}$ kpc (or $69\\ h^{-1}$ kpc at $W(2796)=0.3$ \\AA) with a mean covering factor of $\\langle\\kappa\\rangle=80-86$ \\%.  Together with the scaling relation between halo mass and galaxy luminosity inferred from halo occupation studies, this scaling of $R_{\\rm gas}$ indicates that gas radius is a fixed fraction of the dark matter halo radius.  We compare our results with previous studies and discuss the implications of our analysis for constraining the baryon content of galactic halos and for discriminating between competing scenarios for understanding the nature of the extended gas. ",
        "introduction": "The Mg\\,II $\\lambda\\lambda\\,2796, 2803$ doublets are among the absorption features commonly seen in the spectra of distant quasars that are produced by intervening gaseous clouds along the quasar lines of sight.  Their rest-frame absorption equivalent width $W(2796)$, ranging from $W(2796)\\apll 0.3$ \\AA\\ to $W(2796)>2$ \\AA\\ (e.g.\\ Steidel \\& Sargent 1990; Nestor \\etal\\ 2005; Prochter \\etal\\ 2006a), is found to represent the underlying gas kinematics (e.g.\\ Petitjean \\& Bergeron 1990; Churchill \\etal\\ 2000).  Based on comparisons of the abundance ratios of various associated ions, these Mg\\,II absorption transitions are understood to arise primarily in photo-ionized gas of temperature $T\\sim 10^4$ K (Bergeron \\& Stas\\'inska 1986; Hamann 1997) and neutral hydrogen column density $N(\\hI)=10^{18}-10^{22}$ \\cmjj\\ (Churchill \\etal\\ 2000; Rao \\etal\\ 2006).  In addition, surveys for galaxies near the observed Mg\\,II absorbers have often uncovered luminous galaxies at projected distances ($\\rho \\apll 50\\ h^{-1}$ kpc) and velocity separations ($\\Delta\\,v \\le 250$ \\kms) from the absorbers (Bergeron 1986; Lanzetta \\& Bowen 1990, 1992; Steidel \\etal\\ 1994; Zibetti \\etal\\ 2005; Nestor \\etal\\ 2007; Kacprzak \\etal\\ 2007).  The relatively large associated $N(\\hI)$ and the small separation between Mg\\,II absorbers and galaxies along common lines of sight indicate that Mg\\,II absorbers may offer a sensitive probe of photo-ionized halo gas around galaxies at redshifts from $z=0.3$ to $z=2.3$ in the optical spectral window.  Understanding the physical origin of these Mg\\,II absorbers bears significantly on all efforts to apply their known statistical properties for constraining the baryon content of dark matter halos on different mass scales (Tinker \\& Chen 2008).  Various theoretical models have been developed in the past that describe the nature of the absorbers in the context of gas accretion, including ram pressure stripped gas from accreted satellites (Wang 1993), gravitationally bound cold gas in halo substructures (e.g., Sternberg \\etal\\ 2002), and condensed cold clouds in a thermally unstable hot halo (Mo \\& Miralda-Escude 1996; Maller \\& Bullock 2004; Chelouche \\etal\\ 2007).  Insights to the origin of Mg\\,II absorbers can be obtained from their clustering amplitude on large scales ($\\apg\\,1\\,h^{-1}$ Mpc). The clustering of the absorbers is a consequence of dark matter halos in which they are found.  Measurements of the galaxy and absorber cross--correlation function offer a means of quantifying the mean halo mass of the absorbers. High-mass halos are expected to be highly clustered, while low mass halos on average have weaker clustering strength.  The cross--correlation function of Mg\\,II absorbers and luminous red galaxies (LRGs) have been measured at $\\langle z\\rangle=0.6$ by Bouch\\'e \\etal\\ (2006), who found that Mg\\,II absorbers of $W(2796)>0.3$ \\AA\\ cluster strongly with LRGs and that weaker absorbers with $W(2796)=0.3-1.15$ \\AA\\ on average arise in dark matter halos that are 10 times more massive than those of stronger absorbers with $W(2796)=2-2.85$ \\AA.  Both the observed clustering amplitude and the inverse correlation between the mean halo mass and absorber strength are difficult to interpret.  For example, the frequency distribution function of Mg\\,II absorbers show that there are on average 10 times more $W(2796)=1$ \\AA\\ absorbers than those of $W(2796)=2$ \\AA\\ (e.g.\\ Nestor \\etal\\ 2005; Prochter \\etal\\ 2006a) which, when combined with the clustering measurements, implies that the majority of absorbers reside in massive, highly biased halos.  At the same time, we expect that observed Mg$^+$ ions originate primarily in photo-ionized gas of temperature $T\\sim 10^4$ K and the halo gas in massive dark matter halos becomes too hot for abundant Mg$^+$ to survive.  In Tinker \\& Chen (2008; hereafter TC08), we developed a new technique that adopts the halo occupation framework for understanding the origin of QSO absorption-line systems.  Specifically, the technique adopts a model density profile for the Mg$^+$ ions in individual dark matter halos.  The ``cold'' baryon content of individual dark matter halos, as probed by the presence of Mg$^+$, is then constrained through matching the space density and clustering amplitude of dark matter halos with the observed frequency distribution function of Mg\\,II absorbers and their clustering amplitude.  Our model allows the possibility that a predominant fraction of the gas in massive halos is shock heated to the virial temperature of the halo and becomes too hot to host abundant Mg\\,II absorbing clouds.  Within the hot halo, we further allow the possibility that some dense, cold clouds may penetrate through as seen in high resolution cosmological simulations (Kravtsov 2003).  The result of this halo occupation analysis is the probability function $P(W|M_h)$ that characterizes, for each dark matter halo of $M_h$, the total probability of finding a Mg\\,II absorber of equivalent width $W$.  The halo occupation analysis shows that observations, including both the Mg\\,II absorber frequency distribution function and their clustering amplitude, demand a rapid transition in the halo gas content at $M_h^{crit}\\sim 10^{11.5}$ \\hmsol.  Below $M_h^{crit}$, halos contain predominantly cold gas and therefore contribute significantly to the observed Mg\\,II statistics. Beyond $M_h^{crit}$, the cold gas fraction is substantially reduced and presumably the halo gas becomes too hot to maintain a large contribution to Mg\\,II absorbers.  In order to reproduce the observed overall strong clustering of the absorbers and the inverse correlation between $W(2796)$ and halo mass $M_h$, roughly 5\\% of the gas in halos up to $10^{14}$ \\hmsol\\ is required to be cold.  It is understood under our model that {\\it the $W(2796)$ vs.\\ $M_h$ inverse correlation arises mainly as a result of an elevated clustering amplitude of $W(2796)\\apll 1$ \\AA\\ absorbers, rather than a suppressed clustering strength of $W(2796)\\apg 2$ \\AA\\ absorbers.}  The clustering amplitude of weak absorbers is elevated because of the presence of cold streams in massive hot halos.  The amount of cold gas in clusters is small, therefore these halos can only contribute to weak absorbers.  The initial results of our halo occupation analysis demonstrate that combining known statistics of dark matter halos with a simple model for their gas content can already reproduce the statistical properties known for Mg\\,II absorbers.  It provides a simple prescription for populating baryons in dark matter halos that can be compared directly to results from numerical simulations and offer insights for understanding the physics of gas accretion in halos of all mass scales.  However, some uncertainties remain in the halo occupation analysis.  First, a key parameter that constrains the distribution of Mg$^+$ ions in individual dark matter halos is their incidence rate, $\\kappa(M_h)$, per halo.  It specifies the total probability of detecting an absorber in a halo of mass $M_h$.  In the halo occupation analysis of TC08, we find that dark matter halos of $M_h=10^{11.5-12.5}$ \\hmsol\\ on average have 100\\% incidence rate of Mg$^+$ ions, $\\langle\\kappa(M_h)\\rangle=1$, at $R_{\\rm gas}\\le R_{\\rm 200}/3$ for Mg$^+$ ions and $\\langle\\kappa(M_h)\\rangle=0$ at larger radii.\\footnote{The halo size $R_{\\rm 200}$ corresponds to the radius, within which the enclosed mean density is $200\\times\\bar{\\rho}_m$ and $\\bar{\\rho}_m$ is the background density.  It is motivated by estimates of the virial radius both from the spherical collapse model and from $N$-body simulations that predict $\\approx 180\\times\\bar{\\rho}_m$.}  In halos of lower masses, $\\langle\\kappa(M_h)\\rangle$ declines sharply.  While a 100\\% incidence rate shows that every dark matter halo hosts a uniform gaseous halo of Mg$^+$ ions with size $R_{\\rm gas}$, the low $\\langle\\kappa(M_h)\\rangle$ is more difficult to interpret because $\\langle\\kappa(M_h)\\rangle$ represents a mean value averaged over all halos of mass $M_h$.  A low $\\langle\\kappa(M_h)\\rangle$ may be a result of only a small fraction of dark matter halos containing extended distributions of Mg$^+$ ions or a result of a small covering factor of Mg$^+$ in all halos.  In addition, we have assumed in our initial analysis that the gaseous extent is related to halo mass according to $R_{\\rm gas}= R_{{\\rm gas}*}\\,[M_h/(10^{12}\\,\\hmsol)]^{\\beta}$, and $R_{{\\rm gas}*}=50\\ h^{-1}$ physical kpc and $\\beta=1/3$.  This scaling relation follows the theoretical expectation between virial radius and halo mass, but a higher $R_{\\rm gas}$ would naturally lower $\\kappa$ for a fixed frequency distribution function of Mg\\,II absorbers.  Second, the physical origin of the cold clouds probed by the Mg\\,II absorption transitions is ambiguous.  While the mass scale at which halo gas is found to experience a transition from a cold-mode dominated state to a hot-mode dominated state agrees well with the expectation of theoretical models established to characterize the gas accretion history in dark matter halos (e.g.\\ Birnboim \\& Dekel 2003; Kere{\\^s} \\etal\\ 2005; Dekel \\& Birnboim 2006; Birnboim \\etal\\ 2007), the observed cold-hot transition can also be interpreted as a declining formation efficiency of cool clouds in a hot halo with increasing halo mass (e.g.\\ Mo \\& Miralda-Escude 1996; Maller \\& Bullock 2004).  Both of these scenarios predict that halo gas is primarily cold at $M_h\\apll 10^{11.5}\\,\\hmsol$.  On the other hand, Bouch\\'e \\etal\\ (2006) have argued that the inverse correlation between $W(2796)$ and clustering amplitude is suggestive of the absorbers arising in non-virialized gas flows, such as starburst winds.  While dense clumps in starburst driven outflows are expected to contribute to some fraction of the observed Mg\\,II absorbers, an important quantity to specify is the significance of this fraction. If a large fraction of Mg\\,II absorbers originate in cold dense clumps in starburst winds, then one must also account for the more complex star formation physics in efforts to apply known Mg\\,II statistics for constraining the gas content of dark matter halos. We defer to \\S\\ 5.3 for a more detailed discussion of the caveats in interpreting the known properties of Mg\\,II absorbers as evidence to support their origin in starburst winds.  Direct constraints on the covering fraction of Mg\\,II absorbing gas versus halo mass not only help to break the degeneracy between the fraction of dark matter halos containing cold baryons and the extent of cold gas in individual halos, in principle they also serve to discriminate between the starburst wind scenario and gas accretion. Under the starburst wind scenario, superwind driven outflows are expected to proceed over a finite angular span along the minor axis (e.g.\\ Heckman \\etal\\ 1990; Veilleux \\etal\\ 2005), and result in only partial covering of the halos.  In addition, because at a given epoch only some fraction of galaxies are found at starburst or post-starburst stages, we would expect a fraction of dark matter halos to contain extended Mg\\,II absorbing gas due to outflows.  To obtain empirical constraints on the extent and covering fraction of Mg$^+$ ions around galaxies of different luminosity and mass, we have initiated a survey of Mg\\,II absorbers around known galaxies at small projected distances ($\\rho\\apll 100\\ h^{-1}$ kpc) to a background QSO.  Accounting the presence or absence of Mg\\,II absorbers for an unbiased sample of galaxies at different impact parameters yields a statistical estimate of the gas covering fraction around field galaxies.  A detailed comparison between the absorber strength and galaxy properties further constrains the density profile and extent of the gas.  The primary objectives of our study are (1) to improve the uncertainties of our halo occupation analysis by obtaining empirical measurements of $R_{\\rm gas}$ and $\\beta$, and (2) to examine whether starburst driven outflows are the predominant mechanism for producing the observed Mg\\,II absorbers by obtaining an empirical constraint of the gas covering factor $\\kappa$.  Here we present the initial results of our study.  This paper is organized as follows.  In Section 2, we describe the design of our experiment for constraining the extent and covering fraction of Mg$^+$ ions.  In Section 3, we describe the imaging and spectroscopic data available for our experiment.  In Section 4, we examine the correlation between Mg\\,II absorption strength and galaxy properties, such as the impact parameter and luminosity.  We compare our results with previous studies and discuss the implications of our analysis in Section 5.  We adopt a $\\Lambda$CDM cosmology, $\\Omega_{\\rm M}=0.3$ and $\\Omega_\\Lambda = 0.7$, with a dimensionless Hubble constant $h = H_0/(100 \\ {\\rm km} \\ {\\rm s}^{-1}\\ {\\rm Mpc}^{-1})$ throughout the paper. ",
        "conclusions": "Using the sample of 13 galaxy--Mg\\,II absorber pairs and 10 galaxies at $\\rho\\apll 100\\ h^{-1}$ kpc that do not give rise to Mg\\,II absorption to a sensitive upper limit, we find that galaxies at larger impact parameters produce on average weaker Mg\\,II absorbers.  This anti-correlation is substantially improved, when accounting for the intrinsic luminosity of individual galaxies.  In addition, there exists a distinct boundary at $\\rho=R_{\\rm gas}$, beyond which no Mg\\,II absorbers with $W(2796)\\ge 0.01$ \\AA\\ are found.  A maximum likelihood analysis shows that the observations are best described by an isothermal density profile of the gas and a scaling relation $R_{\\rm gas}=91_{-8}^{+3}\\times (L_B/L_{B_*})^{(0.35\\pm 0.05)}\\ h^{-1}$ kpc with a mean covering factor of $\\langle\\kappa\\rangle=80-86$ \\% at $\\rho\\le R_{\\rm gas}$.  The best-fit profile applies to galaxies of $0.03\\,L_{B_*}$ to $1.6\\,L_{B_*}$ and therefore predominantly cold-mode halos.  For higher mass halos, we expect from TC08 that the corresponding absorption strength for a given $\\rho$ is reduced, namely a smaller $\\bar{W}_0$ in Equation (1).  In this section, we discuss the implications drawn from the results of our analysis.  \\subsection{Comparisons with Previous Studies}  Adopting the best-fit scaling relation and the gaseous profile described in Equation (1), we derive a corresponding halo size of $\\rho\\approx 69\\ h^{-1}$ physical kpc for galaxies of $M_B-5\\,log\\,h=-19.8$ at $W(2796)=0.3$ \\AA.  This halo size is consistent with the finding of Lanzetta \\& Bowen (1990) and the expectation inferred from known Mg\\,II frequency function (Nestor \\etal\\ 2005; see the descriptions in \\S\\ 5.2), but is larger than what was reported in Steidel (1995).  Steidel (1995) reported a characteristic size of Mg\\,II halo of $R_*\\approx 35\\ h^{-1}$ kpc for $L_*$ (corresponding to $M_{B_*}-5\\,\\log\\,h=-19.7$ for the cosmology adopted in this paper) galaxies, but it is not clear at what $W(2796)$ limit the halo size was estimated.  At $\\rho=35\\ h^{-1}$ kpc, we expect $\\bar{W}(2794)=0.9$ \\AA\\ from the best-fit model.  The high gas covering factor derived from our analysis is consistent with what is inferred from the sample of Kacprzak \\etal\\ (2008). Applying the best-fit luminosity scaling relation of gaseous extent, we find that four of the 28 galaxies within the expected gaseous extent for $W_r(2796)> 0.3$ \\AA\\ absorbers in the Kacprzak \\etal\\ sample have a corresponding Mg\\,II weaker than the 0.3 \\AA\\ threshold. The inferred covering fraction is $\\approx 86$\\%.  At the same time, this high covering fraction disagrees with the value reported in Tripp \\& Bowen (2005), the only other on-going survey program designed to constrain gas covering fraction by searching for corresponding Mg\\,II absorbers at the locations of known {\\it field} galaxies\\footnote{We note two earlier surveys by Bechtold \\& Ellingson (1992) and Bowen \\etal\\ (1995) that include a dominant fraction of galaxies in dense cluster environment.  Both studies reported a substantially lower covering fraction, $\\langle\\kappa\\rangle\\apll 20$ \\% for $W(2796)>0.1$ absorbers.  The low covering fraction of strong absorbers is understood, because galaxies clusters reside in massive dark matter halos of $M_h> 10^{14}\\,\\hmsol$ and the fraction of ``cold'' gas in these massive halos is expected to be reduced (see TC08).}.  The authors obtained follow-up spectra of background QSOs that are located within $\\rho=10-39\\ h^{-1}$ physical kpc of 20 known galaxies at $z_{\\rm gal}=0.31-0.55$.  They found that 10 of the 20 galaxies do not produce Mg\\,II absorbers to within a 2-$\\sigma$ upper limit of $W(2796)=0.1$ \\AA, and concluded that the incidence of Mg$^+$ ions is $\\approx\\,50$ \\% at $\\rho\\apll 40\\ h^{-1}$ kpc in individual galactic halos.  Because more luminous galaxies have larger extended Mg\\,II halos according to the scaling relation presented in Equations (2), (6), and (7), this discrepancy can be understood if many of the galaxies in Tripp \\& Bowen are located at impact parameters that are larger than their expected $R_{\\rm gas}$.  Recall that we obtained a similar estimate of $\\langle\\kappa\\rangle\\approx 57$ \\% at $\\rho< 100\\ h^{-1}$ kpc (or $\\langle\\kappa\\rangle\\approx 50$ \\% at $\\rho< 45\\ h^{-1}$ kpc), {\\it prior to} accounting for the scaling relation between $R_{\\rm gas}$ and $L_B$ in individual galaxies (\\S\\ 4).  In addition, the sample in Tripp \\& Bowen includes galaxies of luminosity ranging from $0.3\\,L_*$ to $5\\,L_*$.  Those galaxies that are more luminous than $L_*$ are also expected to reside in more massive halos.  At $z=0$, $5\\,L_*$ galaxies are expected to reside in $> 10^{13}\\,\\hmsol$ halos (Tinker \\etal\\ 2007).  If the host dark matter halos are substantially more massive than $10^{12.0}\\,\\hmsol$, then we expect from the halo occupation analysis in TC08 that they do not contribute significantly to strong Mg\\,II absorbers.  In summary, further inspections of the properties of these galaxies and estimates of the gas covering fraction $\\kappa$ as a function of radius and $W(2796)$ are crucial for resolving this discrepancy.  \\subsection{The Origin of Mg\\,II Absorbers}  The scaling relation and the constraint of $\\langle\\kappa\\rangle$ together also allow us to test whether the Mg\\,II absorbers trace a significant and representative portion of the galaxy population or merely a special class such as those galaxies undergoing a starburst episode.  If extended Mg\\,II halos are a generic feature of field\\footnote{We define field galaxies as those that are not associated with a known cluster.} galaxies at all epochs, then we can derive the expected number density of Mg\\,II absorbers by combining known space density of galaxies (versus luminosity) and the product of the cross section and covering factor of the Mg$^+$ ions.  An agreement with observational results from surveys of Mg\\,II absorbers would challenge the scenario that attributes a large fraction of these absorbers to starburst systems\\footnote{However, we will not to able to rule out a scenario in which the blow-out winds are common in most galaxy histories.  Given enough time, long-lived outflowing gas gets mixed in with existing halo gas and newly accreted material from the IGM.  There is little distinction at this point between outflows and accreted clouds.  It is also not clear whether the outflow material would maintain the same degree of complex gas kinematics as those that are closer to the parent starburst regions.}.  The predicted number density of Mg\\,II absorption systems arising in the extended gaseous halos of galaxies may be evaluated according to \\begin{eqnarray} \\frac{d\\,{\\cal N}(W\\ge W_{lim})}{d\\,z} & = & \\frac{c}{H_0} \\frac{(1 + z)^2}{\\sqrt{\\Omega_{\\rm M}\\,(1+z)^3 + \\Omega_{\\rm \\Lambda}}} \\\\ & &  \\times \\int_0^\\infty d\\left(\\frac{L_B}{L_{B_*}}\\right)\\,\\Phi(L_B,z)\\,\\sigma(L_B,W_{lim})\\,\\kappa, \\nonumber \\end{eqnarray} where $c$ is the speed of light, $\\Phi(L_B,z)$ is the galaxy luminosity function, $\\sigma$ is the gas cross section for producing Mg\\,II absorbers of $W(2796)\\ge W_{lim}$ that scales with galaxy $B$-band luminosity, and $\\kappa$ is the halo covering factor.  Substituting the scaling relationship according to Equations (2), (6), and (7), we find \\begin{eqnarray} \\frac{d\\,{\\cal N}(W\\ge W_{lim})}{d\\,z} & = & \\frac{c}{H_0}\\frac{\\pi\\,\\langle\\kappa\\rangle\\,[R'_{{\\rm gas}*}(W_{lim})]^2\\,(1 + z)^2}{\\sqrt{\\Omega_{\\rm M}\\,(1+z)^3 + \\Omega_{\\rm \\Lambda}}} \\\\ & & \\times\\int_0^\\infty d\\left(\\frac{L_B}{L_{B_*}}\\right)\\,\\left(\\frac{L_B}{L_{B_*}}\\right)^{2\\,\\beta}\\,\\Phi(L_B,z)\\,, \\nonumber \\end{eqnarray} where $R'_{{\\rm gas}*}(W_{lim})$ is the extent of Mg\\,II halos that corresponds to $W(2796)=W_{lim}$.  To evaluate Equation (9), we adopt the luminosity function of the blue galaxy population at $z_{\\rm gal}=0.3-0.5$ from Faber \\etal\\ (2007), which is characterized by $M_{B_*}-5\\,\\log=-19.8$, $\\phi_*=9.6\\times 10^{-3}\\ h^{3}$ Mpc$^{-3}$, and a faint-end slope $\\alpha=-1.3$.  Including galaxies of $L_B\\ge 0.01\\,L_{B_*}$ and our best-fit $R_{{\\rm gas}*}$ and $\\langle\\kappa\\rangle$, Equation (9) yields $d\\,{\\cal N}/dz=0.028-0.03$ for $W(2796)\\ge 2$ \\AA\\ absorbers, $d\\,{\\cal N}/dz=0.21-0.23$ for $W(2796)\\ge 1$ \\AA\\ absorbers, and $d\\,{\\cal N}/dz=0.88-1.07$ for $W(2796)\\ge 0.3$ \\AA\\ absorbers at $z=0.5$.  The predictions based on galaxies more luminous than $0.01\\,L_*$ agree well with the observed $d\\,{\\cal N}/d\\,z$ from Nestor \\etal\\ (2005), lending strong support for the hypothesis that the Mg\\,II absorbers trace a significant and representative portion, rather than a sub-class, of the galaxy population.  The model profile of Equation (1) has been proved to be a good representation of the galaxy and absorber pairs at $\\rho=16.3-100\\ h^{-1}$ kpc, when the scaling relation of Equations (2), (6), and (7) is accounted for.  An important test is to examine whether the model remains a good fit at smaller $\\rho$ and strong absorbers. Identifying galaxy--absorber pairs at small impact parameters is difficult, because the glare of the background QSO often prohibits identification of faint galaxies at angular distances $\\theta<2''$ that corresponds to $\\sim 10\\ h^{-1}$ physical kpc at $z\\apll 0.5$.  Recent studies of intervening absorption-line systems along the lines of sight toward the optical afterglows of distant $\\gamma$-ray bursts (GRB) have uncovered a number of strong Mg\\,II absorbers (e.g.\\ Prochter \\etal\\ 2006b).  Unlike QSOs, GRB afterglows disappear after a while, permitting exhaustive searches of faint galaxies at $\\theta<2''$ that give rise to strong Mg\\,II absorbers along the lines of sight.  A redshift survey of galaxies along the sightline toward GRB\\,060418 ($z_{\\rm GRB}=1.491$) has uncovered the absorbing galaxies for three strong Mg\\,II absorbers at $z=0.603-1.107$ (Pollack \\etal\\ 2008, in preparation), similar to the redshift range covered by our sample.  The rest-frame $B$-band magnitudes of the galaxies range from $M_B-5\\,\\log\\,h=-16.8$ to $M_B-5\\,\\log\\,h=-18.3$; Impact parameters of the galaxies range from $\\rho=7.5\\ h^{-1}$ kpc to $\\rho=16.5\\ h^{-1}$ kpc.  Adopting the best-fit scaling relation, we include the three galaxy--Mg\\,II absorber pairs in the $W(2796)$ versus $\\rho$ plot for comparison (star points in Figure 5).  The good agreement between the best-fit model and the absorbers found along GRB sightlines provides further support for our model at small impact parameters.  \\begin{figure} \\includegraphics[scale=0.45]{f5.eps} \\caption{The $W(2796)$ versus $\\rho$ correlation scaled by galaxy $B$-band luminosity presented in Figure 4, but including pairs from Pollack \\etal\\ (2008; star points) and Bouch\\'e \\etal\\ (2007; points with horizontal errorbars) for comparisons.  The solid curves represent the best-fit model based on an isothermal density profile, with $a_h=0.1\\,R_{\\rm gas}$ (top curve) and $a_h=0.2\\,R_{\\rm gas}$ (bottom curve).  The pairs of Pollack \\etal\\ are for galaxies and Mg\\,II absorbers identified at $z=0.6-1$ along the sightline toward GRB\\,060418.  The errorbars in these data points are smaller than the size of the symbols.  The pairs of Bouch\\'e \\etal\\ are galaxies found at the positions of known Mg\\,II absorbers based on the presence of H$\\alpha$ emission lines.  Errorbars represent the uncertainties of the inferred $M_B$ from the $M_B$ versus $L(\\ha)$ correlation published in Tresse \\etal\\ (2002).  Crosses represent the absorbers in the Bouch\\'e sample that are likely ($\\apg 40$ \\% probability) to be a damped \\lya\\ absorber of $N(\\hI)\\ge 2\\times 10^{20}$ \\cmjj\\ based on the presence of strong Fe\\,II and Mg\\,I transitions (Rao \\etal\\ 2006).} \\end{figure}  \\subsection{Are Strong [$W(2796)>1$ \\AA] Absorbers at $z\\sim 1$ Produced Primarily in Starburst Outflows?}  As described in \\S\\ 1, the large line width of $W(2796)>1$ \\AA\\ Mg\\,II absorbers has recently been interpreted as due to outflow motions in starburst driven winds.  This scenario was initially motivated by the observed inverse correlation between $W(2796)$ and the clustering amplitude of Mg\\,II absorbers at $\\langle z\\rangle=0.6$ (Bouch\\'e \\etal\\ 2006).  In TC08, we have shown that this inverse correlation can be understood as due to an elevated clustering amplitude of $W(2796)\\apll 1$ \\AA\\ absorbers from contributions of survived cold clouds in massive hot halos, rather than a suppressed clustering strength of $W(2796)\\apg 2$ \\AA\\ absorbers.  In addition, we have shown in \\S\\S\\ 4, 5.1, and 5.2 that extended Mg\\,II halos are a generic feature of field galaxies over a wide luminosity range.  The predicted number density of the absorbers from adopting the scaling relation, Equations (2), (6), and (7), and the galaxy luminosity function agrees well with observations.  Specifically, we expect $d\\,{\\cal N}/dz\\approx 0.03$ for $W(2796)\\ge 2$ \\AA\\ absorbers, in comparison to the observed $d\\,{\\cal N}/dz=0.03-0.04$ for $W(2796)\\ge 2$ \\AA\\ absorbers at $z\\approx 0.5$ from Nestor \\etal\\ (2005).  While there is no clear indicator one can apply to unambiguously distinguish between gas inflows and outflows in the distant universe, we discuss several caveats related to the starburst outflow scenario.  If the absorbers are produced in starburst driven outflows, then a large fraction of the absorbing galaxy population might be expected to exhibit disturbed morphology (Mobasher \\etal\\ 2004).  The Mg\\,II absorbing galaxies in our sample exhibit regular disk morphologies in the HST images shown in Figure 2.  In particular, the strong absorber found with $W(2796)=1.55$ \\AA\\ at $z=0.892$ toward 3C336 is associated with an edge-on disk galaxy at $\\rho= 16.3\\ h^{-1}$ kpc.  The three Mg\\,II absorbers of $W(2796)>1$ \\AA\\ found toward GRB\\,060418 also exhibit regular disk structures in high-resolution images obtained both in space using HST and on the ground using LGAO on the Keck telescopes (Pollack \\etal\\ 2008, in preparation).  Kacprzak \\etal\\ (2007) also showed that there is little correlation between $W_r(2796)$ and galaxy asymmetry for Mg\\,II absorbers of $W_r(2796)>1.4$ \\AA.  While some strong Mg\\,II absorbers have been presented to show evidence of arising in expanding superbubble shells (Bond \\etal\\ 2001), available high-resolution images of other strong Mg\\,II absorbers do not support outflow being principally responsible for the observed strong absorbers.  This conclusion is also consistent with the spectral properties of Mg\\,II-selected galaxies published by Guillemin \\& Bergeron (1997)  On the other hand, Bouch\\'e \\etal\\ (2007) presented new observations from a targeted search for corresponding H$\\alpha$ emission in the vicinity of known Mg\\,II absorbers.  These authors successfully identified H$\\alpha$ emission near 14 of 21 Mg\\,II absorbers with $W(2796)>1.9$ \\AA\\ at $z\\sim 1$.  The integrated raw\\footnote{No extinction correction is applied to $L(\\ha)$ in our analysis, because the dust content is expected to vary substantially in $z\\sim 1$ galaxies (e.g.\\ Tresse \\etal\\ 2002) and is unknown for galaxies in the Bouch\\'e \\etal\\ sample.  Bouch\\'e \\etal\\ (2007) applied a constant correction to all observed \\ha\\ flux in their calculation, resulting in a simple scaling of the observed values.  The scatter in their measurements induced by dust remains the same.  For consistency, we consider in the following discussion only observed values without extinction correction both for the absorbing galaxies and for comparison results, such as the \\ha\\ luminosity function and the $M_B-L(\\ha)$ correlation, from the literature.} \\ha\\ luminosities range from $L_{\\ha}=6.8\\times 10^{40}\\ h^{-2}$ erg s$^{-1}$ to $L_{\\ha}=8.4\\times 10^{41}\\ h^{-2}$ erg s$^{-1}$.  The impact parameters range from $\\rho=1.4\\ h^{-1}$ kpc to $\\rho=35\\ h^{-1}$ kpc. Adopting a mean extinction correction $A_V=0.8$ mag, the authors inferred a mean star formation rate of $\\langle{\\rm SFR}\\rangle=5.9\\ h^{-2}$ M$_\\odot$ yr$^{-1}$ for these galaxies.  The \\ha\\ luminosity function at $z\\sim 1$ has been studied by a number of groups (see Doherty \\etal\\ 2006 for a list of references). It is characterized by a Schechter function of $L_{{\\ha}_*}=5\\times 10^{41}\\ h^{2}$ erg s$^{-1}$, $\\phi_*=0.013\\ h^3\\,{\\rm Mpc}^{-3} $, and $\\alpha=-1.3$.  The galaxies detected in \\ha\\ by Bouch\\'e \\etal\\ correspond to $0.14-1.68\\,L_{\\ha_*}$ at $z\\sim 1$, and therefore have only modest $L(\\ha)$.  Integrating the \\ha\\ luminosity function, we obtain a mean space density of 0.034 $h^{3}$ Mpc$^{-3}$ for galaxies of brighter than $0.1\\,L_{{\\ha}_*}$.  This is already 50\\% of all galaxies more luminous than $0.01\\,L_{B_*}$ at $z\\sim 1$, according to the best-fit luminosity function of the blue galaxy population from Faber \\etal\\ (2007).  This exercise shows that \\ha\\ emission is commonly seen in $z\\sim 1$ galaxies and the presence of modest $L(\\ha)$ does not argue for the presence of starburst outflows, unless the outflow material is long-lived.  For comparison, we include the 14 galaxy and absorber pairs from Bouch\\'e \\etal\\ (2007) in the $W(2796)$ versus $\\rho$ plot in Figure 5, applying the same scaling relation described in Equations (2), (6), and (7).  We estimate $M_B$ of these galaxies based on the published $M_B-L(\\ha)$ correlation in Tresse \\etal\\ (2002).  We adopt the observed $L(\\ha)$ and the scaling relation $M_B=73.5-2.27\\times\\log\\,L_(\\ha)$ that best describes the data in Tresse \\etal\\ (2002).  At small impact parameters, the model predictions depend sensitively on the choice of $a_h$ in Equation (1). In addition to the fiducial model, we also include in the figure a model with $a_h=0.1\\,R_{\\rm gas}$.  Nine of the 14 Mg\\,II absorbers are considered likely ($>40$ \\%) damped \\lya\\ absorbers (DLAs) of $N(\\hI)\\ge 2\\times 10^{20}$ \\cmjj\\ based on the presence of strong Fe\\,II and Mg\\,I transitions (Rao \\etal\\ 2006).  DLAs are the high-redshift analog of neutral gas regions that resembles the disks of nearby luminous galaxies (e.g.\\ Wolfe \\etal\\ 2005).  We therefore expect that the impact parameter distribution of galaxy-DLA pairs represents the typical extent of H\\,I disks at high redshifts.  We highlight the potential DLAs in the Bouch\\'e \\etal\\ sample as crosses in Figure 5.  After applying the best-fit scaling relation for Mg\\,II absorbers, we find that the impact parameters of these galaxy and potential DLA pairs range from $8\\ h^{-1}$ kpc to as high as $38\\ h^{-1}$ kpc.  The three largest separation pairs have $\\rho'>28\\ h^{-1}$ kpc, where $\\rho'$ is the luminosity scaled projected distance.  The scaling of impact parameter is justified, because smaller/fainter galaxies have on average smaller \\hI\\ disks (e.g.\\ Cayatte \\etal\\ 1994).  The large impact parameters, together with the likely high $N(\\hI)$, suggest that the galaxies responsible for these absorbers (presumably at smaller angular distances from the QSOs) may still be missing.  For the remaining sample, our model is in reasonably good agreement with their data.  In summary, we find a lack of empirical measurements that argue for starburst outflows to be a dominant mechanism for producing the observed strong Mg\\,II absorbers at $z\\sim 1$.  On the other hand, both the observed number density of Mg\\,II absorbers and the $W(2796)$ versus $\\rho$ correlation can be explained by extended Mg\\,II halos around typical field galaxies.  The fraction of Mg\\,II absorbers originating in starburst outflows is therefore expected to be small at $z\\sim 1$, unless starburst outflows are a common feature in field galaxies.  We caution, however, that the same conclusion cannot be automatically applied to absorbers at $z\\sim 2$ before a similar analysis is performed.  \\subsection{Constraints on the Extended Gaseous Halos Around Galaxies}  The sharp decline of $W(2976)$ found at $\\rho\\approx 90\\ h^{-1}$ kpc in Figure 4 clearly indicates that Mg$^+$ ions have a finite extent in the extended gaseous halos around galaxies.  This is reminiscent of the sharp boundary between the C\\,IV absorbing and non-absorbing regions around $\\langle z\\rangle=0.4$ galaxies reported in Chen \\etal\\ (2001b).  Given the nature of a photo-ionized gas (Bergeron \\& Stas\\'inska 1986; Hamann 1997), we examine whether the finite extent of Mg$^+$ ions can be understood as due to photo-ionization of the absorbing clouds at large radii.  In a two-phase halo model (e.g.\\ Mo \\& Miralda-Escud\\'e 1996), cold clouds that are responsible for producing the absorption features are pressure confined in a hot medium.  As the total gas density declines toward larger radii, the density of the cold clouds also decreases and the clouds become optically thin to the ultraviolet photons that can quickly ionize the Mg$^+$ ions (I.P.\\,$=$\\,15.035 eV).  With a known background radiation field, the location where the ionization transition occurs may be adopted to constrain the density of cold gas at large radii.  We carry out the photo-ionization analysis, using the Cloudy software (Ferland \\etal\\ 1998; version 06.02), and calculate the ionization fractions of Mg$^+$ ions for a grid of models with different total gas density and metallicity.  We adopt the spectral shape of the UV background radiation field of Haardt \\& Madau in Cloudy for the incident ionizing flux, and scale the radiation intensity to the background UV radiation field inferred from QSO proximity effect (Scott \\etal\\ 2002).  We adopt a plane parallel geometry for the clouds with a thickness of 100 pc, which is motivated by the observed coherent length of Mg\\,II absorbers toward lensed QSOs (Rauch \\etal\\ 2002).  The slab of gas is illuminated by the ionization radiation field on both sides.  \\begin{figure} \\begin{center} \\includegraphics[scale=0.4]{f6.eps} \\caption{Ionization fractions of Mg$^+$ ions and H$^0$ atoms versus total gas density $n_{\\rm H}$ for photo-ionized gas.  The values are calculated using the Cloudy software (Ferland \\etal\\ 1998; version 06.02) for gaseous clouds of plane parallel geometry and 100 pc thickness.  The input ionizing radiation intensity is set to be $J_\\nu(912\\AA)=7\\times 10^{-23}$ erg s$^{-1}$ Hz$^{-1}$ cm$^{-2}$ Sr$^{-1}$ suitable for the $z<1$ universe (Scott \\etal\\ 2002).  We perform the calculations, assuming solar (dotted curve for Mg and dot-dashed curve for H) and $1/10$ solar metallicity (solid curve for Mg and dashed curve for H).} \\end{center} \\end{figure}  Figure 6 shows the expected fractions of Mg$^+$ ions (solid and dotted curves) as a function of the total gas density $n_{\\rm H}$.  We include the calculated fraction of H$^0$ atoms (dashed and dash-dotted curves) for comparison.  The solid and dashed curves are for gas of $1/10$ solar metallicity, while the dotted and dash-dotted curves are for gas of solar metallicity.  Our photo-ionization calculations indicate that indeed the fraction of Mg$^+$ declines linearly at $n_{\\rm H}\\apll 0.025\\ {\\rm cm}^{-3}$ for a mean radiation intensity of $J_\\nu(912\\AA)=7\\times 10^{-23}$ erg s$^{-1}$ Hz$^{-1}$ cm$^{-2}$ Sr$^{-1}$ (which is appropriate for the observed background radiation field at $z<1$ Scott \\etal\\ 2002).  However, the transition is not as sharp as what is seen for hydrogen atoms.  Including the estimated ionization fraction from the Cloudy models and adopting the density profile parameterized in Equation (1), we derive the expected correlation between $W_r(2796)$ and $\\rho$ and find that photo-ionization is insufficient for explaining the presence of the sharp decline in the observed Mg\\,II absorption strength.  This extent of Mg\\,II gaseous halos ($\\approx 90\\ h^{-1}$ kpc around typical $L_*$ galaxies) is remarkably similar to the extent of C\\,IV gaseous halo ($\\approx 110\\ h^{-1}$ kpc around typical $L_*$ galaxies, after correcting for the $\\Lambda$-cosmology) reported by Chen \\etal\\ (2001b).  On the basis of 14 galaxy and C\\,IV absorber pairs and 36 galaxies that do not produce corresponding C\\,IV absorption lines to within sensitive upper limits, these authors find that the extent of C\\,IV-absorbing gas around galaxies scales with galaxy $B$-band luminosity as $R_{\\rm gas}({\\rm C\\,IV}) \\propto 110 \\times L_B^{0.5 \\pm 0.1}\\ h^{-1}$ kpc (corrected for the $\\Lambda$ cosmology).  In addition, there exists a sharp boundary between C\\,IV absorbing and non-absorbing regions at $R_{\\rm gas}({\\rm C\\,IV})$.  But the scatter in the observed C\\,IV absorption strengths at $\\rho< R_{\\rm gas}({\\rm C\\,IV})$ between different galaxies appears to be significantly larger than what is seen in Figure 4, and there is a lack of anti-correlation between the observed C\\,VI absorbing strength and galaxy impact parameter.  Given the distinct ionization potentials of the C$^{3+}$ and Mg$^+$ ions, {\\it the location where the abrupt transition between absorbing and non-absorbing regions occurs can be understood as a critical radius below which cool clouds can form and stablize in an otherwise hot halo}.  This two-phase model to interpret QSO absorption line systems was formulated in Mo \\& Miralda-Escud\\'e (1996) and later re-visited by Maller \\& Bullock (2004).  Taking into account the appropriate photo-ionization condition of the cold clouds, Mo \\& Miralda-Escud\\'e presented the expected Mg\\,II and C\\,IV absorption strength versus radius for halos of different mass.  Their models showed a lack of radial dependence on the C\\,IV absorption strength and a tight correlation between Mg\\,II $W_r(2796)$ and $\\rho$ for halos of a wide range of mass.  The differences in the radial dependence of different ions is understood as being due to the photo-ionization condition of cold clouds in the UV background radiation field.  These model expectations agree with observational findings very well.  The agreement between observations and a simple two-phase model is encouraging, although it is not clear how the two-phase model applies to cold-mode halos (e.g.\\ Dekel \\& Birnboim 2006).  Measurements of the relative abundances between C$^{3+}$ and Mg$^+$ as a function of radius will provide additional support for the origin of these metal-line absorbers, further constraining the metallicity and density of halo gas around galaxies.   \\subsection{Implications for the Halo Occupation Analysis and Future Work}  Our initial halo occupation analysis (TC08) included two principal assumptions for the extent of Mg$^+$.  First, we assumed that the extent of the cold gaseous halo is $R_{\\rm gas,12}=50\\,h^{-1}$ physical kpc for $M_h=10^{12}\\,\\hmsol$ halos, which is roughly $1/3$ of the halo radius $R_{200}$ at $z=0.6$.  Recall that the standard definition of a dark matter halo is an object with a mean interior density of 200 times the background density.  Second, the gaseous extent scaled with the halo mass as $M_h^{t}$ and $t=1/3$, following the expectation of the scaling relation between halo mass and the corresponding virial radius.  The results presented in \\S\\ 4 offer a direct test of these assumptions.  To compare the empirical constraints of Equations (6) \\& (7) with our initial model assumptions, we first derive the corresponding dark matter halo mass $M_h$ and their halo radius $R_{200}$ for galaxies of known absolute $B$-band magnitudes.  The adopted $M_{\\rm B_*}$ for the analysis in \\S\\ 4 corresponds to a magnitude of $M_{\\rm b_J}-5\\,\\log\\,h=-19.7$ in the bandpass of the Two-Degree Field Galaxy Redshift Survey (2dFGRS, Colless \\etal\\ 2001).  The 2dFGRS probes galaxies at $z\\sim 0.1$.  Excluding galaxies that exist as satellites in a cluster environment\\footnote{We note that the fraction of satellite galaxies is $\\lesssim 15$ \\% for $L_*$ galaxies (Tinker \\etal\\ 2007; ZCZ07).  In addition, none of the galaxies in our sample are observed to be part of a larger group or cluster.  Here we carry out our calculations, assuming that most of galaxies in our sample are central galaxies in their dark matter halos.}, the mean halo mass for $M_{\\rm B_*}$ galaxies is $10^{12.5}\\,h^{-1}\\,\\hmsol$ at $z\\sim 0.1$ (Tinker \\etal\\ 2007; van den Bosch \\etal\\ 2007).  At $z\\sim 1$, galaxies of this magnitude on average reside in $10^{11.9}\\,\\hmsol$ halos (ZCZ07).  Interpolating in $\\log\\,(1+z)$, we find that galaxies of $M_{\\rm B_*}-5\\,\\log\\,h=-19.8$ are expected to reside in halos of $M_h=10^{12.3}\\,\\hmsol$ at $z=0.4$ and that the halo radius for $M_{\\rm B_*}$ galaxies is $R_{200}=212\\,h^{-1}$ physical kpc.  The best-fit $R_{\\rm gas*}$ from Equation (7) therefore implies $R_{\\rm gas}\\approx 0.4\\,R_{200}$.  Next, we derive the expected scaling relation between $R_{\\rm gas}$ and halo mass $M_h$.  The halo occupation analysis of 2dFGRS galaxies from Tinker \\etal\\ (2007) showed that the relationship between $M_h$ and luminosity for galaxies of $L_{b_J}\\lesssim L_{*}$ at $z\\sim 0.1$ is $M_h = 10^{12.5}\\,(L_B/L_{B_*})^{1.3}\\,\\hmsol$.  this monotonic relationship is appropriate for galaxies that reside at the {\\it centers} of their dark matter halos and are the brightest galaxy in the halo.  A similar scaling relation was obtained from DEEP2 data by ZCZ07 for galaxies at $z\\sim 1$.  Because there is a monotonic relationship between halo mass and galaxy luminosity, the scaling relation of $R_{\\rm gas}$ with $L_B$ in Equation (2) is equivalent to a scaling relation between $M_h$ and $R_{200}$ that are related by $M_h\\propto R_{200}^3$.  Adopting $\\beta=0.35$ from Equation (6), we derive $R_{\\rm gas}=R_{\\rm gas*}\\times [M_h/(10^{12.3}\\,\\hmsol)]^{0.3}$ $h^{-1}$ kpc at $z\\sim 0.4$.  The calculations above show that observations support the initial assumptions in TC08 that the gaseous radius is a constant fraction of the halo radius for all halos and that the gaseous radius scales with halo mass according to $M_h^{t}$ with $t\\approx 1/3$.  We note that our galaxy sample therefore probes halo mass range from $10^{10.5}\\,\\hmsol$ to $10^{12.8}\\,\\hmsol$, and that the best-fit scaling relation is driven by galaxies in the cold-hot transition regime.  Next, we compare the observed covering fraction of Mg\\,II absorbers with the predicted mean of TC08.  As summarized in \\S\\ 1, we find in the halo occupation analysis of TC08 that dark matter halos of $M_h=10^{11.5-12.5}$ \\hmsol\\ on average have a covering fraction of unity $\\langle\\kappa(M_h)\\rangle=1$ at $R_{\\rm gas}\\le R_{\\rm 200}/3$ for Mg$^+$ and $\\langle\\kappa(M_h)\\rangle=0$ at larger radii.  This halo mass range corresponds to a luminosity range of $L_B=0.2-1.4\\,L_{B_*}$ at $z\\sim 0.4$, following the mass-to-light scaling relation discussed above. Sixteen galaxies in our sample have $L_B\\ge 0.2\\,L_{B_*}$, eleven of which are at $\\log\\,\\rho +0.14\\times (M_B-M_{B_*}) < \\log\\,R_{\\rm gas*}$.  We identify a corresponding Mg\\,II absorber for every one of these galaxies, indicating a mean covering fraction of $\\langle\\kappa\\rangle=100$ \\% at $R<R_{\\rm gas}$.  This empirical result agrees with our model expectation.  To assess the significance of detecting a Mg\\,II absorber in every galaxy in the sample of eleven, we adopt the binomial distribution function and find that we can constrain the underlying mean covering factor at $\\langle\\kappa\\rangle>0.76$ with greater than 95\\% confidence level.  Three of the seven remaining $L_B<0.2\\,L_{B_*}$ galaxies in our sample have $\\log\\,\\rho +0.14\\times (M_B-M_{B_*}) < \\log\\,R_{\\rm gas*}$.  Two of them do not have corresponding Mg\\,II absorbers identified to a sensitive upper limit, implying $\\langle\\kappa\\rangle=33$ \\% at $R<R_{\\rm gas}$ for fainter galaxies (presumably lower-mass halos). This is also consistent with the expectation of TC08, where we found a quickly declining $\\langle\\kappa\\rangle$ for halos of $M_h<10^{11.5}$ \\hmsol, although the observations have a much lower statistical significance.  In addition, we note that there exists an intrinsic scatter in the fiducial scaling relation between $M_h$ and $L_B$. This intrinsic scatter is expected to smear out the exact edge of gaseous halos, particularly when considering an ensemble of objects with a wide range of luminosity.  In summary, the initial results from the on-going survey of Mg\\,II absorbers in the vicinity of known galaxies close to QSO lines of sight have allowed us to study in detail the properties of extended gas around galaxies.  The best-fit scaling relation between the extent of Mg$^+$ and galaxy luminosity and the observed mean gas covering fraction together provide strong empirical evidence to support the halo occupation model of TC08.  To improve the statistical significances of various empirical parameters and to obtain a better understanding of the incidence of cold gas in dark matter halos, such as $\\kappa$ vs.\\ $M_h$ and $\\kappa$ vs.\\ $\\rho$, a larger sample of galaxy and Mg\\,II absorber pairs established for a wide range of galaxy luminosity and morphology is necessary.  In particular, only three galaxies in our current sample would be considered in a ``cold-mode'' halo based on their luminosity and none shows a corresponding Mg\\,II absorber.  Deep surveys targeting at fainter dwarfs are necessary to probe these ``cold-mode'' halos.  Finally, we encourage observations of Mg\\,II absorbers along the lines of sight toward close QSO pairs for obtaining a direct measure of the incidence of Mg$^+$ in individual halos."
    },
    "0801/0801.4673_arXiv.txt": {
        "abstract": "{ In this paper, we estimate the gamma-ray and neutrino fluxes coming from dark matter annihilation in a Milky Way framework provided by a recent N-BODY HORIZON simulation. We first study the characteristics of the simulation and highlight the mass distribution within the galactic halo. The general dark matter density has a typical $r^{-3}$ power law for large radii, but the inner behaviour is poorly constrained below the resolution of the simulation ($\\sim 200$ pc). We identify clumps and subclumps and analyze their distribution, as well as their internal structure. Inside the clumps, the power law is rather universal, $r^{-2.5}$ in the outer part with again strong uncertainties for smaller radii, especially for light clumps. We show a full-sky map of the astrophysical contribution to the gamma-ray or neutrino fluxes in this N-body framework. Using quite model independent and general assumptions for the high energy physics part, we evaluate the possible absolute fluxes and show some benchmark regions for the experiments GLAST, EGRET, and a km3 size extension of ANTARES like the KM3NeT project. While individual clumps seem to be beyond detection reach, the galactic center region is promising and GLAST could be sensitive to the geometry and the structure of its dark matter distribution. The detection by a km3 version of ANTARES is, however, more challenging due to a higher energy threshold. We also point out that the lack of resolution leaves the inner structure of subhalos poorly constrained. Using the same clump spectrum and mass fraction, a clump luminosity boost of order ten can be achieved with a steeper profile in the inner part of the sub-halos.}   \\vspace{1.cm} ",
        "introduction": "An ever increasing number of observational and theoretical results strongly suggest, or even require the existence of dark matter, whose enigma becomes thus crucial for the understanding of our universe.  Let us mention amongst others the WMAP results on CMB \\cite{Spergel:2006hy}, the rotation curves of disk galaxies \\cite{Bosma:2003yv}, the formation of large scale structures \\cite{White:1987yr}, the bullet cluster observation \\cite{Clowe:2003tk}, merger modeling and lensing results {\\it e.g.} \\cite{Ferreras:2007na}. Finally, the possibility of numerous extensions of high energy physics beyond the standard model (BSM) to provide new weakly interacting massive particle (WIMP) candidates for dark matter makes the hypothesis very appealing.  Nevertheless, the nature, the identification and the distribution of the dark matter are still open questions, intimately linked with the proof of its existence. Present and near future instrumental projects could bring welcome input to these questions. Namely, both astroparticle physics and collider searches will reach higher sensitivities with experiments like LHC (accelerator); EDELWEISS II, superCDMS and ZEPLIN (direct detection); ANTARES and ICECUBE (neutrino telescopes); GLAST (gamma space telescope) and PAMELA (charged particle search satellite). Amongst the different possibilities, indirect detection is particularly promising. Indeed, relic dark matter particles can accumulate in cosmic storage rings and annihilate. The decay of their annihilation products will give rise to secondary particle fluxes ($\\gamma, \\nu,e^+,\\bar{p}$), which could be detected by dedicated experiments indirectly indicating the presence of dark matter.  In this paper, we will focus on indirect detection of dark matter through gamma rays and neutrinos. Galaxies are thought to be interesting sources for this kind of detection, seen the amount of dark matter they are believed to harbour. As we will discuss later, the detectability of such gamma rays or neutrinos depends strongly on both the astrophysical assumptions on the dark matter distribution in the  halo and on the assumed high energy physics BSM scenario. Some studies concerning different particle physics models can be found in the literature (see \\cite{Jungman:1995df,Bertone:2004pz} for reviews). The popular BSM dark matter scenarios are typically supersymmetric models, models with extra dimensions, light dark matter, little Higgs model, inert doublet model ... or any extensions providing WIMP. The Milky Way astrophysical framework is commonly simplified with assumptions of spherical symmetry, now known to be incorrect \\cite{Athanassoula:2003yw,Allgood:2006,Athanassoula:2007ex,Heller:2007tr}, and typical smooth dark matter density functions extracted from N-body simulations \\cite{Navarro:1996gj,Kravtsov:1997dp,Moore:1999gc}. Few recent works \\cite{Stoehr:2003hf,Diemand:2006ik,Kuhlen:2007wv} and especially \\cite{Kuhlen:2008aw} with an impressive resolution treat in detail the astrophysical aspects of the gamma ray fluxes coming from dark matter annihilation in realistic simulation frameworks. Other works {\\it e.g} \\cite{Tasitsiomi:2002vh,Pieri:2003cq,Koushiappas:2003bn,Bi:2005im,Pieri:2007ir} consider also more sophisticated parametrization inspired by extrapolations of simulation results.  Typical simulation canvas consist of 1-100 millions of particles with mass around $10^5-10^6$ solar masses. The results reproduce well the large-scale structure formation and have now shown that virialized systems  are still left with surviving subhalos, also called {\\it  clumps} . The results are more and more promising with computing upgrades and resolution improvements. Nevertheless, some questions are still open with regard to observations. For instance, the radial density profiles predicted for the innermost region of galactic halos are quite cuspy, whereas observations suggest flat cores (see \\cite{Bosma:2003yv} for a review). Furthermore, simulations predict more numerous galactic satellites than observed for the Milky Way. Even if  N-body calculations may generate too concentrated objects, the simulated haloes are the only realistic or advanced dark matter distribution framework. Specifically, the estimation of dark matter detectability in our neighborhood depends on both the dark matter distribution in the Milky Way -- especially in the innermost region -- and on the number, the size and the concentration of the clumps. Depending on the assumptions or results on these key points, different results have been proposed in previous works \\cite{Koushiappas:2003bn,Evans:2003sc,Pieri:2005pg,Oda:2005nv,Koushiappas:2006qq,Bi:2005im,Baltz:2006sv,Strigari:2006rd,Pieri:2007ir,Strigari:2007at}.  The present article is devoted mainly to the astrophysical contribution concerning dark matter gamma and neutrino indirect detection. We calculated the possible gamma and neutrino fluxes sky map in a N-body simulation framework provided by a HORIZON project simulation \\cite{horizonweb}. The paper is organized as follows: section 2 gives the analysis of the numerical simulation and highlights the resulting dark matter distribution. In section 3, the gamma ray and neutrino flux calculation is shortly reviewed and a comparison of our estimates with regard to GLAST and ANTARES reach is presented. Conclusion and perspectives are given in section 4. ",
        "conclusions": "We evaluated the gamma and neutrino fluxes from dark matter annihilation in a galactic halo framework extracted from a cosmological N-body simulation of the HORIZON project. Although such simulations are the most elaborate and realistic framework for this kind of studies,  there are very few works concerning dark matter detection in N-body frameworks. With reasonable assumptions on the new physics beyond the standard model and on the cosmological scenario for dark matter particles, we proposed an absolute evaluation of gamma and neutrino fluxes. This allowed to test the galactic simulation framework with regard to dark matter indirect detection in future experiments like GLAST and the proposed km3 size neutrino telescope in the mediterranean sea extrapolated from ANTARES sensitivity and preliminary KM3NeT studies.  Our framework and even other cosmological simulations have a resolution limit which is lower than what would be desirable for this type of studies. Nevertheless, it was possible to reach a number of interesting results in our framework. In particular, we showed that the galactic center region is a good benchmark part  of the sky with regard to the GLAST sensitivity. Even if it is more challenging, this region should also be studied by future neutrino telescopes, especially a km3 size telescope. Individual clumps stand out clearer in the direction of the galactic anticenter. The concentration parameters of the clumps are then the crucial information that will determine the flux. Unfortunately, our simulation does not provide enough information on scales below 100 pc to clearly conclude about the detectability of individual subhalos. Furthermore, it should be stressed that our results for the central region of the Galaxy and for the clumps should be considered as a lower bound. A resolution improvement will increase the values of the central density and decrease the size of the central region within which we have little or no information. This will increase the signal in the direction of the Galactic center, as well as in the directions of the clumps. We can hope that specific geometry, non sphericity and structures especially in the central region will be highlighted by future gamma and neutrino observations.  Cosmological simulations made very important progress in the last few years, and a study of gamma-ray induced by dark matter annihilation in the most precise N-BODY framework can be found in \\cite{Kuhlen:2007wv} with a resolution allowing the identification of $\\lesssim 10^6 \\ M_{sun}$ clumps. Though the quality of this work is impressive, this minimal scale is still more than 10 orders of magnitude above the typical WIMP free-streaming scale ($\\sim 10^{-12-4} M_{sun}$ \\cite{2006PhRvL..97c1301P}) and also considerably bigger than the minimal surviving clump mass, which anyway is still under debate."
    },
    "0801/0801.3786_arXiv.txt": {
        "abstract": "We address the problem of particle acceleration in solar flares by fast modes which may be excited during the reconnection and undergo cascade and are subjected to damping.  We extend the calculations beyond quasilinear approximation and compare the acceleration and scattering by transit time damping and gyroresonance interactions. We find that the acceleration is dominated by the so called transit time damping mechanism. We estimate the total energy transferred into particles, and show that our approach provides sufficiently accurate results We compare this rate  with energy loss rate. Scattering by fast modes appears to be sufficient to prevent the protons from escaping the system during the acceleration. Confinement of electrons, on the other hand, requires the existence of plasma waves. Electrons can be accelerated to GeV energies through the process described here for solar flare conditions. ",
        "introduction": "The mechanism of energy  release and the process of its transfer to heating and acceleration of nonthermal particles in many magnetized astrophysical plasmas in general, and solar flares in particular, are still matter of considerable debate. Recent research shows turbulence may play an essential role in these processes. In the case of solar flares, it is believed that the energy comes  from release of stored magnetic energy via reconnection (see Priest \\& Forbes 2000). Turbulence is expected to develop since both ordinary and magnetic Reynolds number are very large. Both observational evidence and theoretical arguments suggest most of the energy is dissipated through turbulence (Biskamp 2000; Lazarian \\& Vishniac 1999, Lazarian et al. 2004; Shay et al. 2004). Recent high resolution observations of solar flares by {\\it Yohkoh} and {\\it RHESSI} satellites have supplied ample evidence that, at least from the point of view of particle acceleration, plasma turbulence and plasma waves appear to be the most promising agent  not only for the acceleration mechanism but also the general energizing of flare plasma (see {\\it e.g.} Petrosian 2007 and Petrosian \\& Liu 2004 and references cited there). This may also be true in other situations (Lazarian et al. 2002, Liu, Petrosian \\& Melia 2004). A substantial progress in understanding of incompressible (Shebalin, Matthaeus \\& Montgomery 1983; Higdon 1983; Goldreich \\& Sridhar 1995, hereafter GS95; Matthaeus et al 1996, Cho, Lazarian \\& Vishniac 2003; Biskamp 2003) and compressible MHD turbulence (see Lithwick \\& Goldreich 2001; Cho \\& Lazarian 2002, 2003), as well as MHD-turbulence-particle interactions (Chandran 2000, Yan \\& Lazarian 2002, 2004, 2008, henceforth YL02, YL04, YL08, respectively) has clarified many aspects of this problem.  It was demonstrated that the often adopted Alfv\\'en modes undergo anisotropy cascade mainly in  the direction perpendicular to the underlying magnetic field $B_0$ with a Kolmogorov spectrum. As a result they are inefficient in acceleration of particles. Slow modes which mix with Alfv\\'enic modes are also negligible for the same reason. Fast modes, on the other hand, develop on their own, as their phase velocity is only marginally affected by mixing motions induced by Alfv\\'en modes.   According to Cho \\& Lazarian (2002) fast modes follow an {\\it isotropic} ``acoustic\" cascade\\footnote{We discuss in \\S 5 the limitations of the model of turbulence in Cho \\& Lazarian (2002, 2003). On the basis of weak-turbulence theory, Chandran (2005) has argued that high-frequency fast waves with $k_\\|>>k_\\bot$ interacting with other fast waves generate high-frequency Alfven-waves with $k_\\|>>k_\\bot$. We expect that the scattering by thus generated Alfven  modes to be similar to scattering by fast modes that created them. Therefore, within the simplified approach adopted in the paper, we do not consider this type of interactions.} where the velocity scales as  $v_k\\propto k^{-1/4}$ and in each wave-wave collision a small fraction of energy equal to $v_{ph}/v_k$ is transferred to smaller scales. In low $\\beta$ medium, the three dimensional energy spectrum is (YL02) \\bea W(k)= \\frac{nm_i\\delta V^2}{8\\pi}k^{-\\frac{7}{2}}L^{-\\frac{1}{2}}\\left(\\begin{array}{cc}k_ik_j/k_\\bot^2&\\\\ &0\\end{array}\\right), \\label{fastspec} \\eea with a cascade time scale of \\be \\tau_{cas}=(v_{ph}/v_k)(kv_{k})^{-1}=(L/\\delta V)M_A^{-1}(kL)^{-1/2}. \\label{tcasfast} \\ee Here $k_{i,j}$ refers to the x,y components of the wave vector ${\\bf k}$, $n$ is the density of the plasma with ions of mass $m_i$ ($\\sim$ proton mass), $\\delta V$ is the initial perturbation injected at the outer scale $L\\equiv k_{\\rm min}$), $v_{ph}=\\omega/k$ is the phase speed of fast mode with frequency $\\omega$ and wave vector $k$, and $M_A=\\delta V/v_A$ is the Alfv\\'enic Mach number for the Alfv\\'en speed $\\simeq v_A=B/\\sqrt{4\\pi nm_i}$.  On small scales, the spectrum of turbulence is affected by damping. Damping becomes important at a wave vector $k_c$ when the damping time $\\Gamma_{k}^{-1}$ becomes comparable or shorter than the  cascading time $\\tau_{cas}$. Beyond this wavevector the turbulence spectrum falls off rapidly. In fully ionized plasma, there are basically two kinds of damping. Collisional damping is important on scales greater than the Coulomb collision mean free path $\\lambda_{\\rm Coul}$. In solar flares $\\lambda_{\\rm Coul}\\sim 5\\times 10^{7}{\\rm cm}\\left(\\frac{T}{10^7{\\rm K}}\\right)^{2}\\left(\\frac{10^{10}{\\rm cm}^{-3}}{n}\\right)$ and the relevant scales are shorter so that collisionless damping is dominant. Taking into account interactions with thermal and nonthermal particles, Petrosian, Yan \\& Lazarian (2006, henceforth PYL06) studied damping of fast modes in solar corona condition and showed that most the damping is also highly anisotropic. For most angles of propagation the inertial range is truncated at large scales. Only quasi-parallel and quasi-perpendicular waves reach short scales comparable to ion gyroradius% \\footnote{The quasi-perpendicular most likely will be damped because of magnetic field wanderings (see PYL06).}. For plasma $\\beta_p=P_{gas}/P_{mag}\\lesssim 0.1$, the damping due to electrons dominates as it's easier for electrons to catch up with the waves. The corresponding damping rate is  (Ginzburg 1961, YL02) \\begin{eqnarray} \\Gamma_{L} & = & \\frac{\\sqrt{\\pi\\alpha\\beta_p}}{2}\\omega\\frac{\\sin^{2}\\theta}{\\cos\\theta}\\exp\\left(-\\frac{\\alpha}{\\beta_p\\cos^2\\theta}\\right), \\label{Ginz} \\end{eqnarray} where $\\alpha=m_e/m_H$, $\\theta$ is the angle between the wave vector and the magnetic field. By equating the above equation and eq.(\\ref{tcasfast}), we obtain the cutoff scale of turbulence owing to the collisionless damping, \\be k_c L=\\frac{4M_A^4\\cos^2\\theta}{\\pi\\alpha\\beta_p\\sin^4\\theta}\\exp\\left(\\frac{2\\alpha} {\\beta_p\\cos^2\\theta}\\right). \\label{landauk} \\ee In what follows we shall use this spectrum of fast modes to calculate the acceleration and confinement of particles by fast modes in solar flare conditions.    \\begin{table} \\caption{Notations in this paper.} \\begin{tabular}{|l l|} \\hline a & power law index of nonthermal particle distribution\\\\ A(E) & acceleration rate\\\\ $E_0$ & lower energy limit of nonthermal particles\\\\ {\\bf k}& wave vector\\\\ $\\omega$ & wave frequency\\\\  $k_c$& turbulence cutoff wavenumber due to damping\\\\ L& energy injection scale\\\\ $\\delta V$& injection speed\\\\ $M_A$ & Alfv\\'enic Mach number $\\delta V/v_A$\\\\ n & number density of the corona\\\\ N(E) & number density of nonthermal particles per energy bin\\\\ $N_0$ & N(E) at the lower energy limit $E_0$\\\\ T & corona temperature\\\\ $v_A$ & Alfv\\'en speed\\\\ v & particle speed\\\\ $\\mu$ & cosine of particle pitch angle\\\\ W({\\bf k})& turbulence spectral energy density\\\\  $\\alpha$& ratio of electron mass $m_e$ to ion mass $m_i$\\\\ B& magnetic field\\\\ $\\beta_p$& ratio of gas pressure to magnetic pressure\\\\ $\\beta$& v/c\\\\ $\\beta_A$& $v_A/c$\\\\ $\\Gamma$ & wave damping rate\\\\ $\\tau_{cas}$ & cascading time of the turbulence\\\\ $\\tau_{acc}$ & acceleration time scale\\\\ $\\tau_{loss}$ & energy loss time scale of particles\\\\ $\\tau_{esp}$ & escaping time of CRs from the system\\\\ $\\eta$ &$\\cos\\theta$\\\\ $\\eta_c$ &$v_A/v$\\\\ $\\lambda_{Coul}$& Coulomb collisional mean free path\\\\ $\\lambda_\\|$ &parallel mean free path of CRs\\\\ $\\theta$& pitch angle of a fast modes\\\\ $\\delta \\theta$& variation of $\\theta$ in turbulence\\\\ \\hline \\end{tabular} \\end{table}  \\begin{table} \\caption{The physical parameters of solar flares we adopted.} \\begin{tabular}{|c|c|c|c|c|} \\hline T(KeV)&n(cm$^{-3}$)&$\\beta_p$&L(cm)&$M_A$ \\\\ \\hline 1&$10^{10}$&$0.01, 0.04, 0.1$&$10^9$&$0.3$\\\\ \\hline \\end{tabular} \\end{table}  In view of various difficulties with quasilinear theory (QLT), a number of nonlinear theories (NLT) have been proposed (see Dupree 1966; V\\\"olk 1973; Jones, Kaiser \\& Birmingham 1973; Goldstein 1976; Felice \\& Kulsrud 2001; Matthaeus et al. 2003; Shalchi 2005). Based on particle trapping due to large scale magnetic perturbations (V\\\"olk 1975), we developed a nonlinear formalism in YL08 to treat cosmic ray scattering in MHD turbulence. In view of these progresses, we believe the time is ripe to investigate the stochastic acceleration by the tested model of MHD turbulence in solar flares. In  \\S 2 we describe how the nonlinear effects are formulated. In \\S 3 and \\S 4 we present results on acceleration and confinement of the particles, and in \\S 5 and \\S 6 we present a brief discussion and summary of the results.  Some mathematical details are given in the appendix. ",
        "conclusions": "In this paper we have discussed the effects that the cascade of turbulence has on acceleration and heating of Solar corona. There the energy is injected at large scales much larger than any plasma scale concerned and this justifies a magnetohydrodynamic treatment of the large scale motions. Due to recent insights into the physics of MHD cascade and its interaction with charged particles we reduced a complex problem of acceleration and heating to a more manageable problems of interactions of Alfv\\'en, slow and fast modes with plasma and energetic particles.  Alfv\\'enic turbulence is inefficient in scattering and accelerating particles because of its anisotropy (Chandran 2000, YL02). Fast modes, instead, have been identified as the MHD turbulence modes dominating the interaction with cosmic rays (YL02,YL04). In this paper, we apply the result to the stochastic acceleration in solar flares.  We assume that the MHD turbulence is strong, i.e. that the critical balanced condition is satisfied for Alfv\\'enic modes. Therefore, to describe fast mode of MHD turbulence we appeal to the results by Cho \\& Lazarian (2002, 2003) on the scaling and coupling of Alfv\\'enic and fast modes. The corresponding papers claim the isotropy of the fast modes. One may expect to see deviations from isotropy, however. For instance, slab Alfv\\'en modes created by streaming instabilities are subject to non-linear damping by the ambient Alfv\\'enic turbulence (YL02, YL04, Farmer \\& Goldreich 2004, Lazarian \\& Beresnyak 2006). Beresnyak \\& Lazarian (2008) showed that the corresponding damping of the quasi-parallel modes depends on the angle between their ${\\bf k}$ vector and the direction of magnetic field. One might expect quasi-slab fast mode to be subject to a similar damping by strong Alfv\\'enic turbulence. This, however, has not been demonstrated so far.  If at the injection scale $\\delta B\\ll B$, the Alfv\\'enic turbulence is weak (see Galtier et al 2000) and develop a cascade with $k_{\\|}=const$. The interaction of the weak Alfv\\'enic turbulence with other modes can be very different from that of the strong Alfv\\'enic turbulence. For instance, it was shown by Chandran (2005) that fast modes develop anisotropy (i.e. the energy in the quasi-slab modes is reduced) owing to their interaction with Alfv\\'en modes in the weak  regime. However, such the Alfvenic weak turbulence has a limited inertial range (see discussion in Cho, Lazarian \\& Vishniac 2003) and at sufficiently large $k$ transfers into a strong turbulence. Moreover, magnetic reconnection in Solar Flares should produce perturbations $\\delta B\\sim B$, which should induce strong turbulence from the very beginning.  Strong Alfv\\'enic turbulence is characterized by $k_{\\bot}\\gg k_{\\|}$ with the GS95 relation $k_{\\|}\\sim k_{\\bot}$ defining a cone in the Fourier space where most of the turbulent energy resides. However, the aforementioned relation should not be understood too literally. The energy outside the cone is not zero and the modes with $k_{\\|}>k_{\\bot}$ are present. Such Alfv\\'enic modes are weakly interacting even being a part of the strong Alfv\\'enic turbulence. Therefore, the Chandran (2005) model is applicable to them.  Nevertheless, according to Cho, Lazarian \\& Vishniac (2002) the energy in these modes is exponentially reduced\\footnote{Because of this exponential reduction the rates of Alfv\\'enic scattering in YL02, which was the first study to take into account the modes outside the GS95 cone, were still grossly subdominant to the fast modes.}. Therefore, assuming that the Alfv\\'enic turbulence is injected at a scale much larger than the typical gyroradius of the energetic particles, we did not consider the anisotropies, introduced by the process of the fast wave cascading induced by Alfvenic modes with large $k_{\\|}$ and small $k_{\\bot}$ (cf. Chandran 2005).\"  The model for MHD turbulence that we adopted in the paper is the turbulence where the back-reaction of energetic particles is limited to changing the cut-off of the turbulence. A more fundamental modifications of turbulence are conceivable, however. For instance, Lazarian \\& Beresnyak (2006) argued that compressions of the fluid of energetic particles may result in the gyroresonance instabilities that can induce an additional quasi-parallel component of the Alfvenic waves. If true, these waves would interact with fast modes in the manner described by Chandran (2005), which would affect the fast mode isotropy. However, if substantial portion of energy resides within these quasi-slab modes, their major effect will be the direct gyroresonance acceleration of energetic particles. We felt that the modification of MHD turbulence arising from the instabilities within energetic particles, e.g. by  the process in Lazarian \\& Beresnyak (2006), is beyond the scope of the present paper.  Apart from the issue of isotropy, there are potential issues related to the exact scaling of fast modes. One dimensional numerical simulations in Suzuki, Lazarian \\& Beresnyak (2007) indicate that fast modes may develop a shock-like cascade, which differs from the finding in 3D MHD calculations in Cho \\& Lazarian (2002, 2003). We adopted the model from the latter works, but wait for higher resolution 3D numerical runs to rectify the scaling.  In addition to being strong, MHD turbulence that we considered was balanced in the sense that the equal flux of energy was assumed in every possible direction. The properties of imbalanced turbulence (see Cho, Lazarian \\& Vishniac 2002; Lithwick, Goldreich \\& Sridhar 2007; Beresnyak \\& Lazarian 2007, Chandran 2008) can be very different from the balanced one. However, we expect the flow of Alfv\\'en waves, which constitute the weak turbulence, to be  subjected to reflection within a Solar corona environment that we deal with. As the result we expect only marginal imbalance for the problem that we deal with.  A threshold for the TTD acceleration is $v\\gtrsim v_A$ set by the Cherenkov resonance condition. In the low $\\beta_p$ environment, the thermal protons can not be accelerated. The low energy threshold would be $\\sim 2(T/10^7)/\\beta_p$KeV. Thermal electrons, instead, can have TTD interactions unless the plasma beta is too low $\\beta_p\\lesssim m_e/m_p$. Our calculations show that TTD acceleration dominates over gyroresonance for the energy range we consider. The acceleration efficiency decreases with the plasma $\\beta_p$ as damping of the fast modes increase with $\\beta_p$.   Particle acceleration rate depends on the wave spectrum and the wave damping rate are partially determined by the particle spectrum. In general, it is required that a self-consistent treatment of the evolution of turbulence and particle acceleration. The calculations above assumes that the spectrum of the turbulence is determined by the cascade and damping by thermal particles. Our numerical calculation of these integrals shows that up to $\\sim$ 10\\% of total energy of turbulence is being transferred to nonthermal particles for the given set of parameters. In most cases, the back reaction of nonthermal particles is negligible as the damping due to the interaction with these nonthermal particles is smaller than thermal damping taking into account field line wandering. In some large flares, however, a large amount of particles need to be accelerated from the thermal reservoir to energies $\\gg k_BT$ (PYL06). In this case, the damping by nonthermal particles can be significant and one needs to solve the coupled equation of the evolution of turbulence and particles simultaneously.  Isotropization is important for the process of acceleration. If there is not enough isotropization, the acceleration by TTD will stop quickly as only parallel velocity is increased during the process. Moreover, without enough scattering, particles will leave the system before they get accelerated. Scattering by fast modes is shown to be adequate for isotropization of protons. Different from the acceleration, the scattering by gyroresonance, even smaller than the TTD scattering rate, plays an essential role in determining the scattering of parallel moving particles, which can be a substantial portion of the particles because of the TTD acceleration.  For electrons, due to their small gyro-radii, gyroresonance does not occur with MHD turbulence except for those high energy electrons. The actual scattering can happen with the plasma turbulence. The work by Cho \\& Lazarian (2004) shows that whistler modes are even more anisotropic than Alfv\\'en modes. We know gyroresonance is very inefficient with anisotropic turbulence.  Analogous to the MHD regime, parallel propagating modes may be generated by kinetic instabilities (see Tsytovich 1977) and be a candidate for interaction with electrons. The TTD interaction itself may induce gyroresonance instability through creating an anisotropic distribution of particles with respect to the magnetic field (Lazarian \\& Beresnyak 2006). Our estimate shows that the whistler waves (for moderate energy) and Alfv\\'en wave (for high energy) generated by the anisotropy instability can provide effective scattering and isotropization for the electron acceleration.  Acceleration of particles by fast modes were previously studied by a number of authors, including Miller, Larosa \\& Moore (1996), Schlickeiser \\& Miller (1998). In their studies, turbulence is assumed isotropic with either Kolmogorov or Kraichnan spectrum. Although coupled equations of wave and particle evolutions were solved in Miller, Larosa \\& Moore (1996), we feel that the one dimensional treatment they adopt is problematic, as the damping caused by the TTD interaction with particles is anisotropic. In this paper, we start with a more physically motivated and numerically tested of turbulence. We deal with fast modes, which are subject to much more efficient linear dissipation.  On the small scales, however, because of the dissipation above, fast modes develop anisotropy. In our treatment, this anisotropy is strongly affected by field line wandering which is determined by the Alfv\\'enic Mach number and plasma $\\beta$ and the efficiency of the acceleration is substantially influenced accordingly. In addition, scattering and confinement are treated using the nonlinear theory we have recently developed (Yan \\& Lazarian 2008). We believe that the approach that we developed is applicable beyond pure Solar Flare problems, e.g. it can be modified to study turbulent acceleration in the medium within clusters of galaxies (see Brunetti \\& Lazarian 2007)."
    },
    "0801/0801.4029_arXiv.txt": {
        "abstract": "We present the results from a multiwavelength campaign on the TeV blazar 1ES 1959+650, performed in May, 2006. Data from the optical, UV, soft- and hard-X-ray and very high energy (VHE) gamma-ray (${\\rm E} > 100$ GeV) bands were obtained  with the \\suzaku and \\swift satellites, with the \\magic telescope and other ground based facilities. The source spectral energy distribution (SED), derived from  \\suzaku and \\magic observations at the end of May 2006, shows the usual double hump shape, with the synchrotron peak at a higher flux level than the Compton peak. With respect to historical values, during our campaign the source exhibited a relatively high state in X-rays and optical, while in the VHE band it was at one of the lowest level so far recorded. We also monitored the source for flux-spectral variability on a time window of 10 days in the optical-UV and X-ray bands and 7 days in the VHE band. The source varies more in the X-ray, than in the optical band, with the 2-10 keV X-ray flux varying by a factor of $\\sim 2$. The synchrotron peak is located in the X-ray band and moves to higher energies as the source gets brighter, with the X-ray fluxes above it varying more rapidly than the X-ray fluxes at lower energies. The variability behaviour observed in the X-ray band cannot be produced by emitting regions varying independently, and suggests instead some sort  of ``standing shock'' scenario. The overall SED is well represented by an homogeneous one-zone synchrotron inverse Compton emission model, from which we derive physical parameters that are typical of high energy peaked blazars. ",
        "introduction": "It is widely accepted that the spectral energy distribution (SED) of blazars is dominated by a non-thermal continuum, produced within a relativistic jet closely aligned with the line of sight, making these objects very good laboratories to study the physics of relativistic jets. The overall emission, from radio to $\\gamma$--rays and, in some cases, to the multi--TeV band, shows the presence of two well--defined broad components (von Montigny et al. 1995; Fossati et al. 1998). Usually, for the blazars that are detected in the TeV bands, the first components peaks in the UV -- soft-X--ray bands (HBL: high energy peaked blazars, Padovani \\& Giommi (1995) and the second one in the GeV--TeV region. The blazar emission is very successfully interpreted so far in the framework of Synchrotron Inverse Compton models. The lower energy peak is unanimously attributed to synchrotron emission by relativistic electrons in the jet, while the second component is commonly believed to be Inverse Compton emission (IC) from the same electron population (e.g., Ghisellini et al. 1998), although different scenarios have been proposed (e.g. B\\\"ottcher 2007).  Since the discovery of the first blazar emitting TeV radiation, Mrk~421 (Punch et al. 1992), TeV blazars have been the target of intense observational and theoretical investigations. Indeed, the possibility of coupling observations of the emission produced by very high energy electrons, in the VHE (very high energy) band (up to Lorentz factors of the order of $10^7$), with observations in the soft and hard X--ray bands offers a unique tool to probe the cooling and acceleration processes of relativistic particles. In fact the synchrotron peak of these sources is usually located in the soft X--ray, while it is in the hard X--ray band that the synchrotron emission by the most energetic electrons can be studied and that the low energy part of the Compton emission component can start to dominate.  Studies conducted simultaneously in the soft and hard X--ray and in the VHE bands are of particular importance, since in the simple Synchrotron-Self Compton (SSC) framework one expects that variations in X--rays and TeV should be closely correlated, being produced by the same electrons (e.g. Tavecchio et al. 1998).  In fact even the first observations at X--ray and TeV energies yielded significant evidence of correlated and simultaneous variability of the TeV and X--ray fluxes (Buckley et al. 1996; Catanese et al. 1997). During the X-ray/TeV 1998 campaign on Mrk\\,421 a rapid flare was detected both at X--ray and TeV energies (Maraschi et al. 1999). Subsequent observations confirmed these first evidences also in other sources.  Note however that the correlation seems to be violated in some cases, as indicated by the observation of an ``orphan'' (i.e. not accompanied by a corresponding X--ray flare) TeV event in 1ES 1959+650 (Krawczynski et al. 2004). In the case of PKS\\,2155-304 a giant TeV flare recorded by HESS (Aharonian et al. 2007), with a TeV flux \"night-average\" intensity of a factor of $\\sim 17$ larger than those of previous campaigns,  was accompanied by an increase of the X-ray flux of only a factor of five without a significant change of the X-ray spectrum (Foschini et al. 2007). In the one-zone SSC scenario this can be accomplished with an increase of the Doppler factor and the associated relativistic electrons together with a decrease of the magnetic field. Therefore, it is important to obtain simultaneous observations over the largest possible UV and X-ray range together with simultaneous VHE observation to probe the correlation between the synchrotron and VHE emission.  To this end we organised a multiwavelength campaign to observe the blazar 1ES\\,1959+650 in the optical, UV, soft and hard X-ray up to the VHE gamma-ray (${\\rm E} > 100$ GeV) bands. This is a bright and flaring X-ray and VHE source that has already been observed many times in these bands. It was discovered in the radio band as part of a 4.85 GHz survey performed with the 91 m NRAO Green Bank telescope (Gregory \\& Condon 1991; Becker, White \\& Edwards 1991). In the optical band it is highly variable and shows a complex structure composed by an elliptical galaxy (M$_R =-23$, $z$=0.048) plus a disc and an absorption dust lane (Heidt et al. 1999). The mass of the central black hole has been estimated to be in the range $1.3-4.4 \\times 10^8 \\, {\\rm M_\\odot}$ as derived either from the stellar velocity dispersion or from the bulge luminosity (Falomo et al. 2002). The first X-ray measurement was performed by {\\it Einstein}-IPC during the Slew Survey (Elvis et al. 1992). Subsequently, the source was observed by \\rosat, {\\it Beppo}SAX, {\\it RXTE, ARGOS, XMM-Newton}. In particular two {\\it Beppo}SAX pointings, simultaneous with optical observations, were triggered in May-June 2002 because the source was in a high X-ray state. These data showed that the synchrotron peak was in the range 0.1-0.7 keV and that the overall optical and X-ray spectrum up to 45 keV was due to synchrotron emission with the peak moving to higher energy with the higher flux (Tagliaferri et al. 2003). The overall SED, with non simultaneous VHE data could be modelled with a homogeneous, one-zone synchrotron inverse Compton model. The results of a multiwavelength campaign performed in May-June 2003 are presented in Gutierrez et al. (2006). This campaign was triggered by the active state of the source in the X-ray band and it was found that the X-ray flux and X-ray photon index are correlated. A similar result was found by Giebels et al. (2002) using{\\it RXTE} and {\\it ARGOS} data. This correlation shows that the X-ray spectrum in the 1-16 keV band is harder when the source is brighter. In the VHE band the source  was detected by the HEGRA, Whipple and MAGIC telescopes (Aharonian et al. 2003; Holder et al. 2003a, Albert et al. 2006). One of the most important results of these observations is probably the ``orphan\" flare mentioned above, seen in the VHE band and not in X-rays (Krawczynski et al. 2004).  1ES\\,1959+650 is therefore one of the most interesting and frequently observed high energy sources of recent years. With the aim of obtaining a better description of the broad band X-ray continuum and in particular of observing simultaneously the synchrotron and IC components, we asked for simultaneous \\suzaku and \\magic observations that were carried out in May 23-25, 2006. Around the same epoch we obtained various Target of Opportunity (ToO) short pointings with {\\it SWIFT} and observed the source also in the optical R-band from ground. A preliminary analysis of these data is reported in Hayashida et al. (2007). In the following we report the data analysis (Sec.2) and the results (Sec.3). The discussion and conclusions are given in Sec. 4, where we model the SED in the framework of a homogeneous, one-zone SSC model. Throughout this work we use $H_{\\rm 0}\\rm =70\\; km\\; s^{-1}\\; Mpc^{-1} $, $\\Omega _{\\Lambda}=0.7$, $\\Omega_{\\rm M} = 0.3$. ",
        "conclusions": "The full SED of 1ES 1959+650 as measured at the end of May 2006 is reported in Fig.~\\ref{sed}, together with other historical data.  During our multiwavelength campaign we simultaneously observed the SED from the optical, to the UV, soft and hard X-rays and VHE bands, monitoring both the synchrotron and Compton components. The historical data in this figure show very strong changes in the X-ray band, while in the optical this is much more attenuated. A behaviour that is also found in the results obtained from our observing campaign.  During our multiwavelegth campaign the source is found to be in a high state with respect to the historical behaviour both in X-ray and optical (e.g. Tagliaferri et al. 2003 and Fig. \\ref{lc_optical}), although not at the highest state as observed in the X-ray (e.g. Holder et al. 2003b, see Fig. \\ref{sed}). In the VHE band, instead, the source is at one of the lowest state so far recorded. We also found that the X-ray fluxes at energies above the synchrotron peak vary more rapidly than the X-ray fluxes below the peak.  Also the VHE band shows historical strong variability, in particular if we consider that in this band there are fewer observations than in the optical or X-ray ones. However, from our data we do not see strong (i.e. a factor of 2-3) variability in the VHE band. Our \\magic data are probably monitoring the part of the SED slightly above the peak of the Compton component. Therefore, one would expect to see a high level of variability. The lack of variability in the \\magic data and the low flux level recorded both indicate that the source was not very active in this band. Overall we can say that during our campaign the source was quite stable (i.e. did not vary by more than a factor of 2) from the optical to the VHE band.  The observed X--ray variability behaviour allows a few interesting considerations about the properties of the emitting regions. First, note that the variability is not random, but follows a raising/decay trend on a timescale of $\\sim$10 days (see the \\swift-XRT results). In this observed time $\\Delta t$, a single blob moving with a bulk Lorentz factor $\\Gamma\\sim 18$ (see below) moves by a distance $\\Delta z \\sim c\\Delta t \\Gamma^2\\sim 2.7$ pc. Therefore we cannot assume that we are observing a single moving blob travelling that far, since the blob would expand, loose energy by adiabatic losses, and change (decrease) its magnetic field. This in turn would decrease the produced flux and would lengthen the variability timescale. Also the internal shock model (Spada et al. 2001, Guetta et al. 2004) can not explain the variability we are observing. In fact, in this model the radiation is produced in a shock resulting from the collision of two shells moving at slightly different velocities. In this case the variability is predicted to be erratic, therefore to explain the variability we are seen we have to finely tune the different $\\Gamma$ of the shells. We are thus led to consider the possibility that the observed radiation originates in the same region of the jet, through some kind of ``standing shock\". For instance, we might think to the interaction of a fast spine and a shear layer occurring at about the same distance from the central powerhouse (see Ghisellini, Tavecchio \\& Chiaberge 2005 for mode details, including the possibility of radiative deceleration of the spine through the ``Compton rocket\" effect in TeV blazars). A ``standing shock\" scenario has already been proposed by Krawczynski et al. (2002) in order to explain the tight correlation between X-ray and TeV flares observed in Mrk501 and it is discussed in some detail also by Sokolov et al.( 2004).  As we did with the previous multiwavelength observing campaigns on 1ES\\,1959+650 that we organised based on the {\\it Beppo}SAX observations (Tagliaferri et al. 2003), we can try to fit our SED with a homogeneous, one-zone synchrotron inverse Compton model. During the {\\it Beppo}SAX campaigns, in order to derive the SSC physical parameters we had to assume a value for the Compton component, that we derived by rescaling a non-simultaneous VHE spectrum based on the X-ray flux.  This time we have also the VHE observations, therefore both SSC components are constrained by real data. As shown by Fig. \\ref{sed}, the X-ray spectrum as observed by \\suzaku and \\swift is about a factor of 2 higher than the one measured with \\sax and also the synchrotron peak has moved to somewhat higher energy, confirming the previous results of a higher energy peak with higher fluxes (e.g. Tagliaferri et al. 2003), that is typical for HBLs (see the dramatic case of MKN\\,501, Pian et al. 1998). The optical fluxes are similar to the one reported for the 2002 SED. The observed VHE spectrum is similar, but lower, to that one of the 2002 SED. In summary the 2006 SED has optical fluxes that are similar to that ones of the 2002, the X-ray fluxes are a factor of 2 higher and the VHE fluxes are a factor of $\\sim 2$ lower. In the assumed one-zone SSC model, the source is a sphere with radius $R$ moving with bulk Lorentz factor $\\Gamma$ and seen at an angle $\\theta$ by the observer, resulting in a Doppler factor $\\delta$. The magnetic field is tangled and uniform while the injected relativistic particle are assumed to have a (smooth) broken power law spectrum with normalisation $K$, extending from $\\gamma _{\\rm min}$ to $\\gamma _{\\rm max}$ and with indexes $n_1$ and $n_2$ below and above the break at $\\gamma _b$. Assuming this model, the SED of May, 2006 can be well represented using the following parameters: $\\delta=18$, R$=7.3 \\times 10^{15}$ cm, B=0.25 G, K=$2.2 \\times 10^3$ cm$^{-3}$ and an electron distribution extending from $\\gamma _{\\rm min} = 1$ to $\\gamma _{\\rm max} = 6.0 \\times 10^5$, with a break at $\\gamma _{\\rm b}= 5.7 \\times 10^4$ and slopes $n_1=2$ and $n_2=3.4$. The intrinsic luminosity is $L' = 5.5 \\times 10^{40}$ erg s$^{-1}$. If we compare these values with the one we derived for the 2002 SED (though in that case we use a slightly different emission model), we saw that the parameters are very similar, with a source that is slightly more compact, a lower magnetic field and an almost identical Doppler factor. Similar values are also found to explain the SED of PKS2155-304 during and after the strong TeV flare observed in July, 2006; although in that case we found less steep slopes for the electrons and an higher value of $\\delta$ (see Foschini et al. 2007). Once again, the physical parameters that we derived assuming a one-zone SSC model are typical of HBL objects. Finally, the historical SEDs of 1ES1959+650 shows that in this source the synchrotron emission is dominating above the Compton one."
    },
    "0801/0801.1110_arXiv.txt": {
        "abstract": "Using the Gemini Near-InfraRed Spectrograph (GNIRS), we have completed a near-infrared spectroscopic survey for $K$-bright galaxies at $z\\sim2.3$, selected from the MUSYC survey. We derived spectroscopic redshifts from emission lines or from continuum features and shapes for all 36 observed galaxies. The continuum redshifts are driven by the Balmer/4000\\,\\AA\\ break, and have an uncertainty in $\\Delta{}z/(1+z)$ of $<0.019$. We use this unique sample to determine, for the first time, how accurately redshifts and other properties of massive high-redshift galaxies can be determined from broadband photometric data alone. We find that the photometric redshifts of the galaxies in our sample have a systematic error of 0.08 and a random error of 0.13 in $\\Delta{}z/(1+z)$. The systematic error can be reduced by using optimal templates and deep photometry; the random error, however, will be hard to reduce below 5\\%. The spectra lead to significantly improved constraints for stellar population parameters. For most quantities this improvement is about equally driven by the higher spectral resolution and by the much reduced redshift uncertainty. Properties such as the age, $A_V$, current star formation rate, and the star formation history are generally very poorly constrained with broadband data alone. Interestingly stellar masses and mass-to-light ratios are among the most stable parameters from broadband data. Nevertheless, photometric studies may overestimate the number of massive galaxies at $2<z<3$, and thus underestimate the evolution of the stellar mass density. Finally, the spectroscopy supports our previous finding that red galaxies dominate the high-mass end of the galaxy population at $z=2-3$. ",
        "introduction": "The extensive use of photometric redshifts has greatly enhanced our knowledge of the $z=2-3$ universe. While color criteria, such as the Lyman break technique \\citep[][]{st96a,st96b}, distant red galaxy selection \\citep[DRGs,][]{fr03,vd03} and BzK selection \\citep{da04} provide an easy identification of high-redshift galaxies, photometric redshifts allow the study of apparently unbiased samples. Our current understanding of the evolution of the mass-density \\citep[e.g.,][]{ru01,ru03,ru06,di03,dr05} and the luminosity function \\citep[e.g.,][]{dah05,sa06,ma07}, the nature of massive high-redshift galaxies \\citep[e.g.,][]{fo04,la05,vd06,pa06}, and galaxy clustering \\citep[e.g.,][]{da03,qu07a,fo07} essentially all rely on photometric redshifts.  Ideally, all these studies would have been based on spectroscopic redshifts. However, obtaining spectroscopic redshifts for the required samples is hampered by several obstacles. Due to their faintness, obtaining redshifts of high-redshift galaxies requires long integrations on 8-10m class telescopes. The largest samples of spectroscopically confirmed high redshift galaxies number in the 1000s \\citep[e.g.,][]{st03}, several orders of magnitude short of the largest photometric samples. Furthermore, and more fundamentally, current spectroscopic samples are strongly biased towards blue, star forming galaxies which are bright in the observer's optical (Lyman break galaxies, or LBGs). However, it has become clear that the typical massive galaxy at this redshift range is red in the rest-frame optical and faint in the rest-frame ultra-violet (UV), and blue LBGs constitute only $\\sim$20\\% of massive galaxies at $z=2-3$ \\citep{vd06}.  Obtaining spectroscopic redshifts for typical massive galaxies requires deep spectroscopy at near-infrared (NIR) wavelengths. Unfortunately, NIR observations are complicated by the combination of the high sky brightness, numerous bright and variable night sky lines and strong atmospheric absorption bands, and the limited field of view of current and planned NIR spectrographs. Thus, obtaining spectroscopic redshifts for hundreds or thousands of galaxies with $K\\sim21$ (the typical brightness of galaxies with $M>10^{11}\\,M_{\\odot}$ at $z\\sim2.5$) will not be feasible in the foreseeable future. Until the next generation of space missions and $>20$\\,m ground-based telescopes we remain largely dependent on photometric redshifts for studies of large and faint galaxy samples beyond $z>1.5$.  Our provisional dependency on broadband photometric studies requires a more accurate calibration and understanding of the involved systematics. The current spectroscopic samples used for calibration of photometric high-redshift studies are based primarily on optical spectroscopy. As these samples are biased towards un-obscured star-forming galaxies, their calibration may not be representative for the total sample of massive galaxies. Photometric properties of red, massive galaxies at high redshift are poorly calibrated, and since red galaxies dominate the high mass end at $2<z<3$, systematics may have large effects on the final results.  NIR spectroscopy on a substantial, unbiased sample of massive, high-redshift galaxies is needed to test our photometric studies, and obtain insights regarding possible systematic effects. The Gemini Near-InfraRed Spectrograph \\citep[GNIRS,][]{el06} is especially well-suited for spectroscopy of $z\\sim2.3$ galaxies, due to the large wavelength coverage offered by the cross-disperser (0.9-2.5 $\\mu$m). This coverage offers two advantages: there is a large probability of finding emission lines, and it allows the characterization of the continuum emission, which is particularly important for galaxies without detected lines. In \\cite{kr06b} we found that a substantial fraction of the massive galaxies at $z\\sim2.3$ have no detected emission lines. For these galaxies we derived spectroscopic redshifts by modeling the stellar continuum, driven by the Balmer/4000 \\AA\\ break \\citep{kr06a}.  In this paper we present our full survey of $K$-selected galaxies at $z\\sim2.3$, conducted with GNIRS on Gemini-South between September 2004 and March 2007. In total we integrated more than $\\sim$80 hours, divided over 6 observing runs, on a sample of 36 galaxies. In previous papers based on preliminary results of this survey we discussed the stellar populations \\citep{kr06b} and the origin of the line emission \\citep{kr07}. In this paper we will give an overview of the total survey (\\S~\\ref{dat}), compare photometric and spectroscopic redshifts (\\S~\\ref{red}) and stellar populations properties (\\S~\\ref{imp}), and discuss the implications for photometric studies (\\S~\\ref{imp}).  Throughout the paper we assume a $\\Lambda$CDM cosmology with $\\Omega_{\\rm m}=0.3$, $\\Omega_{\\rm \\Lambda}=0.7$, and $H_{\\rm 0}=70$~km s$^{-1}$ Mpc$^{-1}$.  The broadband magnitudes are given in the Vega-based photometric system unless stated otherwise. Furthermore, we will measure the scatter and the offset between various properties using the normalized biweight mean absolute deviation and the biweight mean, respectively \\citep{be90}. As biweight statistics are less sensitive towards outliers than the normal mean, and more efficient than the median, they are most appropriate for the small sample sizes in this work.  \\figa % ",
        "conclusions": "\\label{sum}  In this paper we present our complete NIR spectroscopic survey for $K$-selected galaxies at $z\\sim2.3$. We acquired high quality NIR spectra (1-2.5 $\\mu$m) with GNIRS for a total sample of 36 galaxies. The galaxies were selected to be bright in $K$ ($<$19.7) and have photometric redshifts in the range $2.0\\lesssim z\\lesssim 3.0$. All galaxies have photometric stellar masses $>10^{11} M_{\\odot}$. The distribution of observed $J-K$, $R-K$ and photometric rest-frame $U-V$ colors are similar as those of a photometric mass-limited sample extracted from the deep MUSYC fields. This suggests that our spectroscopic sample is representative for a mass-limited ($>10^{11} M_{\\odot}$) sample at $2<z_{\\rm phot}<3$. Our sample is not representative for a $K$-bright galaxy sample, as our redshift distribution is slightly higher. Nevertheless, the distribution of photometric rest-frame $U-V$ colors is similar for the GNIRS sample and the full $K$-bright photometric sample.  We successfully derived spectroscopic redshifts for all galaxies. For 19 of the galaxies we detected rest-frame optical emission lines, which provided us with accurate redshift measurements. For the remaining galaxies we derived redshifts by modeling the continuum spectra in combination with the optical photometry. We tested this method using the emission-line galaxies, and found a scatter of $\\Delta z/(1+z)=0.019$, and no significant systematic offset. The continuum redshifts are more accurate for optically red galaxies, with large Balmer/4000 \\AA\\ breaks, and for redshifts for which the break falls in an atmospheric window.  Comparison with spectroscopic redshifts shows that the original photometric redshifts are generally overestimated by $\\Delta z/(1+z)=0.08$, and have a scatter of $\\Delta z/(1+z)=0.13$. Both the systematic and the scatter are worse for galaxies with shallow NIR photometry ($K\\sim20$). For the galaxies with deeper NIR photometry ($K\\sim21$) we find a systematic offset of $\\Delta z/(1+z)=0.03$ and a scatter of $\\Delta z/(1+z)=0.08$.  In addition to shallow NIR photometry, the lack of well calibrated galaxy templates is another important cause for the large systematics and scatter of photometric redshifts. Especially, dusty and/or young galaxies have high \\dzs. This may not be surprising as the original template set of our photometric redshift code lacked these galaxies. We almost completely remove the systematic offset when repeating the fitting procedure replacing the original templates by 6 SEDs chosen from fits to the spectroscopic sample. We examine other possible causes for the scatter and systematics in \\dzs\\ as well. Emission line contamination is only of minor importance for this sample and may affect the \\zp\\ for some individual galaxies. Furthermore, we find a correlation with $K$-band magnitude, such that $K$-bright galaxies in general have larger \\dzs\\ than those faint in $K$. This effect may be diminished by using a luminosity prior in the photometric redshift code. Although improved photometric redshift codes and better templates may remove the systematics and decrease the scatter, photometric redshifts remain limited by the low resolution spectral shape, and are not likely to surpass a scatter of \\dzs$=0.0$5 at the targeted redshift range -- even for the galaxies with deep NIR photometry. We note that obtaining photometry in narrower bands and for a larger number of filters may lead to better results than has typically been the case to date.  The spectroscopic stellar population properties, as derived by modeling the spectral continuum shape together with optical photometry, allow us to examine the significance and reliability of the stellar population properties derived from the photometry alone. Strikingly, the properties age, SFR, $\\tau$ show a scatter of a factor of$\\sim3-4$ between the photometrically and spectroscopically derived properties. For $A_V$ this scatter is $\\sim1$ mag. The improvement in age and $\\tau$ is driven primarily by the higher resolution spectral shape, and less so by the better constrained redshift. For SFR and $A_V$ both effects are about equally important. The uncertainties in the photometric modeling seem reasonable, as for $\\sim3/4$ of the galaxies the photometric and spectroscopic properties are consistent within 1$\\sigma$. Nevertheless, the photometric and spectroscopic distributions are quite different, accounting for their errors. This may be caused by the fact that age, $A_V$ and $\\tau$ suffer from systematics, due to degeneracies with the photometric redshift. For example, the photometric stellar population properties for several galaxies converge to the properties of best-fit templates, used in the photometric redshift code. This all implies that the best-fit photometric properties and especially their distribution are strongly biased.  The photometrically derived stellar mass and mass-to-light ($M/L_V$) ratio are best determined and show a scatter of a factor of $\\sim2$ with their spectroscopic analogues. However, the systematic offset in \\dzs\\ results in a systematic offset in $M_{\\rm phot}/M_{\\rm spec}$, such that $M_{\\rm phot}$ is about 30\\% too high for the galaxies with deep NIR photometry. This systematic is canceled out in $M/L_V$, as the $L_{V,\\rm phot}$ suffers from similar systematics as $M_{\\rm phot}$. The improvement for both stellar mass and $M/L_V$ is about equally driven by the higher resolution spectral shape and the better constrained redshift.  Rest-frame $V$ magnitudes and $U-V$ colors are the derived properties which are most sensitive to photometric redshift uncertainties and low S/N photometry.  The scatter between the photometrically and spectroscopically derived properties is 0.4 and 0.2 mag for $V$ and $U-V$ respectively for galaxies with deep NIR photometry, and both increase by 50\\% when we include the galaxies with shallow photometry. The systematics in the photometric redshifts cause a systematic offset of -0.2 and -0.1 mag for $V$ and $U-V$ for galaxies with deep NIR photometry. Thus, for both the scatter and the systematics, the improvement is almost completely driven by better constrained redshifts.  In order to assess the impact of the identified systematics on complete photometric samples, we compared photometric and spectroscopic samples in the same mass and redshift range. Surprisingly, we found that the distribution of $J-K$ and $R-K$ color barely changes for massive galaxies ($>10^{11}M_{\\odot}$) at $2.0<z<2.7$. Thus, the spectroscopy supports previous studies that red galaxies dominate the high mass end at $z\\sim2.5$. However, the combination of photometric redshift and stellar mass systematics also affect the number of massive galaxies. The photometric redshift code and template set used in this work overestimates the number of massive galaxies by factor of $\\sim$1.3 at $2<z<3$, resulting in an underestimation of the stellar mass density between $z\\sim2$ and $z\\sim0$.  Overall, our spectroscopic survey demonstrates the large uncertainties of best-fit photometric properties, and the necessity for accurate calibration of photometric studies. Although this study is a step forward in understanding the systematics in photometric studies of massive galaxies at $z>2$, larger spectroscopic samples over a larger redshift range are needed to fully map the systematics and accurately calibrate photometric studies. Furthermore, this study only applies to the high-mass end of the high-redshift galaxy population. Less massive galaxies may have other systematics, although that remains to be explored."
    },
    "0801/0801.0927_arXiv.txt": {
        "abstract": "We present deep $I$ and $z'$ imaging of the colour-selected cluster RzCS 052 and study the color-magnitude relation of this cluster, its scatter, the morphological distribution on the red sequence, the luminosity and stellar mass functions of red galaxies and the cluster blue fraction. We find that the stellar populations of early type galaxies in this cluster are uniformly old and that their luminosity function does not show any sign of evolution other than the passive evolution of their stellar populations. We rule out a significant contribution from mergers in the buildup of the red sequence of RzCS 052. The cluster has a large ($\\sim 30\\% $) blue fraction and and we infer that the evolution of the blue galaxies is faster than an exponentially declining star formation model and that these objects have probably experienced starburst episodes. Mergers are unlikely to be the driver of the observed colour evolution, because of the measured constancy of the mass function, as derived from near-infrared photometry of 32 clusters, including RzCS 052, presented in a related paper. Mechanisms with clustercentric radial dependent efficiencies are disfavored as well, because of the observed constant blue fraction with clustercentric distance. ",
        "introduction": "The existence of a tight, and apparently universal, color-magnitude relation for galaxies in nearby clusters (e.g. Bower, Lucey \\& Ellis 1992; Andreon 2003; Lopez-Cruz, Barkhouse \\& Yee 2004; McIntosh, Rix \\& Caldwell 2005; Eisenhardt et al. 2007 and references therein) implies that the majority of the stellar populations of early-type cluster galaxies were formed at  $z \\gg 1$ over relatively short timescales. Studies of clusters at high redshift, then, should allow us to witness the earlier stages of galaxy evolution, leading to the establishment of the present day luminosity function, color-magnitude relation and morphological mixtures. A classical example  of this kind of studies is the detection of a blueing trend among galaxies in clusters at $z > 0.3$ by Butcher \\& Oemler (1984).   Until large samples of high redshift ($z \\ga 0.8$) clusters become available, detailed `case studies' of individual objects may provide useful clues to the evolution of galaxy populations at half the Hubble time and beyond. Several studies have analyzed a number of such objects in detail (Blakeslee et al. 2003, 2006; Homeier et al. 2005, 2006; Holden et al. 2006; Mei et al. 2006a,b), using both ground-based imaging and spectroscopy and high-resolution imaging with the Advanced Camera for Surveys (ACS) on the Hubble Space Telescope (HST). These observations reiterate that the cluster early-type populations appear to be composed of old stellar populations which were probably in place at high redshift.  Most of the studied high redshift clusters are selected from the X-ray catalogs, which may pre-select objects that have already formed a deep potential well. An alternative strategy is to use clusters selected via the prominent red sequence of early type galaxies (Gladders \\& Yee 2000). Several $z\\geq 1$ clusters have already been identified using the galaxy colours or their spectral energy distributions (Andreon et al. 2005; Stanford et al. 2005).  Here we focus on RzCS 052 (J022143-0321.7), a rich (Abell class 2 or 3) cluster selected via a modified red sequence method, with a measured redshift of $1.016$ and a velocity dispersion of $710\\pm150$ km s$^{-1}$ and a modest X-ray luminosity of $(0.68 \\pm 0.47) \\times 10^{44}$ ergs s$^{-1}$ in the [1-4] keV band. Details about this objects and its properties may be found in Andreon et al. (2007). Because RzCS 052 is less X-ray luminous than clusters at similar redshift, and therefore does not possess a massive X-ray atmosphere, this object allows us to carry out a study of galaxy evolution in a different cluster environment and isolate the effects of gas on galaxy properties (Moran et al. 2007).  The layout of the paper is as follows. We present the data reduction and analysis in the next section. Section 3 discusses the red sequence galaxies. Section 4 deals with the blue galaxies in RzCS 052. Finally, we summarize our results in Section 5. We adopt the concordance cosmological parameters $\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$ and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. Magnitudes are quoted in their native photometric system (Vega for $RI$, SDSS for $z'$ and instrumental for Megacam data). Results of our stochastical computations are quoted in the form $x\\pm\\sigma$ where $x$ is the posterior mean and $\\sigma$ is the posterior standard deviation.   \\begin{figure} \\hbox{% \\psfig{figure=gal6.ps,width=2truecm,clip=}% \\psfig{figure=gal11.ps,width=2truecm,clip=}% \\psfig{figure=gal17.ps,width=2truecm,clip=}% } \\hbox{% \\psfig{figure=gal1.ps,width=2truecm,clip=}% \\psfig{figure=gal4.ps,width=2truecm,clip=}% \\psfig{figure=gal8.ps,width=2truecm,clip=}% \\psfig{figure=gal9.ps,width=2truecm,clip=}% } \\hbox{% \\psfig{figure=gal10.ps,width=2truecm,clip=}% \\psfig{figure=gal12.ps,width=2truecm,clip=}% \\psfig{figure=gal13.ps,width=2truecm,clip=}% \\psfig{figure=gal14.ps,width=2truecm,clip=}% } \\hbox{% \\psfig{figure=gal15.ps,width=2truecm,clip=}% \\psfig{figure=gal18.ps,width=2truecm,clip=}% } \\hbox{% \\psfig{figure=gal7.ps,width=2truecm,clip=}% \\psfig{figure=gal16.ps,width=2truecm,clip=}% \\hskip 2 truecm \\psfig{figure=gal5.ps,width=2truecm,clip=}% } \\caption[h]{Gallery of HST images of spectroscopically confirmed galaxies. The top row shows early-type galaxies. The second to fourth rows show red ($I-z'>-0.12$ mag) late-type galaxies. Last row shows blue ($I-z'<-0.12$ mag) late-type galaxies (two left-most panels) or unclassified galaxy (righmost panel). Each panel is $5''$ wide.} \\end{figure} ",
        "conclusions": "We presented new results on morphology, mass assembly history, role of environment and colour bimodality of galaxies in the colour selected, modest x-ray emitter, cluster RzCS 052 at $z=1.02$, as derived from VLT, HST, and  optical data, supplemented by a coarser analysis of a large sample of about 45 clusters, from $z=0$ to $z=1.22$ (16 in the context of the evolution of red galaxies shown in Fig. 4, and 32 used to constraint the fate of the blue population in Sec 4).  We found that the colour distribution of RzCS 052 is bimodal. Analysis of the morphological mix, slope and intercept of the colour-magnitude relation, and of the mass function shows that RzCS 052 red galaxies differs only by age from their local counterparts and that mergers play a minor role in building massive (down to~$M^*+2$) red galaxies in studied clusters, from $z=1$ to today.  The situation is remarkably different for blue galaxies. The blue fraction, once accounted for the younger age of stellar populations at high redshift and for the higher star formation rate there, is larger in RzCS 052 than in nearby similar clusters, highlighting perhaps for the first time that something, in addition to the flow of the time, is making galaxies bluer at high redshift.  Mergers are unlikely to be the driver of the observed colour evolution between $z=1$ and $z=0$, because of the measured constancy of the mass function, as derived from Spitzer photometry of 32 clusters in Andreon (2006). Mechanisms requiring a substantial intracluster medium, such as ram pressure stripping, are ruled out as well as direct driver, because of the very modest x-ray emission in RzCS 052. Mechanisms with a substantial different efficiency at the center and at one virial radius are strongly disfavored, because of the observed constant blue fraction."
    },
    "0801/0801.2113_arXiv.txt": {
        "abstract": "We study the evolution of dwarf ($L_H$ $<$ 10$^{9.6}$ L$_{H \\odot}$) star forming and quiescent galaxies in the Virgo cluster by comparing their UV to radio centimetric properties to the predictions of multizone chemo-spectrophotometric models of galaxy evolution especially tuned to take into account the perturbations induced by the interaction with the cluster intergalactic medium. Our models simulate one or multiple ram pressure stripping events and galaxy starvation. Models predict that all star forming dwarf galaxies entering the cluster for the first time loose most, if not all, of their atomic gas content, quenching on short time scales ($\\leq$ 150 Myr) their activity of star formation. These dwarf galaxies soon become red and quiescent, gas metal-rich objects with spectrophotometric and structural properties similar to those of dwarf ellipticals. Young, low luminosity, high surface brightness star forming galaxies such as late-type spirals and BCDs are probably the progenitors of relatively massive dwarf ellipticals, while it is likely that low surface brightness magellanic irregulars evolve into very low surface brightness quiescent objects hardly detectable in ground based imaging surveys. The small number of dwarf galaxies with physical properties intermediate between those of star forming and quiescent systems is consistent with a rapid ($<$ 1 Gyr) transitional phase between the two dwarf galaxies populations. These results, combined with statistical considerations, are consistent with the idea that most of the dwarf ellipticals dominating the faint end of the Virgo luminosity function were initially star forming systems, accreted by the cluster and stripped of their gas by one or subsequent ram pressure stripping events. ",
        "introduction": "Dwarf galaxies ($M_B>$-18) are the most common objects in the nearby universe (Ferguson \\& Binggeli 1994). Their importance resides in the fact that they represent in cold dark matter models the building blocks of hierarchical galaxy formation (e.g. White \\& Rees 1978, White \\& Frenk 1991). Their study is thus fundamental for constraining models of galaxy formation and evolution.\\\\ Observations of dwarf galaxies, necessarily limited to the nearby universe, revealed however a more complex origin than that predicted by models, making this class of galaxies even more interesting. Among dwarf galaxies, dwarf ellipticals (dE) and the less luminous dwarf spheroidals (dSph)\\footnote{Unless specified, in the following we indicate with dE all quiescent dwarf galaxies, including dSph} are more common than star forming Im and BCDs (Ferguson \\& Binggeli 1994). Firstly thought as the low-luminosity extension of bright ellipticals, several observational evidences indicate that they form an independent class of objects (Bender et al. 1992). Dwarf and giant ellipticals have Sersic light profiles of index $1/n$, with $n$ progressively increasing with luminosity (Graham \\& Guzman 2003; Gavazzi et al. 2005a). The colour magnitude relations measured using resolved stars in local group dwarfs spheroidals (Mateo 1998; Grebel 1999), the optical (Conselice et al. 2003a) and UV (Boselli et al. 2005a) integrated color magnitude relations, the subsolar [$\\alpha$/Fe] ratios (Van Zee et al. 2004a) and the UV to near-IR spectral energy distribution of dE in the Virgo cluster (Gavazzi et al. 2002a), however, all indicate a more gradual star formation history with recent episodes of activity in these low-luminosity quiescent systems than in massive ellipticals. This constitutes the first evidence against their very old origin. A further disagreement with model predictions is that, although more frequent than bright galaxies, their number density is still significantly smaller than predicted by hierarchical models of galaxy formation for the field luminosity function (Kauffmann et al. 1993; Cole et al. 1994; Somerville \\& Primack 1999, Nagashima et al. 2005) or for the local group (Klypin et al. 1999; Bullock et al. 2000).\\\\ The strong similarities in the structural properties of star forming and quiescent dwarf galaxies in the nearby universe, namely their similar optical morphology and light profiles (both roughly exponentials), suggested that dwarf ellipticals might result from gas removal and subsequent suppression of star formation in gas rich dwarf galaxies. Gas removal might result either from its blowout due to the kinetic energy injected into the ISM by supernova explosion following a strong burst of star formation (Dekel \\& Silk 1986, Vader 1986, Yoshii \\& Arimoto 1987), gas exhaustion through subsequent episodes of star formation (Davies \\& Phillipps 1988), or from external perturbations induced by the hostile environment in which galaxies evolve.  External perturbations include tidally induced mass loss in high speed encounters (Moore et al. 1998), tidal stirring (Mayer et al. 2001a, 2001b) and ram-pressure stripping induced by nearby companions (Lin \\& Faber 1983) or by the hot intergalactic medium in massive clusters (van Zee et al. 2004b). \\\\ Several observational evidences favor the environmental scenario against the gas blowout due to supernova explosions. The clearest indication that the environment plays an important role in the formation and evolution of dwarfs is the morphology segregation effect (Dressler 1980) which extends to low-luminosity systems (Binggeli et al. 1988; 1990; Ferguson \\& Binggeli 1994). Furthermore it has been recently claimed that the removal of the ISM through supernova winds is quite difficult in low-luminosity, dark matter dominated systems (Mac Low \\& Ferrara 1999; Ferrara \\& Tolstoy 2000; Silich \\& Tenorio-Tagle 1998, 2001). The discrete star formation history of several dwarf galaxies in the local group, as deduced by the analysis of their colour magnitude relation (Mateo 1998; Grebel 1999), is a further indication that dwarf spheroidals can retain their ISM through several episodes of star formation.\\\\ Another observational evidence favoring the transformation of star forming galaxies into quiescent dwarf ellipticals is the presence of rotationally supported (Pedraz et al. 2002; Geha et al. 2003; van Zee et al. 2004b) and/or HI gas rich (Conselice et al. 2003b; van Zee et al. 2004b) dE in the Virgo cluster. Recent studies based on SDSS imaging and spectroscopic data have shown that $\\sim$ 50 \\% of the bright end of the dE galaxy population in the Virgo cluster is characterized by disk features such as spiral arms or bars, this fraction decreasing down to $\\sim$ 5\\% at lower luminosities (Lisker et al. 2006a; see also Graham et al. 2003). Meanwhile $\\sim$ 15\\% of the bright dE in Virgo have blue centers revealing a recent activity of star formation. From a statistical point of view, the line-of-sight velocity distribution of dE inside clusters is similar to that of late-type galaxies suggesting a recent infall (Binggeli et al. 1993; Conselice et al. 2001).\\\\ Not all observational evidences, however, are consistent with the transformation of dwarf star forming galaxies into dwarf ellipticals under the effect of the environment. To reproduce the color magnitude relation of dwarf ellipticals, it has been claimed that magellanic irregulars should have faded $\\sim$ 1.5 magnitudes in the B band thus reaching surface brightnesses weaker than $\\mu(B)_e$ $=$ 25 mag arcsec$^{-2}$, values significantly smaller than the observed ones (Bothun et al. 1986). Most of the bright dwarf ellipticals in the Virgo cluster have a bright nucleus (Binggeli et al. 1985; Ferguson \\& Binggeli 1994) of small size, as shown by HST observations ($\\sim$ 4 pc, C\\^ote et al. 2006) while dwarf irregulars do not. While both dwarf irregulars and dwarf ellipticals follow different metallicity-luminosity relations, dwarf spheroidals are more metal-rich than dwarf irregulars at the same optical luminosity (Grebel et al. 2003 and references therein). The flattening distribution of non nucleated dE is similar to that of late-type spirals, Im and BCD (Binggeli \\& Popescu 1995). This class of objects, however, is significantly less round than nucleated systems. Although they exist, the fraction of objects belonging to the intermediate dIrr/dE transition class is too small. Furthermore a simple transformation of Im galaxies recently infalling in the cluster into dE does not seem to reproduce the observed difference in the cluster and field luminosity functions (Conselice 2002). Other strong constraints against the transformation of star forming into quiescent dwarfs are given by the studies of globular clusters, taken as a probe of the early phases of galaxy formation. The specific frequency of globular clusters in dwarf ellipticals, in fact, is significantly higher than that of star forming galaxies. This result has been interpreted as an evidence for a different formation scenario for the two galaxy populations (Miller et al. 1998; Strader et al. 2006).\\\\ With the aim of explaining these evidences, slightly different evolutionary scenarios have been proposed: cluster dwarf ellipticals might have been formed from higher surface brightness BCD galaxies (Bothun et al. 1986), or from tidally induced mass loss in multiple high speed encounters of massive galaxies (galaxy harassment; Mastropietro et al. 2005). Alternatively, Lisker et al. (2006a; 2006b; 2007) proposed that the cluster dwarf elliptical galaxy population is composed of different subcategories of objects, of which not all have been formed from gas stripped star forming dwarfs. As here emphasized the issue is still hotly debated.\\\\ A few years ago we started collecting multifrequency data covering the whole electromagnetic spectrum for galaxies in nearby clusters in order to study the effects of the environment on galaxy evolution. Up to now our research was primarily focused on the bright end of the luminosity function. Our interest covered the present and past star formation activity (Gavazzi et al. 1991, 1998, 2002b, 2006a), the atomic (Gavazzi et al. 2005b, 2006b) and molecular gas content (Boselli 1994; Boselli et al. 1994, 1997a, 2002), the radio continuum and IR properties of cluster galaxies (Gavazzi \\& Boselli 1999; Gavazzi et al. 1991). The results of our analysis, combined with those obtained by other teams, are summarized in a recent review article (Boselli \\& Gavazzi 2006).\\\\ The aim of the present paper is to extend the multifrequency analysis to the low-luminosity end of the luminosity function with the purpose of studying the possible transformation of star forming objects into dwarf ellipticals. This exercise is here done using a complete sample of galaxies in the Virgo cluster. The novelty of this work compared to previous investigations is twofold: the analyzed sample has an unprecedented multifrequency spectral coverage from the UV to near-IR imaging data. Furthermore this unique dataset is compared to the predictions of multizone chemo-spectrophotometric models of galaxy evolution here adapted to take into account the perturbations induced by the cluster environment. The type of perturbation simulated by the models are ram-pressure stripping (Gunn \\& Gott 1972) and starvation (Larson et al. 1980).  The combination of multifrequency data with these models on the galaxies NGC 4569 (Boselli et al. 2006) and NGC 4438 (Boselli et al. 2005b) indeed indicated how powerful this method is for studying and constraining the evolution of cluster galaxies. \\\\ Our previous investigations have emphasized that the most important parameter governing the evolution is the total mass as traced by the near-IR H band luminosity (Gavazzi et al. 1996; 2002a; Boselli et al. 2001). Below a certain mass ($L_H$ $<$ 10$^{9.6}$ L$_{H \\odot}$) we identify the sequence of dwarf galaxies that we subdivide into quiescent and star forming disregarding their detailed morphology.\\\\ A major uncertainty in the interpretation of the multifrequency data, and in particular of those at short wavelength, is the extinction correction, which however is expected to be minor in low-metallicity star forming dwarf systems (Buat \\& Xu 1996) and probably negligible in dwarf ellipticals. For these reasons we decided to limit the present analysis to the cluster dwarf galaxy population, leaving the discussion relative to massive galaxies to a future communication. With the aim of driving the reader's attention to the scientific results of this work, the presentation of the dataset and the general description of the models are given in Appendix. ",
        "conclusions": "We can conclude that models and observations are consistent with an evolution of star forming, low-luminosity late type galaxies recently accreted in Virgo into quiescent dwarfs because of the ram pressure gas stripping and the subsequent stopping of their star formation activity. For consistency with surface brightness measurements, we show that high surface brightness star forming dwarf galaxies (low luminosity spirals and BCD) might be at the origin of the optically selected dwarf ellipticals (both normal and nucleated) analyzed in this work, although we expect that the low surface brightness magellanic irregulars (Sm and Im) once stripped of their gas reservoir, produce quiescent dwarfs with surface brightnesses below the detection limit of the VCC, as those observed in Virgo by Sabatini et al. (2005). The process of transformation is extremely rapid and efficient, since it works on all dwarfs and last on average less than 150 Myr. On longer time scales galaxies get structural and spectrophotometric properties similar to that of dwarf ellipticals. The whole star forming dwarf galaxy population dominating the faint end of the field luminosity function (Blanton et al. 2005), if accreted, can be totally transformed by the cluster environment into dwarf ellipticals on time scales as short as 2 Gyr and thus be at the origin of the morphology segregation observed also at low luminosities.\\\\ This interesting result is of fundamental importance even in a cosmological context because it shows that the majority of dwarf galaxies (if not all of them) are ``young'' even in clusters, and not old as expected in a hierarchical galaxy formation scenario (see also Nelan et al. 2005). For comparison with De Lucia et al (2006), we give in Table \\ref{Tabage} the lookback time when 50\\% and 80\\% of the stars were formed for a model galaxy of $V_C$ $=$ 55 km s$^{-1}$ for a (or two) ram pressure stripping event of efficiency $\\epsilon_0$ = 1.2 M$\\odot$ kpc$^{-2}$ yr$^{-1}$.   \\noindent Although the values given in Table \\ref{Tabage} are not directly comparable to those given by De Lucia et al. (2006) since limited to relatively high mass objects (2.5 10$^9$ M$\\odot$, while our model galaxy is of only 2.4 10$^8$ M$\\odot$), the mean age of the stellar population of low luminosity, quiescent galaxies is significantly younger than that predicted by hierarchical models of galaxy formation (lookback times $\\sim$ 10 and 8.5 Gyr for 50\\% and 80\\% of the stellar population).\\\\ Despite their morphological type, the star formation activity of dwarf systems has been abruptly interrupted by the interaction with the environment in quiescent systems. This conclusion can be extended to the local group, where the study of the stellar color magnitude relation of dwarf spheroidal systems revealed that the star formation activity, although in an episodic manner, lasted for several Gyr (Mateo 1998; Grebel 1999). We can add that the only cluster galaxy population likely to be issued by major merging events, as those predicted by hierarchical models of galaxy formation, is that of massive ellipticals, whose origin is probably very remote ($z$ $\\geq$ 2-3, Dressler 2004; Treu 2004; Nolan 2004; Franx 2004; Renzini 2006). Indeed this is the only Virgo cluster galaxy population with a virialized velocity distribution (Conselice et al. 2001). If we consider clusters of galaxies as those regions where merging events were more frequent at early epochs since the big bang just because characterized by an high galaxy density, we can conclude that the hierarchical formation scenario was the principal driver of galaxy evolution only in massive objects at very early epochs. All observational evidences are consistent with a secular evolution afterward."
    },
    "0801/0801.0266_arXiv.txt": {
        "abstract": "We study the alignment of grains subject to both radiative torques and pinwheel torques while accounting for thermal flipping of grains. By pinwheel torques we refer to all systematic torques that are fixed in grain body axes, including the radiative torques arising from scattering and absorption of isotropic radiation. We discuss new types of pinwheel torques, which are systematic torques arising from infrared emission and torques arising from the interaction of grains with ions and electrons in hot plasma. We show that both types of torques are long-lived, i.e. may exist longer than gaseous damping time. We compare these torques with the torques introduced by E. Purcell, namely, torques due to H$_2$ formation, the variation of accommodation coefficient for gaseous collisions and photoelectric emission. Furthermore, we revise the Lazarian \\& Draine model for grain thermal flipping. We calculate mean flipping timescale induced by Barnett and nuclear relaxation for both paramagnetic and superparamagnetic grains, in the presence of stochastic torques associated  with pinwheel torques, e.g. the stochastic torques arising from H$_2$ formation, and gas bombardment. We show that the combined effect of internal relaxation and stochastic torques can result in fast flipping for sufficiently small grains and, because of this, they get thermally trapped, i.e. rotate thermally in spite of the presence of pinwheel torques. For sufficiently large grains, we show that the pinwheel torques can increase the degree of grain alignment achievable with the radiative torques by increasing the magnitude of the angular momentum of low attractor points and/or by driving grains to new high attractor points. ",
        "introduction": "Polarization of radiation arising from emission or absorption by aligned grains is widely used to study magnetic fields in the diffuse interstellar medium (see Goodman et al. 1995), molecular clouds (see Hildebrand et al. 2000, 2002) and in prestellar cores (Crutcher et al. 2004). The grain alignment involves the alignment of the axis of major inertia $\\ba_{1}$ with angular momentum $\\bJ$ and the alignment of $\\bJ$ with respect to the ambient magnetic field.  An understanding of grain dynamics is extremely important for both understanding of dust properties and dust alignment. The latter is essential for determining situations when polarization of starlight passing through dusty interstellar gas, as well as polarization of dust emission can be reliably interpreted in terms of magnetic field direction. In terms of dust properties, fast rotation should disrupt loose aggregates, placing limits for the fractal dimensions of dust particles.  As dust scatters or absorbs photons, it experiences radiative torques. These torques can be stochastic or systematic. Stochastic radiative torques arise from, for instance, a spheroidal grain randomly emitting or absorbing photons. The latter process, for instance, was invoked by Harwit (1970) in his model of grain alignment based on grains being preferentially spun up in the direction perpendicular to the photon beam. However, Purcell \\& Spitzer (1971) showed that randomization arising from the same grain emitting thermal photons makes the achievable degree of grain alignment negligible. In fact, they showed that random radiative torques arising from grain thermal emission is an important process of grain randomization. More recently, grain thermal emission was analyzed as a source of the excitation of grain rotation as well as its damping in relation to the rotation of tiny spinning grains that are likely to be responsible for the so-called anomalous foreground emission (Draine \\& Lazarian 1998).  Systematic radiative torques were first introduced by Dolginov (1972) in terms of chiral, e.g. hypothetical quartz grains. Starlight passing through such grains would spin them up. Later, Dolginov \\& Mytrophanov (1976) considered an irregular grain model, consisting of two twisted spheroids, made of more accepted materials, e.g. silicate grains, and claimed that these grains will be both spun up and aligned by radiative torques arising from the anisotropic component of radiation field.\\footnote{Using so-called  Rayleigh-Gans approximation, Dolginov \\& Silantiev (1976) provided calculations of the radiative torques for a model of an irregular grain consisting of two ellipsoids twisted with respect to each other, but our calculations of radiative torques using the Discrete Dipole Approximation code (Draine \\& Flatau 2004) for the same set parameters and the same model as in the aforementioned work are in conflict with their analytical findings (see more in Hoang \\& Lazarian 2008b).} Further on, following the convention we adopted in Lazarian \\& Hoang (2007a, hereafter LH07a), we shall call these torques RATs. The work by Dolginov \\& Mytraphanov (1976) was, unfortunately, mostly ignored for 20 years. A possible explanation of this may be due to the fact, that for a long time, the magnitude of RATs for realistic irregular grains remained unclear.  A renewed interest to RATs was induced by the possibility of calculating them for arbitrary grain shapes. This occurred after Bruce Draine modified correspondingly his publicly-available Discrete Dipole Approximation code (hereafter DDSCAT; Draine \\& Flatau 2004). Moreover, Draine \\& Weingartner (1996, 1997, henceforth DW96, DW97, respectively) conjectured that RATs may provide the primary alignment mechanism for interstellar grains. Support for this claim came through later research in the field, by better understanding the dynamics of grains subject to RATs from anisotropic radiation fields, calculating grain alignment with an analytical model (LH07a), as well as including important physical processes like grain wobbling (Lazarian 1994; Lazarian \\& Roberge 1997; Weingartner \\& Draine 2003; Hoang \\& Lazarian 2008a), and accounting within the model for gaseous bombardment and uncompensated torques arising from H$_2$ formation, as was done in Hoang \\& Lazarian (2008a, henceforth HL08a). As it was described in a review by Lazarian (2007), RAT alignment mechanism has become not only the leading candidate to explain interstellar grain alignment, but also to explain grain alignment in many other astrophysical environments, including circumstellar regions (Aitken et al. 2002), accretion disks (Cho \\& Lazarian 2007), comet atmospheres (see Rosenbush et al 2007), and molecular clouds (Whittet et al. 2008).  In LH07a, we subjected to scrutiny the properties of RATs. Using a simple analytical model (AMO) of a helical grain we studied the properties of RATs and the RAT alignment. The results obtained by the AMO were shown to be in good correspondence with numerical calculations obtained by DDSCAT for irregular grains. Invoking the generic properties of the RAT components, we explained the RAT alignment of grains in both the absence and presence of magnetic fields. Intentionally, for the sake of simplicity, in LH07a we studied a simplified dynamical model to demonstrate the effect of RATs. Within the latter model we showed that RATs can align grains at attractor points with low angular momentum (i.e., $J$ of the order of thermal angular momentum $J_{th}$, hereafter, low-$J$ attractor points), and/or attractor points with high angular momentum (i.e., $J\\gg J_{th}$, hereafter high-$J$ attractor points). The high-J attractor points mostly correspond to a perfect alignment of $\\bJ$ with respect to $\\bB$, while the low-J attractor points occur at perfect alignment angle or at some angle in the vicinity of the perfect alignment.   One of the effects that was not discussed in LH07a was the effect of thermal fluctuations.  Lazarian (1994) noticed that in spite of the fast rates of internal relaxation, rotating grains wobble. This conclusion was a consequence of the Fluctuation-Dissipation Theorem (see Landau \\& Lifshitz 1976), which states that any dissipation process, e.g. internal relaxation, should be accompanied by the proportionally fluctuation process. It is possible to show that grains rotating thermally are expected to wobble with larger amplitude compared to their counterparts rotating at suprathermal (much greater than thermal) rates (see Lazarian 1994, Lazarian \\& Roberge 1997). Therefore the zero value of angular momentum at low-J attractor points obtained in LH07a stemmed from the assumption of the perfect coupling of grain axis of the maximal moment of inertia $\\ba_{1}$ with $\\bJ$, which is violated for sufficiently low $J$. In fact, in HL08a, we treated fully the dynamics of the grain alignment by taking into account thermal fluctuations (see also Weingartner \\& Draine 2003), and showed that thermal fluctuations within irregular grains can increase the angular momentum of the low-J attractor points from $J=0$ to $J\\sim J_{d}$, i.e. the angular momentum of grains corresponding to dust grain temperature. In addition, HL08a showed that gas bombardment and other randomizing torques can contribute significantly to the RAT alignment by moving grains from low-$J$ to high-$J$ attractor points, whenever the latter are present. However, we feel that a comprehensive study on the degree of alignment as a function of grain size and radiation intensity is very essential for polarization modeling. Note that earlier papers which attempted to introduce grain alignment into polarization modeling (Cho \\& Lazarian 2005; Pelkonen et al. 2007; Bethell et al. 2007) assumed  perfect alignment for grains larger than some critical size, and no alignment for smaller grains. The latter was inferred using the criterion of whether RATs for radiation field under interest, which are  sharp function of grain size, can induce grains to rotate several times faster than the thermal rotation rate. This is a rather crude criterion because it does not take into account the type of RAT alignment, i.e with or without high-J attractor points. For instance, if RATs align all grains at low-$J$ attractor points, then  assumption above leads to overestimates for the degree of alignment.  In addition to RATs, other uncompensated torques act on grains (Purcell 1979; LD99a; Roberge \\& Ford 2000, henceforth RF00). Purcell (1979) proposed three processes that can produce uncompensated torques: hydrogen formation, photoelectric effect and the variation of accommodation coefficient. These pinwheel torques together with the paramagnetic dissipation were thought to be the major mechanism leading to the alignment of $\\bJ$ with $\\bB$. However, the paramagnetic dissipation was found to be unable to explain the alignment of paramagnetic grains in the weak interstellar magnetic field. In any case, in the presence of RATs, the alignment arising from paramagnetic dissipation for ordinary paramagnetic grains is negligible compared to that arising from RATs\\footnote{Unfortunately, the misconception that RATs act as proxies of Purcell's torques and the alignment arises from paramagnetic dissipation is well entrenched in the grain alignment literature. This claim is based on the DW96 study and, in spite of the fact, that the efficient RAT alignment was demonstrated already in DW97 and later works, the claim persisted. This may be due to the fact that, DW97 did not provide, unlike DW96, the predictions for the degree of alignment. The problem of predicting the degree of alignment happened to be a tough one. The present paper is an attempt in this direction. Interestingly enough, Dolginov \\& Mytrophanov (1976) did consider RATs as a mechanism capable of aligning grains irrespectively of the presence or absence of paramagnetic relaxation. For superparamagnetic grains with sufficient number of iron atom per cluster, Lazarian \\& Hoang (2008) showed that the alignment always happen with high-$J$ attractor points and is perfect. However, even in this case, RATs are more important than the enhanced paramagnetic relaxation. The latter effect causes the stabilization of the high-$J$ attractor point created by RATs.}.  The efficiencies of the Purcell torques decrease if grains wobble thermally as we discussed above. The effect of thermal wobbling was quantified in Lazarian \\& Roberge (1997) and the results of this study were used by Lazarian \\& Draine (1999ab, henceforth LD99ab) to predict new effects of grain dynamics, namely, thermal flipping and thermal trapping. The thermal flipping is the effect of occasional increase of the wobbling angle beyond 90 degrees. When this occurs, the torques that are fixed in grain body coordinates change their direction. In the case thermal flipping occurs fast enough, the Purcell torques get averaged out and the grain rotates thermally in spite of the presence of the uncompensated pinwheel torques, i.e it is thermally trapped.  However, this picture was challenged recently by Weingartner (2008). His study has several new mathematical points. The most important one is that he used a dimensionless variable $q=2I_{1}E/J^{2}$ with $I_{1}$ being the inertia moment along the axis of major inertia $\\ba_{1}$ and $E$ being the total rotational energy to describe the internal relaxation, instead of the angle $\\theta$ between $\\ba_{1}$ and $\\bJ$. He suggested an integration constant for diffusion coefficient of internal relaxation, which differs from that in the Lazarian \\& Roberge (1997). We agree with the choice of the integration constant. For this choice in Weingartner (2008) the diffusion coefficient vanishes when  $\\ba_{1}$ perpendicular to angular momentum $\\bJ$. As a result, he found that grains do not experience thermal flipping as a result of internal relaxation. We agree with this choice of the integration constant as it corresponds to the nature of the Fluctuation-Dissipation Theorem employed in Lazarian (1994) to describe the wobbling of the grains. Indeed, according to the theorem no dissipation should correspond to no fluctuations. For an oblate grain the internal dissipation goes to zero as $\\theta \\rightarrow \\pi/2$. Thus the Fluctuation-Dissipation Theorem suggests that the at this point the fluctuations should also go to zero. One can check that the choice of the constant in Weingartner (2008), which was obtained from more formal considerations, corresponds to this physical requirement. However, as it clear from the rest of the paper, we disagree that this modification of the diffusion coefficients will preclude physical processes of thermal flipping and trapping introduced in Lazarian \\& Draine (1999) from happening.  In what follows, we consider a more realistic picture of suprathermal rotation, namely, we take into account  the fact that any pinwheel torque is accompanied by stochastic torques of the same nature. For instance, apart from systematic torques, the process of H$_2$ formation induces stochastic torques. In this model it is clear that the diffusion is present even when the angle between ${\\bf a_1}$ and ${\\bf J}$ is 90 degrees, which means that the thermal flipping is possible. In this paper we revisit the LD99a study and confirm that the thermal flipping and thermal trapping predicted in LD99a are important effects of grain dynamics.  The structure of the present paper is as followings. In \\S 2 we present pinwheel torques, including long-lived and short-lived torques. We calculate the maximal value of angular momentum that these pinwheel torques can achieve for the ISM. In \\S 3 we revise the problem of thermal flipping and calculate the thermal mean flipping time induced by internal relaxation in the presence of external stochastic torques. Critical size of flipping grain and trapping size for the ISM are estimated in this section. We study the effects of H$_{2}$ pinwheel torques and thermal flipping by stochastic torques accompanied with H$_{2}$ formation on the RAT alignment for the diffuse ISM in \\S 4. We also calculate the value of $J$ at low attractor points as a function of the magnitude of H$_{2}$ pinwheel torques for different grain sizes in this section. We summarize our findings in \\S 5. ",
        "conclusions": "\\subsection{Grains are flipping and get trapped}  New elements of grain dynamics, namely, thermal flipping and thermal trapping were reported in LD99a. The validity of these findings was challenged in Weingartner (2008), who found that grains can not experience thermal flipping as a result of internal relaxation. In the present work we showed that although internal relaxation can not result in grain thermal flipping, but in the presence of impulses due to H$_{2}$ formation and gas bombardment, grains can flip, and get trapped. The impulses play the role of an engine to bring grains through the boundary $\\ba_{1}\\perp\\bJ$.  The thermal trapping is an essential process that, according to LD99a, explains why, in accordance with observations (see Kim \\& Martin 1995), small grains are not aligned by the Purcell (1979) paramagnetic relaxation process. The alternative explanations to this observational fact (see Lazarian 1995) are much less appealing. The sizes of the thermally trapped grains are not different from those which can be obtained using the LD99a treatment, but smaller than those discussed in LD99b. As a result, pinwheel torques are more important for interstellar grains than one could infer from LD99b.  \\subsection{Extending theory of crossovers}  Crossovers are times when grains subject to pinwheel torques slow down, then flip and get accelerating. The original theory of crossovers was suggested by Spitzer \\& McGlynn (1979). Lazarian \\& Draine (1997) improved the theory by accounting for the value of the residual thermal angular momentum, associated with thermal wobbling of the grain (see Lazarian 1994). LD99b showed that with nuclear relaxation taken into account the theory for regular crossovers in the spirit of the aforementioned works is applicable only for grains larger than $\\sim 10^{-4}$~cm, which excludes most of grains in diffuse interstellar gas. LD99a, instead, introduced the concepts of thermal flipping and trapping for smaller grains.  The present paper shows that the flipping and trapped grains are smaller than those discussed in earlier papers. Therefore, in the absence of RATs, there is a range of grains for which the crossovers happen over the time scale larger than the internal relaxation time. In this situation the thermal value of the angular momentum can decrease below  $J_d$ a result of the action of the pinwheel torques. This increases the randomization of grains during crossovers for the range of sizes larger than the flipping size.  \\subsection{How common is suprathermal rotation?}  Suprathermal, i.e. much faster than thermal, rotation was assumed to be default for most of the interstellar grains after the Purcell (1979) study. Studies of RATs in DW96, DW97 only made this point stronger. However, later works that reported thermal trapping of grains (LD99a, LD99b) and low-J attractor points (Weingartner \\& Draine 2003, LH07a) made one wonder whether most grains rotate thermally. The study in HL08a revealed that in the presence of gaseous bombardment or other randomization processes, the grains tend to diffuse from the low-J to high-J attractor point, provided that such high-J attractor point coexist with the low-J one. According to the parameter space study in LH07a (see Fig. 24 in LH07a) this corresponded to a fraction of grains that could vary depending on the angle between the radiation direction and the magnetic field, as well as on the $q^{max}$ ratio.  In the present study we showed that for sufficiently large grains the pinwheel torques, provided that they are comparable with or stronger than the RATs, can induce suprathermal rotation of grains. Interestingly enough, this can happen even in the situations when only low-$J$ attractor points exist. Therefore, the presence of pinwheel torques can increase the percentage of the grains rotating suprathermally. However, it is worth noting, that for a substantial portion of the parameter space radiative torques act against pinwheel torques attempting to decrease the grain rotation rate.    \\subsection{Implications to polarization modeling and magnetic field diagnostics}  A quantitative theory of grain alignment is very important to polarization simulations. A number of works dealing with polarization simulations assume the perfect alignment of the grain axis with respect to the magnetic field (see CL07; Bethell et al. 2007; Pelkonen et al. 2007; Falceta-Goncalves, Lazarian \\& Kowal 2008). A detailed comparison between the polarization efficiency predicted by the theory of grain alignment and observations of dense cloud in Whittet et al. (2008) shows that a high degree of alignment is required to explain the observed polarization degree. How can this be accomplished?  In Lazarian \\& Hoang (2008) we proposed one way of increasing the degree of grain alignment, namely, we noticed that in the situation when grains have superparamagnetic inclusions the high-$J$ attractor points are present for all the angles between the radiation direction and the magnetic field, as well as on the $q^{max}$ ratio. Thus the gaseous bombardment is expected to move all grains to the high attractor points, making the alignment perfect. In the present paper we show an alternative way of increasing the degree of grain alignment. In particular, we show that RATs plus pinwheel torques can enhance the degree of alignment by increasing the value of angular momentum at the low-$J$ attractor point. However, the circumstances for which the pinwheel torques affect the RAT alignment are rather restrictive: the pinwheel torques are required to be comparable or stronger than RATs.  With these theoretical results one can attempt to distinguish between the two alternatives above and simultaneously get insight into the grain composition. First of all, more studies similar to those done by Whittet et al. (2008) are required to make sure that the high alignment is not a result of favorable illumination geometry. Alternatively, one may make measurements in situations where the illumination geometry is known. Comparing these results with the analytical and numerical predictions (see LH07a, HL08a) one can establish with higher certainty that the grain alignment is enhanced compared to what is expected from ordinary paramagnetic grains subject to RATs in the absence of the pinwheel torques. In parallel, one can search for the correlation of grain alignment enhancement with the environments where we expect the higher amplitudes of the pinwheel torques. If such an enhancement is present, it will indicate that grains are {\\it not} superparamagnetic. If they were superparamagnetic, we would expect the perfect alignment for typical conditions of the diffuse ISM, which one cannot be improved further by the action of the pinwheel torques. No correlation, but high degrees of alignment mean, on the contrary, the existence of superparamagnetic grains.  One can conclude by observing an interesting cycle in the development of the grain alignment theory. The initial theory by Davis-Greenstein (1951) appealed to paramagnetic dissipation in ordinary paramagnetic grains. The mechanism was shown not to be sufficiently strong, however. Then to enhance the efficiency of the initial Davis-Greenstein process both superparamagnetism  (Jones \\& Spitzer 1967) and pinwheel torques (Purcell 1979) were appealed to. At the moment we may face the situation that the RATs are not strong enough to explain the alignment and again we appeal to superparamagnetism (Lazarian \\& Hoang 2008) and pinwheel torques (this paper). Nevertheless, there is a substantial difference between the two situations. The paramagnetic relaxation in the diffuse interstellar medium could provide a few percent alignment at most. The RAT alignment provides the degrees of alignment about 20\\% even in the cases of unfavorable directions of illumination (see HL08a). In addition, while the Davis-Greenstein process favored the alignment of small grains, the RAT alignment, in accordance with observations, favors the alignment of large grains.   \\subsection{Summary} Our main results are summarized as follows: \\begin{itemize}  \\item We extended the definition of the pinwheel torques to include three new processes, namely, radiative torques arising from grain infrared emission, torques due to interactions of grain with electrons in plasma, as well as, torques arising from gas flow interacting with a helical grain. We proposed to consider the radiative torques arising from the action of isotroptic radiation on a grain as pinwheel torques. All these torques we identified as long-lived, as their life-time is expected to be longer than the typical grain rotational damping time. We used H$_2$ torques as a proxy of for general pinwheel torques, but, unlike earlier studies, we allowed the pinwheel torques to be both short-lived and long-lived. We found that in the ISM, the pinwheel torques due to infrared emission is in general weaker than that due to H$_{2}$ formation, and RATs.  \\item We proved that flipping and trapping of astrophysical dust grains is the active and important process. We did this by augmenting the torques arising from internal relaxation by the stochastic torques that inevitably accompany the action of pinwheel torques and gas bombardment. . We calculated the thermal flipping induced by the Barnett, nuclear relaxations, as well, as internal relaxation in superparamagnetic grains in the presence of impulses due to H$_{2}$ formation and gas bombardment. Adopting the revised diffusion coefficients for internal relaxation from Weingartner (2008), we found that flipping is fast for small grains, and increases when superparamagnetic inclusions are accounted for. We obtained the critical size of flipping and trapping size of the thermally trapped grains for the ISM.  \\item The most important result of this study is related to the increase of the expected degree of grain alignment when dust grains are subject to both RATs and pinwheel torques. This increase stems from the increase by pinwheel torques of the value of the angular momentum at the low-$J$ attractor point. The increase of angular momentum decreases the wobbling of grains at the low-$J$ attractor point and therefore increase the degree of internal alignment as the internal alignment (of $\\ba_{1}$ with $\\bJ$) is nearly perfect when the grain rotates with rates much larger than the thermal rotation rate. The joint pinwheel plus RATs alignment is different from the Purcell (1979) alignment. In the latter mechanism the alignment is due to paramagnetic dissipation, which induces torques that are very weak compared to RATs. \\end{itemize}"
    },
    "0801/0801.2749_arXiv.txt": {
        "abstract": "We present Very Long Baseline Array (VLBA) observations of the TeV blazars H\\,1426+428, 1ES\\,1959+650, and PKS\\,2155$-$304 obtained during the years 2001 through 2004.  We observed H\\,1426+428 at four epochs at 8\\,GHz, and found that its parsec-scale structure consisted of a $\\sim 17$\\,mJy core and a single $\\sim 3$\\,mJy jet component with an apparent speed of $2.09c\\pm0.53c$.  The blazar 1ES\\,1959+650 was observed at three epochs at frequencies of 15 and 22\\,GHz.  Spectral index information from these dual-frequency observations was used to definitively identify the core of the parsec-scale structure. PKS\\,2155$-$304 was observed at a single epoch at 15\\,GHz with dual-circular polarization, and we present the first VLBI polarimetry image of this source. For 1ES\\,1959+650 and PKS\\,2155$-$304, the current observations are combined with the VLBA observations from our earlier paper to yield improved apparent speed measurements for these sources with greatly reduced measurement errors.  The new apparent speed measured for component C2 in 1ES\\,1959+650 is $0.00c\\pm0.04c$ (stationary), and the new apparent speed measured for component C1 in PKS\\,2155$-$304 is $0.93c\\pm0.31c$. We combine the new apparent speed measurements from this paper with the apparent speeds measured in TeV blazar jets from our earlier papers to form a current set of apparent speed measurements in TeV HBLs. The mean peak apparent pattern speed in the jets of the TeV HBLs is about 1$c$. We conclude the paper with a detailed discussion of the interpretation of the collected VLBA data on TeV blazars in the context of current theoretical models for the parsec-scale structure of TeV blazar jets. ",
        "introduction": "\\label{intro}  The field of TeV gamma-ray astronomy has grown rapidly over the past several years with the advent of sensitive new gamma-ray telescopes such as H.E.S.S.\\ and MAGIC.  Particularly interesting results have been obtained in blazar astronomy, as the total number of detected TeV blazars has increased from six only a few years ago to 17 this year (Wagner 2007b).  The majority of the detected TeV blazars (16 of 17) belong to the class of high-frequency peaked BL Lac objects, or HBLs --- so named because the two peaks in their spectral energy distribution (SED) occur at relatively high UV/X-ray and GeV/TeV energies. This two-peaked spectral energy distribution in HBLs is most commonly interpreted as the result of relativistic electrons (and possibly positrons) radiating in a jet which is undergoing bulk relativistic motion at a small angle to the observer's line of sight, with the low-frequency peak due to synchrotron radiation, and the high-frequency peak due to inverse-Compton scattering of the jet's own synchrotron photons (synchrotron self-Compton, or SSC emission).  The TeV source list is mostly limited to relatively nearby blazars ($z\\lesssim0.2$), because of the absorption of TeV gamma-rays on the extragalactic background light.  The TeV observations have shown dramatic variability in a number of these blazars. Among the most remarkable variability events are the 200\\,s variability timescale detected for PKS 2155$-$304 in July 2006 by H.E.S.S.\\ (Aharonian et al.\\ 2007), and the 3-minute variations detected for Mkn\\,501 in June 2005 by MAGIC (Wagner 2007a).  Such rapid variations suggest extremely small emitting volumes and/or time compression by large relativistic Doppler factors, e.g., $\\delta\\gtrsim100$ for PKS 2155$-$304 (Aharonian et al.\\ 2007). High Doppler factors are also sometimes invoked in specific SSC models of TeV blazar spectra and variability; for example, Fossati et al. (2007) consider two SSC models to explain the X-ray/TeV variability of Mkn 421: one with $\\delta\\sim20$ (scattering in the Klein-Nishina regime), and one with $\\delta\\sim100$ (scattering in the Thomson regime). Such high Doppler factors are at the upper limit of what is expected in relativistic jets, and challenge our understanding of these objects if they do in fact occur.  Complementary observations that are crucial to unraveling the physics of TeV blazar jets are provided by the VLBI technique, which yields radio images of the relativistic jets with sub-parsec resolutions for these nearby blazars.  Apparent jet speeds, brightness temperatures, and limits on jet/counterjet brightness ratios can all be measured from VLBI images, and these quantities all provide constraints on fundamental jet parameters such as the bulk Lorentz factor and the angle of the jet to the line-of-sight (subject to some caveats discussed at length in $\\S$~\\ref{disc}).  Our understanding of the parsec-scale radio properties of the HBL class in general has been increased by the recent work of Giroletti et al.\\ (2004a, 2006). Those authors studied a sample of low-redshift BL Lac objects with a variety of radio instruments, and demonstrated that the radio properties of the HBLs are consistent with them being the beamed versions of nearby low-luminosity FR I radio galaxies.  The TeV blazars thus likely have an intrinsically different parent population with weaker jets, compared to the more distant powerful blazars, which likely have FR II parents (e.g., Urry \\& Padovani 1995).  The mean derived parsec-scale Lorentz factor for the HBL class by Giroletti et al.\\ (2004a), including TeV sources, is only $<\\Gamma>\\sim3$, much lower than the Lorentz factors suggested by the TeV gamma-ray emission.  We have previously published a series of papers investigating the parsec-scale kinematic properties of TeV-detected HBLs through multi-epoch high-resolution VLBI observations, predominantly with the National Radio Astronomy Observatory's Very Long Baseline Array (VLBA) \\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}, and we extend that study in this paper. Previous papers in this series have included: early VLBA and space VLBI observations of Mkn\\,421 (Piner et al.\\ 1999), VLBI observations of Mkn\\,501 (Edwards \\& Piner 2002), VLBA observations of 1ES\\,1959+650, PKS\\,2155$-$304, and 1ES\\,2344+514 (Piner \\& Edwards 2004), and new VLBA polarimetry observations of Mkn\\,421 (Piner \\& Edwards 2005). This fifth paper in the series adds new VLBA observations of the TeV blazars \\object{1ES 1959+650} and \\object{PKS 2155-304} whose kinematics were relatively poorly determined in Piner \\& Edwards (2004) (hereafter Paper~I), as well as a four-epoch series of VLBA observations of \\object{H 1426+428} --- the sixth TeV blazar to be discovered.  Altogether we present six new images of 1ES\\,1959+650, four new images of H\\,1426+428, and one new polarization image of PKS\\,2155$-$304, for a total of eleven new datasets. These datasets comprise the observations from our TeV blazar monitoring program over the years 2001-2004 (excepting the observations of Mkn\\,421, which were published separately in Piner \\& Edwards [2005]).  Note that except for the two brightest sources (Mkn\\,421 and Mkn\\,501), the TeV blazars are too faint in the radio ($\\lesssim$100\\,mJy) to be included in other VLBA monitoring programs such as MOJAVE, and they typically require long observations to obtain images of sufficient dynamic range. Observational backgrounds on the three specific sources studied in this paper are presented in the results section for the specific source.  Our earlier work on the parsec-scale structure of TeV blazars has shown a noticeable lack of superluminal components in their jets, which contrasts with the rapid superluminal motions observed in the jets of more powerful blazars, and with the high Lorentz factors derived from modeling the TeV emission. We have previously interpreted the general lack of superluminal components in HBLs as evidence for a lower bulk Lorentz factor in the parsec-scale radio-emitting region compared to the TeV-emitting region.  Models that have been invoked to explain this `bulk Lorentz factor crisis' include: jets that are decelerated along their length (Georganopoulos \\& Kazanas 2003; Wang, Li, \\& Xue 2004; Bicknell et al.\\ 2005), jets with transverse velocity structures consisting of a fast spine and a slower layer (Giroletti et al.\\ 2004b; Ghisellini, Tavecchio, \\& Chiaberge 2005; Henri \\& Saug\\'{e} 2006), or jets with opening angles large enough that unintentional averaging over multiple viewing angles (because of limited resolution) causes the apparent conflict (Gopal-Krishna, Dhurde, \\& Wiita 2004; Gopal-Krishna, Wiita, \\& Dhurde 2006).  Finally, Gopal-Krishna et al.\\ (2007) consider a combination of the last two models; i.e., large opening angle jets with transverse velocity structures.  Some of these models may be distinguished through the observed statistical distribution of apparent speeds in TeV blazar jets (e.g., Gopal-Krishna et al.\\ 2006), so we conclude the paper by presenting our current best set of TeV blazar apparent speed measurements from the complete series of five papers, which consists of sixteen component speeds in six sources, and discussing the various models in the context of the current observations.  In this paper we use cosmological parameters $H_{0}=71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m}=0.27$, and $\\Omega_{\\Lambda}=0.73$ (Bennett et al.\\ 2003).  When results from other papers are quoted, they have been converted to this cosmology. ",
        "conclusions": "We have presented new multi-epoch VLBI images for the three TeV blazars H~1426+428, 1ES~1959+650, and PKS~2155$-$304 obtained during the years 2001 to 2004. The results for H~1426+428 are the first multi-epoch VLBI results to be presented for this source. The results for 1ES~1959+650 and PKS~2155$-$304 are combined with earlier results for these sources from Paper I to yield a longer time baseline for measurement. The major observational findings for these three sources are: \\begin{enumerate} \\item{H~1426+428's parsec-scale structure during this time range was well-modeled by a $\\sim 17$\\,mJy core and a single $\\sim 3$\\,mJy jet component at a position angle of approximately $-25\\arcdeg$, with an apparent speed of $2.09c\\pm0.53c$.} \\item{1ES~1959+650 consisted of a compact core and a nearby stationary ($0.00c\\pm0.04c$) jet component at a position angle of about $125\\arcdeg$ and a separation of about 0.35 mas at 15 GHz. On larger scales of a few mas the jet is diffuse and directed to the north, so that this source shows an extreme apparent misalignment of about $130\\arcdeg$ on parsec scales.} \\item{PKS~2155$-$304 was observed with dual-circular polarization at 15 GHz. The fractional polarization at the position of the core was 3\\%, and the electric vector position angle was $131\\arcdeg$, about $30\\arcdeg$ misaligned from the innermost jet position angle. The measured apparent speed of the single jet component was $0.93c\\pm0.31c$.} \\end{enumerate}  We combined the new apparent speed measurements from this paper with the apparent speeds measured in TeV blazar jets from our earlier papers to form a current set of apparent speed measurements in TeV HBLs (Table~\\ref{speedtab}). The mean peak apparent pattern speed in the jets of the TeV HBLs is about 1$c$. The statistical result noted in Paper I that the TeV HBLs have significantly slower apparent jet pattern speeds compared to radio-selected or GeV-selected blazars is strengthened by the new results of this paper. The Discussion section ($\\S$~\\ref{disc}) presented a thorough analysis of these results in the context of other radio observations and theoretical models for TeV blazar jets. Conclusions from that analysis were: \\begin{enumerate} \\item{The peak apparent speeds of order 0.5-2$c$ in four of the six studied TeV blazars (Table~\\ref{speedtab}), when taken together with other observed radio properties discussed in $\\S$~\\ref{disc}, are consistent with radio jets with bulk Lorentz factors of $\\Gamma\\sim3$ and viewing angles of a few degrees.} \\item{The very slow peak apparent speeds in two of the six studied TeV blazars in Table~\\ref{speedtab} (Mkn~421 and 1ES\\,1959+650) are likely pattern speeds that are unrelated to the bulk apparent speeds of the jets, which are likely to be similar to those mentioned above.} \\end{enumerate}"
    },
    "0801/0801.2864_arXiv.txt": {
        "abstract": "{The Australian SKA Pathfinder (ASKAP) will be a powerful instrument for performing large-scale surveys of galaxies. Its frequency range and large field of view makes it especially useful for an all-sky survey of Local Volume galaxies, and will probably increase the number of known galaxies closer than 10 Mpc by a factor of two and increase, by at least an order of magnitude, the number detected in HI. Implications for our knowledge of the HI mass function for the very faintest galaxies and for the structure and dynamics of the Local Volume are discussed. } ",
        "introduction": "\\label{sec:intro}  The Local Volume is a key region for the study of the properties of galaxies, including: (1) their internal structure and dynamics; (2) their spatial distribution and dynamics in an environment which lies in the outskirts of a supercluster; and (3) their complete evolutionary history, by virtue of our ability to resolve individual stars. The Local Volume is particularly useful for studying the faintest galaxies and is well-served by having a wealth of accurate redshift-independent TRGB distances through recent surveys by the HST \\cite{Riz07}.  As shown elsewhere in these proceedings, detailed studies of Local Volume galaxies in the 21cm line of neutral hydrogen have been particularly fruitful.  These have recently been rejuvenated by surveys at other wavebands including those of {\\em Spitzer} and {\\em GALEX}.  Due to the sensitivity and resolution, such surveys have often been drawn towards the luminous galaxy population.  However, the Local Volume also offers a unique opportunity to study the faintest galaxies observable and many studies (e.g. {\\em SINGG, LVHIS, THINGS}) have also been careful to select their samples across a range of intrinsic luminosity.  The Square Kilometre Array (SKA) will be a radio telescope of unprecedented power to observe galaxies in the radio continuum and in the 21cm line. Its sensitivity will easily surpass existing telescopes and allow Local Volume galaxies to be studied at the highest spatial resolution. However, it's not due to come on-line for at least a decade. Nevertheless, the next generation of so-called `SKA pathfinders' are around the corner. Their purpose is to test SKA technologies, yet provide sufficient sensitivity to obtain useful science and provide a valuable source of survey material for the SKA itself. Examples of proposed pathfinders are the Allen Telescope Array (ATA) and the Apertif upgrade to WSRT in the northern hemisphere and MeerKAT and the Australian SKA Pathfinder (ASKAP) in the southern hemisphere. In this brief paper, I will look at the implications for our knowledge of the Local Volume of proposed ASKAP surveys. ",
        "conclusions": "The high number of galaxies that, in all likelihood, remain to be discovered in the Local Volume will allow a remarkably dense sampling of the extragalactic environment of the Local Group. This will allow an accurate mapping of the large-scale filamentary features joining the Local Group with Sculptor and other groups. Combined with the redshift-independent distances that are possible for such nearby objects, it will also allow a study of the Hubble flow, infall towards filaments, and tidal stretching owing to nearby overdense regions such as the Local Supercluster and underdense regions such as the Local Void."
    },
    "0801/0801.0050_arXiv.txt": {
        "abstract": "In this talk we introduce our recent results of global 1D MHD simulations for the acceleration of solar and stellar winds. We impose transverse photospheric motions corresponding to the granulations, which generate outgoing Alfv\\'{e}n waves. The \\Alfven waves effectively dissipate by 3-wave coupling and direct mode conversion to compressive waves in density-stratified atmosphere. We show that the coronal heating and the solar wind acceleration in the open magnetic field regions are natural consequence of the footpoint fluctuations of the magnetic fields at the surface (photosphere). We also discuss winds from red giant stars driven by \\Alfven waves, focusing on different aspects from the solar wind. We show that red giants wind are highly structured with intermittent magnetized hot bubbles embedded in cool chromospheric material. ",
        "introduction": "The \\Alfven wave, generated by the granulations or other surface activities, is a promising candidate operating in the heating and acceleration of solar winds from coronal holes. It can travel a long distance so that the dissipation plays a role in the heating of the solar wind plasma as well as the lower coronal plasma, in contrast to other processes, such as magnetic reconnection events and compressive waves, the heating of which probably concentrates at lower altitude.  While high-frequency ioncyclotron waves are recently highlighted for the preferential heating of minor heavy ions (Axford \\& McKenzie 1997) the protons, which compose the main part of the plasma, are supposed to be mainly heated by low-frequency ($\\lesssim 0.1$Hz) components in the MHD regime. because (1) the low-frequency wave is expected to have more power, and (2) the resonance frequency of the proton is higher than those of heavier ions so that the energy of the ioncyclotron wave is in advance absorbed by heavy ions (Cranmer 2000). In this paper, we focus on roles of such low frequency \\Alfven waves.  When considering low-frequency \\Alfven waves in solar and stellar atmospheres, the stratification of the density due to the gravity is quite important because the variation scale of the density, and accordingly \\Alfven speed, is comparable to or shorter than the wavelengths; the WKB approximation is no longer applicable. Also, the amplitude is amplified because of the decrease of the density so that waves easily become nonlinear. Recently, we have extensively studied the heating and acceleration of solar and stellar winds by self-consistent MHD simulations from the photosphere to sufficiently outer region (Suzuki \\& Inutsuka 2005; 2006; hereafter SI06; Suzuki 2007). We review these works in this contribution talk. ",
        "conclusions": "We have performed 1D MHD numerical simulations of solar and stellar winds from the photosphere. The low-frequency \\Alfven waves are generated by the footpoint fluctuations of the magnetic field lines. We have treated the wave propagation and dissipation, and the heating and acceleration of the plasma in a self-consistent manner. Our simulation is the first simulation which treats the wind from the real surface (photosphere) to the (inner) heliosphere with the relevant physical processes.  We have shown that the dissipation of the low-frequency \\Alfven waves through the generation of the compressive waves (decay instability) and shocks (nonlinear steepening) is one of the solutions for the heating and acceleration of the plasma in the coronal holes. However, we should cautiously examine the validity of the 1-D MHD approximation we have adopted. There are other dissipation mechanisms due to the multidimensionality, such as turbulent cascade into the transverse direction (Goldreich \\& Sridhar 1995; Oughton et al.2001) and phase mixing (Heyvaerts \\& Priest 1983). % If \\Alfven waves cascade to higher frequency, kinetic effects (e.g. Nariyuki \\& Hada 2006) becomes important.    We have also extended the solar wind simulations to red giant winds. With stellar evolution, the steady hot corona with temperature, $T \\approx 10^6$ K, suddenly disappears because the surface gravity becomes small; hot plasma cannot be confined by the gravity. Thermal instability also generate intermittent magnetized hot bubbles in cool chromospheric winds."
    },
    "0801/0801.3187_arXiv.txt": {
        "abstract": "Recent observations of the solar corona with the LASCO coronagraph on board of the SOHO spacecraft have revealed the occurrence of triple helmet streamers even during solar minimum, which occasionally go unstable and give rise to large coronal mass ejections. There are also indications that the slow solar wind is either a combination of a quasi-stationary flow and a highly fluctuating component or may  even be caused completely by many small eruptions or instabilities. As a first step we recently presented an analytical method to calculate simple two-dimensional stationary models of triple helmet streamer configurations. In the present contribution we use the equations of time-dependent resistive magnetohydrodynamics to investigate the stability and the dynamical behaviour of these configurations. We  particularly focus on the possible differences between the dynamics of single isolated streamers and triple streamers and  on the way in which magnetic reconnection initiates both small scale and large scale dynamical behaviour of the streamers. Our results indicate that small eruptions at the helmet streamer cusp may incessantly accelerate  small amounts of plasma without significant changes of the equilibrium configuration and might thus contribute to the non-stationary slow solar wind. On larger time and length scales, large coronal eruptions can occur as a consequence of large scale magnetic reconnection events inside the streamer configuration. Our results also show that triple streamers are usually more stable than a single streamer. ",
        "introduction": "Recent observations of the corona with the LASCO coronagraph \\cite{schwenn97} on board of the SOHO spacecraft showed that the corona can be highly structured even during the solar activity minimum. The observations revealed a triple structure of the streamer belt which was existent for several consecutive days. The observations further showed that these triple structures occasionally go unstable leading to a seemingly new and extraordinarily huge kind of coronal mass ejection (global CMEs). Natural questions arising from these observations are whether the helmet streamer triple structure is directly connected with or responsible for the occurrence of global CMEs and what is the physical mechanism of their formation.  The structure of helmet streamers and their stability has been studied both observationally and theoretically for a long time (e.g. \\opencite{pneuman:kopp71}; \\opencite{cuperman:etal90}; \\opencite{cuperman:etal92}; \\opencite{koutchmy:livshits92}; \\opencite{wang:etal93}; \\opencite{cuperman:etal95}; \\opencite{wu:etal95}; \\opencite{bavassano:etal97}; \\opencite{noci}; \\opencite{hundhausen99}). There seems to be a natural association of helmet streamers with coronal eruptions and coronal streamers are assumed to be the source region of the slow solar wind. The traditional view towards the origin of the slow solar wind is that it is a more or less stationary plasma flow on open field lines around the closed field lines of a helmet streamer. Recent observations (e.g. \\opencite{habbal98}; \\opencite{noci}) challenge this traditional view and indicate that the slow solar wind is non-stationary and seems to be produced and accelerated by small eruptions in the helmet streamer stalk above the cusp. This acceleration process of the slow solar wind has been compared with the rise of smoke above a burning candle (Schwenn, private communication).  Pre-SOHO observations of multiple streamer configurations during the maximum phase of solar activity and of the multiple current sheet structure of the heliospheric plasma sheet \\cite{crooker:etal93,woo:etal95} have initiated several studies of the dynamics and stability of multiple current sheets with variations in only one spatial dimension \\cite{otto:birk92,yan:etal94, dahlburg:karpen95,birk:etal97,wang:etal97}. \\inlinecite{einaudi:etal99} have recently presented a model for the generation of the non-stationary slow solar wind based on linear and non-linear stability calculations for a single one-dimensional current sheet with field-aligned flow.  All these models do, however, in a strict sense only apply to the streamer stalk, i.e.\\ to the open field lines of the heliospheric current sheet. Here we aim to investigate both closed and parts of the adjacent open field line regions.  Models of multiple arcade and loop structures have been investigated before by e.g.\\ \\inlinecite{mikic:etal88}, \\inlinecite{biskamp:welter89} and most recently by \\inlinecite{antiochos:etal99}. Our model differs from these models basically by the possibility of having a flexible analytical initial condition for the time-dependent calculations allowing the investigation of different types of structures.  As a first step towards improving the theoretical understanding of the above mentioned phenomena in triple streamer configurations, we have calculated analytic two-dimensional static models of triple helmet streamers in a previous paper (\\opencite{paper1}; further cited as {\\it paper I}). The aim of the present paper is to undertake the next step in this investigation and to study the stability of the stationary state helmet streamer configurations calculated in {\\it paper I}. We will carry this out with the help of time-dependent numerical experiments using the equations of resistive magnetohydrodynamics.  The outline of the paper is as follows. In Section \\ref{basics} we discuss the basic equations and briefly describe our numerical method. Section \\ref{model} outlines our main model assumptions. In Section \\ref{results} we present the results of numerical experiments and in Section \\ref{conclusions} we discuss our results and give an outlook on future work. ",
        "conclusions": "\\label{conclusions} In this paper we have tried to make a step towards a better theoretical understanding of the dynamics of helmet streamers with triple structure. We investigated the possible role of triple streamers for the development of coronal mass ejections and as a possible source for plasma and magnetic field for the wind emanating from streamer regions. In previous works (e.g. \\opencite{steinolfson94}; \\opencite{linker94}; \\opencite{wu:etal95}; \\opencite{wu97} ) only single helmet streamer where modelled and these models assumed the slow solar wind as a stationary plasma flow on open field lines in the streamer region. To get a start equilibrium these authors solved the ideal time-dependent MHD equations numerically until a stationary state was reached. These works showed that helmet streamers can become unstable and produce coronal mass ejections.  The  present investigations where motivated by the new observations with the LASCO coronagraph on SOHO \\cite{schwenn97}. These observations showed, that the streamer belt in the solar activity minimum typically has a triple structure. The observations also gave further strong evidence that a stationary slow solar wind may not exist but is produced by many small eruptions. Apart from these continuously occuring small eruptions, also large, however rarely occuring coronal mass ejections are generated in the triple streamer belt.  To take these observations into account in a helmet streamer model, we developed an analytical stationary  model of triple helmet streamers using the ideal MHD equations in {\\it paper I}. The initial states have to possess a non-vanishing free energy to allow their instability with respect to magnetic reconnection. As discussed in {\\it paper I} we took into account the observation that the streamer configurations are very extended in the radial direction to simplify the calculation of the initial states. We emphasize that  such radially extended configurations cannot be modeled by potential fields. In the present paper we investigated the stability of these stationary state configurations with the help numerical experiments in the framework of time-dependent resistive MHD. We used three ad hoc models for the resistivity, a constant resistivity, a resistivity localized at the thin current sheets and a current-dependent resistivity. We first investigated the ideal stability of our triple streamer configurations and found that they are stable on the time-scale of our simulations.  Next we investigated the resistive stability of a triple streamer configuration without cusp structure. We found that our triple streamer configuration is resistively unstable. Reconnection takes place and plasmoids form in each of the three closed field line regions. By comparing the time evolution of the triple streamer model with a single streamer model we found that for a triple streamer configuration without cusp structure the time-evolution is usually slower than the corresponding time evolution of a single streamer. The triple streamer evolution also shows characteristic differences in comparison with the single streamer case concerning the location of the reconnection sites. We could explain these differences by the influence the three streamers exert on each other.  For triple streamer configurations with cusp structure we found quite similar results for reconnection processes inside the closed field line regions but in addition  we found that the helmet streamer stalk above the cusp is highly unstable to reconnection. This reconnection process leads to the formation of a dome above the triple structure, i.e. a region of closed field lines which encloses the triple structure completely. The resistive instability of the streamer stalk current sheet could be a possible source of plasma and magnetic field for the non-steady solar wind emanating from the streamer regions. In the present paper we have only been able to demonstrate that this mechanism works with a preexisting cusp structure. In later stages of our simulations the cusp is replaced by an X-point at which reconnection can take place. The inclusion of flow on open field lines would also allow for the possibility to generate a new cusp structure making a repetition of the process possible.  Furthermore we found   interaction between the two outer streamers just below the cusp region. This interaction can also contribute to the formation of the dome. This dome makes it more difficult for plasmoids to escape and thus streamers with cusp structure within our model are less likely to eject material than the configurations without cusp.  These results are consistent  with the observational finding that the triple streamer configuration is observed to be stable for several days. One may also speculate about the fact that the observations usually show three streamers which approximately have the same radial extension. A possible explanation on the basis of our model is that if one of the streamers grows and becomes much larger than the other streamers, it becomes prone to instability and a coronal mass ejection occurs similarly to the case of a single streamer. In this process the streamer looses energy, mass and magnetic flux and returns to its original state.  The numerical experiments presented here can only be considered as a very first step towards a complete model of these interesting phenomena. We already mentioned above that although the models with cusp structure seem to be matching the observed streamer structure best of all our models, it is not possible to include regions of open field lines  between the streamers in our models. One way to overcome this shortcoming would be to use three-dimensional models, which is a natural next step. Other possible improvements of the present work are the inclusion of gravity and the use of spherical geometry."
    },
    "0801/0801.1527_arXiv.txt": {
        "abstract": " ",
        "introduction": "The standard model (SM) provides an accurate description of particle physics below the electroweak scale but it is generally though not to be valid up to arbitrary energies. Extensions of this celebrated scheme, invoked to cure diseases like the strong CP, hierarchy or flavor problems, often involve higher gauge symmetries and further matter content. Moreover, string theory is a preferred candidate for the unification of quantum mechanics and general relativity where additional gauge and matter fields are assured. At low energies, some of these new fields can arrange into a ``hidden sector\" if only very massive particles (or gravity) mediate interactions between them and the SM ``visible sector\".  Of course, depending on the scalar content of the theory, gauge symmetries can either be spontaneously broken or remain exact. Then, the corresponding ``hidden\" bosons could have in principle an arbitrary mass. If this is small enough, these hidden bosons can have a very rich phenomenology at present affordable energy scales.  The simplest case concerns just a novel $U(1)_\\mathrm{h}$ symmetry and its corresponding gauge boson, henceforth called ``hidden\" photon. The interplay between this hidden photon and the SM photon modifies the predictions of quantum electrodynamics \\cite{Okun:1982xi}, often claimed to be the most accurate of all physical theories so far, thus constraining the hidden photon parameters. We can turn this argument in the opposite direction: the constraints on hidden photon parameters give us information about how accurate is the QED description of nature at low energies.  A number of laboratory experiments has been devoted to the search of hidden photons, the resulting bounds being strongly dependent on the hidden photon mass. For masses corresponding to macroscopic length scales, experiments testing the Coulomb law \\cite{Williams:1971ms,Bartlett:1988yy} set strong constraints on hidden photons, but still they could be largely improved by experiments dealing with high quality microwave cavities \\cite{Jaeckel:2007ch}. In the microscopic range, laser experiments are also becoming very powerful probes of hidden sector particles \\cite{Cameron:1993mr, Gies:2006ca,Ahlers:2006iz,Ahlers:2007rd,Jaeckel:2007gk,Robilliard:2007bq,Chou:2007zz,Ahlers:2007qf}. At atomic distances, comparison of the Rydberg constant for different atomic levels gives interesting but weak bounds \\cite{Beausoleil:1987,Garreau:1987}. Finally, particle colliders extend the mass range until typical electroweak scales \\cite{Kors:2004dx,Feldman:2006ce,Chang:2006fp,Kumar:2006gm,Feldman:2007nf}.  On top of that, the evolution of stars turns out to be the most sensitive ``laboratory\" to study properties of novel low mass weakly interacting particles \\cite{Raffelt:1996wa,Raffelt:1999tx}. Even with tiny couplings to electrons and protons, they might be still copiously created in the interior of hot and dense stars. Because of their weak interactions they might abandon the star without further scattering, accelerating the consumption of nuclear fuel, and therefore the stellar evolution \\cite{Dicus:1978fp,Frieman:1987ui}. Our present observational data on stellar evolution can strongly constrain this novel luminosity, although the bounds can be relaxed in some concrete models \\cite{Masso:2005ym,Masso:2006gc,Jaeckel:2006id,Jaeckel:2006xm,Brax:2007ak}.  Interestingly enough, in this case the Sun itself could be a copious emitter of weakly interacting particles, that could eventually be detected at Earth inside a sensitive detector \\cite{Sikivie:1983ip} (as this is actually the case with neutrinos). Several of these so-called ``helioscopes\" \\cite{vanBibber:1988ge} have been built with the aim of detecting solar axions \\cite{Lazarus:1992ry,Moriyama:1998kd,Zioutas:2004hi} and remarkably, the CAST collaboration has even recently surpassed the sensitivity of the energy loss arguments for the axion coupling to two photons \\cite{Andriamonje:2007ew}. As we will see, helioscopes can also detect hidden photons so the CAST limits can also be used to constrain the solar hidden photon flux.  The energy loss argument and helioscope bounds for hidden photons were already studied in  \\cite{popov:1991,Popov1999}. However, in this paper the author does not consider either the possibility of a resonant production, which can enhance enormously the hidden photon flux, nor the emission of transversely polarized hidden photons.  In this paper we compute the energy loss bounds using the latest solar data \\cite{Bahcall:2004pz} and derive the CAST helioscope bounds. In particular, it is shown that the new CAST results are extremely sensitive to solar hidden photons, providing the strongest constraint of their existence in the mass range $\\muu\\sim 0.01-1$ eV. Moreover, accounting for the resonant production improves the energy loss bounds in \\cite{Popov1999} roughly up to 1 order of magnitude.  The paper is organized as follows: In Sec.\\ref{Sproduction} the hidden photon solar emission is derived while the principles of the CAST helioscope detection are reviewed in Sec.\\ref{Shelioscope}. We compute the energy loss and CAST bounds in Sec.\\ref{Sbounds} and finally the conclusions are presented. ",
        "conclusions": "I have addressed the calculation of the solar emission of a hypothetical hidden sector photon $B^\\mu$ mixing kinetically with the standard model ordinary photon. I have shown that a resonant effect is possible when the dispersion relation of solar plasmons fits the particle-like dispersion relation of the hidden photon. This happens for transverse plasmons if the hidden photon mass $\\muu$ lies in the range $1\\sim295$ eV (the range of the plasma frequency in the solar model used) and for longitudinal plasmons as long as $\\muu\\lesssim295$ eV.  The conservative requirement that the hidden photon luminosity should not exceed the solar standard luminosity bounds the amount of kinetic mixing up to $\\mix\\lesssim10^{-14}$ depending on the mass and the polarization. At masses beyond $1$ eV, where the strongest bound is reached, the emission of transversally polarized hidden photons dominates over the emission of longitudinal ones. Below this mass both polarizations contribute in a similar amount.  At low masses the bounds are weaker, relaxing proportionally to $\\muu^2$. However, the non observation of a signal in the CAST axion helioscope improves the bounds in this region up to 2 orders of magnitude. Altogether, these are the best limits on the mixing parameter $\\mix$ in the range $3$ meV$<\\muu<40$ keV.   A small room for improvement is available for large masses $\\muu\\gg 10$ keV if the CAST detectors were to rise their top energy threshold.  It should be interesting to extend this study to other stellar objects like supernovae, white dwarfs, red giants and horizontal branch stars, since these can provide stronger bounds than the Sun, specially at masses where a resonant production is possible."
    },
    "0801/0801.0467_arXiv.txt": {
        "abstract": "We argue that the typical energy density of a light scalar field should not be less than $H^4$ in the inflationary Universe. This requirement implies that the non-Gaussianity parameter $f_{NL}$ is typically bounded by the tensor-scalar ratio $r$ from above, namely $f_{NL}\\lesssim 518\\cdot r^{1\\over 4}$. If $f_{NL}=10^2$, inflation occurred around the GUT scale. ",
        "introduction": "Inflation model \\cite{Guth:1980zm,Albrecht:1982wi,Linde:1982zj} provides an elegant mechanism to solve the horizon, flatness and primordial monopole problem due to a quasi-exponential expansion of the universe before the hot big bang. The temperature anisotropies in cosmic microwave background radiation (CMBR) and the large-scale structure of the Universe are seeded by the primordial quantum fluctuations during inflation \\cite{Mukhanov:1990me,Lyth:1998xn}. Since the density perturbation is roughly $10^{-5}$, it is good enough to apply the linear perturbation theory to calculate the quantum fluctuations during inflationary epoch. Within this approach, the Fourier components of fluctuations are uncorrelated and their distribution is Gaussian. That is why the non-Gaussianity from the simplest inflation models is very small ($|f_{NL}|<1$). For useful discussions on non-Gaussianity see \\cite{Salopek:1990jq,Salopek:1990re,Falk:1992sf,Gangui:1993tt,Acquaviva:2002ud,Maldacena:2002vr}, and for a nice review see \\cite{Bartolo:2004if}.   The non-Gaussian perturbation is governed by the n-point correlation function of the curvature perturbation \\e \\langle \\Phi({\\bf{k}_1})\\Phi({\\bf{k}_2})\\cdot\\cdot\\cdot\\Phi({\\bf{k}_n})\\rangle=(2\\pi)^3\\delta^3(\\sum_{i=1}^{n}{\\bf{k}_i})F_n(k_1,k_2,...,k_n),\\q where $\\Phi({\\bf{k}})$ is the Fourier mode of Bardeen's curvature perturbation. The leading non-Gaussian features are known as the bispectrum (three-point function) and trispectrum (four-point function), with their sizes conventionally denoted as $f_{NL}$ and $\\tau_{NL}$ respectively. The non-linearity parameter $f_{NL}$ defined in \\cite{Komatsu:2001rj} takes the form \\e \\Phi({\\bf x})=\\Phi_L({\\bf x})+f_{NL}[\\Phi_L^2({\\bf x})-\\langle\\Phi_L^2({\\bf x})\\rangle], \\q where $\\Phi_L$ denotes the linear Gaussian part of the perturbation in real space. The non-Gaussianity parameter $f_{NL}$ characterizes the amplitude of the primordial non-Gaussian perturbations. The value of $\\Phi_L$ is roughly $10^{-5}$. If $f_{NL}=10^2$, the distribution of $\\Phi$ is still consistent with a Gaussian distribution to $0.1\\%$. It is not easy to detect a small non-Gaussianity in the experiments. The shape of non-Gaussianity is also very important. If $F(k_1,k_2,k_3)$ is large for the configurations in which $k_1\\ll k_2,k_3$, it is called local, ``squeezed\" type; if $F(k_1,k_2,k_3)$ is large for the configurations in which $k_1\\sim k_2\\sim k_3$, it is called non-local, ``equilateral\" type.   Recently a large primordial local-type non-Gaussianity was reported to be marginally detected from the data of WMAP3 in \\cite{Yadav:2007yy}: \\e 27<f_{NL}^{local}<147\\q at $95\\%$ C.L., with a central value of $f_{NL}=87$. See also \\cite{Komatsu:kitpc}. In \\cite{Komatsu:2008hk} WMAP group recently reported their new result as \\e -9<f_{NL}^{local}<111. \\q If a large local-type non-Gaussianity is confirmed by the future cosmological observations and data analysis at high confidence level, it strongly implies that the physics of the early Universe is more complicated than the simple single-field slow-roll inflation.  There are several mechanisms to generate a large local-type non-Gaussianity:\\\\ (1) a non-linear relation between inflaton and curvature perturbations \\cite{Gangui:1993tt}; \\\\ (2) curvaton scenario \\cite{Linde:1996gt,Enqvist:2001zp,Lyth:2001nq,Moroi:2001ct,Lyth:2002my,Bartolo:2003jx}; \\\\ (3) Ekpyrotic model \\cite{Creminelli:2007aq,Buchbinder:2007at,Koyama:2007if,Lehners:2007wc}.\\\\ In this paper we focus on curvaton scenario. As we know, there might be many light scalar fields in supersymmetric theory or string theory. These light scalar fields which are subdominant during inflation are called curvaton. The energy density during inflation is still dominated by the the potential of inflaton. In the usual inflation model the fluctuations of inflaton dominate the curvature perturbations. The energy density and the perturbations caused by these light scalar fields can be ignored during inflation. However it is possible that the fluctuation of curvaton becomes relevant and causes a large local-type non-Gaussianity when its energy density is a significant fraction of the total energy after the end of inflation. In curvaton scenario the perturbations from the inflaton field are considered to be negligible. A large local-type non-Gaussianity may shed light on these light scalar fields.  The non-Gaussianity produced by curvaton is in inverse proportion to the fraction of curvaton energy density in the energy budget at the epoch of curvaton decay. The non-Gaussianity increases as the curvaton energy density decreases. However the energy density for a homogeneous scalar field in the quasi-de Sitter space is typically is given by $H_*^4$. Therefore the non-Gaussianity has an upper bound \\e f_{NL}\\lesssim 518\\cdot r^{1\\over 4}, \\q where $r$ is the tensor-scalar ratio. Present bound on the tensor-scalar ration \\cite{Komatsu:2008hk} is $r<0.20$ and then $f_{NL}\\lesssim 346$.  Our paper is organized as follows. We review curvaton scenario and the curvature perturbation generated by curvaton in Sec. 2. In Sec. 3, we argue that the energy density of curvaton is typically estimated as $H_*^4$ during inflation and then the non-Gaussianity parameter $f_{NL}$ is bounded by the tensor-scalar ratio from above. In Sec. 4, we point out the challenge for getting a red-tilted primordial power spectrum in curvaton scenario. Section 5 contains some discussions including how to distinguish the curvaton scenario from Ekpyrotic scenario. ",
        "conclusions": "An unambiguous detection of $f_{NL}>10$ will rule out most of the existing inflation models. The uncertainty of $f_{NL}^{local}$ from WMAP will shrink to be 42 for 8 years data, and 38 for 12 years data \\cite{Komatsu:kitpc}. The accuracy of PLANCK will be roughly $6$. If the local-type non-Gaussianity is of order $10^2$, it will be detected at high confidence level in the near future.  In this paper we investigate the curvaton scenario in detail and find that the inflation scale is roughly at GUT scale in order to get a large local-type non-Gaussianity. Ekpyrotic model can also provide a large local-type non-Gaussianity. However unlike the slightly red-tilted gravitational wave spectrum in the inflation/curvaton model, the gravitational wave spectrum in Ekpyrotic model is strongly blue and then the amplitude is exponentially suppressed on all observable scales \\cite{Boyle:2003km}. Gravitational wave perturbation can be used to distinguish curvaton scenario from Ekpyrotic model.   Here we also want to clarify two points in this paper. One is that maybe we miss some order-one coefficients in (\\ref{bsg}) and (\\ref{ggg}), and then the bound in (\\ref{fnlr}) should be modified to be a little looser or more stringent. But the order of magnitude of our result can be trusted. The other is that the bound in (\\ref{fnlr}) depends on the WMAP normalization $P_{\\zeta,{WMAP}}=2.457\\times 10^{-9}$. If we let the normalization of the density perturbation free, we find \\e f_{NL}\\lesssim \\gamma^{-1}\\cdot 518\\cdot r^{1\\over 4}, \\q and eq. (\\ref{fnlp}) becomes\\e f_{NL}\\lesssim \\gamma^{-1}\\cdot 871 \\cdot \\({1\\over \\Delta N_e}{|\\Delta \\phi|\\over M_p}\\)^\\half, \\q where $\\gamma=(P_\\zeta/P_{\\zeta,{WMAP}})^\\half$. Larger the amplitude of the density perturbation, smaller the non-Gaussianity. If we also consider $|\\Delta \\phi|/M_p<1$ for $\\Delta N_e=50$, $f_{NL}\\lesssim 123\\cdot \\gamma^{-1}$. The density perturbation cannot be larger than 10 times of its observed value in our Universe if the non-Gaussianity parameter is not smaller than 10.   The equilateral-type non-Gaussianity has not been detected. Many models \\cite{ArkaniHamed:2003uz,Chen:2006nt,ArmendarizPicon:1999rj,Chen:2007gd,Li:2007} were suggested to generate large equilateral-type non-Gaussianity as well. Other mechanisms concerning the large non-Gaussianity are discussed in \\cite{Dvali:2003em,Dvali:2003ar,Matsuda:2007tr,Matsuda:2007ds,Berera:1995ie,Berera:1996nv,Gupta:2002kn}. Anyway, non-Gaussianity will be an important issue for cosmology and string theory.  \\vspace{.5cm}  \\noindent {\\bf Acknowledgments}  We would like to thank M.~Sasaki and P.~J.~Yi for useful discussions."
    },
    "0801/0801.4378_arXiv.txt": {
        "abstract": "The annihilations of neutralino dark matter (or other dark matter candidate) generate, among other Standard Model states, electrons and positrons. These particles emit synchrotron photons as a result of their interaction with the Galactic Magnetic Field. In this letter, we use the measurements of the WMAP satellite to constrain the intensity of this synchrotron emission and, in turn, the annihilation cross section of the lightest neutralino. We find this constraint to be more stringent than that provided by any other current indirect detection channel. In particular, the neutralino annihilation cross section must be less than $\\approx 3\\times 10^{-26}$cm$^3$/s ($1\\times 10^{25}$cm$^3$/s) for 100 GeV (500 GeV) neutralinos distributed with an NFW halo profile. For the conservative case of an entirely flat dark matter distribution within the inner 8 kiloparsecs of the Milky Way, the constraint is approximately a factor of 30 less stringent. Even in this conservative case, synchrotron measurements strongly constrain, for example, the possibility of wino or higgsino neutralino dark matter produced non-thermally in the early universe. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.4187_arXiv.txt": {
        "abstract": "The chemical compositions of seven Carbon-Enhanced Metal-Poor (CEMP) turn-off stars are determined from high-resolution spectroscopy. Five of them are selected from the SDSS/SEGUE sample of metal-poor stars. Another star, also chosen from the SDSS/SEGUE sample, has only a weak upper limit on its carbon abundance obtained from the high-resolution spectrum. The effective temperatures of these objects are all higher than 6000~K, while their metallicities, parametrized by [Fe/H], are all below $-2$; the star with the lowest iron abundance in this study has [Fe/H] = $-3.1.$ Six of our program objects exhibit high abundance ratios of barium ([Ba/H] $> +1$), suggesting large contributions of the products of former AGB companions via mass transfer across binary systems. One star in our study ({\\objf}) exhibits a rapid variation in its radial velocity, which is a strong signature that this star belongs to a close binary. Combining our results with previous studies provides a total of 20 CEMP main-sequence turn-off stars for which the abundances of carbon and at least some neutron-capture elements are determined. Inspection of the [C/H] ratios for this sample of CEMP turn-off stars show that they are generally higher than those of CEMP giants; their dispersion in this ratio is also smaller. We take these results to indicate that the carbon-enhanced material provided from the companion AGB star is preserved at the surface of turn-off stars with no significant dilution, which appears counter to expectations if processes such as thermohaline mixing have operated in unevolved CEMP stars. In contrast to the behavior of [C/H], a large dispersion in the observed [Ba/H] is found for the sample of CEMP turn-off stars, suggesting that the efficiency of the s-process in very metal-poor AGB stars may differ greatly from star to star. Four of the six stars from the SDSS/SEGUE sample exhibit kinematics that are associated with membership in the outer-halo population, a remarkably high fraction. ",
        "introduction": "\\label{sec:intro}  Abundance studies of very metal-poor (VMP; [Fe/H] $ < -2.0$)\\footnote{[A/B] = $\\log(N_{\\rm A}/N_{\\rm B})- \\log(N_{\\rm A}/N_{\\rm B})_{\\odot}$, and $\\log\\epsilon_{\\rm A} = \\log(N_{\\rm A}/N_{\\rm H})+12$ for elements A and B.} stars have been pursued over the past few decades in order to constrain models of nucleosynthesis, stellar evolution, and early chemical enrichment in the Galaxy \\citep[e.g., ][]{beers05}. One important result of these studies is the discovery of Carbon Enhanced Metal-Poor (CEMP) stars, which appear with increasing frequency at lower metallicity \\citep{beers92, beers05, lucatello06, marsteller07}. These stars may be closely related to carbon stars in the Galactic halo, known as CH stars \\citep{keenan42} and subgiant CH stars \\citep{bond74}.  Recent chemical abundance studies for CEMP stars have revealed that most (70--80\\%) CEMP stars also exhibit excesses of s-process elements such as Ba (the CEMP-s stars, according to Beers \\& Christlieb 2005), indicating that the origin of the carbon excesses in these stars is likely to be the triple-$\\alpha$ reaction in Asymptotic Giant Branch (AGB) stars \\citep[e.g., ][]{aoki07}. The CEMP stars that are observed at present are likely to have been polluted by the transfer of carbon-enhanced material from a (former) AGB companion across a binary system, while the AGB star itself has now evolved to an unseen white dwarf \\citep[e.g.,][]{lucatello05}. Thus, the abundance patterns of heavy elements in these stars provide useful constraints on models for s-process nucleosynthesis in AGB stars. On the order of 20\\% of CEMP stars exhibit no significant enhancement of their neutron-capture elements (the CEMP-no stars, according to Beers \\& Christlieb 2005), suggesting the existence of other possible origins for their carbon excesses \\citep[e.g., ][]{norris97b, aoki02a}.  \\citet{aoki07} have shown that the CEMP-no stars generally occur at very low [Fe/H]; extreme examples of this class of stars include HE~0107--5240 and HE~1327--2326, two hyper metal-poor (HMP) stars with [Fe/H] below $-5.0$ \\citep{christlieb02, frebel05} and having very large carbon excesses ([C/Fe] $\\sim +4$), as well as the recently identified ultra metal-poor (UMP; [Fe/H] = $-4.8$) star HE~0557-4840, with [C/Fe] $= +1.6$ (Norris et al. 2007).  Among the CEMP stars for which chemical compositions have been obtained from high-resolution spectroscopy, main-sequence turn-off stars are expected to be of particular importance. In the case of mass transfer in binary systems, the accreted material from the primary AGB star has been mixed at least by the first dredge-up in red giants, while turn-off stars might preserve pure material accreted from the primary at their surfaces. In such cases, one can investigate the efficiency of the production of carbon and neutron-capture elements in AGB stars from abundance measurements of the secondary star. Another interesting view arises from the suggested influence of so-called thermohaline mixing (Charbonnel \\& Zahn 2007; Stancliffe et al. 2007; Denissenkov \\& Pinsonneault 2007), which provides the possibility of mixing the accreted surface material while the observed star is still on the main-sequence or only slightly evolved, prior to first dredge-up.  In this scenario, the contrast of the observed surface abundances of CEMP turn-off stars with more evolved CEMP stars also provides valuable clues to the nature of this proposed extensive mixing process.  Very large new samples of CEMP stars have recently become available, discovered during the course of the Sloan Digital Sky Survey (SDSS; York et al. 2000; Adelman-McCarthy et al. 2007). Although originally designed as an extragalactic survey, SDSS has also discovered large numbers of VMP stars \\citep{beers06}.  Although some of the CEMP stars are the result of directed studies (Margon et al. 2002; Downes et al. 2004), many of them have appeared among the calibration objects used by SDSS to obtain spectrophotometric and telluric corrections for other spectroscopic data.  These calibration stars are primarily brighter, metal-poor main-sequence turn-off F- and G-type stars. The ongoing extension to SDSS, SDSS-II (which includes the program SEGUE: Sloan Extension for Galactic Understanding and Exploration), is expected to provide tens of thousands of additional VMP stars, at least several thousand of which are expected to be CEMP stars. This paper reports the first application of abundance measurements obtained with high-resolution spectroscopy for CEMP star candidates found by the SDSS/SEGUE surveys.  In \\S 2 we discuss the identification of our sample stars and the observations that were carried out. A description of our analysis techniques and results is provided in \\S 3. In \\S 4 we present a discussion of our findings. The interesting kinematics of the SDSS/SEGUE CEMP turn-off stars are discussed in \\S 5.  We conclude with a few remarks in \\S 6. ",
        "conclusions": "Chemical compositions of seven CEMP turn-off stars are determined. Six stars among them exhibit a large excess of Ba, signature of a contribution by the nucleosynthesis in an AGB star. The distribution of carbon abundances in these stars suggest that the surface of such stars preserves the material transferred from the AGB star that was the erstwhile primary star in a binary system. If this is the case, the relatively wide distribution of Ba abundances ([Ba/H]) indicates a diversity of the efficiency of the s-process in metal-poor AGB stars. Further studies to identify the physical mechanism that produces such diversity are clearly desired.  The present study is the first application of high-resolution spectroscopy to candidate CEMP stars from the SDSS and SEGUE sample. Comparisons of our results on stellar parameters and chemical abundances with the estimates from the SDSS spectra confirmed that the selection of metal-poor stars works well in general. The SDSS/SEGUE survey is providing a large sample of candidate metal-poor stars. High-resolution spectroscopy for such stars in the near future will reveal the chemical abundance trends in the lowest metallicity range, as well as be useful for exploring the possible dependence of their chemical properties on their derived kinematics."
    },
    "0801/0801.3591_arXiv.txt": {
        "abstract": "{Deep optical CCD images of the supernova remnant G 15.1$-$1.6 were obtained and filamentary and diffuse emission has been discovered. The images, taken in the emission lines of \\hnii, \\sulfur~and \\oiii, reveal filamentary and diffuse structures all around the remnant. The radio emission at 4850 MHz in the same area is found to be well correlated with the brightest optical filaments. The IRAS 60$\\mu$m emission may also be correlated with the optical emission but to a lesser extent. The flux calibrated images suggest that the optical emission originates from shock-heated gas (\\sulfur/\\ha\\ $>$ 0.4), while there is a possible \\HII\\ region (\\sulfur/\\ha\\ $\\sim$ 0.3) contaminating the supernova remnant's emission to the east. Furthermore, deep long--slit spectra were taken at two bright filaments and also show that the emission originates from shock heated gas. An \\oiii\\ filamentary structure has also been detected further to the west but it lies outside the remnant's boundaries and possibly is not associated to it. The \\oiii\\ flux suggests shock velocities into the interstellar $``$clouds'' $\\sim$100 \\vel, while the \\siirat\\ ratio indicates electron densities up to $\\sim$250 cm$^{-3}$. Finally, the \\ha\\ emission has been measured to be between 2 to 7 $\\times$ \\fluxa, while the lower limit to the distance is estimated at 2.2 kpc. ",
        "introduction": "Supernova remants (SNRs) play an important role to understand the SN mechanism, the interstellar medium (ISM) and their interaction. Most of the SNRs have been detected in radio from their non--thermal synchrotron emission. Observations of SNRs in X--rays allow us to directly probe the hot gas inside the primary shock wave, while optical observations offer an important tool for the study of the interaction of the shock wave with dense material found in the ISM. New searches in optical waveband continue to identify Galactic SNRs (e.g. Boumis et al. \\cite{bou02}, \\cite{bou05}; Mavromatakis et al. \\cite{mav02}, \\cite{mav05}) while in the last decade, observations in X--rays have also detected new Galactic SNRs (e.g. Seward et al. \\cite{sew95}; see also Green \\cite{gre06} for a complete catalogue).  G 15.1--1.6 is not a well known SNR, and was first detected by Reich et al. (\\cite{rei88}) in the Effelsberg 2.7 GHz survey and the radio image was published by Reich et al. (\\cite{rei90}). It is classified as a shell--type SNR, with a spectral index of $\\sim$0.8. Its angular size is 30\\arcmin $\\times$ 24\\arcmin and using the brightness--to--diameter ($\\Sigma$--D) relationship the distance of the remnant was calculated at 5.7 kpc (Green \\cite{gre06}). Radio surveys of the surround area do not reveal any pulsar to by associated with G15.1--1.6 while it has not been detected optically in the past.  In this paper we report the optical detection of G 15.1--1.6. We present images of the remnant in the \\hnii, [S~{\\sc ii}] and [O~{\\sc iii}] emission lines. Deep long slit spectra were also acquired in a number of selected areas. In Sect. 2 we present information about the observations and data reduction, while the results of the imaging and spectroscopic observations are given in Sect. 3. In Sect. 4 we discuss the optical properties of this SNR, while in Sect. 5 we summarize the results of this work. ",
        "conclusions": "The faint supernova remnant G 15.1$-$1.6 was observed for the first time in major optical emission lines. The images show filamentary and diffuse emission structures. The bright filaments are very well correlated with the remnant's radio emission at 4850 MHz suggesting their association, while correlation evidence also shown with the IRAS 60$\\mu$m map. The flux calibrated images and the long--slit spectra indicate that the emission arises from shock heated gas. Finally, an upper limit for the electron density of 250 \\dens and a lower limit for the distance of 2.2 kpc are calculated."
    },
    "0801/0801.1594_arXiv.txt": {
        "abstract": "{ We present a compilation of measurements of the stellar mass density as a function of redshift. Using this stellar mass history we obtain a star formation history and compare it to the instantaneous star formation history. For $z<0.7$ there is good agreement between the two star formation histories. At higher redshifts the instantaneous indicators suggest star formation rates larger than that implied by the evolution of the stellar mass density. This discrepancy peaks at $z=3$ where instantaneous indicators suggest a star formation rate around $0.6$ dex higher than those of the best fit to the stellar mass history. We discuss a variety of explanations for this inconsistency, such as inaccurate dust extinction corrections, incorrect measurements of stellar masses and a possible evolution of the stellar initial mass function. } ",
        "introduction": "Much contemporary research in extragalactic astronomy has revolved around the determination of the instantaneous cosmic star formation history (SFH, Madau et al. 1996; Lilly et al. 1996). However, measuring this quantity from observations requires a number of assumptions, with the form of the dust obscuration corrections and stellar initial mass function (IMF, see Kroupa 2007a for a recent overview) being among the most important.  Integration of the instantaneous star formation history over redshift, making appropriate corrections for stellar evolution processes, yields the current stellar mass density. This quantity can be independently measured, typically using extensive galaxy surveys such as the 2dFGRS or SDSS, often combined with near infrared (NIR) measurements. Numerous studies have attempted comparisons of these quantities. Madau, Pozzetti \\& Dickinson (1998), Cole et. al (2001), Fontana et al. (2004) and Arnouts et al. (2007) all found good agreement between the SFH with a low dust content and measured values of the stellar mass density. On the other hand there have been a number of studies (Eke et al. 2005; Hopkins \\& Beacom 2006, hereafter HB06) which claim that the instantaneous SFH overpredicts the low redshift stellar mass density. We attempt to investigate this possible discrepancy using a compilation of the most up to date measurements of the stellar mass density history (SMH, the $\\Omega_{*}$-redshift relation). This relation is intricately connected to the instantaneous star formation history, but in this context it has some important advantages. The principal advantage is that estimates of stellar mass typically probe a range of the stellar mass function that is somewhat more representative of the stellar mass whereas instantaneous indicators probe only the most massive stars. Furthermore instantaneous measurements can be subject to a greater uncertainty introduced by the effects of dust obscuration.  In \\S 2 we present a compilation of both low and high redshift measurements of the stellar mass density. Using these values, in \\S 3 we derive a best fitting star formation history. We then compare this estimate in \\S 4 to other estimates of the star formation history and highlight any discrepancies. Finally, in \\S 5 we present a discussion of our results and end in \\S 6 with a summary. Throughout this work, we assume a flat $\\Lambda$ CDM cosmology with $\\Omega_{\\Lambda}=0.7$, $\\Omega_{\\rm matter}=0.3$ and $H_{0} = 70\\,\\, {\\rm kms}^{-1}\\,Mpc^{-1}.$ ",
        "conclusions": "In this work a compilation of stellar mass density measurements over the range $0<z<4$ was compiled and converted into a modified form of the Salpeter IMF. A simple analytical parameterisation of the form $\\rho_{*}(z)=a\\,\\times\\,e^{-bz^{c}}$ with $a=0.0023$, $b=0.68$ and $c=1.2$ was found to fit the observations well. Using this compilation a best fitting star formation history was the obtained. This star formation history is well described by the Cole et al. (2001) parameterisation with $a=0.014$, $b=0.11$, $c=1.4$, and $d=2.2$. This star formation history was compared to other indicators of the star formation history including instantaneous measures, the observed rates of core collapse supernovae and the results of recent analysis of the fossil record from the SDSS. Below a redshift of $0.7$  there is good agreement between these estimates and the $1\\sigma$ uncertainty region of our star formation history. At progressively higher redshifts, the stellar mass density and instantaneous indicator inferred star formation histories become inconsistent. Instantaneous measures at $z=3$ imply best fit star formation rates $4$ times larger than those inferred from the stellar mass density.  There are a number of possible causes of this tension. These include principally, uncertainty in the effects of dust on both stellar mass estimates and high redshift star formation rate estimates as well as modeling uncertainties and issues regarding the completeness of the galaxy stellar mass function.  In addition, there are more speculative solutions such as an effective temporal evolution of the initial mass function. We have identified a simple, non-unique model for an evolving IMF that reconciles both the SFH and stellar mass history. Other recent evidence for an evolving IMF has been explored by Dav{\\'e} (2008), and Van Dokkum (2008), who provide different parameterisations. A more vigorous investigation of the implications of an evolving IMF, and suitable choices of parameterisation, is currently underway.  Given the importance of both stellar mass density and star formation rate density measurements to our understanding of the galaxy formation process it is crucial that this discrepancy be resolved. To achieve this, improvements need to be made to measurements, and extensions, such as an evolving IMF, to galaxy formation models need to be implemented. Specifically measurements of the stellar mass density can be improved through the use of larger and deeper surveys to minimise incompleteness, a deeper understanding of the effect of dust attenuation and improvements in both the template fitting procedure and the underlying population synthesis models used to create the templates. Improvements in the star formation history can also be made by refining models of dust attenuation and using larger deeper surveys. Further constraints on the star formation history can also come from improvements in core collapse supernovae rates studies, and more detailed observations of the extragalactic background light, diffuse supernovae neutrino background and the chemical evolution.  \\vskip 20pt  \\noindent We would like to thank the referee Lucia Pozzetti for many helpful comments during the preparation of this document. We would also like to thank again Lucia Pozzetti as well as Leda Sampson and Niv Drory for the provision of an electronic version of their stellar mass density estimates and Ben Panter for providing us with star formation rate estimates. SMW acknowledges support of an STFC studentship and of King's College, NT acknowledges support provided by STFC and AMH acknowledges support provided by the Australian Research Council in the form of a QEII Fellowship (DP0557850)."
    },
    "0801/0801.3905_arXiv.txt": {
        "abstract": "We have studied a sample of 415 associated ($z_{abs}\\sim z_{em}$; relative velocity with respect to QSO $<$ 3000 \\kms) {Mg~II} absorption systems with 1.0$\\le z_{abs}\\le$1.86, in the spectra of Sloan Digital Sky Survey, Data Release 3 QSOs, to determine the dust content, ionization state, and relative abundances in the absorbers.  We studied the dependence of these properties on the properties of the backlighting QSOs and also, compared the properties with those of a similarly selected sample of 809 intervening systems (apparent relative velocity with respect to the QSO of $>$ 3000 \\kms), so as to understand their origin.  Normalized, composite spectra were derived, for absorption line measurements, for the full sample and for several sub-samples, chosen on the basis of the line strengths and other absorber and QSO properties.  Composite absorption lines differ in small but measurable ways from those in the composite spectra of intervening absorption line systems, especially in the relative strengths of {Si~IV}, {C~IV} and {Mg~II}. From the analysis of the composite spectra, as well as from the comparison of measured equivalent widths in individual spectra, we conclude that the associated {Mg~II} absorbers have higher apparent ionization, measured by the strength of the {C~IV} absorption lines compared to the {Mg~II} absorption lines, than the intervening absorbers.  The ionization so measured appears to be related to apparent ejection velocity, being lower as the apparent ejection velocity is more and more positive. Average extinction curves were obtained for the sub-samples by comparing their geometric mean QSO spectra with those of matching (in \\zem $~$and \\imag) samples of QSOs without absorption lines in their spectra.  There is clear evidence for dust-like attenuation in these systems, though the 2175 {\\AA} absorption feature is not present: the extinction is similar to that found in the Small Magellanic Cloud. The extinction is almost twice that observed in the similarly selected sample of intervening systems. We reconfirm with our technique that QSOs with non-zero FIRST (Faint Images of the Radio Sky at Twenty-cm) radio flux are intrinsically redder than the QSOs with no detection in the FIRST survey. The incidence of associated {Mg~II} systems in QSOs with non-zero FIRST radio flux is 1.7 times that in the QSOs with no detection in the FIRST survey. The associated absorbers in radio-detected QSOs which comprise about 12\\% of our sample, cause three times more reddening than the associated absorbers in radio-undetected QSOs. The origin of this excess reddening in the absorbers is indicated by the correlation of the reddening with the strength of Mg II absorption. This excess reddening possibly suggests an intrinsic nature for the associated absorbers in radio-detected QSOs. ",
        "introduction": "Many of the first detected, narrow-line QSO absorption systems had redshifts close to those of their QSOs (e.g. Stockton \\& Lynds 1966; Burbidge, Lynds \\& Burbidge 1966). Such systems, having a relative velocity with respect to the QSO, in units of the speed of light\\footnote{$\\bet={{(1+\\zem)^2-(1+\\zab)^2}\\over{(1+\\zem)^2+(1+\\zab)^2}}$}, $\\beta$, smaller than 0.02, are now termed associated systems. Study of the associated systems is important from the point of view of understanding the energetics and kinematics near the central black hole and also for understanding the ionization structure, dust content and abundances in material directly exposed to the radiation from the QSOs, and in some cases, possibly ejected from, the QSOs or the accretion disks. Associated systems have been suggested to arise (a) in the outer parts of QSO host galaxies (e.g. Heckman et al. 1991; Chelouche et al.  2007) which may have gas properties similar to those in the outer parts of inactive galaxies (Steidel et al. 1997, Fukugita \\& Peebles 2006); (b) by material within 30 kpc of the AGN, accelerated by starburst shocks from the inner galaxy (e.g. Heckman et al. 1990, 1996; D'Odorico et al.  2004; Fu and Stockton 2007a); (c) in the core of the AGN, within 10 pc of the black hole (Hamann et al. 1997a; Hamann et al. 1997b; Barlow \\& Sargent 1997).  In case (a), these are possibly ``halo\" clouds as in normal galaxies, possibly lit up by the QSO to produce extended regions of Lyman alpha emission, but not thought to be moving at the high velocities necessary to explain the dispersion of associated system velocities with respect to the QSO itself. However, using spectra of a set of angularly close QSO pairs, Bowen et al. (2006) confirmed that gas in the outer parts of some QSO host galaxies is detectable as absorption in the spectra of background QSOs, but the spectra of the foreground QSOs did not reveal associated absorption systems: evidently, the appearance of the absorption is dependent on the angle of the line of sight to the spin axis of the accretion disk. The low velocities expected from such gas led to the postulate that clusters of galaxies near the QSO showing associated absorption were responsible for the absorption and the dispersion of cloud velocities.  However, QSOs are not only associated with clusters of galaxies: they appear in a wide range of galaxy types and masses (Jahnke et al. 2004) and, while found in slightly higher density environments (Serber et al. 2006), do not require cluster type densities on large scales (Wake et al. 2004).  In case (b), the gas is typically found to have densities of a few hundred particles cm$^{-3}$, high enough to produce excited fine structure lines of {Si~II} and {C~II}.  There are suggestions in the above references that this is material ejected by a galactic wind,  possibly a superwind from a starburst (Heckman et al. 1990), then lit up by the QSO. (These are often called extended emission line regions, EELR). Cooling flows (Crawford \\& Fabian 1989) no longer seem to be considered in most cases (Fu \\& Stockton 2007b).  In case (c),  absorption is inferred to be very close to the black hole because of variability in the absorption lines, and/or because of the presence of clouds that do not cover the source. Theoretical investigations of case (c) (Arav et al. 1994; Konigl \\& Kartje 1994; Murray et al. 1995; Krolik \\& Kriss 2001; Proga et al. 2000; Everett 2005; Chelouche \\& Netzger 2005 (for a Seyfert galaxy with a much lower radiation field)),  have reinforced the plausibility of thermal or hydromagnetic, radiation assisted flows, both parallel to the AGN jet axis, and perpendicular to it, along the accretion disk. See, for example, Figure 13 of Konigl \\& Kartje (1994), Figure 1 of Richards et al.  (1999), Figure 1 of Murray et al.  (1995) and Figure 1 of Everett (2005).  Early studies concentrated on systems with C IV doublets. Weymann et al. (1979) and Foltz et al. (1986) found a statistical excess of such systems as compared to what was expected if these were randomly distributed in space. Other studies (Young et al. 1982; Sargent et al. 1988) could not confirm these observations. Later it was shown that {Mg~II} systems (Aldcroft et al. 1994) and C IV systems (Anderson et al. 1987; Foltz et al. 1988) with \\bet $<$ 0.0167 are preferentially found in steep-spectrum radio sources, often thought of as sources dominated by emission from lobes rather than the core of the source. Ganguly et al. (2001) showed that high ionization systems (having lines of C IV, N V and O VI) with \\bet $<$ 0.0167 with $z_{abs}<1.0$ are not present in radio-loud QSOs that have compact radio morphologies, flat radio spectra (core dominated sources) and  C IV lines with mediocre FWHM ($\\le 6000$ \\kms). Baker et al.  (2002) from a study of a near-complete sample of low-frequency selected, radio-loud QSOs corroborated trends for C IV associated absorption to be found preferentially in steep-spectrum and lobe-dominated QSOs, suggesting that the absorption is the  result of post-starburst activity and that the C IV lines weaken as  the radio source grows, clearing out the gas and dust. Vestergaard (2003, hereafter V03) studied a sample of high ionization associated systems with $1.5<z_{abs}<3.5$ and found that the occurrence, or not, of such systems is independent of radio properties of the QSOs. V03 did not find the relationship between radio source  size and C IV line strength of Baker et al.  (2002), nor the absence of absorbers in low line width systems of Ganguly et al (2001). V03 did identify weak correlations among QSO properties and the line strengths of the associated absorption systems, that differed from those of the intervening systems, but concluded that differences in the results among the three studies (Baker et al.  2002; Vestergaard 2003 and Ganguly et al. 2001) could probably be attributed to various differences in selection of the relatively small samples (~50-100 systems): among them, the most obvious difference is the optical luminosity of the samples, progressively more luminous in order of the references just cited.  Steep spectrum radio sources, from the above references, have an excess of associated (\\bet $<0.02$) absorbers and are thought to be viewed at a large angle to the jet axis, so the sightline passes over the accretion disk (or torus). Broad absorption line (BAL) systems are often thought to be viewed from the same aspect, but are generally devoid of radio flux. There are suggestions that the associated systems are related in some way to the broad absorption line systems (Wampler et al. 1995; Baker et al. 2002; V03), summarized by V03.  There have also been  suggestions that narrow line QSO absorption line systems (QSOALS) that have high velocities with respect to the QSO, i.e.  apparent velocities of ejection up to tens of thousands of \\kms $~$may also be associated with the QSOs i.e.  they are close to the QSO ($<<$ 10 pc) but have high relative velocity (Vanden Berk et al. 1996; Januzzi et al. 1996; Richards et al. 1999; Richards 2001; V03; Misawa et al. 2007); candidate systems are especially noticeable in the brightest QSOs and in radio sources with flat radio spectra (often said to be core-dominated sources). Richards et al. (2001), using the largest sample of radio QSOs available to date, showed that there is an excess of C IV absorbers at high \\bet $~$values in flat spectrum radio sources (thought to be viewed close to the jet axis), strengthening arguments of Richards et al. (1999) and Richards (2001) that up to 36\\% of the apparently high velocity material is associated with material intrinsic to the background AGN. As shown by Misawa et al. (2007), such material may evidently reach very high outflow velocities, up to and exceeding 20,000 \\kms. Their evidence for the intrinsic  nature of such components is the small filling factor of the absorption lines in covering the background UV radiation source, presumably the full accretion disk (Pereya et al. 2006). A primary question is how to tell if a given system is intrinsic or not, whatever its \\bet $~$value.  In this paper, we use a large and homogeneous sample of associated ($z_{abs}\\sim z_{em}$; relative velocity with respect to the QSO $<$ 3000 \\kms) {Mg~II} systems with $1.0<z_{abs}<1.86$ compiled from the Sloan Digital Sky Survey (SDSS) Data Release 3 (DR3) catalog to determine the average dust content, ionization and relative abundances in these absorbers. Our aim is to determine if these systems are indeed intrinsic to the QSOs by (a) studying the dependence of their average properties on QSO properties and (b) comparing the properties of associated, {Mg~II} systems with those of intervening systems (\\bet $>$ 0.01) selected with similar criteria. In particular, we are looking for a spectroscopic signature to distinguish associated from intervening systems.  We make use of the composite spectra of the sample (and various sub-samples thereof), following the method recently advocated by York et al. (2006, hereafter Y06).  In Section 2, we describe the criteria used for sample selection, various sub-samples generated from the main sample and the method of generating composite spectra. In section 3, we present our results which include (i) a comparison between the properties of the samples of the intervening and associated systems, using several statistical tests; (ii) a discussion of the line strengths in the composite spectra of associated systems compared to those for the intervening sample and of the state of ionization in these systems; (iii) the measured extinction for various sub-samples; (iv) a detailed analysis of the dependence of extinction and other properties on the radio properties of the QSOs; (v) discussion of the abundances in associated sample; and (vi) possible scenarios for the origin and location of the absorbers.  A few systematic effects  that may be hidden in our data are noted in section 4. Conclusions are presented in section 5. ",
        "conclusions": "We have studied a sample of 415 intermediate redshift, {Mg~II} associated systems in the spectra of SDSS QSOs. Our main conclusions are as follows. \\begin{enumerate} \\item There is definite evidence of dust extinction in the associated systems. The average extinction is two times that found in a sample of similarly selected, intervening systems. This larger extinction could be attributed to higher average N$_{\\rm H\\;I}$ +N$_{\\rm H\\;II}$ in associated compared to intervening systems, or to differences in the dust to gas ratio. \\item There is no evidence for the 2175 {\\AA} bump in the extinction curves of the associated absorbers. \\item The extinction curve for these absorbers is similar to that of the SMC, which has also been found to be the case for intervening absorbers. \\item Associated absorbers are in a higher state of ionization compared to the intervening absorbers. \\item In the relative velocity range 0.01$>$\\bet$>$-0.004 studied, the ionization conditions and the total extinction in the associated systems is a function of their apparent relative velocity with respect to the QSO, systems with lower relative velocity being more ionized and more highly reddened. \\item There is no obvious evidence for higher abundances in the associated systems than those in the intervening systems. \\item About a third of the absorption systems do have redshifts higher than the emission redshift of QSOs and thus appear to be infalling. The direct inference would be that lower \\bet $~$systems are closer to the ionization sources. No clues are found that distinguish the material as ambient material (shocked?) in the QSO host galaxies or as returning gas from QSO jet or accretion disk outflows. \\item The QSOs with non-zero radio flux in the FIRST survey are intrinsically redder than QSOs with null detection in the FIRST survey by relative \\ebv $\\sim$0.04, on average, for an SMC extinction curve, even when there are no absorption lines present. \\item Associated systems in QSOs with non-zero radio flux in the FIRST survey have three to four times as much dust as those in QSOs with null detection in the FIRST survey. This excess reddening is correlated to the strength of Mg II absorption lines and appears to originate in the absorbers themselves. These absorbers are thus, significantly different from the intervening systems and may possibly be intrinsic to the QSOs. \\item No clear discriminant between associated and  intervening systems has been found that would definitely work on a  case by case basis, despite the clear but subtle differences in ionization and reddening noted above. ``Intrinsic\" but high ejection velocity systems may be hard to discern among the intervening systems, except through non-detection of a galaxy  at the absorber redshift. \\end{enumerate} \\appendix"
    },
    "0801/0801.2903_arXiv.txt": {
        "abstract": "We study the spectral energy distribution of gamma rays and neutrinos in the precessing microquasar SS433 as a result of $pp$ interactions within its dark jets. Gamma-ray absorption due to interactions with matter of the extended disk and of the star is found to be important, as well as absorption caused by the UV and mid-IR radiation from the equatorial envelopment. We analyze the range of precessional phases for which this attenuation is at a minimum and the chances for detection of a gamma-ray signal are enhanced. The power of relativistic protons in the jets, a free parameter of the model, is constrained by HEGRA data. This imposes limits on the gamma-ray fluxes to be detected with instruments such as GLAST, VERITAS and MAGIC II. A future detection of high energy neutrinos with cubic kilometer telescopes such as IceCube would also yield important information about acceleration mechanisms that may take place in the dark jets. Overall, the determination of the ratio of gamma-ray to neutrino flux will result in a key observational tool to clarify the physics of heavy jets. ",
        "introduction": "The famous and enigmatic microquasar SS433 has been matter of investigation for more than two decades. Consisting of a donor star feeding mass to a black hole, it presents two oppositely directed, precessing jets with hadronic content\\footnote{Iron lines with a shift corresponding to a velocity of $v\\sim 0.26 c$ have been detected, for instance, by Migliari et al. (2002).}. We refer to the relativistic collimated outflows as `dark' jets \\citep{gallo} since the very high kinetic luminosity $L_{\\rm k}\\sim 10^{39}$ erg s$^{-1}$ \\citep{dubner} appears as the dominant power output of the ejected material, having imprinted a deformation on the supernova remnant W50. %  Most of the radiative output of the system is observed in the UV and optical bands, whereas the X-ray emission detected is $\\sim 1000$ lower than the kinetic energy of the jets, probably due to a screening effect with an equatorial outflow \\citep{fabrikaXrays,marshall}. The gamma-ray emission above $0.8$ TeV has been constrained by HEGRA to be $\\Phi_\\gamma<8.93 \\times 10^{-13} {\\rm cm}^{-2}{\\rm s}^{-1}$ \\citep{HEGRA} whereas the neutrino flux upper limit according to AMANDA-II data is $\\Phi_\\nu<0.21 \\times 10^{-8} {\\rm cm}^{-2}{\\rm s}^{-1}$ \\citep{Halzen06}.  In previous hadronic models for high energy emission from microquasars, relativistic protons in the jets interact with target protons from the stellar wind of the companion star \\citep{rom03,hugo02,Or07}. Since in the case of SS433 there is no evidence of such a strong stellar wind, in this work we investigate the possible production of gamma rays and neutrinos resulting from $pp$ interactions between relativistic and cold protons within the jets themselves.   \\section[]{Preliminaries}  The binary SS433, distant $5.5$ kpc from the Earth, displays two mildly relativistic jets ($v_{\\rm b}\\approx 0.26$c) that are oppositely directed and precess in cones of half opening angles of $\\theta\\approx 21^\\circ$. The line of sight makes an angle $i=78^\\circ$ with the normal to the orbital plane and a time-dependent angle $i_{\\rm j}(t)$ with the approaching jet (see Fig. {\\ref{Figrender}}). Assuming that $\\psi(t)$ is the precessional phase of the approaching jet, we shall follow the convention that when $\\psi=0$ the mentioned jet points closer to the Earth. Then, when $\\psi=0.5$, it has performed half of the precession cycle and it makes its largest angle with the line of sight. The mass loss rate in the jets is $\\dot{m}_{\\rm j}= 5\\times 10^{-7}{ M}_\\odot {\\rm yr}^{-1}$, the period of precession is $162$ d and the orbital period is $13.1$ d \\citep{Fabrika100}. The donor star and the compact object are thought to be embedded in a thick expanding disk which is fed by a wind from the supercritical accretion disk around the black hole \\citep{Zwitter}. This equatorial envelope is perpendicular to the jets and according to \\citet{Fabrika100} we assume that it has a half opening angle $\\alpha_{\\rm w} \\approx 30^\\circ $, a mass loss rate $\\dot{M}_{\\rm w}\\approx 10^{-4}{ M}_\\odot{\\rm yr}^{-1} $ and a terminal velocity $v_{\\rm w}\\sim 1500 {\\rm \\ km \\ s}^{-1}$. Also, this extended disk has been recognized as the origin of both the UV and mid-IR emission \\citep{Gies02a,Fuchs05} which can cause significant absorption of gamma-rays as discussed in \\citet{last}.   The spectral identification of the companion star has been difficult due to the presence of the extended disk, since the star is often partially or totally obscured by it. After convenient observations at specific configurations of precessional and orbital phases it has became quite clear that the star is an A-supergiant \\citep{Hillwig,Barnes,Chere}. We assume the masses of the components as derived from INTEGRAL observations \\citep{Chere}, $M_{\\rm bh}= 9 { M}_\\odot$ and $M_\\star= 30 { M}_\\odot$ for the black hole and star respectively. This corresponds to an orbital separation $a \\simeq 79 \\ { R}_\\odot$ for a zero-eccentricity orbit as it is the case for SS433. Since the star is believed to fill its Roche lobe, the implied radius according to \\citet{Eggleton} is $R_L\\simeq 38 { R}_\\odot$.        \\begin{figure} \\includegraphics[trim = 0mm 0mm 0mm 0mm, clip, width=8cm,angle=0]{MN-08-0074-MJ-Fig1.eps} \\caption{Schematic view of the SS433. The \\textit{approaching} jet is most of the time closest to our line of sight and the \\textit{receding} one is oppositely directed.} \\label{Figrender} \\end{figure}   \\subsection{Outline of the jet model} \\textbf{} We assume that a magneto-hydrodynamic mechanism for jet ejection operates in SS433, that is, ejection is realized through the conversion of magnetic energy into matter kinetic energy. The magnetic energy density is supposed to be in equipartition with the kinetic energy density of the ejected particles, so that the corresponding magnetic field along the jet is given by \\be B(z_{\\rm j})= \\sqrt{8\\pi e_{\\rm j}}, \\ee where the kinetic energy density is \\be e_{\\rm j}=\\frac{\\dot{m}_{\\rm j}E_{\\rm k}}{m_{\\rm p}v_{\\rm b}\\pi R_{\\rm j}^2(z_{\\rm j})}. \\ee Here $E_{\\rm k}$ is the classical kinetic energy of a jet proton with velocity $v_{\\rm b}$ and $R_{\\rm j}(z)$ is the jet radius at the height $z_{\\rm j}$ along the jet axis.  The jets are modeled as cones with a half opening angle $\\xi_{\\rm j}\\approx 0.6^\\circ$ \\citep{marshall}. Assuming an initial jet radius $R_0=R_{\\rm j}(z_0)\\approx 5 R_{\\rm Sch}$, where $R_{\\rm Sch}=2GM_{\\rm bh}/{c^2}$, we find the injection point as $z_0= R_0/\\tan{\\xi_{\\rm j}}\\simeq 1.3 \\times 10^{9}$cm. Since the jets are heavy as compared to other similar objects, it is reasonable to admit that they are cold matter dominated. In this case, we assume that a small fraction of relativistic or hot particles are confined by the cold plasma. According to \\citet{broadband} the pressure of cold particles is greater than that of the relativistic ones if the ratio of cold to hot particles is less than 1/1000, and this condition will be greatly satisfied provided that the luminosity carried by relativistic particles is required to be smaller than the total kinetic luminosity of the jet.  Particle acceleration is supposed to take place via diffusive acceleration by internal shocks converting bulk kinetic energy into random kinetic energy. According to the standard model for non-relativistic shock acceleration (e.g. Blandford \\& Eichler 1987 and references therein) we expect that the relativistic proton spectrum is given by a power-law, $N'_p({E'}_p)= K_0 {E'}_p^{-\\alpha}$ at $z_{\\rm j}=z_0$, where the spectral index is the standard value for first order diffusive shock acceleration, $\\alpha=2$. The flux of these protons hence evolves with $z_{\\rm j}$ as $$J'_p({E'}_p)= \\frac{cK_0}{4\\pi} (z_0/z_{\\rm j})^2 {E'}_p^{-\\alpha}$$ in the jet frame, which transformed to the observer frame \\citep{Purmohammad} reads  \\begin{multline} J_p(t,E_p,z_{\\rm j})=\\frac{c K_0}{4 \\pi} \\left(\\frac{z_0}{z_{\\rm j}}\\right)^ 2 \\times  \\label{Jplab} \\\\ \\frac{\\Gamma^{-\\alpha+1} \\left(E_p-\\beta_{\\rm b} \\sqrt{E_p^2-m_p^2c^4} \\cos i_{\\rm j} \\right)^{-\\alpha}}{\\sqrt{\\sin ^2 i_{\\rm j} + \\Gamma^2 \\left( \\cos i_{\\rm j}(t) - \\frac{\\beta_{\\rm b} E_p}{\\sqrt{E_p^2-m_p^2 c^4}}\\right)^2}}  \\\\ \\equiv \\left(\\frac{z_0}{z_{\\rm j}}\\right)^2 \\tilde{J}_p(E_p,t), \\end{multline} where $i_{\\rm j}(t)$ is the angle between the jet axis and the line of sight, $\\beta_{\\rm b}= 0.26$, and $\\Gamma=\\left[ 1-\\beta_{\\rm b}^2 \\right]^{-1/2}$ is the jet Lorentz factor. The normalization constant $K_0$ is obtained by specifying the fraction of power carried by the relativistic protons, $q_{\\rm rel}$, \\be \\pi R_0^2\\int_{E_p'^{\\rm(min)}}^{E_p'^{\\rm(max)}} J'_p(E'_p) E'_p dE'_p= q_{\\rm rel}L_k, \\ee so that \\be K_0 = \\frac{4 q_{\\rm rel} L_k}{c  R_0^2\\ln \\left(\\frac{{E'}_{p}^{\\rm(max)}}{{E'}_p^{\\rm(min)}}\\right)}, \\label{K0norm} \\ee where we take ${E'}_p^{\\rm(min)}\\approx 1$ GeV and the maximum proton energy ${E'}_p^{\\rm(max)}$ will be determined in the next section. We shall adopt, for the illustrative predictions of neutrino and gamma-ray fluxes, a tentative value $q_{\\rm rel}= 10^{-4}$, but a full discussion of the possible range for this parameter will be presented in Sect. \\ref{neutrinos}. ",
        "conclusions": "We have studied the high-energy emission originated in the dark jets of the microquasar SS433. A small fraction of its particle contents are relativistic protons that collide with the cold ions within the jets, producing gamma rays and neutrinos after pion decay processes. We found that up to distances $\\sim 10^{12}{\\rm cm}$ from the black hole, protons with energies below $\\sim 3\\times 10^{6}$GeV will cool dominantly via $pp$ interactions. The ratio of the power carried by relativistic protons to the total kinetic power of the jet, $q_{\\rm rel}$, was kept as a free parameter of the model (for illustrative purposes we have used $q_{\\rm rel}= 10^{-4}$ in the figures).  We have calculated the spectrum of gamma rays and high-energy neutrinos based on the formulae given by \\citet{Kelner06}. We have considered the cooling of the charged pions and muons produced, and we have found that the high-energy neutrino emission is attenuated by synchrotron losses. Adding the contribution from both jets, we have obtained the total gamma-ray spectral intensity of SS433. The gamma radiation will be largely absorbed while leaving the inner regions the system by means of several processes. This will mainly occur through interaction with matter of the star and extended disk, leading to significant photopair production.  UV and mid-IR photons originated in the extended disk are also expected to cause important absorption via $\\gamma\\gamma$ annihilations.    The total optical depth is found to depend on the precessional phase in such a way that when the approaching jet is pointing away from the Earth, at $\\psi\\sim 0.5$, the extended disk blocks the emitting region and the absorption is strongest. In particular, in the range of precessional phases between $\\psi \\gtrsim 0.91$ and $\\psi \\lesssim 0.09$, the gamma rays originated at the base of the jets will travel to the Earth without having to pass through the equatorial disk. With the conservative assumption that this outflowing disk presents a large half opening angle $\\alpha_{\\rm w}=30^\\circ$, the mentioned range of favorable precessional phases gives a total of $\\sim 29$ d for optimal detectability. Since, according to \\citet{Gies02b}, $\\psi=0$ occurred on 2002 June 5, it follows that the next upcoming opportunities to achieve detection will be centered around the following dates every $162$ d: 2008 August 20, 2009 January 29, 2009 July 10, etc. As mentioned above, the exact duration along which the favorable conditions may hold, depends on the half opening angle of the extended disk, which might be smaller than what was assumed. In that case, the observational window would be broader  The observations by HEGRA imply a maximum value for the free parameter $q_{\\rm rel}$ in $\\sim 3 \\times 10^{-4}$. Given the expected sensitivity of the next km$^3$ neutrino telescopes generation, it will be possible to test our model % down to values $\\sim 5\\times 10^{-5}$ in three years of operation. An extended range of this parameter will be probed by gamma-ray observations with GLAST and Cherenkov telescopes, especially if performed on the favorable dates.  We conclude that its dark jets can be the possible site for both gamma-ray and neutrino production in SS433. Since most of the high-energy flux is generated in the inner jets, gamma-ray absorption will make detection with Cherenkov telescopes like VERITAS and MAGIC-II difficult but not impossible if attempted at the favorable dates when the approaching jet is closest to the line of sight. Actually, there are much better prospects for gamma-ray observation with GLAST and neutrino detection with IceCube also seems promising. % The determination of the gamma-ray to neutrino flux ratio would allow to estimate $q_{\\rm rel}$ unambiguously, yielding crucial information about acceleration mechanisms taking place in the jets."
    },
    "0801/0801.2768_arXiv.txt": {
        "abstract": "Mass segregation stands as one of the most robust features of the dynamical evolution of self-gravitating star clusters. In this paper we formulate parametrised models of mass segregated star clusters in virial equilibrium. To this purpose we introduce mean inter-particle potentials for statistically described unsegregated systems and suggest a single-parameter generalisation of its form which gives a mass-segregated state. We describe an algorithm for construction of appropriate star cluster models. Their stability over several crossing-times is verified by following the evolution by means of direct $N$-body integration. ",
        "introduction": "Observations show quite often an increased concentration of massive stars towards the centres of young star clusters (e.g. ONC -- Hillenbrand \\& Hartmann~1998; NGC 2157 -- Fischer et al.~1998; NGC 3603 -- Stolte et al.~2006). This tendency, known as mass segregation, can be of different origin: Initial mass segregation is sometimes considered (e.g. Murray \\& Lin~1996, Bonnell \\& Bate~2006) as a consequence of the formation of massive stars preferably in the densest regions (i.e. the cores) of the parent gas clouds. On the other hand, the process of mass segregation is also known to be one of the most robust features of the two-body relaxation driven evolution of self-gravitating star clusters (Chandrasekhar~1942, Spitzer~1969).   Several approaches were developed to setup a star cluster in the state of mass segregation. Gunn \\& Griffin~(1978), Capuzzo Dolcetta et. al~(2005) and others based their setup on multi-component King models (King~1965, Da~Costa \\& Freeman~1976) with stars separated into several mass classes which interact with each other via smoothed potentials. This approach relies on solving of non-linear set of Poisson equations, which is possible for limitted number of components. A multimass models of star cluster with exact energy equipartition in the core, which also leads to mass segregation, was introduced by Miocchi~(2006). Another approach used e.g. by McMillan \\& Vesperini~(2007) relies on segregation produced by N-body integration of initially unsegregated systems towards the segregated state, i.e. it is equivalent to a simple redefinition of time $t=0$.  In this paper we describe a new class of models of star clusters with continuous stellar mass distributions and a parametrised degree of mass segregation. The models are motivated by a study of the process of mass segregation during dynamical evolution of a self-gravitating cluster, which is briefly described in the following Section. We show that mass segregation strongly manifests itself in the energy space. In Section~\\ref{sec:model} we introduce convenient characteristics of a statistically described ensemble and derive their form for the unsegregated state. We further introduce in a heuristic manner an alternative, single-parameter form of these quantities that gives constraints on the distribution function of a mass segregated system. Afterwards, we describe an algorithm for construction of the corresponding star cluster. In Section~\\ref{sec:tests} we demonstrate the stability of the models by means of $N$-body integrations. Finally, Section~\\ref{sec:conclusions} contains our conclusions. ",
        "conclusions": "\\label{sec:conclusions} We have introduced a way of parametrisation of self-gravitating systems in terms of mean inter-particle potentials. We have demonstrated that this approach can be used for construction of quasi-stationary models of mass segregated star clusters. For the sake of simplicity, we have performed tests with simple power-law mass function. Nevertheless, the approach does not depend on a particular form of the mass function and the standard IMF (Kroupa~2001, 2007) can be used as an input. Notice also that even for a cluster of equal mass stars the algorithm will lead to a system with a desired level of `energy segregation' which is in general the process that drives the star clusters towards core collapse.  Finally, let us remark that the index of mass segregation is likely to be related to the entropy. For $S=0$ the system is highly symmetric in terms of $\\mmean{U^{ij}}$. On the other hand, in the limit of $S=1$ it is required that all binding energy is carried by the two most massive particles, which is usually considered a state of maximal entropy of the self-gravitating system. We suggest that the statistical approach based on characterisation of the system by mean values of suitable physical quantities related to subsets of stars deserves further investigation, providing us, hopefully, with a deeper understanding of the thermodynamics of star clusters."
    },
    "0801/0801.2898_arXiv.txt": {
        "abstract": "Recent evidence of a young progenitor population for many Type-Ia SNe (SNe-Ia) raises the possibility that evolved intermediate-mass progenitor stars may be detected in pre-explosion images. NGC~1316, a radio galaxy in the Fornax cluster, is a prolific producer of SNe-Ia, with four detected since 1980. We analyze Hubble Space Telescope (HST) pre-explosion images of the sites of two  of the SNe-Ia that exploded in this galaxy, SN2006dd (a normal Type-Ia) and SN2006mr (likely a subluminous, 1991bg-like, SN-Ia). Astrometric positions are obtained from optical and near-IR ground-based images of the events. We find no candidate point sources at either location, and set upper limits on the flux in $B$, $V$, and $I$ from any such progenitors. We also estimate the amount of extinction that could be present, based on analysis of the surface-brightness inhomogeneities in the HST images themselves. At the distance of NGC~1316, the limits correspond to absolute magnitudes of $\\sim -5.5$, $-5.4$, and $-6.0$ mag in $M_B$, $M_V$, and $M_I$, respectively. Comparison to stellar evolution models argues against the presence at the SN sites, 3 years prior to the explosion, of normal stars with initial masses $\\ga 6M_{\\odot}$ at the tip of their asymptotic-giant branch (AGB) evolution, young post-AGB stars that had initial masses $\\ga 4M_{\\odot}$, and post-red-giant stars of initial masses $\\ga 9 M_{\\odot}$. ",
        "introduction": "Although there is general consensus that Type-Ia supernova (SN-Ia) explosions are the result of a thermonuclear runaway in a degenerate, approximately Chandrasekhar-mass ($M_{\\rm ch}$), carbon-oxygen stellar core (e.g., Hillebrandt \\& Niemeyer 2000; Badenes et al. 2006; Mazzali et al. 2007), the progenitor systems of  SNe-Ia are still unknown. The two main models leading to a SN-Ia explosion are the single-degenerate (SD) model, in which a white dwarf (WD) accretes matter from a close companion until it approaches the Chandrasekhar limit (Whelan \\& Iben 1973; Nomoto 1982), and the double-degenerate (DD) model, involving the merger of two WDs (Iben \\& Tutukov 1984; Webbink 1984). However, both scenarios face theoretical and observational challenges. The outcome of a DD merger may be accretion-induced core collapse, rather than a SN-Ia (e.g., Nomoto \\& Iben 1985; Saio \\& Nomoto 2004, Guerrero et al. 2004), although there are opposing views that invoke rotation of the stellar surface to prevent this outcome (e.g., Piersanti et al. 2003). An observational search for DD progenitor systems (Napiwotzki et al. 2004; Nelemans et al. 2005) has turned up, among $\\sim 1000$ WDs surveyed, few or no potential DD systems with a total mass exceeding $M_{\\rm ch}$ that will merge within a Hubble time.  The SD scenario, in turn, has been criticized (e.g., Cassisi et al. 1998; Piersanti et al. 1999, 2000) for its assumptions about the existence of {\\it ad hoc} mechanisms that regulate the accretion flow on to the WD (e.g. Hachisu et al. 1996). Observationally, Badenes et al. (2007) have noted the absence, in seven nearby SN-Ia remnants, of the signatures of the strong wind from the accretor that supposedly stabilizes the accretion flow, and permits reaching $M_{\\rm ch}$. Prieto et al. (2007) have found no evidence for a low-metallicity threshold in SN-Ia hosts, in contrast to the predictions by Kobayashi et al. (1998; see also Kobayashi \\& Nomoto 2008) of such a threshold, due to a minimum metallicity that is required for the wind regulation mechanism to be effective. While evidence for circumstellar material, consistent with expectations from a wind from a red-giant companion, has been found in one recent normal SN-Ia (Patat et al. 2007), such material is not observed in other events (Mattila et al. 2005; Simon et al. 2007; Leonard 2007). For the remnant of Tycho's SN, which was a type-Ia (Badenes et al. 2006), there have been conflicting claims about the identification and nature of a remaining companion star (Ruiz-Lapuente et al. 2004;  Fuhrmann 2005; Ihara et al. 2007). If, in the end, no companion is found, this would be another problem for the SD picture. Recent metal abundance measurements suggest a type-Ia explosion also in Kepler's SN remnant (e.g., Blair et al. 2007; Reynolds et al. 2007) but the velocity of the remnant away from the Galactic plane, together with evidence for a circumstellar medium from a massive progenitor, are problematic the context of a binary progenitor system (Reynolds et al. 2007).  For core-collapse SNe, which derive from massive ($\\ga 8 M_{\\odot}$) stars, direct searches of pre-explosion images for progenitors have had a number of successes -- in the case of SN1987A (West et al. 1987; White \\& Malin 1987), and several other events with suitable {\\it Hubble Space Telescope} (HST) data (Maund \\& Smartt 2005; Li et al. 2006; Hendry et al. 2006; Gal-Yam et al. 2007). In contrast, the expectation, based on the SD and DD models, that SNe-Ia explode in old, low-mass, systems, has discouraged such direct searches for progenitors of SNe-Ia in the HST era. While there have been some studies of SN-Ia environments using post-explosion HST data (Van Dyk et al. 1999), the reverse has not been done, largely due to a preconception that the progenitors would be undetectable. (Interestingly, Van Dyk et al. 1999 noted young stellar populations in the vicinity of the four SNe-Ia they studied, but subtracted them off the images, arguing that the SNe-Ia were necessarily derived from an old population).  However, several recent observational developments, in addition to the problems faced by the popular models, suggest that such a search may be useful after all. Mannucci et al. (2005) have measured stellar-mass-normalized type-Ia SN rates as a function of galaxy Hubble type and galaxy colors, and have found that the SN-Ia rate in star-forming galaxies traces the star formation rate. Such a correlation directly implies that at least some of the progenitors are young stars. This study and several subsequent ones (Scannapieco \\& Bildsten 2005; Mannucci et al. 2006; Sullivan et al. 2006; Dilday et al. 2008) have demonstrated that there must be a wide distribution of delay times between star formation and SN explosion, as predicted by some progenitor models (e.g., Greggio \\& Renzini, 1983; Greggio 2005). The spread in the delay times can be described in terms of two populations of progenitors: a ``prompt'' one, which dominates the SN-Ia rate in star-forming galaxies and has  a typical delay times of less than $\\sim 10^8$~yr; and ``tardy'' population, having a SN-Ia delay time of $\\ga 10^9$~yr (or a large-delay tail to the same progenitor population above), needed to produce also the SN-Ia rate measured in old stellar populations with no current star formation. An outstanding puzzle is the possible decrease in the SN-Ia rate at redshifts $z>1$ measured by Dahlen et al. (2004, 2008), at a cosmic time when the star-formation rate was an order of magnitude higher than today. The SN-Ia rate, which would therefore have then been dominated by the prompt population, would be expected to track the star formation rate, as the latter flattens or continues to rise to high $z$. On the other hand, reanalysis of some of these data by Kuznetsova et al. (2008), and independent high-$z$ rates by Poznanski et al. (2007), have questioned the significance of the claimed decrease in SN-Ia rate at high redshift, pointing to the need for further observations to resolve the issue.  In any event, the formation of a degenerate carbon stellar core within $10^8$~yr for a significant fraction of SN-Ia progenitors points to stars with zero-age main sequence (ZAMS) masses of $\\ga 5 M_\\odot$ (e.g. Girardi et al. 2000). During their post-main-sequence evolution, such intermediate-mass stars may reach luminosities high enough to make them potentially detectable in pre-explosion HST images of nearby host galaxies. Such evolved stars could be present in SN-Ia progenitor systems, either in the role of mass donors to the WD (previously produced by a more massive primary), where the evolutionary timescale of the secondary dictates the onset of the explosion; or in the role of single-star SN-Ia progenitors, where degenerate carbon-core ignition occurs in an evolved star that has somehow lost its hydrogen envelope (e.g. Tout 2005; Waldman, Yungelson, \\& Barkat 2007).  Furthermore, recent measurements of extragalactic SN-Ia rates at various redshifts have yielded high rates, especially in star-forming environments. Maoz (2008) has recently compiled, and compared self-consistently, different measurements of various observables related to SN rates, and used them to estimate the fraction of intermediate-mass close binaries that explode as SNe-Ia through the prompt and tardy channels. He shows that the high SN-Ia rates (as well as other independent observables such as intracluster abundances) indicate that most or all close, intermediate-mass, binaries must explode as SNe-Ia, with, possibly, an excess of SNe-Ia over progenitor systems. This conclusion holds despite the uncertainties in the initial mass ranges that lead to viable WD SN-Ia progenitors in the SD and DD scenarios, and the uncertainties in the initial parameters of binaries (binarity fraction, mass ratio distributions, separation distributions). The conclusion stands in contrast to detailed SD and DD models that predict a small exploding fraction among the intermediate-mass population. The independent observation that a major part of the SN-Ia population derives from a young population, indicating initial masses in a narrow range of $\\sim 5-8 M_\\odot$ (see above), further culls the pool of potential progenitors, and makes the case for an excess of SN-Ia explosions compared to progenitors almost unavoidable. A possible solution could, again, be evolved, single-star, stripped envelope, SN-Ia progenitors. Direct searches for such stars in pre-explosion images of the sites of SNe-Ia in nearby galaxies can test this scenario.  In this paper, we perform such a search in the massive nearby galaxy NGC~1316, a prolific SN producer that hosted two SN-Ia events in 2006, by analyzing deep pre-explosion HST images.  After compiling some facts about this galaxy and its SNe, we describe the archival HST and ground-based data we have used, and their analysis for setting flux limits on individual progenitors for each of the two SNe. We then discuss the physical limits on SN-Ia progenitors imposed by these results. ",
        "conclusions": "In this section, we now examine, based on stellar evolution models, which stellar progenitors present 3 years prior to the explosion can be ruled out by our limits. Metallicity that is close to Solar has been measured in this galaxy, both for the stars Goudfrooij et al. (2001a,b) and for the gas Kim \\& Fabbiano (2003).  \\begin{figure*} \\includegraphics[width=0.85\\textwidth]{n1316fig3color.ps} \\caption{Model stellar evolution tracks from ZAMS to AGB (solid curves, from Girardi et al. 2000, for $3-6M_\\odot$, and from Salasnich et al. 2000, for $8-12 M_\\odot$), for the TP-AGB phase (Marigo \\& Girardi 2007; dotted curves), and for post-AGB evolution (Bloecker 1995b; dashed curves), shown in the plane of absolute $I$-band AB magnitude, $M_I$, versus $V-I$ and $B-I$ AB colours. Models are labeled according to their ZAMS masses, in Solar units, with smaller labels for the post-AGB tracks. The $3M_\\odot$ post-AGB model is one which, among the models presented by Bloecker (1995b), has a post-AGB mass of $0.625 M_\\odot$. The grey areas show the regions of parameter space with luminosities greater than our absolute magnitude limits in $B$, $V$, and $I$, for stars at the positions of the SNe in NGC~1316 (assuming a distance of 19~Mpc and no extinction).} \\label{tracks} \\end{figure*}  Figure~\\ref{tracks} shows, in a color-absolute-magnitude diagram, various Solar-metallicity stellar evolution tracks. In the range $3-6 M_\\odot$ ZAMS mass, the tracks are from Girardi et al. (2000), from the main sequence to the onset of thermal pulsations on the AGB branch. The models shown include convective overshooting. Recent models by Marigo \\& Girardi (2007) for the thermal-pulsation stage on the AGB are shown for $3-5 M_\\odot$ ZAMS masses. For $8-12 M_\\odot$ stars, the models are by Salasnich et al. (2000) between the ZAMS and carbon core ignition, assuming Solar composition. The dashed curves show post-AGB tracks by Bloecker (1995b) for stars of ZAMS masses $3-7 M_\\odot$. Bolometric luminosities and effective temperatures of the models were converted to broad-band AB magnitudes by integrating the surface fluxes in stellar-atmosphere model spectra over the appropriate bandpasses. The solar-metallicity stellar atmospheres for a grid of temperatures and surface gravities were interpolated to the temperatures and gravities of the stellar evolution tracks. The stellar atmospheres are from the BaSeL 2.2 database (Lejeuene et al. 1997, 1998; Westera et al. 2002), and from Kurucz (1992) for temperatures above $4000 K$. For the post-AGB models, monochromatic luminosities were obtained assuming a blackbody spectrum, an approximation which is accurate to a few percent for these hot stars and the optical bands considered here. Models are labeled according to their ZAMS masses, in Solar units. The grey-shaded areas show the regions of parameter space with luminosities greater than our absolute magnitude limits in $B$, $V$, and $I$, assuming no extinction, as listed in Table~1. Adoption of the plausible extinctions listed in Table~1 would shift the grey region upward by 0.1~mag, and to the left by 0.15~mag (left panel) and 0.3~mag (right panel).  We see that our upper limits on the stellar luminosities at the sites of the two SNe, do probe some of the parameter space, even if small, of the late-time evolution of intermediate-mass stars. Specifically, according to the models shown, stars with ZAMS masses $\\ga 6M_\\odot$ reach luminosities higher than our limits at the end of their AGB phases. Stars with initial masses $\\ga 9 M_\\odot$ spend much of the time from the first ascent to the red-giant branch and onward at luminosities above the limits. Similarly, post-AGB stars of ZAMS masses $\\ga 4 M_\\odot$, at the beginning of their constant luminosity path to high temperatures on the HR diagram (which results in a monotonic decrease in optical luminosity, see Fig.~\\ref{tracks}), may be above our upper limits.  However, the uncertainties that affect models of the later stages of stellar evolution, particularly uncertainties due to convective overshooting and the details of mixing and mass loss (e.g., Herwig 2005), make the conclusions based on comparison to such models tentative. Indeed, Pauldrach et al. (2004), when modeling the ultraviolet spectra of nine central stars of planetary nebulae (CSPNs), have found severe discrepancies between the masses deduced for five of the stars, compared with those predicted by the mass-luminosity relation from theoretical post-AGB stellar models. Remarkably, the masses of these stars are close to the Chandrasekhar limit. Conversely, Napiwotzki (2006) and Gesicki \\& Zijlstra (2007) have concluded, based, on different analysis methods, that CSPNs, including some of those analyzed by Pauldrach et al. (2004), have masses of $\\sim 0.6 M_\\odot$, similar to most WDs, and with mass-luminosity relations consistent with post-AGB model predictions. In any case, in standard SD models of SNe-Ia, complex binary evolution precedes a SN-Ia explosion, and hence using single-star models to predict the colours and the luminosities of the donor stars could be inaccurate. And, if SNe-Ia  evolve from single stars within some initial mass and metallicity ranges, which have somehow (e.g. via mass loss, or previous binary interactions) lost their hydrogen envelopes, then the stellar tracks for normal AGB evolution that we have used are unlikely to provide reliable predictions for the properties of such progenitors. It should also be kept in mind that AGB stars, which are prodigious dust producers during their thermal-pulse stages, are likely enshrouded in dust shells at the end of their evolution, and the exact point at which they become visible again on the post-AGN branch is uncertain (e.g., Bloecker 1995a,b).  In summary, we have searched pre-explosion images of the sites of two SNe-Ia for potential luminous stellar progenitors, with null results. As always is the case in such studies, our null result does not conclusively rule out any of the scenarios or mass ranges for SN-Ia progenitors. Under any of the possibile scenarios of strong obscuration by dust (by unresolved clumps, or by a screen that is behind the bulk of the stellar mass of the galaxy, or by a circumstellar shell), more luminous progenitors could be hidden. Furthermore, model predictions for the later stages of stellar evolution are still quite uncertain, all the more so if binary interactions take place. Finally, SNe 2006dd and 2006mr may have, by chance, belonged to the tardy SN-Ia population, in which case one would not expect them to be associated with a young population.  Nevertheless, our study shows that deep HST observations of nearby SN hosts can graze some interesting regions of stellar parameter space that may be relevant for young-population SN-Ia progenitors. Future SNe-Ia will explode also in more nearby galaxies that have HST imagery. This will improve the statistics of such searches, will make accessible lower stellar luminosities, and will perhaps eventually reveal an actual progenitor. In view of the many unknowns behind SN-Ia formation, such observational studies may provide valuable clues."
    },
    "0801/0801.2917_arXiv.txt": {
        "abstract": "We report limits in the planetary-mass regime for companions around the nearest single white dwarf to the Sun, van Maanen's star (vMa\\,2), from deep $J$-band imaging with Gemini North and {\\it Spitzer} IRAC mid-IR photometry. We find no resolved common proper motion companions to vMa\\,2 at separations from $3 - 45\\arcsec$, at a limiting magnitude of $J\\approx23$. Assuming a total age for the system of $4.1\\pm1$\\,Gyr, and utilising the latest evolutionary models for substellar objects, this limit is equivalent to companion masses $>7\\pm1\\,\\mjup$ ($T_{\\rm eff}\\approx300$\\,K). Taking into account the likely orbital evolution of very low mass companions in the post-main sequence phase, these $J$-band observations effectively survey orbits around the white dwarf {\\it progenitor} from $3 - 50$\\,AU. There is no flux excess detected in any of the complimentary {\\it Spitzer} IRAC mid-IR filters We fit a DZ white dwarf model atmosphere to the optical $BVRI$, 2MASS $JHK$ and IRAC photometry. The best solution gives $T_{\\rm eff}=6030\\pm240$K, log~g$=8.10\\pm0.04$ and, hence, $M= 0.633\\pm0.022\\Msun$. We then place a $3\\,\\sigma$ upper limit of $10\\pm2\\,\\mjup$ on the mass of any unresolved companion in the $4.5\\mu$m band. ",
        "introduction": "Direct imaging of extra-solar planetary-mass companions to solar-type stars is complicated by problems of contrast and resolution. At the time of writing, no planet has been directly imaged around a solar-type star. An alternative solution to these difficulties is to target intrinsically faint stars instead, such as white dwarfs. Stellar evolution lends two huge advantages when searching for very faint companions: white dwarfs are up to $\\sim10^4$ times fainter than their main sequence progenitors, and the orbits of any planetary-mass companions that lie outside the stellar envelopes during the giant phases will expand outwards as mass is lost from the central star, increasing the projected separation by a maximum factor $M_{\\rm MS} / M_{\\rm WD}$ \\citep{jeans}. Thus, the problems of contrast and resolution are greatly reduced. The possible evolution of planetary systems in the post-main sequence phase is discussed in more detail by \\citet{duncan}, \\citet*{burleigh02}, \\citet{Debes02} and \\citet{villaver}.  The direct detection of such low mass companions to white dwarfs opens up the possibility for spectroscopic investigation of a previously unobserved class of object: evolved low-mass brown dwarfs and planetary-mass gas giants as low in temperature as $T_{\\rm eff}\\approx300$\\,K. In contrast, the directly imaged planetary-mass ($\\approx5\\mjup$) companion to the brown dwarf 2MASSW\\,J1207 has the spectrum of a mid-L~dwarf \\citep{chauvina,chauvinb}, because it is still young ($\\sim10^7$~years old). The coolest known brown dwarf, ULAS\\,J$003402.77-005206.7$, has a temperature $600$K$< T_{\\rm eff} < 700$K \\citep{warren07} and a spectral type T\\,8.5. The letter ``Y'' has been suggested for the next, cooler, spectral type \\citep{Kirkpatrick}, which might include evolved planetary-mass companions to white dwarfs. Alternatively, if no obvious spectral change triggers the use of a new letter, then the T classification will need to be extended beyond T\\,8.5.     The idea of using white dwarfs to find intrinsically faint, low mass companions is not new. \\citet{probst83} and \\citet{becklin88} used the low luminosity of white dwarfs to search for brown dwarf companions as near-infrared photometric excesses, and indeed the latter achieved success with GD\\,165\\,B (L\\,4). More recently, \\citet{fbz05} have conducted a comprehensive search for brown dwarf companions to several hundred white dwarfs but detected only one new pair (GD\\,1400, L\\,$6-7$, \\citealt{GD1400}, \\citealt{dobbie05}), while \\citet{maxted06} and \\citet{wd0137b} have detected a close L\\,8 companion to the white dwarf WD\\,0137$-$349. No other detached substellar companions to white dwarfs are known.    The most intriguing circumstantial evidence for the existence of old planetary systems that have survived to the final stage of stellar evolution comes from the discovery of metal-rich circumstellar dust and gas disks around a growing number of white dwarfs \\citep{zuckerman87,becklin05,gd362,gd56,sdss1228,vh07, Jura07,WD1150}. {\\it Spitzer} mid-infrared spectroscopy has now revealed that the dust disks are composed largely of silicates \\citep{Reach,JuraGD362}. The favoured explanation for the origin of this material is the tidal disruption of an asteroid that has wandered into the Roche radius of the white dwarf \\citep{Jura1}, perhaps through interaction with planets in a solar system that has become destabilized in the wake of the planetary nebula phase \\citep{Debes02}. \\citet{Jura2} suggests that at least 7\\% of white dwarfs possess asteroid belts. If that is indeed the case, it is likely that at least this number of white dwarfs also possess planetary systems.  The recent discovery of a $M {\\rm sin} i = 3.2\\mjup$ planet in a $1.7$\\,AU orbit around the extreme horizontal branch star V391\\,Pegasi by \\citet{V391Peg} proves that such objects can survive red giant branch evolution, strongly suggesting that it is simply a matter of time before a planet is discovered around a white dwarf.  \\citet{burleigh02}  made predictions concerning the likely near-infrared brightness of putative resolved, giant planetary companions to nearby white dwarfs, based on their likely total ages and distances. In 2002 we initiated a programme to search for wide, spatially resolved very low mass common proper motion companions to white dwarfs via direct imaging. In particular, we aim to find substellar companions with masses greater than a few~$\\,\\mjup$ around white dwarfs within $\\approx20$~pc of the Sun. Such companions are expected to have near-IR magnitudes brighter than $J\\sim23.5$, commensurate with the expected sensitivity of an 8m telescope in a one hour exposure.  \\begin{sloppypar} We christened our project {\\it ``DODO''} -- {\\it D}egenerate {\\it O}bjects around {\\it D}egenerate {\\it O}bjects. Preliminary results and progress reports have been published elsewhere \\citep*{dodoa,dodob,dodoc}. Here, we report our results for the nearest single white dwarf to the Sun in our sample, van Maanen's star (vMa\\,2, WD\\,0046$+$051, d$\\,= 4.41$\\,pc, \\citealt{Hipparcos}). We combine two epochs of deep, ground-based $J$-band imaging from the 8m Gemini North telescope with {\\it Spitzer} mid-infrared photometry to place limits on the masses and temperatures of any common proper motion and unresolved ultra-cool substellar and planetary-mass companions. \\end{sloppypar}  \\subsection{vMa\\,2 and previous infrared observations}  vMa\\,2 was discovered serendipitously by \\citet{vMa2} in a survey for common proper motion companions to an unrelated star, HD\\,4628. It has a cool, helium-rich atmosphere and strong resonance Ca\\,II H \\& K lines in its optical spectrum, and is classified as a DZ white dwarf. The heavy elements are expected to sink below the photosphere on a timescale much shorter than the white dwarf cooling time. Their presence has been explained previously in terms of episodic accretion from the interstellar medium \\citep{Dupuis92}, but this scenario requires the accretion rate of hydrogen to be at least two orders of magnitude lower than the metals (\\citealt{Wolff02, Dupuis93, Dufour}). Alternatively, the presence of metals in DZs might be explained from accretion of cometary material or tidally disrupted asteroids and planets. However, no DZ has been identified with an infrared excess due to dust emission, possibly because the diffusion timescale for heavy elements in cool white dwarf helium-rich atmospheres is long and the material may have been accreted as much as $\\sim10^6$~years ago. Nonetheless, DZ white dwarfs are a speculative candidate for hosting  old planetary systems and vMa\\,2, as the nearest single white dwarf to the Sun, is an ideal target for such a search.  Through analysis of {\\it Hipparcos} data, \\citet{Makarov} claimed to have detected astrometrically a $0.06 \\pm 0.02\\msun$ substellar companion to vMa\\,2, with an orbital period of 1.57~years and a maximum separation on the sky of $0.3\\arcsec$. \\citet*{fbm} carried out a search for this companion by direct imaging with adaptive optics in the mid-infrared $L'$ band. They also looked for an unresolved companion as a near- and mid-infrared photometric excess in ground-based and {\\it ISO} photometry. They refuted the existence of Makarov's companion, and of any excess emission due to dust, and placed a limit of $T_{\\rm eff} \\la 500$\\,K on the temperature of any putative substellar companion.  \\begin{table*} \\caption{Adopted parameters for van Maanen's star} \\begin{center} \\begin{tabular}{cccccccccc} \\hline $\\mu\\,^a$ & $\\theta\\,^a$ & d$\\,^a$ & $T_{\\rm eff}\\,^b$ & log~$g\\,^b$ & $M_{\\rm WD}\\,^b$ & $t_{\\rm WD}\\,^c$ & $M_{\\rm MS}\\,^d$ & $t_{\\rm MS}\\,^e$ & $t_{\\rm total}$ \\\\ (mas/yr) & (mas/yr) & (pc) & (K) & & ($\\msun$) & (Gyr) & ($\\msun$) & (Gyr) & (Gyr) \\\\ \\hline \\hline 1231.72 & $-2707.67$ & 4.41 & 6030 (240) & 8.10 (0.04) & 0.633 (0.022) & 3.17 (0.29) & 2.6 & 0.9 & 4.1 \\\\ \\hline \\end{tabular} \\end{center} $^a$~Hipparcos measurements, $^b$~from our new fit to the optical, near-IR and IRAC mid-IR photometry (see section~4.2), $^c$~estimated using evolutionary models appropriate for He-atmosphere white dwarfs with C/O core compositions, see http://www.astro.umontreal.ca/~bergeron/CollingModels/, $^d$~derived from the initial-final mass relation of \\citet{Dobbie06}, $^e$~\\citet{Wood} \\end{table*} ",
        "conclusions": "We have placed limits on ultra-cool substellar and planetary-mass objects around the nearest single white dwarf, vMa\\,2, through a search for common proper motion companions and mid-infrared photometric excesses. The red giant progenitor to vMa\\,2 had a maximum radius of $\\sim1000R_\\odot$ ($4.6$\\,AU, \\citealt*{Hurley}). Therefore, taking into account the post main sequence evolution of the orbits of any companions, our $J$-band images have covered all orbits of substellar objects that originally lay beyond the maximum extent of the red giant envelope, up to 50\\,AU distance. The complimentary {\\it Spitzer} IRAC photometry places slightly higher limits on the masses and temperatures of any spatially unresolved substellar companion whose orbit fortuitously lies along the line of sight to vMa\\,2, or that currently has an orbit within $\\approx13$\\,AU of the white dwarf. These limits are significantly lower than those reported by previous studies, e.g. \\citet{fbm} and \\citet{fbz05}.  A variety of searches for planetary companions to white dwarfs are currently underway (e.g.~\\citealt{dodob}; \\citealt{Debes1}; \\citealt{MullallyB}). Further, extensive results from the {\\it DODO} survey are in preparation \\citep{Hogan}."
    },
    "0801/0801.0409_arXiv.txt": {
        "abstract": "{ Simbol-X is a French-Italian mission, with a participation of German laboratories, for X-ray astronomy in the wide 0.5-80 keV band. Taking advantage of emerging technology in mirror manufacturing and spacecraft formation flying, Simbol-X will push grazing incidence imaging up to $\\sim80$ keV, providing an improvement of roughly three orders of magnitude in sensitivity and angular resolution compared to all instruments that have operated so far above 10 keV. This will open a new window in X-ray astronomy, allowing breakthrough studies on black hole physics and census and particle acceleration mechanisms. We describe briefly the main scientific goals of the Simbol-X mission, giving a few examples aimed at highlighting key issues of the Simbol-X design. ",
        "introduction": "A seminal result obtained with HEAO-1 at the end of the 70' is the precise measure of the spectrum (from a few keV up to $\\sim100$ keV) of the isotropic, extragalactic Cosmic X-Ray Background (CXB), discovered by Riccardo Giacconi, Bruno Rossi and collaborators during one of the first rocket-borne X--ray experiments in 1962. The HEAO1 data showed that the CXB energy density has a broad maximum around 30 keV, where it is about 5 times higher than at 1 keV and 50\\% higher than at 10 keV.  It was soon realized that the CXB is most likely due to the contribution of many discrete sources at cosmological distances \\citep{sw79}. Most of these sources are active galactic nuclei, AGN, implying that the CXB energy density provides an integral estimate of the mass accretion rate in the Universe, and therefore of the super-massive black hole (SMBH) growth and mass density. Unfortunately, the integrated light from all sources detected in HEAO1 all-sky survey could directly explain only less than 1\\% of the CXB.  Indeed, the use of collimated detectors on board first UHURU and Ariel-V, and then HEAO1 in the 1970 decade led to the discovery of $<1000$ X-ray sources in the whole sky.  X-ray imaging observations, performed first by {\\it Einstein} and ROSAT in the soft X-ray band below $\\sim3$ keV, and then by ASCA, BeppoSAX, XMM--Newton and Chandra up to 8-10 keV, detected tens of thousands of X-ray sources, and resolved nearly 100\\% of the CXB below a few keV and up to 50\\% at 6-8 keV. These observations increased by orders of magnitude the discovery space for compact objects (both Galactic neutron stars and black holes and AGN) and for thermal plasma sources. However, they still leave open fundamental issues, such as what is making most of the energy density of the CXB at $\\sim30$ keV.  Above 10 keV the most sensitive observations have been performed so far by collimated instruments, like the BeppoSAX PDS, and by coded masks instruments, like INTEGRAL IBIS and Swift BAT. Only a few hundred sources are know in the whole sky in the 10-100 keV band, a situation recalling the pre--{\\it Einstein} era at software energies. A new window in X-ray astronomy above 10 keV must be opened, producing an increase of the discovery space similar to that obtained with the first X-ray imaging missions. This will be achieved by Simbol-X, a formation flight mission currently under preparation by France and Italy, with a participation from German laboratories. Very much like the {\\it Einstein} Observatory, this mission will have the capabilities to investigate almost any type of X-ray source, from Galactic and extragalactic compact sources, supernova remnant (SNR), young stellar objects and clusters of galaxies, right in the domain where accretion processes and acceleration mechanisms have their main signatures.  This paper summarizes the main scientific goals of the Simbol-X mission, putting the emphasis on the core science objectives. ",
        "conclusions": "Thanks to the emerging technology in X-ray mirrors (e.g. multilayer coating, see Pareschi et al. these proceedings) and spacecraft formation flying, Simbol-X will provide a large collecting area (of the order of 100-1000 cm$^{-2}$) from a fraction of keV up to $\\sim80$ keV, thus overcoming the ``10 keV limit'' for high accuracy imaging and spectroscopy of all past and current X-ray observatories.  This, together to the good image quality (PSF HPD$<20$ arcsec, FWHM$<10$ arcsec), relatively large field of view (12 arcmin diameter), good detector quantum efficiency, resolution and low internal background, will allow breakthrough studies on black hole physics and census, and particle acceleration mechanisms."
    },
    "0801/0801.1914_arXiv.txt": {
        "abstract": "We give a  closer look  at  the   Central Limit Theorem  (CLT) behavior  in  quasi-stationary states  of  the  Hamiltonian Mean Field model, a paradigmatic one  for long-range-interacting classical many-body systems. We present  new calculations  which  show  that, following their time evolution,  we can observe and classify  three kinds  of long-standing quasi-stationary states (QSS) with different  correlations. The  frequency of occurrence of each  class  depends on the size of the system.  The  different   microsocopic nature of the QSS    leads  to different dynamical correlations and therefore to different    results  for the observed CLT behavior. ",
        "introduction": "Very recently  there has been  a lot of interest in   generalizations  of  the  Central Limit Theorem (CLT) \\cite{clt0,umarov,barkai} and on  their   possible (strict or numerically approximate) application   to   systems  with long range correlations \\cite{luis,hil}, at the edge-of-chaos \\cite{ucc}, nonlinear dynamical systems  the maximal   Lyapunov exponent  of which is either  exactly zero or  tends to vanish in the thermodynamic limit (increasingly large systems) \\cite{anteneodo}, hindering  this way  mixing  and thus the application  of standard statistical mechanics.  A possible application of  nonextensive statistical mechanics  \\cite{cost1,clt1} has  been  advocated  in these cases. Along this  line   we discuss  in the present  paper  a   detailed  study of   a  paradigmatic \\textit{toy model}   for  long-range  interacting  Hamiltonian   systems \\cite{hmf,liap,epn-rap,chavanis,fulvio1,ruffo,reply07}, i.e.  the     Hamiltonian  Mean Field (HMF) model  which has been intensively   studied  in the last years. In a recent article \\cite{rapis-ctnext07}, we presented  molecular dynamics numerical results for the HMF model  showing  three kinds of quasi-stationary states (QSS) starting from the same water-bag initial condition with unitary magnetization ($M_0=1$). In the following  we present   how the applicability of the   standard or $q$-generalized CLT is influenced by  the different microscopic dynamics observed in the  three classes of QSS.  In general, averaging   over the threee  classes  can be misleading. Indeed, the frequency of appearence of each of these classes depends on the size of the system under investigation, and there is no clear evidence that  a predominant class exists. ",
        "conclusions": "On the  basis  of the new calculations presented here, one should distinguish among different classes  of QSS for a given size  of  the HMF system. Our tests do confirm that correlations can be different for different dynamical realizations of the same system, starting from the same class of initial conditions, and therefore also the  Central Limit  Theorem behavior  can change.   According to the  class  considered  we can have a Gaussian Pdf, a $q$-Gaussian one  or a mixture  between  the two. In this respect, the presence of  a $q$-Gaussian curve  confirms the results  previously published about the indications  for a generalized  CLT in long-range  Hamiltonian systems  when ergodicity is broken  and correlations  are strong enough. The present results, where we verify that the time evolution of the system is quite sensitive to the specific initial conditions, reminds what occurs in other anomalous systems, like the Hydrogen atom as studied  recently in Refs. \\cite{Oliveira1,Oliveira2}. We  have  checked   also that inequivalence between time averages and ensemble  averages continues to hold within the different  classes  for $M_0=1$ and  for $M_0=0$ initial conditions. Finally, although  it is true that   the frequency of occurrence of events of class 1 and 2  for very large  sizes, and in particular  the  attractor  for the events of class  2, remain to be investigated  with more accuracy  (calculations  with  better   statistics are in progress)  before advancing definitive conclusions,  from these  simulations  one could  conjecture  that   the events  of class 1 tend to  disappear  in the thermodynamic limit. However, even if this  was the case, their probability of occurrence  (as well as the probability of occurrence of the first metastable plateau of events of class 2) can be significantly different from zero for  large  but finite systems and therefore  the  applicability of $q$-statistics in metastable states of real complex systems remains a valid and interesting possibility."
    },
    "0801/0801.3857_arXiv.txt": {
        "abstract": "We present evidence of Fe fluorescent emission in the \\cha\\ \\hetgs\\ spectrum of the single G-type giant \\hr\\ during a large flare. In analogy to solar X-ray observations, we interpret the observed Fe~K$\\alpha$ line as being produced by illumination of the photosphere by ionizing coronal X-rays, in which case, for a given Fe photospheric abundance, its intensity depends on the height of the X-ray source. The \\hetgs\\ observations, together with 3D Monte Carlo calculations to model the fluorescence emission, are used to obtain a direct geometric constraint on the scale height of the flaring coronal plasma. We compute the Fe fluorescent emission induced by the emission of a single flaring coronal loop which well reproduces the observed X-ray temporal and spectral properties according to a detailed hydrodynamic modeling.  The predicted Fe fluorescent emission is in good agreement with the observed value within observational uncertainties, pointing to a scale height $\\lesssim 0.3$\\rstar. Comparison of the \\hr\\ flare with that recently observed on II~Peg by {\\em Swift} indicates the latter is consistent with excitation by X-ray photoionization. ",
        "introduction": "\\label{s:intro}   Since the early 1970's, spatially resolved observations of the solar corona have revealed a high degree of structuring over all scales, from of order of the solar radius down to instrumental resolution limits (e.g., \\citealt{Vaiana73,Vaiana73b}).  This spatial structuring is intimately linked to the characteristics of magnetic field generation and interaction, and to plasma heating mechanisms.  On other stars, especially very active ones (with X-ray luminosity $L_{\\rm X}$ up to $10^4 \\times L_{\\rm X \\odot}$), coronal structure and its relation with stellar parameters such as mass, rotation and evolutionary phase, remains very uncertain.  Techniques used to date to investigate the morphology of stellar coronae comprise rotational modulation \\citep[e.g.,][]{MarinoL03,Flaccomio05}, flare modeling (e.g., \\citealt{Reale04}; \\citealt{Testa07a}, hereafter \\ta), eclipse mapping \\citep[e.g.,][]{White90}, spectroscopic density and radiation field diagnostics \\citep[e.g.,][]{Testa04b,Ness04}, resonance scattering \\citep[e.g.,][]{Testa04a,Matranga05,Testa07b}, velocity modulation \\citep[e.g.,][]{Brickhouse01,Chung04,Hussain05,Huenemoerder06}, and simultaneous Doppler imaging and X-ray spectroscopy \\citep{Hussain07}.  The fluorescent iron line at $\\sim$6.4~keV \\footnote{The Fe fluorescent feature actually consists of two components, at 6.391 and 6.404~keV (for Fe\\,{\\sc i}), which are unresolved by present instruments, therefore hereafter we will refer to the feature as to a single emission line.}  presents a further, potentially powerful diagnostic of stellar coronal geometry.  Extensively used in study of active galactic nuclei and X-ray binaries (see e.g.\\ reviews \\citealt{George91,Reynolds03}), and often seen on the Sun during flares \\citep{Culhane81,Parmar84,Tanaka84,Zarro92,Phillips94}, the line is produced by electron cascade after one of the two K-shell ($n$=1) electrons of an iron atom (or ion) is ejected following photoelectric absorption of X-rays.  In the solar and stellar case, for a given irradiating flare spectrum the line strength depends essentially on three parameters: the flare height, heliocentric angle, and photospheric Fe abundance (e.g., \\citealt{Bai79}; \\citealt{Drake07b}, hereafter D07b). If it can be observed in stars, Fe~K fluorescence therefore presents a means of estimating flare and coronal scale height.  Fe fluorescence has now been detected on a number of pre-main sequence stars with disks \\citep[e.g.,][]{Tsujimoto05,Favata05}, where it is most likely produced by coronal X-ray irradiation of the cold disk material.  Here we present evidence for {\\em photospheric} Fe fluorescent emission in the \\cha\\ \\hetgs\\ spectrum of the X-ray active single G1 giant \\hr\\ (see \\ta\\ and references therein for a description of the characteristics of the target).  The observations and analysis are briefly described in \\S\\ref{s:obs}. In \\S\\ref{s:model} we present 3D Monte Carlo calculations of the Fe K$\\alpha$ fluorescence.  The results are presented in \\S\\ref{s:results}, where we derive an estimate for the coronal scale height and compare it with the prediction of the loop hydrodynamic model that succeeds in describing the observed flaring spectrum (\\ta). To our knowledge, the \\hr\\ {\\em Chandra} spectrum represents only the second detection of Fe photospheric fluorescence.  Very recently, the line was observed by the {\\em Swift} X-ray Telescope during a ``superflare'' on the RS~CVn binary system II~Peg by \\citet{Osten07}, who ascribed the excitation mechanism to electron impact ionization of photospheric Fe instead of photoionization.  We discuss our result and the II~Peg observations in \\S\\ref{s:discuss}, and show that the latter is also consistent with, and more plausibly associated with, photoionization than electron impact. ",
        "conclusions": "\\label{s:discuss}  We have presented the first detailed analysis of observed photospheric fluorescence to probe the geometric properties of stellar coronal emission.  For the flare on \\hr\\ we find the observed Fe~K$\\alpha$ intensity completely consistent with production through photospheric irradiation by flare X-rays.  Our estimate for the flare scale height of $h \\lesssim 0.3$\\rstar\\ provides a cross-check for the results of the hydrodynamic modeling of the flare observed on \\hr\\ that can be satisfactorily reproduced by a single flaring loop of semi-length $L =0.5$\\rstar.  The agreement between the observed $\\mathcal{E}$ and the value computed directly from this model provides further confidence both in the hydrodynamic modeling approach and X-ray fluorescence as the excitation mechanism for the Fe~K$\\alpha$ line. We note that our model based on \\cha\\ data might underestimate the high energy continuum flux because we cannot detect any additional (either thermal or non-thermal) emission possibly present at energies outside the \\cha\\ band. However, if additional high energy emission were present, the higher flux above 7.11~keV would make the efficiency lower bringing it in even better agreement with the predictions of the model.  As noted in \\S\\ref{s:intro}, Fe K$\\alpha$ emission has been extensively observed in solar X-ray spectra during flares \\citep{Culhane81,Parmar84,Tanaka84,Zarro92,Phillips94}.  Two main production mechanisms were initially explored to explain the line: electron impact from a non-thermal electron population, and X-ray photoionization.  The observed Fe K$\\alpha$ evolution with respect to the thermal and non-thermal emission, and the observed center-to-limb variations seen on the Sun, strongly support the X-ray fluorescence mechanism \\citep{Culhane81,Parmar84,Tanaka84}, although in rare cases there is indication that electron impact might contribute during the early impulsive phase \\citep[e.g.,][]{Emslie86}. \\citet{Ballantyne03} have also shown that Fe~K production in accretion disks by non-thermal electron bombardment is extremely unlikely, and requires 2-4 orders of magnitude greater energy dissipation in the electron beam than is required for an X-ray photoionization source. In this context, the recent observations of an extremely large Fe~K$\\alpha$ equivalent width in the PMS star V1486~Ori \\citep{Czesla07}, and variability in Fe~K$\\alpha$ emission uncorrelated with the observed X-ray continuum on Elias 29 \\citep{Giardino07}, seem to challenge the photoionization excitation mechanism.  However, D07b note that fluorescence from PMS disks excited by X-rays originating from the unseen stellar hemisphere will also be observed: in $\\leq 50$\\%\\ of cases, Fe~K$\\alpha$ can then appear anomalous when compared only to the {\\em observed} thermal continuum.  To date, the {\\em Swift} observation of II~Peg \\citep{Osten07} is the only other case where Fe~K$\\alpha$ emission has been observed for a star (other than the Sun) lacking substantial circumstellar material. \\cite{Osten07} dismissed X-ray fluorescence as the excitation mechanism based on arguments for fluorescence in an optically-thin medium, and instead attributed the Fe K$\\alpha$ to electron impact with non-thermal electrons.  As noted in earlier discussions \\citep[e.g.][]{Parmar84,Ballantyne03}, one major problem for an electron impact excitation mechanism is the inefficiency of this process.  Extremely large energies must be dissipated in accelerated electron beams to produce continuum and fluorescence emission that might be observed on a star.  Indeed, \\citet{Osten07} showed that the hard X-ray continuum observed in the II~Peg flare would require an energy of $3 \\times 10^{40}$~ergs in non-thermal electrons if interpreted in terms of thick-target bremsstrahlung.  Over the time of the duration of the flare, this corresponds to an energy dissipation rate in accelerated electrons of more than 100 times the stellar bolometric luminosity.  The hard X-rays are also interpreted as lasting for timescales comparable with the flare soft X-ray emission, at variance with the Sun where this non-thermal component, when present, is generally impulsive. Regardless of the true nature of the hard X-ray flux, the equivalent width measured from the II~Peg {\\em Swift} spectra ranges between 18 and 60~eV, corresponding to efficiency values $\\mathcal{E}$ between 1\\% and 2\\%.  These values are well within the range found from the theoretical calculations presented here, and comparable with the measured value for \\hr.  Since the non-thermal hard X-ray component in the II~Peg flare is not well-constrained by the {\\em Swift} data and could also be explained by thermal emission, the observed Fe~K$\\alpha$ line seems more easily explained as arising from photoionization by flare X-rays (see also the discussion of D07b)."
    },
    "0801/0801.3585_arXiv.txt": {
        "abstract": "We present a new non-parametric deprojection algorithm DOPING (Deprojection of Observed Photometry using and INverse Gambit), that is designed to extract the three dimensional luminosity density distribution $\\rho$, from the observed surface brightness profile of an astrophysical system such as a galaxy or a galaxy cluster, in a generalised geometry, while taking into account changes in the intrinsic shape of the system. The observable is the 2-D surface brightness distribution of the system. While the deprojection schemes presented hitherto have always worked within the limits of an assumed intrinsic geometry, in DOPING, geometry and inclination can be provided as inputs. The $\\rho$ that is most likely to project to the observed brightness data is sought; the maximisation of the likelihood is performed with the Metropolis algorithm. Unless the likelihood function is maximised, $\\rho$ is tweaked in shape and amplitude, while maintaining positivity, but otherwise the luminosity distribution is allowed to be completely free-form. Tests and applications of the algorithm are discussed. ",
        "introduction": "\\label{sec:intro} \\noindent The preliminary step involved in the dynamical modelling of galaxies, concerns the deprojection of the observed surface brightness distribution into the intrinsic luminosity density, as has been practised by [\\refcite{krajnovic04}], [\\refcite{kronawitter00}], [\\refcite{magog98}], among others. Deprojection though, is a non-unique problem, unless performed under very specific configurations of geometry and inclination, as discussed by [\\refcite{binneygerhard}], [\\refcite{kochanekrybicki}], [\\refcite{rybicki87}], [\\refcite{vandenbosch}] and others.  Over the years, several deprojection schemes have been advanced and implemented within the purview of astronomy; these include parametric formalisms designed by [\\refcite{bendinelli}], [\\refcite{palmer}] and [\\refcite{cappellari}], as well as non-parametric methods, such as the Richardson-Lucy Inversion scheme developed by [\\refcite{richardson}] and [\\refcite{lucy}] and a method suggested by [\\refcite{romkoch}]. While the parametric schemes are essentially unsatisfactory owing to the dependence of the answer on the form of the parametrisation involved, the non-parametric schemes advanced till now have suffered from the lack of transparency and in the case of the Richardson-Lucy scheme, lack of an objective convergence criterion.  Here, we present a new, robust non-parametric algorithm: Deprojection of Observed Photometry using an INverse Gambit (DOPING). DOPING does not need to assume axisymmetry but can work in a triaxial geometry with assumed axial ratios, and is able to incorporate radial variations in eccentricity. Although the code can account for changes in position angle, this facet has not been included in the version of the algorithm discussed here. Also, here we present the 1-D results obtained with DOPING but the code provides the full 3-D density distribution that projects to observed brightness map. In a future contribution, (Chakrabarty $\\&$ Ferrarese, {\\it in preparation}), DOPING will be applied to recover the intrinsic luminosity density of about 100 early type galaxies observed as part of the ACS Virgo Cluster Survey, as reported in [\\refcite{coteacs04}].  The paper has been arranged as follows. The basic framework of DOPING is introduced in Section~\\ref{sec:method}. This is followed by a short discourse on a test of the algorithm. An application to the observed data of the galaxy vcc1422 is touched upon in Section~\\ref{sec:1422}. Another application of DOPING is discussed in Section~\\ref{sec:clusters}. The paper is rounded up with a summary of the results. ",
        "conclusions": "\\noindent In this paper, we have introduced a new non-parametric algorithm DOPING that is capable of inverting observed surface brightness distributions of galaxies and galaxy clusters, while taking into account variations in the intrinsic shapes of these systems. The potency of DOPING is discussed in the context of a test galaxy in which the eccentricity is made to change radically with radius. The code is also successfully applied to obtain the luminosity density distribution of the faint nucleated galaxy vcc1422. Lastly, a novel use is made of the capability of DOPING to deproject in general geometries, in determining the intrinsic shape and inclination of a galaxy cluster. It is envisaged that implementing a measure of the LOS extent of a cluster from Sunyaev Zeldovich measurements, will help tighten the estimates of cluster inclination and the intrinsic axial ratios of triaxial clusters."
    },
    "0801/0801.1255_arXiv.txt": {
        "abstract": "The nature of the extended hard X-ray source XMMU~J061804.3+222732 and its surroundings is investigated using {\\sl XMM-Newton}, {\\sl Chandra}, and {\\sl Spitzer} observations. This source is located in an interaction region of the IC~443 supernova remnant with a neighboring molecular cloud. The X-ray emission consists of a number of bright clumps embedded in an extended structured non-thermal X-ray nebula larger than 30\\arcsec in size. Some clumps show evidence for line emission at $\\sim$ 1.9 keV and $\\sim$ 3.7 keV at the 99\\% confidence level. Large-scale diffuse radio emission of IC~443 passes over the source region, with an enhancement near the source. An IR source of about 14\\arcsec $\\times$ 7\\arcsec size is prominent in the 24 $\\mu$m, 70 $\\mu$m, and 2.2 $\\mu$m bands, adjacent to a putative Si K-shell X-ray line emission region. The observed IR/X-ray morphology and spectra are consistent with those expected for J/C-type shocks of different velocities driven by fragmented supernova ejecta colliding with the dense medium of a molecular cloud. The IR emission of the source detected by {\\sl Spitzer} can be attributed to both continuum emission from an HII region created by the ejecta fragment and line emission excited by shocks. This source region in IC~443 may be an example of a rather numerous population of hard X-ray/IR sources created by supernova explosions in the dense environment of star-forming regions. Alternative Galactic and extragalactic interpretations of the observed source are also discussed. ",
        "introduction": "The energy release and the ejection of nucleosynthesis products by supernovae (SNe) events are of great importance for our understanding of the physics of the interstellar medium (ISM). The mixing of the ejected metals with the surrounding matter is of special interest when a SN occurs in a molecular cloud, which may cause further star-forming activity.  Optical and UV studies of the structure of SN remnants (SNRs) have revealed a complex metal composition of ejecta and the presence of isolated high-velocity ejecta fragments interacting with surrounding media. The most prominent manifestations of this phenomena are the fast moving knots observed in some young ``oxygen-rich'' SNRs, such as the Galactic SNRs Cas A  (e.g., Chevalier \\& Kirshner 1979; Fesen\\etal 2006), Puppis A (Winkler \\& Kirshner 1985), G292.0+1.8 (e.g. Winkler \\& Long 2006), and also N132D in the LMC and 1E 0102.2--7219 in the SMC (e.g. Blair et al. 2000).  Ballistically moving ejecta fragments of SNRs can be considered as a class of hard X-ray sources. The prototype was observed in the Vela SNR (Aschenbach, Egger, \\& Tr\\\"umper 1995; Miyata et al.\\ 2001). A massive individual fragment moving supersonically through a molecular cloud can have a luminosity $L_x \\gsim$10$^{31} \\ergs$ in the 1--10 keV band, and is observable with \\xmm\\ and \\chan\\ at a few kpc distance (Bykov 2002, 2003). Its X-ray emission is expected to consist of two components. The first one is thermal X-ray emission from the hot shocked ambient gas behind the fragment bow shock, with a spectrum of an optically thin thermal plasma of an ISM-cloud abundance. The second emission component is nonthermal; the interaction of fast electrons accelerated at the fragment bow-shock with the fragment body produces a hard continuum as well as line emission (X-ray and IR), including the K-shell lines of Si, S, Ar, Ca, Fe, and other elements ejected by SN. Detection of the X-ray line emission would help distinguish an ejecta fragment from the other possible source of hard continuum emission associated with a SNR, namely, a pulsar wind nebula (PWN).  A young SNR of an age of a few thousands years interacting with a molecular cloud can produce hundreds of X-ray sources associated with isolated ejecta fragments. They should be particularly numerous in starforming regions like those in the Galactic center region, where young core-collapsed supernovae in or near molecular clouds are expected to be present in abundance.  The expected observational appearance of isolated ejecta fragments in a molecular cloud differs from what is seen in the Vela SNR. Ejecta fragments interacting with a dense molecular cloud are slowed down and crushed, and they are generally more bright. We will argue here that the X-ray emission spectra of ejecta fragments in a molecular cloud may be dominated by hard non-thermal components, because a powerful but very soft thermal component could be heavily absorbed. On the other hand, the spectra of fast supernova ejecta fragments propagating in a tenuous plasma, as it is the case in the Vela SNR, would be long-lived, less luminous and dominated by thermal emission.   The present paper focuses on IC~443. This is a SNR of a medium age, estimated by Chevalier (1999) to be  $\\sim$ 30,000 years, for which the number of X-ray sources from ejecta fragments should be {\\it much smaller\\ } than in a young SNR (possibly, only a few). It is, however, the best and most reliable laboratory to study this phenomenon since there are only very few examples of clearly established SNR-cloud interactions.  IC~443 (G189.1+3.0) is an evolved SNR of about 45$\\arcmin$ size at a distance of 1.5 kpc (e.g. Fesen \\& Kirshner 1980). Radio observations of IC~443 (e.g. Braun \\& Strom 1986; Green 1986; Leahy 2004) show two half-shells.  This appearance is probably due to interaction of the SNR with a molecular cloud that seems to separate the two half-shells.  The molecular-cloud material has a torus-like structure (Cornett, Chin \\& Knapp 1977; Burton\\etal 1988; Troja, Bocchino, \\& Reale 2006), that can be interpreted as a sheet-like cloud first broken by the expanding pre-supernova wind and then by the SNR blast wave.  Plenty of evidence for shock-excited molecules in this region has been found (e.g. DeNoyer 1979; Burton et al.\\ 1988; Dickman\\etal 1992; Turner\\etal 1992; van Dishoeck, Jansen, \\& Phillips 1993; Tauber\\etal 1994; Richter, Graham, \\& Wright 1995; Cesarsky\\etal 1999; Snell\\etal 2005). The complex structure of the interaction region, with evidence for multiple dense clumps, is seen in 2MASS images (e.g. Rho et al.\\ 2001). Three OH (1720 MHz) masers were found in IC~443 (Claussen \\etal 1997; Hewitt \\etal 2006, and references therein).  Soft X-ray maps of IC~443 based on {\\sl ROSAT} data (Asaoka \\& Aschenbach 1994) and recent radio observations (Leahy 2004) suggest that another SNR, G189.6+3.3, is seen in the IC~443 field (see also the \\xmm\\ study by Troja, Bocchino, \\& Reale 2006). This makes the multiwavelength observational picture even more complex to interpret.  The field of IC~443 was observed in X-rays with {\\sl HEAO 1\\/} (Petre et al.\\ 1988), {\\sl Ginga\\/} (Wang et al.\\ 1992), {\\sl ROSAT\\/} (Asaoka \\& Aschenbach 1994), {\\sl ASCA\\/} (Keohane et al.\\ 1997; Kawasaki et al.\\ 2002), {\\sl BeppoSAX\\/} (Preite-Martinez et al.\\ 2000; Bocchino \\& Bykov 2000), \\chan\\ (Olbert \\etal 2001; Bykov, Bocchino, \\& Pavlov 2005; Gaensler \\etal 2006; Weisskopf \\etal 2007), \\xmm\\ (Bocchino \\& Bykov 2001, 2003; Troja, Bocchino, \\& Reale 2006), and {\\sl RXTE} (Sturner, Keohane, \\& Reimer 2004).  The X-ray emission of IC~443 below 4 keV is dominated by a number of thermal components (e.g., Petre \\etal 1988; Asaoka \\& Aschenbach 1994; Kawasaki et al.\\ 2002; Troja, Bocchino, \\& Reale 2006). The thermal-emission morphology is center-filled, with soft emission filaments visible at energies below 0.5 keV. A gradient of X-ray surface brightness at the SNR limb was found, as well as strong variations of absorbing column density \\nh, which indicates the complex molecular-cloud environment of IC~443 in the southern part of the remnant (e.g. Asaoka \\& Aschenbach 1994).  {\\sl ASCA} observations have established that the hard X-ray emission of IC~443 (above 4 keV) is dominated by localized sources in the southern part of the remnant (Keohane et al.\\ 1997). In \\xmm\\ observations Bocchino \\& Bykov (2003; BB03 hereafter) found 12 sources with fluxes over 10$^{-14}~\\enf$ in the 2--10 keV band. Six of the detected sources are located in a relatively small, of $15^\\prime\\times 15^\\prime$ size, region projected onto the molecular cloud in the South-Eastern part of IC~443. {\\sl BeppoSAX\\/} MECS observations (4--10 keV) showed two sources, 1SAX~J0617.1+2221 and 1SAX~J0618.0+2227, with evidence from the {\\sl BeppoSAX\\/} PDS for the presence of hard emission up to 100 keV for the former (Bocchino \\& Bykov 2000). Observations of this source by \\chan\\ (Olbert\\etal 2001; Gaensler\\etal 2006; Weisskopf\\etal 2007) and \\xmm\\ (Bocchino \\& Bykov 2001) established its plerionic nature. Leahy (2004) argued that the pulsar that powers this plerion is associated with G189.6+3.3 rather than IC~443. The nature of the second hard source -- 1SAX~J0618.0+2227 -- remained unknown. This source, the brightest in the region (excluding the plerion), was resolved with \\xmm\\ into two sources -- the extended XMMU~J061804.3+222732 (of $\\sim$ 20\\arcsec size) and the point-like XMMU~J061806.4+222832. We will call them Src~1 and Src~2 respectively (note that the sources were listed as Src~11 and Src~12 in BB03). The position of XMMU~J061804.3+222732 in the remnant is illustrated in Figures~\\ref{xmm_spitzer} and \\ref{chan_xmm}. \\begin{figure*} \\includegraphics[width=0.99\\textwidth]{f1.eps} \\caption{ Wide-field views of SNR IC~443. {\\em Left:} {\\sl XMM-Newton} 2--8 keV image with {\\sl Spitzer} MIPS 24 $\\mu$m contours overlaid. {\\em Right:} {\\sl Spitzer} MIPS image at 24 $\\mu$m with {\\sl VLA} 1.4 GHz contours overlaid. The images are produced from the {\\sl XMM-Newton} observations 0114100101--0114100601 and 0301960101, and {\\sl Spitzer} MIPS observations r4616960, r4617216, and r4617472. The white arrow points to the studied region.} \\label{xmm_spitzer} \\end{figure*} A dedicated \\chan\\ observation of Src~1 has revealed a complex structure of a few bright clumps embedded in extended emission of $>$ 20\\arcsec size (Bykov, Bocchino, \\& Pavlov 2005; BBP05 hereafter). The brightest clumps are the extended Src~1a and the point-like Src~1b. The apparent position of the source in a SNR -- molecular cloud interaction region naturally leads to SNR-related interpretations. The observed X-ray morphology of Src\\,1 and the spectra of its components are consistent with expectations for a SN ejecta fragment interacting with a dense ambient medium. Alternatively, Src\\,1 could be interpreted as a PWN associated with either IC~443 or G189.6+3.3 (BBP05). However, one cannot exclude the extragalactic origin of the source, that is discussed in some detail in Section~\\ref{altern}.  IC~443 is a candidate counterpart of the EGRET  $\\gamma$-ray source 3EG~J0617+2238, with a flux of about 5$\\times$10$^{-7}$ cm$^{-2}$ s$^{-1}$ above 100 MeV (Esposito et al.\\ 1996). The spectrum of Src~1 extrapolated into the EGRET range is consistent with that of 3EG~J0617+2238 (BBP05). Also the position of Src~1 is consistent (albeit marginally) with that of 3EG~J0617+2238. Such a $\\gamma$-ray luminosity can be expected for both the fragment and PWN interpretations. The forthcoming {\\sl GLAST\\ } mission (e.g. Johnson 2006) will be able to provide an accurate position and spectrum of 3EG~J0617+2238, thus helping to solve the issue. Src~1 lies far away from the 99\\% error circle of the TeV-regime source recently reported by {\\sl MAGIC} (Albert et al. 2007) in the Western part of IC~443 field. The apparent position of TeV {\\sl MAGIC} source is close to the 1720 MHz OH maser detected by Claussen et al. (1997).  We present here new results of a deep 80 ks observation of the region with \\xmm , imaging of the region with the {\\sl Spitzer} infrared observatory, and a new analysis of VLA radio observations. In \\S\\,2 a combined analysis of the new \\xmm\\ observations and all the previous high-resolution X-ray data from \\chan\\ and \\xmm\\ is presented, including images, spectra, and time variations in the X-ray domain. In \\S\\,3 archival radio (VLA), IR (2MASS and {\\sl Spitzer} MIPS), and optical (POSS-II) data are used to constrain the nature of Src~1. A discussion of the obtained results and future prospects are presented in \\S\\,4. ",
        "conclusions": "\\label{concl}  The multi-wavelength observations presented here indicate a possible physical connection of the X-ray source J0618 with the neighboring IR {\\sl Spitzer} and 2MASS sources. That connection, if real, can be understood in a scenario where J0618 originates in an interaction of the IC~443 SNR with the adjacent molecular cloud. The correlation would require the presence of both fast and slow shocks in the clumpy molecular-cloud material. The X-ray line features apparent in the spectra of the clumps favor a scenario in which the shocks are produced by a fast ballistically moving SN ejecta fragment penetrating into a structured molecular cloud. The model provides a physical picture coherent with the current observational data, although alternative scenario cannot be rejected yet.  Alternatively, Src 1 can be interpreted as a massive X-ray cluster of galaxies at a redshift $z >$ 0.5. High-resolution arcsecond-scale observations and fine spectroscopy are required to distinguish between these very different scenarios.  If the SNR-ejecta interpretation is confirmed by further observations, the source J0618 in the IC~443 could be a prototype of a rather numerous population of hard X-ray -- IR sources created by SN explosions in the dense environments of star-forming regions. Such sources would be particularly abundant in the Galactic Centre region."
    },
    "0801/0801.3250_arXiv.txt": {
        "abstract": "The Alpha Magnetic Spectrometer (AMS), whose final version AMS-02 is to be installed on the International Space Station (ISS) for at least 3 years, is a detector designed to measure charged cosmic ray spectra with energies up to the TeV region and with high energy photon detection capability up to a few hundred GeV. It is equipped with several subsystems, one of which is a proximity focusing RICH detector with a dual radiator (aerogel+NaF) that provides reliable measurements for particle velocity and charge. The assembly and testing of the AMS RICH % is currently being finished % and the full AMS detector is expected to be ready by the end of 2008. The RICH detector of AMS-02 is presented. Physics prospects are briefly discussed. ",
        "introduction": "The Alpha Magnetic Spectrometer (AMS)\\cite{bib:ams}, whose final version AMS-02 is to be installed on the International Space Station (ISS) for at least 3 years, is a detector designed to study the cosmic ray flux by direct detection of particles above the Earth's atmosphere, at an altitude of \\mbox{$\\sim$ 400 km}, using state-of-the-art particle identification techniques. AMS-02 is equipped with a superconducting magnet cooled by superfluid helium. The spectrometer is composed of several subdetectors: a Transition Radiation Detector (TRD), a Time-of-Flight (ToF) detector, a Silicon Tracker, Anticoincidence Counters (ACC), a Ring Imaging \\CK\\ (RICH) detector and an Electromagnetic Calorimeter (ECAL). A preliminary version of the detector, AMS-01, was successfully flown aboard the US space shuttle Discovery in June 1998.   \\begin{figure}[htb]%  \\center  \\vspace{-0.5cm}  \\begin{tabular}{cc}  \\mbox{\\epsfig{file=amsrich.eps,width=0.47\\textwidth,clip=}} & \\mbox{\\epsfig{file=RICHmaravilhoso_new.ps,width=0.43\\textwidth,clip=, bbllx=-53,bblly=100,bburx=667,bbury=666}}  \\end{tabular}  \\vspace{-0.2cm}  \\caption{The RICH detector of AMS-02 \\emph{(left)}. View of the assembled RICH detector at CIEMAT \\emph{(right)}.\\label{richdet}}  \\vspace{-0.2cm}  \\end{figure}   The main goals of the AMS-02 experiment are: (i) a precise measurement of charged cosmic-ray spectra in the rigidity region between \\mbox{$\\sim$ 0.5 GV} and \\mbox{$\\sim$ 2 TV} and the detection of photons with energies up to a few hundred GeV; (ii) a search for heavy antinuclei \\mbox{($Z \\ge$ 2)}, which if discovered would signal the existence of cosmological antimatter; (iii) a search for dark matter constituents by examining possible signatures of their presence in cosmic ray spectra. The long exposure time and large acceptance \\mbox{(0.5 m${}^2\\cdot$sr)} of AMS-02 will enable it to collect an unprecedented statistics of more than $10^{10}$ nuclei.   \\vspace{-0.3cm} ",
        "conclusions": "AMS-02 will provide a new insight on the cosmic-ray spectrum by collecting precise data for an unprecedented number of particles above the Earth's atmosphere. The RICH detector will play a key role in the operation of AMS due to its capabilities for velocity reconstruction, charge determination and albedo rejection. Extensive testing has been performed on the RICH detector and its components. The assembly of the RICH detector is currently being finished and the full AMS detector is expected to be ready by the end of 2008.   \\vspace{-0.4cm}"
    },
    "0801/0801.0645_arXiv.txt": {
        "abstract": "Observations from the Space Telescope Imaging Spectrograph define the flux of the DBQ4 star LDS749B from 0.12--1.0~$\\mu$m with an uncertainty of $\\sim$1\\% relative to the three pure hydrogen WD primary \\emph{HST} standards. With $T_\\mathrm{eff}=13575~K$, $\\log g=8.05$, and a trace of carbon at $<$1$\\times10^{-6}$ of solar, a He model atmosphere fits the measured STIS fluxes within the observational noise, except in a few spectral lines with uncertain physics of the line broadening theory. Upper limit to the  atmospheric hydrogen and oxygen fractions by number are 1$\\times10^{-7}$ and 7$\\times10^{-10}$, respectively. The excellent agreement of the model flux distribution with the observations lends confidence to the accuracy of the modeled IR fluxes beyond the limits of the STIS spectrophotometry. The estimated precision of $\\sim$1\\% in the predicted IR absolute fluxes at 30~$\\mu$m should be better than the model predictions for Vega and should be comparable to the absolute accuracy of the three primary WD models. ",
        "introduction": "The DBQ4 star LDS749B (WD2129+00) has long been considered for a flux standard (e.g.,  Bohlin et~al.\\ 1990). To establish the flux on the \\emph{Hubble Space Telescope} (\\emph{HST}) white dwarf (WD) flux scale, STIS spectrophotometry was obtained in 2001--2002. The virtues of LDS749B as a flux standard include an equatorial declination and a significantly cooler flux distribution than the 33000--61000~K primary DA standards GD71, GD153, and G191B2B. Full STIS wavelength coverage is provided from 0.115--1.02~$\\mu$m, and the peak in the SED is near 1900~\\AA. At $V=14.674$ (Landolt \\& Uomoto 2007), LDS749B is among the faintest \\emph{HST} standards and is suitable for use with larger ground-based telescopes and with the more sensitive \\emph{HST} instrumentation, such as the ACS/SBC and COS. The bulk of the STIS data was obtained as part of the FASTEX (Faint Astronomical Sources Extention) program. Finding charts appear in Turnshek et~al.\\ (1990) and in Landolt \\& Uomoto (2007); but there is a large proper motion of 0.416 and 0.034 arcsec/yr in right ascension and declination, respectively.  The absolute flux calibration of \\emph{HST} instrumentation is based on models of three pure hydrogen WD stars GD71, GD153, and G191B2B (Bohlin 2000; Bohlin, Dickenson, \\& Calzetti 2001; Bohlin 2003). In particular, the NLTE model fluxes produced by by the Tlusty code (Hubeny \\& Lanz 1995) determine the shape of the flux distributions using the known physics of the hydrogen atom and of stellar atmospheres. If there are no errors in the basic physics used to determine the stellar temperatures and gravities from the Balmer line profiles, then the uncertainty of 3000~K for the effective temperature of G191B2B means that the relative flux should be correct to better than 2.5\\% from 0.13 to 1~$\\mu$m and to better than 1\\% from 0.35 to 1~$\\mu$m.  A model that matches the observations serves as a noise free surrogate for the observational flux distribution and provides a reliable extrapolation beyond the limits of the observations for use as a calibration standard for \\emph{JWST}, \\emph{Spitzer}, and other IR instrumentation. Currently, the best IR absolute flux distributions are found in a series of papers from the epic and pioneering work of  M.~Cohen and collaborators, i.e., the Cohen-Walker-Witteborn (CWW) network of  absolute flux standard (e.g., Cohen, Wheaton, \\& Megeath 2003; Cohen 2007). The  CWW IR standard star fluxes are all ultimately based on models for Vega and Sirius (Cohen et~al.\\ 1992). More recently, Bohlin \\& Gilliland (2004) observed Vega and published fluxes on the \\emph{HST}/STIS WD flux scale. A small revision in the  STIS calibration resulted in excellent agreement of the STIS flux distribution with  a custom made Kurucz model with $T_\\mathrm{eff}=9400$~K (Bohlin 2007), which  is the same $T_\\mathrm{eff}$ used for the Cohen et~al.\\ (1992) Vega model.  The model presented here for LDS749B and archived in the CALSPEC database\\footnote{The absolute spectral energy distributions discussed in this paper are available in digital form at http://www.stsci.edu/hst/observatory/cdbs/calspec.html.} should have a better precision than the Kuruzc $T_\\mathrm{eff}=9400$~K model for Vega, especially beyond $\\sim$12~$\\mu$m, where the Vega's dust disk becomes important (Engleke, Price, \\& Kraemer 2006). Vega is also a pole-on rapid rotator, which may also cause IR deviations from the flux for a single temperature model. Our modeled flux distribution for LDS749B should have an accuracy comparable to the pure hydrogen model flux distributions for the primary WD standards GD71, GD153, and G191B2B. ",
        "conclusions": "In the absence of any interstellar reddening, a helium model with $T_\\mathrm{eff}=13575~K\\pm50$, $\\log g=8.05\\pm0.7$, and a trace of carbon at $<$1$\\times10^{-6}$ of solar fits the measured STIS flux distribution for LDS749B. The noise-free, absolute flux distribution from the model after normalization to the observed broadband visual flux is preferred for most purposes. This normalized model SED is a high fidelity far-UV to far-IR calibration source; and the flux distribution is available via Table~3 in the electronic version of the \\emph{Journal}. Both the observed flux distribution and the modeled fluxes are also available from the CALSPEC database.\\footnote{http://www.stsci.edu/hst/observatory/cdbs/calspec.html/.}"
    },
    "0801/0801.2476_arXiv.txt": {
        "abstract": "{ The chemical properties of galaxies and their evolution as a function of cosmic epoch are powerful constraints of their evolutionary histories. } { This work provides a grid of numerical models of galaxy evolution over an extended cosmic epoch. The aims are to assess how well current models reproduce observed properties of galaxies, in particular the stellar mass versus gas phase metallicity relation, and to quantify the effect of the merging histories of galaxies on their final properties. } { We use 112 N-body/hydrodynamical simulations in the standard Cold Dark Matter universe, to follow the formation of galaxy-sized halos and investigate the chemical enrichment of both the stellar component and the interstellar medium of galaxies, with stellar masses larger than $\\sim10^9$~M$_{\\odot}$. } { The resulting chemical properties of the simulated galaxies are broadly consistent with the observations. The predicted relationship between the mean metallicity and the galaxy stellar mass for both the stellar and the gaseous components at $z=0$ are in agreement with the relationships observed locally. The predicted scatter about these relationships, which is traced to the differing merging histories amongst the simulated galaxies with similar final masses, is similar to that observed. Under the hierarchical formation scenario, we find that the more massive galaxies are typically more evolved than their low mass counterparts over the second half of the age of the Universe. The predicted correlations between the total mass and the stellar mass of galaxies in our simulated sample from the present epoch up to $z\\sim 1$ agree with observed ones. We find that the integrated stellar populations in the simulations are dominated by stars as old as $4-10$~Gyr. In contrast with massive galaxies, for which the luminosity-weighted ages of the integrated stellar populations in the simulated sample agree with those derived from the modeling of observed spectral energy distributions, simulated galaxies with stellar masses $\\sim10^{9}$~M$_{\\odot}$ at $z=0$ tend to be older than the local galaxies with similar stellar masses. } { The stellar mass versus metallicity relation and its associated scatter are reproduced by the simulations as consequences of the increasing efficiency of the conversion of gas into stars with stellar mass, and the differing merging histories amongst the galaxies with similar masses. The old ages of simulated low mass galaxies at $z=0$, and the weak level of chemical evolution for massive galaxies suggest however that our modelling of the supernova feedback may be incomplete, or that other feedback processes have been neglected. } ",
        "introduction": "The chemical enrichment histories of galaxies provide insight into various processes involved in galaxy formation and evolution. The chemical composition of stars and gas within a galaxy depends on a number of physical processes, such as the star formation history, gas outflows and inflows, stellar initial mass function, etc. Although it is a complicated task to disentangle the effects of these processes, the galactic chemical abundances at various epochs place tight constraints on the likely evolutionary histories of galaxies.  The correlation between galaxy metallicity and luminosity in the local universe is one of the most significant observational results in galaxy chemical evolution studies. Lequeux et al. (1979) first revealed that the oxygen abundance increases with the total mass of irregular galaxies. The luminosity-metallicity relation for irregulars was later confirmed by Skillman et al. (1989) and Richer \\& McCall (1995), amongst others. Subsequent studies have extended the relation to spiral galaxies (e.g., Garnett \\& Shields 1987; Zaritsky et al. 1994; Garnett et al. 1997; Pilyugin \\& Ferrini 2000), and to elliptical galaxies (Brodie \\& Huchra 1991). The luminosity correlates with metallicity over $\\sim10$ magnitudes in luminosity and a factor of $\\sim100$ in metallicity, with indications suggesting that the relationship may be independent of environment (Vilchez 1995; Mouhcine et al. 2007) and morphology (Mateo 1998). More recently, large samples of star-forming galaxies drawn from galaxy redshift surveys, e.g. 2dF Galaxy Redshift Survey and the Sloan Digital Sky Survey (SDSS hereafter), have been used to confirm the existence of the luminosity-metallicity relation over a broad range of luminosity and metallicity (Lamareille et al. 2004; Tremonti et al. 2004). Lee et al. (2006) have extended the mass-metallicity relation to dwarf irregular galaxies, and found that the dispersion is similar over five orders of magnitude in stellar mass, and that the relation between the integrated stellar mass and the oxygen abundance of the interstellar medium is similarly tight from high stellar mass to low stellar mass galaxies. {\\bf The gas phase oxygen abundance vs. stellar mass has been understood either as a depletion sequence or a sequence in astration (see e.g. Tremonti et al. 2004). K\\\"oppen, Weidner, \\& Kroupa (2007) have presented an alternative explanation of the mass-metallicity relation arguing it could be the consequence of the integrated galactic initial mass function depending on the star formation rate, leading to higher oxygen yield in systems where the star formation rate is high. More massive galaxies are expected to have higher yields and thus have a tendency to have higher metallicities. Dalcanton (2007) presented a series of closed-box chemical evolution models including infall and outflow. She has shown that neither simple infall nor outflow can reproduce the observed low effective yields in low-mass galaxies (see Lee et al. 2006 for a similar conclusion), but metal-enriched outflows can do. That is, effectively the freshly synthesized elements need to be removed from the matter cycle, which is, as noted by K\\\"oppen et al. (2007) in principle, equivalent to reducing the number of massive stars.}  Many recent studies in galaxy evolution trace changes in scaling relations of galaxies at earlier epochs. In this context, the galactic chemical abundances at different cosmic epochs can assist in constraining the likely scenarios of galaxy evolution. Different groups have used the classical nebular diagnostic techniques developed to study the properties of {H{\\sc ii}} regions and emission line galaxies in the local universe to probe the properties of the interstellar gas in intermediate ($0<z<1$: Hammer et al. 2001; Lilly et al. 2003; Liang et al. 2006; Maier et al. 2004; Kobulnicky \\& Kewley 2004; Maier et al. 2005; Mouhcine et al. 2006ab; Lamareille et al. 2006) and high-redshift galaxies ($1.5<z<4$: Pettini et al. 1998; Kobulnicky \\& Koo 2000; Mehlert et al. 2002; Lemoine-Busserolle et al. 2003; Erb et al. 2006; Maier et al. 2006). The oxygen abundances for the interstellar medium for luminous star-forming galaxies at intermediate redshift are found to cover the same range as their local counterparts, and most of them fall on the local luminosity-metallicity relation. However, a minority appear to have significantly lower oxygen abundances than the local galaxies with similar luminosities (Kobulnicky et al.2003; Lilly et al. 2003; Liang et al. 2004; Mouhcine et al. 2006a). These luminous, massive, intermediate redshift, star-forming galaxies with low oxygen abundances have bluer colours than the higher metallicity ones, and exhibit physical conditions, i.e., emission line equivalent width and ionization state, very similar to those of the local faint and metal-poor star-forming galaxies (Mouhcine et al. 2006a; Maier et al. 2005). The intermediate redshift massive and large galaxies with low gas-phase oxygen abundances are most likely immature galaxies that will increase their metallicities and their stellar masses to the present epoch. The diversity of the properties of the massive and large galaxies at intermediate redshifts supports the scenario whereby galaxies are still assembling their baryonic content between $z\\sim 1$ and $z=0$ (Hammer et al. 2005). For a sample of mostly disk, intermediate redshift field and cluster galaxies, Mouhcine et al. (2006a) have shown that oxygen abundances do not correlate with either the maximum rotation velocity, i.e. a proxy of the total mass, or with the emission scale length size of the galaxy (see also Lilly et al. 2003). {\\bf Savaglio et al. (2005) have shown that the observed redshift evolution of the mass-metallicity relation could be well reproduced by a simple closed-box model where the star formation timescale is proportional to the galaxy total baryonic mass. }  As these observational data for more statistically significant samples of galaxies accumulate, numerical simulations become important to understand galaxy formation in the context of the currently favored hierarchical clustering scenario. To be compared with the observational data -- which provides not only dynamical information, such as stellar and gas mass and rotation velocity, but also properties of stellar content, such as metallicity and age -- such simulations are required to follow both the dynamical and chemical evolution of galaxies (e.g.: Mosconi et al. 2001; Lia et al. 2002; Kawata \\& Gibson 2003; Kobayashi et al. 2006). Using chemodynamical numerical simulations, we have built an ensemble of simulated late-type galaxies, spanning a factor of 50 in total mass, and sampling a range of assembly histories at a given total mass. We compare our large sample of simulated galaxies with recent observational results, and examine how the current simulations can explain the observations. We pay particular attention to the observed correlations between the galaxy mass and metallicity of both the gaseous and stellar components, key observables in disentangling the formation history of galaxies. In Section~\\ref{simul} we describe our numerical simulations. Section~\\ref{resl} provides comparisons between simulated galaxies and observations. We summarize our conclusions in Section~\\ref{concl}. ",
        "conclusions": "\\label{concl}  We have presented an analysis of the chemical properties over the last half of the age of the Universe for a sample of 112 simulated late-type galaxies, with stellar masses larger than $\\sim 10^9$~M$_{\\odot}$, formed under the hierarchical clustering scenario. The simulations include hydrodynamics, star formation, metal-dependent cooling, supernova feedback, and a chemical enrichment of the interstellar medium that relaxes the instantaneous recycling approximation. Stellar feedback is taken into account by injecting energy into the gas that surrounds regions of recent star formation. We compared the observed properties of the simulated galaxies with the recent observational data at various redshifts. A particular emphasis has been placed upon the relationship between the stellar mass and the metallicity of both the gaseous and the stellar components. The main results can be summarised as follows.  \\begin{itemize} \\item[-] In the hierarchical clustering scenario, even after the system has collapsed, the histories of the growth of the stellar mass in the central region can be different (Fig.~\\ref{mass_assembly}) due to the variety in the minor merger and gas fueling histories, which also leads to the diverse fraction of the gas-to-stellar mass ratio among the galaxies with similar stellar or total mass (Fig.~\\ref{gf_ms}).  \\item[-] The stellar mass in the inner region correlates well with the total mass of galaxies up to $z\\sim 1.2$ (Fig.~\\ref{mtot_ms_obs}), which broadly agrees with the observed relation, especially for higher mass (${\\rm M_{tot}}>10^{10}$ M$_{\\sun}$) galaxies. Lower mass galaxies have greater scatter. For some simulated galaxies, star formation is heavily suppressed and leads to lower stellar mass compared with the total mass.  \\item[-] We have analyzed the metallicity for both stellar and gas components. Both are correlated with the stellar mass up to $z\\sim 1.2$, and show an agreement with recent observations (Figs.~\\ref{oh_ms_gas} and \\ref{zs_mstar}). This relation arises from the fact that higher mass galaxies convert gas more efficiently into stars.  \\item[-] Higher mass galaxies reach a high metallicity earlier. As a result, the slope of the mass-metallicity relation becomes flatter with decreasing redshift (Figs.~\\ref{oh_ms_gas} and \\ref{zs_mstar}). Lower mass galaxies (${\\rm M_{tot}}<10^{10}$ M$_{\\sun}$) evolve chemically more slowly, and have lower metallicity than that predicted from the extrapolation of the mass-metallicity relation for higher mass galaxies at lower redshifts, which is also consistent with the observational data.  \\item[-] The predicted scatter in the mass-metallicity relation also reproduces qualitatively the observed scatter (Figs.~\\ref{oh_ms_gas} and \\ref{zs_mstar}). This suggests that the observed scatter could be due to the difference in mass assembly histories, a natural prediction of the hierarchical clustering scenario (Fig.~\\ref{mass_assembly}).  \\end{itemize}  The primary disagreement between our simulations and observations is as follows:  \\begin{itemize} \\item[-] Our simulations predict almost no correlation between the stellar age and the stellar mass, which cannot explain the observed significantly younger population in lower mass galaxies, e.g. stellar masses lower than ${\\rm \\sim 10^{9} M_{\\odot}}$ (Fig.~\\ref{agestar_mstar}). As a result, the simulations also predict too flat of a colour-luminosity relation (Fig.~\\ref{cmr}). \\end{itemize}  This suggests that the simulations need a mechanism to sustain a low level of star formation activity and also keep the stellar and gas metallicity low. A combination of stronger supernovae feedback and the UV background radiation may help to cause such self-regulated star formation in the lower mass galaxies. This demonstrates that the quantitative comparisons made in the paper are invaluable to improve numerical simulation models, which eventually will aid in completing our understanding of the physical processes governing galaxy formation and evolution."
    },
    "0801/0801.0968_arXiv.txt": {
        "abstract": "Clusters of galaxies are self-gravitating systems of mass $\\sim 10^{14}-10^{15} h^{-1}$ M$_\\odot$ and size $\\sim 1-3 h^{-1}$ Mpc. Their mass budget consists of dark matter ($\\sim 80\\%$, on average), hot diffuse intracluster plasma ($\\lesssim 20\\%$) and a small fraction of stars, dust, and cold gas, mostly locked in galaxies. In most clusters, scaling relations between their properties, like mass, galaxy velocity dispersion, X-ray luminosity and temperature, testify that the cluster components are in approximate dynamical equilibrium within the cluster gravitational potential well. However, spatially inhomogeneous thermal and non-thermal emission of the intracluster medium (ICM), observed in some clusters in the X-ray and radio bands, and the kinematic and morphological segregation of galaxies are a signature of non-gravitational processes, ongoing cluster merging and interactions. Both the fraction of clusters with these features, and the correlation between the dynamical and morphological properties of irregular clusters and the surrounding large-scale structure increase with redshift.  In the current bottom-up scenario for the formation of cosmic structure, where tiny fluctuations of the otherwise homogeneous primordial density field are amplified by gravity, clusters are the most massive nodes of the filamentary large-scale structure of the cosmic web and form by anisotropic and episodic accretion of mass, in agreement with most of the observational evidence. In this model of the universe dominated by cold dark matter, at the present time most baryons are expected to be in a diffuse component rather than in stars and galaxies; moreover, $\\sim 50\\%$ of this diffuse component has temperature $\\sim 0.01-1$ keV and permeates the filamentary distribution of the dark matter. The temperature of this Warm-Hot Intergalactic Medium (WHIM) increases with the local density and its search in the outer regions of clusters and lower density regions has been the quest of much recent observational effort.  Over the last thirty years, an impressive coherent picture of the formation and evolution of cosmic structures has emerged from the intense interplay between observations, theory and numerical experiments. Future efforts will continue to test whether this picture keeps being valid, needs corrections or suffers dramatic failures in its predictive power. ",
        "introduction": "The present chapter provides the general framework of the physics of clusters of galaxies: it outlines how clusters connect to several astrophysical issues, from cosmology to the formation of galaxies and stars. Some of these topics are discussed in the chapters of this volume; we refer to these chapters when appropriate. When the topic we mention here is not dealt with elsewhere in this volume, we refer to some of the most recent reviews. ",
        "conclusions": ""
    },
    "0801/0801.0703_arXiv.txt": {
        "abstract": "{} {We study the visibility of sunspots and its influence on observed values of sunspot region parameters.} {We use Virtual Observatory tools provided by AstroGrid to analyse a sample of 6862 sunspot regions. By studying the distributions of locations where sunspots were first and last observed on the solar disk, we derive the visibility function of sunspots, the rate of magnetic flux emergence and the ratio between the durations of growth and decay phases of solar active regions.} {We demonstrate that the visibility of small sunspots has a strong centre-to-limb variation, far larger than would be expected from geometrical (projection) effects. This results in a large number of young spots being invisible: 44\\% of new regions emerging in the west of the Sun go undetected. For sunspot regions that are detected, large differences exist between actual locations and times of flux emergence, and the apparent ones derived from sunspot data. The duration of the growth phase of solar regions has been, up to now, underestimated. } {} ",
        "introduction": "The birth of a new spot on the solar disk indicates the emergence of magnetic flux through the photosphere, a process which is key to the solar cycle \\cite{Sol2003,Fis2000} and the study of stellar magnetic dynamos. Sunspots also cause variations in the total solar irradiance, an important parameter in determining the Sun's influence on climate \\cite{Fou2006}. The presence of a sunspot is key to a solar region being assigned an active region number by the NOAA Space Weather Prediction Center (http://www.swpc.noaa.gov) so that its evolution and activity can be tracked \\cite{Gal2007}. The formation and evolution of active regions are fundamental to solar dynamic phenomena such as flares and Coronal Mass Ejections and their effect on the Earth environment.  The visibility of sunspots is currently thought to be limited only by geometrical effects arising from projection of the solar sphere onto a 2D image, effects referred to as foreshortening. In this paper we present results obtained serendipitously while analysing sunspot data by means of Virtual Observatory tools, showing that the visibility of small sunspots is much poorer than predicted by the foreshortening model. ",
        "conclusions": "Figure 3-b demonstrates that the visibility of small sunspots is much poorer than expected from geometrical effects associated with foreshortening. It shows that the minimum area required for a spot to be detected at $\\lambda$=$\\pm$30$^{\\circ}$ is more than twice the threshold area at  $\\lambda$=0$^{\\circ}$ if $k$=14 msh/day and almost 4 times if $k$=30 msh/day.  The centre-to-limb variation in visibility is remarkably large.  We investigated whether the asymmetries shown in Fig. 1 display any solar cycle dependence, and found no evidence of it. The USAF/Mount Wilson locations and times of first appearances agree with those from SOHO/MDI continuum data (as verified manually for a sample of regions). It is known that seeing associated with ground-based data does not cause significant reduction of visibility \\cite{Gyo2004}.  How many spots are affected by the visibility effects here described? The distribution of sunspot areas measured at a single longitudinal location on the solar disk is lognormal \\cite{Bog1988} and the number of spots with area at Central Meridian around 10 msh is more than 2 orders of magnitude larger than the number of spots with area of about 100 msh. Therefore the majority of sunspots will cross the visibility curve shown in Fig. 3-b.  Our results have a number of important implications. The first is that the radiative processes that make a small region of strong magnetic field appear as a dark spot, have a strong centre-to-limb variation. This may prove important for the study of sunspots' 3D structure and will require further investigation. Whether larger spots are also affected by the same process will also need further study. Reports of centre-to-limb variations of corrected sunspot areas have appeared in the literature \\cite{Gyo2004, Hoy1983}. Faculae have a large centre-to-limb variation, and their contrast changes sign as one moves towards the disk centre, resulting in their being darker than the surrounding photosphere at the disk centre \\cite{Law1993}. This demonstrates that the appearance of photospheric magnetic flux tubes strongly depends on the viewing angle.  The second implication  is that actual distributions of sunspot lifetimes and areas may differ from the apparent ones derived from observations. The latter have been used to constrain mechanisms of sunspot formation and decay \\cite{Sol2003, Pet1997, Mar1993}. A large number of regions reported  of short duration may in fact have longer lifetimes, and be crossing in and out of the visibility curve. The Gnevyshev-Waldmeier law, stating that a sunspot's lifetime increases linearly with its maximum size, may need to be reassessed in light of our results. Poor visibility of region emergence means that the actual time of magnetic flux emergence can be much earlier than the apparent time, e.g.~for a region seen to emerge at $\\lambda$=$-$50$^{\\circ}$, by approximately 2 days. On the other hand, a large fraction of new emergences in the western portion of the solar disk go undetected. By using the data in the inset of Fig.~1-a for $\\lambda$$>$0 and the value of actual number of emergences $N_1$=160 obtained from the data of Fig. 2, we obtain that 44\\% of new spots emerging in [0$^{\\circ}$, $+$60$^{\\circ}$] were invisible.  The presence of a sunspot is key to a solar active region being assigned an Active Region Number by the NOAA Space Weather Prediction Center (see e.g. Dalla et al. (2007) for the full list of criteria). A region that has been given a NOAA number is monitored and its activity tracked \\cite{Gal2007}. New regions emerging in the west of the Sun with their spots being invisible are missed and not tracked. We conclude that current criteria for assigning Active Region Numbers may need to be revised and that EUV solar images may need to be routinely used to supplement white-light information.  Sunspots are well known to cause depletions in the total solar irradiance (TSI) \\cite{Fou2006} and many models of TSI variation need as input information on the number and areas of sunspots. % While the sunspots that most affect TSI are the largest ones, of area typically well above the visibility threshold shown in Fig. 3-b, our results impact TSI studies because they demonstrate that the apparent age and stage of development of a sunspot may not correspond to the actual ones. The latter information is required when studying the time dependence of sunspot effects on TSI and whether the age of a region is an important factor in determining the magnitude of TSI decrease.  The asymmetry in the distribution of emergences was obtained, unexpectedly, during a study aiming at cross correlating catalogues of sunspot regions and flares, by means of AstroGrid workflows \\cite{Dal2007}. This demonstrates the usefulness of VO tools in making new science possible, by provision of better tools for analysis of large datasets."
    },
    "0801/0801.1828_arXiv.txt": {
        "abstract": "Two recent issues realted to nucleosynthesis in early proton-rich neutrino winds are investigated. In the first part we investigate the effect of nuclear physics uncertainties on the synthesis of $^{92}$Mo and $^{94}$Mo. Based on recent experimental results, we find that the proton rich winds of the model investigated here can not be the only source of the solar abundance of $^{92}$Mo and $^{94}$Mo. In the second part we investigate the nucleosynthesis from neutron rich bubbles and show that they do not contribute to the nucleosynthesis integrated over both neutron and proton-rich bubbles and proton-rich winds. ",
        "introduction": "Over the past decade improvements in neutrino-transport and multi-dimensional computer simulations have lead to a new understanding of the conditions that lead to nucleosynthesis of the elements above iron in core-collapse supernovae. Immediately following the bounce on the proto-neutron star, the shock fully photodisintegrates the infalling material turning it into electron--position pairs, neutrons, and protons. As the nascent neutron star continues to collapse it liberates $10^{53} \\textrm{ergs}$ over the span of $\\sim 10$ seconds primarily in the form of neutrinos. This enormous neutrino flux is deposited in the low density region of photodisintegrated matter inside the gain radius between the neutron star and the accretion shock of the still infalling material and heats it to temperatures in excess of 10 billion K while driving mass away in the form of a neutrino wind theoretically leading to the explosion of the supernova \\cite{Qian96}. The strong flux of neutrinos and anti-neutrinos results in a detailed balance between protons and neutrons that favors the lighter mass protons depending on the respective neutrino spectra leading to an electron fraction that is proton-rich ($Y_e>0.5$) \\citep{Liebendoerfer03, Froehlich06b}. These protons and neutrons recombine into alpha particles that proceed via the $\\alpha(\\alpha n,\\gamma)$ ${}^{9}\\textrm{Be}(\\alpha,n)$ ${}^{12}\\textrm{C}$-reactions followed by a series of $(\\alpha,\\gamma)$-reactions or combined $(\\alpha,p)(p,\\gamma)$-reactions along $N=Z$ into the iron group, primarily ${}^{56}\\textrm{Ni}$ and ${}^{60}\\textrm{Zn}$ which form the seeds of the subsequent nucleosynthesis.  From this point the resulting nucleosynthesis in the neutrino-driven wind essentially depends on the number of seed nuclei to the number of excess neutrons or protons that were frozen out and did not turn into seed nucleii ($Y_e$), the entropy per baryon, the expansion timescale of the ejecta and the amount of the ejecta. As the explosion evolves, an ejected mass element inherits some combination of these parameters and below $\\sim 0.5\\textrm{MeV}$ they remain fairly constant as the matter proceeds to freeze out.  In this paper, we consider the early times when the wind still contains a proton excess because the rates for neutrino and positron captures on neutrons are faster than those for the inverse captures on protons. We consider two interesting problems which are discussed in the following two sections. ",
        "conclusions": ""
    },
    "0801/0801.3745.txt": {
        "abstract": "We describe the HST ACS Coma cluster Treasury survey, a deep two-passband imaging survey of one of the nearest rich clusters of galaxies, the Coma cluster (Abell 1656).  The survey was designed to cover an area of 740 arcmin$^2$ in regions of different density of both galaxies and intergalactic medium within the cluster. The ACS failure of January 27th 2007 leaves the survey 28\\% complete, with 21 ACS pointings (230 arcmin$^2$) complete, and partial data for a further 4 pointings (44 arcmin$^2$).  Predicted survey depth for 10$\\sigma$ detections for optimal photometry of point sources is g$^{\\prime}$ = 27.6 in the F475W filter, and I$_C$=26.8 mag in F814 (AB magnitudes). Initial simulations with artificially injected point sources show 90\\% recovered at magnitude limits of g$^{\\prime}$ = 27.55 and I$_C$ = 26.65. For extended sources, the predicted 10$\\sigma$ limits for a 1 arcsecond$^{2}$ region are g$^{\\prime}$ = 25.8 mag arcsec$^{-2}$ and I$_C$ = 25.0 mag arcsec$^{-2}$.  We highlight several motivating science goals of the survey, including study of the faint end of the cluster galaxy luminosity function, structural parameters of dwarf galaxies, stellar populations and their effect on colors and color gradients, evolution of morphological components in a dense environment, the nature of ultra compact dwarf galaxies, and globular cluster populations of cluster galaxies of a range of luminosities and types. This survey will also provide a local rich cluster benchmark for various well known {\\it global} scaling relations and explore new relations pertaining to the {\\it nuclear} properties of galaxies. ",
        "introduction": "\\label{sec:introduction}  The Coma cluster of galaxies, Abell 1656, is along with the Perseus cluster the nearest rich, and dense cluster environment. Unlike the Perseus cluster, it is at high galactic latitude (b = 87.9$^{\\circ}$) and it has been a popular target for study at all wavelengths. Progressively deeper wide-area photometric surveys of Coma have become available over the past 30 years, and waveband coverage has spread from the original B and V band surveys into the near ultra-violet and infra-red (Godwin \\& Peach 1977; Godwin, Metcalfe \\& Peach 1983 (GMP); Komiyama et al.\\ 2002; Adami et al.\\ 2006; Eisenhardt et al.\\ 2007). A larger area around Coma is covered by the imaging part of Data Release 5 of the Sloan Digital Sky Survey (Adelman-McCarthy et al. 2007). From these surveys, samples of galaxies have been selected for spectroscopic study, which has resulted in an understanding of the internal dynamics of the cluster (Colless \\& Dunn 1996;  Mobasher et al.\\ 2001; Edwards et al.\\ 2002; Gutierrez et al.\\ 2004), the internal dynamics of cluster members (Lucey et al. 1991; Jorgensen et al.\\ 1996; Smith et al.\\ 2004; Matkovic \\& Guzm\\'an 2005; Cody et al. 2007), and their mean luminosity weighted stellar ages, abundances and $\\alpha$-enhancement (Bower et al. 1992a,b; Guzm\\'an et al. 1992; Caldwell et al. 1993; Jorgensen 1999; Poggianti et al. 2001; Moore et al. 2002; Nelan et al. 2005; Smith et al. 2006; Matkovic et al. 2007). Coma presents us with the best opportunity to study large samples of galaxies of different luminosity, environment and morphological type, but at a common distance, and with a common Galactic extinction.  There is good agreement on the distance to the Coma cluster, with independent studies using six different methods yielding values in the range 84 -- 108 Mpc, as summarised in Table \\ref{tab:Distances}. These values fit well with the current concordance cosmology: assuming $H_0$ = 71 km/s/Mpc; $\\Omega_{\\lambda}$ = 0.73; $\\Omega_m$ = 0.27, and a redshift z = 0.0231 gives a distance of 99.3 Mpc. In this paper we assume a distance of 100 Mpc, equivalent to a distance modulus of 35.00. \\clearpage \\begin{deluxetable}{lccl} \\tablecolumns{4} \\tablewidth{0pc} \\tablecaption{\\label{tab:Distances}The Distance to Coma determined by different techniques} \\tablehead{ \\colhead{Technique}&\\colhead{Distance (Mpc)}&\\colhead{Distance Modulus}&\\colhead{Reference}} \\startdata I-band Tully-Fisher&86.3$\\pm$6&34.68$\\pm$0.15&Tully \\& Pierce (2000)\\\\ K$^{\\prime}$-band SBF&85$\\pm$10&34.64$\\pm$0.27&Jensen et al. (1999)\\\\ I-band SBF&102$\\pm$14&35.04$\\pm$0.32&Thomsen et al. (1997)\\\\ $D_n-\\sigma$&96$\\pm$6&34.90$\\pm$0.14&Gregg (1997)\\\\ Fundamental Plane&108$\\pm$12&35.16$\\pm$0.25&Hjorth \\& Tanvir (1997)\\\\ Globular Cluster LF&102$\\pm$6&35.05$\\pm$0.12&Kavelaars et al. (2000)\\\\ \\enddata \\end{deluxetable} \\clearpage  Coma lies in a rich region of the large-scale distribution of galaxies, at the intersection of a number of filaments. Figure \\ref{fig:coma_xyz} shows two projections of the distribution of galaxies in supergalactic co-ordinates in the region of Coma and the nearby richness class 2 cluster Abell 1367. The Great Wall (Geller \\& Huchra 1989) a vertical structure in the two panels of Figure \\ref{fig:coma_xyz}, runs through these two clusters, other filaments intersect the Great Wall at the Coma cluster. \\clearpage \\begin{figure*} \\begin{center} \\includegraphics[width=12cm]{f1.ps} \\caption{\\label{fig:coma_xyz}The location and environment of the Coma cluster in supergalactic co-ordinates. The left panel shows a projection onto supergalactic Y, close to the plane of the sky. The right panel shows a projection onto supergalactic X, in this panel the horizontal (Y) axis is close to measured redshift. All axes are in equivalent cz. Positions of the points are derived from measured sky positions and redshifts. Members of Coma (centre) and Abell 1367 (bottom) are plotted as filled red circles, for these galaxies the redshift used to compute the distance is the cluster redshift, with a random velocity offset chosen to make the cluster appear round. For non cluster members, measured cz is used, and the positions are plotted as filled black circles if they have velocities within $\\pm$600 km/s of the Coma mean in the depth dimension in that particular panel, and as open grey circles otherwise. The Great Wall is the structure at SGY $\\sim 6800$ km/s in the right panel and is seen face on in the left panel} \\end{center} \\end{figure*} \\clearpage There is a uniquely rich multi-wavelength dataset on the  Coma cluster. X-ray observations covering a large area of the cluster have been made with ROSAT (White et al. 1993) and XMM-Newton (Briel et al.\\ 2001), showing the distribution and properties of the hot intra-cluster medium (ICM), and X-ray properties of individual galaxies have been studied by Finoguenov et al.\\ (2004) and by Hornschemeier et al. (2006). Coma has been shown by INTEGRAL to be an extended hard X-ray/soft $\\gamma$-ray source (Renaud et al.\\ 2006). At soft X-ray and Extreme Ultraviolet wavelengths there is a thermal excess  (Lieu et al. 1996; Kaastra et al. 2003; Bonamente et al. 2003; Bowyer et al.\\ 2004). GALEX has been used to observe the cluster in the mid and near Ultraviolet  and has sufficient spatial resolution to measure the UV properties of individual galaxies. In the infra-red, studies of the galaxies have been made using SPITZER, both with IRAC at 3.6 - 8 $\\mu$m (Jenkins et al.\\ 2007) and with MIPS at 24 and 70 $\\mu$m (Bai et al.\\ 2006). At radio wavelengths, VLA continuum maps cover much of the cluster (Miller et al.\\ in preparation), and in the HI 21 cm line there are extensive VLA imaging surveys (Bravo-Alfaro et al.\\ 2000, 2001), and single-dish spectra and fluxes for samples of spiral galaxies (Gavazzi et al.\\ 2006; Vogt et al.\\ 2004).  This wealth of existing data makes Coma a prime target for studies of the origin and evolution of the galaxy content of clusters, and of its interaction with the other components (gas and dark matter). Moreover, as the richest and best studied local cluster, Coma is the zero-redshift baseline for many studies of high-redshift clusters (e.g. Jorgensen et al.\\ 2006). Comparison between low- and high-redshift clusters is vital for our understanding of their evolution, which in turn is essential if we are to disentangle evolutionary effects from the properties which tell us about their formation.  We describe here the HST/ACS Coma Cluster Treasury survey, which aims to provide an unparalleled database of high spatial resolution images of a sample of cluster galaxies. At the distance of the Coma cluster ($\\sim$100 Mpc), the resolution of HST/ACS (0{\\arcsec}.1) corresponds to $\\sim$50 pc. This gives essentially the same physical resolution as ground-based observations have in Virgo and Fornax. Thus the HST/ACS Coma database provides a valuable comparison between high- and low-density clusters, for studies of the effect of environment on galaxy components.  Whilst the HST observations are the prime data upon which this survey is based, it has already generated numerous ancillary observations with facilities such as Subaru, Keck, MMT, UKIRT. It is concurrent with surveys of the cluster in other wavebands, including the ultra-violet (GALEX), infra-red (SPITZER), X-ray (Chandra and XMM-Newton) and radio (VLA). ",
        "conclusions": ""
    },
    "0801/0801.4316_arXiv.txt": {
        "abstract": "We report millimeter interferometric observations of polarized continuum and line emission from the  massive star forming region G34.4. Polarized thermal dust emission at 3 mm wavelength and CO $J=1 \\rightarrow 0$ line emission were observed using the Berkeley-Illinois-Maryland Association (BIMA) array. Our results show a remarkably uniform polarization pattern in both dust and in CO J=$1 \\rightarrow 0$ emission. In addition, the line emission presents a consistent uniform polarization pattern over most of the velocity channel maps. These uniform polarization patterns are aligned with the north-south main axis of the filament between the main millimeter source (MM) and the ultra-compact H {\\scriptsize II} region, which are the central sources in G34.4, suggesting a magnetic field orthogonal to this axis. This morphology is consistent with a magnetically supported disk seen roughly edge-on. ",
        "introduction": "It is generally accepted that magnetic fields play an important role in the process of star formation; magnetic fields are involved in cloud support, fragmentation, and transfer of angular momentum. However, the magnetic field is the least observed physical quantity involved in such process. Magnetic field observations of molecular clouds are divided into measurements of the line-of-sight component of the magnetic field strength through the Zeeman effect and observations of the field in the plane of the sky through linear polarization of dust emission and spectral-line emission. The alignment of dust grains by a magnetic field is physically complicated and is still a matter of intense research. It is accepted, though, that aligned dust grains  will produce polarized emission perpendicular to the projection of the magnetic field onto the plane of the sky. For a recent review of alignment theories see \\citet{Lazarian2007}.   Spectral-line linear polarization has been suggested to arise from molecular clouds under anisotropic conditions, like large velocity gradients (or LVGs) \\citep{Goldreich1981}. The prediction suggests that a few percent of linearly polarized radiation should be detected from molecular clouds and circumstellar envelopes in the presence of a magnetic field. This polarization will be either parallel or perpendicular to the projection of the field onto the plane of the sky. To obtain a qualitative understanding about this effect, consider the CO molecule emitting unpolarized radiation. Under a magnetic field, a CO molecule will develop a small splitting in its rotational, $J$, energy  levels. These magnetic sub-levels will produce radiation components labeled $\\sigma$ for the $\\left|  M-M \\acute{} \\right| $ $=1$ transitions, and $\\pi$ for the $\\left|  M-M \\acute{} \\right|  $ $=0$ transitions, where both components can be linearly polarized either perpendicular or parallel to the magnetic field. If the gas behaves under isotropic conditions (e.g. no velocity gradients) for any direction the $\\sigma$ and $\\pi$ will populate equally, so the radiation components emerging from the radiative decays of these states will combine to give zero net polarization. Now on the other hand, large velocity gradients present in molecular clouds will produce anisotropies in the optical depths for the CO molecular transitions at different directions. If the velocity gradients are smaller in directions parallel to the magnetic field than in those perpendicular to the field,  the optical depths parallel to the field will be larger than the optical depths perpendicular to them. Therefore, the escape of radiation involved in de-exciting the upper $J$ state will  be then reduced more in directions along the field lines than in directions perpendicular to them,  which will lead to populations of the magnetic $\\sigma$ sub-states that are larger than the populations of the $\\pi$ sub-states due to the difference in the angular distributions of both radiation components. The angular distribution of the $\\sigma$ component peaks in directions along the field lines whereas the $\\pi$ component peaks in directions perpendicular to the magnetic field. In this way, the rate of de-excitation for $\\sigma$ will have a larger decrease, due to photon trapping, than the rate of de-excitation for the $\\pi$ radiation component. Because, in this picture, the $\\sigma$ component will have a larger population, its emission will be stronger relative to the $\\pi$ component giving raise to a small amount of linear polarization in the CO emission with the polarization of the $\\sigma$ component, or perpendicular to the magnetic field. The effect was first detected by \\citet{Glenn1997} for the CS molecules while \\citet{Greaves1999} detected CO polarized emission for $(J=2 \\rightarrow 1)$ and $(J=3 \\rightarrow 2)$ transitions.   In order to efficiently map polarized emission and infer detailed information about the magnetic field morphology, high resolution observations are required. The BIMA millimeter interferometer has been used previously to obtain high-resolution polarization maps in several star forming cores \\citep{Lai2001,Lai2002, Lai2003,Cortes2005,Cortes2006a,Cortes2006b}. These previous results show fairly uniform polarization morphologies over the main continuum sources, suggesting that magnetic fields are strong, and therefore cannot be ignored by star formation theory. However, the number of star formation regions with maps of magnetic fields remains small, and every new result is statistically significant. In this work we  present polarization maps of the massive star forming region G34.4, obtained with the BIMA array.  We measured continuum polarization at 3 mm and CO $J=1\\rightarrow 0$ line polarization, obtaining interferometric maps for both line and continuum. The remainder of this paper is divided into five  sections. Section 2 reviews information about the source, section 3 describes the observation procedure, section 4 presents the results, section 5 gives the discussion, and section 6 the conclusions and summary. ",
        "conclusions": "The G34.4+0.23 MM massive star forming core was observed in the 3 mm band in polarized continuum and in CO $J=1 \\rightarrow 0$ polarized line emission with the BIMA array. The data show a uniform polarization pattern in both emissions. The P.A. obtained from the continuum data has an average value of $<\\phi> = -8^{\\circ} \\pm 5^{\\circ}$, and  from the CO polarization $<\\phi> = -2^{\\circ} \\pm 8^{\\circ}$. These results  suggest a magnetic field perpendicular to the main axis of the filament, between the MM source and the UC H {\\scriptsize II} region, in G34.4. The morphology suggests a flattened disk with the magnetic field along the minor axis, as predicted by the theory of magnetically supported molecular clouds. From our  3 mm continuum observations, we estimate  a total core mass of 400 $M_{\\sun}$, in agreement with previous observations. Both line and continuum emission agrees in morphology with previous work by \\citet{Shepherd2004,Shepherd2007}. Finally, additional observations, particularly higher resolution interferometric mapping along the filament (including the most northern source seen in Figure \\ref{1}),  will help to constrain, in greater detail, the morphology of the field and to obtain estimates of its strength in the plane of the sky.  P. C. Cortes acknowledges support from the ALMA-CONICYT fund for development of Chilean Astronomy through grant 31050003. P. C. Cortes would also like  to acknowledge the support given by NCSA and the Laboratory for Astronomical Imaging at University of Illinois at Urbana-Champaign during this research. Finally, P. C. Cortes would like to acknowledge the contribution by Patricio Sanhueza in making Figure \\ref{1}. R. M. Crutcher acknowledges support from NSF grants AST 05-40459 and 06-06822. L. Bronfman acknowledges support from the Chilean Center for Astrophysics FONDAP 15010003.  \\altaffiltext{1}{National Radio Astronomy Observatory, P.O. Box O, 1003 Lopezville Rd, Socorro, NM 87801.} \\altaffiltext{2}{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}"
    },
    "0801/0801.4299_arXiv.txt": {
        "abstract": "{The coming Planck and Herschel missions will survey the sky at unprecedented angular scales and sensitivities. Simulations are needed for better interpretating the results of the surveys and for testing new methods of, e.g., source extraction and component separation.} {We present new simulations of the infrared and submillimeter cosmic background, including the correlation between infrared galaxies. The simulations were used to quantify the source-detection thresholds for Herschel/SPIRE and Planck/HFI, as well as to study the detectability of the cosmic infrared background correlated fluctuations.} {The simulations are based on an empirical model of IR galaxy evolution. For these correlations, we only included the linear clustering, assuming that infrared galaxies are biased tracers of the dark-matter fluctuation density field.} {We used the simulations with different bias parameters to predict the confusion noise for Herschel/SPIRE and Planck/HFI and the completeness levels.  We also discuss the detectability of the linear clustering in Planck/HFI power spectra, including the foreground and backgrounds components.} {Simulated maps and catalogs are publicly available online at http://www.ias.u-psud.fr/irgalaxies/simulations.php. } ",
        "introduction": "The cosmic infrared background (CIB) ($\\lambda\\geq8\\mu m$) is the relic emission of the formation and evolution of galaxies. The first observational evidence of this background was reported by \\citet{1996A&A...308L...5P} and then confirmed by \\citet{1998ApJ...508...25H} and \\citet{1998ApJ...508..123F}. The discovery of a surprisingly high amount of energy in the CIB has shown the importance of studying its sources to understand how the bulk of stars was formed in the Universe. Deep cosmological surveys have been carried out thanks to ISO \\citep [see] [for reviews] {2000ARA&A..38..761G,2005SSRv..119...93E} mainly at 15 $\\mu m$ with ISOCAM \\citep [e.g.] [] {2002A&A...384..848E}; at 90 and 170 $\\mu m$ with ISOPHOT \\citep [e.g.] [] {2001A&A...372..364D}; to SPITZER at 24, 70, and 160 $\\mu m$ \\citep [e.g.] [] {2004ApJS..154...70P,2004ApJS..154...87D} and to ground-based instruments such as SCUBA \\citep [e.g.] [] {1998astro.ph..9121H} and MAMBO \\citep [e.g.] []{2000astro.ph.10553B} at 850 and 1300 $\\mu m$, respectively. These surveys have allowed for a better understanding of the CIB and its sources \\citep [see] [for a general review] {2005ARA&A..43..727L}. Some of the results include: the energy of the CIB is dominated by starbursts although AGN (active galactic nucleus) contribute too, and the dominant contributors to the energy output are the LIRGs (luminous IR galaxies) at $z\\sim 1$ and ULIRGs (ultra luminous IR galaxies) at $z\\sim 2-3$. \\\\   Determination of the CIB by the COBE satellite has been hindered by the accuracy of subtracting the foreground by only providing just upper limits at 12, 25, and 60 $\\mu m$ \\citep {1998ApJ...508...25H}, lower limit has been derived at 24 $\\mu m$ by \\citet{2004ApJS..154...70P} as well as the contribution of 24 $\\mu m$ galaxies to the background at 70 and 160 $\\mu m$ \\citep {2006A&A...451..417D}. The contribution of the galaxies down to 60 $\\mu Jy$ at 24 $\\mu m$ is at least 79\\% of the 24 $\\mu m$ background, and 80\\% of the 70 and 160 $\\mu m$ background. For longer wavelengths, recent studies have investigated the contribution of populations selected in the near-IR to the far-infrared background (FIRB, $\\lambda > 200 \\mu m$): 3.6 $\\mu m$ selected sources to the 850 $\\mu m$ background \\citep {2006ApJ...647...74W} and 8 $\\mu m$ and 24 $\\mu m$ selected sources to the 850 $\\mu m$ and 450 $\\mu m$ backgrounds \\citep{2006ApJ...644..769D}. Similar studies with Planck and Herschel will provide even more evidence of the nature of the FIRB sources.\\\\   Studying correlations in the spatial distribution of IR galaxies as a function of redshift is an essential observation (parallel to the studies of individual high-redshift, infrared, luminous galaxies), to understand the underlying scenario and physics of galaxy formation and evolution. A first study has been done using the 850 $\\mu$m galaxies \\citep{2004ApJ...611..725B}. Although the number of sources is quite small, they find evidence that submillimiter galaxies are linked to the formation of massive galaxies in dense environments destined to become rich clusters. This has now been directly supported by the detection of the clustering of high-redshift 24 $\\mu$m selected ULIRGs and HyperLIRGs \\citep{2006ApJ...641L..17F,2007MNRAS.375.1121M}. Studying correlations with individual IR galaxies is very hard due to either high confusion noises, instrumental noises, or small fields of observation. It has been shown that the IR-background anisotropies could provide information on the correlation between the sources of the CIB and dark matter for large-scale structures \\citep [] [hereafter HK] {2001ApJ...550....7K,2000ApJ...530..124H} and on the large-scale structure evolution. First studies at long wavelengths have only detected the shot-noise component of the fluctuations: \\citet{2000A&A...355...17L} at 170 $\\mu m$, \\citet{2000A&A...361..407M} at 90 and 170 $\\mu m$, \\citet{2001phso.conf..471M} at 60 and 100 $\\mu m$.  \\citet{2007ApJ...665L..89L} and \\citet{2007A&A..4512G} report first detections of the correlated component using Spitzer/MIPS data at 160 $\\mu m$. \\citet{2007ApJ...665L..89L} measured a linear bias $b \\sim 1.7$. \\\\   Future observations by Herschel and Planck will allow us to probe the clustering of IR and submm galaxies. Nevertheless these experiments will be limited, the confusion and instrumental noises will hinder detections of faint individual galaxies. Clustering thus has to be analysed in the background fluctuations \\citep [e.g.] [] {2007astro.ph..3210N}. The need for a prior understanding of what could be done by these experiments has motivated us to develop a set of realistic simulations of the IR and sub-mm sky. \\\\  In Sect. 2 we present the model on which are based our simulations. In Sect. 3 we discuss how the simulations are done and present a set of simulated sky maps and their corresponding catalogs. Different catalogs are created for 3 different levels of correlation between the IR galaxy emissivity and the dark-matter fluctuation density field (strong, medium, and no correlation). For each of these catalogs, we can create maps of the sky at any given IR wavelength and simulate how different instruments will see them. We focus in this paper on Planck/HFI and Herschel/SPIRE. In Sect. 4 we use the simulated maps to give predictions for the confusion noise, the completeness, and the detection limits for each of the study cases, including  the instrumental noise. In Sect. 5 we present the power spectra of the CIB anisotropies for Planck/HFI and discuss their detectability against the significant sources of contamination (shot noise, cirrus, and cosmic microwave background (CMB)). \\\\  Throughout the paper the cosmological parameters were set to $h=0.71,\\Omega_{\\Lambda}=0.7,\\Omega_{m}=0.27$. For the dark-matter linear clustering we set the normalization to $\\sigma_{8}=0.8$. ",
        "conclusions": "This paper presented new simulations of the cosmic infrared and submillimeter background. The simulations are based on an empirical model of IR galaxy evolution (the LDP model) combined with a simple description of the correlation. The IR galaxy spatial distribution follows that of the dark-matter halo density field, with a bias parameter accounting for the possibility of the luminous matter being more or less correlated with the dark matter. The simulated maps and their catalogs are publicly available at http://www.ias.u-psud.fr/irgalaxies/simulations.php. Other maps are available upon request. These maps are intended to be a useful tool for planning future large IR and submillimeter surveys. In this paper, we used the maps to predict the confusion noise and completeness levels for Planck/HFI and Herschel/SPIRE. We also predicted the power spectra of correlated CFIRB anisotropies for Planck/HFI which will be a major advance in the study of the CFIRB anisotropies at large scales (i.e. $\\ell<2000-5000$ depending on the wavelengths). Further analysis of the CFIRB anisotropies, including the use of stacking analysis to isolate the anisotropies in different redshift ranges, will be presented in a second paper, now in preparation."
    },
    "0801/0801.2467.txt": {
        "abstract": "We present results on integral-field optical spectroscopy of five luminous Blue Compact Dwarf galaxies. The data were obtained using the fiber system \\textsc{integral} attached at the William Herschel telescope. The galaxies Mrk~370, Mrk~35, Mrk~297, Mrk~314 and III~Zw~102 were observed. The central $33\\farcs6\\times29\\farcs4$ regions of the galaxies were mapped with a spatial resolution of $2\\farcs7$~spaxel$^{-1}$, except for Mrk~314, in which we observed the central $16\\arcsec\\times12\\arcsec$ region with a resolution of $0\\farcs9$~spaxel$^{-1}$. We use high-resolution optical images to isolate the star-forming knots in the objects; line ratios, electron densities and oxygen abundances in each of these regions are computed. We build continuum and emission-line intensity maps as well as maps of the most relevant line ratios: [\\ion{O}{3}]~$\\lambda5007$/\\Hb, [\\ion{N}{2}]~$\\lambda6584$/\\Ha, and \\Ha/\\Hb, which  allow us to obtain spatial information on the ionization structure and mechanisms. We also derive the gas velocity field from the \\Ha\\ and [\\ion{O}{3}]~$\\lambda5007$ emission lines.  We find that all the five galaxies are in the high end of the metallicity range of Blue Compact Dwarf galaxies, with oxygen abundances varying from $Z_{\\sun}\\simeq0.3$ to $Z_{\\sun}\\simeq1.5$. The objects show \\ion{H}{2}-like ionization in the whole field of view, except the outer regions of III~Zw~102 whose large [\\ion{N}{2}]~$\\lambda6584$/\\Ha\\ values suggest the presence of shocks. The five galaxies display inhomogeneous extinction patterns, and three of them have high \\Ha/\\Hb\\ ratios, indicative of a large dust content;  all galaxies display complex, irregular velocity fields in their inner regions. ",
        "introduction": "Blue Compact Dwarf (BCD) galaxies are low-luminosity and compact objects, which show optical spectra similar to those presented by \\ion{H}{2} regions in spiral galaxies \\citep{ThuanMartin81}. They are also low-metallicity systems \\citep{Searle72,Lequeux79,KunthSargent83,Masegosa94} with intense star formation activity (rates ranging between 0.1 and $1\\;M_{\\sun}$ yr$^{-1}$; \\citealt{Fanelli88}). In most BCDs the gas consumption time scale is much shorter than the age of the Universe, which, together with their low metal content, raised initially the hypothesis that they could be truly young galaxies, forming their first generation of stars \\citep{Searle72,Kunth88}. Nowadays it is clear that the great majority of them are not young, but rather old systems (\\citealt{Papaderos96a}; \\citealt[][=C01a]{Cairos01a}; \\citealt[][=C01b]{Cairos01b}; \\citealt{BergvallOstlin02}; \\citealt[][=C03]{Cairos03}; \\citealt{Gildepaz03}) undergoing short starbursts followed by longer quiescent periods.  An essential requisite to comprehend the formation and evolution of BCDs, and to elaborate coherent pictures of dwarf galaxy evolution, is to understand the process of star-formation in these galaxies ---~how their starburst activity ignites and how it propagates afterward~---, and to derive their star-forming (SF) histories. Even though both issues have been the subject of a considerable observational and interpretative effort during the last years, they are still far from being well understood.  Different scenarios have been invoked to explain the onset of star formation in these galaxies. Some of them favored internal processes, such as the Stochastic Self-Propagating Star Formation proposed by \\cite{Gerola80}, or the hypothesis of cyclic gas re-processing, that is, the cyclic expulsion and later accretion of the Interstellar Medium \\citep{Davies88}. Alternatively, interactions and/or mergers have been proposed as the mechanism responsible for triggering the star formation in dwarfs. From optical surveys \\citep{Campos93,TellesTerlevich95,TellesMaddox00} it was first concluded that gravitational interactions with optically bright galaxies are too rare to play a relevant role.  However, studies at radio wavelengths \\citep{Taylor93,Taylor95,Taylor96} established that \\ion{H}{2}/BCDs do have companions, most of them gas-rich, faint objects, which may go undetected in the optical. In agreement with these findings, more recent optical studies that extended the searches at lower masses and luminosities also found that a substantial fraction of SF dwarfs possess low mass companion galaxies \\citep{Noeske01,Pustilnik01}. \\cite{Ostlin01} and \\cite{BergvallOstlin02} put forward the idea that a merger between two galaxies with different metallicities, one gas-rich and one gas-poor, or infall of intergalactic clouds, could be the most plausible explanation for the starburst activity, at least in the most luminous BCDs. Recent studies, focused on individual objects, have shown that interactions do play a decisive role in the evolution of these systems \\citep{Johnson04,BravoAlfaro04,BravoAlfaro06}.  Much work must still be done to elucidate the star formation histories of BCDs. The first step is to derive the properties of their stellar populations, quite a difficult task in such complex objects. The great majority of them cannot be resolved into stars, and the only way to study their stellar populations is by comparing their integrated properties with the predictions of evolutionary synthesis models. At any location in the galaxy, the emitted flux is the sum of the emission from the local starburst, the flux produced by the nebula surrounding the young stars, and the emission from the underlying, old stellar population, all possibly modulated by dust (\\citealt{Cairos00}; \\citealt[][=C02]{Cairos02}; \\citealt[][=C07]{Cairos07}).  Another issue that was recently brought into discussion is the possible heterogeneity of the BCD class. Among the galaxies classified as BCDs we find a wide variety of objects spanning a large interval in luminosities ($-21.5 < M_{B} < -14.0$) and morphologies. \\cite{LooseThuan86} developed a classification scheme for BCDs based on the appearance of the starburst and the shape of the external envelopes, distinguishing four subclasses: \\textit{iE} galaxies, which show a complex inner structure with several SF regions over-imposed on an outer regular envelope; \\textit{nE}, objects with a clearly defined nucleus and regular isophotes; \\textit{iI}, which have irregular outer and inner parts and, finally, \\textit{iO} galaxies, in which an outer envelope is not found. C01b have grouped the BCDs in four categories according to the position and morphology of the SF knots: {\\em nuclear starburst}, which are galaxies with a prominent central starburst;  {\\em extended starburst}, galaxies with star formation spread over the entire galaxy, {\\em chain starburst}, objects in which the SF knots are aligned to form a chain and finally, {\\em cometary starburst}, galaxies with a \"cometary\" appearance, that is the star formation concentrated in one side of the galaxy.  This finding introduces an additional complication, because it opens the possibility that different star formation mechanisms may operate in BCDs and that galaxies classified as BCDs may actually have different star formation histories and evolutionary paths.  Spectrophotometric studies that put together high quality data in a large wavelength range are fundamental to get insights into the star-formation process and history of BCDs. However, very few such studies have been carried out so far, and all of them focused on individual galaxies (\\citealt{Papaderos99,GildePaz00}; C02; \\citealt{Guseva03a,Guseva03b,Guseva03c}; C07). These studies rely on conventional observing techniques: they combine good resolution broad-band frames in different bands (to spatially isolate the different stellar populations and map the dust distribution), with narrow-band imaging (needed to map the gas distribution, isolate the  starburst regions and to derive their physical properties) and spectroscopy data (to derive the internal extinction, compute physical parameters and chemical abundances of the gas, and remove the contribution from emission lines); a sequence of long-slit spectra sweeping the regions of interest is usually taken. Therefore, these analyses require a great amount of observing time; besides, varying instrumental and  atmospheric conditions make combining all the data together quite complicated. Long-slit spectroscopy has the additional problem of the uncertainty on the exact location of the slit.  We have thus undertaken a project to carry out an Integral Field Spectroscopy (IFS) mapping of a large sample of BCD galaxies. IFS is the ideal observational technique to study BCDs: each single exposure contains both spatial and spectral information in a large area of the galaxy, so just in one shot we collect information for all the SF regions as well as for the low surface brightness (LSB) stellar component. Besides, the kinematical information also allows us to investigate what mechanisms ignite the star-formation in BCDs. In terms of observing time, IFS observations of BCDs are an order of magnitude more efficient than traditional observing techniques, and provide simultaneous data for all spatial resolution elements under the same instrumental and atmospheric conditions, which guarantees a greater homogeneity of the dataset.   Here we present the first results from this project: we study five galaxies representative of the group of BCDs populating the high luminosity and metallicity range, and having rather complex morphologies. These objects, often referred to as Luminous Blue Compact galaxies (LBCGs, see \\citealt{Ostlin98,Kunth00}; C01b) have a special significance in the general scenario of galaxy formation and evolution, as they are claimed to be the local counterparts of the luminous, compact, SF galaxies detected at higher redshifts \\citep{Garland04}. The proximity of these systems allows us to study their structure, stellar populations, star formation processes and chemical abundances with an accuracy and spatial resolution that cannot be achieved at intermediate/high redshift. ",
        "conclusions": "We present results from what is, as far as we know, the first Integral Field Spectroscopy analysis of a sample of galaxies catalogued as Blue Compact Dwarfs. % % Esto lo quite: \"From this study we highlight the following results:\" %  \\begin{itemize}  \\item With the help of high resolution optical images, we define several regions of interest in each object, and for each of them we measure emission line fluxes and compute the most relevant line ratios. These data show that:  \\begin{itemize}  \\item The strength of emission lines and absorption features, the shape of the continuum, and the presence of other relevant lines significantly vary across each galaxy. This indicates varying physical conditions and/or stellar content.  \\item All the regions studied have \\ion{H}{2}-like ionization (i.e. star-formation).  \\item All the identified regions in the observed galaxies have low electron densities, ranging from $\\leq 100$ to 360 cm$^{-3}$, typical of classical \\ion{H}{2} regions.  The electron density also shows significant spatial variations in four out of the five galaxies studied.  \\item The derived oxygen abundances are relatively large in all the cases, ranging from $Z_{\\sun}\\simeq0.3$ to $Z_{\\sun}\\simeq1.5$. No significant oxygen abundances variations within a same object are found.  \\end{itemize}   \\item We build maps of the stellar continuum, emission line fluxes and excitation ratios, finding that:  \\begin{itemize}  \\item Continuum and emission line morphologies are generally different.  \\item The excitation ratios [\\ion{O}{3}]~$\\lambda5007$/\\Hb\\ and [\\ion{N}{2}]~$\\lambda6584$/\\Ha\\ are typical of HII regions in the whole observed fields. Only in the outer regions of III~Zw~102 the [\\ion{N}{2}]~$\\lambda6584$/\\Ha\\ ratio may suggest the presence of shocks.  \\item All the galaxies present a complex extinction pattern. Assuming that the extinction coefficient is constant across the whole galaxy can lead to considerable errors in the derivation of magnitudes and colors and in the determination of the star formation rate and ages.  \\item The \\Ha/\\Hb\\ ratios are higher that the theoretical values, indicating the presence of significant amounts of dust in these galaxies.  \\end{itemize}  \\item In all the five galaxies the central regions display complex, distorted ionized gas velocity fields, though large scale ordered motions are present in three of them. With the current data we cannot determine whether these perturbed velocity fields are the signature of interaction/mergers episodes or they are the result of shocks, stellar winds or supernovas.  Further IFS observations are essential to investigate this issue: higher resolution spectra, in order to measure the velocity dispersion of the gas, as well as deeeper observations which allow to map the kinematics of the stars, are indeed fundamental.  \\item The galaxies studied here, although classified as BCDs, all show characteristics very different from those of \"genuine\" BCD galaxies: they have relatively high metallicities and significant amounts of dust; they show also variations of such properties as electron densities, extinction and ionization degrees, across the galaxies. These properties make them an especially attractive area, as they could be the local counterparts of the blue, high metallicity, vigorously starbursting galaxies, detected at intermediate redshift.  \\end{itemize}"
    },
    "0801/0801.3646_arXiv.txt": {
        "abstract": "We perform numerical simulations of solid particle motion in a shearing box model of a protoplanetary disc. The accretion flow is turbulent due to the action of the magnetorotational instability. Aerodynamic drag on the particles is modelled using the Epstein law with the gas velocity interpolated to the particle position. The effect of the magnetohydrodynamic turbulence on particle velocity dispersions is quantified for solids of different stopping times $t_{\\rm{s}}$, or equivalently, different sizes. The anisotropy of the turbulence is reflected upon the dispersions of the particle velocity components, with the radial component larger than both the azimuthal and vertical components for particles larger than $\\sim $ 10 cm (assuming minimum-mass solar nebula conditions at 5 AU). The dispersion of the particle velocity magnitude, as well as that of the radial and azimuthal components, as functions of stopping time, agree with previous analytical results for isotropic turbulence. The relative speed between pairs of particles with the same value of $t_{\\rm{s}}$ decays faster with decreasing separation than in the case of solids with different stopping time. Correlations in the particle number density introduce a non-uniform spatial distribution of solids in the 10 to 100 cm size range. Any clump of particles is disrupted by the turbulence in less than one tenth of an orbital period, and the maximally concentrated clumps are stable against self-gravitational collapse. ",
        "introduction": "In order for planetesimals to form in a circumstellar disc, the parent dust grains must meet a variety of kinematic conditions if they are to grow in mass and size. The formation of kilometre-sized solid bodies by gravitational instability requires, through the Toomre criterion, that the dust spatial density be $\\rho_{d}>\\Omega^{2}/\\pi G$, where $\\Omega$ is the disc angular velocity and $G$ the gravitational constant. Such densities could be achieved in a mid-plane dust layer after sedimentation has occurred. On the other hand, growth by collisional sticking may be accomplished if the grain relative speeds acquire certain values: experiments have shown that, for micron-sized grains, collision speeds less than $\\sim$0.2 m/s lead to aggregates with a fractal geometry (Blum \\& Wurm 2000), and above $\\sim$ 1 m/s disruption of the colliding agglomerates takes place.  While binary collisions of the smallest dust grains are caused by Brownian motion, other sources of relative velocities in a protoplanetary nebula may come into play as the interaction between solids and gas decreases. Differential sedimentation, radial drift and turbulence can provide the necessary kinetic energies to drive the collisional motion of solids, in addition to changing their spatial distribution throughout the nebular disc. In this paper we focus on the effect of a turbulent gas flow on solid particle velocities and spatial arrangement.  The formalism of V\\\"{o}lk et al. (1980) makes it possible to calculate the velocity dispersions of particles in isotropic turbulence. The effect of the fluctuating gas velocity on individual solids is approximated by a sum over two types of turbulent eddies: those with spatial frequencies $k$ and turnover times $t_{k}>t_{\\rm{s}}$, where $t_{\\rm{s}}$ is the particle stopping time, and which merely advect the particle; and those with $t_{k}<t_{\\rm{s}}$, which decay before the particle has had time to cross them. The result is a velocity dispersion that decreases with particle stopping time as $(\\Omega t_{\\rm{s}})^{-1/2}$, with a sharp steepening for particles with $\\Omega t_{\\rm{s}}\\sim$ a few times $10^{-1}$. Under minimum mass solar nebula conditions at 5 AU, this corresponds to particle radii of tens of centimetres. The method is nicely summarised and applied by Cuzzi \\& Hogan (2003) to chondrule-sized objects.  The accumulation of a large number of solids in small regions of the turbulent protoplanetary disc can lead to feedback effects on the gas flow if the solid mass density is greater than the local gas density. Such particle agglomerations can also lead to dust growth if the associated collision rates are conducive to sticking. One recently suggested clumping mechanism is a streaming instability that arises due to the relative motion between solids and gas in a Keplerian disk (Youdin \\& Goodman 2005). The feedback of solids onto the gas can generate exponential growth of particle density perturbations as well as turbulence, and particle clumps can develop overdensities up to a factor of $\\sim$ 50 relative to the mean gas density (Youdin \\& Johansen 2007, Johansen \\& Youdin 2007).  The process of turbulent concentration (Maxey 1987) has been explored by Cuzzi et al. (2001), who use the ratio of local particle density to its global average as the measure of a multifractal distribution of concentrated chondrules. This distribution resembles that of the dissipation of the turbulent kinetic energy on the Kolmogorov scale (Chhabra 1989). The connection between turbulent concentration and turbulent dissipation suggests that the local structure of the chondrule concentration field may be independent of flow properties. For larger particles, an accurate characterisation of their number statistics in the turbulent circumstellar environment is still lacking.   The recognition that turbulence due to the magnetorotational instability (MRI; Balbus \\& Hawley 1991; Hawley, Gammie \\& Balbus 1995, henceforth HGB) can provide the viscous stresses necessary to allow inward accretion in discs, has prompted investigations of different dynamical aspects of dust evolution in the presence of magnetohydrodynamic (MHD) turbulence. Planetesimal formation by gravitational instability was addressed by Johansen et al. (2006, henceforth JKH), whose numerical studies concluded that groups of at least 1000 particles in their simulations could become gravitationally bound, assuming specific values of disc column densities (150 and 900 g cm$^{-2}$) and a ratio of disc scale height to radius of 0.04. Fromang \\& Nelson (2005) examine the accumulation of individual solids due to vortical structures in a global protoplanetary disc model. They follow the radial migration of up to 3000 bodies, whose drag interaction with the gas corresponds to objects of approximately 1 m in size. Those particles able to migrate inwards do so at different rates, with reductions in semi-major axes ranging from a factor of 1.6 in approximately 50 orbits to a factor of 2.2 in 200 orbits. The particles that become trapped in vortices maintain an approximate constant radial position until the end of the simulation.  In the following we present a study of the kinematics of solid particles in a local MHD model of a protoplanetary disc. In Section 2 we describe the numerical method employed to model the accretion flow and the particles. Section 3 contains results obtained from measurements of particle velocities and spatial distribution. In Section 4 we discuss their significance in the context of particle growth in a protoplanetary nebula, and concluding remarks are presented in Section 5. ",
        "conclusions": "We have performed numerical MHD calculations of solid particle velocities and spatial distribution in a turbulent protoplanetary disc, in the context of the shearing box model. The results suggest that, even though the turbulent flow has an anisotropic spectrum, the dispersion of the particle velocity magnitude (normalised by the corresponding dispersion for the gas) as a function of particle stopping time is consistent with analytical models that assume isotropy. Nevertheless, the dispersions of the individual velocity components for each stopping time are different. In the case of particles with $\\Omega t_{\\rm{s}} \\gtrsim 1$, the radial component of the solid-to-gas velocity dispersion ratio is larger than the azimuthal and vertical components typically by factors of $\\sim 2.3 - \\sim$ 2.5. The relative speed between pairs of particles with the same stopping time decreases towards small values at small separations, whereas for pairs of different stopping times the relative speed reaches a finite value that depends on the sizes of the pair members.  The clustering induced on the particle spatial distribution by the turbulent flow gives rise to short-lived clumps. Under conditions of a minimum-mass solar nebula, these clumps are at least nine times larger than their respective Jeans radii, as a result of high particle velocity dispersions. Gravitational collapse would therefore seem to be an unlikely mechanism in the formation of more massive objects, under the conditions that we assumed.  Future calculations should increase the numerical resolution used in the measurement of the relative velocities between particles, in order to improve fittings to existing analytical models.   \\bigskip"
    },
    "0801/0801.1565_arXiv.txt": {
        "abstract": "Dark energy and dark matter are the dominant sources in the evolution of the late universe. They are currently only indirectly detected via their gravitational effects, and there could be a coupling between them without violating observational constraints. We investigate the background dynamics when dark energy is modelled as exponential quintessence, and is coupled to dark matter via simple models of energy exchange. We introduce a new form of dark sector coupling, which leads to a more complicated dynamical phase space and has a better physical motivation than previous mathematically similar couplings. ",
        "introduction": "Observations are providing increasingly compelling evidence that the expansion of the Universe is accelerating, driven by ``dark energy\" (see e.g.~\\cite{Spergel:2006hy} for recent results). The simplest model of dark energy is a cosmological constant $\\Lambda$, representing the vacuum energy density, and this model provides a very good fit to a range of independent observations. However, there is no satisfactory theoretical explanation for the very small value of $\\Lambda$. Furthermore, the $\\Lambda$ model suffers from a fine-tuning, or ``coincidence\", problem -- why is the dark matter density comparable to the vacuum energy density now, given that their time evolution is so different?  If the dark energy evolves with time, this may alleviate the coincidence problem. The simplest models of evolving dark energy are light scalar fields, known as ``quintessence\". If the quintessence is coupled to the dark matter, then this may be able to account for the similar energy densities in the dark sector today. A decisive way of achieving similar energy densities is if the coupling leads to an accelerated scaling attractor solution, with \\begin{equation} \\label{scal} {\\Omega_{\\text{dark energy}}\\over \\Omega_{\\text{dark matter}}}=O(1) ~~\\mbox{ and }~~ \\ddot a >0\\,. \\end{equation} In this case, the coincidence problem is reduced to a simple choice of parameters to match $ \\Omega_{\\text{dark energy}}/ \\Omega_{\\text{dark matter}}$ to observations. Since the accelerated scaling solution is an attractor, no fine-tuning of initial conditions is needed. However, the dynamics that produces such scaling in the dark sector may have other undesirable consequences.  Here we study quintessence with an exponential potential, \\begin{equation} V(\\varphi)= V_0 \\exp \\left(-\\kappa \\lambda  \\varphi\\right)\\,,~~ \\kappa^2:=8\\pi G\\,, \\end{equation} where $\\lambda$ is dimensionless and $V_0>0$. The dynamics of a universe with exponential quintessence and an uncoupled perfect fluid have been studied, and these models do not admit late-time accelerated scaling attractors that satisfy Eq.~(\\ref{scal})~\\cite{quin}. When we introduce a coupling between the quintessence and dark matter, accelerated scaling attractors are possible~\\cite{Wetterich:1994bg,Amendola:1999qq}. However, in some models, this is achieved at the expense of introducing other problems which can rule out the model~\\cite{Amendola:2006qi}.  A general coupling between a quintessence field $\\varphi$ and dark matter (with density $\\rho_c$) may be described in the background by the balance equations, \\begin{eqnarray} \\label{relation} \\dot \\rho_c  &=& - 3H\\rho_c-Q\\,,\\label{cc}\\\\ \\dot \\rho_{\\varphi} &=& - 3H(1+w_{\\varphi})\\rho_{\\varphi}+ Q\\,. \\label{kg1} \\end{eqnarray} Here $Q$ is the rate of energy density exchange in the dark sector, and \\be \\label{sq} Q~\\left\\{ \\begin{array}{l} >0\\\\ <0 \\end{array} \\right. ~~ \\Rightarrow ~~ \\mbox{energy transfer}~\\left\\{ \\begin{array}{l} \\mbox{dark matter $\\to$ dark energy}\\\\ \\mbox{dark energy $\\to$ dark matter}\\end{array} \\right. \\ee The dark energy equation of state parameter is \\be w_\\varphi:={ p_\\varphi \\over \\rho_\\varphi}={{1\\over 2}\\dot\\varphi^2-V(\\varphi) \\over {1\\over 2}\\dot\\varphi^2+V(\\varphi)}\\,.\\label{wp} \\ee The modified Klein-Gordon equation follows from Eq.~(\\ref{kg1}): \\begin{equation} \\ddot \\varphi +3H \\dot \\varphi + \\frac{dV}{d\\varphi}= {Q \\over \\dot\\varphi}\\,. \\label{kg} \\end{equation} When we include baryons ($ \\rho_b$) and radiation ($\\rho_r $), the remaining evolution equations are \\begin{eqnarray} \\dot \\rho_b &=& - 3H\\rho_b \\,,\\\\ \\dot {\\rho}_r &=& - 4H\\rho_r \\,,\\\\ \\dot H  &=&  -\\frac{\\kappa^2}{2}\\left[\\rho_c+\\rho_b+ \\frac{4}{3}\\rho_r+\\dot\\varphi^2 \\right], \\label{ray} \\end{eqnarray} subject to the Friedman constraint, \\begin{eqnarray} \\label{fried} \\Omega_c+\\Omega_b+\\Omega_r +\\Omega_{\\varphi}=1\\,,~~ \\Omega := \\frac{\\kappa^2 \\rho}{3H^2}. \\end{eqnarray}  We can define effective equation of state parameters for the dark sector, which describe the equivalent uncoupled model in the background: $\\dot \\rho_c +3H(1+ \\wc)\\rho_c=0 $, $\\dot \\rho_\\varphi +3H(1+ \\wv)\\rho_\\varphi=0 $. By Eqs.~(\\ref{cc}) and (\\ref{kg1}), \\be \\label{weff} \\wc={Q \\over 3H\\rho_c}\\,,~~ \\wv= w_\\varphi -{Q \\over 3H \\rho_\\varphi}\\,. \\ee It follows that \\bea && Q >0 ~~ \\Rightarrow ~~ \\left\\{ \\begin{array}{ll} \\wc>0 & ~~\\mbox{dark matter redshifts faster than}~a^{-3}\\\\ \\wv<w_\\varphi & ~~\\mbox{dark energy has more accelerating power}\\end{array} \\right. \\\\ && Q <0 ~~ \\Rightarrow ~~ \\left\\{ \\begin{array}{ll}\\wc<0 & ~~\\mbox{dark matter redshifts slower than}~a^{-3}\\\\ \\wv>w_\\varphi & ~~\\mbox{dark energy has less accelerating power} \\end{array} \\right. \\eea When $Q>0$ it is possible that $\\wv<-1$ (see~\\cite{Huey:2004qv} for specific examples). This means that the coupled quintessence behaves like a ``phantom\" uncoupled model -- but without any negative kinetic energies.  Equations~(\\ref{cc})--(\\ref{ray}) are an autonomous system of the form \\begin{eqnarray} \\dot {\\mathbf{x}} = \\mathbf{f}(\\mathbf{x})\\,, \\label{math1} \\end{eqnarray} and the critical points satisfy $\\mathbf{f}(\\mathbf{x}_*)=0$. In order to study the stability of the critical points, we expand about them, $\\mathbf{x} = \\mathbf{x}_* +\\mathbf{u}$, and Eq.~(\\ref{math1}) yields \\begin{equation} \\dot {\\mathbf{u}} = {\\mathbf{f}'}(\\mathbf{x}_*) \\mathbf{u} + \\mathbf{g}(\\mathbf{x})\\,. \\end{equation} Here $\\mathbf{g}(\\mathbf{x})/||\\mathbf{x}|| \\rightarrow 0$ as $\\mathbf{x} \\rightarrow \\mathbf{x}_*$, and \\begin{equation}\\label{eigv} f'_{ij}(\\mathbf{x}_*) = \\frac{\\partial f_{i}}{\\partial x_{j}}(\\mathbf{x}_*)\\,, \\end{equation} is a constant non-singular matrix, whose eigenvalues encode the behaviour of the dynamical system near the critical point.  If a component of $\\mathbf{f}$ can be written as a fraction $u(\\mathbf{x})/v(\\mathbf{x})$, then a critical point requires the vanishing of the numerator, $u(\\mathbf{x}_*)=0$. If the denominator also vanishes at the critical point, $v(\\mathbf{x}_*)=0$, then care is needed in obtaining the eigenvalues of the linearized system~(\\ref{math1}). Strictly speaking, the fraction $u(\\mathbf{x}_*)/v(\\mathbf{x}_*)$ may not be well defined. However, it is still possible to obtain analytical results via analysis of the behavior of the eigenvalues of $\\mathbf{f}'$ in the limit $v(\\mathbf{x}) \\rightarrow 0$. ",
        "conclusions": "\\label{sec:concl}  We considered the background dynamics of a universe dominated by dark energy (in the form of exponential quintessence) and cold dark matter, where there is energy exchange in the dark sector, as in Eqs.~(\\ref{cc}) and (\\ref{kg1}), \\[ \\dot \\rho_c + 3H\\rho_c=-Q=-\\left[ \\dot \\rho_{\\varphi} + 3H(1+w_{\\varphi})\\rho_{\\varphi}\\right]. \\] For the previously introduced forms of $Q$, given in Eqs.~(\\ref{I}) and (\\ref{II}), \\[ \\mbox{(I):  }~~ Q=\\sqrt{ 2 /3}\\, \\kappa\\,\\beta \\rho_c\\dot\\varphi\\,,~~~ \\mbox{(II):  }~~  Q= \\alpha H \\rho_c\\,, \\] the phase space remains two-dimensional, as in the uncoupled case $Q=0$. We found the properties of the critical points for all signs of $\\lambda, \\alpha, \\beta$. The results are summarized in Tables~\\ref{crit} and \\ref{crit2}, and slightly extend previous work~\\cite{Amendola:1999qq,Holden:1999hm,Billyard:2000bh} in the case of pressure-free matter ($w_c=0$). In both models, critical point D is the cosmologically relevant point, because for nonzero coupling it includes accelerated scaling attractor solutions, \\[ 0<\\Omega_{c *}\\,,\\Omega_{\\varphi *}<1 ~\\mbox{  and  }~ w_{\\text{tot} *} <-{1\\over3}\\,. \\] The stability behaviour of critical point D was investigated numerically, and is shown in Fig.~\\ref{DI} for model (I) and Fig.~\\ref{DII} for model (II).  Our main results are for a new coupling model~\\cite{val}, defined in Eq.~(\\ref{III}), \\[ \\mbox{(III):  }~~Q=\\Gamma \\rho_c\\,. \\] This has a similar form to model (II), but is more physical since the transfer rate $\\Gamma$ is determined only by local properties of the dark sector interaction at each event, and is not dependent on the universal expansion rate. When $\\Gamma>0$, this new model has the same form as simple models for the decay of dark matter particles to radiation~\\cite{Cen:2000xv}, or to quintessence~\\cite{Ziaeepour:2003qs}, and for the decay of the curvaton field into radiation~\\cite{Malik:2002jb}.  Model (III) requires a three-dimensional phase space, since the Hubble rate cannot be eliminated from the equations for $x',y'$. This makes the dynamics more complicated than for models (I) and (II). In particular, a new set of late-time critical points arises in (III), and considerable analytical effort is required to identify these points and determine their stability properties. Our results are summarized in Tables~\\ref{crit3} and \\ref{eigen1}. We performed numerical integrations of the dynamical system in order to confirm the analytical results, and examples of these integrations are shown in Figs.~\\ref{lambda_one} and \\ref{lambda_four}.  The cosmologically relevant critical point is G, which allows for an accelerated critical solution (when $\\lambda^2<2$). However, this is not a scaling solution, since \\[ \\Omega_{c *}=0\\,,~~\\Omega_{\\varphi * }=1\\,, \\] which is similar to the asymptotic behaviour of the standard $\\Lambda$CDM model. This accelerated critical solution is an attractor when $\\gamma>0$, i.e., for the case when dark matter is decaying to dark energy. Note that in model (II), the decaying dark matter case, $\\alpha>0$, does not lead to any accelerated attractor (see Table~II). Model (III) with $\\Gamma>0$ produces an interesting class of models where dark matter decays to dark energy -- so that the primordial universe may have no dark energy -- and where this decay eventually leads to dark energy dominance, independent of initial conditions (since there is an attractor). Although such models do not solve the coincidence problem in the standard way, they may provide a new approach to the broader problem of explaining why dark energy dominates over dark matter only late in the universe's evolution.  The background dynamics for coupling model (III) show new features not present in the previously investigated models (I) and (II). In order to confront this model with observations, the cosmological perturbations with a dark sector coupling of form (III) need to be investigated (see Ref.~\\cite{val}).   \\[ \\] {\\bf Acknowledgements:}\\\\ We thank  Luis Ure\\~na-L\\'opez, Elisabetta Majerotto,  Luca Parisi, Israel Quir\\'os,  Jussi V\\\"aliviita and Shinji Tsujikawa for useful discussions. GCC is  supported by the Programme Alban, the European Union Programme of High Level Scholarships for Latin America, scholarship no.~E06D103604MX and the Mexican National Council for Science and Technology, CONACYT, scholarship no.~192680. RL is supported by the University of the Basque Country through research grant GIU06/37, and by the Spanish Ministry of Education and Culture through research grants FIS2004-01626 and FIS2004-0374-E. The work of RM is supported by STFC."
    },
    "0801/0801.1279_arXiv.txt": {
        "abstract": "We compare the X-ray spectra and luminosities, in the 2-8 keV band, of known and suspected cataclysmic variables (CVs) in different environments, assessing the nature of these source populations.  These objects include nearby CVs observed with ASCA; the Galactic Center X-ray source population identified by Muno et al.; and likely CVs identified in globular clusters.  Both of the latter have been suggested to be dominated by magnetic CVs.  We find that the brighter objects in both categories are likely to be magnetic CVs, but that the fainter objects are likely to include a substantial contribution from normal CVs.  The strangely hard spectra observed from the Galactic Center sources reflect the high and variable extinction, which is significantly greater than the canonical $6\\times10^{22}$ cm$^{-2}$ over much of the region, and the magnetic nature of many of the brightest CVs.  The total numbers of faint Galactic Center sources are compatible with expectations of the numbers of CVs in this field. ",
        "introduction": "The unprecedented spatial resolution of the {\\it Chandra X-ray Observatory} allows us to study populations of faint X-ray sources at distances of kiloparsecs. Large numbers of X-ray sources of moderate luminosities ($10^{31}<L_X<10^{34}$ ergs/s) have been discovered in several Galactic environments (e.g. star-forming regions, globular clusters, the Galactic Center) in recent years. For instance, \\citet{Grindlay01a,Pooley02a} and others have found large numbers of X-ray sources in globular clusters, many of which are optically identified as CVs. \\citet{Muno03} identified a large population of $\\sim$2000 faint, spectrally hard X-ray sources associated with the central 40 pc around Sgr A*.  Interpretation of these observations in terms of known categories of X-ray sources is sometimes hampered by the different observing bands and models used for studies of nearby, well-known sources versus these distant populations. The Galactic Center sources have been inferred \\citep{Muno03,Muno04a}, by their X-ray luminosity and spectral hardness, to be composed largely of the class of CVs known as intermediate polars \\citep{Patterson94}. Intermediate polars have white dwarf accretors with relatively high magnetic fields, which force the accreting material to follow the magnetic field lines of the accretor down onto the magnetic poles. The brightest CVs in globular clusters have been argued to be too X-ray luminous for normal CVs \\citep{Verbunt97}, and several pieces of evidence point to a magnetic nature of at least some CVs in globular clusters, including He II $\\lambda$4686 emission observed in NGC 6397 CVs \\citep{Grindlay95}, X-ray periodicities in some of the brightest CVs in 47 Tuc \\citep{Grindlay01a}, enhanced $N_H$ columns towards the brightest CVs \\citep{Heinke05a}, and the lack of dwarf nova outbursts from globular clusters \\citep{Shara96,Edmonds03b,Dobrotka06}.  However, the question of whether magnetic CVs are more common than normal CVs in globular clusters or the Galactic Center has not yet been explored. We have directly compared, for the first time, the spectra of Galactic Center X-ray sources, and likely globular cluster CVs (those with optical counterparts), with archival ASCA X-ray spectra of well-known nearby CVs.  Full details will be published in Heinke et al. (2008, in prep). ",
        "conclusions": "Studies of faint ($10^{31}<L_X<10^{34}$ ergs/s) hard X-ray sources in globular clusters and the Galactic Center have identified them as primarily CVs, on both observational and theoretical (lack of a sufficiently numerous alternative population) grounds.  They have also suggested that many or most of them are magnetic systems, particularly intermediate polars.  We have compared the spectra of nearby CVs, observed with ASCA, to the low-count spectra or colors of globular cluster CVs and Galactic Center X-ray sources.  We find that significant (although poorly constrained) fractions of the observed cluster and galactic center populations are consistent in their X-ray spectra and fluxes with nonmagnetic CVs.  For the Galactic Center, we require a somewhat higher average extinction than typically assumed ($N_H=10^{23}$ instead of $6\\times10^{22}$ cm$^{-2}$). We do not find evidence for significant differences between the X-ray properties and fraction of magnetic systems of CVs in local space, vs. those of CVs in globular clusters or the Galactic Center.   \\begin{quote} \\verb''"
    },
    "0801/0801.4340_arXiv.txt": {
        "abstract": "{The operational conditions found by $BeppoSAX$ in observing nearby (z$\\leq$0.1) Seyferts were reproduced for $Simbol$-$X$ in order to simulate a realistic final database of the mission. The results indicate that, even in the worst conditions, the $Simbol$-$X$ archive of pointings will allow to fully characterize the high-energy spectrum of nearby Seyferts and, most importantly, to obtain solid results on R and Ec (fundamental to model the cosmic X-Ray background, CXB). The measurement of the inclination angle of the accretion disk will be possible for $\\sim$15 objects allowing to directly test the unified models for AGN. Finally, the time-dependent characteristics of the reflected component will be studied in at least $\\sim$25 objects. ",
        "introduction": "In X-ray astronomy each mission has been an important step forward in the comprehension of the physics of known sources and in the opening of new interesting fields of research.  This is expected to happen also for $Simbol$-$X$ (see, for instance the presentations by Fiore in these proceedings) thanks to its unprecedented sensitivity above 10 keV ($\\sim$2-3 orders of magnitude with respect the previous missions). Moreover $Simbol$-$X$ will also work as an observatory that will leave a rich but, as usual, inhomogeneous archive of data. As such,trying to estimate on how many objects deep spectral or timing studies will be performed, so as to fully exploit the astrophysical potentials of the new instruments, is not trivial. To make an attempt of tackling this issue for $Simbol$-$X$, it has been chosen to test the new mission reproducing the observational conditions (essentially the flux of the sources and the used exposure) really found by the only past mission comparable with $Simbol$-$X$: $BeppoSAX$. The astrophysical field inspected here is the one concerning the nearby (z$\\leq$0.1) Seyfert galaxies, for which $BeppoSAX$ obtained important results on the X-rays emission mechanism, on the geometry of the reflector and on the geometry of the cold absorber (Perola et al. 2002, Risaliti et al. 1999). These results were obtained  investigating the properties of the high-energy cut-off (Ec), and of the reflection component (R) which study is possible only using broad-band observatories. These quantities are hardly measurable and they are known only for few bright objects. They were mainly measured from $BeppoSAX$ which operated above 10 keV with a smaller ($\\sim$2-3 order of magnitude less) sensitivity with respect $Simbol$-$X$ but with a broader band ($\\sim$0.1-200 keV, instead of $\\sim$0.5-80 keV). The capability of inspecting Ec and R is here used as indicator of quality of the data assuming that, if R and Ec are measured, than it will be possible to deeply investigate, both via spectral and timing analysis, the physical conditions close the super-massive black-hole. The goal of this work is to estimate a realistic number of nearby Seyfert that $Simbol$-$X$ will observe in this detail.   \\scriptsize \\begin{table} \\caption{Grid of possible values for the interesting parameters of the simulated $Simbol$-$X$ spectra.} \\label{} \\vspace{-0.55cm} \\begin{center} \\begin{tabular}{cccccc} \\hline\\hline \\\\ N$_{H}^{a}$&$\\Gamma$&R&Ec$^{b}$&F$_{2-10 keV}^{c}$&Exp.$^{d}$\\\\ &&&&&\\\\ \\hline &&&&&\\\\ 0&&0.5&150&1&100\\\\ &&&&&\\\\ 5&1.85$^{f}$&1&200&0.1&50\\\\ &&&&&\\\\ 50&&1.5&250&0.01&10\\\\ &&&&&\\\\ \\hline \\\\ \\end{tabular} \\end{center} \\vspace{-0.55cm} $^{a}$ in units of 10$^{22}$ cm$^{-2}$; $^{b}$ in units of keV; $^{c}$ in units of 10$^{-11}$ erg s$^{-1}$ cm$^{-2}$; $^{d}$ in units of ks; $^{f}$ fixed  \\end{table} \\normalsize   \\vspace{-0.35cm} ",
        "conclusions": "This work is a first attempt to determine a realistic number of nearby (z$\\leq$0.1) Seyfert galaxies for which $Simbol$-$X$ will carry deep spectral and timing analysis. The basic assumption here are: 1) that these studies will be possible when the measurements of R and Ec are possible; 2) that Simbol-X will observe sources  with a distribution of F similar to what done by $BeppoSAX$. The last assumption is quite strong and it is probable that the characteristics of Simbol-X will be used to inspect sources that were barely detectable with $BepoSAX$, as the Compton-thick sourcces. Thus, what obtained here may be considered as lower limits. The results can be summarized as  follows: i) the fraction of nearby Seyfert galaxies detected above 10 keV should be $\\sim$100\\%; ii) for $\\sim$80 Seyfert galaxies in the local universe it will be possible to measure Ec and R; iii) $\\sim$50\\% of these sources are expected to be type II sources;  iv) for $\\sim$25 objects it will be possible to perform meaningful time dependent studies of R; v) for $\\sim$15 type I sources is will be possible to study the inclination angle of the accretion disk inspecting the properties of the reflected components.  \\begin{figure} \\includegraphics[width=2.2cm,height=6.4cm,angle=270]{dadina_f6.ps} \\includegraphics[width=6.8cm,height=2.2cm,angle=0]{dadina_f7.ps} \\caption{Upper panel: Confidence contours for the inclination angle ($\\Theta$) for a type I source with F=1, R=1 and Ec=200. Lower panel: histogram of the F parameter. The shaded region correspond to the observations for which information on the inclination angle is available studying the reflected component.} \\vspace{-0.3cm} \\end{figure}   These results, thus, indicate that $Simbol$-$X$ is expected to leave an archive of pointed observations that will allow to fully characterize the average spectrum of nearby Seyfert galaxies on solid statistical basis (for instance, the distribution of N$_H$ for the Seyfert galaxies in the local Universe was obtained with a sample of 74 objects, Risaliti et al. 1999). Most importantly, it is expected to obtain solid description of the distributions of R and Ec, fundamental to synthesize the CXB. Moreover it will be possible to test the UM for the brightest Seyferts checking if the inclination angle of the system is in accordance with the optical classification of the sources. Finally, the astrophysics of the accretion will be studied using timing techniques at the energies where the bulk of the reflection peaks.   \\vspace{-0.45cm}  \\subsection{Caveats}  \\vspace{-0.2cm}  Some basic assumption and simplifications have been made in this work and these must be kept in mind when evaluating the the opportunities offered by $Simbol$-$X$ in the inspected astrophysical field. The sample of objects/observations used here have been accumulated in six years while $Simbol$-$X$ is expected to fly for 2-5 years (Ferrando et al., 2006). Nonetheless, the higher observational efficiency (up to $\\sim$90\\%, Ferrando et al., 2006) should permit $Simbol$-$X$ to observe the nearby Seyferts for $\\sim$10 Ms as BeppoSAX did. The baseline model assumed here does not include warm absorber or ionized reflection, nor include soft excess in type I sources. All these components introduce additional spectral complexities not accounted for. On the contrary, the capabilities of $Simbol$-$X$ in measuring R and Ec were here tested only on pure statistical basis: i.e. no physical assumption were made during the fitting and the significance of the measurements of R and Ec were calculated leaving all the other parameters free to vary.   \\vspace{-0.2cm}  \\small"
    },
    "0801/0801.3756_arXiv.txt": {
        "abstract": "Emission lines in X--ray spectra of clusters of galaxies reveal the presence of heavy elements in the diffuse hot plasma (the Intra Cluster Medium, or ICM) in virial equilibrium in the dark matter potential well.  Thanks to the X--ray satellites {\\sl Chandra} and {\\sl XMM--Newton} we are now able to measure with good accuracy the distribution and evolution of Iron up to redshift $z\\sim 1.3$.  The capability of studying the chemical and thermodynamical properties of the ICM in high redshift clusters is an efficient tool to constrain the interaction processes between the cluster galaxies and the surrounding medium.  We confirm that the ICM is already significantly enriched at a look-back time of 9~Gyr, and find that the Iron abundance change with redshift as $\\propto (1+z)^{-1.25}$, implying an increase of a factor of $\\sim 2$ with respect to $z\\simeq 1.3$.  This result can be explained by a prompt enrichment by star formation processes in massive ellipticals at $z\\geq 2$, followed by a slower release of enriched gas from disk galaxies into the ICM, associated to a morphological transition from disk to S0. ",
        "introduction": "Clusters of galaxies are the largest virialized objects in the Universe.  They form via gravitational instability from the initial perturbations in the matter density field.  Most of their baryonic content is in the form of diffuse gas.  The diffuse baryons are heated to temperatures between $kT \\sim 1$ keV and $kT \\sim 10$ keV (roughly corresponding to masses ranging from $10^{14}$ to $5 \\times 10^{15} M_\\odot$) and constitute the Intra Cluster Medium (from now on ICM). Thanks to the very low density of the electrons (typically $n_e \\sim 10^{-3}$ cm$^{-3}$) the ICM is optically thin and it is in a state of collisional equilibrium established between the electrons and the heavy ions.  The resulting emission from the ICM in the X--ray band is described by a continuum component due to thermal Bremsstrahlung (roughly scaling as $\\propto T^{1/2} n_e^2$) plus line emission from K--shell and L--shell transitions of heavy ions (Iron being the most prominent).  The collisional equilibrium allows one to derive directly from the X--ray spectrum of the ICM both the electron temperature and the abundances of heavy elements.  X--ray cluster samples have always been considered one of the best tools for cosmology, thanks to the simple link between the X--ray spectral properties (namely the electron temperature) and the total dynamical mass, and to the simple selection function depending only on the flux limit.  Their number density as a function of the mass scale and of the cosmic epoch, strongly depend on cosmological parameters. The number density of massive (hot) clusters is expected to decrease at high redshift at a rate sensibly depending on cosmological parameters (see \\cite{r02}).  In this respect, the detection of high redshift clusters has a very strong impact.  About fifteen years ago, before results from deep surveys with the ROSAT satellite had been published \\cite{R98}, the presence of massive, hot clusters at redshift as low as $z\\sim 0.2$ was strongly questioned.  Now, the presence of a significant number of massive clusters at redshift as large as $z\\sim 1.4$ is firmly established, and this represents a strong hint in favour of a low density, $\\Omega_0\\sim 0.3$, $\\Lambda \\sim 0.7$ CDM Universe.  The highest-z cluster selected so far in the X--ray band is at redshift $z\\sim 1.4$ \\cite{stan06}.  In particular, we would like to remark that the value of the normalization of the initial density perturbation field $\\sigma_8\\sim 0.7$ measured with X--ray clusters (\\cite{bor99}, see Figure \\ref{sigma8}), initially at variance with the measure $\\sigma_8 = 0.84 \\pm 0.04$ in the first year of WMAP, is now in agreement with the new WMAP results \\cite{spe07}. Once again, this is a good example of the complementarity of X--ray clusters with respect to other geometrical cosmological tests, like high--z SNe and CMB experiments.  \\begin{figure}[t!] \\centering \\includegraphics[height=6cm]{tozzip_fig1.ps} \\caption{\\footnotesize Confidence levels in the $\\sigma_8$--$\\Omega_m$ plane from the X--ray luminosity function of distant clusters of galaxies observed with ROSAT.  Different panels correspond to different parametrization of the $L$--$T$ relation (see \\cite{bor99}).} \\label{sigma8} \\end{figure}  Now that we are in the maturity of the {\\sl Chandra} and {\\sl XMM--Newton} era, we realize that clusters harbor an unexpected complexity, that demands a much deeper understanding of the thermodynamics of the ICM and its relation with the other mass components, such as member galaxies and dark matter. One of the biggest evidence of this complexity is given by the cool core problem. In more than half of local clusters the cooling time within 10 kpc can be significantly less than 1 Gyr \\cite{pet06}.  It seems unavoidable to predict that baryons are flowing to the cold phase at a rate of the order of 100, sometimes 1000 $M_\\odot$ yr$^{-1}$.  The simplest model based on isobaric cooling, predicts a spectrum rich in emission lines, which are strongly increasing at low temperatures due to the higher number of ion species.  However, grating spectroscopy of the brightest central regions of clusters, with the {\\sl XMM--Newton} satellite, provided a surprising result.  Many of the lines expected in cooling flows were missing from the observed spectra \\cite{pet01}.  It can be shown that the lowest temperature in the center is of the order of 1/3 of the virial one.  This discovery has a strong impact: it implies that the ICM is kept above a temperature floor, not too far from the virial one, by some heating mechanism.  Which is the process that keep the temperature of the ICM above $1/3 \\, T_{vir}$?  The presence of some extra--energy in the ICM was already known from the study of the scaling relations between X--ray observables and from the observation of an entropy excess in the ICM with respect to the self--similar scaling \\cite{pon99}.  However, there is no consensus on the sources which constantly heat the ICM.  This is an open problem, relevant not only to the ICM, but also for general framework of galaxy formation and evolution.  The main problem with feedback, is that any process we may think of, scales with volume (and then with density $\\rho$), while cooling is a runaway process proportional to $\\rho^2$.  Therefore there is not an obvious, mechanism for self--regulation. Understanding feedback is nowaday the most compelling goal for structure formation models.  Main candidates are SNe explosions and stellar winds (as confirmed by the presence of heavy elements in the ICM) and the much more energetic output from AGN.  A spectacular example that favours AGN as the main heating sources, is the X--ray image of the Persues cluster \\cite{fab06}, where hot bubbles created by the jets are pushing the ICM apart, with a total mechanical energy sufficient to heat significantly the diffuse baryons. Still, how and on which time scale the energy is thermalized into the ICM is still a matter of debate.  It is worth mentioning that the presence of the non--gravitational energy input does not hamper the use of clusters as a tool for cosmology.  The relation between X--ray observables and the dynamical mass is still valid, and maybe it is even tighter once we provide a comprehensive model for the ICM thermodynamics.  A recent example is the use of the parameter $T_X M_{gas}$ which appears to be tightly correlated to the dynamical mass \\cite{kr06}.  Therefore, this increasing complexity is not an impediment, but, on the contrary, is giving us a more powerful tool to investigate at the same time many aspects of structure formation, from subgalactic to cluster scales. The price to pay is, of course, a bigger effort in describing the ICM physics.  In this Contribution, I will show how the study of the evolution of Iron abundance with redshift through detailed X--ray spectroscopy of the ICM may help to shed light on the interaction processes between cluster galaxies and the ICM itself. ",
        "conclusions": "Precise measurements of the Iron content of clusters over large look--back times provide a useful fossil record for the past star formation history of cluster baryons.  In a recent work \\cite{bal07} we confirm the presence of a significant amount of Iron in high-$z$ clusters.  We find significant evidence of a decrease in $Z_{Fe}$ as a function of redshift, which can be parametrized by a power law $\\langle Z_{Fe}\\rangle \\simeq (0.54\\pm 0.04)Z_\\odot \\,(1+z)^{-1.25}$ (where solar abundances refer to \\cite{angr89}). This result can be explained by a prompt enrichment by star formation in massive ellipticals at $z\\geq 2$, followed by a slower release of enriched gas from disk galaxies, associated to a morphological transition from disk to S0.  To investigate further the interaction processes between ICM and cluster galaxies, we need to understand at the same time the distribution and the evolution of heavy elements in the ICM, and therefore we rely on high--resolution, spatially resolved X--ray spectroscopy.  Next future X--ray clusters studies must face a double challenge: first, understanding the physics of the ICM, exploring in greater details the cool--cores and the low surface brightness regions in the outskirts of clusters; second, finding more clusters and groups at high redshifts. To capitalize and extend what we have learned so far, we must have soon both a wide area medium--deep survey as well as a mission devoted to the properties of the ICM.  Looking further into the future, the next generation of large X--ray telescopes must achieve a spatial resolution comparable to that of {\\sl Chandra}, which proved to be crucial to avoid confusion limit and to isolate genuine diffuse emission at very high redshifts."
    },
    "0801/0801.4030_arXiv.txt": {
        "abstract": "We model the X-ray properties of millisecond pulsars (MSPs) by considering hot spot emission from a weakly magnetized neutron star (NS) covered by a hydrogen atmosphere.  We investigate the limitations of using the thermal X-ray pulse profiles of MSPs to constrain the mass-to-radius ($M/R$) ratio of the NS. The accuracy is strongly dependent on the viewing angle and magnetic inclination but is ultimately limited only by photon statistics. We demonstrate that valuable information regarding NSs can be extracted even from data of fairly limited photon statistics through modeling of archival observations of the nearby isolated PSRs J0030+0451 and J2124--3358. The X-ray emission from these pulsars is consistent with the presence of an atmosphere and a dipolar field configuration.  For both MSPs, the favorable geometry allows us to place limits on the allowed $M/R$ of NSs. Assuming 1.4 M$_{\\odot}$, the stellar radius is constrained to be $R > 9.4$ km and $R > 7.8$ km (68\\% confidence) for PSRs J0030+0451 and J2124--3358, respectively. We explore the prospects of using future observatories such as \\textit{Constellation-X} and \\textit{XEUS} to conduct X-ray timing searches for MSPs not detectable at radio wavelengths due to unfavorable viewing geometry. We are also able to place strong constraints on the magnetic field evolution model proposed by Ruderman. The pulse profiles indicate that the magnetic field of an MSP does not have a tendency to align itself with the spin axis nor migrate towards one of the spin poles during the low-mass X-ray binary phase. ",
        "introduction": "Recent X-ray studies have revealed that a number of known rotation-powered millisecond pulsars (MSPs) exhibit predominantly thermal soft X-ray emission \\citep{Grind02,Zavlin06,Bog06a,Zavlin07}. The infered effective emission radii $R_{\\rm eff}$ indicate that this radiation is localized in regions on the neutron star (NS) surface much smaller than the expected stellar radius ($R_{\\rm eff}\\ll R$) but comparable to the classical radius of the pulsar magnetic polar cap $R_{pc}=(2\\pi R/cP)^{1/2}R$. This finding is in agreement with theoretical models of pulsars, which indicate that the conditions in the magnetosphere of a typical MSP favor heating of the polar caps to $\\sim$$10^6$ K by a return flow of energetic particles along the open magnetic field lines \\citep[see, e.g.,][for details]{Hard02,Zhang03}. As this heat is restricted to a small fraction of the NS, study of the X-ray spectra and pulse profiles of MSPs can reveal important information about the star such as the radiative properties of the NS surface, magnetic field geometry, and NS compactness ($R/R_S$, where $R_S=2GM/c^2$). This approach, originally proposed by \\citet{Pavlov97} in the context of radio MSPs, can serve as a valuable probe of key NS properties that are inaccessible by other observational means (e.g.~radio pulse timing). As shown by \\citet{Pavlov97}, \\citet{Zavlin98}, and \\citet{Bog07}, a model of polar cap thermal emission from an optically-thick hydrogen (H) atmosphere provides an excellent description of the X-ray pulse profiles of PSR J0437--4715, the nearest known MSP. On the other hand, a blackbody model is inconsistent with the X-ray timing data and can be definitively ruled out. Furthermore, there is compelling evidence for a magnetic dipole axis offset from the NS center \\citep{Bog07}. Finally, the compactness of PSR J0437--4715 is constrained to be $R/R_S>1.6$ (99.9\\% confidence).  Thus, modeling of X-ray data of MSPs appears to be a very promising approach towards answering long-standing questions regarding the fundamental properties of MSPs and NSs, in general.  The present paper represents an extension of the work described in \\citet{Bog07}. Herein we explore the detailed properties of the MSP X-ray emission model with particular emphasis on its use for constraints on the NS equation of state (EOS). The work is organized as follows. In \\S2 we examine the properties of our model; in \\S3 we discuss an application of our model to archival X-ray observations of PSRs J0030+0451 and J2124--3358. In \\S4 we present a discussion and end with conclusions in \\S5.    \\begin{figure*}[t!] \\includegraphics[width=0.98\\textwidth]{f1.ps} \\caption{(\\textit{Left}) Representative synthetic hydrogen atmosphere lightcurves for a rotating $M=1.4$ M$_\\odot$, $R=10$ km NS with two antipodal hot spots for the four lightcurve classes (I-IV, from top to bottom, respectively), defined by Beloborodov (2002). The dashed lines show the individual flux contribution from the two hot spots while the solid line shows the total observed flux.  All fluxes are normalized to the value corresponding to a face-on hot spot ($\\alpha=\\zeta=0$). Two rotational cycles are shown for clarity. (\\textit{Right}) Orthographic map projection of the MSP surface for the four pulse profiles. The dashed line shows the magnetic axis while the dotted line shows the line of sight to the observer. The hatched region corresponds to the portion of the star not visible to the observer.} \\end{figure*} ",
        "conclusions": "We have examined the properties of our model of thermal emission from hot spots on the surface of a neutron star covered by a hydrogen atmosphere, relevant for MSPs.  Our investigation has demonstrated that energy-resolved modeling of the thermal X-ray pulse profiles and phase-resolved spectroscopy of the continuum emission can, in principle, be used to determine $M/R$ and the pulsar geometry to high accuracy, given observational data of sufficient quality and favorable values of $\\alpha$ and $\\zeta$. As shown in \\S2.1 the thermal pulse profiles are surprisingly insensitive to the details of the polar cap size and shape, the distance to the pulsar, and the uncertainty in the instrument effective area.  This method represents the only feasible approach of studying MSP magnetic field, surface properties, and compactness, especially for isolated MSPs.  Barring any deleterious effect such as additional hidden spectral components \\citep[see, e.g,][and references therein]{Bog06b} this approach can, in principle, lead to unprecedented insight into NS properties.   Our model is found to be in agreement with the observed emission from the nearby solitary MSPs J0030+0451, and J2124--3358.  As with PSR J0437--4715 \\citep{Bog07}, the relatively large pulsed fractions observed in PSRs J0030+0451 and J2124--3358 require the existence of a light-element atmosphere on the stellar surface and cannot be reproduced by a blackbody model for realistic NS radii.  For J0030+0451 and J2124--3358 we are able to place interesting limits on the allowed compactness of $R>9.4$ km and $R>7.8$ km (68\\% confidence) assuming $M=1.4$ M$_{\\odot}$.  Based on our findings in \\S2, we expect deeper observations of this MSP to lead to much tighter constraints on $M/R$, which in turn may firmly rule out certain families of NS EOS. The available thermal X-ray pulse profiles also provide useful constraints on magnetic field models of MSPs. Specifically, the positions of the magnetic polar caps on the NS surface implied by the X-ray data indicate that the magnetic field closely resembles the conventional oblique dipole model of pulsars rather than more exotic field configurations.  The success of this approach towards elucidating crucial NS properties motivates further studies of the nearby sample of MSPs with both \\textit{Chandra} and \\textit{XMM-Newton}. Moreover, this makes MSPs particularly important targets for upcoming X-ray mission such as \\textit{Constellation-X} and \\textit{XEUS}. The great increase in throughput of these facilities will allow searches for new MSPs, detailed observations of a larger sample of known radio MSPs, and unprecedented constraints on key NS properties, especially the NS EOS."
    },
    "0801/0801.4801_arXiv.txt": {
        "abstract": "We describe the construction of GROND, a 7-channel imager, primarily designed for rapid observations of gamma-ray burst afterglows. It allows simultaneous imaging in the Sloan $g'r'i'z'$ and near-infrared $JHK$ bands. GROND was commissioned at the MPI/ESO 2.2\\,m telescope at La Silla (Chile) in April 2007, and first results of its performance and calibration are presented. ",
        "introduction": "Simultaneous imaging in different filter-bands is of interest in a variety of astrophysical areas. The primary aim is to measure the spectral energy distribution or its evolution in variable objects in order to uncover the underlying emission mechanism. Examples are, among others, (1) monitoring of all kinds of variable stars (flare stars, cataclysmic variables, X-ray binaries) to determine the outburst mechanisms and differentiate between  physical state changes and changes induced by geometrical variations, like eclipses; (2) monitoring of AGN to understand the physical origin of the observed variability; (3) determining the inclination of X-ray heated binaries \\cite{orosz} (4) mapping of galaxies to study the stellar population; (5) multi-color light curves of supernovae \\cite{tpm06}; (6) differentiate achromatic microlensing events \\cite{pac86} from other variables with similar light curves; (7) identifying objects with peculiar spectral energy distributions, e.g. photometric redshift surveys for high-$z$ active galactic nuclei, or identifying brown dwarfs; (8) follow-up observations of transiting extrasolar planets \\cite{jha00}; or (9) mapping of reflectance of solar system bodies as a function of their rotation to map their surface chemical composition \\cite{jew02}.  A new need for multi-band imaging arose with the observation of a large number of gamma-ray burst (GRB) afterglows with the {\\it Swift} satellite \\cite{geh04}. With its much more sensitive instruments it detects GRBs over a very wide redshift range. Since intermediate to high-resolution spectroscopy to measure the physical conditions of the burst environment \\citep[e.g.][]{vre07} is constrained to the first few hours after a GRB explosion, a rapid determination of the redshift became important. This is best done with multi-band photometry (until integral field units have grown to several arcmin field-of-views) and deriving a photometric redshift based on the Ly$\\alpha$ break \\cite{lr00}.  Previous and current instruments with simultaneous imaging capability in different filter bands include ANDICAM (A Novel Double-Imaging CAMera; two channels, one for visual, the other for near-infrared \\cite{dep98}, presently operated at a 1.3\\,m telescope), BUSCA (Bonn University Simultaneous CAmera; four visual channels \\cite{rei99}, presently operated at the 2.2\\,m telescope at Calar Alto), HIPO (High-speed Imaging optical Photometer for Occultations; two visual channels \\cite{dun04}, to be operated on SOFIA, the Stratospheric Observatory For Infrared Astronomy), MITSuME (Multicolor Imaging Telescopes for Survey and Monstrous Explosions; three channels with fixed bands g\\amin, R$_{\\rm C}$ and I$_{\\rm C}$ \\cite{kky07}, operated at a 50 cm telescope), TRISPEC (Triple Range Imager and Spectrograph; three channels with one CCD and two near-IR detectors and wheels for filters, grisms and Wollaston prisms \\citep{wny05}), SQIID (Simultaneous Quad Infrared Imaging Device; $JHK$ and narrow-band $L$ filters in front of  individual 512X512 quadrants of an ALADDIN InSb array, designed for the f/15 Cassegrain foci of the KPNO 2.1-m and 4-m telescopes \\citep{edf92}), ULTRACAM (ULTRAfast, triple-beam CCD CAMera for high-speed astrophysics \\cite{dhi07}; portable instrument which has been used, among others, at the Very Large Telescopes at ESO, or the William Herschel Telescope, Canary Islands).  Here we describe the design ($\\S$2) and performance ($\\S$6) of a 7-channel imager, called GROND ({\\bf G}amma-{\\bf R}ay Burst {\\bf O}ptical and {\\bf N}ear-Infrared {\\bf D}etector), which was specifically designed for GRB afterglow observations. We also mention some basics of the operation scheme ($\\S$3), related software ($\\S$4), and the changes to the telescope infrastructure which were implemented to use GROND for rapid follow-up observations ($\\S$5). ",
        "conclusions": "GROND is an imaging system capable of operating in seven colors simultaneously. It has been designed and built at MPE Garching, and commissioned at the 2.2\\,m MPI/ESO telescope on La Silla, Chile. First observations show that all properties are according to specifications/expectations. The  first observations of gamma-ray bursts with GROND have also been obtained (Greiner \\etal\\ 2007a,b; Primak \\etal\\ 2007, Kr\\\"uhler \\etal\\ 2007). Fine tuning of the operations strategy as well as scheduling and analysis software in the upcoming weeks is expected to bring GROND into a fully operational condition, thus allowing the commencement of normal science operations.   \\bigskip"
    },
    "0801/0801.2616_arXiv.txt": {
        "abstract": "We present a preliminary analysis of ASTRO-F data of a complete sample of $\\sim$ 150 EROs (R-K$>$5) down to K$_{Vega}<$19, for which reliable photometric redshifts are available, in the range 0.8$<$z$<$2, selected over two fields (S7 and S2) of the MUNICS survey. The area covered is about 420 arcmin$^2$. We have imaged this area with AKARI telescope in N3 (3.4 $\\mu$m), N60 (65 $\\mu$m) and WL (150 $\\mu$m) down to 12 $\\mu$Jy in the N3 filter, in order to detect the rest frame H or K-band emission, thus providing an excellent sampling of the SED of our EROs. From a first analysis we have an identification rate of $\\sim$ 63\\% in the N3 filter over the S7 field. These data allow us to distinguish starburst from passive early type phenomena, to meseaure the SFR of the starburst component and to constrain the mass assembly of early type galaxies. ",
        "introduction": "While the general build-up of cosmic structures seems to be well described by hierarchical models of galaxy formation ($\\Lambda$CDM), the assembly of the baryonic mass on galactic scale still represents a weak point of the galaxy formation models. One of these difficulties consists in explaining the large population of Extremely Red Objects (R-K$>$5) EROs. This population is a mixture of early type galaxies (ETG) and dusty star forming galaxies (SFG). Hierarchical models should reproduce the abundance of massive red galaxies at 1$<z<$2 and the balance among the number of early type and obscured star-forming galaxies \\cite{ref:Nagamine}. Mid/Far-IR observations can play a key role in the study of this pupulation: such data can allow to disantangle between the early type and star forming phenomena and, in particular, mid-IR data allow to assess a more reliable estimate of stellar masses of EROs, as they allow to sample the rest-frame near-IR emission. Since the mid-IR observations are usually not available for large samples of EROs, the rest frame H-K emission is extrapolated from the observed optical emission (partially affected by extinction) which is dominated by young stars ($<$0.5-1Gyr), whose contribution to the near-IR emission (dominated by older stars) is negligible. As the EROs are mainly composed of old stars, the reliability of stellar masses estimates is extremely uncertain. The way to provide the models with robust observational constraints passes through mid-IR observations of a large ($>$100), bright (K$<$19) and complete sample of EROs.  \\begin{figure} \\centering \\includegraphics[width = 6 cm, height = 5cm]{mignanoa_fig1.jpg} \\includegraphics[width = 7 cm, height = 5.3cm]{mignanoa_fig2.jpg} \\label{detection} \\caption{({\\it left panel}) The S7F5 field imaged with N3 filter (IRC camera). The yellow circle represent the EROs selected from the MUNICS catalogue. ({\\it right panel}) N3 counterpart distribution vs. r. Red histogram shows the counterparts included in the identification sample.} \\end{figure} ",
        "conclusions": ""
    },
    "0801/0801.0878_arXiv.txt": {
        "abstract": "We present medium resolution spectropolarimetry and long term photo-polarimetry of two massive post-red supergiants, IRC~+10420 and HD~179821.  The data provide new information on their circumstellar material as well as their evolution. In IRC~+10420, the polarization of the H$\\alpha$ line is different to that of the continuum, which indicates that the electron-scattering region is not spherically symmetric.  The observed long term changes in the polarimetry can be associated with an axi-symmetric structure, along the short axis of the extended reflection nebulosity. Long term photometry reveals that the star increased in temperature until the mid-nineties, after which the photospheric flux in the optical levelled off. As the photometric changes are mostly probed in the red, they do not trace high stellar temperatures sensitively. And so, it is not obvious whether the star has halted its increase in temperature or not. For HD~179821 we find no polarization effects across any absorption or emission lines, but observe very large polarization changes of order 5\\% over 15 years. During the same period, the optical photometry displayed modest variability at the 0.2 magnitude level. This is unexpected, because large polarization changes are generally accompanied by strong photometric changes. Several explanations for this puzzling fact are discussed. Most of which, involving asymmetries in the circumstellar material, seem to fail as there is no evidence for the presence of hot, dusty material close to the star.  A caveat is that the sparsely available near-infrared photometry could have missed periods of strong polarization activity.  Alternatively, the variations can be explained by the presence of a non-radially pulsating photosphere. Changes in the photometry hint at an increase in temperature corresponding to a change through two spectral subclasses over the past ten years. ",
        "introduction": "Due to the steepness of the Initial Mass Function, massive stars ($\\ga$ 8 M$_{\\odot}$) are extremely rare. This is exacerbated by their comparatively short lifetimes. Yet, although rare, these objects have a crucial impact on the interstellar medium due to their strong winds and high mass-loss rates, and can dominate the light output of entire galaxies.  An area of current interest is that massive evolved stars are often surrounded by bi-polar nebulae \\citep{weis:2003}. Later in the evolution of a star, using spectropolarimetry, it is now established that the ejecta of supernovae deviate from spherical symmetry (e.g. \\citealt{Wang_etal:2003,Leonard_etal:2005}).  It is as yet unclear whether this is due to asymmetric explosions, axi-symmetric stellar winds or a pre-existing density contrast in the surrounding material (e.g. \\citealt{Dwarkadas_Balick:1998,Dwarkadas_Owocki:2002}). Wind-axisymmetry may imply fast rotation, and be related to the beaming of SN explosions, which may be the origin of the extremely luminous, beamed gamma-ray bursts (e.g. \\citealt{Meszaros:2003,Mazzali_etal:2003}).  Here, we address the issue by investigating the circumstellar ejecta of two yellow hypergiants, IRC~+10420 and HD~179821. These objects are thought to have evolved off the post-Red Supergiant branch and are still surrounded by mass ejected during a previous mass losing phase. Only a few yellow hypergiants are known (see the review by \\citealt{de_Jager:1998}), and the number of such hypergiants with circumstellar dust is even smaller - only IRC~+10420 and HD~179821 belong to this class (see for example \\citealt{Oudmaijer08}). Therefore, the study of these two unique objects is important in its own right.    As these stars are distant (3-5 kpc), the direct imaging of their innermost regions is currently beyond the reaches of current technology, although interferometry is starting to resolve the winds of evolved stars \\citep{dewit_etal:2008}.  Observing with spectropolarimetry allows us to probe regions much closer to the star still.  Spectropolarimetry was first effectively used in the study of classical Be stars using the presence of `line effects' \\citep{Poeckert:1975, Poeckert_Marlborough:1976}. These are changes in polarization across spectral lines that have an emission component. They occur because emission-line photons arise over a larger volume than the stellar continuum photons. Consequently, the emission-line photons undergo fewer scatterings as they `see' fewer electrons, resulting in a lower polarization than the continuum. We normally only observe a net polarization change if the geometry of the electron scattering region is aspherical. Many authors have confirmed that this technique provides evidence that envelope geometries around Be stars are indeed disk-like \\citep{dougherty_1992, Quirrenbach_etal:1997, Wood_Bjorkman_Bjorkman:1997}.  More recently the technique has been used to investigate the geometry of circumstellar material around Herbig Ae/Be stars. Studies by \\cite{Oudmaijer_Drew:1999} and \\cite{Vink_etal:2002} show that most Herbig stars exhibit line effects, indicating aspherical electron-scattering regions. For a recent review, see \\cite{Oudmaijer:2007}.  With regard to evolved stars, Davies et al. (2005) conducted a study of Luminous Blue Variables (LBVs) using spectropolarimetry. They found that 50$\\%$ of the objects observed exhibited polarization changes across H$\\alpha$, indicating that some asphericity lies at the base of the stellar wind. Furthermore, they found several objects for which the position angle varied randomly with time, leading them to conclude that the wind around these stars is clumpy (see also \\citealt{Nordsieck_etal:2001}).    IRC~+10420 is now well accepted as a massive, evolved object (e.g. \\citealt{Jones_etal:1993, Oudmaijer_etal:1996} and \\citealt{Humphreys_etal:2002}). This is mainly based on its large distance, high outflow velocity (40 kms$^{-1}$) and high luminosity implied from the hypergiant spectrum.  The situation for HD~179821 is less certain. The presence of non-radial pulsations coupled with comparatively modest photometric changes suggest a massive nature \\citep{LeCoroller_etal:2003}. Furthermore, its circumstellar material has a large expansion velocity of 30 km s$^{-1}$, as measured in CO, suggesting the star is a supergiant \\citep{Kastner_Weintraub:1995}. On the other hand, the overabundance of s-process elements and the low metallicity suggest HD~179821 is perhaps a lower mass post-AGB star \\citep{Zacs_etal:1996,Reddy_Hrivnak:1999,Thevenin_etal:2000}.  Although the latter's nature is a bit more uncertain, there are some striking similarities between IRC~+10420 and HD~179821.  When observed as part of a larger sample of post-AGB stars, these two objects are often markedly different from the rest. In particular, their high outflow velocities (the average outflow velocity for post-AGB and AGB stars is 15 kms$^{-1}$) require much higher luminosities if powered by radiation pressure alone \\citep{Habing_etal:1994}. They were the only objects that showed extensive reflection nebulae in a large survey by \\cite{Kastner_Weintraub:1995}. \\cite{Jura_etal:2001} point out the enormous difference between the space velocities of IRC~+10420 and HD 179821 when compared against low mass post-AGB stars. Furthermore, both objects have an exceptionally strong O{\\sc i} 7774 triplet absorption feature indicating a high luminosity (\\citealt{Humphreys_etal:1973} and \\citealt{Reddy_Hrivnak:1999}, based on \\citealt{slowik_1995}). We therefore proceed with both objects and assume they are evolved post-Red Supergiants.  Recently, both objects have been observed at arcsecond resolution in CO by \\cite{Castro-Carrizo_etal:2007} who found, in accordance with previous estimates, that their mass loss rates exceeded 10$^{-4} \\, \\rm M_{\\odot} yr^{-1}$ when they were in the RSG phase. The envelopes show mild deviations from spherical symmetry in their data.  This paper is organised as follows. In Section 2, we review the experimental setup and explain how the data has been reduced. We present our results for each object in Section 3, and use the new spectropolarimetric data together with past polarization measurements to investigate the nature of the circumstellar environments around each of the stars, which is discussed in Section 4.  We conclude in Section 5.  \\begin {table*} \\centering \\begin {tabular}{llcccccccccc} \\hline \\centering Object & Telescope & Date     & Julian Date & $\\%$P & P.A. (Deg)  & $\\%$P (H$\\alpha$) & P.A. (H$\\alpha$) (Deg) \\\\ \\hline HD~179821 \\\\ & AAT  & 15/09/02 & 2452533  & 1.99 $\\pm$ 0.11 & 40 $\\pm$ 2 \\\\ & WHT  & 30/09/04 & 2453279  & 2.00 $\\pm$ 0.15 & 36 $\\pm$ 3 \\\\  IRC~+10420 \\\\ & NOT  & 16/01/98 & 2450830  & 1.95 $\\pm$ 0.05 & 174 $\\pm$ 1 \\\\ & NOT  & 16/05/98 & 2450950  & 1.80 $\\pm$ 0.03 & 174 $\\pm$ 1 \\\\ & AAT  & 18/09/02 & 2452536  & 2.12 $\\pm$ 0.11 & 173 $\\pm$ 2 & 1.28 $\\pm$ 0.11 & 10 $\\pm$ 2\\\\ & AAT  & 15/08/03 & 2452867  & 2.35 $\\pm$ 0.11 & 179 $\\pm$ 1 & 1.61 $\\pm$ 0.11 & 11 $\\pm$ 1\\\\ & WHT  & 30/09/04 & 2453279  & 3.40 $\\pm$ 0.15 & 174 $\\pm$ 2 & 1.93 $\\pm$ 0.15 & 11 $\\pm$ 2\\\\ \\hline \\end{tabular} \\caption{New polarimetric observations of HD~179821 and IRC~+10420 measured in the $R$ band. For the spectropolarimetric data (those taken at the AAT and WHT), the polarization was measured in the continuum region close to H$\\alpha$, while the final columns are the polarizations at the line centers. The H$\\alpha$ line-center polarization of HD~179821 was not calculated as no line effect was observed, and therefore this measurement would be equal to the continuum polarization. The systematic error of the AAT and WHT spectropolarimetric data is estimated to be of order 0.10\\% and 0.15\\%, respectively.  } \\label{T:SUMDATA} \\end {table*}  \\begin {table*} \\centering \\begin {tabular}{llrrrrrrrr} \\hline Julian Date & Telescope  & $B-V$&   $V$& $V-R$& $R-I$  &  $J$ &  $H$ &  $K$  \\\\ % \\hline 2450268     &  CST       &      &      &      &      &    5.62&      &       \\\\ % 2450303     &  CST       &      &      &      &      &    5.36&  4.40&  3.45 \\\\ % 2450643     &  CST       &      &      &      &      &    5.40&  4.44&  3.50 \\\\ % 2450690     &  CST       &      &      &      &      &    5.40&  4.45&  3.49 \\\\ % 2450707     &  CST       &      &      &      &      &    5.37&  4.43&  3.38 \\\\ % 2450950     &  NOT       & 2.76 & 11.06&  2.40&  1.70&        &      &       \\\\ % 2450985     &  CST       &      &      &      &      &    5.38&  4.44&  3.48 \\\\ % 2451037     &  TSAO      & 2.58 & 11.12&  2.42&  1.60&        &      &       \\\\ % 2451039     &  TSAO      & 2.67 & 11.11&  2.41&  1.58&        &      &       \\\\ % 2451042     &  TSAO      & 2.58 & 11.13&  2.42&  1.54&        &      &       \\\\ % 2451043     &  TSAO      & 2.73 & 11.06&  2.49&  1.61&        &      &       \\\\ % 2451047     &  TSAO      & 2.51 & 11.01&  2.41&  1.58&    5.63&  4.55&  3.73 \\\\ % 2451048     &  TSAO      & 2.75 & 11.03&  2.47&  1.60&    5.43&  4.57&  3.58 \\\\ % 2451050     &  TSAO      & 2.67 & 11.01&  2.46&  1.66&        &      &       \\\\ % 2451052     &  TSAO      & 2.75 & 11.03&  2.47&  1.59&    5.35&  4.44&  3.53 \\\\ % 2451057     &  TSAO      & 2.76 & 10.96&  2.47&  1.64&        &      &       \\\\ % 2451063     &  TSAO      & 2.60 & 10.95&  2.43&  1.59&        &      &       \\\\ % 2451075     &  TSAO      & 2.70 & 11.16&  2.52&  1.60&        &      &       \\\\ % 2451082     &  TSAO      & 2.66 & 11.09&  2.47&  1.59&        &      &       \\\\ % 2451083     &  TSAO      & 2.72 & 11.05&  2.48&  1.60&        &      &       \\\\ % 2451087     &  CST       &      &      &      &      &    5.24&  4.29&  3.33 \\\\ % 2451099     &  TSAO      & 2.62 & 11.16&  2.56&  1.66&        &      &       \\\\ % 2451100     &  TSAO      & 2.71 & 11.16&  2.54&  1.66&        &      &       \\\\ % 2451103     &  TSAO      & 2.76 & 11.13&  2.51&  1.54&        &      &       \\\\ % 2451104     &  TSAO      & 2.72 & 11.03&  2.48&  1.59&        &      &       \\\\ % 2451147     &  CST       &      &      &      &      &    5.35&  4.42&  3.46 \\\\ % 2451245     &  TSAO      &      &      &      &      &    5.32&  4.47&  3.56 \\\\ % 2451292     &  CST       &      &      &      &      &    5.40&  4.45&  3.57 \\\\ % \\hline \\end{tabular} \\caption{ Photometry of IRC~+10420. The data come from the Carlos S\\'{a}nchez Telescope (CST), the Nordic Optical Telescope (NOT) and the Tien-Shan Astronomical Observatory (TSAO), see text for details. Typically,  the photometric errors are of order 0.01-0.03 magnitude.} \\label{T:PHOTDATA} \\end {table*} ",
        "conclusions": "We have presented spectropolarimetry and long term photo-polarimetry for two post-Red Supergiants. A strong depolarization across the H$\\alpha$ emission line is found for IRC~+10420, suggesting an electron-scattering region that is not circularly symmetric, confirming the results of \\cite{Davies_etal:2007}, who found such evidence from their integral field spectroscopy of the object. The time evolution indicates that the source has increased in temperature until at least 1995, after which the {\\it J} band photometry indicates this may have levelled off.  If the temperature increase of the object has indeed halted, an increase in mass loss or mass infall rate can be responsible for the observed H$\\alpha$ emission and polarization.  HD~179821, a cooler object, has less H$\\alpha$ emission, and no depolarization across the H$\\alpha$ line could be detected. The photometry implies that this evolved object is also undergoing a change in temperature. Hitherto, temperature determinations of the object were inconclusive, but if the star is a G supergiant, then the observed change in photometry suggests it has become earlier by one or two subclasses. Strong changes at the 5 per cent level in polarization over the past 15 years have been detected. During the same time, the optical photometry has only varied by at most 0.2 magnitudes. The most obvious explanations for this observation such as changes in either electron or dust scattering can be ruled out.  A complication is that the polarization is not correlated with any other observable. The possibility that the star undergoes irregular, asymmetric, mass ejections is discussed. However, there is little evidence for the occurrence of such events. Instead, it is proposed that the star itself is asymmetric and that its anisotropic radiation, which is scattered off the circumstellar dust results in the net polarization observed.   \\subsection*"
    },
    "0801/0801.4083_arXiv.txt": {
        "abstract": "Using the OGLE catalogue of eclipsing binaries, 15 contact binaries were identified towards the SMC and the LMC at vertical distances from the Galactic plane between 300 pc and 10 kpc. Based on the luminosity function calculated for these contact binaries, we estimated a frequency of occurrence relative to Main Sequence stars in the thick disk at roughly $\\frac{1}{600}$. This estimate suffers from the small number statistics, but is consistent with the value previously found for the solar neighbourhood. ",
        "introduction": "During the past few decades, several attempts have been made to determine the local spatial density of contact binaries, with very diversified results. \\citet{ru02} (where full references are given and differences discussed) estimated their local spatial density at $(1.02\\pm0.24) \\times 10^{-5} \\ \\mathrm{pc^{-3}}$. Later, on the basis of the All Sky Automated Survey (ASAS), a relative frequency of occurrence (RFO) of one contact binary among about 500 solar type stars was derived \\citep{ru06}, which is in full agreement with the mentioned spatial density. The RFO for contact binaries at high galactic latitudes, however, has been entirely unknown.  Here we present an estimate for the spatial occurrence of contact binary systems relative to Main Sequence (MS) stars in the thick disk of our Galaxy, far from the Galactic plane. In particular, the luminosity function for contact binaries in two conical volumes towards the Small Magellanic Cloud (SMC) and the Large Magellanic Cloud (LMC) covering parts of the thick disk and the halo is determined. The results are then compared with the luminosity function for MS stars in the same region of the sky. The term ``contact binaries'' is here used as a synonym for W~UMa-type eclipsing binaries with orbital periods in a range of 0.22 -- 1 days. In this paper, for reasons to be explained below, we limit ourselves to a subset with periods $<0.45$ days.  \\section[Identification of contact binaries]% {Identification of short period contact binaries in the OGLE-catalogue of eclipsing binaries}  The Optical Gravitational Lensing Experiment (OGLE) was intended to detect dark matter in the Milky Way Galaxy using the microlensing technique, with the Magellanic Clouds and the Galactic Bulge being the main targets of the survey \\citep{usk97}. As a by-product, OGLE provides high quality, long-term photometry that can be used to analyse eclipsing binary stars. Two online catalogues \\citep{wy03,wy04} were used for this study. They contain data in the standard photometric $BVI$ system for 2850 and 1351 eclipsing binaries, and cover an area of 4.6 and 2.4 square degrees of the central parts of the LMC and the SMC respectively. The OGLE-II survey has a faint limit of roughly $I=20.5 \\ \\mathrm{mag}$ with a corresponding error of $0.3 \\ \\mathrm{mag}$; the error becomes smaller for decreasing magnitude, reaching a bright limit of the survey at about 13 mag \\citep{wy04}.  The differentiation between contact binaries and other eclipsing binary types in the two OGLE catalogues was carried out by applying a contact binary criterion based on Fourier analysis of the light curves, which was introduced by \\citet{ru97}\\footnote{More detailed explanations for this and other techniques used in this paper can be found in \\citep{ru97} and \\citep{ru06}.}. By means of this criterion and using a visual inspection of the remaining light curves, we identified 10 and 5 contact binaries with orbital periods $P<0.45 \\ \\mathrm{d}$ towards the LMC and the SMC, respectively. The light curves are plotted in Figure \\ref{fig_multi_plot}. Furthermore, we used the online database from the MACHO (Massive Compact Halo Object) project \\citep{al97} to confirm the classification of the contact binaries in the OGLE catalogue. A direct comparison with the OGLE catalogue was only possible for the contact binaries towards the LMC, because there is no MACHO photometry available for most survey fields covering the SMC. The MACHO photometry of the 10 contact binaries towards the LMC is in good agreement with the results obtained from the OGLE survey\\footnote{As an exception, the OGLE photometry in $V$-band for the contact binary OGLE050905.22-693315.1 was found to have a gross error of roughly 3 magnitudes.}.   The contact binaries in our sample have orbital periods in a range of 0.22 -- 0.45 days. The short period limit is a natural cut-off for contact binaries \\citep{ru07}, whereas the upper period limit was intentional: On one hand, we wanted to be sure of the contact binary classification and avoid semi-detached binaries, while on the other hand, we wanted to use the simple $M_V \\equiv M_V (\\log P)$ calibration and avoid problems with uncertain or missing colour indices. Because the final RFO estimate is done using absolute magnitude bins of the luminosity function, the period limits signify an intentional restriction to the low brightness end of the contact binary sequence. We will explain the details below. ",
        "conclusions": "We conclude that each of the two samples yields the relative frequency of one contact binary among about $600$ MS stars in the two conical volumes towards the SMC and the LMC at galactic latitudes of $-44\\degr$ and $-33\\degr$ respectively. These estimates have uncertainties of roughly $50 \\%$ due to the large uncertainties arising from small number statistics and to the still uncertain parameters of the spheroidal component of the Galaxy. This component has been explicitly accounted for in the previous section and is estimated to contribute to the total number of stars in the SMC and the LMC search volumes at levels of about 30\\% and 20\\%, respectively. The contribution varies with the $M_V$. Specifically, for the SMC, it changes from 38\\% for the $M_V$-bin with the largest distances to 17\\% for the $M_V$-bin centred at 6 mag; for the LMC, it changes from 29\\% to 10\\% in the same range of $M_V$.  The results on the relative frequency of contact binary stars at large galacto-centric distances is consistent with estimates for the solar neighbourhood in previous work \\citep{ru06}. These are the very first data available on the contact binary distribution at high galactic latitudes.    \\begin{figure} \\begin{center} \\includegraphics[width=84mm]{lf_lmc.eps} \\caption{\\label{fig-lf_lmc}Luminosity function towards the LMC. The continuous line represents the contact binary luminosity function, whereas the dashed line represents the MS luminosity function divided by a factor of 650.} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[width=84mm]{lf_smc.eps} \\caption{\\label{fig-lf_smc} Luminosity function towards the SMC. The same as in figure \\ref{fig-lf_lmc} except for the scaling factor of 600 instead of 650.} \\end{center} \\end{figure}   \\begin{figure*} \\begin{center} \\includegraphics[width=176mm]{multi_plot.eps} \\caption{\\label{fig_multi_plot} The light curves of the 15 close contact binaries. The fitting was done according to the approach of \\citep{ru93} (i.e. using only the first 5 cosine and the first sine Fourier-coefficients).} \\end{center} \\end{figure*}"
    },
    "0801/0801.4560_arXiv.txt": {
        "abstract": "Narrow Line Seyfert 1 galaxies (NLS1s) are generally considered peculiar objects among the broad class of Type 1 active galactic nuclei, due to the relatively small width of the broad lines, strong X--ray variability, soft X--ray continua, weak \\Oiii{}, and strong \\Feii{} line intensities. The mass \\mbh{} of the central massive black hole (MBH) is claimed to be lighter than expected from known MBH--host galaxy scaling relations, while the accretion rate onto the MBH larger than the average value appropriate to Seyfert 1 galaxies. In this {\\it Letter}, we show that NLS1 peculiar \\mbh{} and $L/L_{\\rm Edd}$ turn out to be fairly standard, provided that the broad line region is allowed to have a disc--like, rather than isotropic, geometry. Assuming that NLS1s are rather ``normal'' Seyfert 1 objects seen along the disc axis, we could estimate the typical inclination angles from the fraction of Seyfert 1 classified as NLS1s, and compute the geometrical factor relating the observed FWHM of broad lines to the virial mass of the MBH. We show that the geometrical factor can fully account for the ``black hole mass deficit\" observed in NLS1s, and that $L/L_{\\rm Edd}$ is (on average) comparable to the value of the more common broad line Seyfert 1 galaxies. ",
        "introduction": "Seyfert 1 galaxies (Sy1s) are often divided into two distinct classes,  namely Broad Line Sy1s (BLS1s), whose H$\\beta$ line has FWHM~$\\gsim 2000$ km/s (hence, as standard Type 1 AGN), and Narrow Line Sy1s (NLS1s), with lower velocities (e.g., Goodrich, 1989). NLS1s are a minority, $\\simeq15$\\% of all the Sy1s, according to the optical spectroscopic classification of the SDSS general field (Williams Pogge \\& Mathur, 2002), the fraction depending on the AGN luminosity (with a peak at $M_{g'}\\sim -22$), and on the radio loudness (radio loud NLS1s account only for $\\sim7$\\% of the class, Komossa et al., 2006, but it is still debated if the NLS1s can be considered a peculiar radio-quiet sub-class among Sy1s, see e.g. Sulentic et al. 2007; Doi et al. 2007). NLS1s also show weak \\Oiii{} and strong \\Feii{} emission line (Osterbrock \\& Pogge, 1985), strong variability, and a softer than usual X--ray continuum (Boller Brandt \\& Fink, 1996; Grupe et al., 1999).  Grupe \\& Mathur (2004) found that NLS1s have, on average, lower \\mbh{} than expected from \\mbh{}--host galaxy relations such as \\mbh{}--$\\sigma_*$ (see Tremaine et al., 2002, and references therein), while BLS1 \\mbh{} are in fairly good agreement to the same relation. The estimated low values of \\mbh{} lead to an average Eddington ratio $L/L_{\\rm Edd}$ for the NLS1 population which is almost an order of magnitude larger than the average value of BLS1s ($L/L_{\\rm Edd} \\simeq 1$ to be compared to $\\simeq 0.1$, Grupe, 2004). Further evidence of low \\mbh{} in NLS1s comes from the observed rapid X--ray variability (see., e.g., Green, McHardy \\& Lehto 1993, and Hayashida 2000).  Such results were interpreted as indications of a peculiar role of NLS1s within the framework of the cosmic evolution of MBHs and of their hosts. In a MBH-galaxy co--evolution scenario, NLS1s are thought to be still on their way to reach the \\mbh{}-$\\sigma_*$ relation, i.e., their (comparatively) small MBHs are highly accreting in already formed bulges. Recently Botte et al. (2005) and Komossa and Xu (2007) came to the conclusion that NLS1s have indeed smaller masses and higher $L/L_{\\rm Edd}$ than BLS1, nevertheless they both do follow the $M-\\sigma_*$ relation for quiescent galaxies. The authors argued that the customarily used \\Oiii{} line is not a reliable surrogate for the stellar velocity dispersion $\\sigma_*$.  The Grupe and Mathur's results and interpretation have been recently confirmed and supported by several other groups, see, e.g., Zhou et al. (2006) and Ryan et al. (2007). Ryan et al. (2007) pointed out that IR-based mass measurements might be unreliable because of the extra IR contribute from the circum-nuclear star-forming regions in NLS1s. Notwithstanding, they suggested that this contamination can not significantly affect their data, and thus is insufficient to account for the MBH mass deficit. In the aforementioned papers, \\mbh{} was computed as \\begin{equation}\\label{eq_virial} M_{\\rm BH} = \\frac{R_{\\rm BLR} v_{\\rm BLR}^2}{G}, \\end{equation} where $R_{\\rm BLR}$ is the broad line region (BLR) scale radius, and $v_{\\rm BLR}$ the typical velocity of BLR clouds. $R_{\\rm BLR}$ is found by means of the reverberation mapping technique (Blandford \\& McKee, 1982), or by exploiting statistical $R_{\\rm BLR}$--luminosity relations (see Kaspi et al., 2000, 2005 and 2007); $v_{\\rm BLR}$ can be inferred from the H$\\beta$ width as \\begin{equation}\\label{eq_def_f} v_{\\rm BLR} = f \\cdot {\\rm FWHM}, \\end{equation} where the FWHM refers only to the broad component of the line, and $f$ is a fudge factor which depends upon the assumed BLR model. For an isotropic velocity distribution, as generally assumed, $f=\\sqrt{3}/2$.  Labita et al. (2006) and Decarli et al. (in preparation) found that in QSOs an isotropic BLR fails to reproduce the observed line widths and shapes, and a disc model should be preferred. A disc--like geometry for the BLR has been proposed by several authors in the past (e.g., Wills \\& Browne 1986; Vestergaard, Wilkes \\& Barthel 2000; Bian \\& Zhao, 2004). In this picture, the observed small FWHM of NLS1 broad lines are ascribed to a small viewing angle with respect to the disc axis, and no evolutionary difference is invoked whatsoever.  In this {\\it Letter}, we adopt the disc--like model for the BLR of Seyfert galaxies. We use the observed frequency of NLS1s to estimate their typical viewing angle, and then compute the appropriate geometrical factor $f$. Using eq.~\\ref{eq_virial}, we will show that the new estimates of \\mbh{} for NLS1s nicely agree with the standard \\mbh{}--$\\sigma_*$ relation. In turn, the accretion rate of the class is found to be similar to that of BLS1s. ",
        "conclusions": "In this {\\it Letter}, we have assessed the claimed peculiarity of Narrow Line Seyfert 1 galaxies within the framework of cosmic evolution of massive black holes, and their host bulges. Indeed, the optical properties of NLS1s, their X--ray fast variability and the faintness of their bulges can be accounted for if one admits lower black hole masses and higher accretion rates (in Eddington units) than standard Broad Line Seyfert 1 galaxies (BLS1s), placing NLS1s in an early evolutionary stage (Grupe \\& Mathur, 2004; Grupe, 2004; Botte et al., 2004; Zhou et al., 2006; Ryan et al., 2007). If this is true, by observing local NLS1s we can have hints of the infancy of the ubiquitous population of super--massive black holes.    We have explored an alternative explanation to the narrowness of H$\\beta$ lines in NLS1s, namely, pole--on orientation of a disc--like broad line region. If BLS1s and NLS1s differ only by the observation angle of the BLR disc, the frequency of NLS1s among the Sy1 class gives the limiting viewing angle of NLS1s. Then, assuming $H/R\\lsim 0.1$ for the disc, we computed corrected geometrical factors linking the observed FWHM to \\mbh{}, and found $f_{\\rm NLS1}\\gsim 2$ and $f_{\\rm BLS1}\\simeq 1$, in agreement with recent estimates given by Labita et al. (2006). The idea of a disc--like BLR is not new (e.g., Wills \\& Browne 1986; Vestergaard, Wilkes \\& Barthel 2000; Bian \\& Zhao, 2004), but for the first time, by re-calculating masses and Eddington ratios for a sample of Sy1s, we found that mass and luminosity functions are similar for NLS1s and BLS1s. In a sense, we can say that all Sy1s are normal, but some are more ``normal'' than others.  We note that, though NLS1s seem to lie in the same region of the \\mbh--$\\sigma_*$ plane, the adopted $\\sigma_*$ values can be largely over--estimated (Komossa \\& Xu, 2007; Mullaney \\& Ward, 2007), and then firm conclusions on the \\mbh--$\\sigma_*$ issue can not be drawn at this stage.   Can a simple orientation model, as the one we adopted here, explain the unique observed properties of NLS1s? NLS1s differ from standard Sy1s not just in the width of optical lines, but, more noticeably, in what are their X--ray properties, both spectral and temporal. The X--ray emission of NLS1s has been studied and discussed in great details by, among others, Boller et al. (1996), using {\\it ROSAT} data, and by Brandt, Mathur \\& Elvis (1997) using {\\it ASCA} data. NLS1s have generally both soft and hard X--ray spectra which are steeper than normal Sy1s, and show large amplitude, rapid variability. Boller et al. (1996) showed how different models, invoking extreme values of one or more of the followings: pole--on orientation, black hole mass, accretion rate, warm absorption, BLR thickness, all explain some aspects of the complex NLS1 soft X--ray phenomenology, but, still, all appear to have drawbacks.  If pole--on orientation has to be the main cause of the uniqueness of the X--ray features of NLS1s, then a necessary condition is that the hard power--law emission is not intrinsically isotropic, e.g., a thermal extended corona (as in Haardt \\& Maraschi, 1991; 1993) is not a viable option. Models in which the X--rays of type I radio quiet AGNs are funneled or beamed have been proposed by several authors (e.g., Madau, 1988; Henri \\& Petrucci, 1997; Malzac et al., 1998; Ghisellini, Haardt \\& Matt, 2004). For example, Ghisellini et al. (2004) showed that an aborted jet model, in which X--rays are produced by dissipation of kinetic energy of colliding blobs launched along the MBH rotation axis, can explain the steep and highly variable X--ray power law. The model, in its existing formulation, does not allow clear predictions of spectral and temporal features other than in the X-rays. To assess its relevance for NLS1s would require a much more detailed modeling. In particular, the peculiar \\Feii{} and \\Oiii{} properties must be accounted for.  The statistics of radio-loud NLS1s is low. In several works the existence of differences in the radio properties between NLS1s and BLS1s has been discussed (see, e.g. Komossa et al. 2006; Zhou et al. 2006; Sulentic et al. 2007; Doi et al. 2007). Doi et al. (2007) suggested that $\\sim$ 50 \\% of radio-loud NLS1s are likely associated with jets with high brightness temperatures, requiring Doppler boosting. This interpretation supports the pole--on orientation model (for a different point of view see Komossa et al. 2006).   Our results, if confirmed, indicate that a population of accreting, undermassive MBHs (with respect to the \\mbh{}--$\\sigma_*$ relation) has to be found yet. This may suggest that the \\mbh{}--$\\sigma_*$ relation was established long ago, during the MBH accretion episodes following the first major mergers of the host galaxies. Moreover, Komossa \\& Xu (2007) found that NLS1s do follow the \\mbh{}--$\\sigma_*$ relation of non--active galaxies, but still they have smaller \\mbh{} and larger $L/L_{\\rm Edd}$ than BLS1s. If this is the case, then $\\sigma_*$ of the host bulges of NLS1 needs to evolve accordingly in order to preserve the \\mbh{}--$\\sigma_*$ relation, or, alternatively, NLS1s are the low mass extension of BLS1s, and the NLS1 high $L/L_{\\rm Edd}$ is a short--lived phenomenon. We note here that the interpretation of Komossa \\& Xu (2007), as well as the one of Grupe \\& Mathur (2004), implies that \\mbh{} and $L/L_{\\rm Edd}$, in principle independent quantities, somehow conspire to produce comparable luminosities as observed in NLS1s and BLS1s. Applying our correction to the \\mbh{} as well as the one to the $\\sigma_*$ proposed by Komossa \\& Xu (2007), the NLS1s would be even off--setted towards higher masses with respect to the \\mbh{}--$\\sigma_*$ relation.  There are however two possible problems with the pole--on orientation model.  First, according to the orientation model, the polarization properties of broad emission lines should depend on the inclination angle, in the  sense that  nearly pole--on Sy1s should not be polarized. However, Smith et al. (2004) found polarized broad lines in few NLS1s, and traces of broad H$\\alpha$ polarization were also found by Goodrich (1989) in 6 out of 17 NLS1s. A second issue is discussed by Punsly (2007), who finds larger line broadening in face--on quasars, possibly due to large isotropic gas velocities or winds.  In conclusion, we found that orientation effects can account for the different optical properties of NLS1s compared to the more common BLS1s. The model is particularly appealing, as it naturally sets masses and accretion rates of NLS1 to fairly standard values. To validate this interpretation, orientation must be able to explain the extreme X--ray properties of NLS1. Jetted models for radio quiet AGNs may be promising in this, and we urge a detailed, critical comparison of such models with the bulk of NLS1 data."
    },
    "0801/0801.1035_arXiv.txt": {
        "abstract": "{In this paper we investigate the radio-MIR correlation at very low flux densities using extremely deep 1.4\\,GHz sub-arcsecond angular resolution MERLIN$+$VLA observations of a 8\\farms5$\\times$8\\farms5 field centred upon the Hubble Deep Field North, in conjunction with \\spitzer\\ 24\\,\\mum\\ data.  From these results the MIR-radio correlation is extended to the very faint ($\\sim$microJy) radio source population. Tentatively we detect a small deviation from the correlation at the faintest IR flux densities. We suggest that this small observed change in the gradient of the correlation is the result of a suppression of the MIR emission in faint star-forming galaxies.  This deviation potentially has significant implications for using either the MIR or non-thermal radio emission as a star-formation tracer of very low luminosity galaxies.} ",
        "introduction": "Since 1970s and 1980s studies of the radio and far-infrared (FIR) properties of galaxies have shown there to be a tight correlation between their emission in these two observing bands which extends over several orders of magnitude in luminosity \\citep{vanderkruit73,condon82}. The advent of the {\\it Infrared Astronomical Satellite} (IRAS) All Sky Survey in 1983 \\citep{neugebauer84,soifer87} enabled much larger systematic samples of galaxies to be studied at infrared wavelengths and hence further demonstrated the consistency and tightness of this correlation, albeit for relatively nearby sources \\citep{helou85,dejong85,condon86,condon91,yun01}. Following {\\it IRAS}, deep observations using the {\\it Infrared Space Observatory} ({\\it ISO}) allowed fainter and higher redshift galaxies to be observed at mid-infrared (MIR) wavelengths. These {\\it ISO} observations showed that the MIR emission from galaxies is loosely correlated with radio emission across a wide range of redshifts, tentatively extending to z$\\sim$4 \\citep{cohen00,elbaz02,garrett02}.  More recently the launch of the \\spitzer\\ Space Telescope in August 2003 has greatly increased the sensitivity of MIR observations and hence our ability to study the MIR-radio correlation. Early results, such as from the \\spitzer\\ First Look Survey (FLS), have confirmed that the MIR-radio correlation holds for relatively bright star-forming galaxies (S$_{20\\,{\\rm cm}}>115$\\,$\\mu$Jy) out to at least a redshift of 1 \\citep{appleton04}.  Radio and infrared emission from galaxies in both the nearby and distant Universe is thought to arise from processes related to star-formation, hence resulting in the correlation between these two observing bands. The infrared emission is produced from dust heated by photons from young stars and the radio emission predominately arises from synchrotron radiation produced by the acceleration of charged particles from supernovae explosions. It has however recently been suggested that at low flux density and luminosities there may be some deviation from the tight well-known radio-IR correlation seen for brighter galaxies \\citep{bell03,boyle07}.  \\begin{center} \\begin{figure*} \\setlength{\\unitlength}{.5in} \\begin{picture}(16,5) \\put(-0.5, 0){\\special{psfile=Beswick-fig1a.ps hscale=90 vscale=90}} \\put(6.5, 0){\\special{psfile=Beswick-fig1b.ps hscale=90 vscale=90}} \\end{picture} \\vskip -0.5cm \\caption{{\\it Left-hand panel:} Radio 1.4\\,GHz versus the MIR 24\\,\\mum\\ flux density of all 377 individual sources (small triangle), and median radio flux density logarithmically binned by their 24\\,\\mum\\ flux density (filled circles) within the 8\\farms5$\\times$8\\farms5 field. {\\it Right-hand panel:} Control plot of 1.4\\,GHz radio flux densities plotted against the 24\\,\\mum\\ source flux densities of the sample. The radio flux densities of the control sample have derived in exactly the same manner as the flux densities plotted in the left-hand, panel but have been measured at randomly assigned sky positions in the 8\\farms5$\\times$8\\farms5 radio image where no known sources at any wavelength are located.} \\label{F20vsF24} \\end{figure*} \\end{center}  \\cite{bell03} argue that while the IR emission from luminous galaxies will trace the majority of the star-formation in these sources, in low luminosity galaxies the IR emission will be less luminous than expected considering the rate of star-formation within the source (i.e. the IR emission will not fully trace the star-formation). In this scenario the reduced efficiency of IR production relative to the source star-formation rate (SFR) would be the result of inherently lower dust opacities in lower luminosity sources and consequently less efficient reprocessing of UV photons from hot young stars into IR emission. The simple consequence of this is that at lower luminosities the near linear radio-IR correlation L$_{\\rm radio}\\propto$L$_{\\rm IR}^{\\!\\gamma}$, with $\\gamma >1$ \\citep[e.g][]{cox88,price92} will be deviated from. {\\it Of course such an assertion is dependent upon the radio emission providing a reliable tracer of star-formation at low luminosities which may be equally invalid.}  Recently \\cite{boyle07} have presented a statistical analysis of Australia Telescope Compact Array (ATCA) 20\\,cm observations of the 24\\,\\mum\\ sources within two regions (the {\\it Chandra} Deep Field South (CDFS) and the European Large Area {\\it ISO} Survey S1 (ELAIS)) of the \\spitzer\\ Wide Field Survey (SWIRE). In this work \\cite{boyle07} have co-added sensitive (rms$\\sim$30\\,$\\mu$Jy) radio data at the locations of several thousand 24\\,\\mum\\ sources. Using this method they have statistically detected the microJy radio counterparts of faint 24\\,\\mum\\ sources. At low flux densities (S$_{\\rm 24\\, \\mu m }=100\\,\\mu$Jy) they confirm the IR-radio correlation but find it to have a lower coefficient (S$_{\\rm 1.4\\,GHz}$\\,=\\,0.039\\,S$_{\\rm 24\\,\\mu m}$) than had previously been reported at higher flux densities. This coefficient is significantly different from results previously derived from detections of individual objects (e.g. \\citealt{appleton04}) and is speculated by \\cite{boyle07} to be the result of a change in the slope of the radio-IR correlation at low flux densities.  In this paper, we utilise very deep, high resolution 20\\,cm observations of the Hubble Deep Field North and surrounding area made using MERLIN and the VLA \\citep{muxlow05} in combination with publicly available 24\\,\\mum\\ \\spitzer\\ source catalogues from GOODS to study the MIR-Radio correlation for microjansky radio sources. This study extends the flux density limits of the radio-IR correlation by more than an order of magnitude for individual sources and overlaps the flux density regime studied using statistical stacking methods by \\cite{boyle07}. Additionally we employ statistical stacking methods, similar to those used by \\cite{boyle07}, to extend the correlation further to still lower flux densities.    We adopt H$_0=$75\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\rm m}=0.3$ and $\\Omega_{\\rm \\Lambda}$=0.7 throughout this paper. ",
        "conclusions": "Using one of the deepest high-resolution 1.4\\,GHz observations made to date, in conjunction with deep 24\\,\\mum\\ \\spitzer\\ source catalogues from GOODS, we have investigated the microJy radio counterparts of faint MIR sources. These observations confirm that the microJy radio source population follow the MIR-radio correlation and extend this correlation by several orders of magnitude to very low flux densities and luminosities, and out to moderate redshifts. This extension of the MIR-radio correlation confirms that the majority of these extremely faint radio and 24\\,\\mum\\ sources are predominantly powered by star-formation with little AGN contamination.  Statistically stacking the radio emission from many tens of faint 24\\,\\mum\\ sources has been used to characterise the size and nature of the radio emission from very faint IR galaxies well below the nominal radio sensitivity of these data. Using these methods the MIR-radio correlation has been further extended and a tentative deviation in this correlation at very low 24\\,\\mum\\ flux densities has been identified.  \\subsection*"
    },
    "0801/0801.0343_arXiv.txt": {
        "abstract": "Galaxy disks evolve through angular momentum transfers between sub-components, like gas, stars, or dark matter halos, through non axi-symmetric instabilities. The speed of this evolution is boosted in presence of a large fraction of cold and dissipative gas component. When the visible matter dominates over the whole disk, angular momentum is exchanged between gas and stars only. The gas is driven towards the center by bars, stalled transiently in resonance rings, and driven further by embedded bars, which it contributes to destroy. From a small-scale molecular torus, the gas can then inflow from viscous torques, dynamical friction, or $m=1$ perturbations. In the weakened bar phases, multiple-speed spiral patterns can develop and help the galaxy to accrete external gas flowing from cosmic filaments.  The various phases of secular evolution are illustrated by numerical simulations. ",
        "introduction": "\\subsection{Bar formation and evolution}  Non-axisymmetries and in particular bars are the motor of secular evolution, in transferring the angular momentum in galaxies.  In the 80' and 90', numerical simulations established how bars formed, in pure stellar disks, without any dark matter haloes, or embedded in rigid halo components: the angular momentum was exchanged within the disk, the outer parts gaining the momentum from the inner parts. This exchange occured mainly at bar formation, the bar pattern speed slowing down slightly, while the stellar orbits in the bar were more and more elongated, corresponding to a lower precessing rate. Once formed, the bar was then robust (e.g. Sellwood 1981).  Already, some dynamical heating was observed to produce some feedback in the bar strength evolution: more unstable stellar disks develop a bar sooner, but then end up with a weaker bar than more stable disks (Combes et al 1990).  \\subsection{Bars with gas disks}  The introduction of gas in disks completely changes the stability, due to dissipation. The angular momentum can be exchanged between gas and stars, and from inner to outer parts, again with a negligible halo component. While gas flows towards the center, the bar is destroyed and gives rise to a triaxial bulge (Friedli \\& Benz 1993).  All this evolution can be traced back to the bar gravity torques acting on the gas. The torques are proportional to the phase shift between the gas and the potential wells due mainly to the stars in the bar. Dissipation produces these phase shifts, which are conspicuous through the characteristic leading dust lanes in barred galaxies. The gas then loses angular momentum and mass concentrates towards the center.  The amplitude of the phenomenon has been quantified from observations of barred galaxies. The gravity torques can be computed from the stellar potential deduced from the red image of the galaxy, and the effective angular momentum exchange computed from the gas distribution, obtained through H$\\alpha$, HI or CO emission images. In strongly barred galaxies, the gas inside corotation inflows to the center in one or two rotations. Even in weakly barred galaxies, the gravity torques are efficiently producing gas flows, and fueling the nucleus, as for instance in NGC 6574 (Lindt-Krieg et al 2007) or NGC 3147 (Casasola et al 2008).  \\subsection{Formation of rings}  The secular evolution as described above, the gas inflow driven by bars, finds its confirmation in the frequent observation of resonant rings, where gas is piling up, and form stars in bright knots. The gravity torques change sign at each resonance, and cancel in the rings where the gas disctribution is symmetrically distributed with respect to the stars.  Galaxies often possess multiple-rings, corresponding to the various Lindblad resonances, ILR, UHR, OLR (Buta \\& Combes 1996). With only one bar, the gas stops its inflow at the ILR, in the nuclear ring.  The decoupling of a nuclear bar inside the nuclear ring triggers the AGN fueling (cf NGC 2782, Fig 1).  \\begin{figure} \\begin{center} \\includegraphics[angle=-90,width=12cm]{combes-fig1.ps} \\end{center} \\caption{ Gray scale plot of the stellar component ({\\it left}), the gas component ({\\it middle}), and an expanded version of the gas distribution  ({\\it right}), in a self-consistent  simulation of a double bar formation, meant to reproduce the galaxy NGC 2782 (the axes are in kpc). It can be seen that the gas, which was predominantly in the more external ring, corresponding to the ILR of the primary bar, at T=690 Myr, is falling progressively inward, and is found inside the ILR at T=750 Myr. From Hunt et al (2008).} \\label{n2782} \\end{figure}   \\vspace{-0.3cm} ",
        "conclusions": "\\vspace{-0.2cm}  Secular evolution is important for bars, it controls their strength, their pattern speed, or their vertical thickness through peanut formation. The dominant angular momentum transfer from the disk occurs with the DM halo or with the outer disk, according to the halo to disk mass ratio, or between the stellar and gaseous components, according to the dissipation character of the gas.  Bars are weakened by central mass concentrations and/or gas inflows, driven by bars themselves, implying self-regulation. The gas fraction able to destroy bars depends on the dark matter halo to disk mass ratio. Lopsidedness ($m=1$ global mode) can also weaken the bar. Bars dissolve into lenses, especially in early-type galaxies.    \\vspace{-0.3cm}"
    },
    "0801/0801.2164_arXiv.txt": {
        "abstract": "\\noindent Recent data support the idea that the filaments observed in H$\\alpha$ emission near the centres of some galaxy clusters were shaped by bulk flows within their intracluster media. We present numerical simulations of evaporated clump material interacting with impinging winds to investigate this possibility. In each simulation, a clump falls due to gravity while the drag of a wind retards the fall of evaporated material leading to elongation of the tail. However, we find that long filaments can only form if the outflowing wind velocity is sufficiently large, $\\sim 10^{8}\\,{\\rm cm\\,s^{-1}}$. Otherwise, the tail material sinks almost as quickly as the cloud. For reasonable values of parameters, the morphological structure of a tail is qualitatively similar to those observed in clusters. Under certain conditions, the kinematics of the tail resemble those reported in Hatch et al.(2006). A comparison of the observations with the numerical results indicates that the filaments are likely to be a few tens of Myrs old. We also present arguments which suggest that the momentum transfer, from an outflowing wind, in the formation of these filaments is probably significant. As a result, tail formation could play a role in dissipating some of the energy injected by a central AGN close to the cluster centre where it is needed most. The trapping of energy by the cold gas may provide an additional feedback mechanism that helps to regulate the heating of the central regions of galaxy clusters and couple the AGN to the ICM. ",
        "introduction": "Optical emission-line nebulae commonly surround massive galaxies in the centres of X-ray bright, cool cluster cores \\citep[][]{crawf}. The origin of the H$\\alpha$ filaments has been attributed to a variety of processes including: condensation from an intracluster medium (ICM) evolving as a cooling flow \\citep[][]{fab84, heckman, don91}; accretion of clouds captured in galaxy mergers \\citep[][]{braine}; expulsion from the central galaxy \\citep[][]{burb}.  NGC 1275 at the centre of the Perseus cluster contains the best studied example of such a nebula \\citep[][]{hatch}. Its filaments are typically 50-100 pc thick and up to 30 kpc long, and the majority of them are radial. \\cite{hatch} presented kinematic data that rules out dynamical models of purely infalling filaments. The observed kinematic properties provide strong evidence that the filaments are not in gravitational free fall, because, if they were, their velocities would rise sharply towards the centre of the nebula \\citep[][]{heckman}. The most conclusive evidence lies in the velocity structures of the northern and northwestern filaments. The lower half of the northern filament is redshifted with respect to the galaxy, whilst the upper section is blueshifted. Thus, the upper part of the filament is flowing away from the galaxy while the lower part is flowing into the galaxy.  The filaments may follow streamlines of the flow of more tenuous material around them. In the Perseus cluster, some filaments appear to have been shaped by the wakes of the buoyant bubbles. Given that the data suggest that, at least, some parts of the filaments are outflowing, their origin may lie within the galaxy. NGC 1275 contains a large reservoir of cold molecular gas that could fuel them.  The filaments have morphologies similar to those of laminar jets, and their relatively smooth structures may place constraints on turbulent motions in the ICM. Alternatively, an ordered, amplified magnetic field trailing behind a buoyant bubble interacting with filaments may prevent the destruction of the filaments by turbulent motions \\citep[][]{rusz07}.  The optical nebula around NGC 1275 emits $4.1 \\times 10^{42}{\\rm erg\\,s^{-1}}$ in H$\\alpha$ and N[II]. The nature of the power source remains unclear. Various excitation mechanisms for these lines have been proposed. For instance, although ionisation by the central AGN may be important for the inner regions, it cannot be the dominant source of power for the extended nebula because the H$\\alpha$ intensity does not decrease with distance from the nucleus \\citep[][]{john88}. The nebula may be excited by stellar UV, but there is no spatial correlation between the filaments and the stellar clusters. Excitation by X-rays from the ICM seems unlikely, as the ICM may be as much as a hundred times less luminous in UV than in the X-rays \\citep[][]{fab03b}. Heat conduction from the ICM to the colder filaments has also been proposed \\citep[][]{don00}, but it might also lead to the evaporation of the filaments on too short a timescale \\citep[][]{nb04}. Shocks and turbulent mixing layers might play a role \\citep[][]{crawf92}. Cosmic rays, preferentially diffusing along the magnetic field lines trailing behind rising bubbles, could possibly drive the excitation in those filaments that are located in bubble wakes \\citep[][]{rusz07}. The same magnetic fields lines could also prevent the filaments from evaporating due to thermal conduction.  The same mechanisms that power the emission could also have a profound effect on the morphology of the emitting region. It is impossible to take all of these processes into account. Consequently, we will investigate the effect of gravity and an impinging wind that strips material from a cloud on the morphology and kinematics of the resulting filament.  The main aim of the work reported here is the calculation of the density and velocity structures of material evaporated from clumps embedded in winds from the central galaxies of galaxy clusters and the comparison of model results with observations. Section 2 contains some preliminary considerations, while the model assumptions and numerical approach are treated in section 3.  The simulation results appear in section 4. In section 5 we investigate the development of tails in realistic environments by including the appropriate density and gravity for a galaxy cluster. Section 6 concerns momentum tranfer between clumps and the winds surrounding them. Section 7 discusses the possible trends in optical emission between clusters, and section 8 is a summary. ",
        "conclusions": "This work is intended mainly as a preliminary study into the processes that shape filamentary structures in the ICM, and not an exhaustive study. It is likely that there are many other processes such as magnetic fields, thermal conduction and viscosity that all play a role in the dynamics that produce the filamentary kinematics and morphology. However, we have concentrated on the basic hydrodynamic system and have found several interesting mechanisms that may go some way to explaining the properties of the observed filaments.  The first ensemble of simulations show that the long, relatively straight filaments at the centre of the Perseus cluster may be formed by the interaction of a wind with clumps of cold material. Flows of low Mach number winds generate structure in the filaments which may be comparable with observations. These results also suggest that the amplitude of the velocity fluctuations along the length of the filament grow with increasing density contrast between the cold material and the ambient wind. The morphology of the observed filaments suggests that the density contrast is large ($\\sim 10^{4}$). Interestingly, the fluctuations evolve with time providing a possible diagnostic tool for estimating the ages of filaments from observations. in this set of simulations, the best comparison with observations suggests filament ages of $\\sim 40$ Myrs in the Perseus cluster.  In a more realistic environment with spatially varying gravity, based on the Perseus cluster, filaments only form if the wind velocity is sufficiently large. If there is no significant wind, then the optical emission may be more amorphous, or more likely, filamentary on much shorter length scales. The morphology and kinematics suggests that a density contrast of $\\eta \\sim 10^{3}$ may be more compatible with the observations, while the best agreement in the kinematic data occurs for $\\eta \\sim 10^{4}$. As a compromise it is possible that the real conditions lay somewhere between these two extremes. Alternatively, other physical processes could be important. It should also be noted that these were 2-d simulations, and a full 3-d study would really be required to find the model conditions that best match reality. Regardless of this, the results do point to a relatively narrow range of parameters providing possible information about the state of the ICM, and the age of the filaments. These results suggest that the cold material is of the order of $10^{4}$ times denser than the ambient ICM, there is an outflowing wind of velocity of $\\sim 10^{3}\\,{\\rm km\\,s^{-1}}$, and that the filaments are approximately a few 10's of Myrs old.  The masses, and densities, of the filaments created by these simulations seem to be low compared with values presented in the previous section. The simulations produce filaments of $10^{6}-10^{7}\\,{\\rm M_\\odot}$, while the crude estimates above suggested that $10^{8}\\,{\\rm M_\\odot}$ filaments might be more likely. It should be noted though, that the quantity of cold gas, in the simulated tail, is more than capable of producing the observed optical luminosity. The simulations also show a lot of cold, dense gas in the vicinity of the cloud itself.  We also highlight the possibility that momentum transfer between the ICM and cold clouds could be a dynamically important process in galaxy clusters. This process could effectively couple the energy and momentum of an outflowing wind to the cold material, thereby dissipating some of the energy injected by a central AGN. The presence of cold gas may therefore provide an additional feedback mechanism that traps energy in the central regions of clusters where it is most needed."
    },
    "0801/0801.2022_arXiv.txt": {
        "abstract": "The understanding and modeling of the structure and evolution of stars is based on statistical physics as well as on hydrodynamics. Today, a precise identification and proper description of the physical processes at work in stellar interiors are still lacking (one key point being that of transport processes) while the comparison of real stars to model predictions, which implies conversions from the theoretical space to the observational one, suffers from uncertainties in model atmospheres. That results in uncertainties on the prediction of stellar properties needed for galactic studies or cosmology (as stellar ages and masses). In the next decade, progress is expected from the theoretical, experimental and observational sides. I illustrate some of the problems we are faced with when modeling stars and the possible tracks towards their solutions. I discuss how future observational ground-based or spatial programs (in particular those dedicated to micro-arc-second astrometry, asteroseismology and interferometry) will provide precise determinations of the stellar parameters and contribute to a better knowledge of stellar interiors and atmospheres in a wide range of stellar masses, chemical compositions and evolution stages. ",
        "introduction": "Major goals of stellar structure and evolution studies are (i) to characterize and describe the physics of matter in the extreme conditions encountered in stars and (ii) to determine stellar properties (like age and mass) that trace the history and evolution of galaxies and constrain cosmological models. To achieve these goals, we rely on numerical stellar models based on input physics that integrate the results of recent theoretical studies, numerical simulations and laboratory experiments. The models inputs and outputs are chosen and/or validated by comparison with accurate astronomical observations.  Numerical 2 and 3{\\small D} hydrodynamical simulations of limited regions of stellar interiors and atmospheres are now under reach of computers. They provide valuable constraints and data for current standard (1{\\small D}) stellar models: abundances, convection, rotationally induced instabilities and mixing, magnetic fields, etc. \\citep[see e.g.][for reviews]{2005ARA&A..43..481A,2007arXiv0708.1499T,2007IAUS..239..517Z}. In parallel, the physics of stellar plasmas is studied in the laboratory with (i) fluid experiments \\citep[study of turbulence in rotating, magnetic fluids, etc., see e.g.][]{1999A&A...347..734R}, (ii) particle accelerators (nuclear reaction cross sections, etc.) and, (iii) the so-called high energy-density facilities (based on high power lasers or z-pinches) which aim at exploring the high temperature and high density regimes found in stars, brown dwarfs and giant planets to get  information on the equation of state ({\\small EOS}), opacities or thermonuclear reactions \\citep[see][for a review]{2006RvMP...78..755R}.  Modern ground-based and spatial telescopes equipped with high quality instrumentation are in use or under development ({\\small VLT-VLTI, JWST,} etc.). They provide very accurate data which, after treatment, give access to stellar global parameters (luminosity, radius, mass, effective temperature $T_{\\rm eff}$, gravity $\\log g$, abundances, etc.). On the other hand, seismic data (such as oscillation frequencies or amplitudes) are being obtained in velocity from the ground \\citep[see, e.g.][]{2007CoAst.150..106B} and in photometry by the space missions {\\small MOST} \\citep{2003PASP..115.1023W} and {\\small CoRoT} \\citep{2006corm.book...39M}. In the next decade, valuable observational data are expected. For instance {\\small GAIA} \\citep{ESA2000,2001A&A...369..339P}, to be launched in 2011, will make astrometric measurements, at the micro-arc second level together with photometric and spectroscopic observations of a huge number of stars covering the whole range of stellar masses, compositions and evolution stages while the Kepler mission, to be launched in 2009, will provide the opportunity to make asteroseismic observations on a wide range of stars \\citep{2007CoAst.150..350C}.  In the following, I discuss the different aspects of stellar modeling, the problems encountered and the perspectives. ",
        "conclusions": ""
    },
    "0801/0801.1217_arXiv.txt": {
        "abstract": "We selected a sample of a dozen blazars which are the prime candidates for simultaneous multi-wavelength observing campaigns in their outburst phase. We searched for optical outbursts, intra-day variability and short term variability in these blazars. We carried out optical photometric monitoring of nine of these blazars in 13 observing nights during our observing run October 27, 2006 $-$ March 20, 2007 by using the 1.02 meter optical telescope equipped with CCD detector and BVRI Johnson broad band filters at Yunnan astronomical observatory, Kunming, China. From our observations, our data favor the hypothesis that three blazars: AO 0235$+$164, S5 0716$+$714 and 3C 279 were in the outburst state; one blazar: 3C 454.3 was in the post outburst state; three blazars: S2 0109$+$224, PKS 0735$+$178 and OJ 287 were in the pre/post outburst state; one blazar: ON 231 was in the low-state; and the state of one blazar: 1ES 2344$+$514 was not known because there is not much optical data available for the blazar to compare with our observations. We observed densely sampled 1534 image frames of these nine blazars. Out of three nights of observations of AO 0235$+$164, intra-day variability was detected in two nights. Out of five nights of observations of S5 0716$+$714, intra-day variability was detected in two nights. In one night of observations of PKS 0735$+$178, intra-day variability was detected. Out of six nights of observations of 3C 454.3, intra-day variability was detected in three nights. No intra-day variability was detected in S2 0109$+$224, OJ 287, ON 231, 3C 279 and 1ES 2344$+$514 in their 1, 4, 1, 2 and 1 nights of observations respectively. AO 0235$+$164, S5 0716$+$714, OJ 287, 3C 279 and 3C 454.3 were observed in more than one night and short term variations in all these blazars were also noticed. From our observations and the available data, we found that the predicted optical outburst with the time interval of $\\sim$ 8 years in AO 0235$+$164 and $\\sim$ 3 years in S5 0716$+$714 have possibly occurred. ",
        "introduction": "Blazars represent a small subset of the most enigmatic class of radio-loud active galactic nuclei (AGN), exhibiting strong variability at all wavelengths of the whole electromagnetic (EM) spectrum, strong polarization from radio to optical wavelengths, and usually core dominated radio structures. The radiation of blazars at all wavelengths is predominantly nonthermal. This class includes BL Lacertae (BL Lac) objects and flat-spectrum radio quasars (FSRQs). BL Lacs show largely featureless optical continuum. In a unified model of the radio-loud AGN based on the angle between the line of sight and the emitted jet from the source, blazars jet make angle of $\\sim$ 10$^{\\circ}$ from the line of sight (Urry \\& Padovani 1995). The radiation emitted by the plasma, with bulk relativistic motion in the jet oriented at small viewing angles, is affected by relativistic beaming, which in turn implies a shortening of time scales by a factor $\\delta^{-1}$, where $\\delta$ is the Doppler factor.  From observations of blazars, it is known that they vary on the diverse time scales. Variability time scales of blazars can be broadly divided into 3 classes viz. intra-day variability (IDV) or micro-variability, short term outbursts and long term trends. Significant variations in flux of a few tenths of magnitude over the course of a day or less is often known as IDV (Wagner \\& Witzel 1995). Short term outburst and long term trends can have time scales range from few weeks to several months and several months to years, respectively. In last about two decades, variability of blazars in radio to optical bands on diverse time scales have been reported in a large number of papers (e.g. Miller et al. 1989; Courvoisier, et al. 1995; Heidt \\& Wagner 1996; Takalo et al. 1996; Sillanp$\\ddot{a}\\ddot{a}$ et al. 1996a, 1996b; Bai, et al. 1998, 1999; Fan et al. 1998, 2002, 2007; Xie, et al. 2002a; Gupta et al. 2004; Ciprini et al. 2003, 2007 and references therein).  We selected a sample of a dozen blazars which are prominent candidates for simultaneous multi-wavelength observing campaigns in their outburst phase. The motivation of the present work was to observe these blazars in search for IDV, short term variability and also find out if there is any one being in the outburst state. Blazar emission mechanism in the outburst state and detected IDV is strongly supported by the jet based models of radio-loud AGN. In general, blazar emission in the outburst state is nonthermal Doppler boosted emission from jets (Blandford \\& Rees 1978; Marscher \\& Gear 1985; Marscher et al. 1992, Hughes et al. 1992). There are other models of AGN that can explain the IDV in any type of AGN are optical flares, disturbances or hot spots on the accretion disk surrounding the black hole of the AGN (Mangalam \\& Wiita 1993 and references therein). Models based on the instabilities on accretion disc are mainly supported blazars IDV when the blazar is in the low-state. When a blazar is in the low-state, any contribution from the jets if at all present, is very weak. Recently, it is noticed that, in the low luminosity AGN, accretion disk is radiatively inefficient (Chiaberge et al. 2006; Capetti et al. 2007). So, there will be an alternative way to explain the IDV in the low-state of blazars, in which a weak jet emission will be responsible for the IDV. With this motivation we recently carried out optical photometric observations of the nine blazars: S2 0109$+$224, AO 0235$+$164, S5 0716$+$714, PKS 0735$+$178, OJ 287, ON 231, 3C 279 and 1ES 2344$+$514 in R passband and 3C 454.3 in V and R passbands.  The paper is arranged as follows: section 2 describes observations and data analysis method, in section 3 we mentioned our results, discussion and conclusion of the present work is reported in section 4. ",
        "conclusions": ""
    },
    "0801/0801.1398_arXiv.txt": {
        "abstract": "We study a 5-dimensional $f({\\cal R})$ brane gravity within the framework of scalar-tensor type theories. We show that such a model predicts, for a certain choice of $f({\\cal R})$ and a spatially flat universe, an exponential potential, leading to an accelerated expanding universe driven solely by the curvature of the bulk space. This result is consistent with the observational data in the cosmological scale. ",
        "introduction": "The idea that our world might be a brane embedded in a higher dimensional space-time (the bulk) \\cite{1} has been in the mainstream of cosmological investigations  in the past few years \\cite{3,4}. This approach differs from the usual Kaluza-Klein idea in that the size of the extra dimensions can be large. The concept of large extra dimensions is discussed phenomenologically in \\cite{5}. An important ingredient of the brane world scenario is that the matter is confined to the brane and the only communication between the brane and bulk is through gravitational interaction or some other dilatonic matter. In general, the matter on the brane leads to a cosmological evolution which is different from the usual evolution governed by the Friedmann equation, that is, in brane cosmology the Hubble parameter on the brane is proportional to the square of energy density \\cite{3,4}. This proportionality is a result of the application of the Israel matching condition which is basically a relation between the extrinsic curvature and the energy-momentum tensor representing matter fields on the brane.  Although in brane theories matter fields live on the brane, the possibility of the presence of matter in the form of a scalar field in the bulk has also been investigated in several works. One of the first motivations to introduce a bulk scalar field was to stabilize \\cite{7} the distance between the two branes in the context of the first model introduced by Randall and Sundrum \\cite{1}. A second motivation was the possibility of the resolution of the famous cosmological constant problem \\cite{8}. Several works have studied, in particular, the impact of the presence of a scalar field in the bulk on the cosmological evolution on the brane, without trying to solve the full system of equations in the bulk \\cite{9,10}. In \\cite{12}, the authors have addressed some of the solutions for these equations and studied the corresponding brane evolution. The purpose of the present study is to employ modified gravity \\cite{parry} in the Einstein frame to explain the origin of such a self interacting scalar potential.  An interesting observation made a few years ago was that the expansion of our universe is currently undergoing a period of acceleration which is directly measured from the light curves of several hundred type Ia supernovae \\cite{13} and independently from observations of the cosmic microwave background (CMB) by the WMAP satellite \\cite{15} and other CMB experiments \\cite{16}. However, the mechanism responsible for this acceleration is not well understood and many authors introduce a mysterious cosmic fluid, the so called dark energy, to explain this effect \\cite{17}. Recently, it has been shown that such an accelerated expansion could be the result of a modification to the Einstein-Hilbert action \\cite{18} in the framework of DGP brane cosmology. In the present work we study the general form of the Einstein-Hilbert action for any function of the Ricci scalar, $ f({\\cal R})$, in  5 dimensions. This is done in the framework of a scalar-tensor type theory \\cite{20} where a scalar field is minimally coupled to gravity with a self-interacting potential. In this formulation we obtain explicit solutions using conformal transformations, a technique employed in the case of an empty bulk with a cosmological constant \\cite{21} or a bulk with a scalar field, similar to the present work, but with an exponential potential. We present explicit solutions for a particular choice of $ f({\\cal R})$ which predict a similar exponential potential.  The organization of the manuscript is as follows: in section 2 we briefly review the scalar-tensor formulation in 5-dimensions and write the full system of equations. In section 3 we consider the cosmological equations for $f({\\cal R})$ gravity which, in the Einstein frame, correspond to a self interacting scalar field with a certain potential. Finally, we study the cosmological evolution on the brane for ${\\cal R}^{m}$ gravity which predicts a power law acceleration in section 4. Conclusions are drawn in the last section. ",
        "conclusions": "In this manuscript we have obtained explicit solutions in a brane world scenario where an arbitrary function of the Ricci scalar is taken as the bulk Lagrangian. Using a conformal transformation, the action is converted to that of a scalar-tensor type theory with a scalar field. We have shown that with a suitable choice for the function $f({\\cal R})$ and brane tension $\\lambda$, an accelerated expanding universe emerges. The source of this acceleration is not related to an exotic matter but to a scalar field whose origin can be traced back to geometry of the brane and, specifically, to the curvature scalar $\\cal R$ and depends on two free parameters, namely $\\alpha$ and $\\lambda_c$. Hence, an accelerating universe driven by curvature would certainly seem to be a possibility."
    },
    "0801/0801.1167_arXiv.txt": {
        "abstract": "With the goal to study the physical and chemical evolution of ices in solar-mass systems, a spectral survey is conducted of a sample of 41 low luminosity YSOs ($L\\sim 0.1-10~{\\rm L}_{\\odot}$) using 5--38 \\mum\\ Spitzer Space Telescope and 3--4 \\mum\\ ground-based spectra. The sample is complemented with previously published Spitzer spectra of background stars and with ISO spectra of well studied massive YSOs ($L\\sim 10^5~{\\rm L}_{\\odot}$).  This paper focuses on the origin of the prominent absorption features in the 5-8 \\mum\\ spectral region.  The long-known 6.0 and 6.85 \\mum\\ bands are detected toward all sources, with the Class 0-type low mass YSOs showing the deepest bands ever observed.  In almost all sources the 6.0 \\mum\\ band is deeper, by up to a factor of 3, than expected from the bending mode of pure solid H$_2$O, based on the optical depths of the 3.0 \\mum\\ stretching and 13 \\mum\\ libration modes.  The depth and shape variations of the remaining 5--7 \\mum\\ absorption indicate that it consists of 5 independent components, which, by comparison to laboratory studies, must be from at least 8 different carriers. Together with information from the 3-4 \\mum\\ spectra and the additionally detected weak 7.25, 7.40, 9.0, and 9.7 \\mum\\ features it is argued that overlapping bands of simple species are responsible for much of the absorption in the 5-7 \\mum\\ region, at abundances of 1-30\\% for CH$_3$OH, 3-8\\% for NH$_3$, 1-5\\% for HCOOH, $\\sim$6\\% for H$_2$CO, and $\\sim$0.3\\% for HCOO$^-$ with respect to solid H$_2$O.  The 6.85 \\mum\\ band likely consists of one or two carriers, of which one is less volatile than H$_2$O because its abundance relative to H$_2$O is enhanced at lower H$_2$O/$\\tau_{9.7}$ ratios. It does not survive in the diffuse interstellar medium (ISM), however.  The similarity of the 6.85 \\mum\\ bands for YSOs and background stars indicates that its carrier(s) must be formed early in the molecular cloud evolution. If an NH$_4^+$ salt is the carrier its abundance with respect to solid H$_2$O is typically 7\\%, and low temperature acid-base chemistry or cosmic ray induced reactions must have been involved in its formation.  Possible origins are discussed for the carrier of an enigmatic, very broad absorption between 5 and 8 \\mum.  It shows large depth variations toward both low- and high-mass YSOs.  Weak evidence is found that it correlates with temperature tracers. Finally, all the phenomena observed for ices toward massive YSOs are also observed toward low mass YSOs, indicating that processing of the ices by internal ultraviolet radiation fields is a minor factor in the early chemical evolution of the ices. ",
        "introduction": "~\\label{sec:intro}  The infrared spectra of protostars and obscured background stars show prominent absorption features at 3.0, 4.25, 4.7, 6.0, 6.85, 9.7, and 15 \\mum\\ along with a suite of weaker features (e.g. \\citealt{dhe96}, \\citealt{whi96}, \\citealt{ger99}, \\citealt{sch99}, \\citealt{gib00}, \\citealt{kea01b}, \\citealt{pon03a}, \\citealt{gib04}, \\citealt{kne05}, \\citealt{whi07}; see \\citealt{boo04} for a complete list of features). These are attributed to absorption in the vibrational modes of molecules in ices, except for the 9.7 \\mum\\ band which is mostly due to silicates. At the low temperatures ($T\\leq$20 K) of dense clouds and circum-protostellar environments atoms and molecules freeze out rapidly on dust grains.  Grain surface chemistry efficiently forms new, simple species, such as H$_2$O and H$_2$CO (e.g. \\citealt{tie82}).  Complex species (e.g.  polyoxymethylene [`POM', -(CH$_2$-O)$_{\\rm n}$-], and hexamethylenetetramine ['HMT', C$_6$H$_{12}$N$_4$]) can be formed through the impact of energetic photons or cosmic rays on the ices, as many laboratory studies have shown (e.g.  \\citealt{sch93}, \\citealt{ber95}, \\citealt{gre95}, \\citealt{ger96}, \\citealt{moo98}, \\citealt{pal00}, \\citealt{mun04}). Thus far, only the simple species H$_2$O, CO, CO$_2$, CH$_4$, CH$_3$OH, NH$_3$ and the $^{13}$CO and $^{13}{\\rm CO}_2$ isotopes have been positively identified in the ices toward both low and high mass protostars as well as extincted background stars whose lines of sight trace quiescent dense cloud material.  Reasonable evidence exists for the presence of HCOOH, OCS and the ions NH$_4^+$, OCN$^-$, and HCOO$^-$ as well, although their existence is sometimes debated because of inaccurate fits with laboratory spectra or the lack of multiple bands for independent confirmation.  The identification of the ions in the ices has been particularly controversial. The first ion claimed was OCN$^-$ \\citep{gri87} and was thought to be produced by heavy energetic processing of CO:NH$_3$ ices. Later it was realized that acid-base chemistry in a HNCO:NH$_3$ ice could yield the same products \\citep{nov01}.  Acid-base reactions are very efficient and occur at temperatures as low as 10 K \\citep{rau04b}. The created ions are less volatile than neutral species and would be able to form an interstellar salt after the other species have sublimated.  The presence of complex species formed by ultraviolet (UV) photon and cosmic ray processing of the simple ices is at least as controversial. Many claims were made \\citep{gre95,gib02}, but the observational evidence is not firm.  It is crucial to determine the complexity of the ices in the circumstellar environment of low mass Young Stellar Objects (YSOs), as they may be delivered directly to comets and they may ultimately serve as the source of volatiles in planets.  At sublimation fronts in the warm inner regions of disks they are the starting point of a complex gas phase chemistry.  The relative abundances of N-, C-, and O-bearing species coming from the ices determine directly the type of chemistry in hot cores (or `hot corinos' for low mass objects; e.g.  \\citealt{caz03}).  As a step toward better characterizing the molecular content of icy grain mantles, this work presents spectra of a large sample of low mass protostars and background stars obtained with the InfraRed Spectrometer (IRS; \\citealt{hou04}) at the Spitzer Space Telescope \\citep{wer04} in the 5-35 \\mum\\ wavelength range.  These data are complemented with spectra at 2-4 \\mum\\ obtained with ground-based facilities. Such full coverage over the near and mid-infrared wavelength ranges is essential in determining the composition of the ices; most species have several vibrational modes and more secure identifications can be made when the features are studied simultaneously. The identification process is aided by studying large samples of sight-lines facilitating correlation of band strengths and shapes with each other and, if known, with physical characteristics along the sight-lines, such as luminosity and evolutionary stage of any heating source.  The great sensitivity of the IRS allows for observations of objects with luminosities that are down by orders of magnitude with respect to what could be observed before in the mid-infrared (typically 1 L$_{\\odot}$ versus $10^4$ L$_{\\odot}$).  Most data presented in this paper were obtained from spectral surveys of nearby molecular clouds and isolated dense cores in the context of the Spitzer Legacy Program `From Molecular Cores to Planet-Forming Disks' (`c2d'; \\citealt{eva03}). Initial results from this program emphasize the importance of thermal processing in the evolution of the ices, as the ices surrounding the low mass YSO HH~46 IRS are more processed compared to B5 IRS1 \\citep{boo04_2}. Ices toward the edge-on disk CRBR 2422.8-3423 were presented in \\citet{pon05} and also show signs of heating. Finally, mid-infrared spectra of ices toward background stars tracing quiescent cloud material were published in \\citet{kne05}, showing for the first time that the (unknown) carrier of the 6.85 \\mum\\ absorption band is abundant even at the coldest conditions away from star formation.  This paper specifically addresses questions on the identification of the 6.0 and 6.85 \\mum\\ bands. These prominent absorption features are still not fully identified, despite being detected nearly 30 years ago with the Kuiper Airborne Observatory \\citep{pue79} and commonly observed toward massive YSOs with the Infrared Space Observatory (ISO; \\citealt{sch96}, \\citealt{kea01b}).  The peak position of the 6.85 \\mum\\ band varies dramatically, and is thought to be a function of the ice temperature \\citep{kea01b}. Either one carrier with a pronounced temperature dependence of its absorption bands, or two independent carriers, one much more volatile than the other, must be responsible for the 6.85 \\mum\\ feature.  If the former is the case, the ammonium ion (NH$_4^+$) is considered a promising candidate \\citep{sch03}.  The 6.0 \\mum\\ band was initially thought to be mainly due to the bending mode of H$_2$O (\\citealt{tie84}; however, see \\citealt{cox89}), but analysis of ISO spectra indicated that this is not the case toward several massive YSOs \\citep{sch96, kea01b}. The depth of the 6.0 \\mum\\ band is in some cases significantly (factor of 2) deeper than that expected from the 3.0 \\mum\\ H$_2$O stretching mode.  Part of this `excess' might be due to a strong vibration of the HCOOH molecule. Recent observations of background stars show a small excess ($\\leq$25\\%), of which half could be due to the enhanced band strength of the H$_2$O bending mode in CO$_2$-rich ices \\citep{kne05}. Finally, it was claimed that much of the 6.0 \\mum\\ excess and part of the 6.85 \\mum\\ band is due to highly processed ices, i.e. a mixture of complex species with C--H and O--H bonds produced after irradiation of simple ices \\citep{gib02}. In this case, one would expect the 6.0 \\mum\\ excess to be enhanced in high radiation environments, such as near more evolved protostars.  Thus, fundamental questions remain on the identification of the 6.0 and 6.85 \\mum\\ bands, and with the large sample presented in this work the possible answers can be further constrained.  Subsequent papers will specifically address the CO$_2$ \\citep{pon08}, CH$_4$ \\citep{obe08}, and NH$_3$ (S.  Bottinelli et al., in preparation) species. Further papers in this series will investigate the ices toward background stars behind large clouds (C. Knez et al., in preparation) and isolated cores (A. C. A. Boogert et al., in preparation).  In \\S\\ref{sec:sou} the source sample is presented, while the observations and data reduction of the ground-based L-band and Spitzer 5-38 \\mum\\ spectra are described in \\S\\ref{sec:obs}. In \\S\\ref{sec:res} the continuum level for the multitude of absorption features is determined, and subsequently the best value for the H$_2$O column density derived, followed by an empirical decomposition of the 5-7 \\mum\\ absorption complex. The positively identified species, or correlations of the components with other observables are discussed in \\S\\ref{sec:id}. Finally, the abundances of the simple species, and constraints on the carriers of unidentified components are discussed in \\S\\ref{sec:dis}.    \\begin{deluxetable*}{lllllrlll} \\tabletypesize{\\scriptsize} \\tablecolumns{9} \\tablewidth{0pc} \\tablecaption{Source Sample~\\label{t:sample}} \\tablehead{ \\colhead{Source}& \\colhead{RA\\tablenotemark{a}} & \\colhead{Dec\\tablenotemark{a}} & \\colhead{Cloud} & \\colhead{Type\\tablenotemark{b}} & \\colhead{$\\alpha_{\\rm IR}$\\tablenotemark{c}} & \\colhead{Obs. ID\\tablenotemark{d}} & \\colhead{Module\\tablenotemark{l}} & \\colhead{L-band\\tablenotemark{m}}\\\\ \\colhead{      }& \\colhead{J2000}        & \\colhead{J2000}         & \\colhead{   }   & \\colhead{ } & \\colhead{ } & \\colhead{ } & \\colhead{ } & \\colhead{ }\\\\} \\startdata L1448 IRS1      & 03:25:09.44            &  +30:46:21.7            &   Perseus       & low         & 0.34                 & \\dataset{0005656832} & SL,SH,LH    & NIRSPEC   \\\\ IRAS 03235+3004 & 03:26:37.45            &  +30:15:27.9            &   Perseus       & low         & 1.44                 & \\dataset{0009835520} & SL,LL       & NIRSPEC   \\\\ IRAS 03245+3002 & 03:27:39.03            &  +30:12:59.3            &   Perseus       & low         & 2.70                 & \\dataset{0006368000} & SL,SH       &           \\\\ L1455 SMM1      & 03:27:43.25            &  +30:12:28.8            &   Perseus       & low         & 2.41                 & \\dataset{0015917056} & SL,SH,LL1   &           \\\\ RNO 15          & 03:27:47.68            &  +30:12:04.3            &   Perseus       & low         &$-$0.21               & \\dataset{0005633280} & LL1,SL,SH,LH& NIRSPEC   \\\\ L1455 IRS3      & 03:28:00.41            &  +30:08:01.2            &   Perseus       & low         & 0.98                 & \\dataset{0015917568} & SL,SH,LL1   &           \\\\ IRAS 03254+3050 & 03:28:34.51            &  +31:00:51.2            &   Perseus       & low         & 0.90                 & \\dataset{0011827200} & LL1,SL,SH,LH& NIRSPEC   \\\\ IRAS 03271+3013 & 03:30:15.16            &  +30:23:48.8            &   Perseus       & low         & 2.06                 & \\dataset{0005634304} & LL1,SL,SH,LH& NIRSPEC   \\\\ IRAS 03301+3111 & 03:33:12.85            &  +31:21:24.2            &   Perseus       & low         & 0.51                 & \\dataset{0005634560} & SL,SH,LH    & NIRSPEC   \\\\ B1-a            & 03:33:16.67            &  +31:07:55.1            &   Perseus       & low         & 1.87                 & \\dataset{0015918080} & SL,SH,LL1   & NIRSPEC   \\\\ B1-c            & 03:33:17.89            &  +31:09:31.0            &   Perseus       & low         & 2.66                 & \\dataset{0013460480} & SL,SH,LL1   &           \\\\ B1-b            & 03:33:20.34            &  +31:07:21.4            &   Perseus       & low         & 0.68                 & \\dataset{0015916544} & SL,LL       &           \\\\ B5 IRS3         & 03:47:05.45            &  +32:43:08.5            &   Perseus       & low         & 0.51\\tablenotemark{k}& \\dataset{0005635072} & LL1,SL,SH,LH& NIRSPEC   \\\\ B5 IRS1\\tablenotemark{e} & 03:47:41.61            &  +32:51:43.8            &   Perseus       & low         & 0.78\\tablenotemark{k}& \\dataset{0005635328} & SL,SH,LH    & NIRSPEC   \\\\ L1489 IRS\\tablenotemark{f} & 04:04:43.07            &  +26:18:56.4            &   Taurus        & low         & 1.10                 & \\dataset{0003528960} & SL,SH,LH    & NIRSPEC   \\\\ IRAS 04108+2803B\\tablenotemark{f} & 04:13:54.72            &  +28:11:32.9            &   Taurus        & low         & 0.90                 & \\dataset{0003529472} & SL,SH,LH    & NIRSPEC   \\\\ HH~300\\tablenotemark{f} & 04:26:56.30            &  +24:43:35.3            &   Taurus        & low         & 0.79                 & \\dataset{0003530752} & SL,SH,LH    & NIRSPEC   \\\\ DG Tau B\\tablenotemark{f} & 04:27:02.66            &  +26:05:30.5            &   Taurus        & low         & 1.16                 & \\dataset{0003540992} & SL,SH,LH    & NIRSPEC   \\\\ HH~46 IRS\\tablenotemark{e} & 08:25:43.78            &$-$51:00:35.6            &   HH~46         & low         & 1.70                 & \\dataset{0005638912} & SL,SH,LH    & ISAAC  \\\\ IRAS 12553-7651 & 12:59:06.63            &$-$77:07:40.0            &   Cha           & low         & 0.76                 & \\dataset{0009830912} & LL1,SL,SH,LH&           \\\\ IRAS 13546-3941 & 13:57:38.94            &$-$39:56:00.2            &   BHR92         & low         & $-$0.06              & \\dataset{0005642752} & SL,SH,LH    &           \\\\ IRAS 15398-3359 & 15:43:02.26            &$-$34:09:06.7            &   B228          & low         & 1.22                 & \\dataset{0005828864} & SL,SH,LL1   &           \\\\ Elias 29\\tablenotemark{g} & 16:27:09.42            &$-$24:37:21.1            &   Oph           & low         & 0.53\\tablenotemark{k}& 26700814             & SWS01 sp3   &             \\\\ CRBR 2422.8-3423\\tablenotemark{h} & 16:27:24.61            &$-$24:41:03.3            &   Oph           & low         & 1.36                 & \\dataset{0009346048} & SL,SH,LH    & NIRSPEC   \\\\ RNO 91          & 16:34:29.32            &$-$15:47:01.4            &   L43           & low         & 0.03                 & \\dataset{0005650432} & SL,SH,LH    & ISAAC   \\\\ IRAS 17081-2721 & 17:11:17.28            &$-$27:25:08.2            &   B59           & low         & 0.55                 & \\dataset{0014893824} & SL,SH,LL1   & NIRSPEC   \\\\ SSTc2dJ171122.2-272602 & 17:11:22.16            &$-$27:26:02.3            &   B59           & low         & 2.26                 & \\dataset{0014894336} & SL,LL       &           \\\\ 2MASSJ17112317-2724315 & 17:11:23.13            &$-$27:24:32.6            &   B59           & low         & 2.48                 & \\dataset{0014894592} & SL,LL       & NIRSPEC   \\\\ EC 74           & 18:29:55.72            &  +01:14:31.6            &   Serpens       & low         & $-$0.25              & \\dataset{0009407232} & SL,SH,LH    & NIRSPEC   \\\\ EC 82           & 18:29:56.89            &  +01:14:46.5            &   Serpens       & low         & 0.38\\tablenotemark{k}& \\dataset{0009407232} & SL,SH,LH    &         \\\\ SVS 4-5         & 18:29:57.59            &  +01:13:00.6            &   Serpens       & low         & 1.26                 & \\dataset{0009407232} & SL,SH,LH    & ISAAC   \\\\ EC 90           & 18:29:57.75            &  +01:14:05.9            &   Serpens       & low         & $-$0.09              & \\dataset{0009828352} & SL,SH,LH    &           \\\\ EC 92           & 18:29:57.88            &  +01:12:51.6            &   Serpens       & low         & 0.91                 & \\dataset{0009407232} & SL,SH,LH    & NIRSPEC   \\\\ CK4             & 18:29:58.21            &  +01:15:21.7            &   Serpens       & low         &$-$0.25               & \\dataset{0009407232} & SL,SH,LH    &           \\\\ R CrA IRS 5     & 19:01:48.03            &$-$36:57:21.6            &   CrA           & low         & 0.98                 & \\dataset{0009835264} & SL,SH,LL1   & ISAAC     \\\\ HH 100 IRS\\tablenotemark{i} & 19:01:50.56            &$-$36:58:08.9            &   CrA           & low         & 0.80                 & 52301106             & SWS01 sp4   &             \\\\ CrA IRS7 A      & 19:01:55.32            &$-$36:57:22.0            &   CrA           & low         & 2.23                 & \\dataset{0009835008} & SL,SH,LH    & ISAAC     \\\\ CrA IRS7 B      & 19:01:56.41            &$-$36:57:28.0            &   CrA           & low         & 1.63                 & \\dataset{0009835008} & SL,SH,LH    & ISAAC     \\\\ CrA IRAS32      & 19:02:58.69            &$-$37:07:34.5            &   CrA           & low         & 2.15                 & \\dataset{0009832192} & SL,SH,LL1   &           \\\\ L1014 IRS       & 21:24:07.51            &  +49:59:09.0            &   L1014         & low         & 1.28                 & \\dataset{0012116736} & SL,LL       & NIRSPEC   \\\\ IRAS 23238+7401 & 23:25:46.65            &  +74:17:37.2            &   CB 244        & low         & 0.95                 & \\dataset{0009833728} & SL,SH,LH    & NIRSPEC   \\\\ &                        &                         &                 &             &                      &                      &             &           \\\\ W3 IRS5\\tablenotemark{i} & 02:25:40.54            &$+$62:05:51.4            &                 & high        & 3.53                 & 42701302             & SWS01 sp3   &             \\\\ MonR2 IRS3\\tablenotemark{i} & 06:07:47.8             &$-$06:22:55.0            &                 & high        & 1.66                 & 71101712             & SWS01 sp3   &             \\\\ GL989\\tablenotemark{i} & 06:41:10.06            &$+$09:29:35.8            &                 & high        & 0.52                 & 71602619             & SWS01 sp3   &             \\\\ W33A\\tablenotemark{i} & 18:14:39.44            &$-$17:52:01.3            &                 & high        & 1.92                 & 32900920             & SWS01 sp4   &             \\\\ GL7009S\\tablenotemark{i} & 18:34:20.91            &$-$05:59:42.2            &                 & high        & 2.52                 & 15201140             & SWS01 sp3   &             \\\\ GL2136\\tablenotemark{i} & 18:22:26.32            &$-$13:30:08.2            &                 & high        & 1.48                 & 33000222             & SWS01 sp3   &             \\\\ S140 IRS1\\tablenotemark{i} & 22:19:18.17            &$+$63:18:47.6            &                 & high        & 1.57                 & 22002135             & SWS01 sp4   &             \\\\ NGC7538 IRS9\\tablenotemark{i} & 23:14:01.63            &$+$61:27:20.2            &                 & high        & 2.31                 & 09801532             & SWS01 sp2   &             \\\\ &                        &                         &                 &             &                      &                      &             &           \\\\ Elias 16\\tablenotemark{k} & 04:39:38.88  \t   &$+$26:11:26.6            &   Taurus        & bg          & -                    & \\dataset{0005637632} & SL,SH       &          \\\\ EC 118\\tablenotemark{k} & 18:30:00.62  \t   &$+$01:15:20.1            &   Serpens       & bg          & -                    & \\dataset{0011828224} & SL,SH       &          \\\\ \\enddata \\tablenotetext{a}{Position used in Spitzer/IRS observations} \\tablenotetext{b}{Source type: 'low'=low mass YSO, 'high'=massive YSO, 'bg'=background star} \\tablenotetext{c}{Broad-band spectral index as defined in Eq.~\\ref{eq:alpha}} \\tablenotetext{d}{AOR key for Spitzer and TDT number for ISO observations} \\tablenotetext{e}{Published previously in Boogert et al. 2004} \\tablenotetext{f}{Published previously in Watson et al. 2004} \\tablenotetext{g}{Published previously in Boogert et al. 2000} \\tablenotetext{h}{Published previously in Pontoppidan et al. 2005} \\tablenotetext{i}{Published previously in Keane et al. 2001} \\tablenotetext{j}{Published previously in Knez et al. 2005} \\tablenotetext{k}{$\\alpha$ enhanced due to  foreground extinction. Exclusion $K_{\\rm s}$-band flux gives much lower $\\alpha$: $-$0.16 (Elias 29), $-0.02$ (B5 IRS1), 0.18 (B5 IRS3), and 0.38 (EC 82)} \\tablenotetext{l}{Spitzer/IRS modules used: SL=Short-Low (5-14 \\mum, $R\\sim100$), LL=Long-Low (14-34 \\mum, $R\\sim100$), SH=Short-High (10-20 \\mum, $R\\sim600$), LH=Long-High (20-34 \\mum, $R\\sim600$); ISO SWS modes used (2.3-40 \\mum): SWS01 speed 1 ($R\\sim250$), speed 2 ($R\\sim250$), speed 3 ($R\\sim400$), speed 4 ($R\\sim800$)} \\tablenotetext{m}{Complementary ground-based L-band observations with Keck/NIRSPEC or VLT/ISAAC} \\end{deluxetable*} ",
        "conclusions": "~\\label{sec:concl}  The present work extends the study of ices over the full 3-20 \\mum\\ wavelength range from the previously well studied massive YSOs ($\\sim10^5$L$_{\\odot}$) to low mass YSOs ($\\sim$1L$_{\\odot}$). The following new insights were obtained:  \\begin{itemize}  \\item Absorption features are detected at 6.0 and 6.85 \\mum\\ in all sources, including high- and low-mass YSOs and background stars tracing quiescent cloud material. An empirical decomposition shows that the 5-7 \\mum\\ absorption complex consists, in addition to the H$_2$O bending mode, of a combination of at least 5 independent absorption components.  \\item In a subset of sources additional weak features are detected at 7.25, 7.40, and at 9.0 and 9.7 \\mum\\ on top of the prominent Si-O stretching band of silicates. These features are associated with CH$_3$OH, HCOOH, NH$_3$, and possibly HCOO$^-$ and indicate abundances of 1-30\\%, 1-5\\%, 3-8\\%, and 0.3\\% with respect to H$_2$O, respectively.  The large source-to-source solid CH$_3$OH abundance variations are likely a result of the conditions at the time of grain surface chemistry.  \\item Component C1 (5.7-6.0 \\mum) is mostly explained by solid HCOOH and H$_2$CO at abundances with respect to solid H$_2$O of typically 1-5\\% and $\\sim$6\\%, respectively.  \\item Component C2 (6.0-6.4 \\mum) also likely arises from a blend of several species.  Solid NH$_3$ can account for 10-50\\% of the absorption. A similar amount can be attributed to absorption by monomers, dimers, and small multimers of H$_2$O mixed with a substantial amount of CO$_2$. The long-wavelength side of C2 is potentially due to anions produced by acid-base chemistry (e.g. HCOO$^-$) or energetic processing.  Indeed, a weak correlation with C4 suggests the carriers of these bands (possibly salts) are related.  \\item Components C3 and C4, the 6.85 \\mum\\ band, show the same characteristics for low mass YSOs and background stars as was previously found for massive YSOs by \\citet{kea01b}.  Their ratio is empirically found to be dependent on dust temperature. The carrier of C4 is likely less volatile than that of C3 and than H$_2$O, but it is not observed in the diffuse ISM.  These characteristics are also consistent with both components being due to NH$_4^+$ (ammonium ion). The detection of strong 6.85 \\mum\\ bands toward deeply embedded YSOs and background stars requires a production at low temperature (acid-base chemistry or processing by cosmic rays), given the lack of heating sources and stellar UV fields.  \\item The origin of the very broad component C5 (5-8 \\mum) is least understood. It is quite strong in several low and high mass YSOs, absent in others and so far undetected toward background stars. It is possibly related to thermal processing, as a weak correlation with the ratio of the polar/apolar CO components is observed and its shape resembles that of warm H$_2$O.  A blend of absorption by ions, as proposed by \\citep{sch03}, which would require both heating and energetic processing of the ices, cannot be excluded. The latter could also lead to an organic residue, whose shape resembles that of the C5 component as well.  \\item Weak correlations are found between the absorption components in the 5-8 \\mum\\ range and tracers of physical conditions.  A more thorough understanding of the conditions, the source geometry and local radiation fields in specific lines of sight is required to further constrain the nature of the carriers (in particular for components C2-C5) and the importance of thermal and energetic processing.  To what degree does scattering play a role? What are the dust temperatures, cosmic ray fluxes, UV radiation fields and time scales as a function of distance along the line of sight? Measurements of source fluxes above wavelengths of 100 \\mum\\ with the Herschel Space Observatory will be valuable, as well as ground-based measurements of the 4.62 \\mum\\ band of OCN$^-$ in more lines of sight.  In addition, further laboratory work, in particular on the effect of cosmic rays on the ices, especially for the identification of components C3 and C4, is required. Finally, study of the 5-8 \\mum\\ features toward a larger sample of background stars is crucial to make further progress in the identification and processing history of interstellar ices.   \\end{itemize}"
    },
    "0801/0801.1956_arXiv.txt": {
        "abstract": "Continuous observations were obtained of active region 10953 with the Solar Optical Telescope (SOT) on board the \\emph{Hinode} satellite during 2007 April 28 to May 9. A prominence was located over the polarity inversion line (PIL) in the south-east of the main sunspot. These observations provided us with a time series of vector magnetic fields on the photosphere under the prominence. We found four features: (1) The abutting opposite-polarity regions on the two sides along the PIL first grew laterally in size and then narrowed. (2) These abutting regions contained vertically-weak, but horizontally-strong magnetic fields. (3) The orientations of the horizontal magnetic fields along the PIL on the photosphere gradually changed with time from a normal-polarity configuration to a inverse-polarity one. (4) The horizontal-magnetic field region was blueshifted. These indicate that helical flux rope was emerging from below the photosphere into the corona along the PIL under the pre-existing prominence. We suggest that this supply of a helical magnetic flux into the corona is associated with evolution and maintenance of active-region prominences. ",
        "introduction": "Solar prominences are cool material (10$^4$ K) floating in the corona (10$^6$ K) and located over polarity inversion lines (PILs) of the photosphere. It is known that prominences are supported by coronal magnetic fields against gravity (see references in Martin 1998). Such magnetic fields have dipped shapes, and the cool plasma is sitting at the bottom of the dips (Kippenhahn \\& Schl\\\"{u}ter 1957). Moreover, many observations of prominences show inverse polarity (Leroy et al. 1984; Tandberg-Hanssen 1995; Lites 2005), which means that the direction of horizontal magnetic fields is toward the positive polarity side from the negative polarity side. Hence, it is often thought that magnetic fields in prominences have helical structures (e.g., Kuperus \\& Tandberg-Hanssen 1967; Kuperus \\& Raadu 1974; Hirayama 1985). This topology of magnetic fields is suggested from observations of erupting prominences, or coronal mass ejections (e.g., Dere et al. 1999; Ciaravella et al. 2000; Low 2001).  How is such a helical magnetic field created in association with prominences in the corona ? Two theories have been discussed: flux rope models (e.g., Rust \\& Kumar 1994; Low \\& Hundhausen 1995; Low 1996; Lites 2005; Zhang \\& Low 2005) and sheared-arcade models (e.g., Pneuman 1983; van Ballegooijen \\& Martens 1989, 1990; Antiochos et al. 1994; DeVore \\& Antiochos 2000; Martens \\& Zwaan 2001; Aulanier et al. 2002; Karpen et al. 2003; Mackay \\& van Ballegooijen 2005, 2006). In the former model, an originally-twisted flux rope emerges from below the photosphere into the corona. The helical structure is thought to be made by the convection in the solar interior. In the latter one, potential fields in the corona are sheared by the photospheric motion along PILs. As a result, magnetic reconnection occurs among the sheared fields, and then helical fields are constructed in the corona. Both models have the same final configuration, but the processes leading to it are quite different. However, we have no clear observational evidence to support either model, although emerging twisted magnetic fields related to flares have already been observed (e.g., Tanaka 1991; Lites et al. 1995; Leka et al. 1996; Ishii et al. 1998). We have had difficulties due to seeing and discontinuity of observing time with ground-based observations in revealing the mechanism of prominence formation.  The \\emph{Hinode} satellite (Kosugi et al. 2007) has a sun-synchronous polar orbit, so that we can observe the Sun continuously without interruption by a spacecraft night. The Solar Optical Telescope (SOT, Tsuneta et al. 2007; Suematsu et al. 2007; Ichimoto et al. 2007; Shimizu et al. 2007) on board \\emph{Hinode} has the Spectro-Polarimeter (SP), which provides long sequences of vector magnetic field measurements at high spatial resolution. We have an observation of one active region that had a pre-existing prominence to investigate the evolution of its associated magnetic fields on the photosphere. We report on the results and findings from this observation in this letter. ",
        "conclusions": "The \\emph{Hinode}/SP observations presented here offer clear evidence of an emerging helical flux rope along the PIL under an active-region prominence: The {\\it weak-field region} became first wider and then narrower during the emergence of the flux rope, while the orientation of the horizontal magnetic fields observed at the photosphere changes from the normal-polarity to the inverse-polarity configuration. We note that we distinguish the {\\it weak-field region} from the filament channel since the {\\it weak-field region} appears temporarily in the filament channel. Cheung et al. (2007) observed similar polarity reversal in their radiative MHD simulation. Lites (2005) pointed out in a similar evolutionary event that the abutting opposite polarities may spatially separate under an active-region prominence although the cadence of that observation was one day and the resolution was much lower than that of ours. Gaizauskas et al. (1997) also reported that elongated voids appeared along a filament channel seen in off-band H$\\alpha$ a few days before prominence formation. Our results with \\emph{Hinode}/SOT suggest that such broadening is associated with emergence of the horizontal flux tubes.  We obtain the following physical parameters from the ME inversion of the SP data for our emerging flux rope event. The {\\it weak-field region} had a width of up to 10,000 km, a duration of one and a half days, and an upward velocity of up to 300 m $s^{-1}$. Hence, the diameter of the flux tube is roughly estimated to be 39,000 km if the maximum upward velocity continues for one and a half days. This estimation is consistent with both the maximum width of the {\\it weak-field region} and typical height of active-region prominences on the limb. The filling factor and strength of horizontal magnetic fields are estimated to be 0.15 and 650 Gauss on average in the {\\it weak-field region}, respectively. When this flux is supplied into the corona, the average strength of the magnetic fields is 100 Gauss. This value is consistent with active-region prominences according to previous measurements of magnetic fields (e.g., Tandberg-Hanssen and Malville 1974; Wiehr and Stellmacher 1991; Casini et al. 2003; Okamoto et al. 2007), and this flux rope is strong enough to support prominences. Therefore, we suggest that this emergence of the helical magnetic flux rope is associated with evolution and maintenance of the prominence.  This is the only observation made so far of an active-region prominence, and we do not rule out the sheared arcade model (see Introduction). It will be useful to observe quantitatively both active-region and quiescent prominences with the SP to determine whether the event reported in this paper is common in the solar atmosphere. We suggest that the above observational evidence must be closely related to the evolution of a prominence via the emergence of a twisted magnetic flux rope. The H$\\alpha$ prominence was already seen before the start of the emerging episode reported here. The appearance of the pre-existing prominence was considerably changed, but after the emergence of the flux rope, the prominence became stable. We will report the relationship between the emerging helical flux rope and the prominence activity in a subsequence paper.   \\  The authors thank B. C. Low for useful comments. \\emph{Hinode} is a Japanese mission developed and launched by ISAS/JAXA, with NAOJ as domestic partner and NASA and STFC (UK) as international partners. It is operated by these agencies in co-operation with ESA and NSC (Norway). This work was carried out at the NAOJ Hinode science center, which was 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 donation from the Sun Microsystems Inc., and NAOJ internal funding."
    },
    "0801/0801.2572_arXiv.txt": {
        "abstract": "The existence of a primordial magnetic field (PMF) would affect both the temperature and polarization anisotropies of the cosmic microwave background (CMB).  It also provides a plausible  explanation for the possible disparity  between observations and theoretical fits to the CMB power spectrum. Here we report on calculations of not only the numerical CMB power spectrum from the PMF, but also the correlations between the CMB power spectrum from the PMF and the primary curvature perturbations. We then deduce a precise estimate of the PMF effect on all modes of perturbations. We find that the PMF affects not only the CMB TT and TE modes on small angular scales, but also on large angular scales. The introduction of a PMF  leads to a better fit to the CMB power spectrum for the higher multipoles, and the fit at lowest multipoles can be used to constrain the correlation of the PMF with the density fluctuations for large negative values of the spectral index. Our prediction for the BB mode for a PMF  average field strength $|B_\\lambda| =4.0$ nG is consistent with the upper limit on the BB mode deduced from the latest CMB observations. We find that the BB mode is dominated by the vector mode of the PMF for higher multipoles. We also show that by fitting the complete power spectrum one can break the degeneracy between the  PMF amplitude and its  power spectral index. ",
        "introduction": "Magnetic fields in clusters of galaxies have been observed \\cite{Kronberg:1992pp,Wolfe:1992ab, Clarke:2000bz,Xu:2005rb} with a strength of $0.1-1.0~\\mu$ G. The existence of a  primordial magnetic field (PMF) of order  1 nG  whose field lines collapse as structure forms is one possible explanation for such magnetic fields in galactic clusters.  The origin and detection of the PMF is, hence, a subject of considerable interest in modern cosmology. Moreover, the PMF could influence a variety of phenomena in the early universe \\cite{Grasso:2000wj} such as the cosmic microwave background (CMB) \\cite{ Challinor:2005ye, Dolgov:2005ti, Gopal:2005sg, Kahniashvili:2005xe, Kosowsky:2004zh, Lewis:2004ef, Mack:2001gc, Subramanian:1998fn, Subramanian:2002nh, Yamazaki:2005yd, Yamazaki:2004vq, Yamazaki:2006bq, Yamazaki:2006ah}, or the matter density field \\cite{ Giovannini:2004aw, Tashiro:2005hc, Tsagas:1999ft, Yamazaki:2006mi}.  Temperature and polarization anisotropies in the CMB provide very precise information on the physical processes in operation during the early universe (WMAP \\cite{Spergel:2006hy,Hinshaw:2006ia,Page:2006hz}, ACBAR \\cite{Kuo:2006ya}, CBI \\cite{Readhead:2004gy,Sievers:2005gj}, DASI \\cite{Leitch:2004gd}, BOOMERANG \\cite{Jones:2005yb}, and VSA \\cite{Dickinson:2004yr}).  The CMB power spectrum from ACBAR and CBI, has indicated a potential discrepancy between these observations at higher multipoles $\\ell \\ge 2000$ and the best-fit cosmological model to the WMAP power spectrum. A straightforward extension of the fit \\cite{Spergel:2006hy} to the WMAP data predicts a rapidly declining power spectrum  in the large multipole range due to the finite thickness of the photon last scattering surface and the Silk damping effect.  The  ACBAR and CBI experiments, however,  indicate continued power up to $\\ell \\sim 4000$. This discrepancy is difficult to account for by a simple retuning of cosmological parameters or by the Sunyev-Zeldovich effect \\cite{Spergel:2006hy,Komatsu02,Bond05}. Among other possible explanations, an inhomogeneous cosmological magnetic field generated before the CMB last-scattering epoch provides a plausible mechanism \\cite{Bamba:2004cu} to produce excess power at high multipoles. Such a field excites an Alfven-wave mode in the primordial baryon-photon plasma  and induces small rotational velocity  perturbations. Since this mode can survive on scales below those at which Silk damping occurs during recombination \\cite{Jedamzik:1996wp,Subramanian:1998fn}, it could be a new source of the CMB anisotropies on small angular scales. The present  work, therefore,  is an attempt to more precisely study the evolution of cosmological perturbations with a PMF.  Previous work \\cite{ Durrer:1999bk, Seshadri:2000ky, Campanelli:2004pm, Challinor:2005ye, Dolgov:2005ti, Gopal:2005sg, Kahniashvili:2005xe, Kosowsky:2004zh, Lewis:2004ef, Mack:2001gc, Subramanian:1998fn, Subramanian:2002nh, Yamazaki:2005yd, Yamazaki:2004vq, Yamazaki:2006bq, Yamazaki:2006ah} has shown that one can obtain information about the PMF from the CMB temperature anisotropies and polarization. However, in those works  attention was only given to a subset of the modes of the CMB anisotropies. In the present work, therefore, we  study the comprehensive effect of the PMF on all modes of the CMB perturbations. Furthermore, in order to clarify  the role of the PMF in the CMB, we take into consideration the possible correlation between the CMB fluctuations induced by the PMF and those due to primordial curvature and tensor perturbations.  Also, we numerically evaluate the CMB power spectrum from the stochastic PMF and thereby avoid recourse to analytic approximations.  In this article, we use adiabatic initial conditions for the evolution of primary density perturbations and consider isocurvature (isothermal) initial conditions when estimating effects on the CMB anisotropy induced by the PMF \\cite{Lewis:2004ef,Giovannini:2006gz}. Throughout this article we fix the best fit cosmological parameters of the $\\Lambda$CDM + Tensor model as follows \\cite{Spergel:2006hy}: $h=0.792$, $\\Omega_bh^2=0.02336$, $\\Omega_m h^2=0.1189$, $n_S=0.987$, $r=0.55$, $n_T=-r/8= -0.069$, and $\\tau_c=0.091$ in flat universe models, where $h$ denotes the Hubble parameter in units of 100 km s$^{-1}$Mpc$^{-1}$, $\\Omega_b$ and $\\Omega_m$ are the baryon and cold dark matter densities in units of the critical density, $n_S$ is the spectral index of the primordial scalar fluctuations, $r$ the ratio of the amplitude of the tensor fluctuations to the scalar potential fluctuations, $n_T$ is the spectral index of the primordial tensor fluctuations, and $\\tau_c$ is the optical depth for Compton scattering. ",
        "conclusions": "We have explored effects of a PMF on the CMB for the allowed PMF parameters which were  deduced in our previous work \\cite{Yamazaki:2006bq} (i.e.~$B_\\lambda < 10~ \\mathrm{nG~and~} n_\\mathrm{B} < -2.4$) . The upper panel of Figure 2 illustrates the CMB temperature and polarization anisotropies from the PMF for scalar, vector and tensor modes for the case when $B_\\lambda = 4.0$ nG and $n_\\mathrm{B} = -2.9$ or $-2.5$ as labeled. The scalar mode dominates for lower $\\ell$ of the TT and TE modes as shown by Giovannini \\cite{Giovannini:2006gz}. In particular, it is comparable in power to the primary TT mode for ($B_\\lambda, n_\\mathrm{B}) = (4.0 ~{\\rm nG}, -2.9$). We note that the curves with  ($B_\\lambda, n_\\mathrm{B}, s^\\mathrm{(S)},s^\\mathrm{(T)} ) = (4.0 ~{\\rm nG}, -2.5, 1,1$) give the best fits to the observed power spectra in both the regions of high and low $\\ell$.  This result  is complementary to and consistent with the nucleosynthesis constraints derived in \\cite{Caprini:2001nb} as shown in \\cite{Yamazaki:2007mm}.  For illustration, let us assume the same scale-invariant power spectrum for both the PMF and the primordial curvature perturbations. In this case, the ratio of the density and velocity perturbations induced by the primordial curvature perturbation to those by the PMF is proportional to $k^2$. Therefore, the temperature anisotropies from the PMF are larger for lower $\\ell$ compared with those from primordial curvature perturbations. Furthermore, the power of the CMB temperature anisotropies from the PMF for lower $\\ell$ depends not only on $B_\\lambda$ but also strongly on $n_\\mathrm{B}$ [Panel (1a) of Fig.2].   Note, that the magnetic field which is continually sourcing fluctuations does not spoil the phase coherence (cf.~cosmic defects). The basic behavior of acoustic oscillations is affected by the pressure of the fluid $kc_s\\delta^\\mathrm{(S)}_b$ and the potential $k\\phi$. Since the pressure of the PMF is sufficiently less than the thermal fluid pressure, the PMF dose not affect  the phase coherence significantly.  The PMF also affects the CMB power spectrum on  small angular scales for two reasons. First, the PMF energy density fluctuations depend only on the scale factor $a$ and  can survive below Silk damping scale. Therefore, the PMF continues to  source the fluctuations through the Lorentz force even below the Silk damping scale. Second, the vector mode from the PMF can be larger than the scalar modes both from the PMF and primordial perturbations at small scales. This is because, after horizon crossing, the latter cannot grow due to the photon pressure leading to acoustic oscillations, while the former can keep growing inside the cosmic horizon. This means that, for higher $\\ell$, the vector mode dominates the temperature anisotropies of the CMB over the scalar and tensor modes from the PMF and the contribution from primary anisotropies. The integrated amplitude of the gravitational waves from the PMF can be negligible after horizon crossing.  This is because the homogeneous solution begins to oscillate inside the horizon and decay rapidly \\cite{ 1979ZhPmR..30..719S, 1982PhLB..115..189R, 1985SvA....29..607P, Pritchard:2004qp}. Consequently, gravity waves  only affect the anisotropy spectrum  on scales larger than the horizon at recombination.  The tensor mode from the PMF  therefore decreases at higher $\\ell$. Thus, the vector mode of the CMB polarization from the PMF dominates for higher $\\ell$ [Panel (1d) of Fig.~2] \\cite{Seshadri:2000ky,Lewis:2004ef} \\footnote{We do not consider the magnetic-field-related BB polarization coming from the Faraday rotation effect \\cite{Campanelli:2004pm,Kosowsky:2004zh}.} In the primary spectrum for higher $\\ell$ the absolute value of the EE mode from the PMF is relatively-small [Panel (1c) of Fig.~2] compared to the TT mode [Panel (1a) of Fig.~2].  Hence, even though the EE mode of the primary spectrum damps less than the TT mode, the EE mode remains much smaller than the TT mode at higher $\\ell$ in the final power spectrum. For the TE mode [Panel (1b) of Fig.~2], the contribution from the PMF vector and scalar modes can be comparable to that from the primordial curvature  perturbations for higher $\\ell$.   Except for a small dip region near $\\ell = 40$, the scalar mode is always small compared to the primary spectrum.  Regarding the BB mode, we note that the  BB mode signal described in this paper is due to magnetic-field-induced CMB fluctuations (with the peak around $\\ell \\sim 2000$ as in \\cite{Seshadri:2000ky,Lewis:2004ef}).  We do not include magnetic-field related BB polarization coming from the Faraday rotation effect (with the peak around $\\ell > 15000$) discussed in \\cite{Campanelli:2004pm,Kosowsky:2004zh}. In our model, the BB mode from the PMF can dominate for $\\ell \\gtrsim 200$ if $B_{\\rm \\lambda}\\gtrsim 2.0$ nG. A potential problem in attempting to detect this signal on such angular scales, therefore,  is the contamination from gravitational lensing which converts the dominant EE power into the BB mode \\cite{Pritchard:2004qp}. However, since we already know quite accurately what the spectrum of the lensing signal must be, we can subtract its power directly. After removing the foreground effect, the BB mode from the PMF effect dominates for higher $\\ell$ even for PMF parameters allowed by the CMB temperature constraint \\cite{Yamazaki:2004vq,Yamazaki:2006bq,Caprini:2001nb}. Note, that the there is a change of scale between panels 1c and 1b.  The BB mode and EE modes are of comparable magnitude.  In Panel 2 of Fig. 2 we depict the CMB temperature and polarization anisotropies in the presence of a PMF  taking into account the correlations. Since we obtain the temperature and polarization anisotropies of the CMB with the PMF from isocurvature initial conditions, the phase of the CMB perturbation with the PMF is different by $\\pi /2$ from those without the PMF (on the adiabatic initial condition). The first, third, and odd numbered peaks of the scalar mode of the CMB perturbations rise for a positive correlation between the PMF and the energy density perturbations, while they are suppressed for a negative correlation. These are compared in Panel 2 of Figure 2 with the observed power spectra (WMAP \\cite{Hinshaw:2006ia,Page:2006hz}, ACBAR \\cite{Kuo:2006ya}, CBI \\cite{Readhead:2004gy,Sievers:2005gj}, DASI \\cite{Leitch:2004gd}, BOOMERANG \\cite{Jones:2005yb}, and VSA \\cite{Dickinson:2004yr}). Panels (2a) and (2b) of Fig.~2 show clearly that the power spectral index of the PMF, $n_\\mathrm{B}$, is more effectively constrained from CMB observations for lower $\\ell$ than those for higher $\\ell$. The models with higher $n_\\mathrm{B}$ give better fits to the observations than those with lower $n_\\mathrm{B}$ for the lower $\\ell$ regions of the TT and TE modes. Furthermore, there is no discrepancy at higher $\\ell$ between observations and theories of the CMB polarization for models with a PMF [Panels (2c) and (2d) of Fig.2]. In  our previous work \\cite{Yamazaki:2004vq,Yamazaki:2006bq} there was a problem from the strong degeneracy between $n_\\mathrm{B}$ and $B_\\lambda$. This degeneracy, however,  is broken by the different effects of the PMF on the CMB power spectrum for lower and higher $\\ell$.  The scalar-mode CMB temperature-anisotropy power-spectrum shape ($\\ell$-scaling for different $n_\\mathrm{B}$ at large scales  - low  $\\ell$) agrees with the results of \\cite{Kahniashvili:2007xe}, while the E-polarization power spectrum shape does not follow semi-analytical estimate in that paper. This difference is caused by the fact that we have included the effects of reionization \\cite{Dodelson:2003booka, Lewis02} which were neglected in \\cite{Kahniashvili:2007xe}. If the universe is re-ionized at $z_\\mathrm{re}$, CMB photons are scattered by electrons and the polarization is generated again. Since re-ionization  results in a new scattering surface  at relatively short distances from us at (and a relatively recent era), the viewing angle of the polarization from re-ionization becomes large. Thus, reionization causes some power to shift to lower multipoles \\cite{Kaplinghat:2002vt, Hu:2003gh}.  We also point out that the observed power spectrum of temperature fluctuations in the CMB is likely to depend on frequency \\cite{Kuo:2006ya}.  Such dependence is theoretically expected to originate from foreground effects such as the Sunyaev-Zel'dovich effect at higher multipoles. In contrast, the effects of a PMF are frequency-independent because the PMF affects the primary CMB as a background.  Therefore, the correlation between the PMF and other foreground effects should be weak.  Because of this, one should  be able to eventually distinguish  the PMF from  foreground effects by using more than two observational data sets at different frequencies.  In summary, we have found that we can constrain more precisely the power law index $n_\\mathrm{B}$ and amplitude $B_\\lambda$ of the PMF from all modes of CMB temperature and polarization anisotropies. The strong degeneracy of these parameters \\cite{Yamazaki:2004vq,Yamazaki:2006bq} is broken by the different effects of the PMF on the CMB power spectrum for lower and higher $\\ell$. The scalar mode from the PMF can be a main source for lower $\\ell$, while the vector mode can dominate for higher $\\ell$ in the CMB temperature anisotropies. Furthermore, these calculations suggest that it is possible to place a limit on the correlation parameters $s^\\mathrm{(X)}$ for large negative values of the spectral index.  For example , $s^\\mathrm{(T)}, s^\\mathrm{(S)}< 0$ for $n_\\mathrm{B} = -2.9$ is ruled out from the effects of the TT mode on both the lowest and highest multipoles as shown in panels 2a and 2b of Figure 2.   Such results may constrain models for  the origin of the PMF, along with other PMF parameters."
    },
    "0801/0801.2277.txt": {
        "abstract": "We present new spectroscopic and photometric data of the type Ibn supernovae 2006jc, 2000er and 2002ao. We discuss the general properties of this recently proposed supernova family, which also includes SN 1999cq. The early-time monitoring of SN 2000er traces the evolution of this class of objects during the first few days after the shock breakout. An overall similarity in the photometric and spectroscopic evolution is found among the members of this group, which would be unexpected if the energy in these core-collapse events was dominated by the interaction between supernova ejecta and circumstellar medium. Type Ibn supernovae appear to be rather normal type Ib/c supernova explosions which occur within a He-rich circumstellar environment. SNe Ibn are therefore likely produced by the explosion of Wolf-Rayet progenitors still embedded in the He-rich material lost by the star in recent mass-loss episodes, which resemble known luminous blue variable eruptions. The evolved Wolf-Rayet star could either result from the evolution of a very massive star or be the more evolved member of a massive binary system. We also suggest that there are a number of arguments in favour of a type Ibn classification for the historical SN 1885A (S-Andromedae), previously considered as an anomalous type Ia event with some resemblance to SN~1991bg. ",
        "introduction": "The most massive stars, i.e. those with initial masses larger than 30$M_\\odot$, are thought to end their lives with the core-collapse explosion of a Wolf-Rayet (WR) star. Before this stage, many of them undergo a short period of instability called the luminous blue variable (LBV) phase. This phase probably lasts of order 10$^{4}$-10$^{5}$ years \\citep{mae87,hum94,boh97,vink02,smi08}, during which they lose mass through recurrent mass-loss episodes. Most LBVs undergo fairly minor variability (typically below 1 magnitude) which is commonly called the S Doradus phase. But occasionally LBVs experience giant outbursts in which their luminosity rapidly increases and they can reach absolute visual magnitudes around $-14$ \\citep{dav97,mau05,vd06}, accompanied by the ejection of a large portion of their hydrogen envelope \\citep{smi6}. An overview of our current knowledge of LBVs can be found in \\citet{vge01}. At the time of explosion, which is expected to occur a few $\\times$ 10$^5$ years after the last major LBV outburst \\citep[see e.g.][]{heg03,eld04}, the mass of the WR star is expected to be of the order of 15-20$M_\\odot$. These stars may eventually explode as type Ib supernovae (SNe Ib) if they have lost the whole H envelope, or as type Ic supernovae (SNe Ic) if the stars are also stripped of their He mantles \\citep[for a comprehensive description of the different SN types see][]{fil97,tura07}.  However, recent work \\citep{gal07,kot06,smi06} has suggested the possibility that some very massive stars may explode already {\\sl during} the LBV phase, producing luminous supernovae (SNe). Previously, a post-LBV channel was proposed by \\citet{sal02} to explain the observed properties of the type IIn SN~1997eg. In the proposed scenario, LBVs may eventually produce SNe which show high luminosity, blue colour, narrow H$\\alpha$ line in emission, and very slow spectro-photometric evolution. Their spectra are dominated by very narrow ($\\lesssim$1000 km s$^{-1}$), prominent H emission lines which sometimes show  complex, multicomponent profiles. These objects would therefore belong to the so-called {\\sl type IIn} SN class \\citep{sch90}. The observed properties of these events are consistent with a scenario in which a significant fraction of the kinetic energy of the SN ejecta is converted into radiation via interaction with a dense, H-rich circumstellar medium (CSM).  \\begin{figure} \\includegraphics[width=8.5cm]{figure1.ps} \\caption{Late R-band images of SN~2006jc. Liverpool Telescope image obtained on December 22, 2006 (left) and a very late VLT image (March 22, 2007; right). In this image, the SN is only marginally detectable and it is very close to a luminous star-forming region. North is down, East is to the left. \\label{fig1}} \\end{figure}  A new clue for the death of the most massive stars has recently been proposed for the explosion of the peculiar SN 2006jc \\citep{ita06,pasto07,fol07,smi07,tom07,seppo07}. Extensive data sets spanning many wavelength regions have been also presented by \\citet{imm07b,anu07,elisa07,sak07}. \\cite{pasto07} showed the amazing discovery that a major outburst, reaching M$_R = -14.1$ (i.e. an absolute magnitude which is comparable to that of LBV eruptions), occurred only two years before the explosion of SN 2006jc, at exactly the same position as the SN and inferred that the two events must be physically related. %If the major outburst occurred in an LBV phase it would be surprising that %core-collapse happened so soon afterwords, at least from a theoretical stellar %evolution viewpoint. The physical connection between the two transients is supported by analysis of the SN spectra. These are indeed blue and dominated by relatively narrow ($\\sim$2200 km s$^{-1}$) He I lines in emission \\citep[hence the new classification as {\\sl type Ibn},][]{pasto07}, likely indicating a type Ib/c SN embedded within a dense, massive He-rich envelope. Unlike SNe IIn, the photometric evolution of this object is extremely rapid, while the spectra do not change as quickly as the luminosity. The best studied object of this type is SN~2006jc, but  three additional events which appear to be spectroscopically rather similar have been found: SNe~1999cq, 2000er and 2002ao \\citep[discovered by][respectively]{mod99,cha00,mar02}. \\citet{fol07} first proposed that 2006jc-like events may constitute a distinct class of core-collapse SNe exploding in a He-rich CSM. SN~2005la \\citep{puk05} is also a somewhat similar SN, although with some unique properties. One might reasonably consider it as intermediate between a type Ibn and a type IIn event \\citep[see][]{pasto07b}.  \\begin{figure*} \\includegraphics[width=8.4cm]{figure2a.ps} \\includegraphics[width=8.4cm]{figure2b.ps} \\caption{R-band images of SN 2000er (left) and SN 2002ao (right). They were obtained on November 26, 2000 using the ESO 3.6m Telescope in La Silla (Chile) and  on February 7, 2002 using the IAC 80-cm, in La Palma (Canary Islands, Spain), respectively. \\label{fig2}} \\end{figure*}  This article is the first of a series of three papers \\citep[together with][]{pasto07b,seppo07} in which we will discuss some peculiar, transitional events possibly produced by the explosions of very massive stars. In this work we will focus on the properties of the 4 objects which belong to the type Ibn SN family. Some of the data analysed in this paper have already been published \\citep{mat00,fol07,pasto07}, others have never been shown before. For SN 1999cq we do not present new data, but we consider only those of \\citet{mat00}. The afore mentioned SN~2005la, which has even more peculiar characteristics, will be analysed separately \\citep{pasto07b}.  The paper is organised as follows: in Section \\ref{host} the properties of the host galaxies of the SN sample are introduced, including reddening and distance estimates. Spectroscopic  and photometric observations are presented in Section \\ref{sp} and Section \\ref{lc}, respectively. The quasi-bolometric light curve of SN~2006jc is presented in  Section \\ref{bolo}, while the parameters derived from the bolometric light curve modelling are discussed in Section \\ref{model}. In Section \\ref{prog} we examine two different plausible scenarios for the progenitors of SN~2006jc and similar events, while in Section \\ref{rate} we estimate the frequency of type Ibn events. Finally, a short summary is reported in Section \\ref{end}. ",
        "conclusions": "\\label{end}  We have presented new data for three type Ibn SNe, i.e. SN~2000er, SN~2002ao and SN~2006jc. In particular, the early-time data of SN~2000er show, for the first time, the early spectroscopic behaviour of a type Ibn event. Early-time spectra of SN~2000er show relatively broad wings in the He I lines, with additional much narrower components. The detection of broad He I spectral features is unequivocal evidence for the presence of  He also in the SN ejecta, not only in the CSM. Therefore, a residual envelope of He might be present in the progenitors of some SNe Ibn. Finally, the very late spectra of SN~2006jc still show strong narrow circumstellar He I lines, together with an emerging H$\\alpha$ likely produced in another CSM region.  This data set allows us to highlight some key points concerning this new class of SNe, since we have now a better picture of the overall spectro-photometric evolution of SNe Ibn, which is essential to better constrain the characteristics of this class: i) SNe Ibn show a surprisingly high degree of homogeneity, which is unexpected if the ejecta were strongly interacting with surrounding CSM; ii) the modelling of the bolometric light curve of SN~2006jc (assuming no ejecta-CSM interaction) suggests that an amount of 0.2-0.4M$_\\odot$ of $^{56}$Ni was ejected. The mass of $^{56}$Ni could even be slightly higher for the brightest SNe Ibn (e.g. SN~1999cq). However, we cannot exclude that the interaction between SN ejecta and CSM may power the light curves of SNe Ibn to some degree, hence the $^{56}$Ni mass might be somewhat lower; iii) the observed data can be modelled  without invoking extraordinarily massive ejecta; iv) their intrinsic rarity (being $\\sim$ 1$\\%$ of all core-collapse SNe) is probably due to the fact that they arise from a rather exotic progenitor scenario, consistent with either a post-LBV/early-WR channel of a very massive (60-100M$_\\odot$) star, or a binary (or even multiple) system formed by a normal LBV plus a core-collapsing WR companion."
    },
    "0801/0801.0741_arXiv.txt": {
        "abstract": "% {Since 1999, we have been conducting a radial velocity survey of 179 K giants using the Coud\\'e Auxiliary Telescope at UCO/Lick observatory. At present $\\sim$20$-$100 measurements have been collected per star with a precision of 5 to 8 m s$^{-1}$. Of the stars monitored, 145 (80\\%) show radial velocity (RV) variations at a level $>$20 m s$^{-1}$, of which 43 exhibit significant periodicities.} {Our aim is to investigate possible mechanism(s) that cause these observed RV variations. We intend to test whether these variations are intrinsic in nature, or possibly induced by companions, or both. In addition, we aim to characterise the parameters of these companions.} {A relation between $\\log g$ and the amplitude of the RV variations is investigated for all stars in the sample. Furthermore, the hypothesis that all periodic RV variations are caused by companions is investigated by comparing their inferred orbital statistics with the statistics of companions around main sequence F, G, and K dwarfs.} {A strong relation is found between the amplitude of the RV variations and  $\\log g$ in K giant stars, as suggested earlier by Hatzes \\& Cochran (1998). However, most of the stars exhibiting periodic variations are located above this relation. These RV variations can be split in a periodic component which is not correlated with $\\log g$ and a random residual part which \\textsl{does} correlate with $\\log g$. Compared to main-sequence dwarf stars, K giants frequently exhibit periodic RV variations. Interpreting these RV variations as being caused by companions, the orbital parameters are different from the companions orbiting dwarfs.} {Intrinsic mechanisms play an important role in producing RV variations in K giants stars, as suggested by their dependence on $\\log g$. However, it appears that periodic RV variations are \\textsl{additional} to these intrinsic variations, consistent with them being caused by companions. \\textbf{\\rm If indeed the majority of the periodic RV variations in K giants is interpreted as due to substellar companions, then massive planets are significantly more common around K giants than around F, G, K main-sequence stars.}} ",
        "introduction": "For more than a decade, radial velocity observations with accuracies of order m\\,s$^{-1}$ have been within reach (see for instance \\citet{marbut2000} and \\citet{queloz2001}). Even accuracies of less than 1 m\\,s$^{-1}$ \\citep{pepe2003} are possible now. With these observations, more than 200 sub-stellar companions have been discovered by measuring the reflex motions of their parent stars. Most of these sub-stellar companions have been detected around F, G and K main sequence stars, but detections around an A star \\citep{galland2006} and several subgiants (\\citet{johnson2006}, \\citet{johnson2007}) have also been reported recently.  Moreover, 10 giant stars were reported to have sub-stellar companions ($\\iota$ Draconis (K2III) \\citet{frink2002}, HD104985 (G9III) \\citet{sato2003}, HD47526 (K1III) \\citet{setiawan2003}, HD13189 (K2II-III) \\citet{hatzes2005}, HD11977 (G5III) \\citet{setiawan2005}, Pollux (K0III) \\citet{hatzes2006}, \\citet{reffert2006}, 4UMa (K1III) \\citet{dollinger2007}, NGC2423 No3 and NGC4349 No127 \\citet{lovis2007}, and recently HD17092 (K0III) \\citet{niedzielski2007})\\footnote{For updated information on sub-stellar companions, see http://exoplanet.eu and http://exoplanets.org.}. In addition to searches for extra-solar companions, radial velocity observations prove to be very useful for detecting solar-like oscillations in stars with turbulent atmospheres, such as the dwarf $\\alpha$ Cen A \\citep[e.g.][]{bedding2006}, the subgiant Procyon \\citep[e.g.][]{eggenberger2004,martic2004} and the giant $\\epsilon$ Ophiuchi \\citep[e.g.][]{deridder2006}.  With techniques for accurate radial velocity observations at hand, a survey was started in 1999 to verify whether K giants are stable enough to be used as astrometric reference stars for SIM/PlanetQuest (Space Interferometry Mission) \\citep{frink2001}. This survey contains 179 stars and uses the Coud\\'e Auxiliary Telescope (CAT) at University of California Observatories / Lick Observatory, in conjunction with the Hamilton Echelle Spectrograph. The survey has recently been expanded to about 380 giants and is still ongoing. For the analysis described in the present paper only data from the initial 179 stars are used.  From this survey, companions have been announced for $\\iota$ Draconis \\citep{frink2002} and Pollux \\citep{reffert2006}. Stars with radial velocity variations of less than 20 m\\,s$^{-1}$ have been presented as stable stars by \\citet{hekker2006a}. In addition, some binaries discovered with this survey, as well as an extensive overview of the sample, will be presented in forthcoming papers. %  As almost all of the stars show significant radial velocity variations, we investigate here which mechanism causes these variations. Non-periodic radial velocity variations, of the order of the investigated timescales, are most likely caused by some intrinsic mechanism, while the periodic variability can also be caused by companions. We also investigate the characteristics of these companions.  In Sect. 2, the radial velocity observations are described in detail. In Sect. 3, the relation between the observed radial velocity amplitude and surface gravity is investigated. In Sect. 4, we explore the hypothesis that all periodic radial velocity variations are caused by sub-stellar companions, and we compare the inferred orbital parameters with those obtained for sub-stellar companions orbiting main sequence stars. Our conclusions are presented in Sect. 5.  \\begin{figure} \\begin{minipage}{\\linewidth} \\begin{minipage}{\\linewidth} \\centering \\includegraphics[width=\\linewidth]{8321f1a.eps} \\end{minipage} \\begin{minipage}{\\linewidth} \\centering \\includegraphics[width=\\linewidth]{8321f1b-90.eps} \\end{minipage} \\end{minipage} \\caption{Radial velocity variations as a function of phase for a star (HIP34693) with a highly significant period (top), with its periodogram (bottom). The dashed line in the periodogram indicates the significance threshold.} \\label{vradphase1} \\end{figure}  \\begin{figure} \\begin{minipage}{\\linewidth} \\begin{minipage}{\\linewidth} \\centering \\includegraphics[width=\\linewidth]{8321f2a.eps} \\end{minipage} \\begin{minipage}{\\linewidth} \\centering \\includegraphics[width=\\linewidth]{8321f2b-90.eps} \\end{minipage} \\end{minipage} \\caption{Radial velocity variations as a function of phase for a star (HIP7607) with a period close to the significance threshold (top), with its periodogram (bottom). The dashed line in the periodogram indicates the significance threshold.} \\label{vradphase2} \\end{figure} ",
        "conclusions": "The tight correlation we found between $\\log g$ and half of the peak-to-peak radial velocity variations seems to indicate that a large fraction of the observed radial velocity variations in our sample of K giants is induced by mechanism(s) intrinsic to the stars. We also present evidence that both intrinsic and extrinsic mechanisms play a role. The stars with a significant periodic signal are almost all located above the radial velocity amplitude vs. $\\log g$ relation, but when the periodic signal is removed, the residuals show the same trend as for the non-periodic stars. Furthermore, no correlation is present between the amplitude of the periodic signal and $\\log g$.  Almost all of the lowest $\\log g$ stars show periodic variations. It may be possible that stars with such low surface gravity cannot be constant and that in these dilute atmospheres instabilities can occur very easily, and therefore may be periodic, but not extrinsic.  Based on the evidence that extrinsic mechanism(s) play a role for K giant stars with periodic radial velocity variations we investigated the hypothesis that this periodic signal is caused by the reflex motion of sub-stellar companions orbiting these stars. We presented the characteristics of the orbital parameters of these companions and compared them with the known orbital parameters of sub-stellar companions orbiting F, G and K dwarfs.  About 25$\\%$ of the stars in our sample have radial velocity variations with significant periodicity, and could possibly harbour a sub-stellar companion, while approximately only $8\\%$ of the 1330 F, G and K main sequence stars investigated by \\citet{marcy2005} have a sub-stellar companion. Recently \\citet{johnson2007} and \\citet{lovis2007} showed that the number of companion harbouring stars increases with mass. The giants in the present sample have typical masses between 1 and 4 M$_{\\sun}$ and are in general more massive than the main sequence stars investigated for companions. So, the high percentage is qualitatively in agreement with the results from the literature. Furthermore, \\citet{lovis2007} suggest that more massive stars form more massive planetary systems than lower mass stars. Figs~\\ref{msini} and \\ref{msinilow} show that we find, in general, more massive companions around the more massive K giants than are present around F, G and K dwarfs.  The high percentage of more massive companions around the more massive K giant stars would also be compatible with the core accretion model. This model predicts very few giant planets, but a relatively large number of planets with the mass of Neptune or smaller around M dwarfs \\citep{laughlin2004,ida2005}. This is mainly the result of a much reduced surface density of the disk and the resulting shorter disk evolution timescales compared to those for more massive stars, implying that planet properties vary with the mass of their host stars. In particular, one would expect more sub-stellar companions with higher masses in our giant sample, as is indeed the case if we assume that the companion hypothesis is correct.  The mean metallicity of companion-hosting K giants would be similar to the mean metallicity of the total sample. This is in contrast with the correlation between companion occurrence and metallicity present in F, G and K dwarfs \\citep[e.g.][]{fischer2005}. So far, several groups have investigated the correlation between companion occurrence and metallicity for giants with different results. \\citet{sadakane2005} and \\citet{pasquini2007} agree that companion-ahosting giants are on average not metal-rich, while \\citet{hekker2007} find that giants with announced companions have higher metallicities than their total sample of giants. A detailed discussion about these different results is presented by \\citet{hekker2007}. They conclude that the samples on which the results are based are slightly different and the more metal-rich stars used in their study are lacking in the study by \\citet{pasquini2007}. Furthermore, there is a difference in zero-point correction for the metallicities of announced companion-hosting stars from different surveys. All in all, these inferences are based on small-number statistics and all results have to be taken with caution.  The larger semi-major axis and long periods of the inferred companions orbiting K giant stars compared to companions orbiting dwarf stars are most likely due to the extended atmospheres of K giants. For the eccentricity no significant difference is found in the distribution between companions orbiting dwarfs or giants. Nevertheless, the high number of inferred companions around giants with eccentricities $<$ 0.3 is striking. One could suspect that companion orbits circularise over time and that the companions in circular orbits are older than the eccentric ones, but there is no evidence for this hypothesis.  In principle, nearly sinusoidal radial velocity variations could also be caused by pulsations or spots. Although the periods of the radial velocity variations could well be the rotational periods of the stars, the presence of prominent spots is not very likely. In that case one would also expect photometric variations with periods correlated with the radial velocity variations. From the Hipparcos photometry \\citep{esa1997} such correlations were not found.  In order to distinguish with certainty between companions and pulsations as the cause of the observed radial velocity variations, one needs to perform a spectral line profile analysis. A technique for doing this with very high-resolution spectra ($R \\geq$ 100\\,000) will be presented separately (Hekker et al., in preparation) because the amount of such data at hand today is insufficient to add significantly to the conclusions of this paper.  \\begin{acknowledgement} We thank Debra Fischer, Geoff Marcy and Paul Butler for useful discussions and the development of the instrumentation and software for the determination of the radial velocities at Lick Observatory.  In addition, we thank the entire staff at Lick Observatory for their excellent support. Finally, we would like to thank the anonymous referee for valuable comments and suggestions. \\end{acknowledgement}"
    },
    "0801/0801.0607_arXiv.txt": {
        "abstract": "We present the first attempt to analytically study the nonlinear matter power spectrum for a mixed dark matter (cold dark matter plus neutrinos of total mass $\\sim 0.1$eV) model based on cosmological perturbation theory. The suppression in the power spectrum amplitudes due to massive neutrinos is, compared to the linear regime, enhanced in the weakly nonlinear regime where standard linear theory ceases to be accurate. We demonstrate that, thanks to this enhanced effect and the gain in the range of wavenumbers to which the PT prediction is applicable, the use of such a  nonlinear model may enable a precision of $\\sigma(m_{\\nu, {\\rm tot}})\\sim 0.07$eV in constraining the total neutrino mass for the planned galaxy redshift survey, a factor of 2 improvement compared to the linear regime. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.2602_arXiv.txt": {
        "abstract": "Pulsar winds shocked in the ambient medium produce spectacular nebulae observable from the radio through $\\gamma$-rays. The shape and the spectrum of a pulsar wind nebula (PWN) depend on the angular distribution, magnetization and energy spectrum of the wind streaming from the pulsar magnetosphere, as well as on the pulsar velocity and the properties of the ambient medium. The advent of {\\sl Chandra}, with its unprecedented  angular resolution and high sensitivity, has allowed us not only to detect many new PWNe, but also study their spatial and spectral structure and dynamics, which has significantly advanced our understanding of these objects. Here we  overview recent observational results on PWNe, with emphasis on {\\sl Chandra} observations. ",
        "introduction": "It is generally accepted that all active pulsars lose their spin energy and angular momentum via relativistic winds comprised of relativistic particles and electromagnetic field. Because the relativistic bulk velocity of the wind leaving  the pulsar magnetosphere is obviously supersonic with respect to the ambient medium, such a wind produces a {\\em termination shock} (TS) at a distance $R_{\\rm TS}$ from the pulsar where the bulk wind pressure, $P_w\\sim\\edot/(4\\pi c R_{\\rm TS}^2)$, is equal to the ambient pressure $P_{\\rm amb}$. The TS radius can be estimated as \\be R_{\\rm TS} \\sim \\left(\\edot/4\\pi c P_{\\rm amb}\\right)^{1/2} \\sim 0.05\\, \\edot_{36}^{1/2} P_{\\rm amb,-10}^{-1/2}\\,\\,\\,\\,{\\rm pc}, \\ee where $\\edot = 10^{36}\\edot_{36}$ ergs s$^{-1}$ is the pulsar's spin-down power, and $P_{\\rm amb}=10^{-10} P_{\\rm amb,-10}$ dyn cm$^{-2}$. At the TS, the pulsar wind is being ``thermalized'', and the downstream bulk flow speed becomes subrelativistic \\cite{1984ApJ...283..710K}. As the relativistic particles of the shocked wind move in the magnetic field and ambient radiation field, they emit synchrotron and inverse Compton (IC) radition, which we observe as a {\\em pulsar wind nebula} (PWN). Since the wind is a universal property of any active pulsar, we expect that {\\em all pulsars must be accompanied by PWNe}. Studying PWNe tells us about the properties of pulsar winds and  their parent pulsars, the properties of the ambient medium, and the mechanisms of wind-medium interaction.   The energies of the synchrotron and IC photons span the range from the radio to TeV $\\gamma$-rays. The mere detection of a PWN in some energy band indicates the emission mechanism and the electron energies involved. For instance, detecting a PWN in the X-ray band, where  the synchrotron emission dominates, implies that the wind particles have been accelerated up to $\\gtrsim 100$ TeV energies, either at the TS or on the way to the TS (note that particles with such energies cannot leave the pulsar magnetosphere because of radiative losses), and that the same particles should produce IC emission in the TeV energy range.  Useful information on the pulsar wind and its interaction with the medium is provided by the PWN morphology, which depends on the wind {\\em outflow geometry} and the direction of {\\em pulsar velocity}. If we assume an {\\em isotropic outflow} from a very slowly moving pulsar, the shocked wind is confined between the TS and {\\em contact discontinuity} (CD) spheres, while the shocked ambient medium fills in the space between the CD and the spherical {\\em forward shock} (FS).     If the pulsar moves with a supersonic velocity, $V_p \\gg c_s$, then the TS, CD and FS surfaces acquire a characteristic convex shapes ahead of the moving pulsar but can exhibit  rather different shapes behind it. In particular, the TS acquires a bullet-like shape, with the distance $R_{\\rm TS,h}\\sim (\\edot/4\\pi c P_{\\rm ram})^{-1/2}\\sim 0.04 \\edot_{36}^{1/2}n^{-1/2}(V_p/100\\,{\\rm km\\,s}^{-1})^{-1}$ pc, between the bullet's head ({\\em bowshock}) and the pulsar (cf.\\ eq.\\ 1), where $P_{\\rm ram} = \\rho V_p^2 = 1.67\\times 10^{-10} n (V_p/100\\, {\\rm km\\,s}^{-1})^2$ dyn cm$^{-2}$ is the ram pressure. The numerical simulations \\cite{2003A&A...404..939V, 2005A&A...434..189B} suggest that the shocked wind is channeled into the {\\em tail} behind the TS bullet, confined by the nearly cylindrical CD surface, where it flows with a mildly relativistic velocity. The shocked ISM matter is also stretched along the pulsar trajectory, and it can be seen in spectral lines (e.g., H$_\\alpha$) from the atoms excited at the FS (e.g., \\cite{2002ApJ...575..407C}).  \\begin{figure}[h] \\label{p-pdot} \\hspace{-0.3in} \\includegraphics[height=3.3in,width=3.1in,angle=90]{P_Pdot.pdf} \\caption{$P$-$\\dot{P}$ diagram for the pulsars in the ATNF catalog \\cite{2005AJ....129.1993M}. Pulsars with known X-ray PWNe are marked by circles. } \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=6.5in,angle=0]{tori3_small.pdf} \\caption{X-ray images of PWNe with toroidal components. The image numbers correspond to Tables 1--2. } \\end{figure}  We know from observations of young, subsonically moving pulsars (such as the Crab pulsar), that the PWN is not spherical even in this case, but it rather looks like a torus, sometimes with one or two jets along the pulsar's spin axis. This means that the pre-shock wind is {\\em not isotropic}, but it outflows preferentially in the equatorial plane of the rotating pulsar. Models of such {\\em torus-jet} PWNe have been simulated in \\cite{2004MNRAS.349..779K} and \\cite{2006A&A...453..621D}  (see also Bucciantini, this proceedings). We are not aware of relativistic MHD models for PWNe created by supersonically moving pulsars with {\\em anisotropic} outflows, but we expect that the compact PWN morphology near the pulsar may significantly depend on the orientation of the spin axis with respect to the direction of pulsar's motion, while the distant PWN tail and the FS should not be strongly affected by the wind anisotropy.  As the PWN appearance critically depends on the pulsar's Mach number, ${\\mathcal M}=V_p/c_s$, and $c_s$ in a hot SNR interior is much higher than that in the normal ISM, we expect that young, powerful pulsars, which have not left their host SNRs, generate torus-jet PWNe, while older pulsars would show bowshock-tail PWNe \\cite{2003A&A...397..913V}.      Only a handful of PWNe had been detected before the launch of \\chan, mostly in radio and H$_\\alpha$ observations \\cite{2002ApJ...575..407C}, while X-ray observations of PWNe had been hindered by low angular resolution of X-ray instruments. \\chan, with its unprecedented $\\approx 0.5''$ resolution, has allowed us to reveal the fine structure of the previously known X-ray PWNe, discover many new PWNe, and separate the pulsar and PWN emission in many cases. Some interesting results on X-ray PWNe have also been obtained with {\\sl XMM-Newton}, which lacks the high angular resolution of \\chan\\ but is more sensitive and has a larger field of view. Many exciting results from the \\chan\\ and \\xmm\\ observations of PWNe have been reported in numerous publications, including two reviews \\cite{2006csxs.book..279K, 2006ARA&A..44...17G}. Here, we present an up-to-date overview of X-ray observations of PWNe, including a gallery of spectacular PWN images and a current catalog of these objects, and discuss correlations between various PWN properties as well as the pulsar-PWN correlations. ",
        "conclusions": "The high angular resolutions and sensitivity of \\chan\\ have allowed us to detect many X-ray PWNe and study their structure and spectra. Our current understanding of X-ray PWNe can be briefly summarized as follows.  \\begin{itemize} \\item Most of the detected PWNe are associated with young, powerful pulsars (partly due to selection effect), but some old PWNe have also been detected (e.g., the tail of PSR B1929+10).  \\item The observed PWN morphologies can be crudely classified into the torus-jet and bowshock-tail types, corresponding to sub- and supersonic pulsar motion, respectively. However, the classification is often uncertain, and many morphological features (e.g., in the Vela and Geminga PWNe) remain to be understood.  \\item PWNe radiate up to a few percent of the pulsar spin-down power in X-rays. The X-ray PWN efficiency generally grows with $\\edot$, but the $\\etapwn$-$\\edot$ dependence shows a very large scatter, including some very dim PWNe, $\\etapwn < 10^{-5}$.   \\item The X-ray luminosities of the detected PWNe and their parent pulsars are of the same order of magnitude, with PWNe being, on average, somewhat more luminous.  \\item The photon indices of the PWN spectra are in the range $1\\lesssim \\Gamma \\lesssim 2$. More X-ray efficient PWNe apparently show softer spectra, but this conclusion should be checked in deeper observations of well-resolved PWNe to mitigate the effects of synchrotron cooling.   \\item Spectral maps of the Crab and Vela PWN show strong correlation between the spectral hardness and morphological features. The Crab's spectra immediately downstream of TS are considerably softer than those of the Vela.   \\item X-ray observations of pulsars/PWNe in the vicinity of extended TeV sources show faint extensions of compact PWNe toward the offset centers of the TeV emission, supporting the interpretation of these TeV sources as ``crushed plerions''.   \\item Parsec-scale X-ray tails found behind many pulsars represent ram-pressure collimated, high-velocity flows of relativistic particles, resembling pulsar jets in their properties. \\end{itemize}  Notwithstanding the considerable progress in observations and modeling of PWNe, there still remains a nubmer of important problems to solve. Here we mention a few of them.  \\begin{itemize} \\item What is the origin of the diverse, often very complex, PWN morphologies? For instance, what is the nature of the outer arc, inner jets, and bar in the Vela PWN, the loops in the 3C\\,58 PWN, the shell and axial tail in the Geminga PWN?   \\item Where and how the X-ray emitting particles are accelerated? Does the acceleration occur in the preshock wind or at the TS? Are the winds composed of electrons/positrons or they have some nucleonic component?   \\item What are the physical parameters that determine the X-ray PWN efficiency, in addition to $\\edot$? Why don't we see any PWN around some pulsars? Is it because of some special properties of the pulsar wind or the ambient medium?    \\item Why are the particle spectra so hard in some PWNe (e.g., Vela) and much softer in others (e.g., Crab)? Is such a difference caused by different properties of the parent pulsars' winds (e.g., magnetization), or different efficiencies of particle acceleration between the magnetosphere and the TS? What is the reason for the apparent efficiency-hardness correlation?   \\item Can all the Galactic extended TeV sources with luminous pulsars nearby be interpreted as relic TeV plerions or there are indeed ``dark accelerators''? Are these sources indeed ``crushed PWNe'' formed by the passage of the inverse SNR shock? Can the huge sizes of the TeV plerions and the large offsets from the pulsars be reconciled with the ``crushed PWN'' scenario? Is the TeV radiation due to the IC scattering or the pion decay? If the former, are the seed photons supplied by the CMB or IR radiation from nearby star-forming regions or dust clouds?  \\item What are the flow velocities and magnetic fields in the long pulsar tails? Which role does the magnetic field play in the tail collimation and what is the magnetic field  topology ? How do these collimated flows decelerate and cool? \\end{itemize}  To answer these questions, new observations are badly needed, both in X-rays and other wavelengths, as well as theoretical work and numerical modeling. Particularly important would be to take full advantage of the outstanding \\chan\\ capabilities as long as it is alive because no X-ray observatory with such high angular resolution is expected in the foreseeable future.  \\begin{theacknowledgments} We thank Koji Mori for providing the {\\sl ASCA} images of Vela X. This work was partially supported by Chandra award AR5-606X. \\end{theacknowledgments}"
    },
    "0801/0801.2328_arXiv.txt": {
        "abstract": "Shell galaxies are widely considered the debris of recent accretion/merging episodes. Their high frequency in low density environment suggests that such episodes could be among the driver of the early-type galaxy secular evolution. We present far and near UV  (FUV and NUV  respectively hereafter)  {\\it GALEX} photometric properties  of a sample of shell galaxies. ",
        "introduction": "In a hierarchical evolutionary scenario, galaxies experience accretion/merging events during their lifetime. While early-type  galaxies in nearby clusters appear (homogeneously) old, the field early-type galaxy population seems to contain genuinely, recently {\\it rejuvenated} objects \\citep[see e.g.][]{Clemens06}. Early-type galaxies showing fine structure, like shells, occupy a special position since they are believed to fill the gap between ongoing mergers and normal elliptical galaxies. The UV emission is crucial to test whether these galaxies do host ongoing/recent star formation processes and study  their distribution across the galaxy. We present  new  {\\it GALEX} observations of  three  shell galaxies NGC 1210, MGC -05-07-1 (GI04-0030-0059 PI D. Bettoni) and NGC 5329 (from archive) in addition to those analyzed in \\citet{Rampazzo07} which we use as baseline for our preliminary conclusions. ",
        "conclusions": "Table 1 summarizes the journal and the basic results of the {\\it GALEX} observations.  \\begin{table*} \\scriptsize \\caption{Journal of the {\\it GALEX} observations} \\begin{tabular}{lcclllll} \\hline \\multicolumn{1}{l}{Name}& \\multicolumn{1}{c}{NUV}  & \\multicolumn{1}{c}{FUV}  & \\multicolumn{1}{l}{P.I.}& \\multicolumn{1}{c}{m$_{NUV}^{tot}$}& \\multicolumn{1}{c}{m$_{FUV}^{tot}$} & \\multicolumn{1}{c}{FUV-NUV}\\\\ \\multicolumn{1}{c}{}& \\multicolumn{1}{c}{exposure [sec]} & \\multicolumn{1}{c}{exposure [sec]} & \\multicolumn{1}{c}{}& \\multicolumn{1}{c}{}& \\multicolumn{1}{l}{}& \\multicolumn{1}{l}{}\\\\ \\hline MCG-05-07-1  &1510 & 1531   & D. Bettoni & 18.67$\\pm$0.07 & 19.76$\\pm$0.07 & 1.11$\\pm$0.10  \\\\ NGC 1210  &1558 & 1608 &  D. Bettoni &17.14$\\pm$0.02&20.08$\\pm$0.07&2.95$\\pm$0.07  \\\\ NGC5329  &3889  &2666   &   MIS&18.20$\\pm$0.02&20.20$\\pm$0.04&2.02$\\pm$0.05 \\\\ \\hline \\end{tabular} \\label{table1} \\medskip{FUV - NUV AB magnitudes have been corrected for Galactic extinction.} \\end{table*} Smoothed images and 2D colour maps are shown in Figure 1. FUV emission in NGC 5329 is present only in the central part of the galaxy. In NGC 1210 and MGC-05-07-01 the FUV is quite strong both in the polar ring of the former galaxy and in the debris systems, residual of the accretion events of  both galaxies. The (FUV-NUV) colour in the  north tail of  NGC 1210  ($\\sim$0.31, $\\sim$0.34, $\\sim$0.96), in the  polar ring   ($\\sim$0.36, $\\sim$0.73) and in the nucleus ($\\sim$1.11) of MGC-05-07-01 are quite blue.  \\citet{Neff05} show that tail/bridges produced by interaction may have similar colour indicative of a recent star formation (200-300 Myr). The age of 1-3 Gyr estimated by Whitmore et al. (1987),   responsible for the present structure of MCG-05-07-01, is consistent with age indication coming from (UV - optical) colours. \\begin{figure} \\centerline{\\psfig{figure=marino1.ps,width=2.8cm}\\psfig{figure=marino2.ps,width=2.8cm}}% \\centerline{\\psfig{figure=marino3.ps,width=2.8cm}\\psfig{figure=marino4.ps,width=2.8cm}}% \\centerline{\\psfig{figure=marino5.ps,width=2.8cm}\\psfig{figure=marino6.ps,width=2.8cm}}% \\caption{{\\it GALEX} FUV (left panels), NUV (middle panels)  images  (5\\arcmin $\\times$ 5\\arcmin) and (FUV-NUV) colour maps (right panels) of MCG-05-07-1 (top row), NGC 1210 (mid row) and NGC 5329 (bottom row).  FUV  and NUV  images are smoothed using {\\it Asmooth} \\citep{Ebeling06}.} \\label{fig1} \\end{figure}"
    },
    "0801/0801.3432_arXiv.txt": {
        "abstract": "Difference imaging provides a new way to discover gravitationally lensed quasars because few non-lensed sources will show spatially extended, time variable flux.  We test the method on lens candidates in the Sloan Digital Sky Survey (SDSS) Supernova Survey region from the SDSS Quasar Lens Search (SQLS) and their surrounding fields.  Starting from 20768 sources, including 49 SDSS quasars and 36 candidate lenses/lensed images, we find that 21 sources including 15 SDSS QSOs and 7 candidate lenses/lensed images are non-periodic variable sources.  We can measure the spatial structure of the variable flux for 18 of these sources and identify only one as a non-point source.  This source does not display the compelling spatial structure of the variable flux of known lensed quasars, so we reject it as a lens candidate.  None of the lens candidates from the SQLS survive our cuts.  Given our effective survey area of order 0.71 square degrees, this indicates a false positive rate of order one per square degree for the method.  The fraction of quasars not found to be variable and the false positive rate should both fall if we analyze the full, later data releases for the SDSS fields.  While application of the method to the SDSS is limited by the resolution, depth, and sampling of the survey, several future surveys such as Pan-STARRS, LSST, and SNAP will avoid these limitations. ",
        "introduction": "Gravitational lensing has many applications, from exoplanet searches to large scale structure (see the reviews by Kochanek, Schneider, and Wambsganss in the Saas Fe lectures, 2006).  Galaxy scale lenses can be used to study the dark matter mass profile of the lens galaxy \\citep[e.g.][]{Kochanek91,Rusin05,Koopmans06,Jiang07}.  When the background source is a quasar, microlensing by individual stars in the lens galaxy can be used to probe the structure of the quasar's accretion disk \\citep{Poindexter07,Morgan07} and broad line regions \\citep{Eigenbrod07}.  Galaxy scale lenses can also constrain cosmological parameters \\citep{Refsdal64,Oguri07}.  Unfortunately, most applications require large samples of gravitational lenses to be competitive with other methods, while fewer than a hundred lensed quasars are known.  Moreover, these lenses were discovered using different methods with different biases, which is especially problematic if homogeneous samples with well-understood selection functions are needed \\citep[see][]{Kochanek04}.  The known lenses were found by their morphological structure or the presence of higher redshift features in the spectrum of a lower redshift galaxy.  Morphological surveys examine optical or radio images of quasars for evidence that they are lensed.  This works best for point-like sources like optical quasars, and flat spectrum radio sources.  The two largest searches for lensed AGN are the Cosmic Lens All-Sky Survey of flat spectrum radio sources \\citep{Myers03,Browne03} and the SDSS Quasar Lens Search \\citep[SQLS;][]{Oguri06}.  The radio samples are limited by the number of sufficiently bright radio sources and the difficulties in obtaining source redshifts.  The optical quasar samples suffer from confusion from stars and galaxies and the effects of color changes, either from the starlight or dust in the lenses, on sample selection.  The spectroscopic method made several serendipitous discoveries of lensed quasars, such as Q2237+0305 \\citep{Huchra85} and SDSSJ090334.92+502819.2 \\citep{Johnston03}.  Following theoretical investigations \\citep[e.g.][]{Kochanek92,Miralda92,Mortlock00,Mortlock01}, the SLACS survey \\citep{Bolton04,Bolton06} used the spectroscopic method on massive early-type galaxies in the SDSS to identify $\\sim$50 candidate lensed starforming galaxies, and confirmed 19 of 29 using Hubble Space Telescope images.  The spectroscopic method is limited by its low yield (about 1 lens candidate per 1000 luminous red galaxy spectra) and biases in its mass range (due to the size of the spectroscopic aperture).  Moreover, by selecting targets based on the properties of the lens galaxy, these lenses are mainly useful for studying the lens galaxies rather than for cosmology \\citep[see][]{Kochanek04}.  \\citet{Kochanek06} proposed a new method to find lensed quasars based on difference imaging.  Difference imaging, also known as image subtraction, is a way to measure the variable intensity in a region of sky \\citep{Tomaney96,Alard98,Alard00}.  It has been used to search for a broad range of variable sources, such as planets transiting stars \\citep[e.g.][]{Hartman04}, microlenses \\citep[e.g.][]{Alcock99,Wozniak00}, and supernovae \\citep[e.g. the Sloan Digital Sky Survey II Supernova Survey,][]{Sako07}.  In difference imaging, a reference image is made by averaging a set of the best images for a field.  Then, for each epoch of observation, a difference image is created by subtracting the reference image convolved to the PSF and flux scale of that epoch.  If a source is variable, there will be a flux residual in the difference image corresponding to the variability of the source.  In difference images, lensed quasars are recognizable because they consist of multiple variable images that are close together.  Since most quasars are variable, with 60\\% having $\\sigma_g \\ge 0.05$ variability over two years \\citep[e.g.][]{Sesar07}, each of these images is variable.  The level of variability is then enhanced by microlensing of the quasar images by the stars in the lens galaxy \\citep[see the review by][]{Wambsganss06}.  Thus, lensed quasars look like compact clusters of variable sources or extended variable sources, which allows us to easily search for lensed quasars in difference images.  The number of false positives from pairs of variable stars, pairs of quasars (related or not), variable star-quasar blends, and supernovae-AGN pairs is expected to be low \\citep{Kochanek06}.  Candidates found this way can be confirmed by light curve analysis, since each image will have the same intrinsic variability, but with time delays between the images and some additional uncorrelated variability from microlensing by the stars of the lens galaxy \\citep[e.g.][]{Pindor05}.  Difference imaging can also resolve variable lensed sources blended with non-variable sources, such as the lens galaxies.  This is important because, as we search for fainter lensed quasars, contamination from the lens galaxy becomes a steadily greater problem.  Tests of image subtraction on the lens Q2237+0305 show that the method can find quasars even when they are buried by the flux of an extraordinarily bright foreground galaxy \\citep{Kochanek06}.  The forthcoming, large scale synoptic surveys, like LSST \\citep{Tyson02} and Pan-STARRS \\citep{Kaiser04}, will be ideal for this method, since they will cover large areas of the sky, sample variable sources frequently, and have deep magnitude limits.  In the meantime, the SDSS II Supernova Survey \\citep{Frieman07,Sako07} is the best survey currently available for our goals.  Intended to find supernovae for dark energy studies, the Supernova Survey repeatedly images SDSS Stripe 82, a 2.5\\degr~wide swath along the celestial equator, covering 300 square degrees and stretching across the Southern Galactic Cap from right ascension 300\\degr~to right ascension 60\\degr~\\citep{Frieman07}.  Ten to twenty public epochs were available for a given field when we started this project, with the number rising to about forty in the most recent releases \\citep{Adelman-McCarthy07}.  We applied the method outlined in \\citet{Kochanek06} to the SDSS Supernova Survey fields that contained 26 previously identified candidates for gravitational lenses, with 39 total components, from the initial SQLS candidate list.  These include 15 candidates with 24 components in the main statistical sample \\citep{Inada07}, and 11 additional candidate systems, with 15 images, that were outside the final selection criteria.  These candidates are listed in Table~\\ref{table:InadaLenses}.  We also checked to see if any pair of lens images in a candidate lens appeared on the difference images, since the images can be well-separated.  The lens candidates had already been rejected for other reasons, such as different spectral energy distributions for the postulated lens images, or the lack of a lens galaxy.  Unfortunately, there are no confirmed lensed quasars in the Supernova Survey region.  We had two goals.  First, to use the variability method as an independent check of the SDSS candidates.  Second, to get a sense of the false positive rate from any other variable sources in the fields.  We outline the method, our approach to source selection, and the results when applied to targets in \\S 2.  We discuss the variable sources in the lens candidate fields in \\S 3, including quasars (\\S 3.2) and an extended variable source (\\S 3.3).  Finally, we conclude in \\S 4 by discussing prospects for finding gravitational lenses in future surveys with difference imaging. ",
        "conclusions": "We searched for sources with the spatially extended variability characteristic of gravitationally lensed quasars by applying difference imaging to the Sloan Supernova Survey fields of the SQLS lens candidates and their surroundings.  We found one source, SDSSJ213245.25+000146.5, that passed basic criteria for variability, non-periodic variability, and which appeared extended on its ``absdiff'' image.  However, it did not show the compelling structure seen in difference images of known lensed quasars \\citep{Kochanek06} -- it only barely passed the effective radius criterion, suggesting it is not likely to be a true lensed quasar.  Furthermore, it was not among the SQLS lens candidates in the SDSS Supernova Survey region.  Our criteria successfully rejected the SQLS candidates, as had \\citet{Inada07} using other criteria.  If used with other lens search techniques, like color or morphological criteria \\citep{Oguri06}, difference imaging could substantially reduce the number of false positives, although it would only work for quasars that are variable during the observation period.  Only a third of the quasars detected by ISIS passed our quasar variability tests and could be analyzed for spatial structure in the difference images.  More generally, as a variability survey for lensed quasars covering 0.71 square degrees, we successfully identified one third of the SDSS quasars in the field and had no false positive detections of lenses.  In many ways, our experiment was limited by the data.  First, there are no known lensed quasars in Stripe 82, so we could evaluate no examples of success or set any limits on the role of false negatives.  Furthermore, when we carried out this experiment, the Supernova Survey had few public ($\\la 25$) epochs available.  Since some epochs will be dropped because of bad seeing, or because image subtraction failed, there are some fields where there were too few enough epochs to extract light curve statistics.  Similarly, the released epochs did not cover a large enough time interval.  In some cases (e.g. SDSSJ213245.25+000146.5 in Figure~\\ref{fig:LCs}), all the non-rejected epochs spanned only a few months.  The performance of our approach would improve markedly with survey data spanning several years and with many more epochs.  Additionally, the relatively poor resolution of the SDSS survey is not ideal for identifying quasar lenses.  This is also true for color and morphological techniques, such as in \\citet{Inada07}, where candidate lenses with image separations of less than 1\\farcs0 were rejected.  With better resolution and more epochs, false positives in the method would decrease.  If we examined every SDSS quasar in the supernova region, we would expect to find of order 1000 were variable with $\\sim 50$ false positives if we assume no improvement in the false positive rate with the addition of more epochs.  In reality, we should have a marked reduction in the rate because we could focus on epochs with better seeing and add in the variability information from the ugiz bands as well.  While these shortcomings will be partly solved by the full SDSS Supernova Survey data set, the real future for the method lies with ground-based surveys like LSST \\citep{Tyson02} and Pan-STARRS \\citep{Kaiser04} and proposed surveys from space like the SuperNova Acceleration Probe \\citep[SNAP,][]{Aldering02}.  Pan-STARRS and LSST would repeatedly survey very large areas ($10^3$--$10^4$ rather than $10^2$ square degrees) with improved resolution ($0\\farcs5$--$1\\farcs$ rather than $>1\\farcs0$) and depth ($\\simeq 24$~mag rather than $\\simeq 22$~mag).  \\citet{Kochanek06} estimated that LSST can discover $\\sim 10^3$ lensed quasars with V$<23$.  The space based surveys would cover far less area (15~square degree for SNAP) but with vastly improved resolution ($\\simeq 0\\farcs05$) and depth ($\\simeq 28$~mag).  Since few lenses have separations this small (1\\farcs5 is the expected median, e.g. \\citet{Mitchell05}), any lens identified by variability is easily confirmed by its morphology.  More important for the space missions is that their great depth allows searches for other lensed, time variable sources such as the supernovae themselves."
    },
    "0801/0801.3268_arXiv.txt": {
        "abstract": "We apply the K-correction to the black hole LMXB GX 339-4 which implies $M_{X}\\geq 6 M_{\\odot}$ by only assuming that the companion is more massive than $\\sim 0.17M_{\\odot}$, the lower limit allowed by applying a 'stripped-giant' model. This evolutionary model successfully reproduces the observed properties of the system. We obtain a maximum mass for the companion of $M_2 \\leq 1.1 M_{\\odot} $ and an upper limit to the mass ratio of $q(=M_2/M_{X})\\leq0.125$. The high X-ray activity displayed by the source suggests a relatively large mass transfer rate which, according to the model, results in $M_2 \\ga 0.3M_{\\odot}$ and $M_{X} \\ga 7M_{\\odot}$. We have also applied this scenario to the black hole binary XTE J1550-564, which has a similar orbital period but the donor is detected spectroscopically. The model successfully reproduces the observed stellar parameters. ",
        "introduction": "\\label{introduction} Low Mass X-ray Binaries (LMXB) are interacting binaries harbouring a neutron star (NS) or a black-hole (BH) which accretes matter from a low mass star. A subgroup, the so-called X-ray transients (SXTs), provide a unique opportunity to set dynamical constraints to the the masses of both NSs and BHs because the spectrum of the companion dominates during the quiescent (X-ray off) states. In particular, we can empirically determine the mass function of the system which is usually expressed as: \\begin{equation} \\label{mf} f(M)=\\frac{PK_2^3}{2 \\pi G}=\\frac{M_X\\sin{i}^3}{(1+q)^2} \\end{equation} where $P$ is the orbital period, $K_2$ the semi-amplitude of the velocity curve of companion star, $M_X$ the mass of the compact object, $i$ is the orbital inclination of the systems and $q=\\frac{M_2}{M_X}$  the mass ratio. Therefore, it is possible to establish a secure lower limit to the mass of the compact object by only measuring $P$ and $K_2$. Although this method has successfully been applied to $\\sim 20$ BH X-ray binaries (see \\citealt{C2006} for details), they only represent the 'peak of the ice-berg' of a large population of $10^8-10^9$ stellar-mass BH present in our galaxy (e. g. \\citealt{BB94}). Most quiescent SXTs are optically faint, with $R\\geq 20$, and hence dynamical constraints are usually affected by large errors. Setting new dynamical constraints and refining the existing mass solutions are the only ways to get new insights into the fundamental properties of stellar-mass BHs.\\\\ GX 339-4 (V821 Ara) was discovered by OSO 7 in 1972 (\\citealt{Ma73}). Since then, several X-ray outburst (e. g. 4 in the last decade) have been observed, and all the X-ray states have been detected in this system (See \\citealt{Mvdk97} for the intermediate state and \\citealt{Mi91} for the Very High state). Moreover, \\cite{Cor00} detected a compact jet emission from this system, showing that it is a microquasar. GX 339-4 was early proposed as a BH candidate based on its X-ray properties (\\citealt{S79}) but the spectral features of the companion star could not be detected even during 'X-ray off' states (\\citealt{Tariq01}), preventing a dynamical confirmation (i. e. $M_X > 3M_{\\odot}$). Only during the 2002 outburst \\cite{hynes03} reported the first detection of the donor star thanks to the discovery of NIII/CIII Bowen emission lines arising from the irradiated companion star. The lines are very sharp and swing with a velocity semi-amplitude of $K_{em}$=$317 \\pm 10 $ km s$^{-1}$. These authors also reported an orbital period of $P=1.7557\\pm0.0004$ days which was later confirmed by \\cite{lyc06}. % The combination of both, radial velocity of the companion and orbital period yields $f(M)=5.8 \\pm 0.5 M_{\\odot}$ (and hence $M_X \\geq 5.3M_{\\odot}$), which represents the first dynamical proof for a BH in GX 339-4. Note that narrow high-excitation emission lines originating from the companion star have also been detected in many others LMXBs (e.g. \\citealt{C04}) and demonstrate that this is a common feature in LMXBs with high X-ray activity (i. e. steady systems and transients during outbursts). However, the NIII/CIII emission lines are excited on the inner hemisphere of the donor star by the Bowen mechanism/photoionization (\\citealt{MCT75}) and only provides a lower limit ($K_{em}$) to the true $K_{2}$-velocity of the donor. In \\cite{MCM05} we tackle this problem in a general approach by modeling the deviation between the reprocessed light-center and the center of mass of a Roche lobe filling star (the so-called 'K-correction') including screening effects by a flared accretion disc. In this paper we compute the K-correction for GX 339-4 which provides a more restrictive lower limit to the mass of the BH in this binary. In sections \\ref{companion} and \\ref{observations} we discuss the nature of the companion star and show that all the known observables can be explained by an scenario in which a 'stripped-giant' companion transfers mass onto a BH. We also apply this evolutionary model to XTE J1550-564 which shows observational properties very similar to GX 339-4. The stellar parameters of the companion in this system are in excellent agreement with a stripped-giant scenario. ",
        "conclusions": "We have applied the K-correction to the BH LMXB GX 339-4. By considering the limit case where the emission line is formed at the limb of the irradiated region of the companion we derive a solid lower limit to $M_{X}\\geq 6 M_{\\odot}$ including the error bars in $f(M)$. Here we have only assumed $M_2\\geq0.166 M_{\\odot}$, the lower limit allowed by the stripped giant model. We find that the stripped-giant evolutionary model explains the non-detection of the companion by SFC01 and the X-ray behaviour of the source. In particular, we propose $M_{2}\\ga0.3M_{\\odot}$, for which we predict  $\\dot{M_{2}}$ large enough to explain the frequent X-ray outbursts displayed by this source. This limit results in $M_{X}\\ga7M_{\\odot}$. From the maximum mass solution we find $q\\leq0.125$. On the other hand, we have also shown that the stripped-giant model successfully explains the observable properties ($M_2$, $R_2$, $T_{eff}$ and $\\dot{M}_2$) of the akin LMXB XTE J1550-564.\\\\"
    },
    "0801/0801.4448_arXiv.txt": {
        "abstract": "At the surface of the Sun acoustic waves appear to be affected by the presence of strong magnetic fields in active regions. We explore the possibility that the inclined magnetic field in sunspot penumbrae may convert primarily  vertically propagating acoustic waves into elliptical motion.  We use helioseismic holography to measure the modulus and phase of the correlation between  incoming acoustic waves and the local surface motion within two sunspots. These correlations are modeled assuming the surface motion is elliptical, and we explore the properties of the elliptical motion on the magnetic field inclination.  We also demonstrate that the phase shift of the outward  propagating waves is opposite to the phase shift of the inward propagating waves in stronger, more vertical fields, but similar to the inward phase shifts in weaker, more inclined fields. ",
        "introduction": "\\label{introduction} Helioseismology uses the observed solar surface acoustic wavefield to construct images of the subsurface structure of the Sun. Of particular interest has been the three-dimensional (3D) modeling of time-distance \\cite{DJHP93} observations of travel-time shifts to deduce the subsurface structure of active regions \\cite{KDS00,ZK03}.  Assuming the travel-time shifts are due to perturbations in the sound speed below the spot, a general consensus in the models has emerged consistent with sound-speed reductions (relative to surrounding quiet Sun) near the surface ($\\lesssim 4$ Mm) and enhancements up to 15 Mm below sunspots \\cite{KDS00,CBK06}.  Using ring diagram analysis \\inlinecite{BAB04}  also find a lower sound speed immediately below the surface and an increase in the sound speed below 7 Mm.  A comparison between Fourier-Hankel analysis and time-distance results by \\inlinecite{B97} first prompted caution in the interpretation of acoustic oscillation signals within sunspots. The influences of strong surface magnetic fields have not been explicitly included in most helioseismic models of active regions. \\citeauthor{LB05ii} (2005a) and \\citeauthor{LB05i} (2005b) have shown that helioseismic phase shifts observed with helioseismic holography vanish below a depth of about 5 Mm, when a surface (``showerglass'') phase shift based on photospheric magnetic flux density, is removed from the data.  Other evidence supports the possibility of strong near-surface contributions to the helioseismic phase (or travel-time) shifts. These include the possible contamination of surface perturbations into the 3D inversions (e.g. \\opencite{K06}; \\opencite{CR07}). It has also been shown that the reduction of $p$-mode amplitudes in magnetic regions can cause travel-time shifts \\cite{RBDTZ06}. The suppression of sources of wave excitation within sunspots can also produce measurable shifts \\cite{HCRB07}.  \\inlinecite{SBCL05} and \\inlinecite{SBC07} have found that phase shifts obtained from seismic holography in sunspot penumbrae vary with the line-of-sight angle (from vertical) as projected into the plane containing the magnetic field and the vertical direction.  A similar effect has also been noted by \\inlinecite{ZK06} with time distance measurements. \\inlinecite{SBC07} find that the  effect is dependent upon the strength and/or inclination of the magnetic field. In the penumbra the magnetic field strength decreases as the magnetic field angle from vertical increases, hence the two properties of the magnetic field cannot be extricated. The phase variation with line-of-sight viewing angle is demonstrated most substantially at frequencies around 5 mHz with a strong, almost vertical magnetic field close to the umbra. \\inlinecite{SBC07} also find that the total variation across all lines-of-sight increases with temporal frequency, particularly in the stronger fields in the penumbrae.  Mode conversion of the acoustic waves in the near surface has been explored as the physical cause of the observed absorption of acoustic waves by sunspots.  A fast acoustic wave, propagating towards the surface from the interior, encounters the depth at which the Alfv\\'en speed is equal to the sound speed ($a\\approx c$) which is typically close to the surface in a  sunspot. Under these conditions it is able to transmit to a slow acoustic mode and convert to a fast magnetic mode \\cite{C05}. \\inlinecite{CC03} explore the mode conversion in two-dimensions with a uniform \\emph{inclined} magnetic field relevant to sunspot penumbrae. The inclination of the magnetic field is found to have a significant dependence on the likelihood of conversion  and fits extremely well with the analysis of \\inlinecite{B95} \\cite{CCB03}. Further work \\cite{C05,SC06} using ray theory has since established that it is the angle between the acoustic wave path and the magnetic field (the `attack angle') which is the crucial factor inducing conversion. With a wide attack angle at the $a\\approx c$ level there is maximum conversion from a fast acoustic to a fast magnetic mode. A fine attack angle encourages  transmission  to a slow acoustic mode, which is guided `up' the magnetic field lines to observational heights in the atmosphere. A consequence of mode conversion may be the observational signature of elliptical motion in regions of inclined magnetic field.  The aim of this paper is to model the observations of two sunspots previously analyzed by \\inlinecite{SBC07} in order to determine the properties of velocity ellipses consistent with the data. These models are based on a least-squares-fit of the phase and modulus information of the local ingression control correlation. We explore the properties of these ellipses, observed with waves at different temporal frequencies, as functions of the magnetic field inclination angle. We also examine the variation with line-of-sight angle of the phase shifts in the outgoing waves, using the local egression control correlation, to assess the relation of the phase shift variations between incoming and outgoing waves.  In the following sections we describe the data (Section \\ref{obs}), give an outline of the helioseismic holography technique used (Section 3), describe the results (Section 4) and discuss our results in the context of mode conversion  (Section 5). ",
        "conclusions": "Mode Conversion predicts (among other things \\cite{C07}) that when the attack angle is small most of the observable energy will be in the slow acoustic mode, and when the attack angle is large most of the observable energy will be in the fast magnetic mode. In this case we will be seeing the line-of-sight effect of a combination of waves coming from all directions impinging on magnetic field with a particular orientation. The observations by MDI consist of line-of-sight Doppler signatures of the surface motion which are presumably caused mainly by pressure perturbations. This means that we would expect to be observing only the slow acoustic mode, however it is possible that we are observing a combination of the acoustic and magnetic modes.   The dependence of the phase shift of the observed ingression correlations with azimuthal angle around sunspot penumbrae, as viewed from different observational vantages, shows that the incoming phase shifts must be (at least partly) photospheric in origin and are influenced by the presence of inclined magnetic fields \\cite{SBCL05}. Analysis of the variation of both the amplitude and phase of the surface velocities provides an opportunity to characterize the magnetically influenced acoustic signature as an ellipse with properties determined by the magnetic fields. We found that the ellipses are either nearly vertical (for weaker, more inclined fields) or generally directed \\emph{away} from the magnetic field direction (for stronger, more vertical fields).  Largely consistent results for two  active regions,  AR9026 and AR9057, are found. Some properties of the surface ellipses, e.g. their inclinations, are different for the two sunspots. Some of this variation may be due to  differences in the field properties. For example, the field in AR9057 is on average $\\sim 15\\%$ weaker than AR9026.  Fits of the elliptical motion in Figures~\\ref{ellipse9026} and \\ref{ellipse9057} depend critically on the correlation modulus which is prone to systematic uncertainties. But the trend is that a stronger, less inclined magnetic field produces elliptical motion with smaller amplitude, eccentricity, and deviation angle, and a larger inclination from vertical. These trends exist for both spots and, largely, at all frequencies. The  shorter semi-major axis at strong magnetic field strengths is consistent with previous knowledge of surface acoustic amplitude suppression in magnetic fields. It is curious to note, however, that at 4 mHz the amplitude is consistently smaller than even 3 mHz.   These are the first results to estimate the behaviour of the surface velocity ellipse at the photosphere within sunspots. The results do not immediately suggest an observation of the slow wave as shown by \\inlinecite{C05} or \\inlinecite{SC06}, but are consistent with their expectations of the behaviour of slow waves at this height in the atmosphere. Mode conversion theory states the alignment is dependent upon $a^2/c^2$  which at the observational heights of $\\sim 200$km of the atmosphere may not be large enough to invoke alignment.  Since the ray analysis is somewhat unrealistic we would expect to observe a combination of fast and slow waves at the surface, which will contribute to a clouded view of the surface velocities. \\inlinecite{RSWS07} are currently exploring the possibility that these apparent surface effects are due to the changes in the radiative transfer within active regions and the formation height of the observational Ni 678 nm line. This explanation requires an absorption mechanism, or else some other means of producing a difference between the amplitudes of upward and downward propagating waves. Thus mode conversion may still be important in this proposed mechanism. A test of the mechanism proposed by Rajaguru et al. (2007) would be to repeat the observations performed here in a magnetically insensitive line, where the proposed radiative transfer effects would not be present. In terms of mode conversion, it is suggested that the main effect occurs along the bright radial filaments of the interlocking comb structure as presented in the penumbral models of \\inlinecite{WTBT04}. However,  observational helioseismic spatial resolution cannot currently resolve this.   This is also the first time that the variation of the phase of the local egression correlation has been analysed in the penumbra. It is curious that the egression correlation shows a reverse dependence when the magnetic field is weak and highly inclined.  This is evidence of a reverse ingression dependence on the line-of-sight, but further investigation is required to understand the behavior at high frequencies in the weaker, more inclined fields.  \\vspace{2cm}  This work was supported in part by the \\textit{European Helio- and Asteroseismology Network} (HELAS)."
    },
    "0801/0801.2974_arXiv.txt": {
        "abstract": "General Relativistic (GR) Magnetohydrodynamic (MHD) simulations of black hole accretion find significant magnetic stresses near and inside the innermost stable circular orbit (ISCO), suggesting that such flows could radiate in a manner noticeably different from the prediction of the standard model, which assumes that there are no stresses in that region.  We provide estimates of how phenomenologically interesting parameters like the ``radiation edge\", the innermost ring of the disc from which substantial thermal radiation escapes to infinity, may be altered by stresses near the ISCO.  These estimates are based on data from a large number of three-dimensional GRMHD simulations combined with GR ray-tracing.  For slowly spinning black holes ($a/M<0.9$), the radiation edge lies well inside where the standard model predicts, particularly when the system is viewed at high inclination.  For more rapidly spinning black holes, the contrast is smaller.  At fixed total luminosity, the characteristic temperature of the accretion flow increases between a factor of $1.2-2.4$ over that predicted by the standard model, whilst at fixed mass accretion rate, there is a corresponding enhancement of the accretion luminosity which may be anywhere from tens of percent to order unity.  When all these considerations are combined, we find that, for fixed black hole mass, luminosity, and inclination angle, our uncertainty in the characteristic temperature of the radiation reaching distant observers due to uncertainty in dissipation profile (around a factor of 3) is {\\it greater} than the uncertainty due to a complete lack of knowledge of the black hole's spin (around a factor of 2) and furthermore that spin estimates based on the stress-free inner boundary condition provide an upper limit to $a/M$. ",
        "introduction": "Recent years have seen a rapid growth in the specificity and detail with which we attempt to describe accretion onto black holes.  In the early days of the subject three decades ago, we were content with the simplest of models \\cite[the `standard' model, ][]{Novikov:1973}, hereafter NT, in which the accretion flow was assumed to be time-steady, axisymmetric, geometrically thin, and following circular Keplerian orbits at all radii outside the innermost stable circular orbit (the ISCO, located at radius $r_{ms}$).  Relying on heuristic arguments framed in a purely hydrodynamic context \\citep{Page:1974}, it was further assumed that all stresses ceased inside the ISCO, so that the inner edge of the disc could be described as falling precisely at that radius.  Dimensional analysis was the basis for any link between the inter-ring stresses essential to accretion and local physical conditions \\citep{Shakura:1973}.  Today, there is tremendous effort to obtain and interpret direct observational diagnostics of the innermost parts of accretion flows onto black holes.  Apparently relativistically-broadened Fe K$\\alpha$ profiles can be discerned in numerous examples \\citep{Reynolds:2003}. Extensive efforts are made to fit detailed disc spectral models to observed continuum spectra \\citep{Gierlinski:2004,Davis:2005,Shafee:2006a,Shafee:2006b,Hui:2007}. These spectral fits can be used to infer the mass of the central black hole; both methods may be used to constrain the spin of the black hole \\citep{Makishima:2000,Miller:2004,Miniutti:2004,Brenneman:2006,Shafee:2006a, Shafee:2006b}. Some hope to use relativistic fluid dynamics to define normal modes of disc oscillation that could then also be used to constrain the central black hole's spin \\citep{Wagoner:2001,Rezzolla:2003,KatoS:2004}.  At the same time, there has also been much progress on the theoretical side. A strong consensus now supports the idea that accretion stresses are the result of correlated MHD turbulence, driven by a pervasive magneto-rotational instability \\citep{Balbus:1998}. Numerical simulations developing this idea have given us a detailed and quantitative description of the vertical profiles of pressure and density inside discs \\citep{Miller:2000,Hirose:2005,Hirose:2007}, as well as provided us with detailed pictures of how their properties vary smoothly across the ISCO \\citep{Krolik:2005}.  These simulations have vindicated the prescient remark made by \\citet{Thorne:1974}:  \\noindent ``In the words of my referee, James M. Bardeen (which echo verbal warnings that I have received from Ya. B. Zel'dovich and V.F. Schwartzman), `It seems quite possible that magnetic stresses could cause large deviations from circular orbits in the very inner part of the accretion disk and change the energy-angular-momentum balance of the accreting matter by an amount of order unity'.\"  Indeed, when the black hole rotates, significant magnetic stresses can be found throughout the accretion flow, all the way to the edge of the event horizon \\citep{Krolik:2005}.  Work on specific dynamical pictures of accretion has stimulated a reexamination of the simple picture that discs have sharp edges, within which little of interest happens.  Although it is true that there are qualitative changes across the ISCO, they do not happen discontinuously. As argued by \\cite{Krolik:2002}, what one means by ``edge\" depends on the question asked.  For example, the ``reflection edge\", the edge outside of which most of the observed Fe K$\\alpha$ and Compton reflection photons are created, is likely to lie near, but possibly either inside or outside, the ISCO \\citep{Reynolds:1997,Krolik:2002,Brenneman:2006,Reynolds:2008}---its exact position depends on the density and optical depth of the gas in that region, and on the intensity of ionizing radiation striking it. Similarly, the ``radiation edge\" (the edge inside of which little of the total luminosity emitted by the flow escapes to infinity) should be near the ISCO, but is not necessarily identical to it. Its position depends both on the profile of dissipation and on the ability of photons to escape to infinity, and to do so without excessive loss of energy to gravitational redshift.  In this paper we examine the influence of magnetic stresses at and inside the ISCO on the apparent size, luminosity and characteristic temperature of black hole accretion discs.  This effort is important to black hole phenomenology because so many observational diagnostics depend upon these three parameters.  A prime example is provided by attempts to use spectral fits to constrain black hole mass and spin.  This program rests upon the idea that the observed luminosity, $L$, is essentially thermal, so that it may be characterized by an effective radiating area and a characteristic temperature: $L = A(a/M,\\theta) T^{4}_{char}$.  Because both $T_{char}$ and $L$ can be measured, $A$ can be inferred through this relation.  With a theoretically-supported connection between $A$ and $a/M$ (and some other constraint on $\\theta$), the inference of $A$ leads to an inference of $a/M$.  Current efforts connect $T_{char}$ to $M$, $\\dot M$, and $a/M$ using the \\cite{Novikov:1973} model for the radial dissipation profile. This model depends in an essential way on the assumption that all stresses cease at and within the ISCO, whose location (at Boyer-Lindquist radial coordinate $r/M=r_{ms}$) is a function only of the black hole mass and spin.  This temperature is further adjusted by an appropriate correction for gravitational and Doppler energy shifts and a ``color temperature correction'' due to opacity effects \\cite[see e.g.][]{Done:2008}.  The apparent size of the disc, $A$, is likewise found from the dissipation profile of the Novikov-Thorne model.  As first noted by \\cite{Page:1974}, significant magnetic forces undercut the rationale for the zero-stress boundary condition; the simulation data we discuss here shows quantitatively how these magnetic forces alter the connections between both $A$ and $T_{char}$ and the black hole spin parameter.  Although more work is needed to explore fully these new effects, in this paper we begin the discussion of how they can influence these inferences.  The simulations reported here employ full general relativity and three-dimensional MHD, so that whilst their treatment of angular momentum flow and inflow dynamics is quite accurate, they do not directly track dissipation. In an accretion disc, energy is extracted from orbital motion and transformed into kinetic and magnetic energy on the largest scales of the turbulence.  Subsequently, this energy cascades down to a dissipation scale (either viscous or resistive) where it is finally thermalized.  Current simulations can describe well the first stages of this process, but can mimic only indirectly the last step: grid-level effects intervene at lengthscales far larger than the physical scale of dissipation.  In fact, the simulation code whose data we will use solves only the internal energy equation, and makes no attempt to follow dissipative energy losses except those associated with shocks.  We proceed by instead making a plausible {\\it ansatz} for heating within the disc that can be determined \\textit{a posteriori } from simulation data.  As we shall see, our results depend primarily on the qualitative fact that dissipation continues smoothly across the ISCO, and on the nature of photon trajectories deep in a relativistic potential; for this reason, we believe that a non-rigorous, but physically-motivated, {\\it ansatz} will not be misleading.  We opt for the simplest choice: a connection between the heating rate and the stress that follows from the standard model for energy conservation in an accretion disc \\citep{Page:1974,Balbus:1994,Hubeny:1998}.  We can also investigate the importance of enhanced stress in a semi-analytic fashion that bypasses most of the limitations of the current simulations.  We employ the model formulated by \\cite{Agol:2000} (hereafter the AK model) that allows for non-zero stress at the ISCO but otherwise computes the dissipation profile using the same approach as the standard model.  The AK model cannot be extrapolated into the plunging region, and has in common with the NT model a fixed disc size.  Its unique feature is the enhanced total dissipation due to the nonzero stress at the ISCO; simulation data provides the single parameter needed to calibrate the model. The AK model, therefore, provides an important link between the standard model and the full simulation results and allows us to gauge the appropriateness of the radiation {\\it ansatz} employed for the latter.  \\begin{table*} \\caption{Simulation Parameters} \\label{sims} \\begin{tabular}{@{}lccccccc} \\hline Name          & $a/M$         & $r_{+}/M$ & $r_{\\rm in}/M$ & $r_{\\rm ms}/M$  & Field & $T_{avg}\\times10^{3}M$ & Originally Presented in \\\\ \\hline KD0b & 0.0   & 2.00 & 2.104  & 6.00 & Dipole & 6--8 & \\cite{De-Villiers:2003b}  \\\\ KD0c & 0.0   & 2.00 & 2.104  & 6.00 & Dipole & 8--10 & \\cite{Hawley:2006}  \\\\ QD0d & 0.0   & 2.00 & 2.104  & 6.00 & Quadrupole & 8--10 & This work  \\\\ KDIa & 0.5   & 1.86 & 1.904  & 4.23 & Dipole & 6--8  & \\cite{De-Villiers:2003b}    \\\\ KDIb & 0.5   & 1.86 & 1.904  & 4.23 & Dipole & 8--10  & \\cite{Hawley:2006}   \\\\ KDPd & 0.9   & 1.44 & 1.503  & 2.32 & Dipole & 6--10  & \\cite{De-Villiers:2003b}  \\\\ KDPg & 0.9   & 1.44 & 1.503  & 2.32 & Dipole & 8--10  & \\cite{Hawley:2006}  \\\\ QDPa & 0.9   & 1.44 & 1.503  & 2.32 & Quadrupole & 8--10 & \\cite{Beckwith:2008a}  \\\\ TDPa & 0.9   & 1.44 & 1.503  & 2.32 & Toroidal & 20--22 & \\cite{Beckwith:2008a}  \\\\ KDG  & 0.93  & 1.37 & 1.458  & 2.10 & Dipole & 8--10  & \\cite{Hawley:2006}  \\\\ KDH  & 0.95  & 1.31 & 1.403  & 1.94 & Dipole & 8--10  & \\cite{Hawley:2006}  \\\\ KDJd & 0.99  & 1.14 & 1.203  & 1.45 & Dipole & 8--10  & This work  \\\\ KDEa & 0.998  & 1.084 & 1.175  & 1.235 & Dipole & 6--8  & \\cite{De-Villiers:2003b} \\\\ KDEb & 0.998  & 1.084 & 1.175  & 1.235 & Dipole & 8--10  & This work \\\\ QDEb & 0.998  & 1.084 & 1.175  & 1.235 & Quadrupole & 8--10  & This work \\\\ \\hline \\end{tabular}   \\medskip Here $a/M$ is the spin parameter of the black hole, $r_{+}$ is the horizon radius, $r_{\\rm in}$ is the innermost radius in the computational grid, field is the initial field topology in the torus and $T_{avg}\\times10^{3}$ is the time-interval over which simulation data was averaged. For reference, we note where individual simulations were originally presented. \\end{table*}   To relate dissipation rates to radiation received at infinity, we must perform one additional calculation using three further assumptions: that the dissipated heat is efficiently converted to photons, that these photons emerge from the accretion flow very near where they are created, and that they are radiated isotropically in the fluid frame. The first two assumptions are equivalent to requiring the timescales for dissipation, radiation, and photon diffusion to be shorter than the inflow timescale for all fluid elements. The third, while not strictly justified, is the simplest guess we can make.  Given those assumptions, we use a general relativistic ray-tracing code to relate the luminosity radiated by each fluid element to the luminosity received by observers located at different polar angles far from the black hole.  A further result of this calculation is a new estimate of the radiative efficiency of accretion.  The traditional calculation of this quantity follows directly from a primary assumption that there are no forces inside the ISCO and two additional assumptions, that the radiation is prompt and all of it reaches infinity.  We improve upon this traditional estimate in two ways: we allow for dissipation associated with the stresses we measure at and inside the ISCO; and we calculate the radiated energy (even within the NT model) that actually reaches distant observers. However, we do not regard our result as sufficiently final or complete to give it much weight, as it is likely to be more model- and parameter-dependent than our placement of the radiation edge.  One reason for downplaying this result is that we find it necessary to omit any estimate of dissipation outside the main part of the accretion flow. We have less confidence that our dissipation prescription is appropriate in the jet, or even in the disc corona, than we do when it is applied to the main disc body and plunging region.  Moreover, radiation from these lower-density regions is less likely to be thermalised and contribute to the radiation usually identified with the disc continuum.  The rest of this paper is structured as follows. In \\S\\ref{numerics} we briefly review the numerical scheme employed to solve the equations of GRMHD in the simulations whose data we use, give an overview of the parameters of the simulations included in this work, and describe our general relativistic ray-tracing code.  In \\S\\ref{dissmodels}, we contrast the dissipation rate distributions predicted by the standard model, its Agol \\& Krolik extension, and our simulation-based {\\it ansatz}.  In \\S\\ref{rad} we compute the luminosity at infinity predicted by each of these dissipation profiles and discuss their consequences for the location of the radiation edge and the characteristic temperature of the accretion flow. Finally, in \\S\\ref{summ}, we draw specific conclusions from our results and describe their implications for black hole phenomenology. ",
        "conclusions": "\\label{summ}  That magnetic forces might cause substantial stress at the ISCO was foreseen very shortly after the invention of the standard model \\cite[][]{Page:1974}. This possibility now appears to be an immediate corollary of the well-established result that MHD stresses account for most of the angular momentum transport in the bodies of accretion discs.  Indeed, such stresses were seen in the first generation of three dimensional GRMHD disc simulations.  The goal of this paper has been to begin the linkage of these numerical MHD simulations to the observable properties of accreting black holes even before the simulations are fully equipped to make predictions about how these systems radiate.   To do so, we have followed a path of cautious extrapolation from older methods. We first used simulation data to fix the single parameter of a model (called AK here) that changes the previous standard (the Novikov-Thorne model) only by admitting the possibility of a non-zero stress at the ISCO. Because the AK model is defined in a way that prevents its extrapolation within the ISCO, we used the formalism underlying both it and the NT model (i.e., vertically-integrating and azimuthally- and time-averaging the equation of momentum-energy conservation under the assumption that the four-velocity and the stress tensor are orthogonal) to create an expression for the dissipation  (called $Q_{\\mathrm{MW}}$ here) valid both inside and outside the ISCO.  \\begin{figure} \\begin{center} \\includegraphics[width=0.48\\textwidth]{saavg_tchar} \\end{center} \\caption[]{Plot showing solid angle averaged $\\{ T_{char} \\}$ . Symbols are as in Figure \\ref{dissedge}} \\label{saavgtchar} \\end{figure}  Happily, in the body of the accretion disc well outside the ISCO, both the AK and the MW method agree fairly well with NT more or less independent of black hole spin and magnetic field topology, although the irregularity in the curves of Figure~\\ref{spindiss} reminds one that the simulations are dynamic and time-varying, and that the 26 samples of simulation data we used define only somewhat imperfectly the long-term time-average.  In addition, with the exception of the extreme high-spin example, where the AK and MW methods depart from NT just outside the ISCO, they do so together.  This is noteworthy because, although they are based on closely-related formalisms, they are not identical: perhaps their most significant contrast is that the AK method assumes that $u_r = 0$, while the MW method does not. Lastly, inside the ISCO, where only the MW method is defined, it follows a smooth extrapolation from larger radius.  When the black hole spins slowly ($a/M \\lesssim 0.9$), $Q_{\\mathrm MW}$ extends with hardly any change in logarithmic derivative with respect to radius.  For higher spin, the extension gradually steepens toward smaller radius, but the next step in our formalism shows that this makes little difference to observed radiation: relativistic ray-tracing shows that the volume deep inside the plunging region, particularly at high spin, contributes little energy to the luminosity reaching observers at infinity.  Thus, we are relatively confident that, despite the uncertainties involved, our estimate of the location of the radiation edge is comparatively insensitive to the exact relation between the flow's detailed dynamical properties and the dissipation rate.  \\begin{figure} \\begin{center} \\includegraphics[width=0.48\\textwidth]{tcshape} \\end{center} \\caption[]{Plot showing the region of the $(T_{char},a/M)$-plane consistent with a given value of $\\theta_{o}$ for a $10 M_{\\odot}$ black hole accreting at $0.1L_{edd}$. Different colors correspond to different $\\theta_{o}$ (see key in bottom left corner of plot), dot-dash lines show the $T_{char}$ curve predicted by $Q_{NT}$ at a given $\\theta_{o}$, and the region enclosed by the dot-dash and dashed lines of a given color show the region of the $(T_{char},a/M)$-plane consistent with that value of $\\theta_{o}$.} \\label{tcshape} \\end{figure}   The dependence of the radiation edge on spin may be summarised succinctly: At the highest spin, there is relatively little difference between the different methods of estimating its position because the ISCO is so close to the horizon that the great majority of photons released in the plunging region never reach infinity, or if they do, are severely redshifted.  It moves from 2--3 times the radius of the ISCO when the disc is viewed face-on to almost exactly at the ISCO when the disc is viewed nearly edge-on.   This inward movement of the radiation edge with increasing inclination angle is quite model-independent, as it stems from relativistic photon propagation effects: when photons from the plunging region do reach infinity with substantial energy, it is because they are emitted in the direction of the orbital motion. Probability of escape is then enhanced by a combination of special relativistic beaming and gravitational lensing; energy at infinity is enhanced by special relativistic Doppler boosting.  As the black hole spin decreases, the diminishing depth of the potential immediately inside the ISCO makes it progressively easier for photons to escape from that region and reduces the gravitational redshift they suffer when they do. The result is that the radiation edge moves farther inside the ISCO as {\\it either} the spin diminishes (at fixed viewing angle) or the inclination angle moves toward the equatorial plane (at fixed spin). At its most extreme, the case of $a/M = 0$ and $\\theta_o = 85^{\\circ}$, $r_{re}$ can be $\\simeq 0.5 r_{ms}$.  Figure~\\ref{ofrms} gives additional cause to believe that these results are comparatively insensitive to dissipation model.  Although the radiation edge can move well inside the ISCO at low spin and high inclination, most of the light received by distant observers is generally emitted in the region near and outside the ISCO.  Only at the highest inclinations ($\\theta_o \\gtrsim 75^{\\circ}$) and lowest spins ($a/M \\lesssim 0.5$) does the contribution of the plunging region to the luminosity approach $50\\%$.   Thus, most of the light seen at infinity likely comes from a region where the predictions of the AK and the MW models differ little.  Nonetheless, because it is also true that most of the light is emitted within a radius at most a few times the ISCO (except for the highest spin viewed more or less face-on), the contrast in total luminosity between the AK and MW models on the one hand, and the NT on the other, are order unity for all spins $\\leq 0.9$. For higher spins, the effect may be smaller, but the uncertainties are also greater.  These conclusions have immediate implications for a number of phenomenological issues.  Firstly, as suggested by \\cite{Falcke:2000}, it may be possible to image the nearest supermassive black hole, the one in Sgr A*.  Because its accretion flow, unlike those of intrinsically brighter systems, could well be radiatively inefficient, a simulation scheme that conserves total energy is more appropriate to analysing its emission.  \\cite{Noble:2007}, using such a code (albeit an axisymmetric version), have produced predicted images that illustrate several of the effects emphasised here, although in their work so far they have not reported quantitative descriptions of characteristic emission radii.  Secondly, the enhanced total radiative efficiency due to dissipation in the marginally stable region may affect estimates of population-mean spin parameters \\cite[e.g., as for AGN by][] {Elvis:2002,Yu:2002}.  Because the efficiency rises with increasing prograde spin in the NT model, the spin inferred by this method may overestimate the actual spin of accreting black holes if this enhancement is ignored.  The additional luminosity from enhanced stress in the innermost part of the accretion flow could significantly alter the emergent spectrum. Employing a simple thermal model, we have found that the characteristic temperature of the flow increases by a factor of 1.2 --1.4 over that predicted by the NT model. As a consequence, the thermal peak of the disk spectrum (at $\\sim 1$~keV in Galactic black holes, $\\sim 10$~eV in AGN) may be pushed to somewhat higher energies.  Several caveats must be mentioned, however, in regard to this prediction. First, this number supposes an emergent spectrum that is Planckian, but most estimates of the disc atmosphere's structure suggest that it is scattering-dominated, so that the color temperature of the spectrum is shifted upward from the effective temperature. The magnitude of this shift depends on details of the disc's vertical structure that are not as yet well known \\cite[see e.g.][]{Davis:2005}. Furthermore, \\cite{Blaes:2006} \\citep[using the vertically stratified shearing box simulations of][]{Hirose:2005} show that magnetic pressure support changes the vertical structure of the disk resulting in a noticeable hardening of the emergent disk spectrum compared to the standard Novikov-Thorne picture due to non-LTE effects. Second, it is possible that some of the enhanced dissipation will occur where the density and optical depth are too low to accomplish thermalisation. Strengthening of the ``coronal\", i.e., hard X-ray, emission, rather than hardening the thermal disc spectrum would then be the likely consequence. Third, our treatment ignores those photons emitted deep in the potential that neither escape directly to infinity nor are captured by the black hole, but instead strike the disc. As shown by \\cite{Agol:2000}, this ``returning radiation\" can be a substantial fraction of all photons emitted when $r \\lesssim 5r_g$. Depending on their spectrum and the structure of the disc atmosphere where they strike, these photons may be either reflected (with Doppler shifts) or absorbed and their energy reradiated at a different (in general, lower) temperature. Quantitatively evaluating all three of these effects is well beyond the scope of this paper, but can be done in future work.  There are also implications for attempts to determine black hole spin from spectral fitting.  In all three models, the characteristic radius of emission is always {\\it near} the ISCO, but does not, in general, coincide with it.  Generally speaking, this characteristic radius is largest for the NT model, smaller (but still outside the ISCO) for the AK model, and smaller still, possibly moving into the plunging region inside the ISCO, for the MW model.  Because the ISCO moves to smaller radial coordinate as $a/M$ increases, these characteristic radii always become smaller for faster spin.  However, the fractional amount by which the characteristic emission radius moves inward in the MW model is greatest for the lowest spins, so that in the end, the MW model predicts a relatively slow variation of radiation edge with black hole spin. The AK model, like the NT model, does not radiate from inside the ISCO, but the additional stress at and just outside the ISCO in this model (relative to the NT prediction) produces a systematic increase in the characteristic temperature.  The magnitude of this shift in characteristic temperature rises, of course, with increasing additional stress.  When there is emission from the plunging region, as in the MW model, the characteristic temperature can rise still higher, but the highly relativistic motions there make observed properties more strongly dependent on inclination angle.  In addition, a larger fraction of the emitted photons can be captured by the black hole.  When all these considerations are combined,  we find that, for fixed black hole mass, luminosity, and inclination angle, the uncertainty in the characteristic temperature of the radiation reaching distant observers due to uncertainty in the dissipation profile is {\\it greater} than the that due to a complete lack of knowledge of the black hole's spin. Clearly, our incomplete understanding of accretion disc physics (here specifically the magnitude of the stress at and inside the ISCO) makes it difficult to determine a black hole spin based on continuum model-fitting.  The best one can say is that estimates based on the traditional Novikov-Thorne model can be expected to yield the most rapid spin possible, but the actual spin may be significantly slower.  Our results demonstrate the potential importance of nonzero stresses at and inside the ISCO. But how representative are the specific values obtained in these simulated discs?  There are two considerations: those arising from purely numerical effects, and those limitations arising from the assumptions and parameters of the model used.  First, the results of numerical simulations can be influenced by finite resolution and the limitations of the numerical technique.  All of the simulations presented in this work were performed at a resolution of $192\\times192\\times64$ $(r,\\theta,\\phi)$ grid zones using ideal MHD and an internal energy equation.  The equation of state and the numerical energy dissipation are unlikely to have a direct effect on our conclusions as $Q_{MW}$ is derived directly from the \\emph{physical} Maxwell stresses within the disc, rather than by measuring some \\emph{numerical} dissipation rate.  Low resolution usually causes the Maxwell stress to be {\\it undervalued}; if so, the implications of this paper would be strengthed by improved resolution.  Until available computer power makes better-resolved three-dimensional simulations possible, the best way we have to test the effects of finite resolution is to compute axisymmetric simulations with higher resolution. A variety of such simulations were presented in \\cite{Beckwith:2008a} with resolutions up to $1024^2$.  We observe that greater resolution reduces the rate of numerical reconnection and improves the ability of the simulation to maintain certain field configurations and small-scale field structures.  Overall the amplitude of the turbulent Maxwell stresses in the disc remained largely unchanged as resolution was increased.  We have also calculated ${\\cal W}^{(r)}_{(\\phi)}$ and $Q_{MW}$, and find no significant qualitative differences from the results presented in this work.  \\newpage  Beyond the purely numerical issues, the value of the Maxwell stress at the ISCO may depend on a number of disc properties.  In the ensemble of simulations presented here, the stress levels are determined in part by the initial field topology (dipole versus quadrupolar, poloidal versus toroidal, and the presence or absence of a net vertical field). Indeed, local shearing-box simulations suggest that the saturated field strength can increase substantially when large-scale vertical field threads the disc \\cite[][]{Balbus:1998}.  It is also possible that the saturation stress depends on disc thickness.  To quantify disc thickness, we define the scale-height $H$ as the proper height above the plane at which the time- and azimuthally-averaged density falls by $1/e$ from its similarly averaged value on the equatorial plane (see \\S\\ref{dissmodels}): $H= \\int^{\\theta_h}_{\\pi/2} \\sqrt{g_{\\theta \\theta} (r=3M,\\theta)} d\\theta$. We similarly define $R$ as the proper (as opposed to coordinate) radial distance from the horizon to $3M$ plus the coordinate distance from the origin to the horizon: $R = r_{in} + \\int^{r=3M}_{r_{in}} \\sqrt{g_{rr} (r,\\theta=\\pi/2)} dr$, with $r_{in}$ as given in Table \\ref{sims}).  We then define the disc thickness as the ratio $H/R$.  Even though $r=3M$ lies well within the plunging region for the lower spin cases, we find that there is a slow enough radial variation in this quantity to make $H/R$ a reasonably well-defined parameter.  Measured in this way, our discs are modestly thick, with a characteristic aspect ratio $H/R = 0.06$--0.2 at $r=3M$.  Most of the range in $H/R$ results from the fact that the maximum pressure in the initial condition for these simulations varies somewhat between different $a/M$.  It is difficult to say, however, just what sort of dependence there may be on disc thickness.  There are some arguments suggesting that magnetic effects may increase with increasing $H/R$.  For example, local shearing-box simulations find that the Maxwell stress is proportional to magnetic pressure \\cite[][]{Balbus:1998,Sano:2004}, but there have not yet been any systematic studies of what regulates the magnetic pressure in global, vertically-stratified discs.  \\cite{Afshordi:2003} have argued that inner disc stresses and dissipation may depend on disc thickness as well as on accretion rate, an argument reiterated by \\cite{Shafee:2007}, but their arguments are framed in an essentially hydrodynamic context, and therefore eliminate any possibility of predicting magnetic stresses.  They are also non-relativistic, and therefore eliminate any effects due to frame-dragging.  There are also arguments that any dependence on $H/R$ may be weak.  In the simplest analytic or quasi-analytic MHD models, magnetic torques at the ISCO remain significant even in the limit of a zero pressure disc \\citep{Krolik:1999a,Gammie:1999}.  As discussed in \\cite{Krolik:2005}, processes analogous to the Blandford-Znajek mechanism can readily transport energy and angular momentum from rotating black holes to the accretion flow, and there is no particular reason to think that these processes should be tightly connected to the pressure in the disc. In the end, only direct simulations with the resolution to describe thin discs adequately will answer this question in a satisfactory way. The results of the present investigation provide yet one more reason why it will be important to do so.  The final question that we address in this work is how the results presented here relate to current state of the art measurements of black hole spin via spectral fitting of the disk continuum. The 6 systems with the best data \\cite[see e.g.][]{Davis:2006,Shafee:2006a,Middleton:2006,Liu:2008} all have spins in the range $a/M\\sim0.1-0.8$ based on disk models that assume the stress-free inner boundary condition. From the perspective of the disk stress models these spins are upper limits.  This might indicate that the hole is counter-rotating, but also opens the possibility that spin determinations might themselves constrain the physics near the ISCO. Firstly, the stress levels at the ISCO could be near the value assumed by the stress-free inner boundary condition; secondly, the classical relationship between stress and dissipation might not hold for enhanced magnetic stresses ear the ISCO; thirdly, the density levels at and inside the ISCO could be insufficient to thermalize the dissipated heat; fourthly, the time-scales for thermalization and radiation of the dissipated heat could be longer than the inflow time-scale.  Another uncertainty which we have not examined is that the plane of the disk and the equatorial plane of the hole could be misaligned and so the disk is subject to the Bardeen-Peterson effect \\cite[see e.g.][]{Fragile:2007}. Understanding the roles played by the these effects will be crucial in providing robust estimates of black hole spins via spectral fitting of the disk continuum."
    },
    "0801/0801.0189_arXiv.txt": {
        "abstract": "We study baryon number violating nucleon decays induced by unparticle interactions with the standard model particles. We find that the lowest dimension operators which cause nucleon decays can arise at dimension $6 + (d_s-3/2)$ with the unparticles being a spinor of dimension $d_s=d_\\U +1/2$. For scalar and vector unparticles of dimension $d_\\U$, the lowest order operatoers arise at $6+d_\\U$ and $7+d_\\U$ dimensions, respectively. Comparing the spinor unparticle induced $n \\to O^s_\\U$ and experimental bound on invisible decay of a neutron from KamLAND, we find that the scale for unparticle physics is required to be larger than $10^{10}$ GeV for $d_\\U < 2$ if the couplings are set to be of order one. For scalar and vector unparticles, the dominant baryon number violating decay modes are $n\\to \\bar \\nu + O_\\U (O^\\mu_\\U)$ and $p \\to e^+ + O_\\U (O^\\mu_\\U)$. The same experimental bound puts the scales for scalar and vector unparticle to be larger than $10^{7}$ and $10^{5}$ GeV for $d_\\U <2$ with couplings set to be of order one. Data on $p \\to e^+ \\mbox{invisible}$ puts similar constraints on unparticle interactions. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.0006_arXiv.txt": {
        "abstract": "{ Small universe models predicted a cutoff in large-scale power in the cosmic microwave background (CMB). This was detected by the Wilkinson Microwave Anisotropy Probe (WMAP). Several studies have since proposed that the preferred model of the comoving spatial 3-hypersurface of the Universe may be a Poincar\\'e dodecahedral space (PDS) rather than a simply connected, flat space. Both models assume an FLRW metric and are close to flat with about 30\\% matter density. } { We study two predictions of the PDS model. (i) For the correct astronomical positioning of the fundamental domain, the spatial two-point cross-correlation function $\\ximc$ of temperature fluctuations in the covering space (where the two points in any pair are on different copies of the surface of last scattering (SLS)) should have a similar order of magnitude to the auto-correlation function $\\xisc$ on a single copy of the SLS. (ii) Consider a ``generalised'' PDS model for an {{\\em arbitrary}} ``twist'' phase $\\phi \\in \\left[0,2\\pi\\right]$. The optimal orientation and identified circle radius for a generalised PDS model found by maximising $\\ximc$ relative to $\\xisc$ in the WMAP maps should yield one of the two twist angles $ \\pm 36\\ddeg$. } { Comparison of $\\ximc$ to $\\xisc$ extends the identified circles method, using a much larger number of data points. We optimise the ratio of these functions at scales $\\ltapprox 4.0${\\hGpc} using a Markov chain Monte Carlo (MCMC) method over orientation ($l$, $b$, $\\theta$), circle size $\\alpha$, and twist $\\phi$. } { Both predictions were satisfied: (i) An optimal generalised PDS solution was found for two different foreground-reduced versions of the WMAP 3-year all-sky map, both with and without the kp2 galactic contamination mask. This solution yields a strong cross-correlation between points which would be distant and only weakly correlated according to the simply connected hypothesis. The face centres are $\\{(l,b)\\}_{i=1,6}\\approx \\{ (184\\ddeg, 62\\ddeg), (305\\ddeg, 44\\ddeg), (46\\ddeg, 49\\ddeg), (117\\ddeg, 20\\ddeg), (176\\ddeg, -4\\ddeg), (240\\ddeg, 13\\ddeg) \\}$ (and their antipodes) to within $\\approx 2\\ddeg$ ; (ii) This solution has twist $\\phi= (+39 \\pm 2.5)\\ddeg$, in agreement with the PDS model. The chance of this occurring in the simply connected model, assuming a uniform distribution $\\phi \\in [0,2\\pi]$, is about {\\probifnotPDS}. } { The PDS model now satisfies several different observational constraints. } ",
        "introduction": "\\label{s-intro}  The past decade and a half have shown considerable growth in attempts to determine the global shape, i.e. not only the curvature, but also the topology, of the spatial comoving section of the Universe, i.e. of the 3-manifold to which a 3-hypersurface corresponds, informally known as ``space''. As noted by several authors, in particular \\nocite{Star93}{Starobinsky} (1993) and \\nocite{Stevens93}{Stevens} {et~al.} (1993), a space that is ``small'' compared to the surface of last scattering (SLS) cannot contain eigenmodes, which are used for expressing density perturbations, which are themselves larger than the space itself. This should lead to a cutoff of power in statistics representing these fluctuations, above which power should drop to zero. This prediction was made after COBE data were available, but {\\em before} the WMAP satellite was launched. For practical, observational reasons, spherical harmonic analyses of temperature fluctuations on the 2-sphere are frequently made. However, a physically more natural statistic to use is one in three-dimensional space, e.g. the two-point auto-correlation function.  \\fauto  The predicted cutoff in large-scale power appears to have been confirmed by the first-year observations of the Wilkinson Microwave Anisotropy Probe (WMAP) experiment. With this data, \\nocite{WMAPSpergel}{Spergel} {et~al.} (2003) published a figure approximately equivalent to such a function, i.e. the black ``WMAP data'' curve of Fig.~16 in their paper.  Their figure shows the auto-correlation calculated as a function of angular separation, shown against projected spatial separation for a (first year) template-cleaned V map with the kp0 galactic contamination mask.  The authors note the surprisingly flat correlation on large scales and suggest a multiply connected universe model to match this function.  In Fig.~\\ref{f-auto}, we calculated the auto-correlation function directly as a function of three-dimensional spatial separation, not of angular separation, using the 3-year integrated linear combination (ILC) map with the kp2 cut. \\posteditorchanges{This figure can be approximately compared to the black ``WMAP data'' curve of Fig.~16 of \\protect\\nocite{WMAPSpergel}{Spergel} {et~al.} (2003), except that the present figure shows the auto-correlation calculated as a function of three-dimensional spatial separation, not of angular separation; \\langed{``however'' would be wrong here; 3-d space vs angle already does imply non-linearity, but this needs to be emphasised.} the relation between the two is {\\em not} linear. Also, our figure uses the 3-year ILC map with the kp2 cut, not the 1-year template-cleaned V map with the kp0 cut.}  In spatial comoving units, we confirm that the auto-correlation is very close to flat for separations larger than $\\approx 10 {\\hGpc}$. The relation between angular and spatial scales is, of course, not linear. Equation~(\\ref{e-alpha-tri}) below (for either a spherical covering space or for a flat covering space using the limit $R_C \\gg \\rSLS \\ge d/2$) can be used to calculate this.  If the size of the Universe\\footnote{See e.g. Fig~10 of \\protect\\nocite{LR99}{Luminet} \\& {Roukema} (1999) for a schematic diagram of various definitions of the ``size'' of the fundamental domain, including the injectivity radius, the in-radius, and the out-radius.} is about 10{\\hGpc}, as this figure seems to indicate, then which of the various 3-manifolds correctly describes comoving space? Motivated by indications that the Universe may have positive curvature, and using eigenmode-based simulations to study the spherical harmonic ($C_l$) spectrum of the WMAP data, \\nocite{LumNat03}{Luminet} {et~al.} \\nocite{LumNat03,Caillerie07}({Luminet} {et~al.} 2003; {Caillerie} {et~al.} 2007) argue that the Poincar\\'e dodecahedral space (PDS) is favoured by the WMAP data over an infinite, simply connected flat space. \\nocite{Caillerie07}{Caillerie} {et~al.} (2007) state that, by requiring maximal repression of the quadrupole signal, an optimal total density of $\\Omtot = 1.018$ is favoured (for a non-relativistic matter density $\\Omm \\equiv 0.27$ and Hubble constant $H_0 = 70$\\kms/Mpc) for the PDS model. Several other authors \\nocite{Aurich2005a,Aurich2005b,Gundermann2005}({Aurich} {et~al.} 2005a, 2005b; {Gundermann} 2005) have also compared simulations for PDS models to the observed first-year and three-year Wilkinson Microwave Anisotropy Probe (WMAP) maps of the cosmic microwave background (CMB).  In all these studies, both the infinite flat models and the PDS models are used in the context of a standard, hot big bang model, i.e. where the Universe has a Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) metric, perturbed by fluctuations that collapse gravitationally to form structures such as filaments and clusters of galaxies, and where values of the metric parameters consistent with the consensus obtained during the last decade of observations are adopted: the Universe is close to flat on length scales up to the SLS and has about 30\\% non-relativistic matter density.  In other work, \\nocite{RLCMB04}{Roukema} {et~al.} (2004) used the identified circles principle \\nocite{Corn96,Corn98b}({Cornish} {et~al.} 1996, 1998) to find a specific optimal orientation of the PDS model based on the WMAP first-year ILC map, and published a tentative set of coordinates. \\nocite{KeyCSS06}{Key} {et~al.} (2007) confirmed the presence of a signal at the celestial coordinates, circle radius, and $-36\\ddeg$ twist published in \\nocite{RLCMB04}{Roukema} {et~al.} (2004), but argue that it should be considered a false positive. Using software independent of that used in \\langed{We want to avoid ambiguity between independence of the article Roukema et al. 2004 (true) and independence of the people Roukema et al. (false).} \\nocite{RLCMB04}{Roukema} {et~al.} (2004), updating to the 3-year WMAP data, and using Gaussian simulations, \\nocite{LewRouk2008}{Lew} \\& {Roukema} (2008) find similar conclusions. They find that a local maximum in the statistic \\langed{``'Statistics' is usually plural, unless this is a very specific mathematical or physical sort. Do you mean this as a single piece of data, which is what the singular form means?'' {\\em ``statistic'' is used here as: ``a function of a sample where the function itself is independent of the sample's distribution.'' \\url{http://en.wikipedia.org/wiki/Statistic}}} used for finding matched circles exists for a circle radius $\\sim 11\\ddeg$ and a $-36\\ddeg$ twist, but it is not statistically significant.   \\nocite{Aurich2005circ}{Aurich} {et~al.} (2006) also made a circles analysis of the WMAP first-year data, using their own estimator and weight function, and, in contrast with \\nocite{RLCMB04}{Roukema} {et~al.} (2004), \\nocite{KeyCSS06}{Key} {et~al.} (2007), and \\nocite{LewRouk2008}{Lew} \\& {Roukema} (2008), did not find any signal at $11\\ddeg$. On the other hand, they did find a tentative PDS signal at $\\Omtot \\approx 1.015$, or equivalently, a circle radius $\\alpha \\approx 40\\ddeg$. The signal is weaker than they expected, but the authors note that uncertainties due to foregrounds and noise structures in the data make it premature to draw firm conclusions.  A disadvantage of the identified circles approach is that it is based on the information in a relatively small number of points on the sky map of temperature fluctuations, making it sensitive to small errors in the data or analysis and requiring prohibitively long computations. Is it \\langed{No, the word ``even'' would be wrong here.} possible to generalise from the identified circles principle?  Moreover, leaving aside the debate about matched circles statistics, observably multiply connected models could reasonably be said to have satisfied only one prediction so far, that of a cutoff in the density fluctuation spectrum on a large scale. Can predictions of the PDS model itself be tested, for example, using the identified circles principle or an extension of it?  In searches for matched circles, the statistic used is usually some variation on what can be considered to be the value of the two-point cross-correlation function of observed temperature fluctuations at pairs of points in the covering space, $\\ximc(r)$, where the two points lie on different copies of the SLS, and the separation of the two points in the pair (on the different copies of the SLS) is zero in the covering space.   \\postrefereechanges{  We can write this function as if it were possible to sample temperature fluctuations as point objects at arbitrary spatial points throughout the covering space, i.e. including arbitrarily low redshifts, as well as points beyond the SLS: \\begin{eqnarray} \\ximc(r) &\\equiv& \\left< \\delta T (x_{i_1}) \\delta T \\left[g_j({x_{i_2}})\\right] \\right>_{i_1,i_2,j}  % \\label{e-ximc-defn} \\end{eqnarray} averaging over triples $(i_1,i_2,j)$ satisfying $d\\left( x_{i_1}, \\left[g_j({x_{i_2}})\\right] \\right) = r$ and $g_j\\not= I$ (the identity $I$ is removed because we only want the cross-correlation). Here, $x_{i_k}$ for $k=1,2$ are arbitrary points on the SLS considered to be located at their comoving spatial positions, $g_j$ for  $j=1, \\ldots, 12$ are the 12 holonomy transformations in the binary icosahedral group $\\Gamma = I^*$ operating on the covering space $S^3$ that match opposite faces of the fundamental polyhedron of the PDS, $d(x,y)$ is the comoving spatial geodesic distance between two points $x,y$ in the comoving covering space (i.e. an arc-length on the covering space $S^3$ embedded in $\\mathbb{R}^4$, e.g. \\nocite{Rouk01-4D}{Roukema} 2001). The temperature fluctuations on one copy of the SLS, $\\delta T(x)$, are extended to the whole covering space via the holonomy transformations, i.e., \\begin{equation} \\forall g_j \\in \\Gamma, \\;\\;  \\delta T(x) = \\delta T \\left[g_j(x)\\right] . \\label{e-temp-covspace} \\end{equation}  The latter assumption, i.e. Eq.~(\\ref{e-temp-covspace}), enables rewriting Eq.~(\\ref{e-ximc-defn}) in a way that becomes observationally realistic, \\begin{eqnarray} \\ximc(r) &\\equiv& \\left< \\delta T (x_{i_1}) \\delta T ({x_{i_2}}) \\right>_{i_1,i_2,j},  % \\label{e-ximc-two} \\end{eqnarray} again subject to the conditions that the average is taken over pairs satisfying $d\\left( x_{i_1}, \\left[g_j({x_{i_2}})\\right] \\right) = r$ and $g_j\\not= I$. The auto-correlation function can then be written \\begin{eqnarray} \\xisc(r) \\equiv \\left< \\delta T (x_{i_1}) \\delta T ({x_{i_2}}) \\right>_{i_1,i_2} , \\label{e-xisc-defn} \\end{eqnarray} averaging over pairs satisfying $ d\\left( x_{i_1}, x_{i_2} \\right) = r$. In words, $\\xisc$ is the 3-spatial auto-correlation function for pairs of points on a single copy of the SLS, while $\\ximc$ is the 3-spatial cross-correlation function for pairs of points where the two members lie on different copies of the SLS in the covering space.     Now rewrite \\langed{$S$ was ``written'' in Roukema et al. 2004, so now it is {\\bf re}written.} the statistic $S$ used in matched circles searches as follows [see e.g. Eq.~(9) of \\nocite{RLCMB04}{Roukema} {et~al.} (2004), with the normalisations (denominator and monopole $T$) ignored for simplicity], \\begin{eqnarray} S &\\equiv& { \\left< {\\delta T }_i \\; {\\delta T }_j  \\right>_{(i,j)} } \\label{e-s-rlcmb04}  \\nonumber \\\\ & = & \\ximc (0) \\label{e-ximc-zero} \\end{eqnarray} for the case where pairs of points $(i,j)$ are (hypothetically) multiply-imaged locations located on a pair of matching circles in the covering space. We use $r=0$ here because that is the defining characteristic of matched circles --- a pair of matching points is a match because the two points are the same space-time points, i.e. they are separated by $r=0$ when their topological images located on adjacent copies of the SLS in the covering space are considered.  }  \\newcommand\\rmax{2} Here, generalise from the zero separation cross-correlation $\\ximc(0)$ to cross-correlations at larger separations $r>0$. We can expect that \\begin{equation} \\ximc(r > 0) < \\ximc(0), \\end{equation} i.e. that correlations weaken with separation and that, in general, $\\ximc(r)$ should be approximately equal to the auto-correlation function $\\xisc(r)$ on scales much smaller than \\langed{$\\ll$ and $<$ differ.} the ``size'' of the fundamental domain, e.g. the in-diameter $2r_{-}$ \\nocite{LR99}(e.g. fig~10,  {Luminet} \\& {Roukema} 1999) \\begin{eqnarray} \\ximc(r > 0) \\sim  \\xisc(r > 0) & &, \\mbox{ if } r \\ll 2r_{-}, \\end{eqnarray} since there is no statistical, physical distinction between a pair of points on different copies of the SLS and a pair of points on a single copy of the SLS, apart from effects that are not locally isotropic, such as the Doppler effect.  Since $\\xisc$ is generally large at small $r$ and small at large $r$, obtaining a large \\langed{Using ``strong'' would be confusing, since we use ``large'' a few words earlier in the same sentence. Moreover, ``strong'' has connotations either of physical force or human force, which is not a useful connotation here. ``Close'' would be utterly confusing to the reader.} correlation requires using a range of length scales $r$ that are relatively small, e.g., \\begin{eqnarray} \\ximc(r \\ltapprox \\rmax \\hGpc) &\\sim& \\xisc(r \\ltapprox \\rmax \\hGpc) \\label{e-ximc-pdsgood} \\end{eqnarray} should be a relatively high positive value if the multiply connected model being studied is the correct model.  Another way of saying this is that this test compares the spatial two-point cross-correlation function of mapped (in the sense of the holonomy transformation $g$, which maps one copy of a point in the covering space to one of its images) and unmapped temperature fluctuations in the covering space, where the two points in any pair are on different copies of the SLS, against the auto-correlation function on a single copy of the SLS.  To see yet another way of thinking about this, suppose that the PDS model at a given orientation and implied circle size is correct and that temperature fluctuations occur as point objects, each emitting isotropically. The covering space {\\em cross}-correlation function, based on pairs of points that are {\\em close to one another due to the application of one or more holonomy transformations $g$, but are not close to one another on a single copy of the SLS}, should then be statistically equivalent (apart from the Doppler effect and foreground/projection effects) to a sampling from the spatial {\\rm auto}-correlation function $\\xisc$, which can be estimated using points that lie on a single copy of the SLS.  {\\em This gives us our first prediction to test the PDS hypothesis: Eq.~(\\ref{e-ximc-pdsgood}) should only be expected to hold if the physically correct PDS (or other) 3-manifold is assumed and is modelled at its correct astronomical orientation.} If the PDS model is correct but a wrong orientation is chosen, then the cross-correlation $\\ximc$ inferred from it on small length scales --- in the covering space since it is the cross-correlation --- should represent correlations of pairs of points that in reality are {\\em not} close together, due to the erroneous orientation that causes mappings from one copy of the SLS to another to be incorrect. Since the correlation of distant points is, in general, small, we have for this incorrect PDS model: \\begin{eqnarray} \\ximc(r \\ltapprox \\rmax \\hGpc) &\\sim& \\xisc(r \\gg \\rmax \\hGpc) \\nonumber \\\\ & \\ll &  \\xisc(r \\ltapprox \\rmax\\hGpc) \\label{e-ximc-pdsbad-first} \\end{eqnarray} i.e. \\begin{eqnarray} \\ximc(r \\ltapprox \\rmax \\hGpc) & \\ll &  \\xisc(r \\ltapprox \\rmax\\hGpc), \\label{e-ximc-pdsbad} \\end{eqnarray} in contrast to Eq.~(\\ref{e-ximc-pdsgood}). Similarly, if the PDS model is incorrect, then arbitrary orientations should also yield small cross-correlations on small length scales, as in Eq.~(\\ref{e-ximc-pdsbad}).  In this paper, the initial aim is to maximise $\\ximc(r \\ltapprox \\rmax \\hGpc)$ relative to  $\\xisc(r \\ltapprox \\rmax \\hGpc)$, varying PDS models over the parameter space of different orientations and circle sizes. (In the actual calculations, we used a range of scales below and a little above $\\sim \\rmax \\hGpc$: see \\SSS\\ref{s-scales} and Eq.~(\\ref{e-range-Gpc}) for details.) This should lead to an estimate of the best orientation and circle size for the PDS model, given the observational data.  However, to further test the PDS model, consider a ``generalisation'' from the mathematically correct PDS model to a class of pseudo-models for which the cross-correlation function can be calculated, but for which most\\footnote{The mathematical term ``almost all'' could be used here, but when taking observational and numerical calculation uncertainties into account, the word ``most'' is more realistic.} of the members of the class are physically invalid. Maximising $\\ximc(r \\ltapprox \\rmax \\hGpc)$ (relative to $\\xisc(r \\ltapprox \\rmax \\hGpc)$) over the extended parameter space defined by this class should then have little chance of yielding an optimal model that is valid, unless the PDS model is astronomically correct.  The ``generalisation'' that we define here is to allow the ``twist'' angle $\\phi$ to be arbitrary in $\\left[0,2\\pi\\right]$. The twist angle can be described as follows. In a single-action, spherical 3-manifold \\nocite{GausSph01}(e.g., Sect. 4.1,  {Gausmann} {et~al.} 2001) thought of as embedded in 4-dimensional Euclidean space, $\\mathbb{R}^4$, any holonomy transformation is a Clifford translation that rotates in $\\mathbb{R}^4$ about the centre of the hypersphere in one 2-plane by an angle of $\\pi/5$, and also by $\\pi/5$ in an orthogonal 2-plane.  From a close-up perspective of the SLS rather than looking at the whole hypersphere $S^3$, or in other words, from a projection into $\\mathbb{R}^3$, one 4-rotation can be thought of as a {\\em translation} in $\\mathbb{R}^3$ from one circle in a matched circle pair to the other member of the pair on the opposite side of the SLS, and the second 4-rotation is then thought of as a {\\em twist} around a vector in the translation direction joining the centres of the two identified circles. This full motion is frequently termed a ``screw motion''.  The ``twist'' is the second of these two rotation angles.  Physically, the only two possible twist angles are $\\pm \\pi/5 .$ If the cross-correlation were to be calculated by filling one copy of the fundamental domain with a uniform distribution of points and then mapping this set of points to copies of the fundamental domain in the covering space, then for an invalid twist angle, some regions of space near the first copy would have either zero, two, three, or more times the density of points in the first copy of the fundamental domain. In other words, sharp discontinuities would occur in calculating the cross-correlation.  However, a practical way of estimating the cross-correlation is the method represented in Eq.~(\\ref{e-s-rlcmb04}), i.e. multiplying temperature fluctuations and then averaging them. Using this method rather than multiplying densities of points, the effect mentioned in the previous paragraph should only lead to slightly more or less frequent sampling in some regions of the cross product of the covering space with itself, not to a modification of the correlations themselves.  The continuous nature of this method, as opposed to the discrete nature of starting with a uniform distribution of points in the fundamental domain, can be seen as follows. For a given pair of observed locations $(x_i, x_j)$ on one copy of the SLS, their comoving spatial separation $d[x_i,g(x_j,\\phi)]$ for a given ``generalised'' \\postrefereechanges{ isometry } $g$, which depends on the generalised twist $\\phi$, changes continuously through the values $\\phi = \\pm \\pi/5,$ not discretely. \\langed{Either ``and'' or ``but'' would be wrong. ``i.e.'' would be correct but would sound pedantic.} Moreover, the product of the observed temperature fluctuations remains constant. Intuitively, we could say that as $\\phi$ varies, the geodesics formed by pairs $[x_i,g_1(x_j,\\phi)]$ do not ``know'' when they meet other pairs  $[x_i,g_2(x_j,\\phi)]$, where $g_1$ and $g_2$ are two holonomy transformations mapping the fundamental domain to two copies of the fundamental domain adjacent to one another in the covering space. \\langed{This should be clearer.}   If the PDS model is correct, then calculating $\\ximc(r)$ using temperature fluctuations and applying isometries using an arbitrary twist angle should give $\\ximc (r \\ltapprox \\rmax \\hGpc)\\ll \\xisc (r \\ltapprox \\rmax \\hGpc)$ as in Eq.~(\\ref{e-ximc-pdsbad}), but give $\\ximc (r \\ltapprox \\rmax \\hGpc)\\sim \\xisc (r \\ltapprox \\rmax \\hGpc)$ as in Eq.~(\\ref{e-ximc-pdsgood}) when the twist angle is correct.   {\\em This gives the second prediction of the PDS model. The maximal cross-correlation for this generalised PDS model estimated using correlations of temperature fluctuations should not only exist as a robust maximum, but it should also give a twist angle of either $\\pm \\pi/5$.}  If the simply connected, perfectly flat model is correct, then the PDS model is incorrect, and it should either be difficult to find a robust maximal correlation $\\ximc (r \\ltapprox \\rmax \\hGpc)$ in this extended 5-parameter space, or else an arbitrary twist angle $\\phi$ should result, with only a small chance of $\\phi$ being close to either of the two values expected for the PDS.  An estimate of how close an ``arbitrary'' twist angle should lie to one of the two PDS values can be made as follows. Given the assumption that \\begin{eqnarray} \\xisc(r \\gg \\rmax \\hGpc) & \\ll &  \\xisc(r \\ltapprox \\rmax\\hGpc), \\label{e-xisc-zeroatbigr} \\end{eqnarray} we can expect that this arbitrary angle should be selected from a uniform probability density distribution on $\\left[0,2\\pi\\right]$. In principle, depending on the assumptions made about the complex statistical properties of the WMAP temperature fluctuation maps, it could be possible for this distribution to be non-uniform, even for a non-PDS space model. However, the estimated spatial auto-correlation function $\\xisc$ for $r \\gtapprox 10${\\hGpc} is close to zero, as shown in fig~16 of \\nocite{WMAPSpergel}{Spergel} {et~al.} (2003) as a function of angular separation on the SLS ($S^2$), and explicitly as a function of spatial separation in Fig.~\\ref{f-auto} here, so the assumption does appear to be supported by the \\langed{The only serious direct empirical evidence is that of WMAP, hence ``the'' rather than implicit ``some''.} empirical evidence.  For the uniform distribution assumption, let us define \\begin{eqnarray} \\Delta\\phi &\\equiv & \\min \\left( \\left|\\phi - \\frac{\\pi}{5}\\right|, \\left|\\phi - \\frac{9\\pi}{5}\\right| \\right) \\label{e-defnDphi} \\end{eqnarray} for $\\phi \\in \\left[0,2\\pi\\right]$. The chance of the observational optimal phase $\\phiWMAP$ being close to $\\pi/5$ or $9\\pi/5$ is then \\begin{eqnarray} P\\left( \\Delta\\phi < \\Delta\\phiWMAP  \\right) &=& \\left\\{ \\begin{array}{l l} 2 \\frac{\\Delta\\phi}{\\pi} , & \\mbox{ if }  \\Delta\\phi \\le \\frac{\\pi}{5}  \\\\ \\frac{1}{5} +  \\frac{\\Delta\\phi}{\\pi} , & \\mbox{ if } \\frac{\\pi}{5} \\le \\Delta\\phi \\le \\frac{4\\pi}{5} .  \\\\ \\end{array} \\right. \\label{e-probphi} \\end{eqnarray} The piecewise nature of this function is because; e.g. there are four ways in which $\\phi$ can differ from $\\pm \\pi/5$ by a small angle such as $10\\ddeg$, but only two ways in which it can differ from $\\pm \\pi/5$ by a large angle such as $100\\ddeg$.   To search for an optimal solution in the parameter space, a Metropolis-Hastings version of a Markov-chain Monte Carlo (MCMC) method is used \\nocite{Neal93}(e.g., Sect. 4.2,  {Neal} 1993,  and references therein) over the five-dimensional parameter space \\begin{equation} \\{ (l,b,\\theta,\\alpha,\\phi) \\}, \\label{e-parameterspace} \\end{equation} where the parameters represent dodecahedron orientation ($l$, $b$, $\\theta$), circle size $\\alpha$ and twist phase $\\phi$. All parameters are initialised as arbitrary angles, except that the circle size is constrained to $5\\ddeg \\le \\alpha \\le 60\\ddeg$ both initially and during the whole chain. See \\SSS\\ref{s-mcmc-method} for more details.  We also ran some MCMC chains starting with the parameters of the \\nocite{RLCMB04}{Roukema} {et~al.} (2004) ``hint'' of a PDS solution with circles of size $\\sim 11 \\ddeg$. If this solution is correct, then we would expect the chains to remain localised around that same solution. If the solution is wrong and if a correct solution exists, then the chains should move towards that correct solution, even though starting at a wrong point. If the solution is wrong and if no correct solution exists, then we would expect the chains to move randomly and fail to find a strong maximum.    The correlation function definitions are described in \\SSS\\ref{s-corr-method}, the probability estimator used to compare multiple-SLS cross-correlation and single-SLS auto-correlation functions is presented in \\SSS\\ref{s-prob-method}, and the MCMC method is described in \\SSS\\ref{s-mcmc-method}. Results are presented in \\SSS\\ref{s-results}. Discussion follows in \\SSS\\ref{s-disc} and conclusions are made in \\SSS\\ref{s-conclu}.   For references on cosmic topology in general, please see the first known article on the subject \\nocite{Schw00,Schw98}({Schwarzschild} 1900, 1998), a short beginner's review of cosmic topology \\nocite{Rouk00BASI}({Roukema} 2000), more in-depth reviews \\nocite{LaLu95,Lum98,Stark98,LR99}({Lachi\\`eze-Rey} \\& {Luminet} 1995; {Luminet} 1998; {Starkman} 1998; {Luminet} \\& {Roukema} 1999), workshop proceedings \\nocite{Stark98,BR99}({Starkman} 1998; {Blanl{\\oe}il} \\& {Roukema} 2000), and lists of two-dimensional and three-dimensional methods (Table~2, \\nocite{LR99}{Luminet} \\& {Roukema} 1999; \\nocite{ULL99b,Rouk02topclass,RG04}{Uzan} {et~al.} 1999; {Roukema} 2002; {Rebou\\c{c}as} \\& {Gomero} 2004). For background on spherical, multiply connected spaces, see \\nocite{Weeks2001}{Weeks} (2001), \\nocite{GausSph01}{Gausmann} {et~al.} (2001), \\nocite{LehSph02}{Lehoucq} {et~al.} (2002), \\nocite{RiazSph03}{Riazuelo} {et~al.} (2004), and \\nocite{LumNat03}{Luminet} {et~al.} (2003). For the identified circles principle, of which the present method can be thought of as an extension, see \\nocite{Corn96,Corn98b}{Cornish} {et~al.} (1996, 1998).  One correction to a point made in much of the above literature concerns the independence of the metric and global topology. For example, \\nocite{LR99}{Luminet} \\& {Roukema} (1999) wrote, ``However, general relativity deals only with local geometrical properties of the universe, such as its curvature, not with its global characteristics, namely its topology''. The error here is that global topology {\\em can}, at least in certain cases, generate a local effect, and thus affect local geometrical properties. See the $T^3$, Newtonian weak field limit, heuristic calculation in \\nocite{RBBSJ06}{Roukema} {et~al.} (2007): if the Universe contains inhomogeneities and is not expanding perfectly isotropically, then at least in the $T^3$ case, \\postrefereechanges{ additional accelerations and decelerations exist in the different fundamental directions, in such a way that they tend to equalise the slightly different expansion rates.}  Comoving coordinates are used when discussing distances (i.e. ``proper distances'', \\nocite{Wein72}{Weinberg} 1972, equivalent to ``conformal time'' if $c=1$). We write the Hubble constant as $H_0 \\equiv 100 h$\\kms/Mpc. ",
        "conclusions": "\\label{s-conclu}  It seems hard to avoid the conclusion that cross-correlations of temperature fluctuations on would-be adjacent copies of the SLS, which are distant from one another and are (on average) very weakly correlated according to the WMAP 3-year observations, (i) imply a highly cross-correlated ``generalised'' Poincar\\'e dodecahedral space symmetry (Figs~\\ref{f-ilc_lbth_N}--\\ref{f-toh_lbth_N}, Table~\\ref{t-dodec})  at which these on-average uncorrelated fluctuations happen to be well-correlated with one another (Fig.~\\ref{f-cross}), and, moreover, (ii) this favoured solution is dominated by a signal whose twist phase $\\phi$ lies within a few degrees of one of the two twist phases necessary for a valid PDS model (Table~\\ref{t-alpha-phi}).  These two successful predictions of the PDS model follow the WMAP confirmation of the generic prediction of small universe models, a power cutoff in structure statistics on large scales, and the solution appears to be consistent with quite different PDS analyses of the WMAP data.  Do we really live in a Poincar\\'e dodecahedral space? Further constraints either for or against the model are certainly still needed, but the evidence in favour of a PDS-like signal in the WMAP data does seem to be accumulating."
    },
    "0801/0801.2003_arXiv.txt": {
        "abstract": "This is the first paper of a series focused on investigating the star formation and evolutionary history of the two early-type galaxies NGC~1407 and NGC~1400. They are the two brightest galaxies of the NGC~1407 (or Eridanus-A) group, one of the 60 groups studied as part of the Group Evolution Multi-wavelength Study (GEMS).  Here we present new high signal-to-noise long-slit spectroscopic data obtained at the ESO 3.6m telescope and high-resolution multi-band imaging data from the HST/ACS and wide-field imaging from Subaru Suprime-Cam. We spatially resolved integrated spectra out to $\\sim$~0.6 (NGC~1407) and $\\sim$~1.3 (NGC~1400) effective radii. The radial profiles of the kinematic parameters $v_{rot}$, $\\sigma$, $h_{3}$ and $h_{4}$ are measured. The surface brightness profiles are fitted to different galaxy light models and the colour distributions analysed. The multi-band images are modelled to derive isophotal shape parameters and residual galaxy images. The parameters from the surface brightness profile fitting are used to estimate the mass of the possible central supermassive black hole in NGC~1407. The galaxies are found to be rotationally supported and to have a flat core in the surface brightness profiles. Elliptical isophotes are observed at all radii and no fine structures are detected in the residual galaxy images. From our results we can also discard a possible interaction between NGC~1400, NGC~1407 and the group intergalactic medium. We estimate a mass of $\\sim$1.03~$\\times$~10$^{9}$~$M_{\\sun}$ for the supermassive black hole in NGC~1407 galaxy. ",
        "introduction": "Early-type galaxies are estimated to contribute to at least half, and perhaps as much as three quarters, of the stellar mass in the local Universe \\citep{bell03}. Understanding their formation and evolution is therefore a fundamental objective of astrophysical research.  \\begin{table*} \\begin{center} \\begin{tabular}{ccccccccccc} \\hline \\hline Name & R.A. & DEC. & Type &PA & v & $\\sigma_{0}$ &$r_{e}$ & M$_{B}$ & L$_{X}$\\\\ &(J2000) & (J2000) & &(degrees) &(km s$^{-1}$)& (km s$^{-1}$)& (arcsec) &(mag) & (erg~s$^{-1}$) \\\\ \\hline NGC~1407 & 03:40:11.9 & $-$18:34:49 & E0 &35 &1779 $\\pm$ 9 & 305.0 $\\pm$ 10.0 &72$\\arcsec$ & $-$21.22 & 1.00~$\\times$~10$^{41}$ \\\\ NGC~1400 & 03:39:30.8 & $-$18:41:17 & E$-$S0 & 40 & 558 $\\pm$ 14 & 279.0 $\\pm$ 2.0 &27$\\arcsec$ & $-$19.95 & 1.32~$\\times$~10$^{40}$ \\\\ \\hline \\hline \\end{tabular} \\end{center} \\caption{Galaxies properties. Type, P.A.: morphological type and position angle of the major axis from HyperLEDA; v, $\\sigma_{0}$, $r_{e}$: velocity, central velocity dispersion and effective radius from HyperLEDA; M$_{B}$: absolute B-band magnitude obtained from HyperLEDA; L$_{X}$: X-ray luminosity from \\citep{osullivan01}.} \\label{ga_prop} \\end{table*}  The debate on whether early-type galaxies are formed by monolithic collapse at high redshifts or via hierarchical merging is still not settled. Photometrically, the tightness of the colour-magnitude diagram suggests that the bulk of the stars were formed at high redshifts, z~$>$~1 (e.g \\citealt{bower92}; \\citealt{kodama99}). However, signs of recent mergers and interactions are often present when images of early-type galaxies are examined (e.g. \\citealt{ft92}). Shells and ripples are believed to be formed by mergers (e.g. \\citealt{quinn84}) or by interactions (e.g. \\citealt{thomson90}; \\citealt{thomson91}). Deviations from perfect elliptical isophotes are correlated with the galaxy shape and amount of rotation support (e.g. \\citealt{khoc05}). Furthermore, boxyness is often associated with X-ray and radio activities which may indicate an evolutionary history more affected by mergers (e.g. \\citealt{bender89}). Hence the detailed photometric study of early-type galaxies may yield important clues on their past merger activities.  Clues to the formation of early-type galaxies may also come from their kinematics. Studies on the global kinematics have established that elliptical galaxies as a class are supported by velocity anisotropy (e.g. \\citealt{binney76}; \\citealt{binney78}). Detailed kinematic studies often reveal kinematic distinct cores (e.g. \\citealt{emsellem04}) and non-relaxed structure (e.g. \\citealt{balce99}) which may be related to the way the galaxy is formed and to the merger history. It is possible to measure the shape of absorption lines, hence the Line-Of-Sight Velocity Distribution (LOSVD), which will tell us about the velocity anisotropy and hence constrain the orbital families (\\citealt{bn90}; \\citealt{rix92}; \\citealt{van93}). The LOSVD is often parametrised by the mean velocity $v_{rot}$ and velocity dispersion $\\sigma$, plus higher order moments ($h_{3}$, $h_{4}$ ....) of a Gauss-Hermite series. The $h_{3}$ and $h_{4}$ offer extra information on the asymmetric and symmetric deviation, respectively, away from a perfect Gaussian.  Observationally, we now have the ability to obtain spatially resolved high signal-to-noise integrated spectra out to large galactic radii and hence larger mass fractions. The high-resolution multi-band imaging data provide further insights on the radial distribution of the galaxy light and allow meaningful comparisons with spectroscopically derived results. Therefore, we are able to consider physical mechanisms acting locally and to investigate how their properties vary with the galactocentric radius.  In this series of papers we focus on the internal properties of the two early-type galaxies NGC~1407 and NGC~1400 in an effort to understand their star formation and evolutionary history. In this first paper (hereafter Paper~I), we present spatially resolved radial profiles of the kinematic parameters $v_{rot}$, $\\sigma$, $h_{3}$ and $h_{4}$ out to $\\sim$~0.6 (NGC~1407) and $\\sim$~1.3 times (NGC~1400) the galaxies' effective radii ($r_{e}$; the radius within which half the galaxy light is contained). We also present spatially resolved multi-band high-resolution wide-field imaging data out to $\\sim$~1.4$r_{e}$ for both galaxies. The spatial distributions of the galaxy light profiles are analysed and the isophotal ellipticity radial profiles and deviation parameters recovered. The residual images of the two galaxies are inspected for fine structure (e.g. dust, ripples, tidal tails, boxy or disky structure, etc.). In the second paper (\\citealt{spola07b}, hereafter Paper~II) we will present spatially resolved stellar population parameters from the same high-S/N long-slit spectroscopic data used here. The aim of Paper~II will be to interpret and combine the results from the stellar populations analysis with the kinematic and photometry results.  This paper is organised as follows. In Section~2 we describe the two sample galaxies. In Sections~3 we describe the spectroscopic and photometric observations, together with an explanation of the relevant data reduction procedures. Section~4 presents the spatially resolved radial kinematic profiles. Section~5 presents the spatially resolved surface photometry results. In Section~6 we discuss a possible interaction between NGC~1400 and NGC~1407 and the group intergalactic medium (IGM). In Section~7 we summarise our results. ",
        "conclusions": "This is the first paper of a series with the aim of studying the star formation and evolutionary history of the two early-type galaxies NGC~1407 and NGC~1400. Here, we have presented spatially resolved radial kinematics and surface photometry. The spectroscopic analysis is performed using high signal-to-noise long-slit data obtained at the ESO 3.6m telescope. The imaging study is based on high-resolution multi-band data from the HST/ACS and wide-field imaging from Subaru Suprime-Cam. The high-quality of the data allowed us to focus our study on the properties of spatially resolved galactic regions and to trace the overall radial trends. The radial profiles of the kinematic parameters $v$, $\\sigma$, $h_{3}$ and $h_{4}$ have been recovered. From the imaging analysis we obtained surface brightness and colour index profiles. We also measured the radial profiles of the isophotal shape parameters and created residual galaxy images to search for fine structure.   \\begin{list}{}{} \\item \\textbf{NGC~1407.} The kinematic study suggests a rotationally supported galaxy (or marginally anisotropic) with the presence of a possible kinematically decoupled core; the detection is uncertain and potentially may be caused by a misalignment of the slit with the nucleus. The surface brightness profiles reveal a flat core probably caused by the presence of a central supermassive black hole, with an estimated mass of $\\sim$1.03~$\\times$~10$^{9}$~$M_{\\sun}$. We find NGC~1407 to be an elliptical galaxy (A$_{4}$=0) with a mean position angle of 59$^{\\circ}$ and a small ellipticity value of 0.05. No fine structure is detected in the residual images.  \\item \\textbf{NGC~1400.} The anisotropy parameter and the kinematic profiles indicate NGC~1400 to be rotationally supported with evident minor axis rotation and flattening due to fast rotation. The galaxy was already found to have a flat core in the surface brightness profile (\\citealt{lauer95}). The galaxy isophote modelling shows elliptical isophotes at all radii (A$_{4}$=0) and no fine structure is detected in the residual galaxy images. The ellipticity is more pronounced, 0.11, and the position angle is 36$^{\\circ}$. From our results we would tend to classify NGC~1400 as an E1 galaxy, in contrast with the E$-$S0 morphological classification proposed in literature. \\end{list}  We found no evidence to support an interaction between NGC~1400 and NGC~1407 in our data. We speculate that what we are now witnessing might be the first infall of a subgroup, dominated by the NGC~1400 galaxy, into the NGC~1407 group."
    },
    "0801/0801.2529_arXiv.txt": {
        "abstract": "{ In this paper we present a study of chemical abundances in six star-forming regions. Stellar parameters and metallicities are derived using high-resolution, high S/N spectra of weak-line T-Tauri stars in each region. The results show that nearby star-forming regions have a very small abundance dispersion (only 0.033\\,dex in [Fe/H]). The average metallicity found is slightly below that of the Sun, although compatible with solar once the errors are taken into account. The derived abundances for Si and Ni show that the observed stars have the abundances typical of Galactic thin disk stars of the same metallicity. The impact of these observations is briefly discussed in the context of the Galactic chemical evolution, local inter-stellar medium abundances, and in the origin of metal-rich stars in the solar neighbourhood (namely, stars more likely to harbour planets). The implication for future planet-search programmes around very young, nearby stars is also discussed. ",
        "introduction": "As of November 2007, about 250 extra-solar planets have been discovered orbiting solar-type stars\\footnote{For a continuously updated list see tables at http://www.exoplanets.eu or http://www.exoplanet.eu}, most of which were detected thanks to the development of high precision radial-velocity instruments \\citep[for a review see][]{Udry-2007}.  Complementary high accuracy spectroscopic studies have shown that stars hosting giant  planets are particularly metal-rich when compared with single field stars \\citep[e.g.][]{Gonzalez-1997, Gonzalez-2001,Santos-2001,Santos-2004b,Santos-2005a, Fischer-2005}. This observation is helping to better understand the processes of planet formation and evolution.  The high metal content of some of the planet-hosts has raised a number of interesting questions regarding their origin. It has been proposed that the high metallicities could be the result of the in-fall of planetary (metal-rich) material into the stellar outer convective envelope \\citep[e.g.][]{Gonzalez-1998,Murray-1998}, although current results do not support this hypothesis \\citep[][]{Santos-2004b,Santos-2005a,Fischer-2005}. The excess metallicity observed likely reflects the higher (average) metal abundances of the clouds of gas and dust that originated the star and planet systems. Such conclusions are also supported by the most recent planet formation models \\citep[][]{Ida-2004b,Benz-2006}.  Because the large majority of planet hosting stars are solar-neighbourhood thin-disk objects \\citep[e.g.][]{Ecuvillon-2007}, one would expect nearby T-Tauri stars and star-forming regions [{\\sc sfr}s] that are metal-rich to exist. Their existence is expected if a well-defined age-metallicity relation exists in the solar-neighbourhood \\citep[][]{Pont-2004b,Haywood-2006}. Previous studies \\citep[][]{Padgett-1996,James-2006}, however, have not succeeded in finding evidence of high metallicity, T-Tauri stars in nearby {\\sc sfr}s. Whether this is due to the absence of such regions, or to the relatively small number of stars and {\\sc sfr}s observed was not clear.  Considering the strong metallicity-giant planet connection, the detection of nearby, metal-rich {\\sc sfr}s would provide preferential targets for future planet searches around young T-Tauri stars. The detection of planets orbiting T-Tauri stars would provide observational constraints of paramount importance for the study of planet formation scenarios, since it gives us some hint about the timescales for planetary formation.  In this paper, we continue the metallicity programme started in \\citet[][]{James-2006}, by increasing the number of observed young T-Tauri stars in each {\\sc sfr}, and by adding abundance determinations for objects that have not yet been studied. In all, we present accurate stellar parameters and chemical abundances for 19 weak line T-Tauri stars [{\\sc wtts}s] in six nearby {\\sc sfr}s, namely Chamaeleon, Corona Australis, Lupus, Rho-Ophiuchus, Taurus and the Orion Nebula Cloud. The results of this study are presented and their implications discussed. ",
        "conclusions": "In this paper, we present stellar metallicities for six nearby {\\sc sfr}s, derived from high resolution spectroscopy of young {\\sc wtts}s. The results show that all the studied {\\sc sfr}s have metallicities between $-$0.13 and $-$0.05\\,dex, with a very small scatter of 0.033\\,dex around the mean value of $<$[Fe/H]$>$=$-$0.08\\,dex (compatible to solar within the uncertainties).  The analysis of the abundances of silicon (alpha-element) and nickel (iron-peak element) also suggests that the abundance ratios in the studied {\\sc sfr}s are typical of those found in solar neighbourhood thin disk stars of similar metallicity.  Together with other results from the literature, these observations suggest that the chemical abundances in the nearby {\\sc ism} are extremely uniform and not strongly above solar. Such a conclusion may have important implications for the models of Galactic evolution.  The results presented here have also important implications on planet searches around young (active) stars. Although these searches are severely limited in current radial-velocity surveys \\citep[e.g.][]{Saar-1997,Santos-2000a,Paulson-2002}\\footnote{See however recent detection by \\citet[][]{Setiawan-2008}}, several instruments are now being built that will change this situation. In particular, new generation adaptive optics systems like the {\\sc eso sphere} planet finder or the Gemini {\\sc gpi} project, as well as a new generation of interferometric instruments like {\\sc prima} ({\\sc eso}), will allow us to study, in unprecedented detail, the existence of long period planets around young solar-type stars. The rarity of metal-rich nearby star-forming regions may limit the goals of these projects if the metallicity-giant planet connection is still present for systems with long orbital periods.  In this context it will also be very interesting to study in detail the abundances of nearby young, post-T Tauri stars \\citep[e.g.][]{Zuckermann-2004,Torres-2006}. Such studies would also help to settle the question about the metallicity of the local {\\sc ism} \\citep[e.g.][]{Sofia-2001}.  If stars with planets have originated from inner galactic regions, it may be interesting to understand how planetary evolution (and survival) may suffer with stellar orbital diffusion. This conclusion may further have impact on the understanding of the frequency of planets in the Galaxy.  As a very important by-product, the observations presented in this paper can further be used to investigate the effects of the metallicity on the properties of stellar populations such as rotation, multiplicity degree, or magnetic activity. These studies will be presented in a separate paper, James et al. (in prep.)."
    },
    "0801/0801.4413_arXiv.txt": {
        "abstract": "We investigate the tidal interaction between a low-mass planet and a self-gravitating protoplanetary disk, by means of two-dimensional hydrodynamic simulations. We first show that considering a planet freely migrating in a disk without self-gravity leads to a significant overestimate of the migration rate. The overestimate can reach a factor of two for a disk having three times the surface density of the minimum mass solar nebula. Unbiased drift rates may be obtained only by considering a planet and a disk orbiting within the same gravitational potential. In a second part, the disk self-gravity is taken into account. We confirm that the disk gravity enhances the differential Lindblad torque with respect to the situation where neither the planet nor the disk feels the disk gravity. This enhancement only depends on the Toomre parameter at the planet location. It is typically one order of magnitude smaller than the spurious one induced by assuming a planet migrating in a disk without self-gravity. We confirm that the torque enhancement due to the disk gravity can be entirely accounted for by a shift of Lindblad resonances, and can be reproduced by the use of an anisotropic pressure tensor. We do not find any significant impact of the disk gravity on the corotation torque. ",
        "introduction": "$} \\label{sec:appA}  In this section, we give the expressions of the radial and azimuthal self-gravitating accelerations $g_r$ and $g_{\\varphi}$, smoothed over the softening length $\\varepsilon_{\\rm sg}$. We use the variables ($u=\\log(r/r_{\\rm min})$, $\\varphi$), where $r_{\\rm min}$ denotes the inner edge radius of the grid. With this set of coordinates, $g_r (u,\\varphi)$ reads \\begin{eqnarray} \\left. g_r (u, \\varphi) \\right. = &-&Ge^{-u/2}\\,\\int_{0}^{u_{\\rm max}} \\int_0^{2\\pi} S_r(u^{'}, \\varphi^{'})\\,K_r(u-u^{'},\\varphi-\\varphi^{'})\\,du^{'}d\\varphi^{'}\\nonumber\\\\ &+&G\\Sigma(u,\\varphi)K_r(0, 0)\\Delta u\\Delta\\varphi, \\label{gamr} \\end{eqnarray} where $S_r$ and $K_r$ are defined as \\begin{equation} S_r (u, \\varphi) = \\Sigma (u, \\varphi) ~e^{u/2}~~~~\\rm and~~~~K_r (u, \\varphi) = \\frac{1+B^2 -e^{-u}\\cos(\\varphi)}{\\{ 2(\\cosh(u) -\\cos(\\varphi)) + B^2 e^u \\}^{3/2}}. \\label{SrKr} \\end{equation} In Eqs.~(\\ref{gamr}) and~(\\ref{SrKr}), $G$ denotes the gravitational constant, $u_{\\rm max} = \\log(r_{\\rm max}/r_{\\rm min})$ with $r_{\\rm max}$ the outer edge radius of the grid, $\\Sigma$ is the disk surface density, $\\Delta u$ and $\\Delta\\varphi$ are the mesh sizes, $K_r (0, 0)=1/B$ and $B=\\varepsilon_{\\rm sg} / r$. Since $\\varepsilon_{\\rm sg} \\propto r$ (see section~\\ref{sec:implement}), $B$ is uniform over the grid.  The second term on the R.H.S. of Eq.~(\\ref{gamr}) is an additional corrective term that ensures the absence of radial self-force. Similarly, $g_{\\varphi} (u, \\varphi)$ reads \\begin{equation} g_{\\varphi} (u, \\varphi) = -G e^{-3u/2} ~\\int_{0}^{u_{\\rm max}}  \\int_0^{2\\pi} S_{\\varphi}(u^{'}, \\varphi^{'})~K_{\\varphi}(u-u^{'}, \\varphi-\\varphi^{'}) ~du^{'} d\\varphi^{'}, \\label{gamt} \\end{equation} with $S_{\\varphi}$ and $K_{\\varphi}$ given by \\begin{equation} S_{\\varphi} (u, \\varphi) = \\Sigma (u, \\varphi) ~e^{3u/2}~~~~\\rm and~~~~K_{\\varphi} (u, \\varphi)=\\frac{\\sin(\\varphi)}{ \\{ 2(\\cosh(u) - \\cos(\\varphi)) + B^2 e^u \\}^{3/2} }. \\label{StKt} \\end{equation} In the particular case where only the axisymmetric component of the disk self-gravity is accounted for, which involves the axisymmetric component of the disk surface density $\\overline{\\Sigma}(u) = (2\\pi)^{-1} \\int_0^{2\\pi} \\Sigma(u,\\varphi)d\\varphi$, $g_{\\varphi}$ cancels out and \\begin{equation} g_r (u) = -G e^{-u/2}\\,\\int_{0}^{u_{\\rm max}} \\overline{S_r}(u^{'})\\,\\widetilde{K_r}(u-u^{'})\\,du^{'} + G \\overline{\\Sigma}(u) \\Delta u \\widetilde{K_r}(0), \\label{gamrzm} \\end{equation} where $\\overline{S_r}(u) = (2\\pi)^{-1} \\int_0^{2\\pi} S_r (u,\\varphi)d\\varphi$ and $\\widetilde{K_r}(u) = \\int_0^{2\\pi} K_r (u,\\varphi)d\\varphi$. ",
        "conclusions": ""
    },
    "0801/0801.4139_arXiv.txt": {
        "abstract": "{Determining the evolutionary stage of a Young Stellar Object (YSO) is of fundamental importance to test star formation theories. Classification schemes for YSOs are based on evaluating the degree of dissipation of the  surrounding envelope, whose main effects are the extinction of the optical radiation from the central YSO and re-emission in the far--infrared to millimeter part of the electromagnetic spectrum. Since extinction is a property of column density along the line of sight, the presence of a protoplanetary disk may lead to a misclassification of pre--main sequence stars with disks when viewed edge--on.} {We performed radiative transfer calculations to show the effects of different geometries on the main indicators of YSO evolutionary stage. In particular we tested not only the effects on the infrared colors, like the slope \\alf \\ of the flux between 2.2 and 24~\\mum , but also on other popular indicators of YSO evolutionary stage, such as the bolometric temperature and the optical depth of silicates and ices.} {We used the axisymmetric 3D radiative transfer codes RADMC and RADICAL to calculate the spectral energy distribution including silicates and ice features in a grid of models covering the range of physical properties typical of embedded and pre--main sequence sources.} {Our set of models compares well with existing observations, supporting the assumed density parametrization and the adopted dust opacities. We show that for systems viewed at intermediate angles (25\\degr\\ - 70\\degr ) the ``classical''  indicators of evolution are able to classify the degree of evolution of young stellar objects since they accurately trace the envelope column density, and they all agree with each other. On the other hand, edge-on system are misclassified for inclinations larger than $\\sim 65$\\degr $\\pm 5$\\degr\\ ,  where the spread is mostly due to the range of mass and the flaring degree of the disk. In particular, silicate emission, typical of pre--main sequence stars with disks, turns into silicate absorption when the disk column density along the line of sight reaches $1 \\cdot 10^{22}$\\persc, corresponding e.g. to a $5 \\cdot 10^{-3}$ \\msun \\ flaring disk viewed at 64\\degr. A similar effect is noticed in all the other classification indicators studied: \\alf , \\tbol , and the \\h2o\\ and \\co2\\ ice absorption strengths. This misclassification has a large impact on the nature of the flat--spectrum sources ($\\alpha \\simeq 0$), whose number can be explained by simple geometrical arguments without invoking evolution. A reliable classification scheme using a minimal number of observations is constituted by observations of the millimeter flux with both a single dish and an interferometer.} {} ",
        "introduction": "A correct classification of the evolutionary stages of young stellar objects (YSOs) is fundamental to understand the time--scales involved  in the star formation process. Since protostars gradually accrete and disperse their envelope material during their lifetime, the amount of residual envelope material can be used to measure the degree of evolution of the protostellar object. The presence of a surrounding envelope has several characteristic effects on the observables of the YSO which can then be linked to the evolutionary stage of the system. In particular, the envelope strongly affects the spectral energy distribution (SED) of the protostellar system, extinguishing and reprocessing the protostellar radiation. Thus, instead of a single blackbody emission, the SED of a young protostellar system appears as a combination of two thermal components: one corresponding to the envelope emission ($T_{\\rm eff}< 100$~K) and the other that is the result of the emission of the central protostellar object ($T_{\\rm eff}> 2000$~K) extinguished and scattered by the surrounding envelope. It is evident that measuring the ``relative intensity'' of the two components would furnish the contribution of the envelope to the whole system. Indeed, low--mass protostellar systems were originally classified according to the amount of extinction towards the protostar, using the ratio between near (2~\\mum ) and mid-infrared fluxes (25-60~\\mum ) \\citep[\\alf$=d{\\mathrm log}(\\lambda F_{\\lambda})/d{\\mathrm log}(\\lambda)$,][]{adams1987,greene1994}, or by the relative amount of emission from the envelope, evaluated by the average frequency of emission ($\\left< \\nu \\right>=\\int \\nu F_{\\nu} d\\nu /\\int F_{\\nu} d\\nu$) converted to a blackbody temperature  \\citep[\\tbol,][]{myers1993}. The combination of observations and theory has led to a three-class classification with class I representing the least evolved protostars with a significant fraction of the mass still left in the envelope and showing \\alf$>$0  and \\tbol$<$650 K;  class II, corresponding to the classical T Tauri stars with disks dominating the infrared emission and with $-2<$\\alf$<0$  and 650~K$<$\\tbol$<$2800~K; and class III pre-main sequence stars which show the presence of circumstellar disks only in the mid--infrared and have \\alf$<-2$  and \\tbol$>$2800 K. The earlier class 0 phase was added by \\citet{andre1993} to include the most deeply embedded protostars that are invisible in the infrared. Finally, some authors refer to the sources with $-0.3<$\\alf$<0.3$ as flat--spectrum sources and view them as transitional objects between class I and II \\citep[e.g.,][]{greene1994}.  This classification has had an enormous success and a huge impact on star formation studies, but it is subject to misclassification. In particular, it does not take into account that the dissipation of the envelope is not a spherical process; in fact, while part of the material is dispersed in the outflow, the bulk of it is accreted into a disk, with the result that systems viewed edge--on would present extinction values typical of younger systems with more massive envelopes, thus leading to a misclassification. This effect is clearly shown in the theoretical work of, e.g., \\citet{whitney2003a, whitney2003b} and \\citet{robitaille2006}. To emphasize the difference between the observational classification and the intrinsic properties of the YSO, \\citet{robitaille2006} introduced the alternative nomenclature ``stage I, II, III\", equivalent to the Lada classes, but referring to the true nature of the source regardless of its observed properties. These authors distinguish stage I from stage II  YSOs according to their accretion rate of the envelope, with the dividing line set at \\mdot $_{\\rm env}=10^{-6}$~\\msun ~$\\rm yr^{-1}$ (stage II and III are defined according to the accretion rate of the disk instead of that of the envelope). In the equations for a rotationally flattened envelope, it easily can be seen that \\mdot $_{\\rm env}$ is ultimately a measure of the volume density of the envelope at the centrifugal radius; for a 1~\\msun \\ star with a 300~AU centrifugal radius   \\mdot $_{\\rm env}=10^{-6}$~\\msun ~$\\rm yr^{-1}$ is equivalent to an H$_2$ number density of $1.8 \\cdot 10^5$~\\percc . \\\\ A possibility to observationally disentangle the extinction coming from envelopes to that produced in the disks can come from their very different temperature structures, with disks going from 30 to 1000~K and envelopes from 10 to 200~K. With such different conditions, dust grains are expected to be very different. Specifically, the envelope grains are expected to be coated by ice unlike the fraction of disk grains with temperatures above 100~K. Thus, comparing the evolutionary stage inferred by the infrared colors, like \\alf , with  \\tbol\\, or the silicates and ice optical depths could provide a tool to distinguish between stage I YSO and edge--on disks. Including the presence of ices (\\h2o\\ and \\co2) in our opacity tables enables us to predict not only the infrared colors, as efficiently presented in the recent literature \\citep{robitaille2006}, but also the optical depth of spectral features, in an attempt to broaden the number of observables for YSO evolution diagnostics.  In this paper, we present model predictions for a series of YSOs covering a large range of disk mass, envelope mass and stellar luminosity, and exploring the effect of inclination on the emerging SED and spectral features.  Section~\\ref{mod} presents the radiative transfer model we used and the adopted ice--coated dust properties; Section~\\ref{res} compares the results of the calculation with existing data; Section~\\ref{con} summarizes our conclusions. ",
        "conclusions": "\\label{con} We used a Monte Carlo-based radiative transfer code to calculate the emission for a grid of models representing stage I and II young stellar objects: embedded YSOs and pre--main sequence stars with disks. The aim was to study the behaviour of  classification parameters that are used to distinguish YSOs embedded in envelopes from those surrounded by disks only. We can summarize our results in these main points: \\begin{itemize} \\item The classical indicators of YSO evolution, \\alf ,\\tbol , and the silicate emission/absorption can describe very well the disappearance of the protostellar envelope for those systems seen at intermediate inclinations (approximately from 25\\degr\\ to 65\\degr ), accounting for more than half of the sources. Interestingly the transition between class I and class II measured with \\alf , \\tbol , or with the silicate feature happens for the same line of sight column density of $\\sim1 \\cdot 10^{22}$~\\persc \\  (corresponding to \\menv $\\sim$ 0.2~\\msun ), which means that these methods are well calibrated relative to each other. \\item  Each of these indicators confuses pre--main sequence stars with edge-on disks with heavily embedded YSOs if the column density exceeds $1 \\cdot 10^{22}$~\\persc ; this column density is reached for inclinations greater than 60\\degr --70\\degr , depending on the mass of the disk. \\item Even a pre--main sequence star surrounded by disk of only $5 \\cdot 10^{-4}$ \\msun\\ is classified as a flat--spectrum source for inclinations greater than 70\\degr. This means that 34\\% of the stage II sources are misclassified and could account for the entire population of flat--spectrum sources observed in nearby star--forming regions. If 10\\% of pre--main sequence disk objects are misclassified as class I, the derived time-scale for the embedded phase is overestimated by a factor of two, when using the ratio of class I/II blindly. \\item \\co2\\ and \\h2o\\ ice optical depth correlate well with \\alf\\ and  \\tbol, providing an alternative way to classify YSOs. For edge-on sources, the strong effect of scattering saturates the continuum around the features, thus affecting the feature strength and hiding the effect of temperature structure along the line of sight. \\item A combination of  single dish and interferometer millimeter continuum observations is able to trace accurately, and with the minimum number of observations, the amount of remaining envelope in emission and thus the true evolutionary stage of the YSO. \\end{itemize}"
    },
    "0801/0801.1236_arXiv.txt": {
        "abstract": "{The blazar \\object{AO 0235+164} was claimed to show a quasi-periodic behaviour in the radio and optical bands in the past, with the main outbursts repeating every 5--6 years. However, the predicted 2004 outburst did not occur, and further analysis suggested a longer time scale, according to which the next event would have occurred in the 2006--2007 observing season. Moreover, an extra emission component contributing to the UV and soft X-ray flux was detected, whose nature is not yet clear. An optical outburst was observed in late 2006 -- early 2007, which triggered a Whole Earth Blazar Telescope (WEBT) campaign as well as target of opportunity (ToO) observations by the Swift satellite.} {In this paper, we present the radio-to-optical data taken by the WEBT together with the UV data acquired by the UltraViolet and Optical Telescope (UVOT) instrument onboard Swift to investigate both the outburst behaviour at different wavelengths and the nature of the extra emission component.} {Multifrequency light curves have been assembled with data from 27 observatories; optical and UV fluxes have been cleaned from the contamination of the southern active galactic nucleus (AGN). We have analysed spectral energy distributions at different epochs, corresponding to different brightness states; extra absorption by the foreground galaxy has been taken into account.} {We found the optical outburst to be as strong as the big outbursts of the past: starting from late September 2006, a brightness increase of $\\sim 5$ mag led to the outburst peak in February 19--21, 2007. We also observed an outburst at mm and then at cm wavelengths, with an increasing time delay going toward lower frequencies during the rising phase. Cross-correlation analysis indicates that the 1 mm and 37 GHz flux variations lagged behind the $R$-band ones by about 3 weeks and 2 months, respectively. These short time delays suggest that the corresponding jet emitting regions are only slightly separated and/or misaligned. In contrast, during the outburst decreasing phase the flux faded contemporaneously at all cm wavelengths. This abrupt change in the emission behaviour may suggest the presence of some ``shutdown\" mechanism of intrinsic or geometric nature. The behaviour of the UV flux closely follows the optical and near-IR one. By separating the synchrotron and extra component contributions to the UV flux, we found that they correlate, which suggests that the two emissions have a common origin.} {} ",
        "introduction": "\\object{AO 0235+164} at redshift $z=0.94$ is one of the best-studied BL Lac objects. The analysis of its radio and optical light curves extending over 25 years led \\citet{rai01} to suggest a quasi-periodic occurrence of the main outbursts every $5.7 \\pm 0.5$ years. The next outburst was predicted around February--March 2004, and a multiwavelength campaign was organised by the Whole Earth Blazar Telescope (WEBT)\\footnote{{\\tt http://www.to.astro.it/blazars/webt/} \\\\ \\citep[see e.g.][]{vil06,vil07,boe07,rai07b}} to follow the expected event. The radio-to-optical observations by the WEBT were complemented by the optical--UV and X-ray data acquired by the XMM-Newton satellite during 3 pointings in January and August 2004, and January 2005, and by optical spectroscopic observations with the 3.6 m Telescopio Nazionale Galileo (TNG). The results of this intense observing effort were published by (\\citealt{rai05,rai06b,rai06a,rai07a}; see also \\citealt{hag07b} for an analysis of colour variability). The predicted outburst was not observed, and time-series analysis on the light curves extended to 2005 revealed a possible characteristic variability time scale of $\\sim 8$ years. Moreover, the XMM-Newton observations suggested the presence of an extra emission component in the source spectral energy distribution (SED), in addition to the synchrotron and inverse-Compton ones. The origin of this component, peaking in the UV/soft X-ray frequency range, was ascribed either to thermal emission from an accretion disc, or to synchrotron emission from an inner jet region.  An increased radio activity was detected in the 2005--2006 observing season \\citep{bac07}, and a dramatic optical brightening was finally observed in late 2006 -- early 2007. This triggered a new WEBT multiwavelength campaign as well as ToO observations by the Swift satellite. In this paper, we present the radio-to-optical observations performed from spring 2005 (i.e.\\ the end of the period considered in \\citealt{rai06b}) to October 2007, and the data acquired by the UltraViolet and Optical Telescope (UVOT) instrument onboard Swift. X-ray data from the Swift X-ray Telescope (XRT) and Burst Alert Telescope (BAT) instruments will be presented in another paper (Kadler et al., in preparation).  This paper is organised as follows: the procedures we adopted to treat the WEBT and UVOT data are outlined in Sect.~2, with particular attention paid to the subtraction of the southern AGN contribution from the optical and UV photometry. The multifrequency light curves are presented and discussed in Sect.~3 while in Sect.~4 time lags among variations at different frequencies are derived by means of statistical analysis. Spectral energy distributions with simultaneous data from near-IR to UV are constructed in Sect.~5, with the aim of separating the synchrotron and the extra emission components and of understanding their relationship. Finally, the results of our work are discussed in Sect.~6. ",
        "conclusions": "The claim by \\citet{rai01} of a possible quasi-periodic occurrence of the major radio and optical outbursts of AO 0235+164 every $5.7 \\pm 0.5$ years led to a mobilization of the international blazar community to observe the next event, predicted to peak in early 2004. A WEBT campaign was organised, with a huge international participation \\citep{rai05,rai06b,rai06a,rai07a}, but the source remained in a faint state and time series analysis on the updated light curves suggested a possible longer period of about 8--8.5 years, delaying the occurrence of the next outburst toward the 2006--2007 observing season. An increased radio activity was registered in 2005--2006 \\citep{bac07}, and a major optical outburst was finally observed in late 2006 -- early 2007, about 8.5 years after the previous major outburst peak, thus confirming the \\citet{rai06b} prediction.  We observed the same event simultaneously in the near-IR and UV frequency ranges, and then at millimetric and centimetric wavelengths, with a progressive time delay toward lower frequencies. We estimated that the 1 mm and 37 GHz outbursts lagged behind the $R$-band one by about 3 weeks and 2 months, respectively. This latter lag appears a factor $\\sim 2$ longer than previously estimated \\citep{rai05,rai06b} when considering the historical light curves until spring 2005. This may reflect some real change either in the jet structure at still unresolved scales or in its energetics with respect to the past. Indeed, there are sources that exhibit a characteristic behavior for many years and then suddenly change \\citep[see e.g.][]{smi96,all96}. Another possibility is that lack of data for long time intervals, chiefly in the optical bands because of solar conjunctions, affects the results of the DCF analysis. Even in the case of the 1997--1998 outburst, which was intensively observed thanks to the efforts of the just-born WEBT \\citep{rai01}, the central part of the event was missed in the optical, and the dimming phase was poorly sampled at 37 GHz. As for the 2006--2007 outburst, notwithstanding the exceptional sampling, we cannot rule out that solar conjunction hid a further optical peak. This would shorten the radio lags.  In any case, the delays estimated in the present paper appear rather short if compared, for example, with the approximately twice-longer lags estimated by \\citet{vil07} for the quasar-type blazar 3C 454.3 in correspondence of its 2004--2006 exceptional outburst, and by \\citet{vil04b} and \\citet{bac06} for BL Lacertae. The short time delays in AO 0235+164 had already been noticed in the past \\citep{web00,rai01}, and this was one of the reasons why microlensing was proposed as a possible explanation of the major variability events in this source \\citep[see e.g.][and references therein]{rai07a}. In contrast, when interpreting the multifrequency variability of AO 0235+164 in the framework of the helical jet model by \\citet{vil99}, the short time delays imply that the corresponding jet-emitting regions are only slightly misaligned  \\citep{ost04}. If, besides the geometrical effect, the outburst is also produced by energetic processes inside the jet, short time lags imply that the emitting regions are also closeby.  At lower ($< 37$ GHz) radio frequencies, the outburst behaviour changes: while the rising phase is progressively delayed, as expected, the decaying phase is observed simultaneously at all wavelengths. In particular, the 22 GHz flux density reaches a maximum value similar to the 37 GHz flux density almost at the same time and then suddenly decreases. Although the observed general multifrequency properties of the outburst can be explained by models dealing with shocks propagating along an inhomogeneous jet \\citep[see e.g.][]{hug89,val92}, the emission behaviour during the dimming phase might suggest that the scenario is more complicated than what is depicted in these models. Indeed, it seems as if a kind of shutdown occurred in the jet between the mm emission region and the 37--22 GHz region. We can speculate that either a jet bending or an intrinsic power off of the perturbation may produce this effect, which could also account for the different flavours of events observed in the historical light curves of the source: the harder outbursts, which are stronger at the higher frequencies, suggest that a shutdown (of any origin) has occurred between the higher- and the lower-frequency emitting regions. Understanding whether the geometric or intrinsic scenario would be more plausible is not easy. One can envisage that if the flux fading is due to a jet bending, the perturbation could continue to travel along the jet, following a misaligned path, until it could eventually be observed at radio wavelengths as soon as the jet (helical) path turns again toward the line of sight. For those sources, for which the Very Long Baseline Interferometry (VLBI) resolution can separate these regions, we should be able to observe a brightening of a VLBI knot.  We notice that our analysis of the multifrequency light curves in terms of harder/softer events presents some similarities with the generalized shock model by \\citet{val92}, who distinguish between high- and low-peaking flares. We finally mention that \\citet{hag07a} interpreted the photometric and polarimetric variations of AO 0235+164 during the December 2006 flare as due to the propagation of a transverse shock, accompanied by a small change in the viewing angle of the jet.  Another major issue we investigated is the existence of an extra emission component mostly contributing in the UV and soft X-ray energy range of the source SED. Swift-UVOT observations during the culminating phase of the optical outburst show that the behaviour of the UV emission strictly follows that of the optical one. Moreover, when constructing SEDs with simultaneous near-IR-to-UV data, and separating the synchrotron and extra-component contributions, we found that they are correlated. Although this result is affected by uncertainties due to the decomposition procedure, it nevertheless suggests that an interpretation of the extra component in terms of radiation from the accretion disc is rather unlikely \\citep[see][for the various possible interpretations]{rai06b,rai06a}. Also the hypothesis of two independent synchrotron components  appears now inadequate. A still viable explanation could be that of anomalous absorption of the optical to near-UV emission in a spectrum that would otherwise be power-law-like from the near-IR to UV. Indeed, examples of noticeable intrinsic absorption are found when analysing the spectra of broad absorption line (BAL) quasars, where several lines between $\\sim$ 1000 and 1500 \\AA\\ (rest frame) can heavily reduce the flux in the corresponding spectral regions \\citep[see e.g.][]{tur88}. A further possibility is that the extra component is the result of inverse-Compton scattering of radio photons off the relativistic electrons producing the IR--optical emission. Indeed, as stated above, the close correlation and short time delays between the optical and radio variations suggest that a lot of radio photons are available in the optical emitting region."
    },
    "0801/0801.3469_arXiv.txt": {
        "abstract": "We introduce a differential equation for star formation in galaxies that incorporates negative feedback with a delay.  When the feedback is instantaneous, solutions approach a self-limiting equilibrium state.  When there is a delay, even though the feedback is negative, the solutions can exhibit cyclic and episodic solutions. We find that periodic or episodic star formation only occurs when two conditions are satisfied.  Firstly the delay timescale must exceed a cloud consumption timescale. Secondly the feedback must be strong.  This statement is quantitatively equivalent to requiring that the timescale to approach equilibrium be greater than approximately twice the cloud consumption timescale. The period of oscillations predicted is approximately 4 times the delay timescale.  The amplitude of the oscillations increases with both feedback strength and delay time.  We discuss applications of the delay differential equation (DDE) model to star formation in galaxies using the cloud density as a variable.  The DDE model is most applicable to systems that recycle gas and only slowly remove gas from the system. We propose likely delay mechanisms based on the requirement that the delay time is related to the observationally estimated time between episodic events.  The proposed delay timescale accounting for episodic star formation in galaxy centers on periods similar to $P\\sim 10$ Myrs, irregular galaxies with $P\\sim 100$ Myrs, and the Milky Way disk with $P\\sim 2$ Gyr, could be that for exciting turbulence following creation of massive stars, that for gas pushed into the halo to return and interact with the disk and that for spiral density wave evolution, respectively. ",
        "introduction": "Gas present in a galaxy fuels star formation or nuclear black hole growth.  However both star formation and active galactic nuclei then release energy and momentum into the interstellar medium (ISM). Consequently the activity can suppress subsequent star formation. The process in which part of the output of a system is returned to its input and influences its further output is termed ``feedback.\" Early studies showed that when feedback by radiative heating is taken into account during gas accretion onto a central mass, steady solutions may not exist \\citep{ostriker76} and the feedback process can cause oscillations or periodic bursts of accretion \\citep{cowie78}. Simulations taking into account feedback processes illustrate that gas flows and star formation in galaxies can exhibit episodic or cyclic behavior \\citep{dong03,pelupessy04,ciotti07,stinson07} or alternatively can asymptotically converge onto a self-regulated equilibrium state \\citep{andersen00,robertson08}.  Galaxies display complex star formation histories. Studies of irregular galaxy populations (e.g., \\citealt{tosi91,smeckerhane94,dohmpalmer02,dolphin03,skillman05,young07,dellenbusch08}), the Milky Way disk \\citep{rochapinto00}, galaxy centers \\citep{bland03,walcher06,cecil01} and the statistics of distant galaxies \\citep{glazebrook99} infer that multiple events of vigorous star formation, separated by millions to billions of years, can occur even in isolated galactic systems. Other studies (e.g., \\citealt{vanzee01,skillman05}) find little evidence for episodic star formation. However, theoretical work has primarily focused on self-regulated star formation \\citep{andersen00,silk01,elmegreen02,monaco04,krumholz06,slyz05,li06,dib06,joung06,elmegreen07,booth07,wada07,schaye07,robertson08} and has not explored when episodic rather than a steady rate of star formation is expected.  As gas flows involving energy input, heating and cooling are complex, there is no simple way to predict when behavior is episodic or cyclic. However it is possible that average quantities can be estimated for these flows and relations based on these quantities can be used to classify their behavior. Delay differential equations can exhibit solutions that asymptotically approach a self-limiting equilibrium state and those that are periodic, even when feedback is negative. Consequently these equations can be used to differentiate between these two behaviors. Delay differential equations have been used to model biological systems with delayed negative feedback (e.g., \\citealt{wazewska88,gyori91,gurney80,kulenovic89}) but have not been applied to astrophysical systems. In this paper, using a delay differential equation, we determine when cyclic or periodic behavior is exhibited by the solutions rather than a smooth decay to a self-regulated steady state. We apply the theory to star forming galaxies, identifying delay mechanisms that could account for episodic accretion events inferred from observations. ",
        "conclusions": "It is now widely recognized that a detailed understanding of feedback and accretion processes is essential to progress in many fields of astrophysics and across the entire cosmological hierarchy, from galaxy clusters down to the scales of individual star forming regions. In order to progress, we will need huge improvements in analytic algorithms and computer power, as well as better conceptual tools for classifying complex behavior. Some processes may indeed be episodic or cyclic, while other instances may exhibit quasi-periodic cycles on the way to fully chaotic behavior. A deeper understanding requires that we should to some degree be able to distinguish between these two very different dynamical manifestations for open and closed systems.  Here we have introduced a simple differential equation model that captures some of the complexity exhibited by astrophysical star forming systems with feedback. We introduce a one dimensional DDE for the molecular cloud density that allows cloud formation to depend on the star formation rate but at a previous time.  Thus current star formation only affects the cloud distribution at a future time, we denote the delay time. The feedback is negative, so in the absence of delay there are no cyclic solutions or instabilities and all solutions asymptotically approach a self-limiting value.  We illustrate that even when the feedback is negative a delay can cause cyclic or episodic behavior. The DDE captures phenomena exhibited by astrophysical simulations of this process, including periodic solutions in some cases but not in others.   The DDE allows us for the first time to classify the solutions and predict when an astrophysical system is self-limiting or likely to exhibit periodic behavior based on timescales that are related to physical feedback and star formation processes.  We find that periodic behavior is likely when two conditions are met. First, the delay timescale must exceed the cloud consumption timescale. Secondly, the star formation must be effective at reducing the rate of formation at densities near the self-limiting or steady state value.  This is equivalent to requiring strong feedback or to requiring that the timescale to approach equilibrium be larger than approximately twice the cloud consumption timescale. We find that the amplitude of the oscillations is sensitive to the feedback strength and to a lesser extent on the ratio of the delay time to the consumption timescale.   We focus on the molecular or self-gravitating cloud density in a galaxy as the most likely variable for the DDE. This allows recycling of gas over long periods of time as gas is recycled through clouds much faster than it is depleted by star formation. The consumption timescale is set by  the lifetime of molecular clouds.  When  feedback delay times are longer than this timescale we predict episodic star formation events and with a period approximately 4 times the delay timescale.  At the present time, there are no compelling constraints on either the feedback strength or the delay time, i.e. the two key parameters of the DDE model. Thus, it is difficult to apply the model rigorously although we suggest avenues for further exploration.  There is more than one candidate for the delay time and associated feedback mechanisms, in particular, the timescale for supernovae to contribute to turbulence, the timescale for spiral density waves to evolve, and the timescale for material sent into the halo to return to interact with the disk.  We associate these three candidate delay mechanisms with possible explanations for episodic star formation events in galaxy centers (on 10 Myr timescales), the solar neighborhood (on Gyr timescales) and dwarf galaxies (on 100 Myr timescales), respectively. Using a log normal density distribution we estimate that feedback is likely to be strong enough that the second condition for episodic solutions can be satisfied.  The approach outlined here is potentially powerful framework to interpret and motivate future observations and simulations. Similar models might be applied to other accreting systems with feedback such as cooling flows. With better observationally constrained models we may be able to use similar simple dynamical models as recipes to drive simulations or  interpret statistics of astrophysical objects that exhibit episodic accretion.  Lacking currently are simulations and observational programs that constrain the timescales and strengths of possible feedback mechanisms and their functional form. In view of this uncertainty, we adopted an exponential function for the feedback process, but is this fully justified? Evidence for feedback-influenced star formation could be sought by probing for correlations between turbulence and deviations from empirical star formation laws. Other forms for the feedback function could be used, such as that of the Mackey-Glass model which can exhibit chaotic behavior \\citep{glass}. More sophisticated global theories of star formation could be developed to better predict the form of the feedback and go beyond the self-limiting equilibrium state models. Higher dimensional models could be explored, similar to those used to model predator and prey populations. By going to systems with additional variables it should be possible to model these systems without delays. The period is not strongly dependent on the amplitude of oscillation for the simple model explored here, however, this may not be true for more complex models.  \\vskip 0.5 truein  We thank Adam Frank, Eric Blackman, Jason Nordhaus and Richard Edgar for helpful discussions. Support for this work was in part provided by by NASA through awards issued by JPL/Caltech, National Science Foundation grants AST-0406823 $\\&$ PHY-0552695, the National Aeronautics and Space Administration under Grant No.$\\sim$NNG04GM12G issued through the Origins of Solar Systems Program, and HST-AR-10972 to the Space Telescope Science Institute. JBH is funded by a Federation Fellowship from the Australian Research Council.    \\appendix"
    },
    "0801/0801.4763_arXiv.txt": {
        "abstract": "We investigate the apparent discrepancy between gas and dust outer radii derived from millimeter observations of protoplanetary disks. Using 230 and 345~GHz continuum and CO J=3-2 data from the Submillimeter Array for four nearby disk systems (HD 163296, TW Hydrae, GM Aurigae, and MWC 480), we examine models of circumstellar disk structure and the effects of their treatment of the outer disk edge. We show that for these disks, models described by power laws in surface density and temperature that are truncated at an outer radius are incapable of reproducing both the gas and dust emission simultaneously: the outer radius derived from the dust continuum emission is always significantly smaller than the extent of the molecular gas disk traced by CO emission.  However, a simple model motivated by similarity solutions of the time evolution of accretion disks that includes a tapered exponential edge in the surface density distribution (and the same number of free parameters) does much better at reproducing both the gas and dust emission. While this analysis does not rule out the disparate radii implied by the truncated power-law models, a realistic alternative disk model, grounded in the physics of accretion, provides a consistent picture for the extent of both the gas and dust. ",
        "introduction": "Characterizing the gas and dust distribution in the disks around young stars is important for understanding the planet formation process, as these disks provide the reservoirs of raw material for nascent planetary systems.  A common method of modeling circumstellar disk structure is to use models described by power laws in surface density and temperature that are truncated at a particular outer radius.  This prescription has its historical roots in calculations of the minimum mass solar nebula, which indicated a surface density profile of $\\Sigma \\propto r^{-3/2}$ \\citep[e.g.][]{wei77}, as well as theoretical predictions of a radial power-law dependence of temperature for accreting disks around young stars \\citep{ada86,ada87}. Observationally, the parameterization of temperature and surface density as power-law functions of radius began with early spatially unresolved studies of continuum emission from disks \\citep{bec90,bec91}.  These models have since been refined and applied to spatially resolved observations of many disks with success \\citep[e.g.][]{mun93,dut94,lay94,dut98}, and they have proven useful for understanding the basic global properties of disk structure.  Recently, however, with the advent of high signal-to-noise, multi-frequency observations of gas and dust in protoplanetary disks, these models have begun to encounter difficulties, particularly in the treatment of the outer disk edge.  The extent of the gas and dust distribution in circumstellar disks has implications for our understanding of the planet formation process in our own solar system.  There is some evidence for a sharp decrease in the surface density of Kuiper Belt objects beyond a distance of 50 AU from the Sun \\citep{jew98,tru01,pet06}.  However, the origin of this edge is unclear. \\citet{ada04} note that the observed distance is far interior to the radius at which truncation by photoevaporation would be expected to occur, while \\citet{you02} find that the presence of such an edge in planetesimal density could be explained by drift-induced enhancement.  A compelling possibility is that the Sun formed in a cluster environment, and the early solar disk was truncated by a close encounter with a passing star \\citep[see][and references therein]{rei05}.  A more complete understanding of the outer regions of protoplanetary disks may provide insight into the processes that shape the outer solar system.  \\citet{pie05} present multiwavelength millimeter continuum and CO isotopologue observations of the disk around the Herbig Ae star AB Aurigae and found from fitting models of disk structure described by truncated power laws that the outer radius of the dust derived from continuum emission ($350\\pm30$~AU) was much smaller than that of the gas derived from $^{12}$CO J=2-1 emission ($1050\\pm10$~AU).  They suggest that a change in dust grain properties resulting in a drop in opacity could be responsible for the difference, and note the possible association with a ring feature in the disk at 200~AU. A similar result was obtained by \\citet{ise07} from observations of the disk around the Herbig Ae star HD 163296: they found a significant discrepancy between the outer radius derived for the dust continuum emission ($200 \\pm 15$ AU) and that derived from CO emission ($540 \\pm 40$ AU). These data appeared to require a sharp drop in surface density, opacity, or dust-to-gas ratio beyond 200~AU; however, as they discuss, there is no obvious physical basis for such a discontinuity.  As \\citet{ise07} demonstrate, the discrepancy in outer radii derived from the dust and gas is not simply an issue of sensitivity; the observations were sufficiently sensitive to detect emission from the power-law dust disk if it did extend to the radius indicated by the CO emission.  The underlying issue is that the truncated power law model does not simultaneously reproduce the extent of both the continuum and CO emission for these disks.  Using data from the Submillimeter Array\\footnote{The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics and is funded by the Smithsonian Institution and the Academia Sinica.} we show that the same apparent discrepancy in gas and dust outer radius applies to the circumstellar disks around several more young stars. In an attempt to understand the origin of this discrepancy, we investigate an alternative surface density profile based on work by \\citet{har98}, which is similar to a power law profile in the inner disk but includes a tapered outer edge.  We show that this model, which has a physical basis in similarity solutions of disk evolution with time, is capable of simultaneously reproducing both continuum and CO emission from these disks. The primary difference between this model and the truncated power-law disk is that instead of a sharp outer edge the surface density falls off gradually, with sufficient column density at large radii that CO emission extends beyond the point at which dust continuum emission becomes negligible. ",
        "conclusions": "With the advent of high signal-to-noise interferometer observations that resolve the outer regions of nearby protoplanetary disks, an apparent discrepancy has emerged between the extent of the dust continuum and molecular gas emission \\citep{pie05,ise07}. Using multi-frequency interferometric data from the Submillimeter Array, we have investigated this disparity for four disk systems (HD 163296, TW Hydrae, GM Aurigae, and MWC 480) in the context of two distinct classes of disk structure models: (1) a truncated power law, and (2) a similarity solution for the time evolution of an accretion disk. The primary difference between these models is in their treatment of the disk outer edge: the abruptly truncated outer edge of the power-law disk causes the visibilities to drop rapidly to zero, leading to an inferred outer radius that is small in comparison with the observed molecular gas emission.  The similarity solution, by contrast, tapers off smoothly, creating a broader visibility function and allowing molecular gas emission to persist at radii well beyond the region in the disk where continuum falls below the detection threshold.  The outer radius discrepancy appears to exist only in the context of the power-law models.  In light of this result, it appears that an abrupt change in dust properties for these disks is unlikely, as there is no physical mechanism to explain such a discontinuity. This may imply that a sharp change in dust properties in the early solar nebula is similarly an unlikely explanation for the Kuiper belt edge observed by \\citet{jew98}, and that a dynamical mechanism such as truncation by a close encounter with a cluster member \\citep[][and references therein]{rei05} may provide a more plausible origin.  In this case, we would expect to observe disks with sharp outer edges only in clustered environments, and a model with a tapered edge would be a more realistic prescription for investigating the structure of a typical isolated disk.  The tapered disk models provide a natural explanation for the disparate outer radii observed using different probes of the disk extent, including comparison of continuum and molecular gas observations \\citep{pie05,ise07}, and also comparison of different isotopologues and rotational transitions of a particular molecule \\citep{pie07}.  When predicting CO emission, this simple model does neglect potential variance in the CO abundance due to depletion in the midplane and photodissociation at the disk surface; however, the results presented are intended simply to illustrate the global differences between gas and dust emission from the two model classes, independent of detailed CO chemistry.  While we cannot rule out disparate gas and dust radii in these disks, we show that an alternative disk structure model, grounded in the physics of accretion, resolves the apparent size discrepancy without the need to invoke dramatic changes in dust opacity, dust density, or dust-to-gas ratio in the outer disk."
    },
    "0801/0801.2765_arXiv.txt": {
        "abstract": "We report the discovery of quiescent emission from molecular hydrogen gas located in the circumstellar disks of six pre-main sequence stars, including two weak-line TTS, and one Herbig AeBe star, in the Chamaeleon~I star forming region.  For two of these stars, we also place upper limits on the $2\\rightarrow1~S(1)$/$1\\rightarrow0~S(1)$ line ratios of $\\sim$~0.4 and 0.5. Of the 11 pre-main sequence sources now known to be sources of quiescent near-infrared hydrogen emission, four possess transitional disks, which suggests that detectable levels of H$_2$ emission and the presence of inner disk holes are correlated.  These H$_2$ detections demonstrate that these inner holes are not completely devoid of gas, in agreement with the presence of observable accretion signatures for all four of these stars and the recent detections of [Ne~{\\scshape ii}] emission from three of them.  The overlap in [Ne~{\\scshape ii}] and H$_2$ detections hints at a possible correlation between these two features and suggests a shared excitation mechanism of high energy photons.  Our models, combined with the kinematic information from the H$_2$ lines, locate the bulk of the emitting gas at a few tens of AU from the stars.  We also find a correlation between H$_2$ detections and those targets which possess the largest H$\\alpha$~equivalent widths, suggesting a link between accretion activity and quiescent H$_2$ emission.  We conclude that quiescent H$_2$ emission from relatively hot gas within the disks of TTS is most likely related to on-going accretion activity, the production of UV photons and/or X-rays, and the evolutionary status of the dust grain populations in the inner disks. ",
        "introduction": "\\label{intro}  The reservoirs of gas and dust found orbiting most sun-like stars with ages on the order of a million years --- also known as T Tauri stars (TTS) --- are now subjects of intense study, as they likely hold the secrets to the formation history of the solar system and other planetary systems.  The evolution of the observable characteristics of these planet forming disks and the underlying phenomena responsible for their detectability may provide the necessary insight to further our understanding of planet forming processes and to determine the likelihood of the formation of planetary systems around sun-like stars.  The existence of circumstellar disks accompanying young low mass stars previously has been inferred by the presence of detectable emission from trace constituents of these disks (i.e., dust and gaseous molecules other than H$_2$).  The earliest attempts to determine the fraction of stars that possess disks in young star forming clusters focused on their near-infrared continuum \\citep{stro1989} and submillimeter continuum emission \\citep{wein1989,beck1990}.  Circumstellar disks possessing warm ($T$~$\\sim$~1000~K) micron-sized dust grains within $\\sim$~1~AU of the pre-main sequence stars produce near-infrared emission in excess of that predicted for the stellar photosphere, while circumstellar disks possessing cold, micron-sized dust grains at radii of several tens of AU produce excess radiation at mid-infrared and submillimeter wavelengths.  From surveys such as these, astronomers have determined the lifetime of warm dust grains in the inner disks to be 3-6~Myr \\citep{lada2000,hais2000,hais2001a}.  Similarly, observations of cold ($T$~$\\sim$~10-100~K) gas molecules (e.g., CO, HCN$^+$, etc) from the outer regions of the disks \\citep{beck1993,zuck1995} tend to support the disk lifetime estimates drawn from near-infrared surveys.  More recently, new disk fraction and lifetime surveys of TTS in nearby star forming regions have been conducted utilizing the Infrared Array Camera (IRAC) and Multiband Imaging Photometer (MIPS) aboard the {\\it Spitzer Space Telescope}.  These instruments are capable of detecting emission from dust grains ($T$~$\\sim$~100-300~K) located at intermediate disk radii, a region of these disks to which the previous photometric surveys were insensitive.  \\cite{ciez2007} took a census of 230 sources identified as weak-line TTS -- weakly or non-accreting TTS possessing either a passive, optically thin disk or no circumstellar material -- in the Lupus, Ophiuchus, and Perseus molecular clouds.  They find a higher disk frequency for the sources located in the highest extinction regions of the clouds, suggesting that the wTTS at the edges of the formation regions are slightly older and more evolved.  Although they find $\\sim$~20\\% of these possess some circumstellar material, $\\sim$~50\\% of what appear to be the youngest stars in the sample show no evidence of disks.  These results, which are consistent with other {\\it Spitzer} surveys of low mass star forming regions \\citep{furl2006,geer2006,kess2006,padg2006,damj2007}, lead them to conclude, based on the existence of detectable emission from cool dust grains, that disk lifetimes are less than ten million years, which roughly agrees with the conclusions about disk lifetimes drawn from near-infrared surveys.  Given the chaotic environment near these young stars and the phenomena (i.e., accretion, photoevaporation, stellar and disk winds, close encounters with other cluster members) that conspire to destroy these disks, the lifetimes measured from these surveys can be explained by disk dissipation models including all or some of these processes \\citep{holl2000}.  Despite the growing consensus that disk lifetimes are only a few million years that has been drawn from several different techniques for studying circumstellar disks, astronomers must remain cautious when inferring the presence or absence of disks based upon the detectability of disk tracers (i.e., dust and CO emission) because the majority of the disk mass is composed of molecular hydrogen (H$_2$).  As these disks evolve according to the core accretion model of planet formation, the small micron-sized dust grains accumulate into larger bodies, which will change the wavelength dependence of the dust emissivities and the collective surface area of dust grains emitting at wavelengths in the near-infrared.  Since a planetesimal must reach a mass of $\\simeq$~10 earth masses before it can gravitationally collect H$_2$ gas from the disk, the main component of the disk must persist, while planetesimals accrete into km-sized objects and disk tracers disappear, in order for gas giant planets to be the result of the core accretion scenario of planet formation.  Therefore, in an accretionally-evolved system in which small dust grains have coagulated into larger planetesimals, an evolutionary stage should exist during which both the near-infrared and mid-infrared excesses will have diminished substantially while the massive gaseous component remains present.  In previous work, \\cite{wein2000}, \\cite{bary2002}, and \\cite{bary2003a} searched a variety of TTS, predominantly wTTS, in the nearby Taurus-Auriga, $\\rho$ Ophiuchus, and TW Hya Association star forming regions for emission signatures from H$_2$ persisting in such accretionally evolved circumstellar disks.  Using various high-resolution, near-infrared spectrometers, narrow line emission centered at the $v=1\\rightarrow0~S(1)$ transition rest wavelength of 2.1218~$\\mu$m was detected from several TTS stars; some were known to possess circumstellar disks (TW~Hya, GG~TauA, and LkCa~15), and one was previously thought to be devoid of any circumstellar material (DoAr~21).  The detection of quiescent emission from the disk of a wTTS with no other detectable disk tracers suggests that the gaseous component may indeed persist as the disk evolves.  However, we estimate the mass of hot H$_2$ gas producing the observed features in these sources to be in the range 10$^{-12}$-10$^{-10}$~M$_\\odot$.  For at least three of these four sources, a substantial reservoir of H$_2$ gas exists in their disks suggesting that the near-infrared H$_2$ emission feature, is itself only a tracer of the disk gas.  In the time since these detections were made, several other authors searching for the same emission signature also have detected H$_2$ gas \\citep{itoh2003,taka2004,rams2007}.  In two of the three additional sources toward which H$_2$ was detected (LkH$\\alpha$~264 and ECHA J0843.3-7905), the emission lines were similarly centered at 2.1218~$\\mu$m, with full-width at half of maximum (FWHM) only slightly larger than those reported for the previous four detections.  In the case of DG~Tau, \\cite{taka2004} reported the detection of an H$_2$ emission line blueshifted by 15~km~s$^{-1}$ and possessing a FWHM of $\\sim$~35~km~s$^{-1}$.  The authors also show that the emitting gas is extended along the axis of a strong outflow previously observed from this source.  Therefore,  the emitting gas from this source is most likely associated with the outflowing gas distinguishing this detection from the others.  In addition to these detections of rovibrational emission, \\cite{sher2003}, \\cite{bitn2007} and \\cite{mart2007} have detected H$_2$ emission from the mid-infrared pure rotational transitions from the Herbig AeBe stars, AB~Aur and HD~97048, with a similarly quiescent kinematic signature.  Here, we present five more detections of H$_2$ near-infrared emission from both classical and weak-line TTS, an infrared companion (IRC), and a detection from a Herbig AeBe star in the Chamaeleon I star forming region.  In addition, we place an upper limit on emission from a second line of H$_2$ at 2.2477~$\\mu$m for two of these sources.  We review the possible stimulation mechanisms with respect to recent X-ray and UV irradiated disk models.  We find that two of our past detections and one of the present detections possess transitional disks --- disks with large inner disk holes devoid of emission from small warm dust grains.  Two of these transitional disk sources also possess detectable emission from a mid-infrared [Ne~{\\scshape ii}] emission line predicted to be a tracer of high energy photons and a signature of photoevaporation, a mechanism thought to be potentially responsible for the presence of the inner disk holes of the transitional disks.  Kinematic information from the H$_2$ emission features suggest that the emitting gas extends to within a few AU of each source, demonstrating for the sources with inner disk holes that gas remains despite the lack of dust emission.  We estimate the location of the regions within the disks where potential excitation photons deposit their energies and find agreement with the location of emitting gas determined from the velocity dispersion of the line.  Finally, we relate the detection of quiescent H$_2$ emission to accretion activity and, possibly, the evolutionary status of the dust grain population. ",
        "conclusions": "As part of a high resolution near-infrared spectroscopic survey for H$_2$ emission from TTS in nearby star forming regions, we have made new detections of emission features centered on the 2.1218~$\\mu$m $v=1\\rightarrow0~S(1)$ transition of H$_2$ in the spectra of two cTTS (CS~Cha and CV~Cha), two wTTS (Glass Ia and Sz~33), one IRC (Glass Ib), and one Herbig AeBe star (HD~97048) located in the Chamaeleon~I dark cloud.  The detection of H$_2$ we report toward HD~97048 is the first detection of near-infrared emission from quiescent H$_2$ toward a Herbig AeBe star.  In every case, the central velocity of the emitting gas is found to be coincident with the systemic velocity of the source.  Assuming a circumstellar disk origin for these emission lines, an assumption well-supported by the central velocities of these lines falling at the rest velocities of the respective stars and by the narrow line widths, we used the HWZI and the FWHM velocities of these Gaussian-shaped emission features to estimate the innermost radii to which the gas extends and the radii at which the bulk of the emitting gas resides .  We find that the innermost radii for all of the stars is on the order of $\\sim$~1~AU.  We further find that the radii about which the bulk of the emission is centered is on the order of $\\sim$~10~AU.  For the binary source (Glass I), H$_2$ emission extends between both components and to the south of the primary.  Although the extended emission suggests that the gas could be entrained in an outflow, the lack of a detectable velocity gradient along the would-be axis of the outflow makes this an unlikely scenario.  However, a circumbinary disk origin for the extended emission is also difficult to reconcile with the large FWHM of the emission lines measured for the extended emission.  The initial goals for this survey was to determine if wTTS lacking detectable emission from standard disk tracers still harbor protoplanetary disks capable of or in the process of forming planetary systems.  Our survey revealed two such stars, making a total of three, including DoAr~21 discovered by us in a previous survey.  The range of H$_2$ line fluxes (6.9$\\times$10$^{-16}$~$\\le$~F$_{2.12}$~$\\le$~8.6$\\times$10$^{-15}$~ergs~cm$^{-2}$~s$^{-1}$) and masses (10$^{-10}$~$\\le$~M$_{H_2}$~$\\le$~10$^{-11}$~M$_\\odot$) measured for the Chamaeleon~I sources with detected H$_2$ line emission were found to agree with all previous detections of quiescent H$_2$ emission.  These trace amounts of gas do not directly inform us about the total disk masses harbored by these systems, though almost certainly the mass of H$_2$ detected directly through line emission at 2.1218~$\\mu$m represents only a small fraction of the total mass in these disks.  Our disk models suggest that the total disk masses could be 10$^2$ to 10$^7$ times greater than the mass measured directly in this single emission line.  Assuming that shock excitation is not the most likely mechanism for stimulating narrow line emission from H$_2$ gas which is not Doppler shifted from the systemic velocity of the source, we suggest that UV fluorescence and/or X-rays are the most likely stimulation mechanisms responsible for the observed H$_2$ emission.  We calculated optical depth surfaces for UV photons and X-rays penetrating a disk composed entirely of H$_2$ to determine where in a core-accretionally evolved disk these photons deposit their energy and possibly generate H$_2$ emission.  These models confirm that the radiation is most likely to be absorbed by gas in the upper atmospheres of these disks at intermediate radii and show that these stimulation mechanisms require significantly dense and massive columns of gas to produce the observed emission features at radii coincident with the those determined from the kinematic information.  Many improvements may be made to these calculations by including opacities associated with large dust grains, other gaseous species, or a disk model that better approximates an irradiated transitional disk.  However, the results from these calculations suggest that a gas-rich inner disk is necessary to produce quiescent H$_2$ emission via UV photons and X-rays, and that the stimulated gas resides at high disk altitudes and intermediate disk radii.  All four sources in our survey to-date known to possess transitional disks with optically thin inner disk holes are also H$_2$ detections.  These detections place gas within the inner disk holes similar to the resolved and unresolved detections of [Ne {\\scshape ii}] emission lines from the sources.  The presence of gas in the inner disk holes confirms the presence of gas within these gaps necessary to explain the accretion activity observed for these sources.  The correlation between sources possessing quiescent H$_2$ emission and [Ne {\\scshape ii}] suggests a shared stimulation mechanism for both of these lines, a relationship to the photoevaporation process, and a distinct stage in the evolution of these circumstellar disks.  We also find that most sources in our survey with H$\\alpha$~EWs sufficient to classify them as cTTS have detectable H$_2$ emission.  The only exception is WW~Cha, which may also be a detection, though we have cautiously classified it as a `possible' detection, sine the H$_2$ emission towards this source, if real, is only at the 2.4$\\sigma$ level.  The apparent relationship between H$\\alpha$ emission and accretion activity suggests that the quiescent H$_2$ emission may likewise be correlated with accretion in agreement with the findings of \\cite{carm2007}.  This is not entirely surprising given that the UV photons necessary for stimulating H$_2$ emission are likely generated by accretion hot spots on the surfaces of the stars.  In fact, this is strong evidence that UV fluorescence undoubtedly contributes to the observed H$_2$ emission.  Given the dependence of the UV penetration depths on the sizes of the dust grains in the disk, illustrated by \\cite{nomu2007}, and the overlap between H$_2$ detections and sources possessing transitional disks, we suggest that the evolutionary stage of the dust grain population in the inner regions of these disks is also correlated with the production of quiescent H$_2$ emission.  Based upon these findings, we propose that Sz~33 and DoAr~21, both H$_2$ sources that are borderline c/wTTS, likely harbor transitional disks."
    },
    "0801/0801.3555_arXiv.txt": {
        "abstract": "{} { Observations of shell-type supernova remnants (SNRs) in the GeV to multi-TeV $\\gamma$-ray band, coupled with those at millimetre radio wavelengths, are motivated by the search for cosmic-ray accelerators in our Galaxy. The old-age mixed-morphology SNR W~28 (distance $\\sim$2~kpc) is a prime target due to its interaction with molecular clouds along its northeastern boundary and other clouds situated nearby.} { We observed the W~28 field (for $\\sim$40~h) at very high energy (VHE) $\\gamma$-ray energies ($E>$0.1~TeV) with the H.E.S.S. Cherenkov telescopes. A reanalysis of EGRET $E>$100~MeV data was also undertaken. Results from the NANTEN 4m telescope Galactic plane survey and other CO observations were used to study molecular clouds. } { We have discovered VHE $\\gamma$-ray emission (HESS~J1801$-$233) coincident with the northeastern boundary of W~28 and a complex of sources (HESS~J1800$-$240A, B and C) $\\sim$0.5$^\\circ$ south of W~28 in the Galactic disc. The EGRET source (GRO~J1801$-$2320) is centred on HESS~J1801$-$233 but may also be related to HESS~J1800$-$240 given the large EGRET point spread function. The VHE differential photon spectra are well fit by pure power laws with indices $\\Gamma \\sim 2.3$ to 2.7. The spectral indices of HESS~J1800$-$240A, B, and C are consistent within statistical errors. All VHE sources are $\\sim$10$^\\prime$ in intrinsic radius except for HESS~J1800$-$240C, which appears pointlike. The NANTEN $^{12}$CO(J=1-0) data reveal molecular clouds positionally associating with the VHE emission, spanning a $\\sim$15~km~s$^{-1}$ range in local standard of rest velocity. } { The VHE/molecular cloud association could indicate a hadronic origin for HESS~J1801$-$233 and HESS~J1800$-$240, and several cloud components in projection may contribute to the VHE emission. The clouds have components covering a broad velocity range encompassing the distance estimates for W~28 ($\\sim$2~kpc) and extending up to $\\sim$4~kpc. Assuming hadronic origin and distances of 2 and 4~kpc for cloud components, the required cosmic-ray density enhancement factors (with respect to the solar value) are in the range $\\sim$10 to $\\sim$30. If situated at 2~kpc distance, such cosmic-ray densities may be supplied by SNRs like W~28. Additionally and/or alternatively, particle acceleration may come from several catalogued SNRs and SNR candidates, the energetic ultra compact HII region W~28A2, and the HII regions M~8 and M~20, along with their associated open clusters. Further sub-mm observations would be recommended to probe in detail the dynamics of the molecular clouds at velocites $>$10~km~s$^{-1}$ and their possible connection to W~28. } ",
        "introduction": "The study of shell-type supernova remnants (SNRs) at $\\gamma$-ray energies is motivated by the long-held idea that they are the dominant sites of hadronic Galactic cosmic-ray (CR) acceleration to energies approaching the \\emph{knee} ($\\sim 10^{15}$~eV) (e.g. Ginzburg \\& Syrovatskii \\cite{Ginzburg:1}, Blandford \\& Eichler \\cite{Blandford:1}). CRs (hadrons and electrons) are injected into the SNR shock front, and are then accelerated via the diffusive shock acceleration (DSA) process (for a review see Drury \\cite{Drury:2}). Subsequent $\\gamma$-ray production from the interaction of these CRs with ambient matter and/or electromagnetic fields is a tracer of such non-thermal particle acceleration, and establishing the hadronic or electronic nature of the parent CRs in any $\\gamma$-ray source remains a key issue. Two SNRs, RX~J1713.7$-$3946 and RX~J0852.0$-$4622, have so far established shell-like morphology in VHE $\\gamma$-rays (Aharonian \\etal \\cite{HESS_RXJ1713,HESS_VelaJnr,HESS_RXJ1713_II,HESS_VelaJnr_II,HESS_RXJ1713_III}), with spectra extending to 20~TeV and beyond. In particular for RX~J1713.7$-$3946, particle acceleration up to at least 100~TeV is inferred from the H.E.S.S. observations. Although a hadronic origin of the VHE $\\gamma$-ray emission is highly likely in the above cases (Aharonian \\etal \\cite{HESS_RXJ1713_II,HESS_VelaJnr_II}, Berezhko \\& V\\\"olk \\cite{Berezhko:1}, Berezhko, P\\\"uhlhofer \\& V\\\"olk \\cite{Berezhko:2}), an electronic origin is not ruled out.  Disentangling the electronic and hadronic components in TeV SNRs may be made easier by studying: (1) SNR $\\gamma$-ray spectra well beyond $\\sim$10 TeV, an energy regime where electrons suffer strong radiative energy losses and due to Klein-Nishina effects the resulting inverse-Compton spectra tend to show a cut-off; (2) older SNRs (age approaching 10$^5$ yr) in which accelerated electrons have lost much of their energy through radiative cooling and do not reach multi-TeV energies; (3) SNRs interacting with adjacent molecular clouds of very high densities $n> 10^3$~cm$^{-3}$. It is the latter regard especially (and to a certain degree the second) which makes the SNR W~28 (G6.4$-$0.1) an attractive target for VHE $\\gamma$-ray studies. In this paper we outline the discovery of VHE $\\gamma$-ray emission from several sites in the W~28 field and briefly discuss their relationship with molecular clouds, W~28, and other potential particle accelerators in the region.  W~28 (G6.4$-$0.1) is a mixed-morphology SNR, with dimensions 50$^\\prime$x45$^\\prime$ and an estimated distance between 1.8 and 3.3~kpc (eg. Goudis \\cite{Goudis:1}, Lozinskaya \\cite{Lozinskaya:1}). It is an old-age SNR (age 35000 to 150000~yr; eg. Kaspi \\etal \\cite{Kaspi:1}), thought to have entered its radiative phase of evolution (eg. Lozinskaya \\cite{Lozinskaya:1}) in which much of its CRs have escaped into the surrounding interstellar medium (ISM). We note also that the evolutionary status (Sedov and/or radiative) of shell-type SNRs may depend on the density of their surroundings (see eg. Blondin \\etal \\cite{Blondin:1}).  W~28 is distinguished by its interaction with a molecular cloud (Wootten \\cite{Wootten:1}) along its north and northeastern boundaries. This interaction is traced by the high concentration of 1720~MHz OH masers (Frail \\etal \\cite{Frail:2}, Claussen \\etal \\cite{Claussen:1,Claussen:2}), and also the location of very high-density ($n>10^3$~cm$^{-3}$) shocked gas (Arikawa \\etal  \\cite{Arikawa:1}, Reach \\etal \\cite{Reach:1}). The shell-like radio emission (Long \\etal \\cite{Long:1}, Dubner \\etal \\cite{Dubner:1}) peaks at the northern and northeastern boundaries where interaction with the molecular cloud is established. Further indication of the influence of W~28 on its surroundings is the expanding HI void at a distance $\\sim$1.9~kpc (Vel\\'azquez \\etal \\cite{Velazquez:1}). The X-ray emission, which overall is well-explained by a thermal model, peaks in the SNR centre but has local enhancements in a region overlapping the northeastern SNR/molecular cloud interaction (Rho \\& Borkowski \\cite{Rho:2}).  In the neighbourhood of W~28 are the radio-bright HII regions M~20 (Trifid Nebula at $d \\sim$1.7~kpc Lynds \\etal \\cite{Lynds:1} -- with open cluster NGC~6514), M~8 (Lagoon Nebula at $d\\sim 2$~kpc Tothill \\etal \\cite{Tothill:1} --- containing the open clusters NGC~6523 and NGC~6530) and the ultra-compact HII region W~28A2, all of which are representative of the massive star formation taking place in the region. Further discussion concerning the active star formation in this region may be found in van den Ancker \\etal (\\cite{Ancker:1}) and references therein. Additional SNRs in the vicinity of W~28 have also been identified: G6.67$-$0.42 and G7.06$-$0.12 (Yusef-Zadeh \\etal \\cite{Yusef:1}), G5.55+0.32, G6.10+0.53 and G7.20+0.20 (Brogan \\etal \\cite{Brogan:1}). The pulsar PSR~J1801$-$23 with % spin-down luminosity $\\dot{E} \\sim 6.2\\times 10^{34}$ erg~s$^{-1}$ and distance $d = 13.5$~kpc (based on its dispersion measure) is at the northern radio edge (Kaspi \\cite{Kaspi:1}). More recent discussion (Claussen \\etal \\cite{Claussen:3}) assigns a lower limit of 9.4$\\pm$2.4~kpc for the pulsar distance.  W~28 has also been linked to $\\gamma$-ray emission detected at $E>300$~MeV by COS-B (Pollock \\cite{Pollock:1}) and $E>100$~MeV by EGRET (Sturner \\& Dermer \\cite{Sturner:1}, Esposito \\etal \\cite{Esposito:1}, Zhang \\etal \\cite{Zhang:1}). The EGRET source, listed in the 3rd catalogue (Hartman \\etal \\cite{Hartman:1}) as 3EG~J1800$-$2338, is positioned at the southern edge of the radio shell. We have also performed an analysis of EGRET data, with additional data not included in the 3rd catalogue, and results are discussed later in this paper.  Previous observations of the W~28 region at VHE energies by the CANGAROO-I telescope revealed no evidence for such emission (Rowell \\etal \\cite{Rowell:1}) and upper limits at the $\\sim$0.2 to 0.5 Crab-flux level for energies $E>1.5$~TeV (1.1 to 2.9$\\times 10^{-11}$~erg~cm$^{-2}$~s$^{-1}$) were set for various regions. ",
        "conclusions": "\\label{sec:conclusion} In conclusion, our observations with the H.E.S.S. $\\gamma$-ray telescopes have revealed VHE $\\gamma$-ray sources in the field of W~28 which positionally coincide well with molecular clouds. HESS~J1801$-$233 is seen toward the northeast boundary of W~28, while HESS~J1800$-$240 situated just beyond the southern boundary of W~28 comprises three components. Our studies with NANTEN $^{12}$CO(J=1-0) data show molecular clouds spanning a broad range in local standard of rest velocity $V_{\\rm LSR}=$5 to $\\sim$20~km~s$^{-1}$, encompassing the distance estimates for W~28 and various star formation sites in the region. If connected, and at a distance $\\sim$2~kpc, the clouds may be part of a larger parent cloud possibly disrupted by W~28 and/or additional objects related to the active star formation in the region. Cloud components up to $\\sim$4~kpc distance ($V_{\\rm LSR}>$10~km~s$^{-1}$) however, remain a possibility.  The VHE/molecular cloud association could indicate a hadronic origin for the VHE sources in the W~28 field. Under assumptions of connected cloud components at a common distance of 2~kpc, or, alternatively, separate cloud components at 2 and 4~kpc, a hadronic origin for the VHE emission implies cosmic-ray densities $\\sim$10 to $\\sim$30 times the local value. W~28 could provide such densities in the case of slow diffusion. Additional and/or alternative particle accelerators such as HII regions representing very young stars, other SNRs/SNR candidates and/or several open clusters in the region may also be contributors. Alternatively, if cloud components at $V_{\\rm LSR}>$10~km~s$^{-1}$ are at distances $d\\sim4$~kpc, as-yet undetected particle accelerators in the Scutum-Crux arm may be responsible. Detailed modeling (beyond the scope of this paper), and further multiwavelength observations of this region are highly recommended to assess further the relationship between the molecular gas and potential particle accelerators in this complex region, as well as the nature of the acclerated particles. In particular, further sub-mm observations (eg. at high CO transitions) will provide more accurate cloud mass estimates, and allow to search for disrupted/shocked gas towards the southern VHE sources. Such studies will be valuable in determining whether or not W~28 and other energetic sources have disrupted molecular material at line velocities $>$10~km~s$^{-1}$."
    },
    "0801/0801.4696_arXiv.txt": {
        "abstract": "Chaotic inflation predicts a large gravitational wave signal which can be tested by the upcoming Planck satellite.  We discuss a SUGRA implementation of chaotic inflation in the presence of moduli fields, and find that inflation does not work with a generic KKLT moduli stabilisation potential. A viable model can be constructed with a fine-tuned moduli sector, but only for a very specific choice of K\\\"ahler potential. Our analysis also shows that inflation models satisfying $\\partial_{i} W_{\\rm inf}=0$ for all inflation sector fields $\\phi_i$ can be combined successfully with a fine-tuned moduli sector. \\vskip 0.1in \\noindent {\\bf Keywords:} inflation, cosmology of theories beyond the SM ",
        "introduction": "In chaotic inflation models the energy scale of inflation is high, typically of the order of the grand unified scale~\\cite{linde}. As a consequence  these models give a large tensor contribution to the density perturbations.  This makes them testable by current and future CMB experiments, most notably by the upcoming Planck satellite.  However, chaotic inflation is not easy to implement in a supergravity theory~\\cite{kawasaki,kadota}.  The inclusion of other high energy physics, such as moduli fields, creates further problems~\\cite{brax,kallosh,kallosh2,kss}. Naturally, any realistic inflation model must be part of some full theory, containing all known physics. The effects of other sectors of the theory on inflation can not be ignored.  As shown by Lyth~\\cite{lyth} in the context of slow-roll inflation, a measurable tensor mode requires the inflaton field to change by superplanckian values during inflation. % Examples of such ``large field models'' of inflation  are chaotic and natural inflation~\\cite{natural}.  At present no string theory derivation of a large field inflaton model exists. The displacement of the inflaton in brane models of inflation is bounded by the size of the compactified space, and in all known models less than the Planck scale~\\cite{baumann,bean}.  In all examples of modular inflation the inflationary scale is too low for an appreciable tensor signal~\\cite{kallosh,kallosh2}. N-flation~\\cite{kss,Nflation,grimm}, the stringy realization of assisted inflation~\\cite{assisted}, gives rise to appreciable tensor modes.   However, it is not clear whether all of the underlying assumptions are satisfied in these models \\cite{kallosh}. Despite the negative results, so far there is no  ``no-go'' theorem stating that string theory cannot give large field inflation.  It may very well be that it can be realized in corners of the landscape not yet explored -- after all, the search has only just begun.  In this paper we consider a ${\\mathcal N} = 1$ SUGRA implementation of chaotic inflation, and analyse what happens when it is combined with a KKLT-like moduli sector.  In our set-up the inflaton and moduli sector only interact gravitationally.   Our approach is phenomenological in that we analyse the SUGRA effective field theory, but do not attempt to derive the model from string theory.    It should be noted in this context that moduli fields are not unique to string theory.  Flat directions abound in any SUSY theory.  If SUSY is broken in some hidden sector by non-perturbative physics,  the moduli sector has the same qualitative properties as the KKLT model, and our results apply.  As mentioned,  it is not easy to construct a model of chaotic inflation  in the presence of additional moduli fields, even when they are stable.  First of all, there is the $\\eta$-problem, present in all models of SUGRA inflation~\\cite{copeland,dine}.  The potential during inflation is of the form $V \\sim \\e^K \\tilde V$, which for a canonically normalised inflaton field $\\vp$ gives rise to a large inflaton mass $m_\\vp^2 \\sim H^2$ ruining slow-roll inflation.  This can be solved by fine-tuning the K\\\"ahler potential so that the inflaton mass is accidentally small.  More elegantly, the inflaton mass can be protected by symmetries.  In this paper we will introduce a shift symmetry for the inflaton field that leaves the K\\\"ahler potential invariant to solve the $\\eta$-problem~\\cite{kadota,gaillard,banks}.  Inclusion of moduli fields in the system gives rise to a whole new set of obstacles to implement inflation.  The moduli fixing potential breaks supersymmetry.  Consequently there are soft corrections to the inflaton potential.  The soft terms are small in the limit of low scale SUSY breaking, with a small gravitino mass $m_{3/2}^2 \\lesssim H^2$.  At the same time, the requirement that the moduli fields remain stabilised in their minimum during inflation, and do not run away to infinity, implies that the moduli masses should be sufficiently large. This requirement is usually expressed as a constraint on the Hubble parameter during inflation $H^2 \\ll m_{\\rm mod}^2$~\\cite{KL}.  In a generic potential $m_{\\rm mod}^2 \\sim m_{3/2}^2$ and without fine-tuning (in addition to that required to set the cosmological constant to zero) these requirements are at odds with each other.  This is for example the case in the original KKLT model~\\cite{KKLT}.  It is difficult to embed large field inflation in such a set up.  We see there is tension between keeping the soft corrections to the inflaton potential small, and keeping the moduli fields fixed during inflation. This can be eased if the modulus sector is fine-tuned so that the modulus and gravitino masses are no longer of the same order of magnitude.  This is achieved explicitly in the Kallosh-Linde (KL) set-up~\\cite{KL}, which uses a racetrack potential for the modulus field.  In this case, parameters are tuned so that the modulus mass is much larger than the gravitino mass.  Having the Hubble constant during inflation  between these mass scales $m_{\\rm mod}^2 \\ll H^2 \\ll m_{3/2}^2$ offers a way to solve both problems. Note that it also allows the gravitino mass to be in the phenomenologically favoured TeV range, without the need for low scale inflation (in fact, this was the original motivation for KL).  In this paper we will analyse chaotic inflation in the presence of a single modulus field with a no-scale K\\\"ahler potential. The models we will study have the superpotential \\be W = W_{\\rm mod}(T) + m \\phi_1 \\phi_2 \\, . \\label{w1} \\ee We consider both a generic KKLT potential and a fine-tuned KL potential.   The above inflaton superpotential was first  proposed in~\\cite{kawasaki}. Refs.~\\cite{brax,kallosh,kallosh2} added a moduli sector to the set-up.  We extend their results by an in-depth discussion of the effects of the moduli dynamics, with an emphasis on finding the conditions for successful inflation. As expected, inflation does not work in the KKLT set-up.  Whether KL works depends sensitively on the K\\\"ahler potential for the inflaton fields. Although the moduli corrections are small after inflation due to the fine-tuning in the KL set-up, this is not necessarily true during inflation. During inflation the modulus field $T$ is slightly displaced from its post-inflationary minimum, disrupting the minute fine-tuning of the potential, with potentially large effects.   Indeed,  consider the following K\\\"ahler potentials \\numparts \\bea K_{1} &=& -3\\log[T+\\bar T] -\\frac12 (\\phi_1 -\\bar{\\phi}_1)^2 + \\phi_2 \\bar{\\phi}_2 \\, , \\label{k1}  \\\\ K_{2} &=& -3\\log[T+\\bar T] -\\frac12 (\\phi_1 -\\bar{\\phi}_1)^2 -\\frac12( \\phi_2- \\bar{\\phi}_2)^2 \\, , \\label{k2}  \\\\ K_{\\alpha} &=& -3\\log\\left[T+\\bar T- \\frac13 (T+\\bar T)^\\alpha \\, \\phi_2 \\bar{\\phi}_2\\right]-\\frac12 (\\phi_1 -\\bar{\\phi}_1)^2 \\, . \\label{kalpha} \\eea \\endnumparts All K\\\"ahler potentials have a shift symmetry for the inflaton field $\\phi_1$ to solve the $\\eta$-problem.  However, as we will show, only $K_{1}$ combined with the KL modulus sector gives a viable model.  For all the other models, independent of modular weight $\\alpha$, the coupling between the modulus and inflaton sectors leads to instabilities in the potential, with a runaway behaviour for some of the fields. It is thus crucial to take the dynamics of the modulus field during inflation into account for a correct analysis of the model.   This paper is organised as follows.  The next section provides the background material, with a concise summary of the KKLT and KL moduli stabilisation potential, as well as a discussion of SUGRA chaotic inflation without moduli.  The rest of the paper discusses the combination of chaotic inflation and moduli fields.  In section~\\ref{s:model2} we study the model with $K_{2}$~\\eref{k2}. Although inflation does not work, it is useful to analyse why.  In section~\\ref{s:model1} we consider the model with $K_{1}$~\\eref{k1}.  As mentioned above, this is a viable model of chaotic inflation.  We discuss the inflationary predictions, in particular whether the supergravity corrections can leave a signature in the CMB.  Finally, in section~\\ref{s:KL} we use the insight gained in the previous sections to discuss more generic combinations  of chaotic inflation and KL moduli stabilisation, including models with~\\eref{kalpha}.  We end with some concluding remarks.  Throughout this article we will work in units with the reduced Planck mass $\\mpl = 1/\\sqrt{8\\pi G_N}$ set to unity. ",
        "conclusions": "In this paper we studied SUGRA chaotic inflation in the presence of stabilised moduli fields.  To avoid the usual $\\eta$-problem a shift symmetry for the inflaton field is introduced.  But this is not enough, as the moduli stabilisation sector gives rise to additional contributions to $\\eta$ and $\\epsilon$ which are generically not small.  The moduli sector breaks supersymmetry, and as a result the inflaton fields get soft mass contributions of the order of the gravitino mass.  These corrections need to be small for successful inflation.  But in a generic moduli potential such as KKLT,  the modulus mass is of the same order as the gravitino mass, and it is impossible to keep the corrections to the inflaton small while making sure the modulus remains fixed in its minimum during inflation.  KL addressed this problem by constructing a fine-tuned moduli potential with $m_{3/2}^2 \\ll m_T^2$.  Indeed, calculating the potential in any model in which inflation is combined with a KL moduli sector by adding the respective superpotentials, the moduli corrections to inflation appear small while at the same time the modulus is heavy.  All of the above assumes that  the modulus $T$ is fixed during inflation. However, the modulus is a dynamical field, and this changes the situation drastically.  Although during inflation the modulus is only slightly displaced from its post-inflationary vacuum, this is enough to disrupt the minute fine-tuning of the KL model.  The corrections to the effective inflaton potential are generically large, and whether inflation works is a model dependent question.  Inflation combined with the KL moduli stabilisation scheme works well if the derivative of the inflaton superpotential during inflation vanishes $(W_{\\rm inf})_i =0$ with $i$ running over all inflaton sector fields. This is for example the case for $D$-term hybrid inflation.  On the other hand if $(W_{\\rm inf})_i \\neq 0$, there are large corrections to the masses of the inflaton sector fields, which are missed if the modulus dynamics are not kept. For models with $W_{\\rm inf}$ a polynomial in the shift symmetric inflaton field, these corrections are fatal.  If $W_{\\rm inf}$ is some polynomial of inflaton and ``spectator'' fields, the corrections to the $\\eta$-parameter can be harmlessly small if the spectator fields have a small VEV.  However, one must also check that the masses of the spectator fields are positive definite during inflation to avoid a run away behaviour.  For the chaotic inflation models under consideration this requires the spectator field $\\phi_2$ to have a minimal K\\\"ahler (but note that this model is only ``just'' stable).  It is not sufficient for $\\phi_2$ to appear inside the modulus log with unit modular weight, in which case upon a small field expansion  it will have a minimal K\\\"ahler.  In fact, no matter what the modular weight, if $\\phi_2$ is placed inside the log [see \\eref{Kall}] the spectator field becomes tachyonic during inflation.  Our route to a successful inflation model in this paper was to take a specific choice of K\\\"ahler potential that minimises the impact of the moduli corrections. We calculated the inflationary predictions for the viable model 1, which has a minimal kinetic term for the spectator field $\\vp_2$ \\eref{model1}, \\eref{combine}. Although the spectral index $n_s = 0.967$ is the same as for chaotic inflation with a quadratic potential,  the values of the slow-roll parameters differ from those of a purely quadratic potential.  The difference is largest for those parameters that stabilise $T$ at large values.  The degeneracy between the quadratic model and the model with moduli can be broken if tensor perturbations are observed, as this allows us to extract the values of $\\eta$ and $\\epsilon$ from the CMB data.  Hence, in the future, with the launch of the Planck satellite, we may be able to observe the presence of moduli fields in the sky.  Note that the problems arising from the variation of the modulus $T$ during inflation are not unique to chaotic inflation. Combining moduli with $F$-term hybrid inflation was recently discussed in~\\cite{dp}, where even a careful choice of K\\\"ahler could not save the model. Instead, taking inspiration from~\\cite{anaXW}, the moduli problems were reduced by multiplying the superpotentials of the two sectors, instead of adding them. It would be interesting to see if a similar approach can help chaotic inflation models, although we will leave this for future work.     \\ack SCD thanks the Netherlands Organisation for Scientific Research (NWO) for financial support.  \\appendix"
    },
    "0801/0801.1258_arXiv.txt": {
        "abstract": "{ It  is widely accepted  that the large  obliquity of Uranus  is the result of a great tangential collision (GC) with an Earth size proto-planet at the end of the accretion process.  The impulse imparted by the GC had affected the  Uranian  satellite  system.   Very recently,  nine  irregular  satellites (irregulars) have  been discovered around Uranus.  Their  orbital and physical properties, in particular those of  the irregular Prospero, set constraints on the  GC scenario. }{ We attempt to  set constraints  on the  GC scenario  as the  cause of Uranus' obliquity  as well  as on the  mechanisms able  to give origin  to the Uranian irregulars. }{ Different  capture mechanisms  for irregulars operate  at different stages on the giant planets  formation process. The mechanisms able to capture the uranian  irregulars before and after  the GC are  analysed.  Assuming that they were  captured before the  GC, we calculate  the orbital transfer  of the nine irregulars by the impulse imparted  by the GC.  If their orbital transfer results dynamically implausible, they should have originated after the GC.  We then investigate and discuss the  dissipative mechanisms able to operate later. }{ Very  few transfers exist for  five of the  irregulars, which makes their existence before  the GC hardly expected.  In  particular Prospero could not exist  at the time of  the GC.  Different capture  mechanisms for Prospero after  the GC  are investigated.   Gas drag  by Uranus'envelope  and pull-down capture  are  not  plausible   mechanisms.   Capture  of  Prospero  through  a collisionless interaction  seems to  be difficult.  The  GC itself  provides a mechanism of permanent capture.  However, the capture of Prospero by the GC is a low probable  event.  Catastrophic collisions could be  a possible mechanism for the birth  of Prospero and the other irregulars after  the GC. Orbital and physical clusterings should then be expected. }{ Either Prospero had to originate  after the GC or the GC did not occur.  In the former case, the mechanism for the origin of Prospero after the GC remains an  open question. An observing program able  to look for dynamical and physical  families is  mandatory.  In the  latter case, another  theory to account  for  Uranus' obliquity  and  the  formation  of the  Uranian  regular satellites on the equatorial plane of the planet would be needed. ",
        "introduction": "Very recently, rich systems  of irregular satellites (hereafter irregulars) of the giant  planets have been discovered.   Enabled by the  use of large-format digital  images  on  ground-based  telescopes,  new  observational  data  have increased the  known population  of Jovian irregulars  to 55 (Sheppard  et al. \\cite{shepparda}),   the  Saturnian   population   to  38   (Gladman  et   al. \\cite{gladmanb}, Sheppard et  al.  \\cite{sheppardb}, \\cite{shepparde}) and the Neptunian  population to  7 (Holman  et  al.  \\cite{holman},  Sheppard et  al. \\cite{sheppardd}).   The Uranian  system  is of  particular  interest since  a population  of  9  irregulars  (named  Caliban,  Sycorax,  Prospero,  Setebos, Stephano,  Trinculo, S/2001U2:  XXIV Ferdinand,  S/2001U3: XXII  Francisco and S/2003U3: XXIII  Margaret) has been  discovered around Uranus (Gladman  et al. \\cite{gladman}, \\cite{gladmana}, Kavelaars  et al.  \\cite{kavelaars}, Sheppard et al.  \\cite{sheppardc}).   The discovery of these objects  provides a unique window on  processes operating  in the young  Solar System. In  the particular case of  Uranus, their existence may  cast light on  the mechanism responsible for its  peculiar rotation axis  (Parisi \\& Brunini \\cite{parisi},  Brunini et al.  \\cite{brunini} (hereafter BP02)).   Irregulars of giant planets are characterized by eccentric, highly tilted with respect of  the parent planet equatorial  plane, and in  some case retrograde, orbits.  These objects cannot have  formed by circumplanetary accretion as the regular  satellites  but they  are  likely products  of  an  early capture  of primordial  objects from  heliocentric  orbits, probably  in association  with planet formation  itself (Jewitt \\& Sheppard \\cite{jewitta}).   It is possible for an object circling about the sun to be temporarily trapped by a planet. In terms of the classical three-body problem  this type of capture can occur when the object passes through the  interior Lagrangian point, $L_{2}$, with a very low relative  velocity.  But, without any  other mechanism, such  a capture is not permanent  and the objects will  eventually return to a  solar orbit after several or several hundred orbital  periods.  To turn a temporary capture into a  permanent one  requires a  source of  orbital energy  dissipation  and that particles could remain  inside the Hill sphere long enough  for the capture to be effective.  Although  currently  giant  planets  have  no efficient  mechanism  of  energy dissipation for permanent capture, at their formation epoch several mechanisms may have operated:  1) {\\em gas drag}  in the solar nebula or  in an extended, primordial  planetary atmosphere  or  in a  circumplanetary  disk (Pollack  et al.\\cite{pollack}, Cuk \\& Burns \\cite{cuk}), 2) {\\em pull-down capture} caused by the  mass growth and/or orbital  expansion of the planet  which expands its Hill    sphere    (Brunini     \\cite{bruninia},    Heppenheimer    \\&    Porco \\cite{heppenheimer}),  3) {\\em collisionless  interactions} between  a massive planetary satellite and guest bodies  (Tsui \\cite{tsui}) or between the planet and  a binary  object (Agnor  \\&  Hamilton \\cite{agnor}),  and 4)  collisional interaction between  two planetesimals  passing near the  planet or  between a planetesimal and a regular satellite.  This last mechanism, the so called {\\em break-up} process, leads to the formation of dynamical groupings (e.g. Colombo \\& Franklin  1971, Nesvorny  et al. \\cite{nesvornya}).   After a  break-up the resulting fragments of  each progenitor would form a  population of irregulars with similar surface composition,  i.e.  similar colors, and irregular shapes, i.e.   light-curves  of  wide  amplitude.   Significant  fluctuations  in  the light-curves of  Caliban (Maris et  al.  \\cite{maris}) and Prospero  (Maris et al.   \\cite{marisa}) and  the  time  dependence observed  in  the spectrum  of Sycorax (Romon  et al.  \\cite{romon}) suggest  the idea of  a break-up process for the origin of these bodies.    Several  theories to  account  for the  large  obliquity of  Uranus have  been proposed.   Kubo-Oka  \\& Nakazawa  (\\cite{kubo-oka}),  investigated the  tidal evolution of  satellite orbits and  examined the possibility that  the orbital decay of  a retrograde satellite leads  to the large obliquity  of Uranus, but the large mass required for  the hypothetical satellite makes this possibility very  implausible.    An  asymmetric  infall  or  torques   from  nearby  mass concentrations during the collapse of  the molecular cloud core leading to the formation of the  Solar System, could twist the  total angular momentum vector of the  planetary system.   This twist could  generate the obliquities  of the outer  planets (Tremaine  \\cite{tremaine}). This  model has  the disadvantages that the  outer planets must form before  the infall is complete  and that the conditions for the event that would  produce the twist are rather strict.  The model  itself is  difficult  to  be quantitatively  tested.   Tsiganis et  al. (\\cite{tsiganis}) proposed that the  current orbital architecture of the outer Solar System could have been  produced from an initially compact configuration with  Jupiter and  Saturn  crossing  the 2:1  orbital  resonance by  divergent migration.   The crossing  led to  close encounters  among the  giant planets, producing   large   orbital  eccentricities   and   inclinations  which   were subsequently  damped to the current value by gravitational interactions with planetesimals.   The  obliquity changes  due  to  the  change in  the  orbital inclinations.  Since the inclinations are damped by planetesimals interactions on  timescales much  shorter than  the timescales  for precession  due  to the torques from the Sun, especially for Uranus and Neptune, the obliquity returns to small values if it is small before the encounters (Hoi et al. \\cite{hoi}).  Large stochastic impacts at the  last stage of the planetary formation process have  been  proposed  as  the  possible cause  of  the  planetary  obliquities (e.g. Safronov  \\cite{safronov}).  The large obliquity  of Uranus (98$^\\circ$) is usually attributed to a  great tangential collision (GC) between the planet and an Earth-size  planetesimal occurred at the end of  the epoch of accretion (e.g., Parisi \\& Brunini \\cite{parisi}, Korycansky et al.  \\cite{korycansky}). The collision imparts  an impulse to Uranus and  allows preexisting satellites of the  planet to change  their orbits.  Irregulars  on orbits with  too large semimajor  axis  escape  from  the  system (Parisi  \\&  Brunini  1997),  while irregulars  with a  smaller semimajor  axis may  be pushed  to outer  or inner orbits  acquiring greater  or lower  eccentricities depending  on  the initial orbital elements, the geometry of the impact and the satellite position at the moment of impact.  The orbits  excited by this perturbation must be consistent with  the present  orbital  configuration of  the  Uranian irregulars  (BP02).   In an attempt to clarify the origin of Uranus obliquity and of its irregulars, we are using  in this study the most updated information  on their orbital and physical properties.  In Section  2, we improve  the model developed  in BP02 for the  five Uranian irregulars known at  that epoch and extend our study to  the new four Uranian irregulars recently  discovered by  Kavelaars et al.   (\\cite{kavelaars}) and Sheppard et al. (\\cite{sheppardc}).  The origin of these objects after the GC is  discussed  in Section  3,  where several  mechanisms  for  the origin  of Prospero are investigated.  The discussion of the results and the conclusions are presented in Section 4. ",
        "conclusions": "It is usually believed that the large obliquity of Uranus is the result of a great tangential collision (GC)  with an Earth-sized proto-planet at the end of the accretion process.   We have calculated the transfer of angular momentum and  impulse at impact  and have shown  that the GC  had strongly affected the orbits  of Uranian satellites.  We calculate  the transfer of the  orbits of  the nine  known Uranian  irregulars by  the GC.   Very few transfers exist  for five of  the nine irregulars, making  their existence before the GC hardly expected.  In particular, Prospero could not exist at the time of  the GC.  Then, either Prospero had to  originate after the GC or the  GC did  not occur, in  which case  another theory able  to explain Uranus'  obliquity and  the formation  of the  Uranian  regular satellites would be  needed.  It is usually  believed that the  regular satellites of Uranus have accreted from material  placed into orbit by the GC (Stevenson et al. \\cite{stevenson}).    Within the  GC scenario,  several possible mechanisms  for the  capture of Prospero after the GC were investigated.  If the Uranian irregulars belong to individual captures and relating the origin of the outer uranian system to a common formation process,  gas drag by Uranus' envelope and pull-down capture seem to be implausible.  Three-body gravitational encounters might be  a source of  permanent capture.   However, we  found that  the minimum permanent  orbital radius of  a guest  satellite of  Uranus is  $\\sim$ 955 $R_U$ while the current semiaxis of  Prospero is 645 $R_U$.  The GC itself could provide a mechanism of permanent capture and the capture of Prospero could have occurred from a heliocentric orbit as is required within the GC scenario, but due to  the low rate of incoming objects it  turns out to be difficult.  Break-up  processes could be  the mechanism for the  origin of Prospero and  the other  irregulars in the  frame of  different scenarios. Prospero might  be a fragment  of a primary  KBO fractured by  a collision with another KBO.  The fragment could  have been captured by Uranus if the two KBOs had a minimum orbital  eccentricity of 0.37.  Prospero could be a secondary  member of  a  collisional family  originated  by the  collision between another satellite  of Uranus and a KBO  where the parent satellite of Prospero could have been captured  by any mechanism before or after the GC.   This process  has  the disadvantage  that  it is  unlikely that  the preexisting satellite  were formed from a circumplanetary  disk as regular satellites  given the  large orbital  semiaxis required  for  this object. Since  collisional  scenarios require  in  general  high collision  rates, perhaps the irregulars were originally much more numerous than now.  Then, Prospero  and also  the other  irregulars might  be the  result  of mutual collisions  among  hypothetical preexisting  irregulars  (Nesvorny et  al. \\cite{nesvorny}, \\cite{nesvornyb})  which could have been  captured by any other mechanism before the GC.   The  knowledge of  the  size and  shape  distribution  of irregulars  is important to know their relation  to the precursor Kuiper Belt population.  It could bring  valuable clues to  investigate if they are  collisional fragments from break-up processes occuring at the Kuiper Belt and thus has nothing to do with how they  were individually captured later by the planet,  or if they are collisional fragments produced during or  after the capture event (Nesvorny et al.  \\cite{nesvorny},  \\cite{nesvornyb}).  The differential  size distribution of the  Uranian irregulars approximates a  power law with  an exponent $q$=1.8 (Sheppard et al.\\cite{sheppardc}).  If we assume that the size distribution of the nine  irregulars with  radii greater than  7 km  extends down to  radii of about  1 km,  we  would expect  about 75  irregulars  of this  size or  larger (Sheppard et al.\\cite{sheppardc}).  The nuclei  of Jupiter  family comets are  widely considered to  be kilometer- sized fragments produced collisionally in  the Kuiper Belt (Farinella \\& Davis \\cite{farinellaa}).   Jewitt  et  al.   (\\cite{jewittb})  compared  the  shape distribution  of  cometary  nuclei  in  the  Jupiter  family  with  the  shape distribution of  small main-belt asteroids of  similar size (1 km  -10 km) and with  the  shape  distribution  of  fragments produced  in  laboratory  impact experiments.   They  found that  while  the  asteroids  and laboratory  impact fragments show  similar distribution of axis ratio  ($\\mean{b/a}$ $\\sim$ 0.7), cometary nuclei  are more elongated  ($\\mean{b/a}$ $\\sim$ 0.6).   They predict that if comets reflect their collisional origin in the Kuiper Belt followed by sublimation-driven  mass loss  once inside  the orbit  of Jupiter,  small KBOs should  have average shapes  consistent with  those of  collisionally produced fragments (i.e.,$\\mean{b/a}$  $\\sim$ 0.7). To date, constraints  on the shapes of only  the largest  KBOs are  available.  Prospero being  a bit  larger than cometary nuclei,  displays a variability of 0.21  mag in the R  band (Maris et al.  \\cite{marisa}).   This corresponds  to an axis  ratio projected  into the plane  of  the  sky,  b/a  of  0.8.   The knowledge  of  the  size  and  shape distribution of irregulars would shed light in the size and shape distribution of small KBOs as well as on the irregulars capture mechanism.   Colors are an  important diagnostic tool in attempting  to unveil the physical status and  the origin of the  Uranian irregulars.  In particular  it would be interesting  to  assess  whether  it  is  possible  to  define  subclasses  of irregulars just looking  at colors, and comparing colors  of these bodies with colors of minor bodies in  the outer Solar System.  Avalilable literature data show a dispersion  in the published values larger than  quoted errors for each Uranian irregular  (Maris et al.   \\cite{marisa}, and references  therein). We have concluded in Maris et al.  (\\cite{marisa}), that the Uranian irregulars are slightly red but they are not as red as the reddest KBOs.  An  intensive  search  for  fainter  irregulars  and a  long  term  program  of observations able to recover in  a self consistent manner light-curves, colors and phase effects informations is mandatory."
    },
    "0801/0801.0859_arXiv.txt": {
        "abstract": "In this paper, we investigate the structures of galaxies which either have or have had three BHs using $N$-body simulations, and compare them with those of galaxies with binary BHs. We found that the cusp region of a galaxy which have (or had) triple BHs is significantly larger and less dense than that of a galaxy with binary BHs of the same mass. Moreover, the size of the cusp region depends strongly on the evolution history of triple BHs, while in the case of binary BHs, the size of the cusp is determined by the mass of the BHs.  In  galaxies which have (or had) three BHs, there is a region with significant radial velocity anisotropy, while such a region is not observed in galaxies with binary BH. These differences come from the fact that with triple BHs the energy deposit to the central region of the galaxy can be much larger due to multiple binary-single BH scatterings. Our result suggests that we can discriminate between galaxies which experienced triple BH interactions with those which did not, through the observable signatures such as the cusp size and velocity anisotropy. ",
        "introduction": "According to recent observations, elliptical galaxies can be classified into two groups: ``weak-cusp'' galaxies and ``strong-cusp'' galaxies \\citep{lauer95, faber97}.  The central surface brightness profiles of the weak-cusp galaxies are expressed as $\\Sigma(R) \\propto R^{-\\gamma}$ with $\\gamma \\le 0.3$, and those of the strong-cusp galaxy the same formula with $\\gamma \\ge 0.5$. The slope of the volume density profile of strong-cusp galaxies is around $-2$, being consistent with the isothermal cusp. Such a steep cusp is naturally formed in dissipative process involving gas dynamics and star formation. On the other hand, the weak cusp corresponds to the slope of the volume density shallower than $-1$, which is not likely to be formed through dissipative process.  One possible way to form a weak cusp is the merging of two galaxies containing black holes (BHs). When two galaxies, each with a black hole, merge, two BHs sink toward the center of the merger remnant  to form a binary through the dynamical friction \\citep{bbr80}.  The back reaction of the dynamical friction heats up field stars. As a result, a shallow cusp of stars develops in the central region \\citep{Ebisuzakietal1991,NM99,merritt06}.  One problem with this binary BH scenario is what would be the final fate of the binary BH. \\citet{bbr80} pointed out that the merging timescale of the binary BH might be very long, after the binary BH ejected out the stars which can interact with the binary (loss cone depletion). The stars will be supplied in the timescale of the relaxation time, which is much longer than the Hubble time. Recent $N$-body simulations confirmed this theoretical estimate \\citep{MF04, berczik2005}.  If galaxies are formed through hierarchical clusterings, in many cases, a binary BH is formed after a merger event. If there were sufficient gas left in the merger, interaction with gas might lead to the quick merging of the two BHs. However, in the case of ``dry'' mergers which would result in the formation of giant ellipticals, by definition not much gas is left and it would be difficult for two BHs to merge.  If one galaxy with binary BH and the other with single BH merge, the central BHs form a triple system.  Iwasawa, Funato \\& Makino (2006, hereafter referred to as Paper I) investigated the evolution of triple BH system in the galactic center, using $N$-body simulations.  They found that the strong binary-single BH interaction \\citep{ME94} and the Kozai cycle \\citep{Kozai62, BLS} drives the eccentricity of the BH binary high enough that two BHs merge quickly through gravitational wave radiation. \\cite{HL07} performed statistical simulations of evolution of central BH systems and reached a similar conclusion.  This paper is a follow-up of Paper I. In this paper, we investigate the structure of a galaxy containing three BHs. We also investigated their observational properties which would help us to find galaxies which have (or had) triple BHs.  We performed $N$-body simulations of the evolution of triple (or binary) BH in a host galaxy, in order to study the dynamical evolution of the structure of stellar systems containing BHs.  In our simulations, both dynamical evolution of BHs and that of field stars are integrated consistently.  The interaction between BHs affects not only the evolution of themselves but also the spatial and kinematic structure of field stars around them.  In turn, the distribution of field stars  affects the interaction between BHs. To understand the structure of galaxies containing BHs, a  self-consistent simulation in which the orbits of BHs and stars are treated self-consistently is essential.  The structure of this paper is as follows.  In section 2, we describe the initial models and the method of our numerical simulations.  In section 3, we show the effect of the BH triples dynamics on the structure of the galaxy. Summary and discussion are given in section 4. ",
        "conclusions": "\\subsection{Possibility to Find Triple BH Systems}  In previous sections, we have seen that the density and velocity structure of galaxies with triple BHs is different from that with two BHs.  In Paper I, we showed that two of three BHs in a galaxy merge within several dynamical times.  The standard hierarchical clustering scenario of galaxy formation suggests that galaxies frequently merge and form ellipticals.  Many elliptical galaxies, therefore, might host a binary BH. It is likely that many of them had a triple BH system once upon a time. Furthermore, there may be galaxies with a triple BH system. The lifetime of a triple BH system in a galaxy is several times the dynamical time of the core of the galaxy, which is  about $10^8$ years ($100$ M years). If we assume every elliptical galaxy have binary BH, and all large ellipticals are formed through mergings of elliptical galaxy. Speaking, 1\\% of large ellipticals could have triple BH systems now.  One possible way to discriminate between a galaxy with binary or single BH and that have (or had) triple BH is the measurement of the cusp radius and density.  Our results suggest that the central cusp of a galaxy with three BHs is larger by a factor of few than that with binary BH.  Thus the galaxy with a large cusp and low density in the central region might have (or had) three BHs.  Another possible way is to use the difference of velocity anisotropy. Anisotropy parameter of galaxy with three BHs is positive (radial) in the outer region due to the heating by BHs.  The anisotropy parameter $\\beta$ is estimated by the line of sight velocity and their higher moment, assuming that the galaxy is spherical. However, with this method it is difficult to obtain local anisotropy parameter in the central region of galaxies, unless the central density cusp is steep \\citep{Gerhard93}. Another method is to fit the model with observational data \\citep{Cretton00,Gebhardt00,Gebhardt03}. Figure 10 in \\citet{Gebhardt03} shows anisotropy profile of galaxies. The shapes of some galaxies profile in their paper, for example all ``weak-cusp'' galaxies (NGC3608, NGC4291, NGC4649), are similar to of our result in Figure \\ref{fig9}. These galaxies might contain or have contained triple BHs.  \\subsection{Conclusions}  In this paper, we investigated the effect of two- or three-BH systems on the structure of galaxies. We found that  if the galaxy contains three BHs, (1) multiple three body scattering events reduce the density of the central cusp and increase the cusp radius and (2) in outer region, orbits of stars are likely to be radial. These difference allow us to discriminate the galaxies with two or three BHs."
    },
    "0801/0801.2547_arXiv.txt": {
        "abstract": "% {} {We investigate the present-day photometric properties of the dwarf spheroidal galaxies in the Local Group. From the analysis of their integrated colours, we consider a possible link between dwarf spheroidals and giant ellipticals. From the analysis of the $M_{V}$ vs (B-V) plot, we  search for a possible evolutionary link between dwarf spheroidal galaxies (dSphs) and dwarf irregular galaxies (dIrrs). } {By means of chemical evolution models combined with a spectro-photometric model, we study the evolution of six Local Group dwarf spheroidal galaxies (Carina, Draco, Sagittarius, Sculptor, Sextans and Ursa Minor). The chemical evolution models, which adopt up-to-date nucleosynthesis from low and intermediate mass stars as well as nucleosynthesis and energetic feedback from supernovae type Ia and II, reproduce several observational constraints of these galaxies, such as abundance ratios versus metallicity and the metallicity distributions. The proposed scenario for the evolution of these galaxies is characterised by low star formation rates and high galactic wind efficiencies.}{ Such a scenario allows us to predict integrated colours and magnitudes which agree with observations. Our results strongly suggest that the first few Gyrs of evolution, when the star formation is most active, are crucial to define the luminosities, colours, and other photometric properties as observed today. After the star formation epoch, the galactic wind sweeps away a large fraction of the gas of each galaxy, which then evolves passively. \\\\ Our results indicate that it is likely that at a certain stage of  their evolution, dSphs and dIrrs presented similar photometric properties. However, after that phase, they evolved along different paths, leading them to their currently disparate properties. }{} ",
        "introduction": "The dwarf spheroidal galaxies (dSphs) are among the most common types of galaxies in the universe. They are found normally in groups (C\\^ot\\'e et al. 1997) and clusters (Phillips et al. 1998, Ferguson \\& Sandage 1991). The ones found in the Local Group have become an increasingly important matter of study in the last few years due to their proximity, which enables one to study in detail objects and processes which were formally restricted to our own Galaxy.  Several issues regarding the formation and evolution of the Local Group dSphs can help in the attempt to clarify the whole subject of galaxy formation. For example, how did the dSphs form? When? What mechanism ruled their evolution? Are these galaxies remnants of the building blocks from which larger galaxies assembled? Is their evolution mainly affected by the environment or do internal processes play a major role? Are these galaxies linked to other types of dwarf galaxies in the context of any evolutionary scenario? In order to answer these questions, one should try to understand not only the present day properties of the dSphs, but also try to understand their past evolutionary history. One possible procedure is to make use of models which, being based on present day observational constraints, allow one to trace the past evolution of the dSphs.  Originally, the local dSphs were believed to be very old simple systems, similar to globular clusters (Shapley 1938). However, more recent deep photometric observations of main sequence turn-off stars revealed also intermediate-age populations, as well as a significant metallicity range, different from most globular clusters. In fact, analysis of colour-magnitude diagrams (CMDs) suggested that these galaxies are characterised by complex and different star formation (SF) histories (van den Bergh 1994; Hernandez, Gilmore, Valls-Gabaud 2000; Dolphin et al. 2005). In almost all cases, the mechanisms which trigger and control the SF are yet unknown and several scenarios have been proposed. These scenarios should also explain the complete lack of gas in these galaxies, the low metallicities, the low values of [$\\alpha$/Fe] relative to Galactic stars with the same [Fe/H] (Bonifacio et al. 2000; Shetrone, C\\^ot\\'e, Sargent 2001; Shetrone et al. 2003; Tolstoy et al. 2003, Bonifacio et al. 2004;  Venn et al. 2004; Sadakane et al. 2004; Geisler et al. 2005; Monaco et al. 2005) and the metallicity distributions, including the large metallicity range (Koch et al. 2005, Bellazzini et al. 2002). \\\\ Several attempts to model the properties of dSphs have been performed, generally following different approaches. By means of high-resolution cosmological numerical simulations, Ricotti \\& Gnedin (2005) and Kawata et al. (2006) studied dSphs in a cosmological framework. Their studies are useful to understand the link between the dSphs and the first galaxies, but cannot investigate in detail the chemical properties of dSphs and the relative roles of SNe Ia and II in the chemical enrichment and gas ejection processes. Marcolini et al. (2006), by means of a chemo-dynamical evolution model, studied the evolution of the interstellar medium (ISM) of the Draco dSph. In their picture no galactic wind develops and, in order to deplete the galaxy of its gas and to stop star formation, one must invoke an external mechanism, such as ram pressure stripping or tidal interactions. Fenner et al. (2006) and Ikuta $\\&$ Arimoto (2002) studied the properties of dSphs by means of chemical evolution models including galactic winds. Their results indicate that a small fraction of the ISM is carried away by the SN-driven winds and confirm the need of external gas removing processes to reproduce the present-day gas fractions of dSphs (see also Mori \\& Burkert 2000, Mayer et al. 2006). Lanfranchi $\\&$ Matteucci (2003, 2004, hereinafter LM03, LM04), alternatively, by means of a detailed chemical evolution model with galactic winds, were able to reproduce many observational constraints of six local dSph galaxies (Carina, Draco, Sagittarius, Sculptor, Sextans, Ursa Minor). The scenario proposed by these authors considered low efficiency star formation rates (SFR) derived from colour-magnitude diagrams together with intense galactic winds. \\\\ LM03,04 were able to reproduce not only the chemical properties, but also the lack of central neutral gas and the metallicity  distributions of the studied galaxies. In particular, the  [$\\alpha$/Fe]  vs [Fe/H] and neutron capture element ratios in these galaxies are well reproduced as well as the stellar metallicity distribution (SMD).\\\\ The photometric properties of the local dSphs, however, are rarely addressed in any of these models, even though they could help, not only in constraining the formation and evolution of these galaxies, but also in clarifying the subject of a possible evolutionary connection between the gas poor dSph galaxies and gas rich dwarf irregular galaxies (dIrrs). In such a scenario, a starburst in a dIrr gives rise to a super wind which removes all the gas of the galaxy (which could be removed also by ram pressure stripping or tidal stripping) and halts the SF, giving rise to a dSph galaxy (Lin $\\&$ Faber 1983; Dekel $\\&$ Silk 1986;  van den Bergh 1994; Papaderos et al. 1996; Davies $\\&$ Phillipps 1988). There are, however, several drawbacks in that scenario, from both the chemical and photometric point of view. First, the large scale distribution of dSphs is substantially different from that  of dIrrs. Most of the dSphs cluster around the two giant spirals of the Local Group (Grebel 1998), with very few exceptions. On the other hand, most of the dIrrs lie at large ($>500 kpc$) distances from the large galaxies (Mateo 1998). This is probably  linked to the morphology-density relation for dwarf galaxies and is interpreted as an evidence of environmental effects on galaxy evolution (Grebel 2001). In particular, it may be possible that their proximity to large galaxies had some effects on the evolution of dSphs.  \\\\ Futhermore, the dIrrs and dSphs show rather different observational properties. The integrated colour-magnitude diagram of the Local Group dwarf galaxies exhibits a clear distinction between the dSphs and the dIrrs, which occupy different regions (Mateo 1998). A few galaxies, namely transition type dwarfs, exhibit  intermediate photometric properies between those of dSphs and dIrrs. Also the luminosity/metallicity relation can be used to distinguish between dIrrs and dSphs at the present epoch (Mateo 1998, Grebel et al. 2003). At the same luminosity, the dIrrs tend to exhibit lower metallicities than the dSph galaxies. These two facts could be related to different evolutionary histories for the two types of dwarfs and only a study of their past history could help in clarifying this issue.  In this work we show that a spectro-photometric code coupled with a chemical evolution model allows us to investigate the past evolution of the dSphs and impose constraints on the formation and evolution of these galaxies and also to the possible connection with other types of dwarf and giant galaxies. A chemical evolution code which reproduces successfully several observational constraints can provide the parameters required to predict the evolution of the photometric properties of a sample of six local dSph galaxies.   The paper is organized as follows: in Sect. 2, the chemical evolution models which reproduce the chemical data of these galaxies and their results are described. In Sect. 3 we describe the spectro-photometric code. The results of  our models compared to observational data are shown   in Sect. 4 and finally in Sect. 5 we draw some conclusions. ",
        "conclusions": "By means of a chemical evolution model combined with a spectro-photometric code we were able to predict the evolution of several photometric properties of six Local Group Dwarf Spheroidal galaxies (Carina, Draco, Sagittarius, Sculptor, Sextans and Ursa Minor). The chemical evolution models adopt up-to-date nucleosynthesis from intermediate mass stars, massive stars and SNe type II and Ia and take into account the contribution of SNe to the energetics of the ISM. For the six dSphs, the star formation histories are taken from the  observed colour-magnitude diagrams (CMDs, Dolphin et al. 2005). The proposed scenario for the evolution of these galaxies is characterised by low star formation rates and high galactic wind efficiencies. Such a scenario allows us to predict colours and magnitudes in agreement with observations.  The main conclusions can be summarized as follows:  \\begin{itemize}  \\item the total luminosities of 5 (Sculptor, Carina, Sextans, Ursa Minor and Draco) out of 6 dSphs galaxies analysed here are dominated by the SSPs formed at the lowest metallicities, i.e. $Z\\le 0.0004$. This is a consequence of the low SF efficiency and of the intense wind which removes a large fraction of the gas of the galaxy, almost halting the SF and preventing metal-rich stars to be formed. For each galaxy, the stellar populations dominating the metallicity distribution (LM04) dominate also the total light.  \\item in the case of Sagittarius, the B and K luminosities are dominated by stellar populations with higher metallicities, i.e. $Z\\ge 0.004$,  due to the much higher SF efficiency adopted for this galaxy ($>$ 10 times higher than the ones of the other dSphs);  \\item only during the SF phase, when all galaxies exhibit a high gas content, the effects of dust extinction in the total luminosities in the U, V, B bands are noticeable. After the onset of the galactic winds and the consequent removal of the interstellar gas, the effects of extinction become negligible;   \\item We compared the predicted current integrated (U-V) and (V-K) colours as a function of the absolute V magnitude for the dSphs to the colour-magnitude relation for giant spheroids in clusters. The main differences between dSphs and the giant ellipticals concern their star formation histories. In particular, the formation of dSphs occurred with a lower degree of synchronicity than the formation of giant ellipticals. The differences may also be due to environmental effects, since clusters are environments denser than the Local Group. In the colour-magnitude plots, the six local dSphs  studied here seem to form a sequence similar to the one observed for their counterparts in clusters, although with a slightly steeper slope and a larger dispersion in their ages. Although we can not draw any firm conclusion on this aspect, it is not unlikely that the dSphs can be regarded as the lowest-mass tails of their larger counterparts and not the building blocks from which they were assembled. This issue represents a challenge to all galaxy formation models based on the popular $\\Lambda$ cold dark matter cosmological scenario.  \\item The study of the evolution of the integrated V magnitude as a function of the (B-V) colour suggests that, during their past history, the dSphs may have shown photometric properties similar to the ones of the present-day dIrrs. It is possible that dSphs and dIrrs shared  a common progenitor phase, and then evolved through different paths. After the common phase, dIrrs have experienced little evolution, mantaining most of their gas and blue colours. dSphs have instead experienced a stronger evolution, losing all of their gas and showing present-time colours redder than dIrrs. Transition types might represent a sub-class of dIrrs, which may evolve into dSphs.    \\end{itemize}"
    },
    "0801/0801.2771_arXiv.txt": {
        "abstract": "Void regions of the Universe offer a special environment for studying cosmology and galaxy formation, which may expose weaknesses in our understanding of these phenomena.  Although galaxies in voids are observed to be predominately gas rich, star forming and blue, a sub-population of bright red void galaxies can also be found, whose star formation was shut down long ago.  Are the same processes that quench star formation in denser regions of the Universe also at work in voids?  We compare the luminosity function of void galaxies in the 2dF Galaxy Redshift Survey, to those from a galaxy formation model built on the Millennium Simulation.  We show that a global star formation suppression mechanism in the form of low luminosity ``radio mode'' AGN heating is sufficient to reproduce the observed population of void early-types. Radio mode heating is environment independent other than its dependence on dark matter halo mass, where, above a critical mass threshold of approximately $M_{\\rm vir}\\!\\sim\\!10^{12.5} M_\\odot$, gas cooling onto the galaxy is suppressed and star formation subsequently fades.  In the Millennium Simulation, the void halo mass function is shifted with respect to denser environments, but still maintains a high mass tail above this critical threshold.  In such void halos, radio mode heating remains efficient and red galaxies are found; collectively these galaxies match the observed space density without any modification to the model.  Consequently, galaxies living in vastly different large-scale environments but hosted by halos of similar mass are predicted to have similar properties, consistent with observations. ",
        "introduction": "how does one understand early-type galaxies in voids when void environments are typically very gas rich and slowly evolving, which tends to promote star formation rather than suppress it.  At this point it may be valuable to backtrack somewhat and revisit the two main mechanisms acting in the model that can turn blue galaxies red. The first is important for satellite (i.e. non-central) galaxies.  Upon infall into a more massive system, any extended hot gas around the (now) satellite is stripped by the denser medium and added to the more massive halo. Such ``strangulation'' drives a satellite galaxy to redden rapidly once its remaining cold disk gas is exhausted.  The second is the radio mode heating discussed above and in Section~\\ref{sec:model}. Radio mode AGN operate only in central galaxies where the host halo has grown above a critical mass threshold, approximately $M_{\\rm vir} \\sim 10^{12.5} M_\\odot$.  So what physical processes have occurred to produce in the observed 2dFGRS early-type void population?  It may be that such red galaxies are simply a population of satellites and have survived long enough for an effect on their colours to be seen.  To test this idea using the model we perform a simple and revealing exercise.  From our knowledge of which galaxies are central and which are satellites, we remove those that are satellites and hence those that could have experienced the strangulation effect.  This is shown by the dotted lines in Figure~\\ref{fig:LFcol}.  Although satellites do contribute to the void early-type luminosity function at $M_{\\rm b_J}\\!-\\!5\\log_{10}\\! h \\simgt -18.5$, they are clearly a minor component brighter than this.  Indeed, all bright early-type void galaxies are centrals in our galaxy formation model.  The second of our star formation shut-down mechanisms, i.e. radio mode low luminosity AGN heating, requires that central galaxies reside in sufficiently large dark mater halos.  Such halos are not expected to be common in void regions of the Universe where low mass halos dominate.  In low mass halos the central density of hot gas is not high enough to maintain the low Eddington accretion needed to power a central AGN outflow.  We check this explicitly in Figure~\\ref{fig:MF}, where we plot the halo mass function for halos in all environments (left panel) and void environments (right panel) (wide-dashed lines in both).  In addition, we break the halo mass functions into those halos hosting red central galaxies (solid lines) and those hosting blue central galaxies (dashed lines).  Figure~\\ref{fig:MF} reveals the origin of the void early-type population -- they are simply central galaxies that live in halos massive enough to have the low luminosity radio mode operating.  In other words, the physics assumed by the model that transforms blue galaxies into red operates independent of environment but can still produce the observed environmental trends in voids (note that the transition from blue-dominated to red-dominated dark matter halos begins at approximately the same mass in both panels, $M_{\\rm vir} \\sim 10^{12} M_\\odot$).  Specifically, \\emph{it is the shift in halo mass function with environment that changes, not the transformation mechanism itself.}  This picture has important consequences for galaxy evolution more generally.  First, it indicates the necessity for active galaxies in voids, confirmed in a number of studies (see, e.g., \\citealt{Constantin2008} at $z\\sim0$, and \\citealt{MonteroDorta2008} at $z\\sim1$).  Second, it requires that void galaxies should have similar properties to those in all other environments, assuming the galaxies compared are hosted by halos of similar mass.  This is a theoretical confirmation of numerous observational studies that have shown that environments on scales larger than a few Mpc (i.e. outside the dark matter halo) are not important for galaxy formation \\citep[e.g. see][and references therein]{Blanton2007}. For example, \\cite{Patiri2006} investigate the properties of void galaxies in the Sloan Digital Sky Survey and compare to those of a similarly constructed mock catalogue using the same \\cite{Croton2006} model used here.  They found little difference in the mean properties of void galaxies relative to the field, specifically in colour, specific star formation rate, and morphology.  Theoretically one may expect large-scale environment to be important, e.g. the effect of assembly bias for dark matter halos \\citep{Gao2005}.  However the galaxy population does not appear to be overly sensitive to this \\citep[see, for example,][]{Croton2007}.  A similar conclusion was recently arrived at in a complimentary analysis by \\cite{Tinker2008}.     It has been suggested that void environments pose a problem for galaxy formation theory in a $\\Lambda$CDM universe due to the apparent over-abundance of dark matter halos in voids relative to the observed abundance of galaxies \\citep{Peebles2001}.  In this paper we show that no conflict exists within the current observational uncertainty.  \\begin{itemize} \\item Our model can reproduce the observed abundance of void galaxies, both globally and by colour, without any modification or additional environment dependent physics. \\item Early-type ``red and dead'' red sequence galaxies appear naturally in the voids.  They arises because of a shift in the halo mass function in low density environments combined with an environment independent star formation shut-down mechanism efficient above a critical halo mass (here, radio mode AGN).  Together, these approximately produce the correct observed abundance. \\item Some notable consequences follow from our results.  For example, at a given host halo mass, void galaxies are expected to have similar properties on average to those in the field, because such galaxies will have had similar evolutionary histories to field and cluster galaxies. \\end{itemize}  Voids and void galaxies provide an important probe of both cosmology and galaxy formation.  Future wide-field surveys will produce large-scale maps of the Universe out to high redshift, from which void evolution can be studied in detail.  Such work will further constrain galaxy formation theory, both in under-dense and the wider field environment.",
        "conclusions": "\\label{sec:results}  \\subsection{The void luminosity function} \\label{subsec:LF}  \\begin{figure*} \\plotscaled{./figures/figure3_new.ps} \\caption{The Millennium Simulation halo mass function (left panel) and halo mass function in voids (right panel).  In both panels the halo mass function is broken up into those who host red central galaxies (solid lines) and those that host blue central galaxies (dashed lines).  Red sequence galaxies occupy the most massive halos in all environments -- these halos are subject to the radio mode low luminosity AGN heating that ultimately shuts down star formation. Notably, this shutdown begins at approximately the same mass in both panels, $M_{\\rm vir} \\sim 10^{12-12.5} M_\\odot$.  This implies that the (environment independent) radio mode heating, plus a shift in the halo mass function with environment, is sufficient to reproduce the observed abundance of void early-type galaxies seen in Figure~\\ref{fig:LFcol}.} \\label{fig:MF} \\end{figure*}   We begin with Figure~\\ref{fig:LFall}, which shows the luminosity function of model galaxies for both the population as a whole, and all galaxies in void regions.  Here we define a void as \\cite{Croton2005} did, $\\delta_8\\!<\\!-0.75$, i.e. where the density within an \\eightMP\\ radius of the galaxy is less than $25\\%$ the mean.  Over-plotted are the observational results for the full 2dFGRS \\citep[filled circles,][]{Norberg2002} and void galaxy luminosity function \\citep[open squares,][]{Croton2005}.  As was the focus of \\cite{Croton2006}, the global luminosity distribution of model galaxies is a good match to the local observations.  Important model ingredients that produced this result are the supernova in shallower potentials that expel disk gas to reduce star formation and flatten the faint-end slope, and the inclusion of an AGN heating source in large halos to starve massive central galaxies of star forming fuel, resulting in the observed bright-end exponential cut-off, as described in Section~\\ref{sec:model}.  Without these two critical aspects in the model, the luminosity function instead would have a power-law shape reflecting the underlying mass function of halos \\citep{Benson2003}.  Note that there remain discrepancies in the top curve of Figure~\\ref{fig:LFall}, notably that the model over-predicts the abundance of very bright galaxies.  As discussed in \\cite{Croton2006}, these galaxies represent a population undergoing strong star formation and merger-induced starbursts.  In such ultraluminous infrared galaxies (ULIRGs), nearly all the light from young stars is absorbed by dust and re-radiated in the mid- to far-infrared \\citep{Sanders1996}.  Improved modelling of the effects of dust are required to adequately reproduce the properties of such systems.  The lower solid line in Figure~\\ref{fig:LFall} shows the model result for the void galaxy population in the Millennium Simulation.  The good overall agreement is a result of the physical prescriptions and global parameter choices in the \\cite{Croton2006} model and not additional fine-tuning (of which there was none).  There is some over-prediction of the very brightest void galaxies.  However the discrepancy is not significant enough to expect that improvements to existing aspects of the model, e.g. dust as described above, cannot alleviate such differences.  Additionally, we find a small over-prediction of faint ($M_{\\rm b_J}\\!-\\!5\\log_{10}\\! h \\simgt -18.5$) galaxies. Unfortunately, systematic variations in the observed faint-end galaxy luminosity function exist at a level greater than this (e.g. contrast \\citealt{Cole2001} to \\citealt{Huang2003}).  Figure~\\ref{fig:LFall} alone answers the challenge posed by Peebles and outlined in the Introduction.  Specifically, in a $\\Lambda$CDM Universe the distribution of dark matter halos in voids, coupled with a realistic model of galaxy formation, is consistent with voids as observed in the real Universe.  Such regions are typically not devoid of galaxies as claimed in \\cite{Peebles2001}.  The remainder of this letter will aim to dissect and understand the good agreement shown in Figure~\\ref{fig:LFall} between model and observation.   \\subsection{Galaxy colours in voids} \\label{subsec:col}  The colour of a galaxy tells us a lot about the relevant physics that has been dominant during its evolution.  A red spectrum typically indicates that a star formation shutdown mechanism has been operating for a significant part of the galaxy's lifetime.  We can use this knowledge, in conjunction with our theoretical model, to gain insight into how shutdown may occur and to what degree it may (or may not) be environment specific.  To this end, in Figure~\\ref{fig:LFcol} we again plot the void galaxy luminosity function for both the 2dFGRS \\cite[symbols,][]{Croton2005} and semi-analytic model (lines), but now broken up by galaxy spectral type (early/late) or colour (red/blue).  Early-type (triangles) and late-type (squares) 2dFGRS galaxies are determined using the principal component analysis of \\cite{Madgwick2003}.  For the model we use the bi-modal galaxy colours to separate red (solid line) from blue (dashed line) at $m_{b_{\\rm J}}\\!-\\!m_{r_{\\rm F}}\\!=\\!1.07$ \\citep[see figure~9 of][]{Croton2006}.  Comparing model to observation, Figure~\\ref{fig:LFcol} shows agreement (to within $2\\sigma$) between the two sets of early-type/red and late-type/blue void luminosity functions.  This has occurred as a byproduct of matching the model to the observed \\emph{global} properties and only these -- no special void environment physics was needed.  Adding detail to the physical prescriptions assumed by the model would presumably improve the agreement further, but this is not our focus here.  The early-type 2dFGRS void galaxies in Figure~\\ref{fig:LFcol} return"
    },
    "0801/0801.4294_arXiv.txt": {
        "abstract": "We describe the results of a radio continuum survey of the central 4\\sdeg$\\times$1\\sdeg\\ with the 100 m Green Bank Telescope (GBT) at wavelengths of 3.5, 6, 20, and 90 cm.  The 3.5 and 6 cm surveys are the most sensitive and highest resolution single dish surveys made of the central degrees of our Galaxy.  We present catalogs of compact and extended sources in the central four degrees of our Galaxy, including detailed spectral index studies of all sources.  The analysis covers star-forming regions such as Sgr B and Sgr C where we find evidence of a mixture of thermal and nonthermal emission.  The analysis quantifies the relative contribution of thermal and nonthermal processes to the radio continuum flux density toward the GC region.  In the central 4\\sdeg$\\times$1\\sdeg\\ of the GC, the thermal and nonthermal flux fractions for all compact and diffuse sources are 28\\%/72\\% at 3.5 cm and 19\\%/81\\% at 6 cm.  The total flux densities from these sources are $783\\pm52$ Jy and $1063\\pm93$ Jy at 3.5 and 6 cm, respectively, excluding the contribution of Galactic synchrotron emission. ",
        "introduction": "The central few hundred parsecs of the Milky Way comprise a region in the Galaxy unique for its high stellar density, intense ionizing radiation field, massive black hole, enhanced density of cosmic rays, and unusual magnetized structures \\citep{f04,r01,y07,b87,y04}.  The extent of the GC region is roughly 400 pc in diameter, defined by a region with relatively high gas density \\citep[$n_{\\rm{H_2}}\\gtrsim10^4$ cm$^{-3}$;][]{h98}.  This region, sometimes called the Galactic nucleus or ``central molecular zone'', produces 5\\%--10\\% of the Galaxy's infrared and Lyman continuum luminosity and contains 10\\% of its molecular gas \\citep{b87,m96}.  The density and physical diversity of objects in the GC region make it highly complex and require a wide range of observations to unravel.  At the simplest level, it is important to understand how the basic components of the region interact.  Nonthermal emission, in the form of supernova remnants (SNRs) and nonthermal radio filaments \\citep[NRFs;][]{y84}, dominates the cm-wavelength emission in the region \\citep[e.g.,][]{a79,d95}.  The relativistic component of the ISM has also been observed at TeV energies as a diffuse source tracing the molecular gas distribution \\citep{a06}.  However, it is not clear how these energetic electrons traced by the nonthermal emission affect other components of the GC interstellar medium.  For example, energetic electrons may explain the unusually high ambient gas temperature in the region and subsequent low star formation efficiency \\citep{y07}.  Thermal emission traces star formation regions and photoionized clouds.  An accurate census of thermal emission can constrain the amount of star formation occurring there or find new star forming regions.  Separating the thermal and nonthermal processes also allows an estimate of their basic properties \\citep[$n_e$, $B$;][]{h96}.  Single-dish radio continuum observations are a useful tool for studying the large-scale properties of the GC region.  Several other single-dish surveys of radio continuum emission from the GC region have been conducted \\citep{a79,ha87,r90,h92,d95}.  \\citet{a79} observed the Galactic disk from $l=60$\\sdeg\\ to the GC region near 6 cm with the 100 m Effelsberg telescope.  This was one of the first single-dish surveys of the GC region with a resolution of a few (2\\damin6) arcminutes and it discovered many compact and diffuse sources.  Since that time, the Parkes 64 m and Effelsberg 100 m telescopes have been used to create complete surveys of the Galactic disk in the northern and southern celestial skies near 12 cm, which were sensitive to faint emission on large scales \\citep{r90,d95}.  These studies revealed dozens of new SNR candidates, compact \\hii\\ regions, and Galactic loops and spurs, vividly demonstrating the chaotic structure of the Galactic interstellar medium.  Although most of the world's best radio telescopes have surveyed the radio continuum in the GC region, there has been limited study of spectral indices on arcminute scales.  We were motivated to extend upon previous observations using the largest, fully-steerable telescope, the Green Bank Telescope\\footnote{The GBT is operated by the National Radio Astronomy Observatory, which is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} (GBT).  The high resolution and sensitivity of the GBT observations allow us to separate and quantify flux from all thermal and nonthermal emitters.  Section \\ref{gcsurvey_obs} describes the observations and data reduction.  In \\S\\ \\ref{gcsurvey_res}, the results of the survey are described, including the compilation of compact and extended source catalogs at 3.5 and 6 cm and a detailed discussion of the radio continuum properties for each source in the region.  We calculate the percentage of flux from sources in the central degrees of the Galaxy from thermal/nonthermal processes at 3.5 and 6 cm.  Finally, \\S\\ \\ref{gcsurvey_con} summarizes the results of this analysis.  A study of the continuum emission from the GC lobe \\citep{s84} is not included here, but will be discussed in detail in a later paper \\citep{l08}. ",
        "conclusions": "\\label{gcsurvey_con} This paper has shown results from a new survey of the radio continuum emission from the central degrees of the Galaxy at 90, 20, 6, and 3.5 cm with the GBT.  The 6 and 3.5 cm surveys are the most sensitive, highest resolution, single-dish radio surveys of the central degrees of the GC made at these wavelengths.  The primary products of this study are catalogs of all compact and extended sources, including a spectral index analysis of these sources.  We have shown that the compact sources ($\\theta\\lesssim5$\\arcmin) detected at 6 and 3.5 cm surveys are most likely to be Galactic \\hii\\ regions and are mostly found near well-known \\hii\\ complexes, such as Sgr B and Sgr E.  About one quarter of the sources detected at 3.5 cm are also detected at 6 cm;  most of these sources have thermal spectral indices.  Extended, nonthermal emission in the 3.5, 6, and 20 cm surveys is found associated with GC star-forming regions.  The emission is on size scales of tens of arcminutes and is particularly found for $l=358\\ddeg5-359\\ddeg5$.  The brightness-distribution function of NRFs estimated by high resolution observations indicates that they are much more numerous at low flux densities, so the extended nonthermal emission may yet be resolved as new, nonthermal radio filaments.  Nonthermal emission also traces supernova remnants, and these observations find nonthermal emission in Sgr B, consistent with a recent report of a supernova in that star-forming region.  Another interesting result from our spectral index study is that the 6/3.5 cm spectral index distribution around the G359.1--0.5 SNR is consistent with the idea that it is interacting with a neighboring molecular cloud that may be in the GC region.  We also find a structure south of the Arched filaments that is thermal and seems to be morphologically connected to the well-known Arched filaments.  The Arched filaments and the new southern Arched filaments (G0.07--0.2) form a contiguous structure that is a radio-continuum counterpart to the Radio Arc Bubble.  The Bubble surrounds the brightest part of the nonthermal emission from the Radio Arc and is close to the dense star clusters, the Quintuplet and Arches clusters.  Several recent works has definitively shown that the ionization of the Arched filaments is caused by the Arches star cluster, while the rest of the Bubble is ionized by the Quintuplet star cluster \\citep{r01,la02,si07}.  \\citet{si07} further suggests that the Bubble was formed by the winds and ionizing photons of massive stars the Quintuplet cluster.  This work highlights the fact that the Bubble --- a thermal structure formed by stellar action --- has a clear effect on the brightness of the nonthermal Radio Arc.  The reason for the enhanced brightness inside the Bubble is not clear, but should be considered in models for the formation of NRFs.  Through a combination of slice and integrated flux analysis, all objects detected in the surveys are catagorized as having a thermal or nonthermal origin.  Thus, the distribution of flux between thermal and nonthermal sources can be quantified.  Within the central 4\\sdeg$\\times$1\\sdeg\\ of the Galaxy, the thermal to nonthermal flux fractions for all discrete emission are 28\\%/72\\% at 3.5 cm and 19\\%/81\\% at 6 cm.  This does not include the background synchrotron contribution from the Galactic plane, which begins to dominate the Galaxy's flux density for wavelengths longer than 6 cm.  Also, some of these sources in the field are likely to be in the foreground of the GC region, although the density of gas and stars is generally much higher in the GC region, so most sources are likely to truly be in the central few hundred parsecs of the Galaxy.  The high fraction of nonthermal emission in the radio continuum emission is consistent with the idea that the cosmic ray density is enhanced in the nuclear disk, assuming that the magnetic field is not unusually strong ($\\sim1$ mG) in the region.  Other studies have also found evidence for such a cosmic ray enhancement, which seems heat and ionize molecular gas in the GC region, possibly affecting the star formation process \\citep{y07}."
    },
    "0801/0801.3682_arXiv.txt": {
        "abstract": "{}% {% We present the current status of ongoing searches for molecular hydrogen in high-redshift ($1.8<\\zabs \\le 4.2$) Damped Lyman-$\\alpha$ systems (DLAs) capitalising on observations performed with the ESO Very Large Telescope (VLT) Ultraviolet and Visual Echelle Spectrograph (UVES). }% {% We identify 77 DLAs/strong sub-DLAs, with $\\log N(\\HI)\\ge 20$ and $\\zabs >1.8$, which have data that include redshifted H$_2$ Lyman and/or Werner-band absorption lines. This sample of H\\,{\\sc i}, H$_2$ and metal line measurements, performed in an homogeneous manner, is more than twice as large as our previous sample (Ledoux et al. 2003) considering every system in which searches for H$_2$ could be completed so far, including all non-detections. }% {% H$_2$ is detected in thirteen of the systems, which have molecular fractions of values between $f\\simeq 5\\times 10^{-7}$ and $f\\simeq 0.1$, where $f=2N($H$_2)/(2N($H$_2)+N(\\HI))$. Upper limits are measured for the remaining 64 systems with detection limits of typically $\\log N($H$_2)\\sim 14.3$, corresponding to $\\log f<-5$. We find that about 35\\% of the DLAs with metallicities relative to solar [X/H$]\\ge -1.3$ (i.e., 1/20$^{\\rm th}$ solar), with X~=~Zn, S or Si, have molecular fractions $\\log f>-4.5$, while H$_2$ is detected -- regardless of the molecular fraction -- in $\\sim 50$\\% of them. In contrast, only about 4\\% of the [X/H$]<-1.3$ DLAs have $\\log f>-4.5$. We show that the presence of H$_2$ does not strongly depend on the total neutral hydrogen column density, although the probability of finding $\\log f>-4.5$ is higher for $\\log N(\\HI)\\ge 20.8$ than below this limit (19\\% and 7\\% respectively). The overall H$_2$ detection rate in $\\log N(\\HI)\\ge20$ DLAs is found to be about 16\\% (10\\% considering only $\\log f>-4.5$ detections) after correction for a slight bias towards large $N(\\HI)$. There is a strong preference for H$_2$-bearing DLAs to have significant depletion factors, [X/Fe$]>0.4$. In addition, all H$_2$-bearing DLAs have column densities of iron into dust grains larger than $\\log N($Fe$)_{\\rm dust}\\sim 14.7$, and about 40\\% of the DLAs above this limit have detected H$_2$ lines with $\\log f>-4.5$. This demonstrates the importance of dust in governing the detectability of H$_2$ in DLAs. Our extended sample supports neither the redshift evolution of the detection fraction of H$_2$-bearing DLAs nor that of the molecular fraction in systems with H$_2$ detections over the redshift range $1.8<\\zabs\\le 3$. }% {% }% ",
        "introduction": "Damped Lyman-$\\alpha$ systems (DLAs) were discovered in the seventies \\citep[e.g., ][]{Lowrance72, Beaver72, Carswell75, Wright79} and identified afterwards as redshifted damping absorptions from large column densities of neutral atomic hydrogen \\citep[][]{Smith79}.  Numerous DLAs, with $N(\\HI)\\ge 2\\times 10^{20}$\\,atoms~cm$^{-2}$, have been discovered through large dedicated surveys \\citep[e.g., ][]{Wolfe86} and more recently thanks to the huge number of quasar spectra available from the Sloan Digital Sky Survey \\citep{Prochaska05}. Because DLAs contain most of the neutral hydrogen available for star formation in the Universe \\citep{Wolfe86, Lanzetta91} and are associated with numerous metal absorption lines, they probably arise in the interstellar medium of protogalaxies, progenitors of present-day galaxies \\citep[see, e.g., ][]{Wolfe00, Haehnelt00, Wolfe05}. Our understanding of DLAs is mainly based on the study of low-ionisation metal absorptions \\citep[e.g., ][]{Prochaska02} but also high-ionisation species \\citep{Lu96, Wolfe00, Fox07a, Fox07b} and, in a few cases, molecular absorptions \\citep[e.g., ][]{Ledoux03}. The latter are not conspicuous however in contrast to what is seen in the Galaxy and, for a long time, only the DLA towards Q\\,0528$-$2505 was known to contain H$_2$ molecules \\citep{Levshakov85}. H$_2$-bearing DLAs are nevertheless crucial to understand the nature of DLAs because molecular hydrogen is an important species to derive the physical conditions in the gas \\citep[see, e.g., ][]{Tumlinson02, Reimers03, Hirashita05, Srianand05, Cui05, Noterdaeme07}.  The first systematic search for molecular hydrogen in high-redshift ($\\zabs >1.8$) DLAs was carried out using the Ultraviolet and Visual Echelle Spectrograph (UVES) at the Very Large Telescope (VLT) \\citep{Ledoux03}. It consisted of a sample of 33 DLAs with H$_2$ detected in eight of them. Molecular fractions were found to lie in the range $-3.5<\\log f<-1$ with $f=2N($H$_2)/(2N($H$_2)+N(\\HI))$. Upper limits of typically $N($H$_2)\\sim 2\\times 10^{14}$~cm$^{-2}$ (corresponding to $\\log f<-5$) were measured in the other systems. More recently, we noted a correlation between the presence of molecular hydrogen and the metallicity of high-redshift DLAs \\citep{Petitjean06}. High molecular fractions ($\\log f>-4$) were found in about 40\\% of the high-metallicity DLAs ([X/H$]\\ge 1/20^{\\rm th}$ solar) whilst only $\\sim 5$\\% of the [X/H$]<-1.3$ DLAs have $\\log f>-4$. Other papers by our group focused on specific detections. We presented the analysis of three systems with low molecular fractions, i.e., $\\log f<-4$, one of them having a low metallicity \\citep{Noterdaeme07lf}, and the H$_2$-bearing DLA with, to date, the highest redshift, at $\\zabs =4.224$ towards Q\\,1441$+$2737 \\citep{Ledoux06b}. We note that a possible detection of H$_2$ in a DLA towards a Gamma-ray Burst (GRB) afterglow has been reported recently \\citep{Fynbo06}. However, the origin of DLAs at the GRB host-galaxy redshift is very likely to be different from those observed in QSO spectra \\citep[e.g., ][]{Jakobsson06, Prochaska07a}.  We present here the whole sample of UVES high-redshift QSO-DLAs for which the wavelength range where H$_2$ lines are redshifted is covered by the available spectra. This sample is more than twice as large as in our previous study \\citep{Ledoux03}. We present the observations and the UVES DLA sample in Sect.~2 and provide comments on individual absorbers in Sect.~3. We discuss the overall population in Sect.~4 and results in Sects.~\\ref{mol_f} to \\ref{z}. We conclude in Sect.~\\ref{conclusion}. ",
        "conclusions": "}  We present results of the largest survey of molecular hydrogen in high-redshift ($1.8<\\zabs \\le 4.2$) DLAs compiled to date using high signal-to-noise ratio, high spectral-resolution VLT-UVES data. We analyse data for 77 DLAs/strong sub-DLAs with $N(\\HI)\\ge 10^{20}$~cm$^{-2}$, a dataset more than twice as large as that studied previously by \\citet{Ledoux03}. From the thirteen high-redshift H$_2$-bearing DLAs known to date, nine have been discovered by our group. Due to the superb quality of the Ultraviolet and Visual Echelle Spectrograph, we are able to detect unambiguously the H$_2$ absorption features, measure accurate column densities, and in the cases of non-detections derive stringent upper limits.  A double-sided Kolmogorov-Smirnov test shows that the ten H$_2$-bearing systems with $\\log f>-4.5$ (which is our conservative completeness limit) have \\HI\\ column densities that are compatible with those of the overall DLA population. This may be due however to small number statistics. There is evidence (see Sect.~\\ref{overall_pop}) that the probability of finding large molecular fractions is higher in DLAs with large $N(\\HI)$, as observed in the LMC. About $7$\\% of the systems with $\\log N(\\HI)<20.8$ have $\\log f>-4.5$ while $\\sim 19$\\% of the systems with larger \\HI\\ column densities have similar molecular fractions. There is no sharp transition in the molecular fraction of DLAs at any of the measured total hydrogen ($\\HI+$H$_2$) column density.  It is surprising to see that most of the DLAs have very low molecular fractions, i.e., $\\log f \\la -6$. This is much smaller than observed along lines of sight in the Galactic disk \\citep{Savage77}. However, \\citet{Wakker06} measured similar molecular fractions towards high-latitude Galactic lines of sight.  We confirm that a good criterion to find H$_2$-bearing DLAs is to select high-metallicity systems. Indeed, 35\\% of the systems with [X/H$]\\ge -1.3$ have $\\log f>-4.5$ whilst this is the case for only 4\\% of those with metallicities lower than that. Since there is a correlation between metallicity and depletion factor, the latter being defined as [X/Fe] with X$=$Zn, S or Si, H$_2$ is found in DLAs also having the highest depletion factors. Therefore, clouds with large molecular fractions are expected to be dusty and clumpy \\citep[e.g.,][]{Hirashita03}. They can be missed because of the associated extinction and/or because of their small cross-section.  The presence of H$_2$ is closely related to the dust column density. Indeed, all detections pertain to systems where the column density of iron into dust grains is larger than $5\\times 10^{14}$~cm$^{-2}$, and about 40\\% of these systems have detectable H$_2$. This shows that the presence of dust is an important ingredient in the formation of H$_2$ in DLAs. The low molecular fractions measured for most of the DLAs are probably a consequence of low abundances of metals and dust.  We show that the probability of finding H$_2$ increases with increasing velocity width of low-ionisation metal line profiles. The correlation between velocity width and metallicity \\citep{Ledoux06a}, if interpreted as a mass-metallicity relation, provides a natural explanation in which H$_2$-bearing DLA systems are preferably associated with massive objects where star-formation is enhanced. There is no evidence of systematic outflows in H$_2$-bearing DLAs neither from the H$_2$ profiles nor from the study of the associated high-ionisation phase.  From the comparison between H$_2$-bearing systems and the overall UVES sample, we show that there is no evolution with redshift of the fraction of H$_2$-bearing DLAs nor of the molecular fraction in systems with detected H$_2$ over the range $1.8<\\zabs\\le 3$. This, compared to the large amounts of H$_2$ observed in the local Universe, suggests that a significant increase of the molecular fraction in DLAs could take place at redshifts $\\zabs\\le 1.8$. Ultraviolet observations from space are therefore needed to observe the H$_2$ Lyman and Werner bands in low and intermediate redshift DLAs.  Increasing the sample of H$_2$-bearing DLAs is required to cover a large range in column density and derive the H$_2$ frequency distribution, $f(N($H$_2),${\\sl X}) -- where {\\sl X} is the absorption distance -- in a way similar to studies completed for \\HI\\ \\citep{Lanzetta91}. This would assess whether the steep slope observed in $f(N(\\HI),${\\sl X}) at large column densities \\citep[e.g., ][]{Prochaska05} is due to the conversion \\HI$\\rightarrow$H$_2$ \\citep{Schaye01} rather than to a magnitude-dependent bias from dust obscuration of the background quasars \\citep[e.g., ][]{Boisse98, Vladilo05, Smette05}. Both effects can explain the apparent lack of DLAs of both high metallicity and large \\HI\\ column density, as observed in all DLA samples.  While the exact nature of DLAs is still an open debate, H$_2$-bearing DLAs provide interesting probes of the physical conditions close to star-forming regions. About a decade ago, only one H$_2$-bearing DLA was known. With the criteria revealed by the present survey, it will be possible to select these systems more efficiently and their numbers should increase rapidly. This opens up the exciting prospect of the detailed study of the ISM in distant galaxies.  \\acknowledgement{We thank an anonymous referee and the language editor for useful comments that improved the paper. PN is supported by a PhD studentship from ESO. PPJ and RS gratefully acknowledge support from the Indo-French Centre for the Promotion of Advanced Research (Centre Franco-Indien pour la Promotion de la Recherche Avanc\\'ee) under contract No. 3004-3. }"
    },
    "0801/0801.3966_arXiv.txt": {
        "abstract": "We present some results obtained from the synthesis of Stokes profiles in small-scale flux tubes with propagating MHD waves. To that aim, realistic flux tubes showing internal structure have been excited with 5 min period drivers, allowing non-linear waves to propagate inside the magnetic structure. The observational signatures of these waves in Stokes profiles of several spectral lines that are commonly used in spectropolarimetric measurements are discussed. ",
        "introduction": "The observed wave behaviour in network and facular solar regions, and the role of waves to connect photospheric and chromospheric layers, have drawn a considerable discussion in the recent literature, since it has been proposed that thin magnetic flux tubes in these regions can act as wave guides, supplying energy to the upper layers, in particular for waves with periodicity in the 5-min range \\citep{DePontieu+etal2004}. However, observational studies of the relation between the photospheric and chromospheric signals do not provide a unique picture. It is not clear whether the oscillations remain coherent through the atmosphere \\citep{Lites+Rutten+Kalkofen1993, Krijer+etal2001, Judge+etal2001, Centeno+etal2006b}. There are some hints that the coupling between the photosphere and the chromosphere is within vertical channels \\citep{Judge+etal2001, Centeno+etal2006b}, however, other results point toward inclined wave propagation \\citep{Bloomfield+etal2006, DePontieu+etal2003}. In particular, all these studies point out the important role of the magnetic canopy and the $c_S=v_A$ transformation layer. The dominance of long-period oscillations in the chromosphere of these regions, also remains to be explained.  Here we perform numerical simulations of magneto-acoustic waves in small-scale flux tubes, with a subsequent Stokes diagnostics. We address the questions of coherence of oscillations through the atmosphere, wave behaviour at the canopies and the change of the wave period with height. ",
        "conclusions": "In this paper we have studied the wave behaviour in small-scale flux tubes by means of non-linear numerical simulations and Stokes diagnostics. The following items summarize our conclusions.   (i) If the driver that excites oscillations has a vertical component, velocity oscillations with an amplitude of some 150 \\hbox{m$\\;$s$^{-1}$} and Stokes $V$ amplitude oscillations of 10$^{-3}$, in units of the continuum intensity,  may be detected in low spatial resolution observations. The oscillations are coherent through the whole atmosphere (with the corresponding time delay due the upward propagation).  (ii) If the driver that excites oscillations is purely horizontal, no variations would be detected in observations with reduced spatial resolution, since the vertical velocity variations are produced by the surface mode and are in antiphase on both sides of the tube.  (iii) The LOS velocity shows a weak non-linear behaviour in canopy regions, already at heights of 300--350 km in the photosphere.  (iv) It is important to take into account the radiative losses of oscillations. They can produce a decrease of the effective cut-off frequency, making possible the propagation of the 5-min oscillations in the chromosphere in vertical small-scale magnetic structures."
    },
    "0801/0801.1708_arXiv.txt": {
        "abstract": "We present observations of the dwarf galaxies Draco and Ursa Minor, the local group galaxies M32 and M33, and the globular cluster M15 conducted with the Whipple 10m gamma-ray telescope to search for the gamma-ray signature of self-annihilating weakly interacting massive particles (WIMPs) which may constitute astrophysical dark matter (DM).  We review the motivations for selecting these sources based on their unique astrophysical environments and report the results of the data analysis which produced upper limits on excess rate of gamma rays for each source.  We consider models for the DM distribution in each source based on the available observational constraints and discuss possible scenarios for the enhancement of the gamma-ray luminosity.  Limits on the thermally averaged product of the total self-annihilation cross section and velocity of the WIMP, $\\left<{\\sigma}v\\right>$, are derived using conservative estimates for the magnitude of the astrophysical contribution to the gamma-ray flux.  Although these limits do not constrain predictions from the currently favored theoretical models of supersymmetry (SUSY), future observations with VERITAS will probe a larger region of the WIMP parameter phase space, $\\left<{\\sigma}v\\right>$ and WIMP particle mass ($m_\\chi$). ",
        "introduction": "The existence of dark matter (DM) is supported by a variety of observational data including measurements of the cosmic microwave background \\citep{spergel2007}, the large-scale distribution of galaxies \\citep{tegmark2004}, and gravitational lensing \\citep{clowe2006}.  In the $\\Lambda$CDM cosmological model that is currently favored by these data, DM comprises approximately $\\sim$26\\% of the total energy density of the universe \\citep{spergel2007}. However, the nature of the particles that constitute DM remains unknown.  A popular DM candidate is weakly interacting massive particles (WIMPs), which existed in thermal equilibrium during the early universe and later decoupled as the universe expanded.  Since the time of decoupling, the WIMPs have remained non-relativistic, behaving as a collisionless fluid on all but perhaps the shortest spatial scales.  In order to reproduce the observed relic density of DM, this hypothetical particle would need to have a cross section on the scale of weak interactions.  A stable particle with these properties, the lightest neutralino $\\chi$, can be accommodated in theories of supersymmetry (SUSY).  The mass of the neutralino is constrained to be $\\gtrsim$ 6 GeV by CMB measurements and accelerator searches \\citep{bottino2003} and $\\lesssim$ 100 TeV by the unitarity limit on the thermal relic \\citep{griest1990}.  In the conventional SUSY scenarios, the neutralino is a Majorana particle which can efficiently self-annihilate in astrophysical environments with high DM density producing secondary particles including high-energy gamma rays.  The former and current generation of Air-Cherenkov Telescopes (ACTs), including Whipple, HEGRA, CANGAROO-III, VERITAS, H.E.S.S., and MAGIC are sensitive in the gamma-ray energy range from below 100 GeV to above 10 TeV and can therefore make a substantial contribution to the search for the signatures of DM self-annihilation.  Recently several ACTs have detected gamma rays from the Galactic Center (G.C.) \\citep{albert2006,aharonian2004,kosack2004}.  Although a more traditional astrophysical origin for this signal is currently favored \\citep{atoyan2004,aharonian2005}, DM self-annihilation has been proposed as a possible explanation for these observations \\citep{horns2005}.  We present observations taken with the Whipple 10m telescope of five astrophysical sources with the purpose of detecting the signature of DM self-annihilation.  Section \\ref{sourceReviewSection} summarizes the motivations for selecting each source.  In Section \\ref{dmFluxSection} we discuss the signature of DM self-annihilation into gamma rays and its dependence on the source astrophysics and the particle physics properties of the WIMP.  In Section \\ref{dataSection} we review the atmospheric Cherenkov technique and the methods used to analyze the data.  Results of the data analysis are described in Section \\ref{resultsSection}.  Models for the DM distribution and scenarios for the enhancement of the gamma-ray flux are presented in Section \\ref{interpretationSection}.  We conclude in Section \\ref{susyLimitsSection} by discussing the implications of these observations for the parameter space of allowed SUSY models. ",
        "conclusions": "We have conducted a search for the gamma-ray signature of neutralino self-annihilation from five sources: the dwarf spheroidal galaxies Draco and Ursa Minor, the globular cluster M15, and the local group galaxies M32 and M33.  Each of these sources was chosen as a favorable representative of different astrophysical conditions that could potentially enhance the neutralino density and gamma-ray self-annihilation flux.  For a generic MSSM model of the neutralino, the self-annihilation flux is only detectable by the Whipple 10m telescope if such an enhancement is significant.  A standard analysis of the data revealed no significant excesses, and upper limits on the gamma-ray flux from each source were derived relative to the flux of the Crab Nebula.  We have derived limits on $\\left<{\\sigma}v\\right>$ of the neutralino as a function of its mass using models for the DM profile of each source and the generic differential gamma-ray spectrum of the neutralino self-annihilation constructed from two representative channels, $b\\bar{b}$ and $\\tau^{+}\\tau^{-}$.  Using the upper limit on the gamma-ray rate measured from Draco and the most conservative estimate of the DM distribution in this source, we obtain 95 \\% C.L.  upper limits on $\\left<{\\sigma}v\\right>$ of $<1.9\\times10^{-21}$ cm$^{3}$ s$^{-1}$ and $<1.2\\times10^{-22}$ cm$^{3}$ s$^{-1}$ for a neutralino of mass 1 TeV annihilating exclusively through the $b\\bar{b}$ and $\\tau^{+}\\tau^{-}$ channels, respectively.   We have considered potential enhancements to the DM self-annihilation flux including the effects of DM substructure and the adiabatic compression of the DM halo due to baryonic infall.  These scenarios could enhance the annihilation flux by as much as $10^{4}$ and possibly higher.  However the uncertainties of these estimates are large due to the poorly understood dynamics in the cores of the objects as well as their potentially complex merging histories.  If one assumes that DM is composed of neutralinos with $\\left<{\\sigma}v\\right> \\simeq 3 \\times$ 10$^{-26}$ cm$^{3}$ s$^{-1}$, then some extreme enhancement scenarios for M32 and M33 may be ruled out.  The current generation of ACTs such as VERITAS, H.E.S.S., MAGIC, and CANGAROO-III have the potential to improve significantly the sensitivity of these measurements and thus probe a larger region of the MSSM parameter space.  For a source with a Crab Nebula-like spectrum, VERITAS has a flux sensitivity $\\sim$10 times better than the Whipple 10m telescope and a peak detection rate near 150 GeV.  The lower energy threshold of VERITAS will allow it to be sensitive to the low to intermediate neutralino mass range of 100 GeV--1 TeV.  With the improved sensitivity of these instruments, conventional astrophysical backgrounds may become significant.  The sensitivity of observations of the G.C. to DM annihilations is already limited by the presence of such backgrounds.  Observations of extragalactic sources such as those discussed in this work have the potential to avoid this limitation. Furthermore, the identification of the unique spectral signature of DM self-annihilations in two or more of these sources would effectively rule out a traditional astrophysical process.  Next-generation ACT instruments such as the currently planned Cherenkov Telescope Array (CTA) and Advanced Gamma Ray Imaging System (AGIS) will potentially be 10$^{2}$--10$^{3}$ times as sensitive as the Whipple 10m telescope and could perform dedicated deep observations with 10--10$^{2}$ times longer exposure than the observations presented in this work.  This may allow the exclusion of MSSM models with even the most conservative assumptions for the DM distribution in the sources with the lowest anticipated astrophysical backgrounds such as the dwarf spheroidal galaxies."
    },
    "0801/0801.2017_arXiv.txt": {
        "abstract": "We estimate the amplitude of perturbation in dark energy at different length scales for a quintessence model with an exponential potential. It is shown that on length scales much smaller than  hubble radius, perturbation in dark energy is negligible in comparison to that in in dark matter. However, on scales comparable to the hubble radius ($\\lambda_{p}>1000\\mathrm{Mpc}$) the perturbation in dark energy in general cannot be neglected. As compared to the $\\Lambda$CDM model,  large scale matter power spectrum is suppressed in a generic quintessence dark energy model. We show that on  scales $\\lambda_{p} < 1000\\mathrm{Mpc}$, this suppression is primarily due to different background evolution compared to $\\Lambda$CDM model. However, on much larger scales perturbation in dark energy can effect matter power spectrum significantly. Hence this analysis can act as a discriminator between $\\Lambda$CDM model and other generic dark energy models with $w_{de} \\neq -1$. ",
        "introduction": "Cosmological observations  suggest that about $70\\%$ content of our universe is made of a form of matter which drives the accelerated expansion of the universe\\cite{obs_proof}. These observations include Supernova type Ia observations \\cite{nova_data1}, observations of Cosmic Microwave Background \\cite{boomerang,wmap_params,2003Sci...299.1532B} and large scale structure \\cite{2df, sdss, 2004PhRvD..69j3501T}. The accelerated expansion of the universe can  of course be explained by introducing a cosmological constant $\\Lambda$ in the Einstein's equation \\cite{ccprob_wein,review3}. However, the cosmological constant model is plagued by the fine turning problem \\cite{ccprob_wein}. This has motivated the study of dark energy models to explain the current accelerated expansion of the universe. The simplest model as an alternative to cosmological constant model is to assume that this accelerated expansion is driven by a canonical scalar field with a potential $V(\\phi)$. This class of dark energy models are known as quintessence models and the scalar field is known as a quintessence field. Various quintessence models have been studied in literature \\cite{quint1,ferreira,liddle}. There exists another class of string theory inspired scalar field dark energy models known as tachyon models \\cite{tachyon1,2003PhRvD..67f3504B}. Models of dark energy which allow $w<-1$ are known as phantom models \\cite{2002PhLB..545...23C}. Phantom type dark energy can also be realized in a scalar tensor theory of gravitation. (See for example Ref.\\cite{STG}.) Other scalar field models include k-essence field \\cite{2001PhRvD..63j3510A}, branes \\cite{brane1}, Chaplygin gas model and its generalizations \\cite{chaply}. There are also some phenomenological models \\cite{water}, field theoretical and reorganization group based models  (see e.g. \\cite{tp173}), models that unify dark matter and dark energy \\cite{unified_dedm1}, holographic dark energy models \\cite{HGDE}  and many others like those based on horizon thermodynamics (e.g. see \\cite{2005astro.ph..5133S}). For reviews of dark energy models see for instance Ref.\\cite{DEreview} and for constraining parameters using observations see Ref.\\cite{param_fit}.  Homogeneous dark energy distribution leads to accelerated expansion of the Universe which, in turn, governs the luminosity distance and angular diameter distance. The rate of expansion also influences growth of density perturbations in the universe. This is evident from  the abundance of rich clusters of galaxies and their evolution and the Integrated Sachs Wolfe effect \\cite{isw0}.  In this paper we present a set of argument which lead to the conclusion that inhomogeneous dark matter with homogeneous dark energy at all length scales is inconsistent with Einstein's equation if $p_{de} \\neq -\\rho_{de}$. We further analyze how dark matter power spectrum is influenced by perturbation in dark energy with an evolving equation of state parameter.    Dark energy perturbations have been extensively studied in the linear approximation \\cite{weller_lewis,bean_dore,depert}. It was shown in Ref.\\cite{weller_lewis} that dark energy perturbations affect the low $l$ quadrapole in the CMB angular power spectrum through the ISW effect. This analysis was done for a constant equation of state parameter. For models with $w>-1$ this effect is enhanced while for phantom like models it is suppressed. In these models dark matter perturbations and dark energy perturbations are anti-correlated for large effective sound speeds. This anti-correlation is a gauge dependent effect \\cite{bean_dore}. Detailed studies in dark energy perturbations also include \\cite{chpgas_pert,sph_coll}. Clustering of dark energy within matter over density and voids were studied by Mota et al\\cite{mota}.   In this paper we study the evolution of perturbation in dark energy in a quintessence model which results in an evolving equation of state parameter different from that considered in Refs.\\cite{amendola1, weller_lewis}. We use a specific model of scalar field dark energy with an exponential potential. We find that although for scales much smaller than hubble radius the perturbation in dark energy is small, for scales $\\geq H^{-1}$ the perturbation in dark energy can be comparable to that in matter. Hence, although on small scales ($<< H^{-1}$) the dark energy can be treated as homogeneous, one has to take into account the perturbation in dark energy over scales $\\sim H^{-1}$ if $w_{de} \\neq -1$. Clearly, in the specific case of $w_{de} = -1$ the dark energy is homogeneous at all scales.    We choose to work in the longitudinal gauge as in that case we can directly relate the metric perturbation $\\Phi$ to the gravitational potential perturbation. For a specific model, we investigate how quintessence  dark energy influences matter power spectrum. We show that on scales $\\lambda_{p} < 1000\\mathrm{Mpc}$, the matter power spectrum is not significantly affected whether or not  we include fluctuations in the quintessence field in  perturbation equations. However, on much larger scales, including or excluding fluctuations in quintessence field does result in significant changes in the matter power spectrum.    This paper is organized as follows. In Section \\ref{sec::inhomogeneous DE} we  discuss the background cosmology for matter and the scalar field system and describe the cosmological perturbation equation in longitudinal gauge for this system. In  Section \\ref{sec::Numerical Solutions} we obtain the numerical solution of the perturbation equation to determine the ratio of the perturbations in dark energy to the perturbations in matter at various length scales. Section \\ref{sec::Conclusions} summarizes the results. ",
        "conclusions": "\\label{sec::Conclusions}  In this paper,  we have investigated the perturbations in dark energy. This is motivated by the fact that the assumption that the distribution of dark energy (with $w \\neq -1$) is homogeneous at all length scales is inconsistent with the observational fact that dark matter is distributed inhomogeneously. On length scales comparable to or greater than the Hubble radius ($\\lambda_{p}> 1000\\mathrm{Mpc}$), the perturbation in dark energy can become comparable to perturbation in matter if $w_{de} \\neq -1$. The model parameters we have chosen correspond to $w \\approx -0.8 $ and $w \\approx -0.9$, which are within the range allowed by Supernova observations and WMAP5 observations. Given this range, the evolution of perturbations differs significantly. For scales $\\lambda_{p}< 1000\\mathrm{Mpc}$, the perturbation in dark energy $\\delta_\\phi$ can be neglected in comparison with the perturbation in matter $\\delta_{m}$ at least in the  linear regime. We have demonstrated this using an exponential potential for the quintessence field. This result agrees with those presented in Ref.\\cite{mota} on sub-Hubble scales.  We have further demonstrated that quintessence dark energy results in suppression of matter power spectrum relative  to  $\\Lambda CDM$ model. We found that at length scale of $\\lambda_{p}= 1000\\mathrm{Mpc}$ and for the value of the parameter $\\lambda = 1$, the matter power spectrum is suppressed by about $4\\%$ compared to its value in the $\\Lambda$CDM model for the same set of initial condition. However, at $\\lambda_{p}= 10^{5}\\mathrm{Mpc}$, matter power spectrum is suppressed by about $15\\%$ compared to its value in the $\\Lambda$CDM model. We have demonstrated that on scales $\\lambda_{p}< 1000\\mathrm{Mpc}$ this suppression is primarily due to  different background evolution relative to $\\Lambda$CDM model. The resultant matter power spectrum is nearly invariant even if we assume that quintessence field is homogeneous on these scales. However, on much larger scale $\\lambda_{p}> 1000\\mathrm{Mpc}$, including or excluding fluctuation in the quintessence field results in significant changes in the  matter power spectrum.   All these results emphasize that dark energy can indeed be treated as nearly homogeneous on scales $\\lambda_{p}< 1000\\mathrm{Mpc}$. However, on much larger scales ($\\lambda_{p} > 1000\\mathrm{Mpc}$), if the equation of state parameter deviates from -1, then perturbations in dark energy does influences  matter power spectrum significantly. If a definitive detection of perturbations in dark energy is made, it will certainly rule out the cosmological constant at least as the sole candidate of dark energy."
    },
    "0801/0801.0154_arXiv.txt": {
        "abstract": "% We describe very accurate imaging of radio spectral index for the inner jets in three FRI radio galaxies. Where the jets first brighten, there is a remarkably small dispersion around a spectral index of 0.62. This is also the region where bright X-ray emission is detected. Further from the nucleus, the spectral index flattens slightly to 0.50 - 0.55 and X-ray emission, although still detectable, is fainter relative to the radio. The brightest X-ray emission from the jets is therefore not associated with the flattest radio spectra, but rather with some particle-acceleration process whose characteristic energy index is 2.24.The change in spectral index occurs roughly where our relativistic jet models require rapid deceleration.  Flatter-spectrum edges can be seen where the jets are isolated from significant surrounding diffuse emission and we suggest that these are associated with shear. ",
        "introduction": "\\label{radio}  The detection of X-ray synchrotron emission on kiloparsec scales in the bases of low-luminosity (FR\\,I) radio jets (e.g.\\ \\citealt{Hard02,Hard05,ngc315xray}) requires distributed particle acceleration, consistent with the failure of adiabatic models to reproduce the observed brightness profiles at radio wavelengths \\citep{LB04}. The acceleration mechanism is not understood. Here, we summarize the results of recent work on this problem, using accurate radio spectral mapping (this section) and detailed radio -- X-ray comparisons (Section~\\ref{xray}).  In the course of our jet-modelling programme \\citep[and references therein]{L07}, we have accumulated high-resolution, multi-frequency radio images of the bases of FR\\,I jets and have used these to derive accurate spectral-index distributions.  We have studied NGC\\,315 \\citep{ngc315ls} and 3C\\,296 \\citep{3c296}; a third example, the nearby radio galaxy 3C\\,31 (z = 0.0169; Laing et al., in preparation), is shown in Fig.~\\ref{fig:3c31alpha}. There is no evidence for any significant deviation from power-law spectra in 3C\\,31 within 70\\,arcsec of the nucleus.  Out to $\\approx$7\\,arcsec in both jets, the spectral index at 1.5-arcsec resolution is slightly steeper ($\\langle \\alpha \\rangle = 0.62$)\\footnote{$S(\\nu) \\propto \\nu^{-\\alpha}$} than the average for the inner jets. From 7 -- 50\\,arcsec in both jets, the mean spectral index is in the range 0.55 -- 0.57. Further from the nucleus, there is a gradual spectral steepening. There are also slight, but significant variations in spectral index across both jets within $\\approx$30\\,arcsec of the nucleus in the sense that their West edges tend to have flatter spectra ($\\langle \\alpha \\rangle = $ 0.52 -- 0.54; Fig.~\\ref{fig:3c31alpha}c).  There is a clear spectral separation between the South jet and the surrounding emission, matching the separation of these regions defined by the sharpest brightness gradients (Figs~\\ref{fig:3c31alpha}a and b). This is particularly clear where the jet first enters the diffuse emission. The spectral identity of the jet is evidently maintained even after it bends abruptly about 2\\,arcmin South of the nucleus and remains until it terminates in a region of high brightness gradient.  The outer South jet in 3C\\,31 is therefore a clear example of the type of spectral structure noted in other FR\\,I sources \\citep{KSR97,3c296} wherein a flatter-spectrum `jet' with a distinct spectral identity is superposed on a `sheath' of steeper-spectrum emission.  \\begin{figure} \\plotone{3c31_alphabw.eps} \\caption{Radio images of 3C\\,31. (a) Sobel-filtered, mean L-band image (normalized by total intensity) at a resolution of 5.5\\,arcsec.  (b) and (c) Spectral index, $\\alpha$ from weighted least-squares, power-law fits to the total intensity.  (b) 5-frequency fit between 1365 and 4985\\,MHz at a resolution of 5.5\\,arcsec FWHM.  (c) 6-frequency fit between 1365 and 8440\\,MHz at a resolution of 1.5\\,arcsec FWHM for the inset area in panel (b).\\label{fig:3c31alpha}} \\end{figure} ",
        "conclusions": "These observations contribute to the developing picture of spectral variations in FR\\,I sources. Bright X-ray emission is detected close to the nucleus, in the faint, well-collimated jet bases that precede the sudden radio brightening (e.g.\\ Fig.~\\ref{fig:xprofiles}). There is approximate morphological correspondence between features in the radio and X-ray brightness distributions after the former brightens, although there are differences on small scales (e.g.\\ Fig.~\\ref{fig:ngc315xray}b). In contrast, there are no systematic transverse variations in the X-ray/radio ratio within $\\approx$30\\,arcsec of the nucleus in NGC\\,315 (the best resolved case; \\citealt{ngc315xray}). Particle acceleration appears to be distributed throughout the jet volume, rather than being exclusively associated with discrete knots or with the boundary.  The ratio of X-ray to radio emission decreases where our kinematic models show that the jets start to decelerate from speeds of 0.8 -- 0.9$c$ (Fig.~\\ref{fig:xprofiles}; \\citealt{LB02,Hard02,CLBC,ngc315xray}).  Where the jets first brighten and before they decelerate, there is a remarkably small dispersion around a radio spectral index of $\\alpha = 0.62$ in the three sources we have studied in detail, as well as 3C\\,66B (Fig.~\\ref{fig:3c31alpha}; \\citealt*{HBW,ngc315ls,3c296}).  The average is dominated by emission immediately after the point at which the jets first brighten. This is also the region from which X-ray emission is detected from the main jets in all four sources (Fig.~\\ref{fig:xprofiles}; \\citealt{HBW,Hard05}).  The spectral index of the fainter emission close to the nucleus in 3C\\,449 \\citep{KSR97}, PKS1333$-$33 \\citep{KBE} and 3C\\,66B \\citep{HBW} appears to be slightly steeper than $\\alpha = 0.62$ although the uncertainties are larger. Further from the nucleus, the spectra flatten slightly to $\\alpha = $ 0.50 -- 0.55, contrary to any naive expectation from models in which electrons are accelerated at the brightening point and suffer synchrotron losses as they propagate. X-ray emission is still detected from these regions, but at a lower level relative to the radio (Fig.~\\ref{fig:xprofiles}). The brightest X-ray emission from the jets is therefore not associated with the flattest radio spectra, but rather with some particle acceleration process whose characteristic energy index is $2\\alpha + 1 = 2.24$. A related result is that an asymptotic low-frequency spectral index of 0.55 is common in FR\\,I jets over larger areas than we consider here \\citep{Young}. Flatter-spectrum edges can be seen where the jets are isolated from significant surrounding diffuse emission, most clearly in NGC\\,315 \\citep{ngc315ls}. Our kinematic models \\citep{LB02,CLBC,3c296} show that all of the jets have substantial transverse velocity gradients and it is plausible that the process that produces the flatter spectrum is associated with high shear \\citep{SO}.  In 3C\\,31, the flatter-spectrum regions (Fig.~\\ref{fig:3c31alpha}c) occur predominantly on the outer edges of bends, perhaps consistent with this idea.  As well as a smooth steepening of the jet spectrum at larger distances from the nucleus, as would be expected from synchrotron and adiabatic losses affecting a homogeneous electron population, multiple spectral components are observed (Fig.~\\ref{fig:3c31alpha}b). Jets appear to retain their identities even after entering regions of diffuse emission and are clearly identifiable by their flatter spectra. They are usually also separated from the surrounding emission by sharp brightness gradients (Fig.~\\ref{fig:3c31alpha}a).  This spine/sheath separation is observed in FR\\,I sources with bridges of emission extending back towards the nucleus (e.g.\\ 3C\\,296; \\citealt{3c296}) as well as tailed sources like 3C\\,31 (Fig.~\\ref{fig:3c31alpha}). Although there is an overall trend for the spectrum of the diffuse emission to steepen towards the nucleus in bridges and away from it in tails \\citep{Parma99}, the variations in individual objects are complex.  The termination regions of jets in tailed FR\\,I sources are perhaps best regarded as bubbles which are continually fed with fresh relativistic plasma by the jets and which in turn leak material into the tails. Their spectral steepening would then be governed by a combination of continuous injection, adiabatic, synchrotron and inverse Compton energy losses and escape."
    },
    "0801/0801.0965_arXiv.txt": {
        "abstract": "As part of a large survey of halo and thick disc stars, we found one halo star, HD 106038, exceptionally overabundant in beryllium. In spite of its low metallicity, [Fe/H] = $-$1.26, the star has log(Be/H) = $-$10.60, which is similar to the solar meteoritic abundance, log(Be/H) = $-$10.58. This abundance is more than ten times higher the abundance of stars with similar metallicity and cannot be explained by models of chemical evolution of the Galaxy that include the standard theory of cosmic-ray spallation. No other halo star exhibiting such a beryllium overabundance is known. In addition, overabundances of Li, Si, Ni, Y, and Ba are also observed. We suggest that all these chemical peculiarities, but the Ba abundance, can be simultaneously explained if the star was formed in the vicinity of a hypernova. ",
        "introduction": "The single stable isotope of beryllium, $^{9}$Be, is a pure product of cosmic-ray spallation of heavy (mostly CNO) nuclei \\citep*{RFH70}. Analyses of Be abundances in metal poor stars \\citep{MBCP97,BDKR99} have found a relationship between [Fe/H]\\footnote{[A/B] = log [N(A)/N(B)]$_{\\rm \\star}$ - log [N(A)/N(B)]$_{\\rm\\odot}$} and log(Be/H) with slope close to one, and between [O/H] and log(Be/H) with slope between 1 and 1.5, depending on the oxygen indicator used. Independently of the behaviour of the [O/Fe] ratio at lower metallicities, these results suggest a primary production of Be in the early Galaxy \\citep{Ki01}.  As a primary element, and assuming cosmic-rays to be globally transported across the early Galaxy, Be may show a smaller scatter than the products of stellar nucleosynthesis \\citep{SY01} at a given time, suggesting its potential use as a cosmochronometer.  So far, the linear relations appear to be very tight, showing a surprisingly low scatter comparable to the measurement errors. This picture, however, might change with the increase of the samples analysed, as hinted by the results of \\citet{BN06}. Nevertheless, turn-off stars of the globular clusters NGC 6397 and NGC 6752 were found \\citep{Pas04,Pas07} to have the same Be abundance of field stars of the same metallicity. This strongly support the production of Be to be a global process. Ages derived from these abundances, in a comparison with a model of the evolution of Be with time \\citep{Va02}, show an excellent agreement with ages derived from theoretical isochrones, supporting the use of Be as a cosmochronometer. Moreover, \\citet{Pas05} showed that Be abundances could be used to study the differences in the time scales of star formation in the halo and the thick disc of the Galaxy.  In this letter, we report the discovery of an extremely Be enriched halo star, HD 106038, with an abundance 1.2 dex higher than stars of similar metallicity. This unique star deviates considerably from the observed relations of Be with Fe and O. It was identified during the analysis of a large sample containing near to one hundred halo and thick disc stars (Smiljanic et al. 2008, in preparation).  Neither the standard scenario for Be production, involving spallation of cosmic-rays on nuclei of the interstellar medium \\citep{Va02}, nor the superbubbles (SBs) scenario \\citep{Par00} seem to be able to produce such Be enriched objects. The SBs model predict a scatter in the Be abundance \\citep{PD00} that may explain the stars found by \\citet{BDKR99} and \\citet{BN06} which have similar atmospheric parameters but Be abundances differing by $\\sim$ 0.5 dex. The very high Be abundance in HD 106038, however, would require an extremely poor mixing of the SNe ejecta with the ISM which seems to be difficult to justify \\citep{Par00}.   \\begin{figure} \\begin{centering} \\includegraphics[width=6.5cm]{smiljanic_fig01.eps} \\caption{Comparison between the spectra of HD 106038 (solid line) and of HIP 7459 (dashed line), a star with close atmospheric parameters and similar metallicity, in the Be region. The dominating element of the nearby blended features are also indicated. The V and Ti features have the same strength in the two stars while some difference in the Zr line is noted.} \\label{fig:be} \\end{centering} \\end{figure}  To the best of our knowledge, there is only one other case of extremely Be enhanced star in the literature.\\footnote{We exclude from the discussion the chemically peculiar A or F stars with enhanced Be lines. The peculiar abundances of these stars are thought to be caused by  effects of diffusion. As shown by \\citet*{Ric02}, these effects do not result in overabundances in stars with similar temperature and metallicity as HD 106038.} The star J37 of the open cluster NGC 6633 was found by \\citet{As05} to have log (Be/H) = $-$9.0 $\\pm$ 0.5. The chemical peculiarities of star J37 might be best explained by the accretion of rocky material similar to chondritic meteorites \\citep{As05}. As we shall see below, the accretion of such a material is unlikely for our population II star. ",
        "conclusions": "Since the standard scenario for cosmic-ray spallation does not explain the enhancement of Be in HD 106038, a peculiar and/or rare event may be related to its formation. A combination of two or more rare events to produce the observed features is unlikely, we therefore concentrate on single events.  To reproduce the very particular chemical pattern of HD 106038, a nucleosynthetic site must be able to overproduce Si and Ni without overproducing other $\\alpha$ and iron group elements. Elements in normal halo stars with the same metallicity as HD106038 come mostly from SNe II. It is therefore unlikely that the same SNe II may produce the observed enhanced [Si/Fe] and [Ni/Fe] ratios. Moreover, it has about 16 times more Be than what models involving SN II predict for its metallicity \\citep{Va02}. SNe Ia produce large amounts of Fe, Ni, and other iron group elements but are not expected to produce large amounts of O and Si \\citep{Iw99}. Therefore, in case of a contribution from SNe Ia, ratios such as [O/Fe] and [Si/Fe] would fall below normal halo stars, contrary to what is observed. Moreover, SNe Ia are expected to produce about one order of magnitude less spallation products than SN II \\citep{Fie02}.  All these features, however, may possibly be found in the ejecta of a hypernova (HNe), which can be enriched in both intermediate mass elements (as S and Si) as in iron group elements \\citep{Nak01,Pod02}. Moreover, HNe can produce large amounts of Be and Li by spallation \\citep{Fie02,Nak04}. Therefore, we suggest that the material which formed this star was probably contaminated by the nucleosynthetic products of a HNe.  \\subsection{Hypernova}\\label{sec:hyp}  Hypernovae are core-collapse SNe (usually of type Ic) with exceptionally large kinetic energy production, resulting in spectra dominated by very broad absorption line blends \\citep{Maz00}. The energy released in the explosion can be one order of magnitude larger than that of normal core-collapse SNe \\citep{Iw03}. Some hypernovae, typically the most massive and energetic events, are linked to Gamma-Ray bursts \\citep{Iw98}.  \\citet{Fie02} and \\citet{Nak04} calculated the yields of spallation products resulting from HNe explosions. While \\citet{Nak04} calculate the energy distribution of the ejecta with a hydrodynamic code and solve the cosmic-ray transfer equation, \\citet{Fie02} use an empirical formula for the energy distribution and do not solve the transfer equation but adopt an approximation to have the mass fraction of the ejecta that produces the spallation products. \\citet{Nak04} claim the simplifications adopted by \\citet{Fie02} to overestimate the yields by a factor $\\sim$ 3.  The yield of Be per HNe can be one or two orders of magnitude larger than the one per SNe II \\citep{Fie02,Nak04}. However, as a rare event, they are not major contributors of Be in the Galaxy. \\citet{Fie02} predict $^{7}$Li/$^{9}$Be $\\sim$ 8.6 and Be/O $\\sim$ $5.6 \\times  10^{-7}$ (both ratios by number). The calculations by \\citet{Nak04} predict $^{7}$Li/$^{9}$Be $\\sim$ 4.2, also by number. Both predictions are close to what is observed in HD 106038; the ratio between the observed excess of $^{7}$Li and $^{9}$Be is $^{7}$Li/$^{9}$Be = 5.6 (while the HNe is expected to have produced all the observed Be abundance, it is responsible only for the excess of Li with respect to the primordial plateau). We cannot estimate the contribution of the possible HNe on the observed oxygen abundance, thus only a lower limit can be placed, Be/O \\textgreater $1.9 \\times 10^{-7}$, given by the assumption that all the observed oxygen has been produced by the HNe.  Both models, however, predict much more $^{6}$Li than observed, $^{7}$Li/$^{6}$Li $\\sim$ 1.9 by \\citet{Fie02} and $^{7}$Li/$^{6}$Li $\\sim$ 1.2 by \\cite{Nak04}, by number. The observed ratio between the excess of $^{7}$Li and the $^{6}$Li abundance is $^{7}$Li/$^{6}$Li $\\leq$ 15. We consider this ratio an upper limit since, given its fragility, some $^{6}$Li has probably been destroyed in previous evolutionary phases. The production of Be without a corresponding production of $^{6}$Li would be extremely difficult to understand.  Nucleosynthetic calculations by \\citet{Nak01} find the ejecta of HNe to have smaller amounts of C and O and larger amounts of Si, S, and Ar, when compared to normal SNe. \\citet{Nak01} and \\citet{No06} also note larger [(Zn,Co)/Fe] and smaller [(Mn,Cr)/Fe] ratios. An overabundance of Zn, which is not observed, can be avoided with a deeper mas cut\\footnote{The coordinate, in mass, separating the part of the star that is ejected and the one that forms the remnant.}, which would also result in a larger [Ni/Fe] \\citep{No06}. These are in qualitatively agreement with the observations, supporting the HNe hypothesis.  The weak s-process in massive stars seems to efficiently produce only elements with a mass number\\footnote{The mass number of Y is A = 89 and of Ba is A $\\sim$ 136.} up to 90 \\citep{RH00}. The Y overabundance may require an enhanced flux of neutrons, which would also contribute to the production of Ni. The Ba overabundance, however, is more difficult to understand. It is not clear whether this same mechanism would result in the overproduction of Ba. Moreover, a significant amount of Ba is expected to be produced only by the main s-process in AGBs or by the r-process, usually associated to massive stars. Since pollution by AGBs is not possible (see below), the Ba overabundance is likely a product of a massive star. For example, although not expected, \\citet*{Maz92} found Ba to be overabundant by a factor of 5 in the spectrum of SN 1987A. Although Ba might pose a problem for our scenario, we recall that theoretical predictions for the r-process elements in HNe are not available. Therefore, whether this scenario is able to explain the Ba abundance is still an open question.  The HNe scenario, at least qualitatively, is able to explain most features observed in HD 106038 within a single peculiar event. More work, however, is still needed to show whether the scenario still holds quantitatively. A detailed comparison with nucleosynthetic predictions of theoretical models is necessary to validate or not the HNe hypothesis.  \\subsection{Other scenarios}\\label{sec:oth}  In this subsection we present some alternative scenarios for the origin of the Be enhancement in HD 106038, which were discarded for the reasons presented below.  A pollution by AGB stars, or any kind of evolved star, can at once be discarded. Although these may explain the s-process elements, and maybe Li, they do not explain the Be overabundance. On the contrary, the material ejected by an AGB or by a massive star, after successive mixing events, would be depleted in Be.  The SBs scenario \\citep{Par00} would be able to reproduce the observed Li/Be and Be/O ratios only with an extreme model where particles are accelerated from pure SNe ejecta. The SBs evolution models, however, predict that most material inside of a SB comes from the ISM. Moreover, the remaining chemical peculiarities are not typical of SNe II ejecta and thus can not be explained within the same scenario.  Another possibility is the engulfing of a sub-stellar object, a planet or planetesimals debris as in star J37 of the open cluster NGC 6633 \\citep{As05}. This, however, can be excluded with a robust quantitative argument. The accreted material would be confined to the surface convective layer of the star. In a metal poor star this layer is much shallower than in a solar metallicity star. With the equation given by \\citet{Mu01} we estimate the surface convective layer of HD 106038 to have $\\sim$ $4.5 \\times 10^{-3}$ M$_{\\odot}$. The mass of Be in this layer is $\\sim$ $7.7 \\times 10^{-13}$ M$_{\\odot}$ while in a star with normal abundance of Be it would be $\\sim$ $4.8 \\times 10^{-14}$ M$_{\\odot}$. Assuming the accreted material to have a composition similar to chondrites meteorites \\citep{Lod03} a mass of Fe of $5.3 \\times 10^{-6}$ M$_{\\odot}$ would also be accreted with the required mass of Be. However, the total mass of Fe in the convective layer of the star is $\\sim$ $3.3 \\times 10^{-7}$ M$_{\\odot}$. In this scenario all, or almost all, Fe in the convective layer of the star would come from the accreted material. HD 106038 would then originally be a population III metal free star; an extremely unlikely possibility.  In addition, we note that the most metal poor star found to host planets has [Fe/H] = $-$0.68 \\citep{Co07}.  A longer exposure to EPs would be the natural explanation if HD 106038, for some reason, was younger than halo stars of the same metallicity. Its Be abundance would be a result of the accumulated action of EPs in a cloud where star formation was, somehow, delayed. Its Be abundance higher than the solar photospheric one could be a sign of a solar or younger age, in agreement with the suggestion of \\citet{Pas04} that Be abundances could be used as a cosmic clock. The position of the star in the HR diagram, although not favourable for a good age determination, favours an older age and argues against this hypothesis.  If the star originated in or near the Galactic centre, the enhanced star formation and supernovae events could provide an enhanced EPs flux. This flux might also originate from a non-stellar source such as the central black hole. A bulge origin for the star, however, seems unlikely. In particular the abundances of Ni, Y, and Ba are not compatible with the ones of bulge stars \\citep{MR94}. Since it requires another source for the Ni and s-process overabundances, we also discard this hypothesis."
    },
    "0801/0801.3304_arXiv.txt": {
        "abstract": "{A promising method for the detection of UHE neutrinos is the Lunar Cherenkov technique, which utilises Earth-based radio telescopes to detect the coherent Cherenkov radiation emitted when a UHE neutrino interacts in the outer layers of the Moon. The LUNASKA project aims to overcome the technological limitations of past experiments to utilise the next generation of radio telescopes in the search for these elusive particles. To take advantage of broad-bandwidth data from potentially thousands of antennas requires advances in signal processing technology. Here we describe recent developments in this field and their application in the search for UHE neutrinos, from a preliminary experiment using the first stage of an upgrade to the Australia Telescope Compact Array, to possibilities for fully utilising the completed Square Kilometre Array. We also explore a new real time technique for characterising ionospheric pulse dispersion which specifically measures ionospheric electron content that is line of sight to the moon.}  \\begin{document} ",
        "introduction": "The origin of the most energetic particles observed in nature, the ultra high energy (UHE) cosmic rays (CR), which have energies extending up to at least $2 \\times 10^{20}$ eV, is currently unknown. Finding the origin of these particles will have important astrophysical implications. However, direct detection of UHE neutrinos is very difficult due to their extremely small interaction cross-sections. Instead, they may be detected indirectly via observation of the Askaryan effect \\cite{askaryan62excess} in the lunar regolith. Askaryan first predicted coherent Cherenkov emission in dielectric media at radio and microwave frequencies. Using the Moon as a large volume neutrino detector, coherent radio Cherenkov emission from neutrino-induced cascades in the lunar regolith can be observed with ground based telescopes. This method was first proposed by Dagkesamanskii and Zheleznykh \\cite{dagkesamanskii89radio} and first applied by Hankins, Ekers and O'Sullivan \\cite{hankins96asearch} using the Parkes radio telescope.  Coherent Cherenkov radiation is a linearly polarised broadband emission. The spectrum of coherent Cherenkov emission rises approximately linearly with frequency until a peak value is reached. The peak frequency is determined by de-coherence and/or attenuation in the regolith, and can vary between a few hundred MHz and approximately 5 GHz. The dependence of the peak frequency on shower geometry makes the choice of an optimum observation frequency non-trivial \\cite{clancy07lunar}. ",
        "conclusions": ""
    },
    "0801/0801.3845_arXiv.txt": {
        "abstract": "The formation of disk galaxies is one of the most outstanding problems in modern astrophysics and cosmology. We review the progress made by numerical simulations carried out on large parallel supercomputers. These simulations model the formation of disk galaxies within the current structure formation paradigm in which the Universe is dominated by a cold dark matter component and a cosmological constant. We discuss how computer simulations have been an essential tool in advancing the field further over the last decade or so. Recent progress stems from a combination of increased resolution and improved treatment of the astrophysical processes modeled in the simulations, such as the phenomenological description of the interstellar medium and of the process of star formation. We argue that high mass and spatial resolution is a necessary condition in order to obtain large disks comparable with observed spiral galaxies avoiding spurious dissipation of angular momentum. A realistic model of the star formation history. gas-to-stars ratio and the morphology of the stellar and gaseous component is instead controlled by the phenomenological description of the non-gravitational energy budget in the galaxy. This includes  the energy injection by supernovae explosions as well as by accreting supermassive black holes at scales below the resolution. We continue by showing that  simulations of gas collapse within cold dark matter halos including a phenomenological description of supernovae blast-waves allow to obtain stellar disks with nearly exponential surface density profiles as those observed in real disk galaxies, counteracting the tendency of gas collapsing in such halos to form cuspy baryonic profiles. However, the ab-initio formation of a realistic rotationally supported disk galaxy with a pure exponential disk in a fully cosmological simulation is still an open problem. We argue that the suppression of bulge formation is related to the physics of galaxy formation during the merger of the most massive protogalactic lumps at high redshift, where the reionization of the Universe likely plays a key role. A sufficiently high resolution during this early phase of galaxy formation is also crucial to avoid artificial angular momentum loss and spurious bulge formation. Finally, we discuss the role of mergers in disk formation, adiabatic halo contraction during the assembly of the disk, cold flows, thermal instability and other  aspects of galaxy formation, focusing on their  relevance to the puzzling origin of bulgeless galaxies. ",
        "introduction": "Galaxies occupy a special place in our quest for understanding the Universe. They are large islands in a nearly empty space and contain most of the ordinary baryonic matter,  stars and interstellar gas, that emits radiation and can thus be detected by astronomers$^1$. Galaxies come essentially in two broad categories$^2$, those in which the luminous mass is arranged in a rotating disk of stars and gas, called disk galaxies or spiral galaxies because of the presence of spiral arms of gas and stars (Figure 1), and those in which the luminous mass is distributed in a smooth, featureless spheroidal structure with little or no rotation,  also known as elliptical galaxies (Figure 1). Disk galaxies have a radial light distribution $I(r)$ that is well fit by a decaying exponential law$^3$, $I(r) \\sim exp(-r/r_d)$, where $r_d$ is a characteristic scale length ($r_d \\sim 2-4$ kpc for typical spiral galaxies). Indeed many disk galaxies contain also a spheroidal stellar component at their center, the stellar bulge, which has structural properties similar to an elliptical galaxies albeit being much smaller in size$^2$. Both types of galaxies are known to contain dark matter, namely matter that is not traced by radiation. In disk galaxies dark matter clearly dominates over luminous matter by mass, as inferred from their high rotation speeds which requires the gravitational pull of a massive and extended halo of dark matter$^4$. We live in a galaxy of the first kind, the Milky Way. Indeed disk galaxies are ubiquitous in the local Universe, and also at the largest distances and earliest epochs at which the best ground and space-based telescopes have been able to study the morphology of galaxies reliably$^5$. Only the most massive galaxies in the Universe do not posess a disk component, while this becomes progressively more dominant compared to the spheroidal component as the mass of the galaxy decreases (Figure 2).     \\begin{figure} \\centering{ \\resizebox{12cm}{!}{\\includegraphics{gdesignlr.ps}}\\\\ \\resizebox{7cm}{!}{\\includegraphics{m87lr.ps}}\\\\ \\caption{Top: a typical disk-dominated galaxy, the nearby spiral galaxy M102 in the Ursa Major constellation, 27 million light years from the Sun (the image was obtained by Chris \\& Dawn Schur from Payson, Arizona at 5150 feet elevation with an amateur telescope). A small spheroidal bulge is visible at the center of the disk. Bottom: a typical spheroidal galaxy with no disk component, the elliptcal galaxy M87, located at 60 million light years from us (credits in the picture).}} \\label{fig:feedback} \\end{figure}      \\subsection{The theoretical framework}  The formation of disk galaxies is one of the major unsolved problems of modern astrophysics. The basic theoretical framework states that disk galaxies arise from the gravitational collapse of a rotating protogalactic cloud of gas within the gravitational potential well of the dark halo$^6$. The gas cools via radiative processes during the collapse and eventually settles in centrifugal equilibrium at the center of the halo potential well forming a rotationally supported gas disk provided that some angular momentum is retained during the collapse $^7$. These ideas were developed two decades ago and they still constitute the backbone of disk galaxy formation models $^{8,9,10,11,12,13}$. What has changed dramatically since then is the cosmological context in which such idea is applied, which reflects the remarkable progress that cosmology has undergone in the meantime. After two decades of active debate there is now one cosmological paradigm according to which the energy density of the Universe is dominated by cold dark matter and a cosmological constant, while ordinary baryonic matter contributes only to a few percent level (an even smaller contribution is yielded by neutrinos)$^{14}$. This model is supported by observations of the large scale mass distribution in the Universe traced by galaxies  themselves$^{15}$ and by the power spectrum of density fluctuations inferred from the cosmic microwave background radiation$^{14}$. Cold dark matter interacts only via gravity with itself and with ordinary baryonic matter, and is not subject to any dissipative force. The governing evolutionary equation for dark matter  is the collisionless Boltzmann equation that describes a zero-pressure fluid, also termed a collisionless fluid$^2$. In this model, called $\\Lambda$CDM (CDM stands for \"cold dark matter\", while $\\Lambda$ is the cosmological constant required to explain the observed acceleration of the Universe) structure forms hierarchically in a bottom-up fashion, starting from the amplification via gravitational instability of primordial small density fluctuations in the dark matter$^{16,17}$. Because of the scale-free nature of gravity and the dissipationless nature of cold dark matter, in such a model one expects the formation of self-similar, ellipsoidal collapsed objects, dark matter halos, at all scales$^{18}$. The largest halos are the last to form$^{17}$. Also, direct three-dimensional simulations of structure formation in a CDM Universe predict that halos of any mass and size should contain a swarm of smaller halos, the so-called substructure$^{19,20}$, as shown in Figure 3.   \\begin{figure} \\centering{ \\resizebox{14cm}{!}{\\includegraphics{S2T_vs_Mstar.ps}}\\\\ \\caption{The ratio between the mass of the stellar spheroid ($S$) and the sum of the mass of the stellar  spheroid and the  stellar disk ($T$) as a function of galaxy mass from a galaxy sample of the Sloan Digital Sky Surveys (SDSS)${^12}$ (2007 Blackwell Publishing Ltd). Red points with error bars show the median $S/T$ as a function of stellar mass together with the 10 and 90 percentiles of the distribution.}} \\label{fig:feedback} \\end{figure}  The model predicts quantitatively the size and mass of the dark halo of a galaxy with a given measured rotational velocity (the rotational velocity of stars and gas probes the depth of the galaxy gravitational potential well). For a galaxy like our own Milky Way, for example, its observed rotational velocity of $\\sim 220$ km/s implies a halo mass of about $10^{12}$ solar masses and a halo radius of about $300$ kpc$^{21}$ (by comparison, the disk of our galaxy contains about $6 \\times 10^{10}$ solar masses of stars and about $10^{10}$ solar masses of gas$^2$). Numerical simulations of the growth of dark matter halos also predict that the radial density profiles of such halos diverge near the center and are well described by a power-law, $\\rho \\sim r^{-\\gamma}$ ($\\rho$ being the density and $r$ the spherically averaged radius of the halo), with $\\gamma=1-1.5$ $^{22, 23, 24, 25}$. Since dark matter dominates by mass over ordinary baryonic matter, gas collapses within such halos, eventually forming a galaxy, because it is pulled inward by their gravitational attraction rather than collapsing due to its own gravity. In this scenario there is no such a thing as a galaxy forming in isolation, rather structure, both dark and baryonic, builds up via continous accretion and merging of smaller systems containing a mixture of dark and baryonic matter$^{26}$. This highly dynamical picture emerging from the current cosmology is the main difference compared to earlier attempts to study galaxy formation. Halos gain their angular momentum by tidal torques due to asymmetries in the distribution of matter, and also by acquiring the angular momentum originally stored in their relative orbit as they come together and merge into a larger system$^{27,28,29}$. Prevailing models have then assumed that baryons and dark matter start with the same specific angular momentum before the collapse begins since they are subject to the same tidal torques $^{8}$.      \\begin{figure} \\centering{ \\resizebox{16cm}{!}{\\includegraphics{fornaxlr.ps}}\\\\ \\caption{Example of a massive dark halo ($\\sim 10^{14}$ solar masses) assembled by hierarchical merging in a numerical simulation of the $\\Lambda$CDM structure formation model (simulation performed by L.Mayer and F.Governato with the GASOLINE code  $^{44}$). Many small halos (\"substructure\") are orbiting inside the larger halo and will eventually merge with it as a result of dynamical friction eroding their orbital energy and orbital angular momentum. These dark matter lumps are the sites in which gas collapses and forms galaxies.}} \\label{fig:feedback} \\end{figure}      \\subsection{Computer simulations:the angular momentum problem}  The modern tool used to study the hierarchical growth of structure driven by gravity and the concurrent collapse of baryons within dark halos is represented by three-dimensional computer simulations that solve the gravitational and hydrodynamical forces between parcels of gas and dark matter. We will discuss the methodology employed by such simulations in the next section. For now it suffices to say that in the most popular simulation method both the gas (i.e. the baryonic component) and the dark matter are represented by particles so that structures are discretized in mass and space. The evolutionary equations, such as the collisionless Boltzmann equation for cold dark matter, the Euler equation for baryonic matter (baryonic matter is treated as an ideal gas) and the Poisson equation, which holds for both, are solved for such discrete representation of physical reality. Available methods to discretize physical variables and governing equations are constructed in such a way that they should converge to the exact continuum solution for an infinite number of particles. As we will see in the next section, discretization itself, along with other aspects of the current methods, can introduce spurious effects in the computer models. Simulations take advantage of large parallel supercomputers in which hundreds of processing units are used simultaneously to compute the forces and advance the system to the next timestep.  One important prediction of simulations of a CDM Universe is that halos have a rather universal value of the angular momentum (per unit mass) at any given epoch quite irrespective of their precise mass assembly history. This is parameterized via the dimensionless spin parameter $\\lambda = J E^{1/2}/GM^{5/2}$, where $E$, $M$ and $J$ are the total energy, mass and angular momentum of the dark halo ($G$ is the gravitational constant). One can show that $\\lambda$ is proportional  to the ratio between the rotational kinetic energy and the kinetic energy in disordered motions associated with the halo. Halos have a universal distribution of spin parameters, which peaks at a value $\\lambda \\sim 0.035$ $^{29}$.  Simple spherical one-dimensional models that study disk formation in an isolated CDM halo (namely a halo that does not interact or merge with other halos) predict that  the size of disks resulting from the infall and collapse of baryons matches the size of observed disks in galaxies very well$^8$. The models use mainly two inputs, both coming from cosmological simulations, the halo density profile, which is related to the gravitational pull that drives the gas collapse, and the initial specific angular momentum of gas as implied by the typical values of the spin parameter. They further assume that angular momentum is conserved during the collapse. Hence this result is simple and remarkable at the same time; it says that CDM halos have the right amount of angular momentum to form observed disk galaxies.  For more than a decade researchers have tried to reproduce the latter result with fully three-dimensional computer simulations but have run into several problems. It was soon realized that, once the hypothesis of isolation is removed and hierarchical merging is accounted for, angular momentum can be lost by the gas to the dark matter due to a process known as dynamical friction$^{30,2,31}$. During mergers, previously collapsed clumps of gas and dark matter fall into a larger dark halo and suffer a drag force as they move through the latter. The loss of angular momentum caused by the drag force, called \"dynamical friction\", is more effective when the gas is distributed into cold and dense lumps rather than being smooth and extended$^{32}$. But gas is expected to be clumpy  in a model with collisionless cold dark matter in which collapse can occur at all scales, and large halos grow by accreting smaller halos which bring their own dense collapsed gas. As a result, early simulations$^{32}$ were obtaining improbable small disks with ten times less angular momentum than real ones. Two types of solutions for this \"angular momentum problem\" have been considered since then. The first is very drastic and calls for revising the cosmological model itself. Alternative models in which the dark matter has a non-negligible thermal velocity rather then being \"cold\" would produce collapsed systems only above a characteristic scale because the thermal jittering will tend to smear out short-wavelength density perturbations$^{33}$. These warm dark matter models (WDM) behave like CDM on large scale, thus maintaining its succesful features. The reduced clumpiness of dark matter halos in the WDM model implies that baryons are smoothly distributed rather than arranged in previously collapsed dense lumps when they fall into large galaxy-sized halos, and therefore lose less angular momentum by dynamical friction$^{33,34}$. The second, less exotic possibility, is that baryons do not just follow the merging hierarchy imposed by dark matter but somehow decouple from it  and remain much smoother. This could happen if the thermal energy content of baryons was enough to resist gravitational collapse, at least up to some critical mass scale. This way a fraction of the gas that would have entered a halo in dense clumps within smaller halos would instead enter with a smooth distribution, perhaps avoiding catastrophic loss of angular momentum. Various plausible astrophysical mechanisms can be responsible for increasing the thermal energy content of the baryons, for example the energy injection by supernovae explosions and the ambient radiation field produced by stars, accreting black holes or also external galaxies. There is, however, a third possibility. This is that the baryons clump excessively in computer simulations because the numerical methods adopted can introduce artificial loss of angular momentum. As we argue in this report, a solution lies probably in a combination of the latter two proposals, with no need of revising the standard cosmological structure formation model.  \\begin{figure} \\centering{ \\resizebox{18cm}{!}{\\includegraphics{tobiaslr.ps}}\\\\ \\caption{ Disk size as a function of mass resolution in numerical simulations of disk formation in an isolated CDM halo$^{48}$ (2007 Blackwell Publishing Ltd). The three panels show density maps of gas in a slice through the centre of the gaseous galactic disk after 5 Gyr;  the gas mass resolution decreases from the left to the right by about a factor 8 each snapshot (the maximum resolution is 1 million gas particles). The box side length is 20 kpc for every panel. At all resolutions disks show asymmetries such as central bar-like structures and spiral arms. The bar, however, is a strong and long-lived feature only for sufficiently high spatial resolution (set by the gravitational softening)$^{48}$.}} \\label{fig:feedback} \\end{figure}    The next two sections will be devoted, respectively, to the role of numerical effects in disk formation simulations, and to the modeling of gas thermodynamics and star formation in the simulations. We will then show how the structure of simulated disks is affected by different models of thermodynamics and star formation. Finally, we will summarize the current status of the field and  the major problems that remain to be solved, including the puzzling origin of disk galaxies without a bulge. We will attempt to recall the most important contributions by the various groups actively involved in this field of research while at the same time covering in more detail some recent results of the research group to which the authors of this report belong. ",
        "conclusions": ""
    },
    "0801/0801.1071_arXiv.txt": {
        "abstract": "We summarize and discuss recent work (Fregeau 2007) that presents the confluence of three results suggesting that most Galactic globular clusters are still in the process of core contraction, and have not yet reached the thermal equilibrium phase driven by binary scattering interactions: that 1) the three clusters that appear to be overabundant in X-ray binaries per unit encounter frequency are observationally classified as ``core-collapsed,'' 2) recent numerical simulations of cluster evolution with primordial binaries show that structural parameters of clusters in the binary-burning phase agree only with ``core-collapsed'' clusters, and 3) a cluster in the binary-burning phase for the last few Gyr should have $\\sim 5$ times more dynamically formed X-ray sources than if it were in the core contraction phase for the same time. ",
        "introduction": "This proceedings article briefly summarizes my contributed talk at the ``Population Explosion'' meeting, which itself was a brief summary of the work described in \\citet{fregeau07}.  I have taken the opportunity here to include supplemental material, but first refer the reader to the detailed, complete discussion in \\citet{fregeau07}.  \\subsection{X-Ray Sources and Cluster Dynamics} Although X-ray binaries have been known to be overabundant in Galactic globular clusters (by unit mass relative to the disk) for over 30 years \\citep{1975ApJ...199L.143C,1975Natur.253..698K}, it is only recently that clear evidence for their dynamical origin has revealed itself.  \\citet{2003ApJ...591L.131P} found that the number of X-ray sources with $L_X > 4 \\times 10^{30}\\,{\\rm erg}/{\\rm s}$, $N_X$, in a cluster follows a clear, nearly linear correlation with the encounter frequency $\\Gamma$, a measure of the dynamical interaction rate in the cluster.  There are a few clear outliers to this trend, however.  Terzan 1 is overabundant in $N_X$ by a factor of $\\sim 20$, NGC 6397 by a factor of $\\sim 5$, and NGC 7099 by a factor of $\\sim 2$ \\citep{2006MNRAS.369..407C,2007ApJ...657..286L}. The common thread among these clusters is that they are all classified observationally as ``core-collapsed,'' while the other clusters in the sample are not.  A cluster is observationally termed core-collapsed if its surface brightness profile is consistent with a cusp at the limit of observational resolution.  Since this observational classification is linked with the cluster evolutionary state (as described below), it appears that cluster evolution complicates the $N_X$--$\\Gamma$ correlation.  \\subsection{Understanding Cluster Core Radii} The evolution of a globular cluster, being a bound self-gravitating system, is very similar to that of a star, and comprises three main phases. The cluster first ``core contracts'' on a relaxation timescale.  Once the core density becomes large enough for binary scattering interactions to begin generating energy, the cluster settles into the ``binary-burning'' phase, in which the cluster's core properties remain roughly constant with time. Once the binary population is exhausted in the core, it collapses via the gravothermal instability, leading to extremely high central densities, followed by a series of gravothermal oscillations in which the core expands and contracts repeatedly.  (For a graphical representation of the three main phases of cluster evolution, see Figure 1 of \\citet{1991ApJ...370..567G}, Figure 5 of \\citet{2003ApJ...593..772F}, or Figure 29.1 of \\citet{2003gmbp.book.....H}.) Recent comparison of star cluster evolution simulations with observations of cluster structural parameters have shown that the clusters observationally classified as core-collapsed are in fact most likely in the binary-burning phase, while the rest are most likely still in the core-contraction phase \\citep{2006MNRAS.368..677H,2007ApJ...658.1047F}.  \\subsection{A Refined $\\Gamma$} First introduced by \\citet{1987IAUS..125..187V}, the encounter frequency $\\Gamma$ is an estimate of the {\\em current} dynamical interaction rate in the cluster, which is assumed to be proportional to the current number of observable X-ray sources.  Typically, it is approximated as \\begin{equation}\\label{eq:standardgamma} \\Gamma \\equiv \\frac{dN_{\\rm int}}{dt} \\propto \\rho_c^2 r_c^3 / v_\\sigma \\, , \\end{equation} where $r_c$ is the cluster core radius, $\\rho_c$ is the core mass density, and $v_\\sigma$ is the core velocity dispersion.  More accurate approximations of the interaction rate have been used, including numerical integrals over cluster models \\citep{2003ApJ...591L.131P}, but all are estimates of the {\\em current} interaction rate.  Recent work has shown that dynamically formed X-ray binaries have finite detectable lifetimes ($\\sim 10^5$ to $\\sim 10^9$ yr, depending on binary type), and furthermore that there is a lag time of several Gyr between dynamical interaction and the binary turning on as an X-ray source \\citep{2006MNRAS.372.1043I,2007arXiv0706.4096I}. Since the binary interaction rate is a function of the cluster properties (density, velocity dispersion, etc.), it is clear that the X-ray binaries we see in clusters today were formed several Gyr ago, and thus encode the recent cluster evolution history in their populations.  We thus adopt a refined version of $\\Gamma$ which encodes this history by integrating the interaction rate over the dynamical evolution of the cluster. The resulting quantity is the total number of strong dynamical interactions, which should be roughly proportional to the currently observable number of X-ray binaries.  Plugging in the different phases of cluster evolution will yield different estimates for the number of sources, and possibly allow one to differentiate among the different phases by comparing with the observed number of sources.  Using standard assumptions, and adopting some results from recent $N$-body simulations for the evolution of $r_c$ with time, one finds the ratio of the number of interactions for a cluster in the binary-burning phase to one in the core contraction phase is \\begin{equation} \\frac{N_{\\rm int,bb}}{N_{\\rm int,cc}} = \\frac{t_x}{t_0} \\frac{2.835}{\\left(\\frac{t_0}{t_0+9t_\\ell}\\right)^{0.315}- \\left(\\frac{t_0}{t_0+9t_\\ell+9t_x}\\right)^{0.315}} \\, , \\end{equation} where $t_x$ is the X-ray source lifetime, $t_\\ell$ is the lag time between strong dynamical interaction and X-ray turn-on, and $t_0$ is the current cluster age.  (The gravothermal oscillation phase is excluded from the discussion since in this phase the core binary fraction would be much smaller than what is observed.) This expression has a minimum of $2.0$ and a maximum of $17.8$ in the range $t_x=10^{-4}$--$3\\,{\\rm Gyr}$, $t_\\ell=1$--$10\\,{\\rm Gyr}$, for $t_0=13\\,{\\rm Gyr}$. For the canonical values of $t_x=1\\,{\\rm Gyr}$ and $t_\\ell=3\\,{\\rm Gyr}$ with $t_0=13\\,{\\rm Gyr}$, the value is $5.0$.  Since the number of X-ray sources should scale roughly linearly with the number of interactions, this suggests that if a cluster is in the binary-burning phase (and has been for a time $t_\\ell+t_x$ to the present), it should have $\\sim 5$ times as many X-ray sources than it would if it were in the core contraction phase. The three observationally core-collapsed clusters in the sample previously mentioned are overabundant in $N_X$ by factors compatible with the range of values allowed by this expression, suggesting that the observationally core-collapsed clusters are indeed in the binary-burning phase, while the rest are still in the process of core contraction.  \\subsection{Discussion} Since only $\\sim 20$\\% of Galactic globular clusters are observationally classified as core-collapsed, the conclusion that seems strongly suggested is that most clusters ($\\sim 80$\\%) are currently still in the core contraction phase, while those that are core-collapsed are in the binary burning phase.  This goes counter to the widely held belief that most clusters are currently in the binary burning phase, and complicates the many existing studies that have assumed cluster core properties that are constant with time, including predictions of blue straggler populations, tidal-capture binaries, and the evolution of the core binary fraction \\citep[e.g.,][]{2004ApJ...605L..29M,1994ApJ...423..274D,2005MNRAS.358..572I}. Similarly, this result does away with the need for alternate energy sources in cluster cores to explain currently observed core radii, including intermediate-mass black holes in most clusters, mass loss from stellar mergers, and ongoing mass segregation \\citep{2006astro.ph.12040T,chatterjeeposter,2004ApJ...608L..25M}. ",
        "conclusions": "A few interesting aspects of this analysis did not fit into the original {\\em letter}. It has been pointed out that an important but neglected factor going into the interaction rate is the neutron star retention fraction, since it depends on the cluster escape speed, and since it is neutron star $+$ stellar binary interactions that predominantly lead to X-ray binaries (Rasio, private communication). Since there is no reason to expect that the cluster escape speed at the time of neutron star formation does not vary from cluster to cluster, the correlation between dynamical interaction rate and $N_X$ should be blurred significantly by this effect.  We note that the analysis in \\citet{fregeau07} avoids this issue entirely by comparing a cluster only with a version of itself in a different evolutionary phase.  Another issue is that there is apparently a significant contribution to the X-ray binary population from primordial binaries in clusters with small encounter frequencies.  (This is likely the reason for the sub-linear dependence of $N_X$ on $\\Gamma$ in \\citet{2003ApJ...591L.131P}.)  This has given rise to the use of the interaction rate per unit cluster mass, $\\gamma \\equiv \\Gamma/M_{\\rm clus}$, to help isolate the dynamically-formed population in the analysis. The use of $\\gamma$ has, for example, shown that a significant fraction of cataclysmic variables are dynamically formed \\citep{2006ApJ...646L.143P}.  The results of \\citet{fregeau07} could be further tested in the future with the more sensitive $\\gamma$ as the library of deep X-ray observations of globular clusters grows.  Finally, we mention that in the near future we plan to test our simple overabundance prediction by using the detailed numerical models of \\citet{2005MNRAS.358..572I}. Although there are some hints from that paper that inputting a time-varying core does not noticeably affect the final core binary fraction, we expect that the number of visible X-ray binaries, while noisy, may be more noticeably affected."
    },
    "0801/0801.4462_arXiv.txt": {
        "abstract": "Since GRBs fade rapidly, it is important to publish accurate, precise positions at early times. For \\Swift-detected bursts, the best promptly available position is most commonly the X-ray Telescope (XRT) position. We present two processes, developed by the \\Swift\\ team at Leicester, which are now routinely used to improve the precision and accuracy of the XRT positions reported by the \\Swift\\ team. Both methods, which are fully automated, make use of a PSF-fitting approach which accounts for the bad columns on the CCD. The first method yields positions with 90\\% error radii $<$4.4\" 90\\% of the time, within 10--20 minutes of the trigger. The second method astrometrically corrects the position using UVOT field stars and the known mapping between the XRT and UVOT detectors, yielding enhanced positions with 90\\% error radii of $<$2.8\" 90\\% of the time, usually ~2 hours after the trigger. ",
        "introduction": "For the majority of \\Swift-detected GRBs, the best promptly available position is that of the X-ray telescope (XRT, \\cite{Burrows05}). It is thus desirable to reduce the 3.5\" boresight uncertainty associated with this instrument.  We have developed two techniques to achieve this goal. The first is a fitting technique which accounts for hot columns on the X-ray CCD. This is described in Section~1 and the application of this to promptly available data is detailed in Section~2. The second technique is applied to the full ground dataset, and uses the field stars in the UV/Optical telescope (UVOT, \\cite{Roming05}) to astrometrically correct the XRT position, eliminating the XRT's boresight uncertainty. This is described in Section~3. In Fig.~\\ref{fig:errdist} we show the distribution of position uncertainties produced by these techniques, comparing them with positions determined onboard the XRT, and the `refined' positions produced from the full dataset without astrometric correction. Finally, in Section~4 we discuss forthcoming improvements to the second technique, and the potential for applying it to the prompt data. For an overview of the different positions available from the \\Swift\\ XRT, see \\url{http://www.swift.ac.uk/xrt_pos.php}  \\begin{figure} \\includegraphics[height=8cm,angle=-90]{errdist} \\includegraphics[height=8cm,angle=-90]{errcheck} \\caption{\\emph{Left:} Distribution of the 90\\% confidence error radii produced for XRT positions of GRBs. Distributions shows are those obtained onboard automatically (solid), from SPER data (Section 2, dotted), on the ground from the full dataset (short dashes) and the enhanced positions (Section 3, long dashes).\\emph{Right:} Distribution of the offsets of the UVOT-enhanced positions from the UVOT position for GRBs with both, divided by the position error. As can be seen, 90\\% of the enhanced positions agree with the UVOT positions, confirming the error circle is correctly calibrated.} \\label{fig:errdist} \\end{figure} ",
        "conclusions": ""
    },
    "0801/0801.0594_arXiv.txt": {
        "abstract": "{} {We present $J$, $H$, $K$ spectrally dispersed interferometry with a spectral resolution of 35 for the Mira variable S~Orionis. We aim at measuring the diameter variation as a function of wavelength that is expected due to molecular layers lying above the continuum-forming photosphere. Our final goal is a better understanding of the pulsating atmosphere and its role in the mass-loss process.} {Visibility data of S~Ori were obtained at phase 0.78 with the VLTI/AMBER instrument using the fringe tracker FINITO at 29 spectral channels between 1.29\\,$\\mu$m and 2.32\\,$\\mu$m. Apparent uniform disk (UD) diameters were computed for each spectral channel. In addition, the visibility data were directly compared to predictions by recent self-excited dynamic model atmospheres.} {S~Ori shows significant variations in the visibility values as a function of spectral channel that can only be described by a clear variation in the apparent angular size with wavelength. The closure phase values are close to zero at all spectral channels, indicating the absence of asymmetric intensity features. The apparent UD angular diameter is smallest at about 1.3\\,$\\mu$m and 1.7\\,$\\mu$m and increases by a factor of $\\sim$\\,1.4 around 2.0\\,$\\mu$m. The minimum UD angular diameter at near-continuum wavelengths is $\\Theta_\\mathrm{UD}$=8.1$\\pm$0.5\\,mas, corresponding to $R\\sim$\\,420\\,R$_\\odot$. The S~Ori visibility data and the apparent UD variations can be explained reasonably well by a dynamic atmosphere model that includes molecular layers, particularly water vapor and CO. The best-fitting photospheric angular diameter of the model atmosphere is $\\Theta_\\mathrm{Phot}$=8.3$\\pm$0.2\\,mas, consistent with the UD diameter measured at near-continuum wavelengths.} {The measured visibility and UD diameter variations with wavelength resemble and generally confirm the predictions by recent dynamic model atmospheres. These size variations with wavelength can be understood as the effects from water vapor and CO layers lying above the continuum-forming photosphere. The major remaining differences between observations and model prediction are very likely due to an imperfect match of the phase and cycle combination between observation and available models.} ",
        "introduction": "Mira stars are low-mass, large-amplitude, long-period variable stars on the AGB, evolving toward the planetary nebula and white dwarf phases. They exhibit a mass-loss rate on the order of $\\sim$10$^{-6}$\\,M$_\\odot$/year that significantly affects the further stellar evolution and is one of the most important sources for the chemical enrichment of the interstellar medium. The dust condensation sequence, the wind-driving mechanism, and the role of pulsation are currently not well understood, in particular for oxygen-rich AGB stars (Woitke et al. \\cite{woitke06}; H\\\"ofner \\& Andersen \\cite{hoefner07}). The pulsating atmospheres of Mira stars can become very extended because of dynamic effects including shock fronts, and they are very cool in their outer parts. Here, molecules can form, which for O-rich stars are most importantly H$_2$O, CO, TiO, and SiO (Tsuji et al. \\cite{tsuji97}; Tej et al. \\cite{tej03}; Ohnaka \\cite{ohnaka04}). Wittkowski et al. (\\cite{wittkowski07}) found for the case of S~Ori that Al$_2$O$_3$ dust condenses within the extended atmosphere at phase-dependent distances of 1.8--2.4 photospheric radii. This extended atmosphere, which is characterized by phase-dependent temperature and density stratifications, the presence of molecular layers, and the formation of dust, is thus of particular interest for our better understanding of pulsation and mass loss. Observed radii of Mira stars have been found to differ for different optical and infrared bandpasses (e.g. Thompson et al. \\cite{thompson02}; Mennesson et al. \\cite{mennesson02}; Ireland et al. \\cite{ireland04a}; Perrin et al. \\cite{perrin04}; Eisner et al. \\cite{eisner07}), and this has been attributed to the presence of molecular layers located above the continuum-forming photosphere. Here, we present both a spectro-interferometric observation of the Mira star S~Ori that covers the near-infrared $J$, $H$, and $K$ bands simultaneously at a spectral resolution of 35 and a comparison to recent self-excited dynamic model atmospheres.  S~Ori is a Mira variable star with spectral type M6.5e--M9.5e and $V$ magnitude 7.2--14.0 (Samus et al. \\cite{samus04}). We use a Julian Date of last maximum brightness $T_0=2453190\\ \\mathrm{days}$, a period $P=430\\ \\mathrm{days}$ (as in Wittkowski et al. \\cite{wittkowski07}), and the distance of $480\\ \\mathrm{pc}\\ \\pm\\ 120\\ \\mathrm{pc}$ from van Belle et al. (\\cite{vanbelle02}). The broadband near-infrared $K$ UD angular diameter of S~Ori has been measured by van Belle et al. (\\cite{vanbelle96}), Millan-Gabet et al. (\\cite{millan05}), and Boboltz \\& Wittkowski (\\cite{boboltz05}) to values between 9.6\\,mas and 10.5\\,mas at different phases. Joint VLTI/MIDI and VLBA/SiO maser observations by Wittkowski et al. (\\cite{wittkowski07}) have shown that S~Ori exhibits significant phase-dependencies of the atmospheric extension and dust shell parameters with photospheric angular diameters between 7.9\\,mas and 9.7\\,mas. \\begin{table*} \\caption{Observation log. Night starting 12 October 2007, JD 2454386.} \\label{tab:obslog} \\begin{tabular}{lllllllrrrr} \\hline\\hline Target & Purpose & $\\Theta_\\mathrm{LD}$& DIT & Time & $\\Phi_\\mathrm{Vis}$ & $B_p$ [m] & PA$_p$ & AM & Seeing & $\\tau_0$\\\\ &  & [mas] &[msec] &[UTC] &     & E0-G0/G0-H0/E0-H0   &$\\deg$  && [$\\arcsec$]&[msec]\\\\\\hline \\object{45~Eri} & Calibrator (K3 II-III) &2.15$\\pm$0.04& 25 & 08:03-08:07 & &16.0/31.9/47.9&-107&1.1&1.3&1.3\\\\ 45~Eri & Calibrator (K3 II-III) &2.15$\\pm$0.04& 50 & 08:09-08:12 & &16.0/32.0/48.0&-107&1.1&1.3&1.3\\\\ \\object{$\\gamma$~Eri}& Check star (M0.5 IIIb)&8.74$\\pm$0.09& 25 & 08:22-08:26 & &15.6/31.2/46.8&-104&1.1&1.4&1.2\\\\ $\\gamma$~Eri& Check star (M0.5 IIIb)&8.74$\\pm$0.09& 50 & 08:28-08:32 & &15.5/31.0/46.5&-104&1.1&1.4&1.2\\\\ S~Ori&Science target&& 25 & 08:45-08:49 & 0.78 &15.9/31.8/47.8&-107&1.1&1.3&1.3\\\\ S~Ori&Science target&& 50 & 08:52-08:57 & 0.78 &16.0/31.9/47.9&-107&1.1&1.3&1.3\\\\ $\\alpha$~Hor & Calibrator (K2 III) &2.76$\\pm$0.03& 25 & 09:14-09:18 &&14.5/28.9/43.3&-91&1.1&1.2&1.4\\\\ $\\alpha$~Hor & Calibrator (K2 III) &2.76$\\pm$0.03& 50 & 09:21-09-24 &&14.3/28.7/43.0&-90&1.1&1.4&1.2\\\\ S~Ori & Science target && 25 & 09:36-09:40 & 0.78 &15.9/31.8/47.7&-106&1.1&1.5&1.1\\\\\\hline \\end{tabular} \\end{table*} ",
        "conclusions": ""
    },
    "0801/0801.2558_arXiv.txt": {
        "abstract": "I review measurements of star formation in nearby galaxies in the UV--to--FIR wavelength range, and discuss their impact on SFR determinations in intermediate and high redshift galaxy populations. Existing and upcoming facilities will enable precise cross-calibrations among the various indicators, thus bringing them onto a common scale. ",
        "introduction": "Determinations of star formation rates (SFRs) in galaxies utilize indicators at a variety of wavelengths, from the X--ray to the radio. Many indicators have been defined in response to specific needs. For instance, when new populations of galaxies are discovered using a new wavelength window or improved observing techniques/instruments in a certain waveband, there is a push to investigate whether that waveband can be used to derive SFRs as well, and/or to define the uncertainties and limitations of doing so.  The advent of new facilities (e.g., Herschel, LMT, ALMA, JWST, etc., and the many ground--based telescopes under construction or design) together with existing ones (HST, Spitzer, Chandra, and the vast array of existing ground--based facilities) will cover extensively the electromagnetic spectrum at unprecedented sensitivities. This will offer the opportunity to cross-calibrate many of the SFR indicators across a range of redshifts, and, therefore, on many galaxy populations at various stages of their evolution.  In this brief review, I discuss the current status and the known limitations of SFR indicators in a few wavelength regimes: ultraviolet (UV), optical/near--infrared, and mid/far--Infrared (MIR/FIR). ",
        "conclusions": ""
    },
    "0801/0801.2591_arXiv.txt": {
        "abstract": "The discovery of over 200 extrasolar planets with the radial velocity (RV) technique has revealed that many giant planets have large eccentricities, in striking contrast with most of the planets in the solar system and prior theories of planet formation.  The realization that many giant planets have large eccentricities raises a fundamental question: ``Do terrestrial-size planets of other stars typically have significantly eccentric orbits or nearly circular orbits like the Earth?'' Here, we demonstrate that photometric observations of transiting planets could be used to characterize the orbital eccentricities for individual transiting planets, as well the eccentricity distribution for various populations of transiting planets (e.g., those with a certain range of orbital periods or physical sizes).  Such characterizations can provide valuable constraints on theories for the excitation of eccentricities and tidal dissipation.  We outline the future prospects of the technique given the exciting prospects for future transit searches, such as those to be carried out by the CoRoT and Kepler missions. ",
        "introduction": "\\label{sec_intro}  Theorists have proposed numerous mechanisms that could excite orbital eccentricities.  Some of these mechanisms are expected to affect all planets independent of their mass (e.g., perturbations by binary companions, passing stars, or stellar jets; e.g., Holman et al.\\ 1997; Laughlin \\& Adams 1998; Ford et al.\\ 2000; Zakamska \\& Tremaine 2004; Namouni 2007), while the efficiency of other mechanisms would vary with planet mass (e.g., planet-disk or planet-planet interactions; e.g., Artymowicz 1992; Goldreich \\& Sari 2003; Chatterjee et al.\\ 2007).  If the mechanism(s) exciting eccentricities of the known giant planets also affect terrestrial planets, then Earth-mass planets on nearly circular orbits could be quite rare.  On the other hand, if large eccentricities are common only in systems with massive giant planets and/or very massive disks, then there may be an abundance of planetary systems with terrestrial planets on low eccentricity orbits (Beer et al.\\ 2004).  Thus, understanding the eccentricity distribution of terrestrial planets could provide empirical constraints for planet formation theories (e.g., Ford \\& Rasio 2007) and shed light on the processes that determined the eccentricity evolution in our solar system.  Since the discovery of transiting giant planets with eccentric orbits, authors have begun to consider the implications of eccentricities for transiting planets (e.g., Barnes 2007; Burke 2008).  The CoRoT and Kepler space missions are expected to detect many transiting planets and measure their sizes and orbital periods, including some in or near the ``habitable zone'' (e.g., Kasting et al.\\ 1993).  The Kepler mission aims to determine the frequency of Earth-like planets, and study how the frequency and properties of planets correlates with the properties of their host stars (Basri et al.\\ 2005). Since a significant eccentricity would cause the stellar flux incident on the planet's surface to vary, a planet's eccentricity affects its climate (i.e., equilibrium temperature, amplitude of seasonal variability) and potentially its habitability (Williams et al.\\ 2002; Gaidos \\& Williams 2004).  Our method could be applied to these planets to determine the frequency of terrestrial planets that could also have Earth-like climates, and thus influence the design of future space missions that will attempt to detect and characterize nearly Earth-like planets (e.g., Space Interferometry Mission-PlanetQuest) and search them for signs of life (e.g., Terrestrial Planet Finder).   In \\S\\ref{SecDuration}, we show how the duration of a transit is affected by a planet's orbital eccentricity.  We outline how to interpret the transit duration for transiting planets with both low signal-to-noise light curves (\\S\\ref{SecLoSN}) and high signal-to-noise (\\S\\ref{SecHiSN}).  We compare the magnitude of the effect on the transit duration to the expected precision of eccentricity constraints based on Kepler photometric data (\\S\\ref{SecStatErr}) and also the typical accuracy of stellar parameters (\\S\\ref{SecStarErr}). We demonstrate that Kepler observations could be used to characterize the orbital eccentricities of terrestrial planets.  We describe statistical approaches for analyzing the distribution of transit durations of a population of transiting planets in \\S\\ref{SecStats}. In \\S\\ref{SecRV}, we discuss the role of radial velocity observations for constraining eccentricities of Earth-like transiting planets.  In \\S\\ref{SecDiscuss}, we conclude with a discussion of how the results could contribute to understanding the formation and evolution of terrestrial planets and address fundamental questions, such as ``What is the frequency of terrestrial planets that have Earth-like eccentricities?'' and ``What is the frequency of terrestrial planets that pass through the habitable zone?'' ",
        "conclusions": "\\label{SecDiscuss} We have described how photometric observations can constrain the eccentricities of individual transiting planets and characterize the eccentricity distribution of a population of planets.  For each planet, we can test the null hypothesis that it is on a circular orbit and calculate the posterior probability distribution for the eccentricity.  A combination of such analyses for several transiting planets could be used to characterize the eccentricity distribution of a population of planets.  For example, this method could be used to investigate how the fraction of eccentric orbits varies with the orbital period or physical proprieties of the star and planet. We expect this type of analysis to become increasingly powerful given the rapidly growing number of known transiting planets.  This type of analysis will be particularly valuable for low-mass transiting planets or transiting planets around faint host stars. In both cases, radial velocity follow-up observations will be extremely challenging.  This will be the case for many transiting planet candidates found by space-based transit searches, such as CoRoT and Kepler.  The capability of these missions to discover terrestrial-mass planets is particularly exciting.  Our method could could determine if terrestrial planets with low eccentricities like the Earth are common or rare.  This would provide significant constraints on theories proposed to explain the eccentricities of extrasolar planets. For example, Kepler might find many terrestrial planets with low eccentricity orbits, suggesting that the mechanisms that excite the eccentricities of giant planets are often ineffective for terrestrial mass planets.  In this scenario, those terrestrial planets that do have large eccentricities might typically be accompanied by nearby giant planets, suggesting that it is the giant planets are responsible for exciting the eccentricities of terrestrial planets (Veras \\& Armitage 2005, 2006).  Alternatively, Kepler might find that terrestrial planets are much more common than giant planets, and yet they still commonly have large eccentricities. This could arise from eccentricity excitation mechanisms that do not require giant planets, or due to interactions with previous giant planets that have since been ejected, accreted or destroyed by the star (e.g., Ford et al.\\ 2005; Raymond et al.\\ 2006; Mandell et al.\\ 2007).  Our method also provides a means for studying the tidal evolution of short-period planets.  Tidal effects are likely to circularize planets with sufficiently short orbital periods.  For short-period planets, we can compute the tidal circularization timescale ($t_{\\rm circ}$) based on the properties of the star and planet.  If there is a sharp transition between circular and eccentric orbits, then this could be used to place constraints on tidal theory (e.g., Zahn \\& Bouchet 1989; Melo et al.\\ 2001; Mathieu et al.\\ 2004).  For eccentric planets with relatively short $t_{\\rm circ}$, it may be possible to place a lower limit on the $Q$ factor the is related to the planet's internal structure (e.g., Ford et al.\\ 1999; Bodenheimber et al.\\ 2001; Maness et al.\\ 2007; Mardling 2007).  In a Bayesian framework, one can calculate the posterior probability distribution for the fraction of each orbit during which the planet-star separation is between an inner and outer cut-off.  If the cut-offs are set to be the putative boundary of the habitable zone, then we can then ask, ``What fraction of terrestrial planets are in the habitable zone for some/at least half/all of their orbit?''.  The results of such investigations could have implications for the climates of potentially habitable planets, the frequency of such planets, and the design of future missions that aim to detect and characterize nearby Earth-like planets that could harbor life (e.g., Marcy et al.\\ 2005).  We have demonstrated that it is practical to collect sufficient photons to characterize the eccentricity distribution of terrestrial extrasolar planets, but we assumed that limb-darkening parameters can be well constrained by some combination of stellar modeling and external observations.  Our analytical estimates have not incorporated limb darkening effects or potential systematic uncertainties in stellar models.  Future search should address both of these effects. Multi-wavelength observations (particularly in the infrared) could be particularly useful for addressing both these issues.  In particular, we plan to investigate the potential for combinations of space-based detections and ground-based follow-up observations to improve the characterization of the eccentricities of transiting planets.  Finally, we note that this method for characterizing the eccentricities of terrestrial planets from transit light curves underscores the importance of developing and validating precise and accurate stellar models.  Uncertainties in stellar parameters models are expected to dominate the error budget for bright target stars. The potential for systematic uncertainties due to stellar modeling will make it particularly challenging to study low eccentricity systems as a function of stellar properties.  Fortunately, we these concerns would not preclude our method from being applied to terrestrial-sized planets recognizing the relatively large eccentricities typical for giant planets."
    },
    "0801/0801.0846_arXiv.txt": {
        "abstract": "We consider the bound states of the massive scalar field around a rotating black hole immersed in the asymptotically uniform magnetic field. In the regime of slow black hole rotation, the Klein-Gordon equation allows separation of variables. We show that the growth rate of the instability can be amplified a few times by the magnetic field. The effect occurs because the magnetic field adds the \"effective mass\" term $B |m|$ to the scalar field potential for a Kerr black hole. In addition, and as a by-product, we discuss the behavior of the quasinormal modes for the magnetized rotating black holes. ",
        "introduction": "Interaction of black holes and magnetic fields is an important factor for large astrophysical black holes, and, especially, for supermassive galactic black holes because of enormous magnetic fields in the nucleus of galaxies and near supermassive black holes. Strong magnetic fields are induced also by accretion disks of rotating charged matter near black holes. Yet, interaction of strong magnetic fields and black holes is, apparently, not limited just by the context of large scale astrophysical systems. Recent development in the  brane world theories opens possibility of observing quantum gravity at the Tev scale, so that mini black holes might be created at particle collisions in the Large Hadron Collider. Observing black holes in laboratories could make it possible to test interaction of black holes with external fields. After all, in the early universe strong magnetic fields could effect primordial black holes.   Sufficiently strong magnetic field near a black hole deforms the black hole space-time, so that one cannot consider the magnetic field as a test field on the black hole background, but a consistent solution of the Einstein-Maxwell equations is necessary instead. Fortunately, exact solutions for such a situation exist for both non-rotating (the Ernst solution) \\cite{8} and rotating cases (the Diaz solution) \\cite{9}.   Yet, the essential question is whether such a strong magnetic field, that could deform the black hole space-time, can exist, and, whether the Ernst and Diaz solutions can be considered as at least some approximate approach to a realistic situation?. The Ernst solution is described by the metric \\begin{equation} d s^{2} = \\Lambda^{2} \\left( \\left(1- \\frac{2 M}{r} \\right) d t^{2} - \\left(1- \\frac{2 M}{r} \\right)^{-1} d r^{2} -r^{2} d \\theta^{2} \\right)  - \\frac{r^{2} \\sin^{2} \\theta}{\\Lambda^{2} } d \\phi^{2}, \\end{equation} where the external magnetic field is determined by the parameter $B$, \\begin{equation} \\Lambda = 1 + \\frac{1}{4} B^{2} r^{2} \\sin^{2} \\theta, \\end{equation} and the \"unit\" magnetic field measured in $Gs$ is \\begin{equation} B_{M} = 1/M = 2.4 \\times 10^{19} \\frac{M_{Sun}}{M}. \\end{equation}  The magnetic field in the Ernst solution is poloidal and homogeneous far from the black hole.  This is not a big restriction for our consideration. Indeed, although the large scale magnetic field in the observable universe has both toroidal and poloidal components, it is the poloidal component that dominate in centers of galaxies and thereby in the region near a super-massive black hole. For mini black holes, which according to the brane-world scenarios, are expected to be observed in particles collisions, the homogeneous (poloidal) magnetic field is natural. Even though the super-strong accelerating magnetic field is screened in the region of particles collisions, we can assume, at least in principle, that if mini black holes are created in a laboratory, one could \"immerse\" mini-black holes in the magnetic field with required properties.  In order to make estimations of possible influence of the magnetic field on the supermassive black holes we need the two parameters at hand: the magnetic field parameter $B$ and the mass of the black hole $M$. Modern observations suggest that super-massive black holes in centers of galaxies can have mass $M \\approx 10^6 - 10^9 M_{Sun}$, while observations of the magnetic field are much more complicated and lead sometimes to controversial predictions for the value of the magnetic field in the centers of galaxies \\cite{Zakharov:2002cf}, \\cite{Contra}. At the same time, observations of the Active Galactic Nuclei (AGN) and micro- quasars indicate the existence of wide X-ray emission of lines of heavy ionized elements in their spectra \\cite{Zakharov:2002cf}. This effect is very well described when supposing existence of the super-strong magnetic field $B \\approx 10^{10}- 10^{11}$ $Gs$ \\cite{Zakharov:2002cf}. In addition some other theories \\cite{Tolubaev} imply existence of a strong magnetic field in the AGN. From these data for $M$ and $B$ at hand, by formula (2), one can easily see that for the heaviest galactic black holes, the magnetic field can be of order of fractions  of a \"unity\" ($1/M$), thereby effecting the black hole metric itself.   The Diaz and Ernst solutions, which we shall analyze here, describe the black holes immersed in a magnetic field decaying near the black hole to some asymptotic value, so that far from black hole the magnetic field is uniform and represents the Melvin Universe. Different effects around such black holes have been studied in \\cite{10} and some generalizations were obtained in \\cite{11}. Potentially observable effects, such as quasinormal modes and gravitational lensing were considered in \\cite{12}, \\cite{13},  \\cite{14} for non-rotating magnetized black holes. In this paper, we shall consider scattering of scalar waves around rotating magnetized black holes. An incident wave near the black hole will be partially absorbed (tunneled through the potential barrier) and partially radiated away as a response of a black hole to an external perturbation. When a black hole of mass $M$ and radius $r_{+}$ is rotating, and the real oscillation frequency of the perturbation (for a mode with an azimuthal number $m$) satisfies the inequality \\begin{equation} \\omega < \\frac{a m}{2 M r_{+}}, \\end{equation} then the energy radiated away exceeds the energy of the incident wave. This is the well known effect of super-radiance predicted by Zeldovich and Misner \\cite{15} and computed for Kerr black holes by Starobinsky \\cite{16}. Press and Teukolsky suggested that if the wave radiated away will be reflected by a mirror surrounding a black hole, one could make a kind of \"black hole bomb\", because initial small perturbations would grow without bound \\cite{17}. The role of the mirror can play the potential well, another local minimum, which appears because of the non-vanishing massive term \\cite{18}.   In the limit $M \\mu <<1$, where $M$ is the black hole mass and $\\mu$ is the inverse Compton wavelength of the particle, Detweiler showed that the massive scalar field exerts the superradiant instability with the  maximal growth rate \\cite{20}, \\cite{21} (at $a=M$) given by the formula \\begin{equation} \\gamma = \\frac{1}{24} \\frac{a}{M} (\\mu M)^{9} (G M/c^{3})^{-1}. \\end{equation} This instability has very small growth rate and probably cannot be significant for real black holes. Accurate calculations of instability made recently in \\cite{18_1} allowed to observe the maximal growth \\begin{equation} \\gamma = 1.5 10^{-7} (G M/c^{3})^{-1}. \\end{equation} at $M \\mu =0.42$. Yet, due-to short lifetimes or large particle masses, this instability is tiny for any known massive boson particles, except possibly, the neutral pion $\\pi^{0}$ when $M \\sim 10^{12} kg.$  In this paper we shall consider the bound states of massive scalar field in the vicinity of the magnetized black holes and analyze the corresponding superradiant instability. As a by-product we shall consider the $B$-corrections to the orbital geodesic motion (as that for which the instability is important) of particles around the Ernst black hole. The paper is organized as follows. In Sec. II the massive scalar field equation is decoupled for the regime of \"small\" (in comparison with a \"unity\" $1/M$) magnetic field, and is reduced to the wave-like equation with an effective potential. Sec. III considers the bound states and their instability in the limit $\\mu M \\ll 1$. In particular, we obtain the generalization of the well-known Detweiler formula for the  superradiant instability for a non-zero external magnetic field. In Sec. IV we briefly discuss the quasinormal modes of rotating magnetized black holes. In the Conclusion we summarize the obtained results and outline the number of open questions.   \\vspace{4mm} ",
        "conclusions": "We have separated variables for the Klein-Gordon and Hamilton-Jacoby equations in the limit $B \\ll M^{-1}$. This corresponds to the \"weak\" magnetic field as for the deformation the black hole space-time, but a super-strong field in the astrophysical context. It is shown that the growth rate of the superradiant instability may be amplified by a few times due-to the magnetic field. This enlarges the range of black hole masses for which the instability may be significant in the neutral pion decay. The essential problem which was beyond our investigation is the calculation of the instability growth for the highly rotating black holes. Another important question is a detailed description of the process of filling of the quasi-bound states \\cite{Galtsov2}."
    },
    "0801/0801.0900_arXiv.txt": {
        "abstract": "In this series we construct an effective field theory (EFT) in curved spacetime to study gravitational radiation and backreaction effects. We begin in this paper with a derivation of the self-force on a compact object moving in the background spacetime of a supermassive black hole. The EFT approach utilizes the disparity between two length scales, which in this problem are the size of the compact object $r_m$ and the radius of curvature of the background spacetime ${\\cal R}$ such that $\\varepsilon \\equiv r_m / {\\cal R} \\ll 1$, to treat the orbital dynamics of the compact object, described as an effective point particle, separately from its tidal deformations. The equation of motion of an effective relativistic point particle coupled to the gravitational waves generated by its motion in a curved background spacetime can be derived without making a slow motion or weak field approximation, as was assumed in earlier EFT treatment of post-Newtonian binaries. Ultraviolet divergences are regularized using Hadamard's {\\it partie finie} to isolate the non-local finite part from the quasi-local divergent part. The latter is constructed from a momentum space representation for the graviton retarded propagator and is evaluated using dimensional regularization in which only logarithmic divergences are relevant for renormalizing the parameters of the theory. As a first important application of this framework we explicitly derive the first order self-force given by Mino, Sasaki, Tanaka, Quinn and Wald. Going beyond the point particle approximation, to account for the finite size of the object, we demonstrate that for extreme mass ratio inspirals the motion of a compact object is affected by tidally induced moments at $O(\\varepsilon^4)$, in the form of an Effacement Principle. The relatively large radius-to-mass ratio of a white dwarf star allows for these effects to be enhanced until the white dwarf becomes tidally disrupted, a potentially $O(\\varepsilon^2)$ process, or plunges into the supermassive black hole. This work provides a new foundation for further exploration of higher order self force corrections, gravitational radiation and spinning compact objects. ",
        "introduction": "In two previous papers \\cite{GalleyHu:PRD72,GalleyHuLin:PRD74,Galley:PhD}, using a stochastic field theory approach based on open system concepts,  we derive the scalar, electromagnetic and gravitational self-force to leading order on a particle moving in an arbitrary curved background spacetime. % We begin with the particle following a quantum mechanical path while interacting with a linear quantum field \\cite{JohnsonHu:PRD65,JohnsonHu:FoundPhys35,Johnson:PhD}. The conditions on a stochastic field theory (for open systems) to emerge from a quantum field theory (of closed systems) are that the mass and size of the particle are large enough so the particle worldline is sufficiently decohered from its interactions with the quantum fluctuations of the field that it can be considered as quasi-classical, and yet  sufficiently small that quantum fluctuations manifest as classical stochastic forces \\cite{GellMannHartle:PRD47}.  \\subsection{Quantum, Stochastic and Effective Field Theories}  When there is a significant discrepancy between the two mass (or energy or length) scales in a problem, as in the extreme mass ratio binary systems under consideration here, one could use an open system stochastic description for their dynamics, such as developed in \\cite{GalleyHu:PRD72,GalleyHuLin:PRD74}. When this discrepancy is huge (such as between the QCD and GUT scales in particle physics) the stochastic component is strongly suppressed in the wide range between the two scales, away from the threshold region \\cite{CalzettaHu:PRD55}. Then the stochastic field theory description will give rise to an effective field theory (EFT) description \\cite{Burgess:EFT} for the motion of the small mass subsystem. Due to the large separation in the mass scales quantum loop corrections from the field and the intrinsic quantum mechanical worldline fluctuations are very strongly suppressed. These two factors render a quantum field theory (QFT) into a stochastic field theory  (with sufficiently decohered histories) and in turn (with sufficiently small stochasticity) an effective field theory for the dynamics of the reduced systems. We shall explain the essence and demonstrate the advantages in taking a field theory approach to treat radiation-reaction of classical systems. The application of EFT to the treatment of gravitational radiation from post-Newtonian (PN) binary systems was first introduced by Goldberger and Rothstein \\cite{GoldbergerRothstein:PRD73}. Our formulation of a curved spacetime effective field theory (CS-EFT) goes beyond with two important features: it is for any curved spacetime background and there is no slow motion or weak field restrictions \\footnote{In addition to \\cite{GalleyHu:PRD72, GalleyHuLin:PRD74, GoldbergerRothstein:PRD73} for applying effective field theory techniques to the motion of extended bodies and charges, see also a recent paper of \\cite{KolSmolkin:grqc}.}.   In this and three subsequent papers \\cite{Galley:EFT2, Galley:EFT3, Galley:spin} we construct a CS-EFT and apply it to derive the self-force on a compact object moving in an arbitrary curved background  For concreteness, the background spacetime is assumed to be that of a supermassive black hole (SMBH) with curvature scale $\\cR$ much larger than the size of the compact object $r_m$. The smallness of  the  ratio $\\ve \\equiv r_m / \\cR$ makes it a good expansion parameter for  a perturbation theory treatment describing the extreme mass ratio inspiral (EMRI) of the compact object. These binary systems are expected to be good candidates for detecting gravitational wave signatures using the space-based gravitational wave interferometer LISA \\cite{LISA}. We now give some realistic numbers for astrophysical processes in this category to delimit the range of validity for carrying out these perturbation expansions  for these EMRI sources detectable in LISA's bandwidth . However, we emphasize that the formalism we construct here is of a sufficiently general nature that it can be applied to any compact object  moving in an arbitrary curved background, including those spacetimes sourced by some form of stress-energy and those possessing a cosmological constant.   \\subsection{Relevant scales in EMRIs}  Consider the motion of a compact object (a black hole, neutron star or white dwarf  with a mass $m$ ranging from  about 1 to 100  solar masses) moving through the spacetime of a SMBH with a mass $M \\sim 10^{5-7} M_\\odot$. We have in mind that the compact object moves in a stationary background provided by the supermassive black hole, such as the Schwarzschild or Kerr spacetimes. Such spacetimes are appropriate for a description of the EMRI in which the compact object is bound by the gravitational pull of the SMBH. By emitting gravitational waves the binary system loses energy until the compact object plunges into the SMBH. The emission of gravitational radiation from such a system is expected to be detected with the anticipated construction and launch of the LISA space-based interferometer \\cite{LISA}.  It is believed that most SMBHs lurking in the middle of galaxies, which are thought to host the prime sources of gravitational wave emissions detectable by LISA, are spinning and clean in the sense that most, if not all, of the surrounding material has already fallen into the black hole.  (Active galactic nuclei are a notable exception \\cite{LISAsources}.)  Because of this the Kerr background is perhaps the most astrophysically relevant spacetime for the extreme mass ratio inspiral. The Kerr solution is vacuous ($R_{\\mu\\nu}=0$), stationary and stable under small perturbations \\cite{Whiting:JMathPhys30} and possesses two Killing fields. The first Killing field ($\\xi^\\alpha$) is time-like and describes time-translation invariance everywhere outside of the ergoregion. The second ($\\psi^\\alpha$) is space-like and describes the axial  symmetry  of the spacetime.  The Ernst \\cite{Ernst:JMathPhys17} and Preston-Poisson \\cite{PrestonPoisson:PRD74}  geometries describe  a black hole immersed in an external magnetic field. From an astrophysical viewpoint, the external magnetic fields that a black hole at the center of a galaxy experiences are relatively weak and unlikely to significantly affect the motion of the compact object until a very high order in the perturbation theory.  There are two relevant length scales in EMRIs. The smaller scale is set by the size of the compact object itself, denoted $r_m$. For an astrophysical black hole its radius is $r_{bh} = 2 G_N m \\sim m/ m_{pl}^2$ where $m_{pl}^{-2} = 32\\pi G_N$ in units where $\\hbar = c= 1$ \\footnote{We follow the conventions of \\cite{MTW} so that the metric has mostly positive signature $(-,+,+,+)$.}.  For a neutron star with a mass $\\approx 1.4 M_\\odot$ and a radius of $10-16$ km it follows that $r_{ns} \\approx 4.8 -7.7 \\, G_N m \\sim m/m_{pl}^2$. Therefore, it is to be expected that the size of the compact object, be it a black hole or a neutron star, is of the order of its mass  \\footnote{See Section \\ref{sec:finsize} for more details concerning a white dwarf, which has a typical radius $r_{wd} \\sim 10^3 m / m_{pl}^2$.} .  The second relevant scale is the radius of curvature of the background spacetime, $\\cR$. We take $\\cR$ to be the following curvature invariant \\begin{eqnarray} \\cR = \\big( R_{\\mu\\alpha \\nu \\beta} R^{\\mu\\alpha \\nu \\beta} \\big) ^{-1/4}   .  \\label{curvinvariant0} \\end{eqnarray} For a (possibly rotating) stationary SMBH the radius of curvature is \\begin{eqnarray} \\cR \\sim  \\sqrt{ \\frac{ m_{pl}^2 r^3}{M} } \\end{eqnarray} where $r$ is the typical orbital distance for the compact object away from the central black hole. For example, $r$ is the geometric mean of the semi-major and semi-minor axes of a compact object in an inclined elliptical orbit. In an approximately circular orbit $r$ is the orbital radius and for a particle moving faster than the escape velocity $r$ is the impact parameter.  In the strong field regime where $r \\sim M / m_{pl}^2$ the curvature scale is also $\\sim M / m_{pl}^2$ implying that $r/\\cR \\sim m / M$ whence  a perturbative expansion in $\\ve$ is equivalent to one in $m/M$.  The typical variation in time and space of the background is $\\gsim \\cR$. The wavelength $\\lambda$ of radiated metric perturbations from the compact object in a bound orbit is \\begin{eqnarray} \\lambda \\sim \\sqrt{ \\frac{ m_{pl}^2 r^3 }{ M } } \\sim \\cR , \\end{eqnarray} which shows that the wavelength of the gravitational waves does not provide a separate scale independently from $\\cR$.   \\subsection{The CS-EFT approach: Issues and main features}  The effective field theory approach was first introduced for the consideration of gravitational radiation from post-Newtonian binary systems in \\cite{GoldbergerRothstein:PRD73}, spinning compact objects in \\cite{Porto:PRD73, PortoRothstein:PRL97, Porto:grqc0701106, PortoRothstein:0712.2032}, and dissipative effects due to the absorption of gravitational waves in \\cite{GoldbergerRothstein:PRD73_2, Porto:0710.5150}. See \\cite{Goldberger:LesHouches, PortoSturani:grqc0701105, GoldbergerRothstein:GRG38} for introductory reviews. Let us call these theories PN-EFT: they are constructed to describe the motion of two slowly moving compact objects in a \\emph{flat background}. In particular, the compact objects are treated as effective \\emph{point particles}, the worldlines of which carry non-minimal operators describing the multipole moments from companion-induced tidal deformations as well as possible spin degrees of freedom and other intrinsic moments. Below we describe briefly some general features of EFT and the specific differences between our new CS-EFT approach and the existing PN-EFT.  The use of point particles to source the metric perturbations  about the flat background spacetime (note that the high frequency waves in a quantum description corresponds to massless spin-two particles, the gravitons, in flat space quantum field theory) prompts the appearance of divergences. Fortunately there exists a well-established bank of tools and techniques in quantum field theory for regularizing these divergences and renormalizing the parameters and coupling constants of the theory. A theory is renormalizable in the effective field theory sense if observables are calculated in the low energy limit: the divergences can always be absorbed into a renormalization of the coupling constants of the infinite number of non-minimal worldline operators. The use of dimensional regularization is particularly useful in effective field theories because the renormalization group equations are mass-independent for this scheme indicating that only logarithmic divergences contribute to the renormalization procedure \\footnote{The effective field theory approach has also been used in \\cite{Leibovich:NRGRtalk} to derive the radiation reaction force on an electrically charged extended body interacting with its own electromagnetic radiation, generalizing the Abraham, Lorenz and Dirac (ALD) equation for a point charge. It can also be derived from a stochastic field theory perspective, which contains EFT. See \\cite{JohnsonHu:PRD65, JohnsonHu:FoundPhys35, Johnson:PhD}.}.  Our CS-EFT approach differs from this group of work in two ways. First, we work with an arbitrary \\emph{curved spacetime}. Second, we allow for the compact object to move with \\emph{relativistic speeds} in \\emph{strong field} regions of the background spacetime. The post-Newtonian effective field theory of \\cite{GoldbergerRothstein:PRD73} treats bodies moving slowly through a weak gravitational field.   \\subsubsection{In-In formulation for real and causal equations of motion}  Technically there are also fundamental differences. To derive real and causal equations of motion we emphasize the need to use the closed-time-path  (CTP)  integral formalism based on an in-in generating functional \\cite{Schwinger:JMathPhys2, Keldysh:JEPT20, ZhouSuHaoYu:PhysRep118, Jordan:PRD33, CalzettaHu:PRD35, CalzettaHu:PRD37} (`in' and `out' here refer to the initial and final vacua used to define the vacuum transition amplitude in the generating functional). The authors of \\cite{GoldbergerRothstein:PRD73} use the in-out formalism which is acceptable for field theories in a flat background spacetime since the in- and out-vacua are equivalent up to an irrelevant phase in that special case. In the presence of spacetime curvature, the difference becomes serious, as we will show in \\cite{Galley:EFT2} for the EMRI scenario. The in-out formalism can calculate matrix elements in scattering processes, but not expectation values of physical observables in real-time evolution. The equations of motion for the effective particle dynamics obtained from an in-out formulation are not  generally  causal. An initial value formulation of quantum field theory via the CTP generating functional is the only correct way for the description of the system's evolution -- it guarantees real and causal equations of motion for the particle dynamics \\cite{Jordan:PRD33, CalzettaHu:PRD37}.   \\subsubsection{Effective point particle description}  Another  important ingredient in our construction is an effective point particle description for the motion of the compact object. Going beyond the point particle approximation is necessary to include the effects of tidal deformations induced by the background curvature as well as the effects from spin and other intrinsic moments.  Following \\cite{GoldbergerRothstein:PRD73},  we introduce all possible terms into the point particle action that are consistent with general coordinate invariance and reparameterization invariance (and invariance under $SO(3)$ rotations for a non-spinning spherically symmetric compact object). By implementing a matching procedure using coordinate invariant observables we can match the observables of the effective point particle theory with the long wavelength limit of observables in the full ``microscopic\" theory to determine the values of the coupling constants of the non-minimal terms. As we will show in Section \\ref{sec:finsize} this allows us to deduce the order at which finite size effects affect the motion of the compact object through the statement of an \\emph{Effacement Principle}. To our knowledge this has not been explicitly given in the literature before for the EMRI scenario.   \\subsubsection{The power counting rules}  Power counting is a generalization of dimensional analysis. In our perturbative treatment it is crucial for determining how the Feynman rules scale with the parameter $\\ve$.  Once the scaling of the Feynman rules are known we determine all of the tree-level Feynman diagrams that appear at a particular order. Those diagrams containing graviton loops are safely ignored. We also assemble the diagrams that include the non-minimal worldline operators describing the finite size of the compact object. Significantly, this allows us to determine the order in $\\ve$ at which finite size effects enter the particle equations of motion. With the power counting rules  the EFT approach becomes an efficient and systematic framework for calculating the self-force to any order in perturbation theory. Furthermore, by knowing how each Feynman diagram scales with $\\ve$ we can study a particular physical interaction that is of interest by focusing our attention on a single diagram or on a few diagrams without having to calculate every contribution that appears at that order and at lower orders. For example, the leading order spin-spin interaction (spin here refers to that of the compact object, not of the SMBH) contributes to the self-force at third order in $\\ve$ for a maximally rotating compact object  and can be calculated from the appropriate Feynman diagram \\cite{Galley:PhD, Galley:spin}. The power counting rules, the Feynman rules and their scaling with $\\ve$ are derived in Sections \\ref{ch3:powercounting} and \\ref{ch3:Feynmanrules}.   \\subsubsection{Divergences and Regularization}  In Section \\ref{ch3:reg} we propose a method for regularizing the divergences that appear in the effective action. Our approach utilizes a mixture of distributional and momentum space techniques within the context of dimensional regularization. We know from previous work % that the finite part of the self-force is generally non-local and history dependent. However, the ultraviolet divergences are quasi-local and independent of the history of the effective point particle's motion. To isolate the quasi-local divergence from the non-local finite part we use the method of Hadamard's {\\it  partie finie}, or finite part, from distribution theory. (See Appendix \\ref{app:distro} for a brief review of the definitions and concepts of distribution theory relevant in this work.)  After isolating the divergence from the non-local, finite remainder we then use a momentum space representation for the propagator  in a curved background, first derived for a scalar field by Bunch and Parker \\cite{BunchParker:PRD20},  to calculate the divergent contributions. Their method is straightforward but not efficient for higher spin fields, including gravitons in a curved space. (See also the related work of \\cite{HuOConnor:PRD30}.)  A novel method applicable for any tensor field is developed in \\cite{Galley:momentum} for computing the momentum space representation of the Feynman propagator. The method is sufficiently general to do the same for any quantum two-point function, including the retarded propagator $D^{ret}_{\\alpha \\beta \\gamma^\\prime \\delta^\\prime} (x,x^\\prime)$.  This approach makes use of diagrammatic techniques borrowed from perturbative quantum field theory. In Riemann normal coordinates, we expand the field action in terms of the displacement from the point $x$. The series can be represented in terms of Feynman diagrams, which allows for an efficient evaluation of each term in the expansion. Furthermore, we prove that some of the diagrams are zero to all orders. This identity is not recognized in \\cite{BunchParker:PRD20} even though its relation to certain steps made in their calculations is evident.   \\subsection{MST-QW Equation}  Assembling all these essential ingredients, in Section \\ref{sec:EFTMST-QW} we give a demonstration of how the curved spacetime effective field theory construction is implemented, outline the steps in the regularization of divergences, perform an actual calculation of the effective action and from it derive the equation of motion for the compact object including the effect of gravitational self-force. As an application we work to first order in the mass ratio  (i.e., $\\ve$) and obtain the well-known Mino-Sasaki-Tanaka-Quinn-Wald (MST-QW) equation \\cite{MinoSasakiTanaka:PRD55, QuinnWald:PRD56}.  This also sets the stage for calculating the second order self-force,  the emitted gravitational waves  and the motion of compact objects with spin \\cite{Galley:EFT2, Galley:EFT3, Galley:spin}. ",
        "conclusions": "We develop an effective field theory approach for systematically deriving the self-force on a compact object moving in an arbitrary curved spacetime without the slow motion or weak field restrictions. The EFT is a realization of the open quantum system paradigm  in systems with a large scale separation such that the system's induced fluctuations from the backreaction of the coarse-grained quantum field is utterly negligible \\cite{CalzettaHu:PRD55}. An initial value formulation of quantum field theory is adopted here using the closed-time-path (CTP) formalism for the in-in generating functional, which guarantees real and causal equations of motion for the compact object. As an illustration of the procedures involved in our approach we showed how to derive the MST-QW equation describing the (first order) self-force on a compact object.  The CTP formalism is needed for a calculation of the second order self-force as will be shown in \\cite{Galley:EFT2}.   We describe the compact object as an effective point particle that is capable of accounting for tidally induced finite size effects. % In calculating the effective action we encounter ultraviolet divergences stemming from a point particle interacting with arbitrarily high frequency modes of a graviton field. Using Hadamard's {\\it partie finie} to isolate the non-local finite part from the quasi-local divergences we are able to implement dimensional regularization within a (quasi-local) momentum space representation for the graviton propagator \\cite{Galley:momentum}. As such, all power divergences can be immediately set to zero implying that only logarithmic divergences are relevant for renormalizing the parameters of the theory. At first order, the effective action has a power divergence and may therefore be trivially regularized using dimensional regularization.  In the spirit of an Effacement Principle we find that the finite size of the compact object first affects its motion at $O(\\ve^4)$ for a non-spinning black hole and neutron star. For a white dwarf star we deduce that such effects may be enhanced until the white dwarf is tidally disrupted at $O(\\ve^2)$ in which case the effective point particle description, and in particular the effective field theory developed here, breaks down. One may conceivably construct a new effective field theory by treating the supermassive black hole, the white dwarf and the accreting mass as an effective point particle possessing many relevant non-minimal couplings to the background geometry describing the intrinsic moments of this composite object.   The leading order finite size corrections cause a deviation from the motion of a minimally coupled point particle that is not caused by interactions with gravitons but is due to the torques that develop on the tidally deformed compact object. On the other hand, the self-force is affected by the induced moments of the compact object at $O(\\ve^5)$.  In summary, the EFT approach has at least two major advantages over the existing approaches: It provides a systematic procedure for carrying out a perturbative treatment, and an economical way to treat the ultraviolet divergences. Our CS-EFT improves on the PN-EFT introduced in \\cite{GoldbergerRothstein:PRD73} in that it is valid for a general curved spacetime and not limited to slow motion or weak field conditions. These will prove to be of special benefit for higher order self-force calculations. We will apply these steps to calculate the self-force at second order in $\\ve$ \\cite{Galley:EFT2}, the gravitational radiation emitted by EMRIs \\cite{Galley:EFT3} and the self-force on spinning compact objects \\cite{Galley:spin}."
    },
    "0801/0801.0307_arXiv.txt": {
        "abstract": "{\\bf Abstract.} There is accumulating evidence that (fundamental) scalar fields may exist in Nature. The gravitational collapse of such a boson cloud would lead to a {\\em boson star} (BS) as a new type of a compact object. Similarly as for white dwarfs and neutron stars, there exists a limiting mass, below which a BS is {\\em stable} against complete gravitational collapse to a black hole.  According to the form of the self-interaction of the basic constituents and the spacetime symmetry, we can distinguish mini-, axidilaton, soliton, charged, oscillating and rotating BSs. Their compactness prevents a Newtonian approximation, however, modifications of general relativity, as in the case of Jordan-Brans-Dicke theory as a low energy limit of strings, would provide them with {\\em gravitational memory}.  In general, a BS is a compact, {\\em completely regular} configuration with structured layers due to the anisotropy of scalar matter, an exponentially decreasing 'halo', a critical mass inversely proportional to constituent mass, an effective radius, and a large particle number. Due to the Heisenberg principle, there exists a completely stable branch, and as a coherent state, it allows for rotating solutions with {\\em quantised angular momentum}.  In this review, we concentrate on the fascinating possibilities of detecting the various subtypes of (excited) BSs: Possible signals include gravitational redshift and (micro-)lensing, emission of gravitational waves, or, in the case of a giant BS, its dark matter contribution to the rotation curves of galactic halos. ",
        "introduction": "In this review, we will assume that {\\em fundamental scalar fields} exist in Nature, that --- in the early stages of the universe --- they would have formed absolutely stable soliton-type configurations kept together by their self-generated gravitational field. Theoretically, such configurations are known as {\\em boson stars} (BSs). In some characteristics, they resemble neutron stars; in other aspects, they are different and, thereby, astronomers may have a chance to distinguish BS signals from other compact objects, like neutron stars or black holes (BHs). BSs can be regarded as descendants of self-gravitating photonic configurations called {\\em geons} (gravitational electromagnetic units) and proposed in 1955 by Wheeler \\cite{Wh55}.  In this review, we shall mainly concentrate on the results which are not included in the first reviews from 1992 \\cite{Je92,LP92,LM92} as well as later in the Marcel Grossman meetings in Jerusalem and Rome \\cite{mielkeMG8,MS02}. Since then, the number of publications investigating how BSs could possibly be detected have increased; we intend to focus strongly on these more recent research results. This could provide astronomers a handling on possible signals of BSs. All of these theoretical results are only a humble beginning of the understanding where and in which scenarios a BS could be detected. We hope other researchers may find this review stimulating for gaining new ideas or for refining  older research.  A second intention of this review is to distinguish clearly the different theoretical models underlying the label BS which could lead to different observational consequences.  The first distinction is that the matter part of a BS can be described by either a complex or by a real scalar field. Then, there can be different interactions: (a) self-interactions described by scalar field potentials, (b) minimal coupling to gauge fields, the scalar field can be carrier of a charge (electric or hypercharge, e.g.). Moreover, the BS scalar field can couple, in standard general relativity (GR), minimally, or, in scalar-tensor (ST) or Jordan-Brans-Dicke (JBD) theory, non-minimally to gravity. In the latter case, if the strength of gravitational force is influenced by the JBD scalar field, there is an interaction of the BS scalar field with the real JBD scalar field; this can lead to {\\em gravitational memory} effects in BSs. JBD theory is closely related to low energy limits of superstring theories \\cite{CJ78} which imply the primordial production of scalar fields. Then, relics like the dilaton, the axion, combined as {\\em axidilaton}, or other moduli fields could remain in our present epoch as candidates.  Already the first two papers on BSs provided the two main directions: The history of these {\\em hypothetical} stars starts in 1968 with the work of Kaup \\cite{K68} using a complex massive scalar field with gravitational interactions in a semi-classical manner. The energy-momentum tensor is calculated classically providing the source for gravitation; a very detailed investigation of the solution classes was done in \\cite{FLP87a}. However, one year later, Ruffini and Bonazzola \\cite{RB69} used field quantisation of a real scalar field and considered the ground state of $N$ particles. The vacuum expectation value of the field operators yield the same energy-momentum tensor and thus, not surprisingly, the same field equations. The two different physical constituents, complex versus quantised real scalar field BS, yield the same macroscopical results. It should be noted, however, that the gravitational field $g_{\\mu\\nu}$ is kept {\\em classical} due to the non-renormalisability of standard GR, or, alternatively, treated it as principal low energy part of some renormalisable superstring model.  More recently, a BS using a field quantised complex scalar field has been constructed. BSs consisting of pure semi-classical real scalars cannot exist because their static solutions in flat spacetime are unstable due to Derrick's theorem \\cite{De64}, solutions may arise only if the real scalar field possesses a time-dependence leading to non-static {\\em oscillating} BSs. However, if a fermion star is present as well, a real scalar component can be added and if it interacts with the fermions, then a combined boson-fermion star is the result as in the first calculation by T.D.~Lee and Pang \\cite{LP87}; cf.~Section \\ref{leepang}.  A challenge for the BS model is that, so far, no {\\em fundamental} scalar particle has been detected with certainty in the laboratory. Several are proposed by theory: The Higgs particle $h$ is a necessary ingredient of the standard model. However, the possible discovery of the Higgs boson of mass $m_{\\rm h}=114.5$ GeV/$c^2$ at the Large Electron Positron (LEP) collider at CERN \\cite{Ac00,Ba00} gives the BS strand of investigation a fresh impetus.  For the BS, we need a scalar particle which does not decay; or if it decays (like the Higgs, e.g.), one has to assume that, in the gravitational binding, the inverse process is efficient enough for an equilibrium, as is the case for the $\\beta$-decay inside a neutron star. In the latter case, the direct physical predecessor of that kind of a BS is not clear; but, of course, unstable Higgs particles could not have formed a BS by themselves. In order to explore that unknown particle regime, one can play with the model parameters such as particle mass or interaction constants. Then, BSs of rather different sizes can occur: it could be just a `gravitational atom'; it could be as massive as the presumed BH in the central part of a galaxy; or it could be an alternative explanation for parts of the dark matter in the halo of galaxies.  In 1995, experiments \\cite{AK02} proved the existence of the fifth possible state of matter, the {\\em Bose-Einstein condensate} (BEC); in 2001, the Nobel prize was awarded for its experimental realization in traps. BSs, if they exist, would be an astrophysical realization with a self-generated gravitational confinement, cf.~\\cite{JB01,BLV01}.  Let us sum up some of the properties of complex scalar field BSs (following an earlier version of \\cite{T02}) in Table I as we shall discuss in Section III.   \\begin{center} \\parbox{14cm} {TABLE I. Overview of some complex scalar field BS properties.} \\end{center} $$\\vbox{\\offinterlineskip \\hrule \\halign{&\\vrule#&\\strut\\quad\\hfil#\\quad\\hfil\\cr height2pt&\\omit&&\\omit&&\\omit&\\cr & {\\bf Property}      && {\\bf BS}    &\\cr height2pt&\\omit&&\\omit&&\\omit&\\cr \\noalign{\\hrule} height2pt&\\omit&&\\omit&&\\omit&\\cr & Constituents  \\hfill &&   Scalars  (Bose-Einstein-Condensation)    \\hfill &\\cr height2pt&\\omit&&\\omit&&\\omit&\\cr & Pressure support \\hfill &&   Heisenberg's uncertainty relation    \\hfill &\\cr height2pt&\\omit&&\\omit&&\\omit&\\cr & Size   \\hfill &&  Gigantic up to very compact (few Schwarzschild radii): Table IV, Fig.~\\ref{fig2} \\hfill &\\cr height2pt&\\omit&&\\omit&&\\omit&\\cr & Surface  \\hfill &&   Atmosphere    \\hfill &\\cr height2pt&\\omit&&\\omit&&\\omit&\\cr & Appearance  \\hfill &&  Transparent (if only gravitationally interacting)  \\hfill &\\cr height2pt&\\omit&&\\omit&&\\omit&\\cr & Structure  \\hfill &&  Einstein-Klein-Gordon equation      \\hfill &\\cr height2pt&\\omit&&\\omit&&\\omit&\\cr & Gravitational potential  \\hfill &&  Newtonian weak up to highly relativistic \\hfill &\\cr height2pt&\\omit&&\\omit&&\\omit&\\cr & Last stable orbit \\hfill &&   None    \\hfill &\\cr height2pt&\\omit&&\\omit&&\\omit&\\cr & Rotation \\hfill &&   Differentially; discrete    \\hfill &\\cr height2pt&\\omit&&\\omit&&\\omit&\\cr & Avoiding a baryonic BH by \\hfill &&   Jet, particle dynamics    \\hfill &\\cr height2pt&\\omit&&\\omit&&\\omit&\\cr & Avoiding  a scalar BH by \\hfill &&   Stability, evolution    \\hfill &\\cr height2pt&\\omit&&\\omit&&\\omit&\\cr & Gravitational redshift \\hfill &&   Comparable to neutron star values, but larger due to transparency \\hfill &\\cr height2pt&\\omit&&\\omit&&\\omit&\\cr & Gravitational (micro-)lensing \\hfill &&   Extreme deflection angles possible (Fig.~\\ref{fig3}); MACHOS \\hfill &\\cr height2pt&\\omit&&\\omit&&\\omit&\\cr & Luminosity  \\hfill &&   Larger than luminosity of a BH   \\hfill &\\cr height2pt&\\omit&&\\omit&&\\omit&\\cr & Star disruption (tidal radius) \\hfill &&   Yes    \\hfill &\\cr height2pt&\\omit&&\\omit&&\\omit&\\cr & Distinctive observational  \\hfill &&   Broadening of emission lines    \\hfill &\\cr & \\ \\ \\ signatures    \\hfill &&   Gravitational waves    \\hfill &\\cr &     \\hfill &&   {\\v C}erenkov radiation    \\hfill &\\cr height2pt&\\omit&&\\omit&&\\omit&\\cr height2pt&\\omit&&\\omit&&\\omit&\\cr} \\hrule}$$    \\subsection{Fundamental scalar fields in Nature?}  The physical nature of the  spin-0-particles out of which the BS is presumed to consist, is still an open issue. Until now, no fundamental elementary scalar particle has been detected with certainty in accelerator experiments, which could serve as the main constituent of the BS. In the theory of Glashow, Weinberg, and Salam, a Higgs boson doublet $(\\Phi^+, \\Phi^0)$ and its anti-doublet $(\\Phi^-, \\bar \\Phi^0)$ are necessary ingredients in order to generate masses for the `heavy photons', i.e.~the $W^{\\pm }$ and $Z^0$ gauge vector bosons \\cite{Qu83}. After symmetry breaking, only one real scalar particle, the Higgs particle $h:=(\\Phi^0 + \\bar \\Phi^0)/\\sqrt{2}$, remains free and occurs in the state of a constant scalar field background \\cite{Pe97}. As it is indicated by the rather heavy top quark \\cite{Ab95} of 176 GeV/$c^2$, the mass of the Higgs particle is expected to be below 1000 GeV/$c^2$. This is supported by the possible discovery of the Higgs boson of mass $m_{\\rm h}=114.5$ GeV/$c^2$ at the Large Electron Positron (LEP) collider at CERN \\cite{Ac00,Ba00}. In order to stabilise such a light Higgs against quantum fluctuations, a {\\em supersymmetric} extension of the standard model is desirable. However, then it will be accompanied by additional heavy Higgs fields $H^0$, $A^0$, and a charged doublet $H^{\\pm}$ in the mass range of 100 GeV/$c^2$ to 1 TeV/$c^2$. For the unlike case of a Higgs mass $m_{\\rm h}$ above 1.2 TeV/$c^2$, the self-interaction $U(\\Phi)$ of the Higgs field is so large that any  perturbative approach of the standard model becomes unreliable. Tentatively, a conformal extension of the standard model with gravity included has been analysed, cf.~\\cite{PR94,HMMN95}. Future high-energy experiments at the LHC at {\\sc Cern} should reveal if Higgs particles really exist in Nature.  As free particles, the Higgs boson of Glashow-Weinberg-Salam theory is unstable, e.g.~with respect to the decays $h \\rightarrow W^+ + W^-$ and $h \\rightarrow Z^0+ Z^0$, if it is heavier than the gauge bosons. In a compact object like the BS, these decay channels are expected to be in partial equilibrium with the inverse processes $Z^0+ Z^0 \\rightarrow h$ or, via virtual $W$ triangle graph, $Z^0 \\rightarrow h + \\gamma$, for instance \\cite{CL84}, by utilising gravitational binding energy. This is presumably in full analogy with the neutron star \\cite{HWW94,ST83,We99} or quark star \\cite{KWWG95,GKW95}, where one finds an equilibrium of $\\beta $- and inverse $\\beta $-decay or of a quark-gluon plasma (with similar decay and inverse-decay processes) and thus stability of the macroscopic star with respect to radioactive decay.  The known spin-0-particles in particle physics, although effectively described by a Klein-Gordon (KG) equation, are not elementary, they consist of quarks having spin 1/2. Nevertheless, they possess some physical properties which can help us to understand which kind of particles can occur in the different BS types.  The electrically charged pions $\\pi^{\\pm}$ with a mass of about 140 MeV/$c^2$ are a typical example for complex scalar fields \\cite{IZ80,MS93}. The electrically neutral ${\\rm K}^0$, $\\bar{\\rm K}^0$ mesons are described by a complex KG field as well; ${\\rm K}^0$ and $\\bar{\\rm K}^0$ are distinguished by their hypercharge $Y:=B+S =\\pm 1$, where $B=0$ is the baryon number and $S$ is the strangeness corresponding to a global $U(1)$ symmetry. In both cases,  particles (e.g.~$\\pi^{+}$, ${\\rm K}^0$) and anti-particles (e.g.~$\\pi^{-}$, $\\bar{\\rm K}^0$) occur. Generally, Noether's theorem leads to a conserved charge \\be Q = e ( N^+ - N^- )  \\label{Ncharge} \\ee related to the number $N$ of particles and anti-particles, respectively, where $e$ defines the absolute value of the charge. Let us stress that if the KG equation for a complex scalar field admits a {\\em global} $U(1)$ symmetry, as in the case for the ${\\rm K}^0$, $\\bar{\\rm K}^0$ mesons, ``merely'' a conserved particle number $N$ arises and the charge is the total strangeness $S$. Regarding the high degree of instability of pions, $K$ mesons, and Higgs fields, the effect of gravitational binding is essential to see if these particles had enough time to form a BS.  In string theories, there may exist several fundamental scalar particles having different global or local charges. For example, for complex scalar particles with global $U(1)$ symmetry or for real scalar particles, the BS charge shall be called just {\\em boson number} $N$. For complex scalar particles with local $U(1)$ symmetry, the BS has a charge given by (\\ref{Ncharge}). In a BS with temperature near zero, we will assume that all constituents are particles, i.e.~no anti-particles present, or vice versa.  In his model, Kaup \\cite{K68} used a complex scalar field with global $U(1)$ symmetry so that no gauge bosons are present. Ruffini and Bonazzola \\cite{RB69}, however, took a real scalar field for the general relativistic BS and a complex scalar field with global $U(1)$ symmetry for the Newtonian BS description. In the case that $U(1)$ symmetry is {\\em local}, the complex scalar field couples to the gauge field contribution, cf.~\\cite{IZ80}, p.~31; Jetzer and van der Bij \\cite{JB89} called their hypothetical astrophysical object {\\em charged boson star} (CBS). Within the Glashow-Weinberg-Salam theory before symmetry-breaking of the group $SU(2)\\times U(1)$, this $U(1)$ charge describes the weak hypercharge, and not an electric one. During the earliest stages of the universe, a complex scalar field with different kinds of $U(1)$ charge (than electric or hyperweak) could have been generated.  There is also the particle/field classification with respect to the spatial reflection $P$. Both a real and a complex scalar field can be either a {\\em scalar} or a {\\em pseudo-scalar}.  Thus, particle characterisation will lead to different consequences for the detection of each BS model.     \\subsection{Boson star as a self-gravitating Bose-Einstein condensate \\label{BEC}}  Since Einstein and Bose it is well-known that scalar fields represent identical particles which can occupy the same ground state. Such a {\\em Bose-Einstein condensate} (BEC) has been experimentally realized in 1995 for cold atoms of even number of electrons, protons, and neutrons, see Anglin and Ketterle \\cite{AK02} for a recent review. In the mean-field ansatz, the interaction of the atoms in a dilute gas is approximated by the effective potential \\be U(|\\Psi|)_{\\rm eff}= \\frac{\\lambda}{4} |\\Psi|^4 \\, .  \\label{Ueff} \\ee This leads to a {\\em nonlinear} Schr\\\"odinger equation for $\\Psi$, in this context known as Gross-Pitaevskii equation. In a microscopic approach, one introduces bosonic creation and annihilation operators $b^\\dagger$ and $b$, respectively, satisfying \\be [b,b^\\dagger]  =1 \\ee and finds that every number conserving normal ordered correlation function $\\langle b_1 \\cdots b_n \\rangle$ splits into the sum of all possible products of contractions $\\langle b_i  b_j \\rangle$ as in Wick's theorem of quantum field theory (QFT). For $n=1$ one recovers the Gross-Pitaevskii equation, whereas the next order leads to the Hartree-Fock-Bogoliubov equations, see \\cite{KB01} for details. Therefore, it is gratifying to note that BSs with {\\em repulsive} self-interaction $U(|\\Phi|)$ considered already in Refs.~\\cite{MS81,CSW86} have their counterparts in the effective potential (\\ref{Ueff}) of BEC. Thus, some authors \\cite{JB01,BLV01} advocate consideration of a cold BS as a {\\em self-gravitating} BEC on an astrophysical scale.  Recently, the so-called {\\em vortices}, collective excitations of BECs with angular momentum in the direction of the vortex axis $z$, have been predicted \\cite{WH00} and then experimentally prepared \\cite{MCW00}. Since, quantum-mechanically, the z-component $J_z= a N \\hbar$ of the total angular momentum is necessarily quantised by the {\\em azimuthal} quantum number $a$, it is not possible to ``continuously\" deform this state to the ground state, and by this circumstance, contributing to its (meta-) stability. In 1996, axisymmetric solutions of the Einstein-KG equations have been found by us \\cite{SM96,MS96,SM98,SM99} which are rotating BSs (differentially due to the frame-dragging in curved spacetime) and exhibit, as a collective state, the same relation $J =aN$ for the total angular momentum as in the case of the vortices of an BEC and a meta-stability against the decay into the ground state as well; cf.~Section \\ref{rotaBS}. ",
        "conclusions": ""
    },
    "0801/0801.4304_arXiv.txt": {
        "abstract": "We present an analysis of the 5--8~$\\mu$m \\textit{Spitzer}-IRS spectra of a sample of 68 local Ultraluminous Infrared Galaxies (ULIRGs). Our diagnostic technique allows a clear separation of the active galactic nucleus (AGN) and starburst (SB) components in the observed mid-IR emission, and a simple analytic model provides a \\textit{quantitative} estimate of the AGN/starburst contribution to the bolometric luminosity. We show that AGNs are $\\sim30$ times brighter at 6~$\\mu$m than starbursts with the same bolometric luminosity, so that even faint AGNs can be detected. Star formation events are confirmed as the dominant power source for extreme infrared activity, since $\\sim85\\%$ of ULIRG luminosity arises from the SB component. Nonetheless an AGN is present in the majority (46/68) of our sources. ",
        "introduction": "Ultraluminous Infrared Galaxies (ULIRGs, $L_\\mathit{IR}>10^{12} L_\\odot$) are the local counterparts of the high-redshift objects dominating the cosmic background in the far-infrared and millimetric bands. Unveiling the nature of their energy source is fundamental in order to understand the star formation history and the obscured AGN activity in the distant Universe.  Since their discovery, several infrared indicators have been proposed to determine whether the central engine in ULIRGs is an AGN or a starburst (SB). The presence of high-ionization lines in the mid-IR spectra of ULIRGs points to AGN activity, while intense PAH emission features are typical of starburst environments (Genzel et al. 1998; Laurent et al. 2000). Recently, the absorption feature of amorphous silicate grains centered at 9.7~$\\mu$m has also been used together with the PAH emission to assess the nature of the obscured power source (Spoon et al. 2007). An alternate way to disentangle AGNs and SBs in ULIRGs has been proposed by Risaliti et al. (2006, hereafter R06), based on the separation of the two continuum components in 3--4~$\\mu$m spectra. This method has been successfully applied to a sample of $\\sim$50 nearby ULIRGs (Risaliti, Imanishi \\& Sani 2007, \\textit{submitted}) and provided an estimate of the average AGN/SB contribution to ULIRGs. The key reason for using the continuum emission at $\\lambda \\simeq$3--4~$\\mu$m as a diagnostic is the difference between the 3-$\\mu$m to bolometric ratios in AGNs and starbursts ($\\sim$two orders of magnitude larger in the former). This makes the detection of the AGN component possible even when the AGN is heavily obscured and/or bolometrically weak compared to the starburst. However the original prescription is limited by the low quality of the available L-band spectra of ULIRGs (R06, Imanishi et al. 2006), which makes the results on individual sources highly uncertain, except for the $\\sim$10--15 brightest objects.  At present we have extended the analysis to the 5--8~$\\mu$m spectral band, using the observations of the IRS instrument (Houck et al.~2004) onboard \\textit{Spitzer}. We disentangled the AGN and SB contributions to the observed 5--8~$\\mu$m emission of ULIRGs by combining average spectral templates representing the different properties of the two physical processes at work. The high quality of \\textit{Spitzer}-IRS data, in addition to the relatively low dispersion of the intrinsic continuum properties of both AGNs and starbursts in this spectral range, allows a much more accurate determination of the AGN/SB components than possible at other wavelengths (e.g. X-rays) or with other diagnostic methods based on emission lines. In this paper we present our decomposition method, and discuss a simple analytical model providing a \\textit{quantitative} estimate of the AGN/SB contribution to the bolometric luminosity of each source. ",
        "conclusions": "The use of average templates for AGN and SB emission has allowed us to disentangle the two components in the 5--8~$\\mu$m spectra of 68 local ULIRGs, observed with the \\textit{Spitzer Space Telescope}. We have been able to detect an AGN in more than 60\\% of our sources, and estimate its contribution to the bolometric luminosity. In a statistical sense, we confirm that local Ultraluminous Infrared Galaxies are powered for $\\sim$85\\% by intense star formation and for the remaining $\\sim$15\\% by AGN activity. Our method proves to be successful in unveiling an intrinsically faint or obscured AGN inside a ULIRG. In this context we also put on a sound basis our initial assumption that the wavelength interval 5--8~$\\mu$m is an appropriate spectral range in order to search for AGNs: an AGN turns out to be approximately 30 times more luminous at 6~$\\mu$m than a starburst with the same bolometric luminosity."
    },
    "0801/0801.4903_arXiv.txt": {
        "abstract": "This paper proposes a new phenomenology for strong incompressible MHD turbulence with nonzero cross helicity. This phenomenology is then developed into a quantitative Fokker-Planck model that describes the time evolution of the anisotropic power spectra of the fluctuations propagating parallel and anti-parallel to the background magnetic field~$\\bm{B}_0$.  It is found that in steady state the power spectra of the magnetic field and total energy are steeper than $k_\\perp^{-5/3}$ and become increasingly steep as $C/{\\cal E}$ increases, where $C=\\displaystyle \\int d^3\\! x \\; \\bm{v}\\cdot\\bm{B}$ is the cross helicity, ${\\cal E}$ is the fluctuation energy, and $k_\\perp$ is the wavevector component perpendicular to~$\\bm{B}_0$. Increasing $C$ with fixed ${\\cal E}$ increases the time required for energy to cascade to smaller scales, reduces the cascade power, and increases the anisotropy of the small-scale fluctuations. The implications of these results for the solar wind and solar corona are discussed in some detail. ",
        "introduction": "Much of our current understanding of incompressible magnetohydrodynamic (MHD) turbulence has its roots in the pioneering work of Iroshnikov~(1963) and Kraichnan~(1965). These studies emphasized the important fact that Alfv\\'en waves travelling in the same direction along a background magnetic field do not interact with one another and explained how one can think of the cascade of energy to small scales as resulting from collisions between oppositely directed Alfv\\'en wave packets.  They also argued that in the absence of a mean magnetic field, the magnetic field of the energy-containing eddies at scale~$L$ affects fluctuations on scales~$\\ll L$ much in the same way as would a truly uniform mean magnetic field.  Another foundation of our current understanding is the finding that MHD turbulence is inherently anisotropic.  Montgomery \\& Turner~(1981) and Shebalin, Matthaeus, \\& Montgomery~(1983) showed that a strong uniform mean magnetic field~$\\bm{B}_0$ inhibits the cascade of energy to small scales measured in the direction parallel to~$\\bm{B}_0$. This early work was substantially elaborated upon by Higdon~(1984), Goldreich \\& Sridhar~(1995, 1997), Montgomery \\& Matthaeus~(1995), Ng \\& Bhattacharjee~(1996, 1997), Galtier et al~(2000), Cho \\& Lazarian~(2003), Oughton et al~(2006), and many others. For example, Cho \\& Vishniac (2000) used numerical simulations to show that when the fluctuating magnetic field~$\\delta B$ is $\\ga B_0$ the small-scale turbulent eddies become elongated along the local magnetic field direction.  Goldreich \\& Sridhar (1995) introduced the important and influential idea of ``critical balance,'' which holds that at each scale the linear wave period for the bulk of the fluctuation energy is comparable to the time for the fluctuation energy to cascade to smaller scales. Goldreich \\& Sridhar (1995, 1997), Maron \\& Goldreich~(2001), and Lithwick \\& Goldreich~(2001) clarified a number of important physical processes in anisotropic MHD turbulence and used the concept of critical balance to determine the ratio of the dimensions of turbulent eddies in the directions parallel and perpendicular to the local magnetic field.  Over the last several years, research on MHD turbulence has been proceeding along several different lines. For example, one group of studies has attempted to determine the power spectrum, intermittency, and anisotropy of strong incompressible MHD turbulence using direct numerical simulations. (See, e.g., Cho \\& Vishniac 2000, M\\\"uller \\& Biskamp 2000, Maron \\& Goldreich 2001, Cho et~al~2002, Haugen et~al~2004, Muller \\& Grappin~2005, Mininni \\& Pouquet 2007, Perez \\& Boldyrev 2008).  Another series of papers has explored the properties of anisotropic turbulence in weakly collisional magnetized plasmas using gyrokinetics, a low-frequency expansion of the Vlasov equation that averages over the gyromotion of the particles.  (Howes et al~2006, 2007a, 2007b; Schekochihin et al 2007). These authors investigated the transition between the Alfv\\'en-wave cascade and a kinetic-Alfv\\'en-wave cascade at length scales of order the proton gyroradius~$\\rho_i$, as well as the physics of energy dissipation and entropy production in the low-collisionality regime. Turbulence at scales~$\\lesssim \\rho_i$ has also been examined both numerically and analytically within the framework of fluid models, in particular Hall MHD and electron MHD.  (Biskamp, Schwarz, \\& Drake 1996, Biskamp et~al~1999, Matthaeus et al~2003; Galtier \\& Bhattacharjee 2003, 2005; Cho \\& Lazarian 2004; Brodin et al~2006, Shukla et al~2006).  Another group of studies has investigated the power spectrum, intermittency, and decay time of compressible MHD turbulence.  (Oughton et al 1995, Stone et~al~1998, Lithwick \\& Goldreich~2001, Boldyrev et al~2002, Padoan et al~2004, Elmegreen \\& Scalo 2004).  Additional work by Kuznetsov (2001), Cho \\& Lazarian (2002, 2003), Chandran (2005), and Luo \\& Melrose (2006) has begun to address the way in which Alfv\\'en waves, fast magnetosonic waves, and slow magnetosonic waves interact in compressible weak MHD turbulence. Another recent development is the finding that strong incompressible MHD turbulence leads to alternating patches of alignment and anti-alignment between the velocity and magnetic-field fluctuations. (Boldyrev 2005, 2006; Beresnyak \\& Lazarian 2006; Mason, Cattaneo, \\& Boldyrev 2006; Matthaeus et~al~2007) These studies examined how the degree of local alignment (and anti-alignment) depends upon length scale, as well as the effects of alignment upon the energy cascade time and the power spectrum of the turbulence.  The topic addressed in this paper is the role of cross helicity in incompressible MHD turbulence. The cross helicity is defined as \\begin{equation} C = \\int d^3\\!x \\; \\bm{v}\\cdot {B}, \\label{eq:defC} \\end{equation} where $\\bm{v}$ is the velocity and~$\\bm{B}$ is the magnetic field. The cross helicity is conserved in the absence of dissipation and can be thought of as a measure of the linkages between lines of vorticity and magnetic field lines, both of which are frozen to the fluid flow in the absence of dissipation (Moffatt 1978).  In the presence of a background magnetic field, $\\bm{B}_0 = B_0 \\hat{z}$, the cross helicity is also a measure of the difference between the energy of fluctuations travelling in the $-z$ and $+ z$ directions.  Dobrowolny, Mangeney, \\& Veltri~(1980) showed that MHD turbulence with cross helicity decays to a maximally aligned state, with $\\delta \\bm{v} = \\pm \\delta\\bm{ B}/\\sqrt{4\\pi\\rho}$, where $\\delta \\bm{v}$ and $\\delta \\bm{B}$ are the fluctuating velocity and magnetic field and $\\rho$ is the mass density. Different decay rates for the energy and cross helicity were also demonstrated by Matthaeus \\& Montgomery~(1980).  In another early study, Grappin, Pouquet, \\& L\\'eorat~(1983) used a statistical closure, the eddy-damped quasi-normal Markovian (EDQNM) approximation, to study strong 3D incompressible MHD turbulence with cross helicity, assuming isotropic power spectra.  They found that when~$C\\neq 0$, the total energy spectrum is steeper than the isotropic Iroshnikov-Kraichnan $k^{-3/2}$~spectrum. Pouquet et~al~(1988) found a similar result in direct numerical simulations of 2D incompressible MHD turbulence.  More recently, Lithwick, Goldreich, \\& Sridhar~(2007) and Beresnyak \\& Lazarian (2007) addressed the role of cross helicity in strong MHD turbulence taking into account the effects of anisotropy.  This paper presents a new phenomenology for strong, anisotropic, incompressible MHD turbulence with nonzero cross helicity, and is organized as follows. Section~\\ref{sec:wpc} presents some relevant theoretical background. Section~\\ref{sec:theory} introduces the new phenomenology as well as two nonlinear advection-diffusion equations that model the time evolution of the power spectra.  Analytic and numerical solutions to this equation in the weak-turbulence and strong-turbulence regimes are presented in Sections~\\ref{sec:weak} and~\\ref{sec:strong}.  Section~\\ref{sec:strong} also presents a simple phenomenological derivation of the power spectra and anisotropy in strong MHD turbulence.  Section~\\ref{sec:trans} presents a numerical solution to the advection-diffusion equation that shows the smooth transition between the weak and strong turbulence regimes. Section~\\ref{sec:unequal} addresses the case in which the parallel correlation lengths of waves propagating in opposite directions along the background magnetic field are unequal at the outer scale.  In Section~\\ref{sec:solarwind}, the proposed phenomenology is applied to turbulence in the solar wind and solar corona, and in Section~\\ref{sec:comp} the results of this work are compared to the recent studies of Lithwick, Goldreich, \\& Sridhar~(2007) and Beresnyak \\& Lazarian~(2007). ",
        "conclusions": "This paper proposes a new phenomenology for strong, anisotropic, incompressible MHD turbulence with cross helicity and introduces a nonlinear advection-diffusion equation [equation~(\\ref{eq:FPpm})] to describe the time evolution of the anisotropic power spectra of the $w^+$ and $w^-$ fluctuations.  It is found that in steady state the one-dimensional power spectra of the energetically dominant $w^+$ fluctuations, $E^+(k_\\perp)$, is steeper than $k_\\perp^{-5/3}$, and that $E^+(k_\\perp)$ becomes increasingly steep as the fractional cross helicity~$\\sigma_c$ increases.  Increasing $\\sigma_c$ also increases the energy cascade time of the $w^+$ fluctuations, reduces the turbulent heating power for a fixed fluctuation energy, and increases the anisotropy of the fluctuations at small scales.  Although most of the discussion has focused on forced, steady-state turbulence, the results of this paper can also be applied to decaying turbulence.  For example, equations~(\\ref{eq:tau1}) and (\\ref{eq:tau2}) can be used to estimate the time scale for turbulence to decay.  The resulting prediction is that if the fluctuations are initially excited with $w_{k_f}^+ \\gg w_{k_f}^-$ and with comparable parallel correlation lengths at the outer scale, then the turbulence will decay into a state in which~$w^-$ fluctuations are absent, as in the earlier work of Dobrowolny, Mangeney, \\& Veltri (1980), Grappin et~al~(1983), and Lithwick \\& Goldreich (2003).  This ``maximally aligned'' state will then be free from nonlinear interactions, and will persist for long times until it damps via linear dissipation."
    },
    "0801/0801.4418_arXiv.txt": {
        "abstract": "We present an analysis of the Suzaku observation of the northeastern rim of the Cygnus Loop supernova remnant.  The high detection efficiency together with the high spectral resolution of the Suzaku X-ray CCD camera enables us to detect highly-ionized C and N emission lines from the Cygnus Loop.  Given the significant plasma structure within the Suzaku field of view, we selected the softest region based on ROSAT observations.  The Suzaku spectral data are well characterized by a two-component non-equilibrium ionization model with different best-fit values for both the electron temperature and ionization timescale.  Abundances of C to Fe are all depleted to typically 0.23 times solar with the exception of O.  The abundance of O is relatively depleted by an additional factor of two compared with other heavy elements.  We found that the resonance-line-scattering optical depth for the intense resonance lines of O is significant and, whereas the optical depth for other resonance lines is not as significant, it still needs to be taken into account for accurate abundance determination. ",
        "introduction": "X-ray spectroscopic imaging observations of supernova remnants (SNRs) enable us to investigate the thermodynamic state of the hot plasma they contain.  The spatial structure of electron temperature ($kT_{\\rm e}$), ionization timescale, and elemental abundance have now been derived for many SNRs with ASCA, Chandra, XMM-Newton, and Suzaku.  The X-ray emission from SNRs is usually assumed to be optically thin since the electron density is typically low (0.1$-$10\\,cm$^{-3}$).  This is safely the case for continuum emission, however the optical depth for the resonance lines cannot be assumed to be negligibly small for bright SNRs, as first pointed out by Kaastra \\& Mewe (1995) in the context of  X-ray observations. Resonance transitions are optically allowed transitions from the ground state of an ion.  In the coronal limit, they are more likely to occur than other transitions, as there is a large population of ions in the ground state. If there is a sufficient ion column density along a particular line of sight, resonance line photons can be scattered out of that line of sight to appear at another location. Resonance line scattering, thus, leads to an underestimate of elemental abundances as well as biases in $kT_{\\rm e}$ determined by line intensity ratios. It should be noted that any random or turbulent velocities will tend to decrease the effects of resonance line scattering.  The Cygnus Loop is one of the brightest SNRs in the soft X-ray band and appears as a rim-brightened shell.  The Cygnus Loop is nearby (540\\,pc; Blair et al.~2005) and has a low neutral H column density (several $\\times 10^{20}$\\,cm$^{-2}$) and large apparent size (2.$\\!\\!^\\circ$5$\\times$3.$\\!\\!^\\circ$5; Levenson et al.~1997), which enable us to study its spatially-resolved soft X-ray emission. Analysis of IUE observations of the Cygnus Loop indicates a high shock velocity (130\\,km s$^{-1}$), departures from steady flow behind the shock, and significant optical depths for the UV resonance lines (Raymond et al.~1981).  Subsequent studies (e.g., Raymond et al.~1988; Cornett et al.~1992) have shown that optical depth effects play an important role in shaping the UV and optical spectral and morphological properties of the Cygnus Loop.  In this paper, we report on the very soft X-ray emission of the northeastern  region of the Cygnus Loop as observed with the Suzaku observatory (Mitsuda et al.~2007).  The extended low energy range of the Suzaku X-ray CCD camera (XIS; Koyama et al.~2007) combined with their superior energy resolution allows us to detect highly ionized C and N emission lines from the northeastern region of the Cygnus Loop (Miyata et al.~2007; paper I, hereafter).  These authors determined the abundances of C, N, and O independently for the first time for this source and found that their relative abundances are roughly consistent with solar values. Throughout the present paper, CGS units are used unless specified otherwise. ",
        "conclusions": "We analyzed the very soft X-ray emission from the northeastern region of the Cygnus Loop with the Suzaku Observatory. The X-ray emitting plasma requires at least two spectral components; a high $kT_{\\rm e}$ component with small ionization timescale and a low  $kT_{\\rm e}$ component with large ionization timescale.  Usually, X-ray emitting plasmas of SNRs are assumed to be optically thin because the typical density is quite low. Since the region we analyzed is very bright with a potentially large line-of-sight path length, however, this assumption may not be valid for some resonance lines. The line-center cross section of resonance scattering can be expressed as \\begin{equation}% \\sigma = {\\sqrt{\\pi} e^2 \\over m c}{f \\over \\nu}\\left({v \\over c}\\right)^{-1} \\simeq 1.86 \\times 10^{-9}\\,{f \\over E}\\,v^{-1} \\quad\\mbox{cm}^2\\ , \\end{equation} where $f$ and $\\nu$ are the oscillator strength and frequency, respectively, of the line transition concerned, $v$ is the root-mean-square kinetic velocity of the ion, $m$ is the electron mass and other quantities have their usual meanings.  In the last expression, $E$ is the line energy in keV. The kinetic velocity is composed of thermal and turbulent motions, as \\begin{equation}% v^2 = \\left({2kT_{\\rm i} \\over m_{\\rm i}}\\right)^2 + \\xi^2\\ , \\end{equation} where $kT_{\\rm i}$ is the ion temperature in keV, $m_{\\rm i}$ is the ion mass, and $\\xi$ represents the  root-mean-square turbulent velocity due to motions other than thermal one.  In the following analysis, we assume that the second term (turbulent motion) is negligibly small compared to the first one (thermal motion) in equation (2); the validity is discussed later.  The line-center optical depth for resonance scattering is given by \\begin{equation}% \\tau = n_z \\sigma L = \\left({n_z \\over n_Z}\\right) \\left({n_Z \\over n_{\\rm H}}\\right) \\left({n_{\\rm H} \\over n_{\\rm e}}\\right)\\,n_{\\rm e} \\sigma L \\ , \\end{equation} where $L$ is the path length through the plasma, $n_z$ the ion density, $n_Z$ the element density, $n_{\\rm H}$ the hydrogen density, and therefore $n_z/n_Z$ represents the ionic fraction, $n_Z/n_{\\rm H}$ the elemental abundance relative to hydrogen. We assume $L = 2.5$\\,pc and the emission volume to be 2.5$\\times$0.93$\\times$2.5\\,pc$^3$. The electron density is obtained from the emission measure to be 1.25 and 1.35\\,cm$^{-3}$ for the high and low $kT_{\\rm e}$ components, respectively, and $n_{\\rm e}/n_{\\rm H}$ is taken to be 1.2.  For the ion temperature, we assume $kT_{\\rm i} = kT_{\\rm e}$ = 0.236\\,keV,  which is the density-weighted mean value from table~\\ref{table:vnei}. In the calculation of $\\sigma$ for the He-like K$\\alpha$ lines, we summed the oscillator strengths for the three main lines of the He-like triplet, namely the forbidden ($z$), intercombination ($x+y$), and resonance ($w$) lines. The ionic fractions of O$\\;${\\scriptsize\\rmfamily{VII}}\\,K$\\alpha$ and O$\\;${\\scriptsize\\rmfamily{VIII}}\\,K$\\alpha$ were calculated by taking into account the NEI condition for each component separately (using Masai 1984).  The density-weighted overall ionic fractions of O$\\;${\\scriptsize\\rmfamily{VII}}\\,K$\\alpha$ and O$\\;${\\scriptsize\\rmfamily{VIII}}\\,K$\\alpha$ were then determined to be 0.51 and 0.33, respectively.  The inferred values of $\\tau$ for O$\\;${\\scriptsize\\rmfamily{VII}}\\,K$\\alpha$ and O$\\;${\\scriptsize\\rmfamily{VIII}}\\,K$\\alpha$ turn out to be 0.54 and 0.16. Table~\\ref{table:optical_depth} summarizes the resonance-line-scattering cross sections and optical depths for the K emission lines detected in the Suzaku data and shown in table~\\ref{table:gaus_fit}.  The optical depth of O$\\;${\\scriptsize\\rmfamily{VII}}\\,K$\\alpha$ is the largest whereas those of the other emission lines are not negligibly small.  In order to model the effect of self-absorption in a simple way, we employed the so-called ``escape-factor'' method (Irons 1979).  We assume the global (source- and direction-average) Doppler-profile escape factor for plane-parallel ``slab'' geometry (Bhatia \\& Kastner 2000) which is a reasonable assumption for the shell-brightened appearance of the Cygnus Loop.  Escape factor values for the relevant emission lines are summarized in table~\\ref{table:optical_depth}. All values are less than 0.87, indicating that our line intensities have been underestimated by a factor of 13\\% or more. Optical depth effects, therefore, play a significant  role in the X-ray emission from the Cygnus Loop with the O abundance, in particular, being underestimated by a factor of 20--40\\%.  However, optical depth effects alone cannot account for the entire abundance discrepancy of a factor of two that we observe.  It should be noted as well that the escape factor is different for O$\\;${\\scriptsize\\rmfamily{VII}}\\,K$\\alpha$ and O$\\;${\\scriptsize\\rmfamily{VIII}}\\,K$\\alpha$, implying that the {\\it true} line intensity ratio differs from the observed value. Correcting the line intensity ratio of O$\\;${\\scriptsize\\rmfamily{VIII}}\\,K$\\alpha$ to O$\\;${\\scriptsize\\rmfamily{VII}}\\,K$\\alpha$ for this effect results in a revised ratio of  0.21, which implies $kT_{\\rm e}=0.15$\\,keV based on the APEC model of equilibrium ionization.  This optical depth effect not only reduces  derived elemental abundances but it also biases fitted $kT_{\\rm e}$ measurements from their {\\it true} values.  Furthermore, varying the elemental abundance affects the calculation of the optical depth as shown in equation (3). Future work will be needed to develop more sophisticated emission codes that take optical depth effects into account.      We have made a few assumptions that tend to enhance the effects of resonance line scattering, including (a) temperature equipartition between electrons and ions, (b) no bulk motions, and (c) no turbulence. The temperature equipartition is a reasonable assumption for the region observed, as argued by Ghavamian et al. (2001) based on the optical spectra as well as by Miyata \\& Tsunemi (1999) on X-rays. Bulk motions may not be such a big problem for the Cygnus Loop due to its relatively low shock speed (few hundred km\\,s$^{-1}$) and its large angular size on the sky. The region observed is right at the edge of the remnant's rim and therefore can be approximated as a slab moving uniformly across the sky.  This would not be the case if there was significant curvature of the remnant shell in the line of sight, or if we were viewing through the center of the remnant where the back part of the shell would be moving away from us while the front part was moving toward us.  So the proximity of the Cygnus Loop to the Earth again favors resonance line scattering.  As for (c), very little information is available.  One location where turbulence likely occurs  is at the contact discontinuity of the ejecta in young SNRs.  The instabilities that operate here drive turbulent eddies and vortices that imprint a random pattern of velocity flows onto the overall radial expansion of the ejecta.  Any such turbulence has an effect of increasing the root-mean-square velocity broadening $v$ resulting in a reduction of the scattering cross section.  The forward shock in the adiabatic phase of remnant evolution is hydrodynamically stable and will not suffer this effect.  But radiative cooling at the forward shock causes hydrodynamical instabilities (Blondin et al. 1998).  Ultimately the amount of line broadening due to turbulence of this type will be some fraction of the shock velocity, so that younger SNRs will tend to have a higher line broadening and hence lower cross-section for resonance line scattering.  If the part of the Cygnus Loop we observed is still in the Sedov phase of evolution, the conditions for minimal turbulence are met.  Therefore, according to all three items discussed above, the Cygnus Loop is favored to have a significant effect due to resonance line scattering.   Since the Cygnus Loop appears to be a rim-brightened shell in the X-ray band, optical depth effects may only play an important role at the rim where the plasma path lengths are largest.   The abundance of O relative to other heavy elements has been show to decrease at the shell region of the Cygnus Loop (e.g. Miyata \\& Tsunemi 1999).  Likewise, Leahy (2004) analyzed Chandra data of the bright southwestern region and found that the abundances of O group elements (C, N, and O) is roughly half those of the Ne group (Ne, Na, Mg, Al, Si, S, and Ar) and Fe group (Ca, Fe, and Ni). The depleted O abundance relative to other elements is consistent with our results and suggests that optical depth effects might be important at the southwestern region too. It should be noted that Miyata \\& Tsunemi (1999) and Leahy (2004) assumed the C and N abundances relative to O to be solar because they had neither high enough spectral resolution nor sufficient effective area for detecting C and N lines.  The Suzaku XIS camera enables us to determine the abundances of C, N, and O independently for diffuse X-ray sources so future additional observations with Suzaku are essential for abundance measurements. Katsuda et al.~(2007) analyzed four Suzaku pointings of the northeastern region of the Cygnus Loop and showed that O is relatively depleted by a factor of two compared  with other heavy elements.  This result also supports our view for a significant optical depth effect. The Suzaku observatory has so far performed 8 mapping observations from the northeastern toward the southwestern region of the Cygnus Loop. Detailed analysis of the radial distribution of intensities of O$\\;${\\scriptsize\\rmfamily{VII}}\\,K$\\alpha$  and O$\\;${\\scriptsize\\rmfamily{VIII}}\\,K$\\alpha$ as well as their intensity ratios may help to clarify optical depth effects in the SNR. However, for this to work it will be necessary to separate out the effects of any intrinsic radial variation in the plasma temperature, which can also produce variations in the observed line intensity ratio.  In clusters of galaxies, an anomalous ratio of Fe K$\\beta$ to K$\\alpha$ emission lines has been observed, which is offered as evidence for resonance line scattering in the relatively dense central cluster plasma (e.g., Tawara 1995).  Unfortunately, due to their NEI condition, the intensity ratio of the  K$\\beta$ to K$\\alpha$ emission line complexes is not a good indicator of resonance line scattering in SNRs. For definitive answers,  we need to resolve individual K-shell lines of He-like ions  as well as the  Ly$\\alpha$ and Ly$\\beta$ lines  for a range of elemental species as  a function of position across the extent of diffuse thermal X-ray sources.  Hopefully in the near future X-ray microcalorimeters onboard $XEUS$, Con-X, and/or $NeXT$  will provide this capability and allow X-ray astronomers to use resonance scattering and other techniques to gain a deeper insight into the nature of SNRs.  \\vspace{1cm}  We are grateful to all the other members of the Suzaku team. EM is supported by  Grant-in-Aid for Specially Promoted Research (16002004) and the 21st Century COE Program, \\lq{\\it Towards a new basic science: depth and synthesis}\\rq. JPH acknowledges support from NASA grant NNG05GP87G. This research 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."
    },
    "0801/0801.4768_arXiv.txt": {
        "abstract": "We present results on the interstellar medium (ISM) properties of 29 galaxies based on a comparison of {\\it Spitzer} far-infrared and Westerbork Synthesis Radio Telescope radio continuum imagery. Of these 29 galaxies, 18 are close enough to resolve at $\\la$1~kpc scales at 70~$\\micron$ and 22~cm. We extend the \\citet{ejm06a,ejm06b} approach of smoothing infrared images to approximate cosmic-ray (CR) electron spreading and thus largely reproduce the appearance of radio images.  Using a wavelet analysis we decompose each 70~$\\micron$ image into one component containing the star-forming {\\it structures} and a second one for the diffuse {\\it disk}. The components are smoothed separately, and their combination compared to a free-free corrected 22~cm radio image; the scale-lengths are then varied to best match the radio and smoothed infrared images. We find that late-type spirals having high amounts of ongoing star formation benefit most from the two-component method. We also find that the disk component dominates for galaxies having low star formation activity, whereas the structure component dominates at high star formation activity.  We propose that this result arises from an age effect rather than from differences in CR electron diffusion due to varying ISM parameters. The bulk of the CR electron population in actively star-forming galaxies is significantly younger than that in less active galaxies due to recent episodes of enhanced star formation; these galaxies are observed within $\\sim$10$^{8}$~yr since the onset of the most recent star formation episode. The sample irregulars have anomalously low best-fit scale-lengths for their surface brightnesses compared to the rest of the sample spirals which we attribute to enhanced CR electron escape. ",
        "introduction": "The interstellar medium (ISM) is a complex environment comprised of a diverse mix of extremely tenuous matter (by terrestrial standards) spanning a range of energetic states. Atomic, ionized, and molecular material make up a series of gaseous ({\\it thermal}) phases  usually categorized by their temperature and density; combinations of these components are used to define the so-called two- and three-phase models of the ISM and variations thereof \\citep[e.g.][and references therein]{cm95}, however, this is not the entire picture.  There is an additional phase of the ISM which is often overlooked due to the difficulties associated with making direct observations of its constituents. This is the relativistic ({\\it non-thermal}) phase which is made up of relativistic charged particles known as cosmic rays (CRs) and magnetic fields. The CRs within galaxies have an energy density comparable to that of the gaseous phases \\citep[e.g.][]{bc90}. They fill up the entire volume of galaxies and are important  sources of heating and ionization of the ISM, though characterizing their propagation remains an ongoing astrophysical problem  \\citep[see][for a review]{smp07}.  The relativistic components of the ISM are important dynamically and may play a significant role in the regulation of star formation during the formation and evolution of galaxies \\citep[e.g.][and references therein]{kf01,cox05}. Magnetic fields both help to support interstellar matter against a galaxy's gravitational potential and confine CRs to galaxy disks; thus, magnetic fields and CRs take part in the hydrostatic balance and stability of the ISM while possibly determining the properties of gas spiral arms \\citep{beck07} and even aiding in the triggering of star formation \\citep{msw74,be82}. For sufficiently large CR pressures Parker instabilities \\citep{ep66} can create breaches in magnetic disks, allowing CRs and interstellar material to freely stream into intergalactic space.  To date, most of our knowledge about the relativistic ISM outside of the Galaxy has been obtained indirectly through the detection of synchrotron emission via multi-frequency radio observations \\citep[e.g.][]{nd91,dlg95,ul96,ji99,rb05}. Synchrotron emission arises from CR electron energy losses as these particles are accelerated in the magnetic fields of galaxies. Although the energy density in CR electrons is only $\\sim$1\\% of that for CR nuclei, the similarity between the spatial distributions of gamma-ray and synchrotron emission within the Galaxy suggests that CR electrons and CR nuclei are fairly well mixed on the scales of a few hundred parsecs \\citep[e.g.][]{ch82,jb86,ww91}. The spatial distribution of a galaxy's synchrotron emission is a function of a galaxy's CR electron and magnetic field distributions. Thus, radio synchrotron maps provide only limited insight on the source distribution of the CR electrons as well as the distances the particles may have traveled before ending up in their current location of emission.  Massive stars ($\\ga$8~$M_{\\sun}$) are the progenitors of supernovae (SNe) whose remnants (SNRs), through the process of diffusive shock acceleration \\citep{ab78,bo78}, appear to be the main acceleration sites of CR electrons responsible for a galaxy's observed synchrotron emission. These same young massive stars are often the primary sources for dust heating as they emit photons which are re-radiated at far-infrared (FIR) wavelengths. This shared origin between the FIR and radio emission of galaxies is thought to be the foundation for the observed FIR-radio correlation among \\citep[e.g.][]{de85,gxh85,sn97,nb97,yun01} and within galaxies \\citep[e.g.][]{bg88,Xu92,mh95,hoer98,hip03,ejm06a,ah06}.   \\begin{deluxetable*}{ccccccccccccc} \\tablecaption{Basic Galaxy Data\\label{tbl-galdat}} \\tablewidth{0pt} \\tabletypesize{\\scriptsize} \\tablehead{ \\colhead{} & \\colhead{R.A.} & \\colhead{Decl.} & \\colhead{$D_{25}$} & \\colhead{} & \\colhead{} & \\colhead{$M_{\\rm opt}$} & \\colhead{$W_{20}$/$W_{20}^{0}$} & \\colhead{$V_{r}$} & \\colhead{Dist.} & \\colhead{$i$} & \\colhead{PA} & \\colhead{Distance}\\\\ \\colhead{Galaxy} & \\colhead{(J2000)} & \\colhead{(J2000)} & \\colhead{(arcmin)} & \\colhead{Type} & \\colhead{Nuc.} & \\colhead{(mag)} & \\colhead{(km s$^{-1}$)} & \\colhead{(km s$^{-1}$)} & \\colhead{(Mpc)} & \\colhead{($\\degr$)} & \\colhead{($\\degr$)} & \\colhead{References}\\\\ \\colhead{(1)} & \\colhead{(2)} & \\colhead{(3)} & \\colhead{(4)} & \\colhead{(5)} & \\colhead{(6)} & \\colhead{(7)} & \\colhead{(8)} & \\colhead{(9)} & \\colhead{(10)}& \\colhead{(11)}& \\colhead{(12)}& \\colhead{(13)} } \\startdata NGC~628  & 1 36 41.7 & +15 46 59 & 10.5$\\times$9.5 & SAc   & \\ldots\t       & -20.9\t& 74\t   /175 & 657 & 7.3 &25 & 25 &  1      \\\\ NGC~925  & 2 27 17.0 & +33 34 43 & 10.5$\\times$5.9 & SABcd & H~\\textsc{II}     & -20.6\t& 224\t   /267 & 553 & 9.1 &57 &102 &  2      \\\\ NGC~2403 & 7 36 51.4 & +65 36 09 & 21.9$\\times$12.3& SABcd & H~\\textsc{II}     & -19.7\t& 257\t   /306 & 131 & 3.2 &57 &127 &  2      \\\\ Holmb~II & 8 19 04.0 & +70 43 09 &  7.9$\\times$6.3 & Im    & \\ldots\t       & -17.1\t& 73\t   /121 & 157 & 3.4 &37 & 15 &  3      \\\\ NGC~2841 & 9 22 02.6 & +50 58 35 &  8.1$\\times$3.5 & SAb   & Lin/Sy1\t       & -20.7\t& 611\t   /664 & 638 &14.1 &67 &147 &  4      \\\\ NGC~2976 & 9 47 15.3 & +67 55 00 &  5.9$\\times$2.7 & SAc   & H~\\textsc{II}     & -17.6\t& \\ldots/\\ldots &   3 & 3.6 &64 &143 &  3      \\\\ NGC~3031 & 9 55 33.2 & +69 03 55 & 26.9$\\times$14.1& SAab  & Lin\t       & -21.2\t& 446\t   /515 & -34 & 3.6 &60 &157 &  2      \\\\ NGC~3184 &10 18 16.9 & +41 25 28 &  7.4$\\times$6.9 & SABcd & H~\\textsc{II}     & -19.0\t& 142\t   /396 & 592 &11.1 &21 &135 &  5      \\\\ NGC~3198 &10 19 54.9 & +45 32 59 &  8.5$\\times$3.3 & SBc   & \\ldots\t       & -20.2\t& 318\t   /343 & 663 &13.7 &68 & 35 &  2      \\\\ IC~2574  &10 28 21.2 & +68 24 43 & 13.2$\\times$5.4 & SABm  & \\ldots\t       & -17.7\t& 123\t   /134 &  57 & 4.0 &67 & 50 &  6      \\\\ NGC~3627 &11 20 15.0 & +12 59 30 &  9.1$\\times$4.2 & SABb  & Sy2\t       & -20.8\t& 378\t   /417 & 727 & 9.4 &65 &173 &  2      \\\\ NGC~3938 &11 52 49.5 & +44 07 14 &  5.4$\\times$4.9 & SAc   & \\ldots\t       & -20.1\t& 112\t   /265 & 809 &13.3 &25 &  0 &  7      \\\\ NGC~4125 &12 08 05.8 & +65 10 27 &  5.8$\\times$3.2 & E6p   & \\ldots\t       & -21.6\t& \\ldots/\\ldots &1356 &22.9 &58 & 95 &  8      \\\\ NGC~4236 &12 16 42.1 & +69 27 46 & 21.9$\\times$7.2 & SBdm  & \\ldots\t       & -18.1\t& 176\t   /185 &   0 & 4.5 &72 &162 &  3      \\\\ NGC~4254 &12 18 49.5 & +14 24 59 &  5.4$\\times$4.7 & SAc   & \\ldots\t       & -21.6\t& 272\t   /544 &2407 &16.6 &30 &  0 &  7$^{*}$\\\\ NGC~4321 &12 22 54.9 & +15 49 21 &  7.4$\\times$6.3 & SABbc & Lin\t       & -22.1\t& 283\t   /534 &1571 &14.3 &32 & 30 &  8      \\\\ NGC~4450 &12 28 29.5 & +17 05 06 &  5.2$\\times$3.9 & SAab  & Lin\t       & -21.4\t& 290\t   /433 &1954 &16.6 &42 &175 &  7$^{*}$ \\\\ NGC~4552 &12 35 39.8 & +12 33 23 &  5.1$\\times$4.7 & E     & \\ldots\t       & -20.8\t& \\ldots/\\ldots & 340 &15.9 &23 &  0 &  9      \\\\ NGC~4559 &12 35 57.7 & +27 57 36 & 10.7$\\times$4.4 & SABcd & H~\\textsc{II}     & -21.0\t& 251\t   /273 & 816 &10.3 &67 &150 &  7      \\\\ NGC~4569 &12 36 49.8 & +13 09 46 &  9.5$\\times$4.4 & SABab & Lin/Sy\t       & -22.0\t& 360\t   /397 &-235 &16.6 &65 & 23 &  7$^{*}$\\\\ NGC~4631 &12 42 08.0 & +32 32 26 & 15.5$\\times$2.7 & SBd   & \\ldots\t       & -20.6\t& 320\t   /322 & 606 & 7.7 &83 & 86 &  10     \\\\ NGC~4725 &12 50 26.6 & +25 30 06 & 10.7$\\times$7.6 & SABab & Sy2\t       & -22.0\t& 410\t   /570 &1206 &11.9 &46 & 35 &  2      \\\\ NGC~4736 &12 50 53.0 & +41 07 14 & 11.2$\\times$9.1 & SAab  & Lin\t       & -19.9\t& 241\t   /400 & 308 & 5.0 &37 &105 &  8      \\\\ NGC~4826 &12 56 43.7 & +21 40 52 & 10.0$\\times$5.4 & SAab  & Sy2\t       & -20.3\t& 311\t   /363 & 408 & 5.0 &59 &115 &  7      \\\\ NGC~5033 &13 13 27.5 & +36 35 38 & 10.7$\\times$5.0 & SAc   & Sy2\t       & -20.9\t& 446\t   /501 & 875 &14.8 &63 &170 &  7      \\\\ NGC~5055 &13 15 49.2 & +42 01 49 & 12.6$\\times$7.2 & SAbc  & H~\\textsc{II}/Lin & -19.0\t& 405\t   /489 & 504 & 7.8 &56 &105 &  7      \\\\ NGC~5194 &13 29 52.7 & +47 11 43 & 11.2$\\times$6.9 & SABbc & H~\\textsc{II}/Sy2 & -21.4\t& 195\t   /244 & 463 & 7.8 &53 &163 &  7      \\\\ NGC~6946 &20 34 52.3 & +60 09 14 & 11.5$\\times$9.8 & SABcd & H~\\textsc{II}     & -21.3\t& 242\t   /457 &  48 & 6.8 &32 & 69 &  11     \\\\ NGC~7331 &22 37 04.1 & +34 24 56 & 10.5$\\times$3.7 & SAb   & Lin               & -21.8\t& 530\t   /561 & 816 &14.5 &71 &171 &  2  \\enddata \\tablecomments{Col. (1): ID. Col. (2): The right ascension in the J2000.0 epoch. Col. (3): The declination in the J2000.0 epoch. Col. (4): Major- and minor-axis diameters. Col. (5): RC3 type. Col. (6): Nuclear type: H~\\textsc{II}: H~\\textsc{II} region; Lin: LINER; Sy: Seyfert (1, 2). Col. (7): Absolute R magnitude, when available; otherwise from the V or B bands. Col. (8): Observed/inclination-corrected 21 cm neutral hydrogen line width at 20\\% of maximum intensity, in km~s$^{-1}$, taken from \\citet{tul88} or RC3.  Col. (9): Heliocentric velocity. Col. (10): Distance in Mpc Col. (11): Inclination in degrees. Col. (12): Position Angle in degrees. (13): Distance References}  \\tablerefs{(1)~\\citet{sfe01}; (2)~\\citet{wf01}; (3)~\\citet{ik02}; (4)~\\citet{lm01}; (5)~\\citet{dl02}; (6)~\\citet{ik03}; (7)~K. Masters 2007, in preparation: ($^{*})$ indicates distance set to Virgo Cluster center; (8)~\\citet{jt01}; (9)~\\citet{lf00}; (10)~\\citet{as05}; (11)~\\citet{ik00}} \\end{deluxetable*}  Coupling the shared origin of a galaxy's FIR and radio emission with the fact that the mean free path of dust-heating photons ($\\sim$100~pc) is significantly shorter than the expected diffusion length of CR electrons ($\\sim$1-2~kpc) led \\citet{bh90} to conjecture that the radio image of a galaxy should resemble a smoothed version of its infrared image. Consequently, it appears that the close spatial correlation between the FIR and radio continuum emission within galaxies can be used to characterize the propagation history of CR electrons. This prescription has been shown to hold for galaxies observed at the ``super resolution'' ($\\la$1$\\arcmin$) of {\\it IRAS} HIRES data \\citep{mh98} and, more recently, for high resolution ($\\sim$18$\\arcsec$) {\\it Spitzer} 70$~\\micron$ imaging \\citep[][hereafter M06a]{ejm06a}. This phenomenology has been further corroborated on scales $\\ga$50~pc by \\citet{ah06} who find synchrotron haloes around individual star-forming regions are more extended than FIR-emitting regions within the Large Magellanic Cloud.  \\citet[][hereafter, M06b]{ejm06b} recently studied how the spatial distributions of a galaxy's FIR and radio emission vary as a function of the intensity of star formation. They concluded that CR electrons are, on average, younger and closer to their place of origin within galaxies having higher amounts of star formation activity compared with more quiescent galaxies. Using a wavelet-based image decomposition, we extend this work by attempting to characterize separately CR electron populations associated with a galaxy's diffuse disk and its star-forming complexes. We carry out this study for a sample of galaxies observed as part of the {\\it Spitzer} Infrared Nearby Galaxies Survey \\citep[SINGS;][]{rk03} and the  Westerbork Synthesis Radio Telescope \\citep[WSRT-SINGS;][]{rb07} for which we have the spatial resolution to resolve physical scales $<1$~kpc.  The paper is organized as follows: The galaxy sample is defined in $\\S$2 while the observations and data reduction techniques are discussed in $\\S$3. In $\\S$4 we present and discuss a correlation analysis between FIR/radio ratios and various other physical quantities on sub-kiloparsec scales within galaxies; this section leads us to confirm quantitatively that our favored phenomenological model is the best description for the FIR-radio correlation within galaxies. In $\\S$5 we describe our wavelet-based, two-component image-smearing analysis; the corresponding results are then presented in $\\S$6 and their physical implications are discussed in $\\S$7. We briefly discuss outstanding issues and future prospects in $\\S$8 and summarize our results and conclusions in $\\S$9. ",
        "conclusions": "} Using a two-component image-smearing analysis, we have separated the signatures of CR electron diffusion at spatial scales corresponding to star-forming structures ($<$1~kpc) and galaxy disks ($\\geq$1~kpc) within 18 galaxies observed as part of SINGS and WSRT-SINGS. Our results and conclusions can be summarized as follows: \\begin{enumerate}  \\item We confirm and extend earlier results of M06a,b. Empirically, the dispersion in the FIR-radio correlation within galaxies is most reduced by an image-smearing model; this improvement is significantly better than what can be achieved by fitting correlations and removing linear trends.  \\item The best-fit global scale-lengths decrease as a function of increasing star formation activity as measured by the infrared surface brightness of a galaxy. Our interpretation is that a galaxy's CR electrons are closer to their place of origin within galaxies having intense star formation activity.  \\item The trend of decreasing best-fit {\\it global} scale-length with increasing radiation field energy density is due to higher surface brightness galaxies having undergone a recent enhancement of star formation rather than variations in other ISM parameters. For sufficiently large enhancements, these galaxies are observed within $\\sim$10$^{8}$~yr of the onset of the most recent star formation episode.   \\item Unlike spirals, irregular galaxies lack any well defined diffuse disk component at either 70~$\\micron$, or especially at 22~cm. Presumably, the CR electrons escape these galaxies soon after leaving their parent star-forming regions due the absence of a dense ISM which would keep large-scale interstellar magnetic field locked into place. This conclusion helps to explain why these galaxies have global FIR/radio ratios systematically greater than canonical values.   \\item As infrared surface brightness increases, the characteristic diffusion scale-length of a galaxy's CR electron population begins to transition at \\(\\log~U_{\\rm rad} \\leq -12.5\\), or \\(\\log~\\Sigma_{\\rm SFR} \\leq -2.3\\), from being biased by CR electrons making up its diffuse disk to being biased by those recently injected near star-forming structures. From this we conclude that a galaxy's CR electron population transitions from being dominated by old CR electrons to being dominated by young CR electrons as a function of star formation intensity.  \\item The two-component analysis works better than smearing with a single smoothing kernel for spiral galaxies of type Sb or later which have high amounts of ongoing star formation activity (i.e. $\\sim$40\\% of the sample). This result suggests that star formation must be intense and highly structured for the two-component analysis of these data to differentiate properly between the different CR electron populations.   \\end{enumerate}"
    },
    "0801/0801.1651_arXiv.txt": {
        "abstract": " ",
        "introduction": "Many studies have showcased the great potential of the LHC for producing and discovering supersymmetric particles \\cite{lhc,cmstdr,LHC}, and the ability of experiments at a linear $e^+ e^-$ collider (LC) to measure sparticle properties in detail, if their pair-production thresholds lie within its kinematic reach \\cite{ilc}. Most of these studies have assumed that $R$ parity is conserved, in which case the lightest supersymmetric particle (LSP) may provide the cold dark matter postulated by astrophysicists and cosmologists \\cite{EHNOS}. Further, most studies have been within the framework of the minimal supersymmetric extension of the Standard Model (MSSM) \\cite{mssm}, and assumed that the LSP is the lightest neutralino $\\chi$. We also adopt this framework in this paper. In this case, the classic signature of sparticle pair production is missing energy carried away by the dark matter particles $\\chi$. Studies have indicated that experiments at the LHC should be able to detect gluinos and squarks weighing up to $\\sim 2.5$~TeV \\cite{Baer}, whereas any sparticles weighing less than the beam energy should be detectable at a LC.  One specific supersymmetric version of this framework that has commonly been examined is the Constrained MSSM (CMSSM)~\\cite{funnel,cmssm,efgosi,eoss,cmssmwmap}, in which the soft supersymmetry-breaking mass parameters are assumed to be universal at some high scale, generally taken to be the supersymmetric GUT scale, $M_{GUT} \\sim 10^{16}$ GeV.  Within the CMSSM, renormalization group equations (RGEs) can be used to calculate the weak-scale observables in terms of four continuous and one discrete parameter; the scalar mass, $m_0$, the gaugino mass, $m_{1/2}$, and the trilinear soft breaking parameter, $A_0$ (each specified at the universality scale), as well as the ratio of the Higgs vevs, $\\tbt$, and the sign of the Higgs mixing parameter, $\\mu$. The reaches of colliders such as the LHC or a LC are then often expressed in the $(m_{1/2}, m_0)$ plane for representative values of $A_0, \\tbt$ and the sign of $\\mu$.  However, the mechanism of supersymmetry breaking is not known, and alternative scenarios should also be considered. Rather than postulate that the soft supersymmetry-breaking parameters are universal at some GUT scale, one might consider theories in which this universality assumption for the the soft supersymmetry-breaking parameters is relaxed. One possibility, motivated to some extent by supersymmetric GUT scenarios and the absence of flavour-changing interactions due to sparticle exchanges, would be to relax (for example) the universality assumption for the soft supersymmetry-breaking contributions to the Higgs scalar masses at the GUT scale (the NUHM) \\cite{nonu,nuhm}, and more radical abandonments of universality could also be considered.  We consider here a different generalization of the CMSSM, in which universality of the soft supersymmetry-breaking mass parameters is maintained, but is imposed at some lower input scale $M_{in} < M_{GUT}$ \\cite{eos1,eos2}. Such GUT-less (or sub-GUT) scenarios may arise in models where the dynamics that breaks or communicates supersymmetry breaking to the observable sector has an intrinsic scale below $M_{GUT}$, and switches off at higher scales, much as the effective dynamical quark mass in QCD switches off at scales $> \\Lambda_{QCD}$. Mirage unification scenarios~\\cite{mixed} offer one class of examples in which the low-energy evolution of the gaugino masses is as if they unified at some scale $< M_{GUT}$. In principle, one could consider scenarios in which universality is imposed on the different MSSM soft supersymmetry breaking parameters $m_{1/2}, m_0$ and $A_0$ at different input scales $M_{in}$. However, here we follow~\\cite{eos1,eos2} in studying the simplest class of GUT-less scenarios with identical $M_{in}$ for all the soft supersymmetry-breaking parameters.  As one would expect, the reduction in the universality scale has important consequences for the low-energy sparticle mass spectrum. In particular, the hierarchy of gaugino masses familiar in the GUT-scale CMSSM is reduced with, for example, a substantial reduction in the ratio of gluino and bino masses. Likewise, squark and slepton masses also approach each other as $M_{in}$ is reduced. These effects have important consequences for the $(m_{1/2}, m_0)$ planes in GUT-less scenarios: for example, the boundaries imposed by the absence of a charged ${\\tilde \\tau_1}$ LSP and the generation of an electroweak symmetry breaking vacuum approach each other as $M_{in}$ decreases.  A corollary of the `squeezing' of the sparticle mass spectrum is the observation made in \\cite{eos1} and \\cite{eos2} that, as the universality scale $M_{in}$ is decreased from the GUT scale, there are dramatic changes in the cosmological constraint imposed on the parameter space by the relic density of neutralinos inferred from WMAP and other observations~\\cite{WMAP}. In general, as $M_{in}$ decreases, the regions where the relic neutralino LSP density falls within the range preferred by WMAP and other measurements~\\cite{WMAP} tend to move to larger $m_{1/2}$ and $m_0$. This implies that, whereas in the GUT-scale CMSSM the relic neutralino is {\\it overdense} in most of the region with $m_{1/2}, m_0 < 1$~TeV, as $M_{in}$ decreases to $\\sim 10^{11}$~GeV most of this region becomes {\\it underdense}.  In this paper, we consider the implications of these observations for the prospects for sparticle detection at the LHC and a LC. ATLAS and CMS have estimated their reaches in inclusive supersymmetry searches for multiple jets and missing transverse energy, as functions of the accumulated and analyzed LHC luminosity, which may be expressed as reaches for gluino and squark masses~\\cite{cmstdr}. These may in turn be converted into the reaches in the $(m_{1/2}, m_0)$ planes for different values of $M_{in}$. The masses of weakly-interacting sparticles such as sleptons, charginos and neutralinos are determined across these $(m_{1/2}, m_0)$ planes, and hence the ATLAS/CMS reaches may be converted into the corresponding sparticle pair-production thresholds at a generic LC. These converted reaches may be interpreted in at least two ways. If the LHC {\\it does discover} supersymmetry, then one may estimate, within the CMSSM or any given GUT-less model, the {\\it maximum} centre-of-mass energy that would suffice for a LC to make detailed follow-up measurements of at least some sparticles. Conversely, if the LHC {\\it does not discover} supersymmetry within a given physics reach, one can, within the CMSSM or any given GUT-less model, estimate the {\\it minimum} centre-of-mass energy below which a LC would not provide access to any sparticles. In general, because of the `squeezing' of the sparticle mass spectrum as $M_{in}$ {\\it decreases}, for any given LHC physics reach the required LC centre-of-mass energy {\\it increases} correspondingly.  This argument can be carried through whether one disregards the cosmological density of dark matter entirely, or regards it solely as an upper limit on the relic LSP density, or interprets it as a narrow preferred band. In the third case, the prospects for sparticle detection at the LHC recede with the preferred dark matter regions in the $(m_{1/2}, m_0)$ planes as $M_{in}$ decreases. Within the specific preferred dark-matter regions, the relation between the LHC and LC reaches can be made more precise. For example, if the LHC discovers sparticles with 1~fb$^{-1}$ of data, within the CMSSM a centre-of-mass energy of 600~GeV would suffice for  a LC to to produce pairs of neutralinos, if they provide the cold dark matter, whereas over 1 TeV might be required in a GUT-less model with $M_{in} > 10^{11.5}$~GeV. These required energies increase to 800~GeV in the CMSSM and 1.4~TeV in GUT-less models with $M_{in} > 10^{11.5}$~GeV if the LHC requires 10~fb$^{-1}$ to discover supersymmetry. ",
        "conclusions": "We have discussed in the previous Section how much centre-of-mass energy would be required to `guarantee' the observability of sparticle pair production in $e^+ e^-$ collisions under various hypotheses for the integrated luminosity required for discovering supersymmetry at the LHC and for different values of the universality scale $M_{in}$. We have also discussed how corresponding sparticle exclusions at the LHC would set lower limits on the possible thresholds for producing different sparticle pairs at a LC. To conclude, we now consider the capabilities of LCs with various specific proposed centre-of-mass energies.  Even if supersymmetry were to be found at the LHC with 1~fb$^{-1}$ of integrated luminosity, a LC with $E_{cm} = 0.5$~TeV would not be `guaranteed' to produce $\\chi \\chi$ pairs or other sparticle pairs. However, even if supersymmetry were to be excluded in the LHC's 1~fb$^{-1}$ discovery region, the possibility of observing sparticles at a LC with $E_{cm} = 0.5$~TeV could not be excluded for $M_{in} > 10^{13.5}$~GeV, and such a LC might also pair-produce charged sparticles if $M_{in} > 10^{15}$~GeV and/or produce $\\chi \\chi_2$ in association if $M_{in} > 10^{14.5}$~GeV. On the other hand, if supersymmetry were not even within the 10~fb$^{-1}$ discovery reach of the LHC, a LC with $E_{cm} = 0.5$~TeV might be (barely) above the $\\chi \\chi$ threshold only if $M_{in} \\gtrsim 10^{15.5}$~GeV, and there would be no likelihood of charged-sparticle or $\\chi \\chi_2$ production.  A LC with $E_{cm} = 1$~TeV would be `guaranteed', if supersymmetry were to be found at the LHC with 1~fb$^{-1}$ of integrated luminosity, to produce $\\chi \\chi$ pairs in any GUT-less scenario with $M_{in} > 10^{12}$~GeV. Analogous `guarantees' for charged-sparticle pair production or associated $\\chi \\chi_2$ production could be given only for $M_{in} > 10^{13} (10^{14})$~GeV, respectively. On the other hand, if supersymmetry were not even within the 10~fb$^{-1}$ discovery reach of the LHC, it might still be possible to find $\\chi \\chi$ (charged-sparticle pairs) ($\\chi \\chi_2$) at a LC if $M_{in} > 10^{12.5} (10^{13.3}) (10^{13})$~GeV.  Finally, even if the LHC would require 10~fb$^{-1}$ to discover supersymmetry, a LC with $E_{cm} = 1.5$~TeV would be `guaranteed' to produce $\\chi \\chi$ and $\\chi \\chi_2$ pairs in all the allowed WMAP-compatible scenarios, and charged-sparticle pair production would be `guaranteed' for all except a small range of $M_{in}$ between $10^{12}$ and $10^{13}$~GeV. Hence, a LC with $E_{cm} = 1.5$~TeV would be well matched to the physics reach of the LHC with this luminosity, whereas a LC with a lower $E_{cm}$ might well be unable to follow up on a discovery of supersymmetry at the LHC. However, as already mentioned, even in the absence of any `guarantee', it could still be that the LHC discovers supersymmetry at some mass scale well below the limit of its sensitivity with 10~fb$^{-1}$ of integrated luminosity, in which case a lower-energy LC might still have interesting capabilities to follow up on a discovery of supersymmetry at the LHC.  It is clear that the physics discoveries of the LHC will be crucial for the scientific prospects of any future LC. Supersymmetry is just one of the scenarios whose prospects at a LC may depend on what is found at the LHC. Even within the supersymmetric framework, there are many variants that should be considered. Even if $R$ parity is conserved, the LSP might not be the lightest neutralino. Even if it is, the relevant supersymmetric model may not be minimal. Even if it is the MSSM, supersymmetry breaking may not be universal. Even if it is, the universality scale may not be the same for gauginos and sfermions. Nevertheless, we hope that study serves a useful purpose in highlighting some of the issues that may arise in guessing the LC physics prospects on the basis of LHC physics results."
    },
    "0801/0801.3462_arXiv.txt": {
        "abstract": "We consider a magnetic Bianchi I braneworld, embedded in between two Schwarzschild-AdS spacetimes, boosted equal amounts in opposite directions and compare them to the analagous solution in four-dimensional General Relativity. The efficient dissipation of anisotropy on the brane is explicitly demonstrated, a process we dub braneworld isotropization. From the bulk point of view, we attribute this to anisotropic energy being carried into the bulk by hot gravitons leaving the brane. From the brane point of view this can be interpreted in terms of the production of particles in the dual CFT. We explain how this result  enables us to gain a better understanding  of the behaviour of anisotropic branes already studied in the literature.  We also show how there is evidence of particles being over-produced, and comment on how this may ultimately provide a possible observational signature of braneworlds. ",
        "introduction": "The most recent data from WMAP suggests that the Cosmic Microwave Background (CMB) is isotropic to within one part in $10^5$~\\cite{wmap}. Why is our universe {\\it so} isotropic? This has been an important question in cosmology ever since the CMB's discovery back in the mid 1960s~\\cite{cmb}. The obvious way to approach this is to assume that the universe began in a highly {\\it an}isotropic state, and ask what dynamical mechanism caused the universe to shed nearly all its anisotropy, leaving behind the highly symmetric state we observe today.  The most popular mechanism is, of course, inflation~\\cite{inflation}. In inflation, almost all anisotropy was diluted away during a period of accelerated expansion in the early universe. The accelerated expansion was driven by a scalar field rolling slowly down its potential,  closely resembling a positive cosmological constant. Indeed, it can be shown that in the presence of a positive cosmological constant, all but one of the Bianchi models\\footnote{The exception is  Bianchi IX.}, describing a homogeneous but anisotropic cosmology, evolve asymptotically to an isotropic de Sitter solution~\\cite{wald}. However, there is still an ongoing discussion on how fine-tunned the initial conditions should be to get a sufficiently flat, homogeneous and isotropic universe after inflation (see \\cite{carrol} and references therein).  It is often easy to forget that before inflation, plenty of other ``isotropization''mechanisms were put forward, with varying degrees of success. The earliest was probably the ``phoenix'', or ``oscillatory'',  universe~\\cite{phoenix}, which expands for a while, recollapses towards a bounce, and then starts to expand again. As entropy builds up after each oscillation, it was argued that the period of oscillation increased, and the  universe approached an isotropic, homogeneous state of thermal equilibrium. Although these ideas ran into problems with  inevitable singularities~\\cite{sings}, they have been revived recently with the advent of the cyclic universe~\\cite{cyclic} and a possible resolution to its singularity~\\cite{niz}.  Another means of damping anisotropy is particle production due to quantum effects at very early times~\\cite{pp}. Energetically, this can be understood as anisotropy being converted into thermal radiation. If, for example, the universe were contracting along one direction, but expanding along the others, the momenta of virtual particles would be blueshifted along the contracting direction. This would make the virtual particles more likely to materialize as real particles, drawing energy from the contraction, and lowering the amount of anisotropy.  However, if we assume an arbitrarily large amount of initial anisotropy, a huge amount of thermal radiation would need to be produced to account for the levels of isotropy seen in the CMB, and as a result, the photon-baryon ratio would massively exceed its observed value~\\cite{bm}. Particle production alone cannot, therefore, account for enough dissipation, and we typically assume it to be one of many contributing factors, including neutrino viscosity~\\cite{viscosity}, and, inevitably, inflation.  In this paper, we point out another  mechanism for dissipating anisotropy: {\\it braneworld isotropization}. Consider a Randall-Sundrum type braneworld~\\cite{rs}, and assume it is highly anisotropic. As the brane evolves, the anisotropic energy on the brane leaks into the bulk, and the brane becomes rapidly more isotropic, whilst the bulk becomes { less} homogeneous and/or { less} isotropic. To get a feel for how this works consider the projection of the Einstein equation onto a Randall-Sundrum braneworld~\\cite{sms}, schematically given by \\be G_\\mn(\\gamma)=8 \\pi G_4 T_\\mn+G_5^2(T_\\mn)^2-E_\\mn, \\ee where $G_4$ and $G_5$ are the brane and bulk Newton's constants respectively, and  $T_\\mn$ is the energy-momentum tensor for matter  on the brane.  Note that we have a non-local piece, $E_\\mn$, which is the electric part of the  bulk Weyl tensor, projected onto the brane.  The ``anisotropic energy density'' on the brane, in the sense described in~\\cite{viscosity},   is  stored in the local geometrical piece, $G_\\mn$, and as the brane evolves, this energy is transferred into the bulk through the bulk Weyl term, $E_\\mn$.  The contribution of the bulk Weyl term is often best understood holographically (see, for example~\\cite{bwholo3}) using the braneworld version of the AdS/CFT correspondence~\\cite{adscft, bwholo1}. This states that {\\it  Randall-Sundrum braneworld gravity is dual to a CFT, with a UV cut-off, coupled to gravity in $3+1$ dimensions}. It follows that a {\\it classical} source on the brane behaves like a {\\it quantum} source in $3+1$ dimensons, in the sense that it has been dressed with quantum corrections from the CFT~\\cite{bwholo2}.  The anisotropy dissipation process we have just described can now be interpreted as particle production in the CFT. Energy is drawn from the anisotropy to fuel the production of CFT particles, leading to isotropization on the brane.  There is already evidence in the literature for braneworld isotropization. Consider, for example, the quest to find anisotropic geometries supporting a perfect fluid. This is easy enough in four-dimensional General Relativity, the Bianchi or Kantowski-Sachs metrics being good examples~\\cite{kantowski, exacsols}. In brane cosmology, full solutions have not been so easy to find~\\cite{maar1, orsaylot, frolov}, and only a rather contrived set of kasner-like solutions are known exactly~\\cite{orsaylot, frolov}. Indeed, one can prove that if the bulk is static, an anisotropic brane cannot support a perfect fluid~\\cite{orsaylot}. This result can now be  easily understood: particle production in the CFT leads to  dissipation of anisotropy over time. From the $5D$ point of view, the leaking of anisotropy into the bulk prevents it from being static. Another interesting example of braneworld isotropization lies in the study of anisotropy dissipation in braneworld inflation~\\cite{maar2}. There it was noticed that inflation begins sooner than it would in $4D$ General Relativity. We attribute this to particle production in the CFT drawing its energy from the anisotopy, aswell as the kinetic energy of the inflaton.  In this paper, we demonstrate braneworld isotropization explicitly by means of a concrete example. We generate the initial anisotropy on a Bianchi I braneworld by means of a constant magnetic field. Unlike much of the existing literature on anisotropic braneworlds  the full bulk solution is known, corresponding to the planar limit of two boosted Schwarzschild-AdS black hole spacetimes, cut and pasted together to form the brane in the usual way. Because the brane embedding is highly non-trivial when viewed in global coordinates, it is  convenient to work in a boosted coordinate system on either side of the brane, so that the bulk metric resembles the planar limit of two particular Myers-Perry-AdS spacetimes~\\cite{MP,HHT}. The boosts are equal and opposite on either side of the brane, so one may think of the corresponding Myers-Perry solutions as having equal and opposite angular momentum. We then track the brane's evolution, paying particular attention to the anisotropic shear, and compare it to the corresponding scenario with the same source in four-dimensional General Relativity. As expected, the shear anisotropy dissipates far quicker in the Randall-Sundrum braneworld scenario, than in $4D$ GR.  A study of magnetic fields on a braneworld is also important in its own right. One of the great puzzles in  astrophysics concerns the origin of large magnetic fields in galaxies and galaxy clusters (for an excellent review of magnetic fields in the early universe, see~\\cite{magrev}). It has been argued that these fields were generated by a galactic dynamo process~\\cite{dynamo}, sourced by a smaller pre-galactic ``seed'' field. However, one still has to account for the origin of the seed field, and even then the efficiency of this process has been brought into question~\\cite{dynamobad}. We are naturally led to consider the possibility that galatic fields result from large scale primordial fields left over from the big bang.  In the standard cosmology, one can place bounds on the size of a large scale primordial magnetic field from nucleosynthesis ($B \\lesssim 10^{-7}$ G)~\\cite{BBNbound}, and from CMB temperature fluctuations ($B \\lesssim 10^{-9}$ G)~\\cite{CMBbound}. Given that the CMB bound is the stronger, we might speculate that this bound is weakened in a braneworld context, due to isotropization.  Magnetic fields on branes have been considered in the past~\\cite{magbrane, ecc}, although a full knowledge of the bulk has been absent. Whilst it is true that one can certainly learn a lot {\\it without} full knowledge of the bulk, we believe it is a dangerous game to play. It is not at all obvious that a bulk solution that is regular near the brane evolves into something regular far from the brane. A good example of this is the black string solution~\\cite{bwbh} which is regular near the brane, but singular on the AdS horizon.  The rest of this paper is organised as follows. In section~\\ref{sec:4dgr}, we review the Bianchi I solution for a constant magnetic field in four-dimensional General Relativity. For completeness, and in keeping with supernovae observations~\\cite{supernova}, we will allow for  a positive cosmological constant that can be set to zero if necessary. In section~\\ref{sec:bw}, we introduce a certain Myers-Perry-AdS black hole in five dimensions, and take the planar limit, so that we end up with planar Schwarzschild-AdS in boosted coordinates. We then cut and paste two equal and oppositely boosted black holes onto one another, in the usual way, in order to form a Bianchi I braneworld, supported by brane tension, and a constant magnetic field.  In section~\\ref{sec:compare} we compare the two scenarios, with special emphasis on the  dissipation of anisotropy, and the effect of particle production in the CFT. Finally, in section~\\ref{sec:discuss}, we summarize our results.  Our conventions follow the Landau-Lifshitz notation for the curvature tensors, and we use a $(-,+,...,+)$ signature for the metric. ",
        "conclusions": "\\lab{sec:discuss} In this paper, we have explicitly demonstrated the existence of a new phenomenon, called {\\it braneworld isotropization}, by means of a concrete example. We were able to embed a magnetic Bianchi I braneworld in between two Schwarzschild-$AdS_5$ spacetimes, boosted equal amounts in opposite directions. The magnetic field breaks isotropy on the brane, and leads to the production of thermal graviton radiation, that can carry energy-momentum into the bulk. The boosted coordinates mean that the brane observer sees equal and opposite momentum carried into the bulk by the  hot gravitons leaving the brane. We can view braneworld isotropization as anisotropic energy on the brane being dumped into the bulk.  From a completely $4D$ point of view, we can understand this effect using the AdS/CFT correspondence. In the Randall-Sundrum scenario, the gravitational physics in the bulk is equivalent to a strongly coupled CFT, cut-off in the UV, and minimally coupled to gravity on the brane. When isotropy is broken, the coupling of the CFT to gravity leads to particle production, the required  energy being drawn from the anisotropy.  This phenomenon can now be used to readily explain a number of existing results in the literature~\\cite{maar1, maar2, orsaylot, frolov}. For example, an  anisotropic brane cannot support a perfect fluid in a static bulk~\\cite{orsaylot},  because the leaking of anisotropy off the brane prevents the bulk from being static, at least when viewed by an asymptotic braneworld observer.  Note that in the solution described in this paper, we do not have a perfect fluid on the brane, and although the bulk is static according to an asymptotic observer using Schwarzschild time, it is {\\it not} static according to an asymptotic observer using the boosted time coordinate. It is the latter time coordinate that is used by the asymptotic braneworld observer at large  $r(\\tau)$.  It is worth noting that we have been able to prove that the solution presented here, along with the known isotropic solutions~\\cite{bcg}, are the most general solutions for a Bianchi I brane of the form (\\ref{metric4d}), embedded in Schwarzschild-$AdS_5$ with some form of ``sensible'' matter on the brane. We have not included the proof since it is lengthy, and not particularly illuminating. However, it does tell us that we will have to work a bit harder if we want to, say, include an arbitrary perfect fluid on the brane in addition to the magnetic field. This is a subject for future study. Nevertheless, we anticipate that many of the features seen here will remain, in particular, momentum being generated in the bulk by the magnetic field on the brane, and particle production in the CFT leading to anisotropy dissipation.  Perhaps one of the most interesting features of this work was the possibility that one could detect a braneworld signature in the sky, as discussed near the end of section~\\ref{sec:bw}. In the braneworld picture, particles are actually {\\it over}-produced along the direction of the magnetic field, so that the expansion along that direction is {\\it slower} than in the orthogonal directions, in complete contrast with what happens in standard $4D$ General Relativity. For this reason it would be interesting to extend this work to other, more phenomenological, Bianchi models. If, say, one of the spatial directions were  periodic, we might expect {\\it angular} momentum to be carried into the bulk, and so we would look for an embedding in the general Myers-Perry spacetimes~\\cite{MP, HHT}.  A perturbative study along these lines has been considered~\\cite{Guth}, although a far more complete analysis is clearly required.  In the long term, we would like to consider  a physically realistic scenario with a Bianchi VIIh brane containing a homogeneous magnetic field along with an isotropic perfect fluid made up of matter, radiation and a cosmological constant. Of course, this would require an amalgamation of the ideas outlined in the previous two paragraphs. It would be very interesting to see what effect braneworld isotropization has on the CMB, and in particular the bounds on the size of the primordial magnetic field~\\cite{CMBbound}.  We must also keep in mind constraints coming from nucleosynthesis. For example, if braneworld isotropization is {\\it too} efficient, then we might worry that CFT particle production will result in too much dark radiation, leading to an unacceptably low baryon density parameter."
    },
    "0801/0801.2652_arXiv.txt": {
        "abstract": "{The distance to the Galactic Centre (GC) is of importance for the distance scale in the Universe.  The value derived by Eisenhauer et al. (2005) of 7.62 $\\pm$ 0.32 kpc based on the orbit of one star around the central black hole is shorter than most other distance estimates based on a variety of different methods.} {To establish an independent distance to the GC with high accuracy. To this end Population-{\\sc ii} Cepheids are used that have been discovered in the \\OG-{\\sc ii} and \\OG-{\\sc iii} surveys.} {Thirty-nine Population-{\\sc ii} Cepheids have been monitored with the SOFI infrared camera on 4 nights spanning 14 days, obtaining typically between 5 and 11 epochs of data. Light curves have been fitted using the known periods from the OGLE data to determine the mean $K$-band magnitude with an accuracy of 0.01-0.02 mag. It so happens that 37 RR Lyrae stars are in the field-of-view of the observations and mean $K$-band magnitudes are derived for this sample as well.} {After correction for reddening, the period-luminosity relation of Population-{\\sc ii} Cepheids in the $K$-band is determined, and the derived slope of $-2.24 \\pm 0.14$ is consistent with the value derived by Matsunaga et al. (2006). Fixing the slope to their more accurate value results in a zero point, and implies a distance modulus to the GC of 14.51 $\\pm$ 0.12, with an additional systematic uncertainty of 0.07 mag.  Similarly, from the RR Lyrae $K$-band period-luminosity relation we derive a value of 14.48 $\\pm$ 0.17 (random) $\\pm$ 0.07 (syst.). The two independent determinations are averaged to find 14.50 $\\pm$ 0.10 (random) $\\pm$ 0.07 (syst.), or 7.94 $\\pm$ 0.37 $\\pm$ 0.26 kpc. The absolute magnitude scale of the adopted period-luminosity relations is tied to an LMC distance modulus of 18.50 $\\pm$ 0.07.  } {} ",
        "introduction": "The distance to astronomical objects is a crucial parameter, yet often very difficult to obtain with high precision. The distance to the Galactic Centre (GC) is of special importance, e.g for dynamics (Oort constants, determining distances using a rotation model), or for calibrating standard candles. The classically accepted value comes from the review by Reid (1993) and is $R_0$ = 8.0 $\\pm$ 0.5 kpc.   Over the last few years the distance to the GC based on the orbit of the star called S2 around the central black-hole (BH) has caught attention.  Initially, Eisenhauer et al. (2003) derived a value of 7.94 $\\pm$ 0.42 kpc which was revised by Eisenhauer et al. (2005) to 7.62 $\\pm$ 0.32 kpc having more epochs of data available. The neglect of post-Newtonian physics in these analysis may have lead to an underestimate of the distance by about $0.11 \\pm 0.02$ kpc (Zucker et al. 2006), leading to a current best estimate of 7.73 $\\pm$ 0.32 kpc (corresponding to a distance modulus (DM) of 14.44 $\\pm$ 0.09) to the GC based on the BH.   On the other hand, most other recent distance determinations give a longer distance, more in line with the classical value: (1) High-amplitude delta-scuti stars give 7.9 $\\pm$ 0.3 kpc (McNamara et al. 2000); (2) RR Lyrae stars suggest a value of 8.8 $\\pm$ 0.3 kpc (Collinge et al. 2006), 8.3 $\\pm$ 1.0 kpc (Carney et al. 1995) or 8.0 $\\pm$ 0.65 kpc (Fernley et al. 1987). (3) Earlier work on the Red Clump gave a longer distance of 8.4 $\\pm$ 0.4 kpc (Paczy\\'nski \\& Stanek 1998), although Nishiyama et al. (2006) derive 7.52 $\\pm$ 0.10 (stat) $\\pm$ 0.35 (syst) kpc, and Babusiaux \\& Gilmore (2005) 7.7 $\\pm$ 0.15 kpc; (4) From a comparison of Miras found in the OGLE database in the direction of the Galactic Bulge (GB) to those in the Magellanic Clouds, Groenewegen \\& Blommaert (2005) find a distance in the range 8.5 to 9.0 kpc, in agreement with earlier work on Miras (Catchpole et al. 1999); (5) Analysis of the Hipparcos proper motions of 220 Cepheids lead to $R_0$ = 8.5 $\\pm$ 0.5 kpc (Feast \\& Whitelock 1997); (6) Modelling the observed colour-magnitude diagram in $V,I$ and $J,K$ using a population synthesis code, Vanhollebeke et al. (2008) derive a distance of 8.60 $\\pm$ 0.16 kpc.   With the exception of some Red Clump based distances, the results obtained by Eisenhauer et al. imply a much shorter distance to the GC than found by most other methods, and this calls for an independent investigation of this matter.  In this paper the distance to the GC is determined using \\tc\\ (hereafter P2C) discovered in the OGLE micro-lensing survey, and for which the mean $K$-band magnitude will be determined by infrared monitoring. Comparing to the calibrated P2C period-luminosity (PL) relation in the $K$-band from Matsunaga et al. (2006; hereafter M06) then provides the distance, after correction for reddening.  In addition, the mean $K$-band magnitude will be determined for RR Lyrae stars that are in the field, and compared to the calibrated $K$-band PL-relation from Sollima et al. (2006). The Matsunaga et al. and Sollima et al. relations both imply an LMC DM of 18.50 as detailed in Sect.~5.  In Sect.~2 the sample is discussed, and the observations are presented in Sect.~3 for the P2C and Sect.~3 for the RR Lyrae.  The results are discussed in Sect.~5. ",
        "conclusions": "The $PL$-relation in the $K$-band of P2C in the GB is derived.  The slope is found to be in agreement with that derived by M06. Fixing the slope to their more accurate value implies a DM to the GC of 14.51 with a formal error bar of 0.03.   There is also the systematic error bar to consider. M06 presented $JHK$ period-luminosity relations based on 46 P2C with periods between 1.2 and 80 days in 26 Galactic globular clusters (GCs). For the absolute magnitude scale they adopted a relation between absolute $V$-magnitude of the Horizontal Branch and metallicity (Gratton et al. 2003), which in turn is calibrated using main-sequence fitting to three GCs. This calibration implies an RR Lyrae based LMC DM of 18.50 $\\pm$ 0.09 (Gratton et al. 2003). M06 show that the DM based on P2C in the LMC and their $K$-band $PL$-relation is also compatible with 18.5.  They also show that there is no significant trend with metallicity over the range $-2.2 \\la$ [Fe/H] $\\la -0.5$, in agreement with theoretical predictions (Bono et al. 1997, 2003, Di Criscienzo et al. 2007), and indicating that this $K$-band $PL$-relation should be applicable for GC P2C as well. The metallicity of the P2C in the Bulge is unknown but that of RR Lyrae is estimated to be on average [Fe/H]= $-1.0$ (Walker \\& Terndrup 1991).  Any difference between that metallicity and the mean metallicity of about [Fe/H]= $-1.5$ of the GCs in M06 would result in an uncertainty in the ZP of \\less 0.03 mag.  Di Criscienzo et al. (2007) show that adopting a different $M_{\\rm V}$-[Fe/H] relation has a negligible effect on the derived slope of the NIR $PL$-relations.  The ZP in the calibrating relation by Gratton et al. has a formal error of 0.07 and this has to be considered as a source of systematic uncertainty in the derived distance.  There is other source of (random) error to consider, namely how representative this particular set of 39 stars (minus the one outlier) that defines the $PL$-relation is in view of the fact that they scatter along the line-of-sight due to the intrinsic depth of the Bulge. To simulate this, additional Monte-Carlo simulations were carried out. Random samples of 38 stars were selected from the original sample, and the $PL$-relation re-derived. The dispersion in the ZP is about 0.11 mag. This is likely a slight overestimate as in this approach the randomly drawn samples do not necessarily have the large spread in period that the true sample was selected to have.   The DM to the GC we derive based on the P2C is 14.51 $\\pm$ 0.12 (random) $\\pm$ 0.07 (syst). The random error could be improved further by observing additional systems when the full OGLE-{\\sc iii} database becomes available.  Based on the serendipitously observed RR Lyrae stars in the field a DM of 14.52 $\\pm$ 0.18 is derived. The error bar is for 50\\% due to the uncertainty in the adopted absolute magnitude of RR Lyra itself (Sollima et al.). Their PL-relation led to an LMC distance of 18.54 $\\pm$ 0.15.  If instead we would {\\em assume} the LMC distance to be 18.50 (to be consistent with the P2C calibration) then we would find a DM of 14.48 $\\pm$ 0.13 (random) $\\pm$ 0.07 (syst), were the systematic error comes from the uncertainty in the ZP of the observed LMC PL-relation. Like for the P2C sample, there is an additional 0.11 mag systematic uncertainty due to the limited sample size. The final DM to the GC based on the $K$-band RR Lyrae stars is 14.48 $\\pm$ 0.17 (random) $\\pm$ 0.07 (syst).  As the two distance estimates are derived independently, they can be averaged and the best empirical estimate of the DM to the GC based on the current data is 14.50 $\\pm$ 0.10 (random) $\\pm$ 0.07 (syst).   The theoretical WIK relation gives a formal result of 14.44 (or 14.33 with anomalous reddening) with an internal error bar of 0.04 mag. A random error of 0.11 has to be added to this, due to the limited sample, and as the theoretical relation is tied to the observed relations of M06 a similar systematic error bar of 0.07 has to be considered.  Although within the error bar of the purely empirical results, it brings up the question of reddening and the reddening law. An additional absorption in $K$ of 0.1 mag would bring all three methods in very good agreement. On the other hand, the reddening estimates listed in Table~1 are in excellent agreement and are based on $(J-K)$ colours (Marschall et al. 2006), $(V-I)$ (Sumi 2004), and $(V-R)$ (Popowski et al. 2003). If the absorption in $K$ were underestimated, it would also imply an underestimate of the reddening in the other maps, or a significantly higher selective reddening $A_{\\rm K}/A_{\\rm V} \\sim 0.16$ instead of 0.12.    A final remark is that independent distances to some of these P2C may be obtained using surface-brightness relations (e.g. Groenewegen 2004) and the Baade-Wesselink technique. This would require better sampled $K$-band light curves than were needed for the present study and well-sampled radial velocity curves. Although observationally expensive it would give an improved understanding on the systematic error in the present analysis."
    },
    "0801/0801.2008_arXiv.txt": {
        "abstract": "The generalized Chaplygin gas, which interpolates between a high density relativistic era and a non-relativistic matter phase, is a popular dark energy candidate. We consider a generalization of the Chaplygin gas model, by assuming the presence of a bulk viscous type dissipative term in the effective thermodynamic pressure of the gas. The dissipative effects are described by using the truncated Israel-Stewart model, with the bulk viscosity coefficient and the relaxation time functions of the energy density only. The corresponding cosmological dynamics of the bulk viscous Chaplygin gas dominated universe is considered in detail for a flat homogeneous isotropic Friedmann-Robertson-Walker geometry. For different values of the model parameters we consider the evolution of the cosmological parameters (scale factor, energy density, Hubble function, deceleration parameter and luminosity distance, respectively), by using both analytical and numerical methods. In the large time limit the model describes an accelerating universe, with the effective negative pressure induced by the Chaplygin gas and the bulk viscous pressure driving the acceleration. The theoretical predictions of the luminosity distance of our model are compared with the observations of the type Ia supernovae. The model fits well the recent supernova data. From the fitting we determine both the equation of state of the Chaplygin gas, and the parameters characterizing the bulk viscosity. The evolution of the scalar field associated to the viscous Chaplygin fluid is also considered, and the corresponding potential is obtained. Hence the viscous Chaplygin gas model offers an effective dynamical possibility for replacing the cosmological constant, and to explain the recent acceleration of the universe. ",
        "introduction": "The observations of high redshift supernovae \\cite{Pe99} and the Boomerang/Maxima/WMAP data \\cite{Ber00}, showing that the location of the first acoustic peak in the power spectrum of the microwave background radiation is consistent with the inflationary prediction $\\Omega =1$, have provided compelling evidence for a net equation of state of the cosmic fluid lying in the range $-1\\leq w=p/\\rho <-1/3$. To explain these observations, two dark components are invoked: the pressureless cold dark matter (CDM) and the dark energy (DE) with negative pressure. CDM contributes $\\Omega _{m}\\sim 0.3$, and is mainly motivated by the theoretical interpretation of the galactic rotation curves and large scale structure formation. DE is assumed to provide $\\Omega _{DE}\\sim 0.7$ and is responsible for the acceleration of the distant type Ia supernovae. There are a huge number of candidates for DE in the literature (for recent reviews see \\cite{PeRa03} and \\cite{Pa03}).  One possibility are cosmologies based on a mixture of cold dark matter and quintessence, a slowly-varying, spatially inhomogeneous component \\cite{8}. An example of implementation of the idea of quintessence is the suggestion that it is the energy associated with a scalar field $Q$ with self-interaction potential $V(Q)$. If the potential energy density is greater than the kinetic one, then the pressure $p=\\dot{Q}^{2}/2-V(Q)$ associated to the $Q$-field is negative. Quintessential cosmological models have been intensively investigated in the physical literature \\cite{quint}.  A different line of thought has been followed in \\cite {15,16,17}, where the conditions under which the dynamics of a self-interacting Brans$% -$Dicke (BD) field can account for the accelerated expansion of the Universe have been analyzed. Accelerated expanding solutions can be obtained with a quadratic self-coupling of the BD field and a negative coupling constant $\\omega $ \\cite{15}.  Dissipative effects, including both bulk and shear viscosity, are supposed to play a very important role in the early evolution of the Universe. A cosmic fluid (pressureless and with pressure) obeying a perfect fluid type equation of state cannot support the acceleration \\cite{17}. A solution to this problem, and thus avoiding the necessity of a potential for the BD field, is to assume that some dissipative effects of bulk viscous type take place at the cosmological scale \\cite{16}. A combination of a cosmic fluid with bulk dissipative pressure and quintessence matter can drive an accelerated expansion phase of the Universe and also solve the coincidence problem (the observational fact that the energy density of cold dark matter and of $Q$-matter should be comparable today) \\cite{11}. The dynamics of a causal bulk viscous cosmological fluid filled flat homogeneous Universe in the framework of the BD theory was considered in \\cite{MaHa03}. The bulk viscous pressure term in the matter energy-momentum tensor leads to a non-decelerating evolution of the Universe.  Neither CDM nor DE have direct laboratory observational or experimental evidence for their existence. Therefore it would be important if a unified dark matter - dark energy scenario could be found, in which these two components are different manifestations of a single fluid \\cite{Paetal}. A candidate for such an unification is the so-called generalized Chaplygin gas, which is an exotic fluid with the equation of state $p=-B/\\rho ^{n}$, where $B$ and $n$ are two parameters to be determined. It was initially suggested in \\cite{Ka01} with $n=1$, and then generalized in \\cite{Be02} for the case $n\\neq 1$. The Chaplygin gas also appears in the stabilization of branes in Schwarzschild-AdS black hole bulks as a critical theory at the horizon \\cite% {Ka00} and in the stringy analysis of black holes in three dimensions \\cite% {Ka98}. The Chaplygin equation of state can be derived from Born-Infeld type Lagrangians \\cite{Be02}, \\cite{No05}. This simple and elegant model smoothly interpolates between a non-relativistic matter phase ($p=0$) and a negative-pressure dark energy dominated phase.  The cosmological implications of the Chaplygin gas model have been intensively investigated in the recent literature \\cite{Chco}. The Chaplygin gas cosmological model has been constrained by using different cosmological observations, like type Ia supernovae \\cite{Fa02}, the CMB anisotropy measurements \\cite{Be03}, gravitational lensing surveys \\cite{De03}, the age measurement of high redshift objects \\cite{Al03} and the X-ray gas mass fraction of clusters \\cite{Cu04}. The obtained results are somewhat controversial, with some of them claiming good agreement between the data and the Chaplygin gas model, while the rest ruling it as a feasible candidate for dark matter. In particular, the standard Chaplygin gas model with $n=1$ is ruled out by the data at a 99\\% level \\cite{Cu04}. The exact solutions of the gravitational field equations in the generalized Randall-Sundrum model for an anisotropic brane with Bianchi type I geometry, with a generalized Chaplygin gas as matter source were obtained in \\cite% {MaHa05}.  The possibility of constraining Chaplygin dark energy models with current Integrated Sachs Wolfe (ISW) effect data was investigated in \\cite{GiMe06}. In the case of a flat universe the generalized Chaplygin gas models must have an energy density such that $\\Omega _{c}>0.55$ and an equation of state $w<-0.6$ at 95\\% confidence level. The extent to which the knowledge of spatial topology may place constraints on the parameters of the generalized Chaplygin gas (GCG) model for unification of dark energy and dark matter was studied in \\cite{Ber06}. By using both the Poincar\\'{e} dodecahedral and binary octahedral spaces as the observable spatial topologies, the current type Ia supernovae (SNe Ia) constraints on the GCG model parameters were examined. An action formulation for the GCG model was developed in \\cite% {BaGhKu07}, and the most general form for the nonrelativistic GCG action consistent with the equation of state has been derived. The thermodynamical properties of dark energy have been investigated in \\cite{GoWaWa07}. For dark energy with constant equation of state $w>-1$ and the generalized Chaplygin gas, the entropy is positive and satisfies the entropy bound. Observational constraints on the generalized Chaplygin gas (GCG) model for dark energy from the 9 Hubble parameter data points, the 115 SNLS Sne Ia data and the size of baryonic acoustic oscillation peak at redshift, $z=0.35$ were examined in \\cite{WuYu07}. At a 95.4\\% confidence level, a combination of the three data sets gives $% 0.67\\leq B/\\rho_0^{1+n}\\leq 0.83$ (where $\\rho_0$ is the present day energy density) and $-0.21\\leq n\\leq 0.42$, which is within the allowed parameters ranges of the GCG as a candidate of the unified dark matter and dark energy. However, the standard Chaplygin gas model ($n=1$) is also ruled out by these data at the 99.7\\% confidence level. A geometrical explanation for the generalized Chaplygin gas within the context of brane world theories, where matter fields are confined to the brane by means of the action of a confining potential, was considered in \\cite{HeSe07}.  The evolution of the Universe contains a sequence of important dissipative processes, including GUT (Grand Unified theory) phase transition, taking place at $t\\approx 10^{-34}$ s and a temperature of about $T\\approx 10^{27}$ K, when gauge bosons acquire mass, reheating of the Universe at the end of inflation ($t\\approx 10^{-32}$ s), when the scalar field decays into particles, decoupling of neutrinos from the cosmic plasma ($t\\approx 1$ s, $% T\\approx 10^{10}$ K), when the temperature falls below the threshold for interactions that keep the neutrinos in thermal contact, nucleosynthesis, decoupling of photons from matter during the recombination era ($t\\approx 10$ s, $T\\approx 10^{3}$ K), when electrons combine with protons and no longer scatter the photons etc. \\cite{Ma}.  The first attempts at creating a theory of relativistic dissipative fluids were those of Eckart \\cite{Ec40} and Landau and Lifshitz \\cite{LaLi87}. These theories are now known to be pathological in several respects. Regardless of the choice of the equation of state, all equilibrium states in these theories are unstable and in addition signals may be propagated through the fluid at velocities exceeding the speed of light. These problems arise due to the first order nature of the theory, that is, it considers only first-order deviations from the equilibrium leading to parabolic differential equations, hence to infinite speeds of propagation for heat flow and viscosity, in contradiction with the principle of causality. Conventional theory is thus applicable only to phenomena which are quasi-stationary, i.e. slowly varying on space and time scales characterized by mean free path and mean collision time.  A relativistic second-order theory was found by Israel \\cite{Is76} and developed in \\cite{IsSt76} and  \\cite% {HiLi89,HiSa91} into what is called ``transient'' or ``extended'' irreversible thermodynamics. In this model deviations from equilibrium (bulk stress, heat flow and shear stress) are treated as independent dynamical variables, leading to a total of 14 dynamical fluid variables to be determined. For general reviews on causal thermodynamics and its role in relativity see \\cite{Ma} and \\cite{Ma95}. Causal bulk viscous thermodynamics has been extensively used for describing the dynamics and evolution of the early Universe, or in an astrophysical context \\cite% {ChJa97}.  It is the purpose of this paper to consider the effects of a possible existence of a bulk viscosity of the generalized Chaplygin gas on the cosmological dynamics of the Universe. The viscous effects are described by using the truncated Israel-Stewart theory \\cite{IsSt76}. By using the Laplace transformation and the convolution theorem, the second order differential equation describing the evolution of the Hubble parameter $H$ is transformed into an integral equation. The field equations are solved by means of an iterative scheme. Then the general solutions of the equations are obtained in a parametric form in the zero, first, second and $m$th order approximation, and the relevant cosmological parameters (scale factor, energy density, Hubble parameter, deceleration parameter etc.) are obtained. The scalar field interpretation of the Chaplygin gas is generalized to take into account the viscosity and dissipative effects.  In order to compare the predictions of the model with the observational data we have fitted the luminosity distance-redshift relation with the latest observational data of the type Ia supernovae. The model fits well these data. From the fitting we determine both the equation of state of the Chaplygin gas, and the parameters characterizing the bulk viscosity. Even by taking into account the effect of the bulk viscosity, the $n=1$ Chaplygin gas models are ruled out by the observations.  The present paper is organized as follows. The physical model and the basic equations are presented in Section II. The evolution equation for the Hubble parameter is studied in Section III, and the behavior of the cosmological parameters is obtained. The observational data have been compared with the theoretical predictions of the model in Section IV. In Section V we discuss and conclude our results. In the present paper we use a system of units so that $8\\pi G=c=1$. ",
        "conclusions": "In the present paper we have considered the dynamics of a bulk viscous Chaplygin gas filled flat homogeneous and isotropic universe. We have derived and formulated the evolution equations of the system, we have considered their behavior by using both analytical and numerical techniques, and we have compared the predictions of our model with the supernova data. The most attractive feature of the Chaplygin gas is that it could explain the main observational properties of the Universe without appealing to an effective cosmological constant. Generally, the obtained analytical and numerical solutions of the gravitational field equations describes an accelerating universe, with the effective negative pressure induced by the Chaplygin gas and the bulk viscous pressure driving the acceleration.  From the equation of state of the Chaplygin gas with $\\gamma =0$ it follows that for the critical values $p_{c}$ and $\\rho _{c}$ of the pressure and density the parameter $w_{c}=p_{c}/\\rho _{c}$ is given by $w_{c}=-B/\\rho _{c}^{n+1}-\\Pi _{c}/\\rho _{c}$. Evaluating this relation at the present time when $\\rho _{c}=\\rho _{c0}$ gives $B=-w_{c0}\\rho _{c0}^{n+1}-\\Pi \\left( \\rho _{c0}\\right) \\rho _{c0}^{n}$. The Chaplygin gas behaves like a cosmological constant for $w_{c0}=-1$, which gives the relation between the constant $B$ and the present day value of the bulk viscous pressure as \\begin{equation} B=\\rho _{c0}^{n+1}\\left[ 1-\\frac{\\Pi \\left( \\rho _{c0}\\right) }{\\rho _{c0}}% \\right] =\\left( \\frac{3H_{0}^{2}}{8\\pi G}\\right) ^{n+1}\\left[ 1-\\frac{\\Pi \\left( \\rho _{c0}\\right) }{\\rho _{c0}}\\right] , \\end{equation}% where $H_{0}=3.24\\times 10^{-18}h$ s$^{-1}$, $0.5\\leq h\\leq 1$, is the Hubble constant \\cite{PeRa03,Pa03}. Since $\\Pi \\left( \\rho _{c0}\\right) <0$, the presence of the bulk viscous effects can significantly increase the value of $B$.  By comparing the model with $\\gamma =0$ to the Gold 2006 supernova data, it turns out that a good agreement with these observations can be established for a wide range of the power $\\ s\\in \\left( 0.2,~2\\right) $ which occurs in the phenomenological laws (\\ref{csi}), which characterize the bulk viscosity coefficient. The other viscosity parameter $\\alpha $ can be obtained from the equation $3^{s-1}\\alpha H_{0}^{2s-1}=1$, and by choosing a value for the Hubble parameter. For $h=0.7$ ($% H_{0}=2.268\\times 10^{-18}$s$^{-1}$) we obtain $\\alpha =\\left( 6.2385\\times 10^{-11}\\text{s}^{-0.6},~2.8573\\times 10^{52}\\text{s}^{3}\\right) $ for the above-established range of the parameter $s$. As for the equation of state of the Chaplygin gas, by taking into account the definition of $\\lambda $, $\\lambda =B/3^{n}H_{0}^{2n+2}$, and for the same value of the Hubble parameter, the confrontation with supernova data selects the pairs $\\left( n,~B\\right) $ represented on Fig. \\ref{nB}.  \\begin{figure}[tbp] \\includegraphics[height=5cm]{nB.eps} \\caption{The parameter values $B$ corresponding to the best fit values $% \\left( \\protect\\lambda ,~n\\right) $ represented in a logarithmic scale as a function of $n$.} \\label{nB} \\end{figure}  Scalar fields are supposed to play a fundamental role in the evolution of the early universe. The Chaplygin gas model can be also described from a field theoretical point of view by introducing a scalar field $\\phi $ and a self interacting potential $U(\\phi )$, with the Lagrangian \\cite{Ka01,Be02}, \\cite{Bi02}, \\cite{Fa02}, \\cite{De04} \\begin{equation} L_{\\phi }=\\frac{1}{2}\\dot{\\phi}^{2}-U(\\phi ). \\end{equation}  The energy density and the pressure associated to the scalar field $\\phi $ associated to the bulk viscous Chaplygin gas are given by \\begin{equation} \\rho _{\\phi }=\\frac{\\dot{\\phi}^{2}}{2}+U\\left( \\phi \\right) =\\rho , \\label{rho} \\end{equation} and \\begin{equation} p_{\\phi }=\\frac{\\dot{\\phi}^{2}}{2}-U\\left( \\phi \\right) =\\gamma \\rho -\\frac{B% }{\\rho ^{n}}+\\Pi ,  \\label{p} \\end{equation} respectively.  The scalar field and the potential can be obtained from the equations \\begin{equation} \\phi \\left( t\\right) -\\phi _{0}=\\int_{t_{0}}^{t}\\sqrt{\\left( 1+\\gamma \\right) \\rho -\\frac{B}{\\rho ^{n}}+\\Pi }dt, \\end{equation} and \\begin{equation} U\\left( t\\right) =\\frac{1}{2}\\left[ \\left( 1-\\gamma \\right) \\rho +\\frac{B}{% \\rho ^{n}}-\\Pi \\right] , \\end{equation} respectively, where $\\phi _{0}$ is an arbitrary constant of integration.  The dependence of the potential $U(\\phi )$ on the scalar field $\\phi $ is represented in Fig.~\\ref{FIG8}.  \\begin{figure}[!ht] \\centering \\includegraphics{f8.eps} \\caption{The potential $U(\\phi )$ of the viscous Chaplygin gas associated scalar field as a function of the scalar field $\\phi $ for a dust universe ($\\gamma =0$), $n=0.1$, $s=1/4$, and for different values of $\\lambda _0$: $% \\lambda _0=0.01$ (solid curve), $\\lambda _0=0.03$ (dotted curve), $% \\lambda _0=0.05$ (dashed curve) and $\\lambda _0=0.07$ (long dashed curve). } \\label{FIG8} \\end{figure}  In conclusion, we have found that the viscous Chaplygin gas model offers a real possibility for replacing the effective cosmological constant and to explain the recent acceleration of the universe."
    },
    "0801/0801.0705_arXiv.txt": {
        "abstract": "Early universe equations of state including realistic interactions between constituents are built up. Under certain hypothesis, these equations are able to generate an inflationary regime prior to the nucleosynthesis period. The resulting accelerated expansion is intense enough to solve the flatness and horizon problems. In the cases of curvature parameter $\\kappa $ equal to $0$ or $+1$, the model is able to avoid the initial singularity and offers a natural explanation for why the universe is in expansion. All the results are valid only for a matter-antimatter symmetric universe. ",
        "introduction": "}  The gravitational field, as described by General Relativity, couples to all types of energy: rest masses, kinetic terms and interaction terms. Relativistic cosmology is, in consequence, deeply concerned with such sources. The kinetic terms and the rest-masses are currently taken into account by assuming ideal cosmic fluids constituted by ultrarelativistic and/or non-relativistic matter. Solutions of this type are found in standard texts \\cite{nar,wein} and -- when multi-component fluids are considered -- in several papers, e.g., \\cite{Solutions}. Interaction terms are commonly used only in perturbative models. In fact, using the Boltzmann equation in a Friedmann-Robertson-Walker (FRW) background, the inhomogeneities both in the cosmic microwave background (CMB) and in the matter content \\cite{Dod,kolb} are studied, for comparison with the observational data \\cite% {Wmap,LargeScale}. Notwithstanding, interactions are not considered as direct sources of gravitation in this line of research \\cite{artigo 1}.  Of course, some well-known proposals do consider interaction processes as direct sources of gravitation. They are usually related to the accelerated expansion regimes: present-day dynamics \\cite{SuNo} or inflation \\cite% {Liddle}. Nevertheless, these theories do not actually consider the fundamental interactions (electromagnetic, weak and strong) between particles in the source constituents. Indeed, in the standard inflationary approaches \\cite{Guth,AlbreSten,Linde} an accelerated expansion is obtained through self-interaction processes of scalar \\textit{inflaton} fields $\\phi _{i}\\left( x\\right) $. This self-interaction is chosen so as to produce just those features which are necessary to describe the early accelerated regime \\cite{Liddle}. The phenomenological explanations for the present-day acceleration include, among others, \\emph{(i)} the quintessence models \\cite% {CalDaveStein,Stein,Zlatev,FraRos,PeRa}, which roughly follow the same lines of the inflationary theory; \\emph{(ii)} models of matter and dark-energy unification via Chapligyn-like equations of state (EOS) \\cite{Orfeu} or through equations of the Van der Walls type \\cite{Capozzielo,Kremer}; \\emph{% (iii)} models of mass-varying-neutrino type, which couple neutrinos to a quintessence scalar field \\cite{Motta}; \\emph{(iv)} models introducing interactions in the energy conservation equation \\cite{Gabi,NelsonPinto}.  A formal procedure has been recently proposed for including the (fundamental) interactions as direct sources of gravitation in the cosmological context \\cite{artigo 1}. Imported from equilibrium statistical mechanics, this formalism allows the construction of realistic equations of state for both relativistic \\cite{Dashen,Reichl} and non-relativistic systems \\cite{Pat,Beth}. Our objective here is to find primeval cosmic fluid EOS taking into account \\textit{physically realistic} interaction processes between the constituent particles, in addition to their kinetic and rest-mass terms. In particular, we shall examine their effect on the scale factor evolution. It will be shown that under certain hypothesis and approximations an early accelerated regime can be obtained as a consequence of these interacting processes.  The idea of building realistic equations of state considering interaction between elementary particles in the primeval universe is not new. Actually, during the late 1970's and the beginning os the 80's a series of papers by Bugrii, Trushevsky and Beletsky \\cite{ucraone,ucratwo,ucrathree,ucrafour} discussed the construction of the high-energy EOS and their application to the pre-nucleosynthesis universe.\\footnote{% The authors are thankful to an unknown referee for calling their attention to these works.} Nevertheless, their treatment is different from the one developed here as we will make clear some pages ahead.  The paper is organized as follows. Section \\ref{sec-Pns} reviews the general features of the early universe in its standard presentation \\cite% {nar,wein,Dod,kolb,Liddle}. Section \\ref{sec-SisInt} presents some results from equilibrium statistical theory of interacting systems. Specifically, the coefficients appearing in the perturbative fugacity expansions are expressed in terms of the scattering matrix operator $\\hat{S}$. In addition, the matrix $S_{2}$ describing the two-particle scattering is associated to the experimentally observable phase-shifts. The goal of Section \\ref% {sec-EoS_Pns} is to construct realistic EOS\\ for the pre-nucleosynthesis universe and discuss the hypothesis and approximations undertaken. In Section \\ref{sec-Conseq},\\ the more direct cosmological consequences coming from these equations are examined, including the effect of driving an accelerated expansion (inflationary era). Section \\ref{sec-Final} contains some final comments. Details of a too technical nature, as well as some data, have been relegated to appendices. ",
        "conclusions": ""
    },
    "0801/0801.1190_arXiv.txt": {
        "abstract": "Post-starburst, or E+A galaxies, are the best candidates for galaxies in transition from being gas-rich and star-forming to gas-poor and passively-evolving as a result of galaxy-galaxy interactions.  To focus on what E+A galaxies become after their young stellar populations fade away, we present the detailed morphologies of 21 E+A galaxies using high resolution {\\sl HST}/{\\sl ACS} and {\\sl WFPC2} images.  Most of these galaxies lie in the field, well outside of rich clusters, and at least 11 (55\\%) have dramatic tidal features indicative of mergers. Our sample includes one binary E+A system, in which both E+As are tidally disturbed and interacting with each other. Our E+As are similar to early types in that they have large bulge-to-total light ratios (median $B/T$ = 0.59), high S\\'ersic indices, ($n \\gtrsim 4$), and high concentration indices ($C \\gtrsim 4.3$), but they have considerably larger asymmetry indices ($A \\gtrsim 0.04$) than ellipticals, presumably due to the disturbances within a few $r_e$ caused by the starburst and/or the galaxy-galaxy interaction.  We conclude that E+As will be morphologically classified as early-type galaxies once these disturbances and the low surface brightness tidal features fade. The color morphologies are diverse, including six E+As with compact (0.4\\,--\\,1.4 kpc) blue cores, which might be local analogs of high-$z$ ellipticals with blue-cores.  The large fraction (70\\%) of E+As with positive color gradients indicates that the young stellar populations are more concentrated than the old.  These positive color gradients (i.e., bluer nuclei) could evolve into the negative gradients typical in E/S0s if the central parts of these galaxies are metal enhanced. Our E+As stand apart from the E/S0s in the edge-on projection of the Fundamental Plane (FP), implying that their stellar populations differ from those of E/S0s and that E+As have, on average, a \\ml that is 3.8 times smaller.  The tilt of the E+A FP indicates that the variation among their stellar populations is closely tied to the structural parameters, i.e., E+As follow their own scaling relationships such that smaller or less massive galaxies have smaller {\\sl M/L}. We find a population of unresolved compact sources in nine E+As (45\\%), all of which have merger signatures.  In the four E+As with suitable color data, the compact sources have colors and luminosities consistent with newly-formed star clusters.  The bright end of the cluster LF is fainter in redder E+A's, suggesting that the young star clusters fade or are disrupted as the merger remnant ages. In summary, the morphologies, color profiles, scaling relations, and cluster populations are all consistent with E+As evolving ultimately into early-types, making the study of E+As critical to understanding the origin of the red sequence of galaxies. ",
        "introduction": "If some galaxies evolve from star-forming, gas-rich, disk-dominated galaxies (late-types) into quiescent, gas-poor, spheroid-dominated galaxies (early-types), we should find objects caught in the midst of this transformation.  The best candidates are the so-called ``E+A'', ``K+A'', or \"post-starburst\" galaxies \\citep{Dressler83,Couch87} due to their combination of late- and early-type characteristics, including both a significant young stellar population (age $\\lesssim$ 1 Gyr) and a lack of on-going star formation.  These galaxies have been spectroscopically identified by their strong Balmer absorption lines and absence of emission lines (e.g., [\\ion{O}{2}] and H$\\alpha$) in various environments and at all redshifts \\citep{Zabludoff96, Poggianti99, Goto03, Blake04, Tran03, Tran04}.  While the cause of the abrupt end of their star formation is poorly understood, there is strong evidence that galaxy-galaxy tidal interactions or mergers trigger the starburst in many cases.  First, most E+A galaxies reside in low-density environments, such as poor groups, that are similar to those of star-forming galaxies \\citep{Zabludoff96, Quintero04, Blake04, Balogh05, Goto05, Hogg06, Yan08}. Therefore, many E+As must arise from a process common in the field, such as galaxy-galaxy interactions, instead of a mechanism limited to denser, hotter environments, such as ram pressure stripping \\citep{Gunn&Gott72} or strangulation \\citep{Balogh00}.  Second, a significant fraction of E+As have tidal features \\citep{Zabludoff96, Yang04, Blake04, Tran03, Tran04, Goto05}. Third, optical and {\\sl NIR} colors show that the spectral signatures of E+As require enhanced recent star formation, rather than simply a truncation of star formation in a normal spiral galaxies \\citep{Balogh05}.  What will E+A galaxies become? In general, E+As are bulge-dominated, highly-concentrated \\citep{Quintero04, Tran04, Blake04, Goto05, Balogh05}, relatively gas-poor \\citep{Chang01, Buyle06}, and kinematically hot systems \\citep{Norton01}. Therefore, in a statistical sense, E+As are likely to become E/S0 galaxies.  However, due to a lack of spatial resolution in previous studies, we do not know whether the detailed properties of individual E+As, e.g., their bulge fractions, color gradients, internal kinematics, and newly formed stellar clusters, are consistent with their presumed evolution into early type galaxies. Using {\\sl HST/WFPC2}, \\citet{Yang04} showed that their morphological features are consistent with a  transition from late to early types, but their sample contained only the five bluest E+As galaxies from the Las Campanas Redshift Survey (LCRS; \\cite{Zabludoff96}) and thus was not representative.  As a result, we still do not know whether the entire population of E+As will evolve into E/S0s, whether there is a distinguishable subclass of E+As that evolves into E/S0s, or whether E+As evolve into typical E/S0s.  To answer these questions, we must understand how well E+As match the {\\it full} range of E/S0 properties. Most fundamentally, are the global morphologies (e.g., bulge-to-total light ratios and concentration) of the whole LCRS E+A sample consistent with those of E/S0s? Second, E/S0s in the local universe become redder toward their center; these negative color gradients originate from metallicity gradients \\citep{Peletier90}.  In contrast, E+As exhibit a wide range of color morphologies \\citep{Yang04, Yamauchi05}.  Can the color profiles of E+As evolve into those of the typical early types? Third, the number of globular clusters per unit luminosity is higher in early types than in late types \\citep{Harris&vandenBergh81}. If E+As are in transition from late to early types, one should find new star clusters formed during the starburst. Are there such clusters, and, if so, do their colors and numbers coincide with the expected evolution of the globular cluster systems of present-day E/S0s? Fourth, early-type galaxies lie on the Fundamental Plane (FP), an empirical scaling relation between the effective radius, the central velocity dispersion, and the mean surface brightness, with remarkably small scatter \\citep{Djorgovski87, Dressler87}.  Will E+As lie on the same FP once they evolve?  To determine whether E+As evolve into objects that are indistinguishable from the bulk of E/S0s, we present {\\sl HST/ACS} observations of the 15 remaining E+A galaxies from the \\citet{Zabludoff96} sample.  We combine these with the previous {\\sl HST/WFPC2} observations of five blue E+As \\citep{Yang04} and of one serendipitously discovered E+A \\citep{Yang06}. The resulting high resolution imaging of the 21 confirmed E+A galaxies in the LCRS sample enables us to study the detailed color morphologies and the properties of the newly formed star cluster candidates at the sub-kpc scale.  The detailed morphologies, color profiles, and cluster populations of many E+As are consistent with the galaxy-galaxy interaction scenario.  We discuss how these features will evolve and then relate this evolution to the properties of E/S0s.  Using existing kinematic data \\citep{Norton01} and our {\\sl HST} photometry, we also address whether or not the various scaling relations of E+As are consistent with those of E/S0s.  This paper is organized as follows.  We describe our E+A galaxy sample and the {\\sl HST/ACS} data reduction in \\S \\ref{sec:observation}. In \\S \\ref{sec:morphology}, we examine the general morphology of E+As, including a discussion of tidal features (\\S \\ref{sec:qualitative_morphology}), surface brightness profiles (\\S \\ref{sec:fitting}), structural parameters (\\S\\ref{sec:profile}), and concentration/asymmetry measures (\\S \\ref{sec:ca}).  The color profiles, including a class of E+As with luminous blue cores \\citep{Yang06}, are presented in \\S \\ref{sec:color_profile}.  We compare the scaling relations of E+As with those of E/S0s in \\S \\ref{sec:fp}.  We present the properties of the newly-formed young star clusters in \\S \\ref{sec:cluster}.  We summarize in \\S\\ref{sec:conclusion}.    \\begin{figure*} \\epsscale{0.9} \\plotone{f1.low.ps} % \\caption{ ({\\it Left}) High-contrast $R$ band (\\R{702}) images show the low surface brightness tidal features.  ({\\it Middle})  $R$ band images for the {\\sl WFPC2} sample (EA01AB -- EA05). ({\\it Right}) Residual $R$ band images subtracted from the smooth symmetric model components.  We bound each image with 4 arcsec tickmarks and include a 4 kpc horizontal scalebar.  Note the diverse morphologies of E+A galaxies: tidal and disturbed features, dusty galaxies, blue-cores, bars, and even compact star clusters.  Because EA01A is too disturbed to be modeled by axisymmetric models, we restrict our analysis to EA01B and show the residual image only for EA01B in the top panel. \\label{fig:images_wfpc2}} \\end{figure*}  \\begin{figure*} \\epsscale{0.9} \\plotone{f2a.low.ps} % \\caption{ ({\\it Left})   Same as for Fig \\ref{fig:images_wfpc2}, except these are {\\sl ACS} images of EA06--20. ({\\it Middle}) Two-color composite images from the $B$ and $R$ bands. ({\\it Right})  Same as for Fig \\ref{fig:images_wfpc2}, except for the {\\sl ACS} images. \\label{fig:images_acs} } \\end{figure*}  \\begin{figure*} \\addtocounter{figure}{-1} \\epsscale{0.9} \\plotone{f2b.low.ps} \\caption{Continued.} \\end{figure*}  \\begin{figure*} \\addtocounter{figure}{-1} \\epsscale{0.9} \\plotone{f2c.low.ps} \\caption{Continued.} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusion}  We study the detailed morphologies of 21 E+A galaxies using high resolution {\\sl HST}/{\\sl ACS} and {\\sl WFPC2} images to investigate into what E+A galaxies will evolve after their young stellar populations fade away in a few Gyr.  Our findings are: \\\\   1. The morphologies of E+As are extremely diverse, ranging across train-wrecks, barred galaxies, and blue-cores to relaxed disky galaxies. Most of these galaxies lie in the field, well outside of rich clusters, and at least 11 (55\\%) have tidal or other disturbed features.  Our sample includes one binary E+A system, in which both E+As are tidally disturbed and interacting with each other.  These results support the picture in which galaxy-galaxy tidal interactions or mergers are responsible for triggering the E+A phase in many cases.   2. E+As are bulge-dominated systems (median bulge fraction $B/T$ = 0.59) and their light is highly concentrated (S\\'ersic index $n \\gtrsim 5$). When dust is negligible (at least 67\\% of the time), E+As have high concentration indices ($C \\gtrsim 4.3$) consistent with those of spheroids, but considerably larger asymmetry indices ($A \\gtrsim 0.04$) than ellipticals due to structures within a few $r_e$ that presumably arise from the starburst and/or recent merger.  Thus E+As would be morphologically classified as early-type galaxies once these disturbances relax and the low surface brightness tidal features dissipate or fade.   3. The color morphologies of E+As are as diverse as their structural morphologies. A large fraction (70\\%) have positive color gradients (bluer toward center), indicating that their young stellar populations are more concentrated than their older populations.  We demonstrate that evolution can invert these gradients into the negative gradients typical of E/S0s if the inner parts of E+As have become more metal enriched than the outer parts due to the centralized star formation.   4. Six E+As (30\\%) exhibit compact (0.4--1.4 kpc) blue cores, which might be the local analogs of the high-$z$ elliptical blue-cores \\citep{Menanteau01a}.  We discovered LINERs in three of these blue-core E+As \\citep{Yang06}, and the relationship between LINERs and blue-cores could be an important clue to what stops the star formation in E+A galaxies.   5. E+As stand apart from the E/S0 fundamental plane (FP) in the edge-on projection, implying that the stellar populations of E+As are different from that of E/S0s. E+As have, on average, a \\ml that is 3.8 times smaller than that of E/S0s.  The tilt of the E+A FP indicates that the variation of the stellar populations among E+As is closely tied to their structural parameters, i.e., E+As follow their own scaling relationships such that smaller or less massive galaxies have smaller {\\sl M/L}. Such a trend arises naturally within a merger scenario, where low mass galaxies (the progenitors of low-mass E+As) have higher gas fractions \\citep{Young&Scoville91} and could produce relatively larger populations of young stars.       6. We find a population of unresolved compact sources in at least nine E+A galaxies (45\\%). The colors and luminosities of these young star cluster candidates are consistent with the ages inferred from the E+A spectra (0.01 -- 1 Gyr).  The bright end of the cluster luminosity function fades as the host galaxy becomes redder, suggesting that the newly-formed young star clusters age in parallel to their host. This interpretation is confirmed by the color evolution of the cluster systems.   We have now examined the full set of E+As from the Las Campanas Redshift Survey and so have representative results for local E+A galaxies.  We have used high spatial resolution images to probe their detailed morphologies. We find that their properties are either consistent with those of E/S0s or, if left to evolve passively, will become like those of early-types. The morphologies, color profiles, scaling relations, and young star clusters suggest that E+As galaxies are caught in the act of transforming from late-type to early-type galaxies."
    },
    "0801/0801.4583_arXiv.txt": {
        "abstract": "We study the long term evolution of magnetic fields generated by an initially unmagnetized collisionless relativistic $e^+e^-$ shock. Our 2D particle-in-cell numerical simulations show that downstream of such a Weibel-mediated shock, particle distributions are approximately isotropic, relativistic Maxwellians, and the magnetic turbulence is highly intermittent spatially, nonpropagating, and decaying. Using linear kinetic theory, we find a simple analytic form for these damping rates.  Our theory predicts that overall magnetic energy decays like $(\\omega_p t)^{-q}$ with $q \\sim 1$, which compares favorably with simulations, but predicts overly rapid damping of short wavelength modes.  Magnetic trapping of particles within the magnetic structures may be the origin of this discrepancy. We conclude that initially unmagnetized relativistic shocks in electron-positron plasmas are unable to form persistent downstream magnetic fields.  These results put interesting constraints on synchrotron models for the prompt and afterglow emission from GRBs. ",
        "introduction": "The prompt emission and afterglows of gamma-ray bursts (GRBs) may be manifestations of ultrarelativistic shock waves.  These shock waves may be mediated via the relativistic form of Weibel instability (Weibel 1959; Yoon and Davidson 1987; Medvedev and Loeb 1999; Gruzinov and Waxman 1999).  The free energy from strong plasma anisotropy in the shock transition layer generates strong magnetic fields (with strengths comparable to the available free energy).  However, these fields have very small spatial scales, i.e., the order of the plasma skin depth, $c/\\omega_p$, where $\\omega_p$ is the plasma frequency. These initially small-scale B-fields must survive for tens of thousands to millions of inverse plasma periods to serve as this source of the magnetization for synchrotron models of burst emission and afterglows (Gruzinov \\& Waxman 1999; Piran 2005ab; Katz, Keshet, \\& Waxman 2007). Whether or not these field can is an open question.   Numerical and analytic studies (Kazimura {\\it et al.} 1998; Silva {\\it et al.} (2003); Frederiksen {\\it et al.}  2004; Medvedev {\\it et al.} 2005; Hededal {\\it et al.}  2005; Nishikawa {\\it et al.}  2003, 2005; Spitkovsky this proceedings) have elucidated the basic physics. The instability initially forms filaments of electric current and $B$ fields, which then merge to {\\it inverse cascade} magnetic energy to larger scales, but only in the {\\it foreshock} region.  When the B-fields reach the magnetic trapping limit (Davidson {\\it et al.} 1972; Kato 2005; also see Milosavljevic, Nakar, \\& Spitkovsky 2006; Milosavljevic \\& Nakar 2006a), particle orbits become chaotic, disorganizing the filaments.  The disorganized magnetic fluctuations scatter their supporting particles, which isotropizes and thermalizes the flow, within tens to hundreds of skin depths (Spitkovsky 2005). The magnetic energy peaks in this layer at $\\sim$10-20\\% of the bulk plasma flow energy.  However, present simulations have not deeply followed the flow into the downstream region to explore the long term behavior of these B-fields.  Thus, the question of the structure and long-term survival of the B-fields remains open (see for instance, Gruzinov \\& Waxman 1999; Gruzinov 2001ab; Medvedev {\\it et al.} 2005).  In this proceeding, we discuss recent work (Chang, Spitkovsky, and Arons 2008; hereafter CSA08) which shows that this magnetic energy must rapidly decay in the downstream medium.  We first describe the basic features of the downstream plasma from our numerical simulations. We then calculate the evolution (decay) of the downstream plasma using Vlasov linear response theory and then compare this evolution with simulations.  While linear theory does reasonably well in estimating the decay rate of the total magnetic energy, it overestimates the damping rate of shorter wavelength modes.  We discuss this discrepancy as a result of magnetic trapping. Finally, we summarize our results. ",
        "conclusions": "\\label{sec:discussion}  We have studied the downstream evolution of magnetic turbulence in the context of a collisionless $e^+e^-$ shock both analytically and numerically.  Our simulations show that the downstream region consist of nonpropagating magnetic clumps embedded in quasi-homogenous medium where the background particle distribution function is an isotropic Maxwellian.  In such a background, we showed that magnetic energy will decay like $t^{-q}$ with $q\\sim 1$.  However, linear theory overpredicts the decay rates at short wavelengths compared to simulations.  Magnetic trapping may play an role in resolving this discrepancy.  Rapid field decay puts severe constraints on GRB emission mechanism, but they may not be inconsistent with GRB observations (see Pe'er \\& Zhang 2006).  Finally, if ion-electron collisionless shocks reach roughly equipartition with each other as suggest by recent large scale simulations (Spitkovsky 2008), they would reproduce the physics of the $e^{\\pm}$ shock and their B-fields would decay as well."
    },
    "0801/0801.3530_arXiv.txt": {
        "abstract": "We use semi-analytic models implemented in the {\\it Millennium Simulation} to analyze the merging histories of dark matter haloes and of the galaxies that reside in them. We assume that supermassive black holes only exist in galaxies that have experienced at least one major merger. Only a few percent of galaxies with stellar masses less than $M_* < 10^{10} M_{\\odot}$ are predicted to have experienced a major merger and to contain a black hole. The fraction of galaxies with black holes increases very steeply at larger stellar masses. This agrees well with the observed strong mass dependence of the fraction of nearby galaxies that contain either low-luminosity (LINER-type) or higher-luminosity (Seyfert or composite-type) AGN. We then investigate when the major mergers that first create the black holes are predicted to occur. High mass galaxies are predicted to have formed their black holes at very early epochs. The majority of low mass galaxies never experience a major merger and hence do not contain a black hole, but a significant fraction of the supermassive black holes that do exist in low mass galaxies are predicted to have formed recently. ",
        "introduction": "\\label{sec:intro}  By studying active galactic nuclei (AGN), we learn about the physical mechanisms that trigger accretion onto the central supermassive black holes of galaxies. When a black hole accretes, it increases in mass. By studying populations of AGN at low and at high redshifts, we hope to infer the history of how black holes build up their mass.  It has been established that supermassive black holes most occur in galaxies with bulges \\citep{kormendy1995}, and that the mass of the black hole correlates with the luminosity and the stellar velocity dispersion of the host bulge \\citep{magorrian1998, ferrarese2000, gebhardt2000}. This indicates that the formation of galaxies and supermassive black holes are likely to be closely linked. In the local Universe, the fraction of bulge-dominated galaxies hosting AGN decreases at lower stellar masses \\citep{ho1997, kauffmann2003}. In order to form a black hole, it is necessary for gas to lose angular momentum and sink to the centre of the galaxy \\citep{haehnelt1993, volonteri2003}. The gravitational torques that operate during galaxy-galaxy mergers are known to be a very effective mechanism for concentrating gas at the centers of galaxies \\citep{mihos1996}. Models for AGN evolution have often assumed that black holes are formed and fuelled, and AGN activity is triggered during major mergers of galaxies \\citep{kauffmann2000, wyithe2003, croton2006}.  At low and moderate redshifts, there is no conclusive observational evidence that mergers play a significant role in triggering AGN activity in galaxies. In the local Universe, \\citet{li2006} have shown that narrow line AGN do not have more close companions than matched samples of inactive galaxies. Even at intermediate redshifts (z $\\sim 0.4-1.3$), moderate luminosity AGN hosts do not have morphologies indicative of an ongoing merger or interaction \\citep{hasan2007}. The conclusion seems to be that although major mergers may be responsible for AGN activity in some galaxies, other fueling mechanisms are likely to be most important in the low redshift Universe. It has also been established that high mass black holes have largely stopped growing at early cosmic epochs, whereas low mass black holes are still accreting at significant rates today \\citep{heckman2004}. X-ray observations show that very high-luminosity AGN activity peaked at early cosmic epochs ($z \\sim 2$), while low-luminosity AGN activity peaks at lower redshifts \\citep{steffen2003,barger2005,hasinger2005}.  It has been postulated that this so-called ``anti-hierarchical'' growth of supermassive black holes can be explained if there are two modes of accretion onto black holes that have very different efficiencies \\citep{merloni2004, mueller2007}. The early formation formation of ``new'' black holes may result in very luminous quasar-like events. To form a supermassive black hole, a more violent process such as a major merger may be required to funnel a large amount of gas into the central region of the galaxy. Subsequent accretion of gas onto already existing black holes may be an inefficient process and produce lower luminosity AGN \\citep{haehnelt1993, duschl2002}. The history of accretion after the black hole is formed may not necessarily be tightly linked to the dynamical history of the galaxy, but may be controlled by the accretion and feedback processes occurring in the vicinity of the black hole itself.  In this work, we use the combination of the {\\it Millennium Simulation} and semi-analytic models of galaxy formation to study the fraction of galaxies that have undergone major mergers as a function of mass and cosmic epoch. We investigate whether this can be related to the demographics of black holes in the local Universe and to the apparent disappearance of the most luminous quasar activity in massive galaxies % at late times.  In Sec.~\\ref{sec:simulation}, we briefly introduce the simulation we use and explain how galaxy mergers are tracked in the simulation. In Sec.~\\ref{sec:merger}, we show that if we assume that black holes only form when galaxies undergo major merging events, then most present-day low mass galaxies are predicted not not to contain black holes and hence will not host AGN. In Sec.~\\ref{sec:firstmerger}, we use the simulations to predict when galaxies of different masses have underdone their first major merger. Conclusions and discussions are presented in the final section. ",
        "conclusions": "\\label{sec:discussions}  We analyze the merger histories of dark matter haloes and galaxies in the {\\em Millennium Simulation} and use our results to try to understand the demographics of black holes in nearby galaxies. Black holes are assumed to form only if a major merger occurs. Although a significant fraction of low mass ($< 10^{10} M_{\\odot}$) galaxies have experienced minor mergers, less than a few percent are predicted to have experienced a major merger. If our assumption that a major merger is required in order to form a black hole is correct, the majority of low mass galaxies are predicted not to contain black holes at the present day. This is one possible explanation of the observed lack of AGN in low mass galaxies \\citep{ho1997, kauffmann2003}.  We also investigate when galaxies of different stellar masses are predicted to have formed their first black holes. High mass galaxies form their first black holes at very early epochs. The distribution of formation times is almost flat as a function of lookback time for low mass galaxies. This means that if a low mass galaxy has a black hole, there is a significant probability that it formed in the last few Gigyears. We also compute the number density of newly formed black holes as a function of redshift. We find that the peak number density occurs at $z \\sim2-3$, in good agreement with the observed peak in the quasar space density. More detailed predictions for how AGN of different luminosities are expected to evolve requires a more detailed physical model for how the black holes accrete gas over the history of the Universe. In addition, in certain wavebands AGN activity might be obscured by gas and dust surrounding black hole \\citep{hopkins2006}. More detailed consideration of these issues will form the basis for future work."
    },
    "0801/0801.2911_arXiv.txt": {
        "abstract": "We study the tidal effects of a Kerr black hole on a neutron star in black hole-neutron star binary systems using a semi-analytical approach which describes the neutron star as a deformable ellipsoid. Relativistic effects on the neutron star self-gravity are taken into account by employing a scalar potential resulting from relativistic stellar structure equations.  We calculate quasi-equilibrium sequences of black hole-neutron star binaries, and the critical orbital separation at which the star is disrupted by the black hole tidal field: the latter quantity is of particular interest because when it is greater than the radius of the innermost stable circular orbit, a short gamma-ray burst scenario may develop. ",
        "introduction": "During the past decades, theorists have been modelling various kinds of double compact objects since (1) they are among the most promising gravitational wave sources to be detected by ground-based and space-based laser interferometers \\cite{Detectors} and (2) they have also been invoked as possible engines of short gamma-ray bursts \\cite{GRB1} (see also \\cite{GRB,LeeGRB}) in the case of black hole-neutron star (BH-NS) and neutron star-neutron star (NS-NS) mergers. The remnants of both kinds of mergers, in fact, may result in a black hole with negligible baryon contamination along its polar symmetry axis and surrounded by a hot massive accretion disk: before the disk gas is accreted to the black hole, intense neutrino fluxes are emitted which, through energy transfer, trigger a high-entropy gas outflow off the surface of the accretion disk (``neutrino wind''); at the same time, energy deposition by $\\nu\\bar\\nu$ annihilation in the baryon-free funnel around the rotation axis, powers relativistically expanding $e^\\pm\\gamma$ jets which can give rise to gamma-ray bursts \\cite{GRB1}. The fate of BH-NS binaries, in particular, depends on the relative values of $r_{ISCO}$, the radius of the innermost stable circular orbit, and $r_{tide}$, the orbital separation at which the tidal disruption of the star by the BH occurs: if $r_{ISCO}<r_{tide}$ the star is disrupted and then swallowed by the BH, otherwise it is swallowed without disruption. This is therefore a crucial issue for the SGRB mechanism we have described: only if $r_{ISCO} < r_{tide}$ the merger may result in the black hole + hot massive accretion disk + baryon-free funnel SGRB scenario.  A significant effort in studying BH-NS coalescing systems has been undertaken in order to model them as gravitational wave sources and to understand gamma-ray burst engines. Unfortunately, BH-NS as well as BH-BH binaries (as opposed to NS-NS binaries) have not been observed yet and therefore it is not currently possible to infer their properties from observational data. Thus, in modelling the behaviour of such binaries and in making predictions about them, one has to rely on binary evolution and population synthesis models. Several approaches addressing different issues have been employed to study coalescing binaries. Post-Newtonian (PN) studies of compact binaries (\\cite{Blanchet} and references therein), for example, typically approximate the constituents of the binary as point sources. PN expansions may not converge rapidly enough in the strong-field region and thus are indicated for the inspiral phase: this has lead to a strong effort in stitching together PN studies with other methods that are more indicated for addressing the merging and ringdown, e.g. \\cite{BertiIyerWill}. Other approaches deal with finite size effects due to at least one of the binary constituents; many of them assume Newtonian gravity in some or all aspects of the calculation. Finite size-effects are connected to tidal interactions between the binary constituents: these interactions are present well before the final phases of the coalescence and have been studied in the literature using various approaches and addressing several issues. We refer the reader to \\cite{Christian-articolo} for an overview on the literature.  Nowadays, the mainstream strategy in studying BH-NS binaries is to develop fully relativistic codes, as is being done over the years by several groups (see \\cite{GR_eg},\\cite{TBFS08} for recent work). However the high computational cost of simulations of compact binaries makes it desirable to have frameworks to study certain features and dynamical regimes which admit the use of approximations, thus allowing a drastic reduction of the computational resources required. Motivated by this reason and by the intent of investigating SGRB engines, in this paper we focus on BH-NS binaries and develop a semi-relativistic model in order to describe the behaviour of a NS undergoing tidal interactions with its BH companion.  Our starting point is the model developed in \\cite{WL2000} by Wiggins and Lai: \\begin{itemize} \\item it is essentially the \\emph{affine model} approach of Carter, Luminet and Marck \\cite{CL85},\\cite{LM85}, i.e. the NS is treated as a \\emph{Newtonian} extended object which responds to its self-gravity, to its internal pressure forces and to the \\emph{relativistic} BH tidal field under the constraint that its shape is always that of an ellipsoid (hence the name ``affine model''); this constraint descends from the fact that ellipsoidal figures are solutions to the dynamical equations describing a Newtonian incompressible star responding to the linear terms of a BH tidal field (see \\cite{EFE} for an extensive explanation). \\item NSs on equatorial circular orbits around a BH are considered and the quasi-equilibrium sequences of a \\emph{Newtonian polytropic} NS\\footnote{Wiggins and Lai also consider white dwarfs, but we are not interested in white dwarfs in this paper.} interacting with a Kerr BH are determined from large separation until the star tidal disruption: the critical orbital radius at which the tidal disruption occurs is compared to the radius of the ISCO in the context of SGRB engines. \\end{itemize} A similar approach has been used in \\cite{LS76}.  In this paper we improve their model in two directions by \\begin{enumerate} \\item including general relativistic effects in the response of the NS to the BH tidal field \\item using a general, barotropic equation of state (EOS) for the NS. \\end{enumerate} Obviously, both points arise from the desire of building a model which has as many ``realistic'' features as possible; however there is more. The NS disruption, which may lead to the SGRB mechanism previously described, occurs when the BH tidal force starts prevailing on the NS self-gravity; this condition may be approximately expressed as $\\alpha M_{NS}/R^2_{NS} \\simeq M_{BH}R_{NS}/r_{tide}^3$, where $\\alpha$ is a dimensionless coefficient, $M_{NS}$ and $R_{NS}$ are the NS mass and radius, and $M_{BH}$ is the BH mass.  $r_{tide}$ hence depends essentially on two parameters: the mass ratio $q=M_{BH}/M_{NS}$ and the NS compactness at equilibrium $C=M_{NS}/R_{NS}$. The main advantage of Wiggins and Lai's model as opposed to many fully relativistic formulations, is that it allows one to choose large values of $q$.  On the other hand, the weak point in ellipsoidal models in general, is that Newtonian self-gravity is adopted, and this choice is inappropriate to describe very compact stars.  In our approach we improve this point by approximating the effects of general relativistic gravity by means of an effective scalar gravitational potential $\\Phi_{TOV}$. This potential (see \\ref{app:pseudoTOV}) stems from the Tolman-Oppenheimer-Volkoff (TOV) stellar structure equations in General Relativity, and has been adopted and tested in the context of stellar core collapse and post-bounce evolution simulations: \\begin{itemize} \\item in \\cite{RamppJanka2002} Rampp and Janka present the \\textsc{Vertex} code for supernova simulations; in order to approximate relativistic gravity, this code makes use of the generalised potential in place of the usual Newtonian potential in all Newtonian hydrodynamics equations \\item in \\cite{Liebendorfer05} Liebend{\\\"o}rfer et al. perform a comparison between results of the \\textsc{Vertex} and \\textsc{Agile-BoltzTran} codes, the latter being a fully relativistic (1D) hydrodynamics code; it is shown that both codes produce qualitatively very similar results except for some small (but growing) quantitative differences occurring in the late post-bounce evolution \\item in order to achieve a better agreement than that reported in \\cite{Liebendorfer05}, different improvements of the aforementioned effective relativistic potential are explored and tested by Marek and collaborators in \\cite{Dimmelmeier2006}; the Newtonian equations of hydrodynamics remain untouched. \\end{itemize} The way we apply the main idea of \\cite{RamppJanka2002}-\\cite{Dimmelmeier2006} in the present, completely different, context is the following: the Eulerian hydrodynamics equations --- which govern the star in the affine model and which are Newtonian --- are left formally unchanged, but the pressure profile of the star at equilibrium appearing in them is built with the TOV equations and gravitational potential used is $\\Phi_{TOV}$.  Furthermore, we extend the affine model to include general barotropic equations of state (see $\\S$ \\ref{sec:model}).  In this paper we concentrate on calculating $r_{tide}$ with our approach.  As a test of the model we reproduce the fully relativistic results obtained in \\cite{TBFS08}, where the tidal disruption limit of binaries containing a Schwarzschild BH and a polytropic relativistic NS is calculated. We subsequently compute $r_{tide}$ for binaries composed of a Kerr black hole and a NS, modeled with more realistic equations of state.  The paper is organized as follows. In Section \\ref{sec:model} we describe the basic equations of the affine model suitably modified as explained above; in Section \\ref{sec:results} we test our model and present the results of the numerical simulations; in Section \\ref{sec:conclusions} we draw the conclusions. ",
        "conclusions": "\\label{sec:conclusions} In this paper we have studied the effects of a Kerr BH tidal field on a NS; we have focused on determining the orbital separation at which the NS is tidally disrupted ($r_{tide}$) and the star quasi-equilibrium sequence for several BH-NS binaries.  Our work is based on the affine model, which we have improved with respect to previous works, by describing the NS self-gravity with an effective relativistic scalar potential (Eqs.\\,(\\ref{def:dPhiTOVdr}) and (\\ref{def:mTOV})), and by using more realistic equations of state for the NS matter.  This approach has the advantage of allowing a quick computation of $r_{tide}$ and of the quasi-equilibrium sequences for any value of the binary mass ratio $q$ and of the BH spin parameter $a$.  A comparison of the results obtained using our approach for NS disruption with Full-GR results, which are available in the case of non-rotating BHs and polytropic NSs, shows an excellent agreement up to small values of $q$ (which depend on the NS compactness), beyond which our model starts to \\emph{underestimate} $r_{tide}$; this happens because we assume that the centre of mass of the star moves on along a BH geodesic, a condition which is better satisfied for larger mass ratios.  Our model agrees with relativistic calculations much better than ellipsoidal models in which the neutron star self-gravity is treated at a Newtonian level.  Given the very good results of this test, we have determined the quasi-equilibrium sequences and the tidal disruption radii of several BH-NS binaries in order to evaluate with our model the role played by (1) the BH spin parameter $a$ and (2) the equation of state of the NS fluid.  We found that the soft gamma-ray burst scenario is favoured by \\emph{low} values of the mass ratio $q$ and \\emph{high} values of the BH spin parameter $a$.  In addition, the higher the BH spin parameter is, the more the values of the NS axes $a_2$ and $a_3$ (i.e. the principal axes which do not point at the BH) tend to coincide.  In quantitative terms we have shown that, for the stiffer equation of state APR2 a SGRB may occur if $q\\lesssim 28$ when $a=0.99\\,M_{BH}$ or if $q\\lesssim 5.7$ when $a=0.5\\,M_{BH}$.  For the softer equation of state BALBN1H1 a SGRB may develop if $q\\lesssim 33$ when $a=0.99\\,M_{BH}$ or if $q\\lesssim 6.6$ when $a=0.5\\,M_{BH}$.  As a future development of our work, we intend to study the dynamical behaviour of BH-NS binaries and to determine the gravitational radiation emitted in the late phases of their inspiral. This means taking into account tidal interaction corrections to the orbital dynamics and to the gravitational waveform, and thus will require the inclusion in our model of a precise treatment of the orbit. Moreover we would like to improve the results of our model for low values of the mass ratio $q$.   \\appendix"
    },
    "0801/0801.2120_arXiv.txt": {
        "abstract": "{We analyze the  distribution of G and K  type stars towards the  Galactic poles  using RAVE  and ELODIE  radial  velocities, 2MASS photometric star counts, and UCAC2 proper motions.  The combination of photometric  and  3D  kinematic  data  allows us  to  disentangle  and describe the vertical distribution of dwarfs, sub-giants and giants and their kinematics.  We identify  discontinuities within the kinematics and magnitude counts that separate the  thin disk,  thick disk and  a hotter  component.  The respective scale heights of the thin  disk and thick disk are 225$\\pm$10\\,pc and 1048$\\pm$36\\,pc. We also constrain the luminosity function  and  the kinematic distribution  function. The existence of a  kinematic gap between the thin and thick  disks is incompatible with the thick disk having formed  from the  thin disk by  a continuous process,  such as scattering  of  stars by  spiral  arms  or  molecular clouds. Other mechanisms of formation of the thick disk such as `created on the spot' or smoothly `accreted' remain compatible with our findings. ",
        "introduction": "It  is  now widely  accepted  that  the  stellar density  distribution perpendicular  to  the  Galactic  disk  traces at  least  two  stellar components,  the thin and  the thick  disks.  The  change of  slope in the logarithm of the vertical density distributions at $\\sim$ 700\\,pc (Cabrera-Lavers et  al. \\cite{cab05}) or $\\sim$ 1500 pc (Gilmore  \\&  Reid  \\cite{gil83})  above  the Galactic  plane is  usually explained as  the signature of  a transition between these two  distinct components: the thin and  the thick disks. The thick disk is an  intermediate stellar population between the thin disk and  the stellar halo, and  was initially defined  with the other stellar  populations  by combining  spatial,  kinematic and  abundance properties (see a summary of  the Vatican conference of 1957 by Blaauw \\cite{bla95} and Gilmore \\& Wyse \\cite{gil89}).  Its properties are described  in a long  series of publications with  often diverging characteristics (see the analysis  by Gilmore \\cite{gil85}, Ojha \\cite{ojh01}, Robin et al.  \\cite{rob03} and also by Cabrera-Lavers et al.   \\cite{cab05}, that  give  an overview  of recent  improvements). Majewski (\\cite{maj93}) compared a nearly exhaustive list of scenarios that describe many possible formation mechanisms for the thick disk.  In this  paper, we attempt to  give an answer  to the simple but  still open questions:  are the  thin and  thick disks  really two distinct components?  Is there any continuous transition between them? These questions were  not fully settled by analysis  of star counts by Gilmore \\& Reid (\\cite{gil83}) and later workers. Other  important   signatures  of   the  thick  disk   followed  from kinematics:  the  age--velocity   dispersion  relation  and  also  the metallicity--velocity dispersion relation.  However the identification of  a thin--thick  discontinuity depends  on the  authors, due  to the serious difficulty of assigning accurate ages to stars (see Edvardsson et  al.   \\cite{edv93} and Nordstr\\\"om et al. \\cite{nor04}).   More recently   it  was  found   that  the [$\\alpha$/Fe] versus [Fe/H] distribution  is related to the kinematics (Fuhrmann  \\cite{fur98}; Feltzing  et al.   \\cite{fel03};  Soubiran \\& Girard \\cite{sou05};  Brewer \\& Carney \\cite{bre06}; Reddy et al. \\cite{red06})  and provides an effective  way  to  separate  stars  from  the  thin  and  thick  disk components.  Ages and abundances are important to describe the various disk components  and to depict  the mechanisms of their  formation.  A further complication comes from the recent indications of the presence of  at  least  two   thick  disk  components  with  different  density distributions,   kinematics    and   abundances   (Gilmore    et   al. \\cite{gil02};    Soubiran    et    al.     \\cite{sou03};    Wyse    et al. \\cite{wys06}).  Many of the  recent works favor the presently  prevailing scenarios of thick disk formation by the  accretion of small satellites, puffing up the early stellar Galactic disk or tidally disrupting the stellar disk (see for example Steinmetz  \\& Navarro \\cite{ste02}; Abadi et al. \\cite{aba03};  Brook  et al.   \\cite{bro04}).   We  note however  that chemodynamical  models of  secular Galactic  formation including extended ingredients  of stellar formation  and gas dynamics  can also explain  the formation of  a thick  disk distinct  from the  thin disk (Samland \\& Gerhard \\cite{sam03}; Samland \\cite{sam04}).  In this paper,  we use the recent RAVE  observations of stellar radial velocities, combined  with star counts and proper  motions, to recover and  model the  full 3D  distributions of  kinematics and  densities for nearby   stellar   populations.  In  a   forthcoming   study, metallicities  measured from  RAVE  observations will  be included  to describe  the galactic  stellar populations  and their  history.  The description of data is given in Sect.~2, the model in Sect.~3, and the interpretation  and  results  in  Sect.~4.  Among  these  results,  we identify discontinuities visible  both within the density distributions and the kinematic  distributions.  They  allow to  define more precisely the  transition between the thin and  thick stellar Galactic disks. ",
        "conclusions": "We revisit  the thin-thick disk transition using  star counts and kinematic data towards  the Galactic poles.  Our Galactic modeling of star count, proper motion and radial velocity allows us to recover the LF, their kinematic distribution  function, their vertical density distribution, the relative distribution of giants, sub-giants and  dwarfs, the  relative contribution from  thin  and thick  disk components, the  asymmetric drift  coefficient and the  solar velocity relative to the LSR.  The double exponential fitting of the vertical disk stellar density distribution is not sufficient to fully characterize the thin and thick disks. A more complete description of the stellar disk is given by its  kinematical decomposition.  From  the star  counts, we see a  sharp transition  between the thick and  thin components.  Combining star counts with kinematic data, and applying a  model with 20 kinematic components,  we discover a gap between the vertical velocity dispersions of thin disk components with $\\sigma_W$  less  than  21  km\\,s$^{-1}$  and a  dominant  thick  disk component at $\\sigma_W$=45.5km\\,s$^{-1}$.  The thick disk scale height is found to be 1048$\\pm$36\\,pc. We identify this thick  disk with the intermediate metallicity ([Fe/H] $\\sim$\\,--0.6  to  --0.25)  thick  disk described,  for  instance,  by Soubiran et  al. (\\cite{sou03}).  This  thick disk is also  similar to the thick disk measured by Vallenari et al (\\cite{val06}) who find \"no significant velocity  gradient\" for  this stellar component.   We note that star  counts at  $m_{\\rm K}\\sim15$  suggest a  second thick disk or halo component with  $\\sigma_W\\sim$65km\\,s$^{-1}$.  Due to the separation of the thin and thick components, clearly identified with stars counts and visible within the kinematics, the thick disk measured in this paper cannot be the result of dynamical heating of the thin disk by massive molecular clouds or by spiral arms. We would expect otherwise a continuous kinematic distribution function with significant kinematic components covering without discontinuity the range of   $\\sigma_W$ from 10 to 45 km\\,s$^{-1}$.  We find that, at the solar position, the surface mass density of the thick disk is 27\\% of the surface mass density of the thin disk. The thick disk has velocity dispersions $\\sigma_U=50\\,$km\\,s$^{-1}$,     $\\sigma_W$=45.5\\,km\\,s$^{-1}$,    and asymmetric drift $V_{\\rm lag}=33 \\pm 2$\\,km\\,s$^{-1}$. Although clearly  separated from the  thin disk, this  thick component remains a  relatively `cold' thick  disk and has  characteristics that are close to the thin disk properties. This  `cold'  and  rapidly  rotating  thick disk  is  similar  to  the component identified by many kinematic  studies of the thick disk (see Chiba \\& Beers  \\cite{chi00} for a summary). Its  kinematics appear to be  different  from  the  thick  disk stars  studied  at  intermediate latitudes  in pencil  beam surveys  (eg Gilmore  et  al \\cite{gil02}), which  appear to be  significantly affected  by a  substantial stellar stream with  a large lag  velocity. They interpret  this stellar stream  as  the possible  debris  of  an  accreted satellite  (Gilmore \\cite{gil02};    Wyse    et al. \\cite{wys06}). Maybe some  connections exist with streams identified in  the solar  neighborhood as  the  Arcturus stream  (Navarro et  al \\cite{nav04}).  Some  mechanisms of  formation connecting  a thin  and a   thick components are compatible with our  findings.  It may be, for instance a `puffed-up' thick disk, i.e. an earlier thin disk puffed up by the accretion of  a  satellite (Quinn  et  al.  \\cite{qui93}).  Another  possibility, within the monolithic  collapse scenario, is a thick  disk formed from gas with  a large vertical scale  height before the  final collapse of the gas  in a thin disk,  i.e.  a `created  on the spot' thick disk.   We also notice the Samland (\\cite{sam04})  scenario: a chemodynamical model of formation of  a disk galaxy within  a growing dark  halo that provides both a `cold'  thick disk and a metal-poor `hot' thick disk.  A  popular scenario  is  the  `accreted' thick  disk  formed from  the accretion of satellites.  If the thick disk results from the accretion of just a  single satellite, with a fifth of the  mass of the Galactic disk, this has been certainly a major event in the history of the Galaxy,  and it  is hard  to  believe that  the thin  disk could  have survived this upheaval.  Finally,  from the thick  disk properties  identified in this paper, we  can reject the most improbable  scenario of formation: the  one of type  `heated' thick disk  (by molecular  clouds or spiral arms)."
    },
    "0801/0801.2585_arXiv.txt": {
        "abstract": "The $\\sigma$~Ori cluster is an unbound aggregate of a few hundred young, low--mass stars centered on the multiple system $\\sigma$~Ori. This cluster is of great interest because it is at an age when roughly half of the stars have lost their protoplanetary disks, and the cluster has a very large population of brown dwarfs.  One of the largest sources of uncertainty in the properties of the cluster is that the distance is not well known.  The directly measured Hipparcos distance to $\\sigma$~Ori~AB is 350$^{+120}_{-90}$~pc.  On the other hand, the distance to the Orion~OB1b subgroup (of which $\\sigma$~Ori is thought to be a member), 473$\\pm$40~pc, is far better determined, but it is an indirect estimate of the cluster's distance.  Also, Orion OB1b may have a depth of 40~pc along our line of sight. We use main sequence fitting to 9 main sequence cluster members to estimate a best fit distance of 420$\\pm30$~pc, assuming a metallicity of $-$0.16$\\pm$0.11 or 444~pc assuming solar metallicity.  A distance as close as 350~pc is inconsistent with the observed brightnesses of the cluster members.  At the best fit distance, the age of the cluster is 2--3~Myrs. ",
        "introduction": "The bright O9.5V star \\object{$\\sigma$~Ori} is a trapezium--like system with six known components.  The brightest component, $\\sigma$~Ori~AB (V=3.$^m$80), is a 0.25$^{{\\prime}{\\prime}}$ binary \\citep{horch01} with an O9V primary and a B0.5V secondary \\citep{edwards76} and an orbital period of $\\sim$170~years or $\\sim$158~years \\citep{heintz1974, heintz1997}.  The O9V primary, $\\sigma$~Ori~A, is itself a double lined spectroscopic binary\\footnote{We wish to thank Deane Peterson for sharing unpublished results of his observations of $\\sigma$~Ori~AB with us.} \\citep{peterson07, bolton74, micz1950}.  The spectral type of $\\sigma$~Ori~C is A2V.  The D and E components are B2V stars with V$\\simeq$6.8 and $\\simeq$6.6 respectively (see Table~\\ref{true_table}).  The $\\sigma$~Ori cluster was first recognized as a group of high--mass stars by \\citet{garrison67}, and as cluster of low--mass pre-main-sequence stars by \\citet{wws}.  Continuing work on the cluster has revealed a young cluster of several hundred low--mass stars \\citep{paper1, burn05, kenyon05} and a rich population of brown dwarfs \\citep{bejar99,bejar01,bejar04, zapatero02}.  About one third to one half of the stars retain their accretion disks \\citep{oliveira04, hern2007}.  The cluster is considered part of the Orion OB1b association.  Age estimates for Orion OB1b range from less than 2~Myrs \\citep{brown94} up to 7~Myrs \\citep{blaauw91}.  This is an exceptionally interesting age because it is the age when protoplanetary disks are making the transition from optically thick to optically thin and may be the age when giant planets form.  The more accurately the cluster age can be measured, the tighter the constraint on disk lifetimes and the time available for giant planet formation will be.  See \\citet{handbook} for a recent review of observational data on the low--mass population of the $\\sigma$~Ori cluster.  The most significant source of uncertainty for the age of this cluster is the uncertain distance to the cluster.  Many authors adopt the Hipparcos distance of 350$^{+120}_{-90}$~pc for $\\sigma$~Ori \\citep{hip} as the distance to the center of the cluster.  This has the virtue of being a direct measurement of the distance to the cluster, but it has an uncertainty of 30\\% (see \\citet{schroeder04} for a discussion of biases on Hipparcos parallaxes of O stars).  Others use the Hipparcos distance to the Orion~OB1b subgroup (of which $\\sigma$~Ori is a member), 473$\\pm$40~pc \\citep{dezeeuw99}, as the distance to the cluster.  This value is more precise because it is the average distance to 42 members of the association,   % yet it is only an indirect measurement of the distance to $\\sigma$~Ori.  Similarly, \\citet{hern2005} find a distance of 443$\\pm16$~pc from the Hipparcos parallaxes of the combined Orion~OB1b and 1c subgroups.  The Orion~OBIb association has a size of $\\sim$30--40~pc across the sky, so it is likely to have a similar depth along our line of sight.  The cluster could easily lie $>$20~pc in front of or behind the center of Orion~OBIb.  % It would be preferable to have a direct measure of the cluster's distance that is more precise than the Hipparcos distance to $\\sigma$~Ori.  In a brief abstract, \\citet{garrison67} said that main~sequence fitting to 15 B stars near $\\sigma$~Ori yielded a narrow main~sequence at a distance modulus of 8.2 (440~pc).  Garrison did not correct for the small values of reddening that some of the likely cluster members have.  Garrison does not appear to have ever published a more detailed description of this result.  In this paper we re--examine the main~sequence fitting distance for the $\\sigma$~Ori cluster using published spectroscopy and photometry for the stars that lie within 30$^{\\prime}$ of $\\sigma$~Ori~AB and have spectral types earlier than F0. ",
        "conclusions": "From Figure~\\ref{true_cmd} it is clear that the $\\sigma$~Ori cluster must be more distant than the nominal 350~pc Hipparcos distance for $\\sigma$~Ori.  We estimate a distance of 420$\\pm$30~pc for the cluster, assuming an [Fe/H] % of $-$0.16, or 444~pc assuming solar metallicity.  This is consistent with, but significantly more precise (7\\%) than the Hipparcos distance (30\\%).  Our distance estimate is consistent with the estimated distance to Orion OB1b.  Most of the older age estimates for the cluster assumed a distance of 350~pc.  Our more tightly constrained distance shows that the cluster age must be closer to the young end of the range of the estimated ages.  This places the age of the cluster in the range of 2--3~Myrs.   Sixteen of the 19 stars in our sample are probable members of the $\\sigma$~Ori cluster. HD~37564 is too bright to be a cluster member. HD~294279 is too faint to be a cluster member. If HD~37333 is in fact an equal mass binary, it is likely to be a cluster member.  %  The existence of a tight main sequence among the O, B, and A stars of the $\\sigma$~Ori cluster suggests that any other clusters within the Orion OB1a and Orion OB1b groups (such as the 25~Ori cluster \\citep{briceno05}) should also have tight main sequences."
    },
    "0801/0801.4165.txt": {
        "abstract": "We report observations of the nova RS Ophiuchi (RS Oph) using the Keck Interferometer Nuller (KIN), approximately 3.8 days following the most recent outburst that occurred on 2006 February 12.  These observations represent the first scientific results from the KIN, which operates in N-band from 8 to 12.5 $\\mu$m in a nulling mode.    The nulling technique is the sparse aperture equivalent of the conventional coronagraphic technique used in filled aperture telescopes.  In this mode the stellar light itself is suppressed by a destructive fringe, effectively enhancing the contrast of the circumstellar material located near the star.  By fitting the unique KIN  data, we have obtained an angular size of the mid-infrared continuum of 6.2, 4.0, or 5.4 mas for a disk profile, gaussian profile (FWHM), and shell profile respectively.  The data show evidence of enhanced neutral atomic hydrogen emission and atomic metals including silicon located in the inner spatial regime near the white dwarf (WD) relative to the outer regime.  There are also nebular emission lines and evidence of hot silicate dust in the outer spatial region, centered at $\\sim$ 17 AU from the WD, that are not found in the inner regime.  Our evidence suggests that these features have been excited by the nova flash in the outer spatial regime before the blast wave reached these regions.  These identifications support a model in which the dust appears to be present between outbursts and is not created during the outburst event.  We further discuss the present results in terms of a unifying model of the system that includes an increase in density in the plane of the orbit of the two stars created by a spiral shock wave caused by the motion of the stars through the cool wind of the red giant star. These data show the power and potential of the nulling technique which has been developed for the detection of Earth-like planets around nearby stars for the Terrestrial Planet Finder Mission and Darwin missions. ",
        "introduction": "Classical novae (CN) are categorized as cataclysmic variable stars that have had only one \\textit{observed} outburst - an occurrence typified by a Johnson V-band brightening of between six and nineteen magnitudes \\citep{war95}.  These eruptions are well-modeled as thermonuclear runaways (TNR) of hydrogen-rich material on the surface of white dwarf (WD) primary stars that, importantly, remain intact after the event.  Current theory tells us that CN can be modeled as binary systems in which a lower-mass companion -- the secondary -- orbits the WD primary such that the rate of mass transfer giving rise to the observed eruption is very low.  Recurrent novae (RN) are a related class of CN which have been observed to have more than one eruption.  Like CN, RN events are well-represented as surface TNR on WD primary stars in a binary system, but are thought to have much higher mass transfer rates commensurate with their greater eruption frequency.  There are two types of systems that produce recurrent novae -- CVs, in which the WD accretes from a main sequence star that orbits the WD on a time scale of hours, and symbiotic stars, in which the WD accretes from a red-giant companion that orbits that WD on a time scale of years. %Dwarf novae (DN) are distinct from either of these broad classes in that their eruptions, while high in %frequency, are well-modeled as a sudden release of potential energy from a large clump of material as %it is transfered from secondary to the primary.  CN and RN produce a few specific elemental isotopes by the entrainment of metal-enriched surface layers of the WD primary during unbound TNR outer-shell fusion reactions.  In contrast, type Ia supernovae produce most of the elements heavier than helium in the Universe through fusion reactions leading to the complete destruction of their WD primary.  Some theoretical models indicate that RN could be a type of progenitor system for supernovae.  Importantly, these theories are predicated on two critical factors: 1.) the system primary must be a compact carbon-oxygen core supported solely by electron degeneracy pressure and 2.) there must be some mechanism to allow the WD mass to increase secularly towards the Chandrasekhar limit. %It does appear that the RS Ophiuchi (RS Oph) system is a member of this subclass of RN and we are %very fortunate to have it as a nearby laboratory to advance our understanding of the astrophysics of %these great engines of nucleosynthesis.  The nova RS Oph has undergone six recorded episodic outbursts of irregular interval in 1898 \\citep{fle04}, 1933 \\citep{ada33}, 1958 \\citep{wal58}, 1967 \\citep{bar69}, 1985 \\citep{mor85} and now 2006.  There are also two possible outbursts in 1907 \\cite[]{sch04} and 1945 \\cite[]{opp93}.  All outbursts have shown very similar light curves.  This system is a single-line binary, symbiotic with a red giant secondary characterized as K$5.7 \\pm 0.4$  I-II \\citep{ken87} to a K7 III \\citep{mur99} in quiescence and a white dwarf primary in a $455.72\\pm0.83$ day orbit about their common center of mass as measured using single-line radial velocity techniques \\cite[]{fek00}.  \\citeauthor{opp93} examined all the outbursts and found that V band luminosity of RS Oph decreased averaged 0.09 magnitudes/day for the first 43 days after outburst. A 2-magnitude drop would then require on average 22 days, establishing RS Oph as a \u00d2fast\u00d3 nova based on the classification system of \\cite{pay57}. %:speed of rsoph % see page 130 in notebook 3 for speed calculation using AAVSO data.  The most recent outburst of the nova RS Oph was discovered at an estimated V-band magnitude of 4.5 by H. Narumi of Ehime, Japan on 2006 February 12.829 UT \\cite[]{nar06}.  This is 0.4 mag brighter than its historical average AAVSO V-band {\\it peak} magnitude so it is reasonable to take Feb 12.829 (JD 2453779.329) as day zero.  The speed of an outburst is characterized by its $t_2$ and $t_3$ times which are  the intervals in days from the visible maximum until the system has dimmed by 2 and 3 magnitudes, respectively.   For this outburst the $t_2$ and $t_3$ times are 4.8 and 10.2 days, respectively. %As the WD approaches the Chandrasekhar limit the amount of mass initially ejected by the %thermonuclear runaway (TNR) decreases.   It would not, therefore be unreasonable to expect the %initial V-band speed of the nova to gradually increase.  Indeed it must be the case that the mass of the %WD has grown over time. If earlier in its secular evolution it was a 1.0 $M_{\\sun}$ WD and still had a %20 year recurrence time, its accretion rate would have had to be unacceptably high.  %:distance to rsoph The distance to the RS Oph system is of  importance to the interferometry community as it  effects interpretation of astrometric data (cf. \\citet{mon06}).  There has been a good deal of disagreement   in the literature with a surprisingly broad range, from as near as  0.4 kpc \\citep{hac01} to as far as 5.8 kpc \\citep{pot67}. %As late as 2001, \\citeauthor{hac01} had estimated the nearer distance while early analysis of X-ray %data from the Rossi X-ray Timing Explorer confirmed a distance of 1.6 kpc \\citep{sok06}.    Distance %estimates using absorption line  calculations in the nova's last epoch (\\cite{hje86} \\& \\cite{sni85}) and, %using envelope expansion parallaxes from new radio observations from the current epoch \\citep %{rup08} establish a distance at the mid range. \\citet{bar08} have recently undertaken a thorough review of the various techniques that have been used to derive a distance to RS Oph  and obtain a distance  of $1.4^{+ 0.6}_{-0.2} $ kpc. It is this value that we adopt for  astrometric calculations in this paper.  %The $t_3$ and $t_2$ times noted above do constrain the distance to RS Oph and this is worth noting here. If we restrict ourselves to the use of the $t_3$ time as a more settled figure, and apply the Maximum Magnitude Rate of Decline (MMRD) expression attributed to \\citet{dow00} we obtain a range of distances of $D_{RS Oph} = 2.51 \\pm 1.28$ kpc.  To this broad range we apply an additional constraint. As first noted by \\citet{cas85}, the distance to RS Ophiuchi cannot exceed approximately 2.0 kpc because it's spectrum does not contain evidence of absorption by material in the Carina arm of the Galaxy.  Inclusion of the Galactic upper constraint and the MMRD lower constraint calculated above yields a range of distances to the nova RS Ophiuchi of $1.61 \\pm 0.39$ kpc.  %Note: see journal #3, pages 130 - 134  The structure of this paper is as follows. We report high-resolution N-band observations of  RS Oph using the nulling mode of the 85 meter baseline Keck Interferometer, beginning with a discussion of the nulling mode itself in Section 2.  We discuss the observations in Section 3, and the data and analysis in Section 4.   In Section 5 we  introduce a new physical model of the system, which unifies many of the observations into a coherent framework.  The results of this paper and those of other recent observations of RS Oph are discussed in the context of this model in Section 6. Finally, Section 7 contains a summary of our major results and conclusions.  % note to Bill - we don't specifically take Lane et al into account so I changed the words \"all of the observations\" to \"many...\" ",
        "conclusions": "We  analyzed data from  the recurrent nova RS Oph for the epoch at $\\sim$ 4 days post-outburst using the new KIN instrument. These data allowed us to determine the size of the emitting region around the RS Oph at wavelengths from 8-12 $\\mu$m. By fitting the unique KIN inner and outer spatial regime data, we  obtained an angular size of the mid-infrared continuum of 6.2, 4.0, or 5.4 mas for a disk profile, gaussian profile (FWHM), and shell profile respectively. The data show evidence of enhanced neutral atomic hydrogen emission located in the inner spatial regime relative to the outer regime.  There is also evidence of a 9.7 $\\mu$m silicate  feature seen outside of this region, which is consistent with dust that had condensed prior to the outburst, and which has not yet been disturbed by the blast wave from the nova. Our analysis of the observations, including the new ones presented in this paper, are most consistent with a new physical model of RS Oph, in which spiral shock waves associated with the motion of the two stars through the cool wind from the red giant create density enhancements within the plane of their orbital motion.  Further observations are needed to clarify this new picture of the RS Oph system.  One issue that has not been fully resolved is whether or not the red giant star really overflows its Roche lobe. If so a hot spot would be expected where the material streaming from the red giant envelope hits the accretion disk, and UV or X-ray observations could search for this effect. Another approach would be to observe  RS Oph over several orbital cycles using infrared photometry to look for  variations of the light curve due to the departure of the red giant star from spherical symmetry.  High resolution spectra could also help.   Confirmation of the rotational velocity of the red giant star measured by Zamanov et al. (2007) would be worthwhile, and could provide another estimate of the absolute size of the red giant star, assuming that it is co-rotating with the orbit.  Further KIN observations within the next few years would also be helpful as they could show evidence of the re-establishment of the spiral shock wave, and perhaps some information about the shape of the circumstellar material and dust formation.  Another epoch of HST observations would also determine the deceleration of the outflow in the two directions, i.e., within the plane of the orbit of the two stars and along the poles.  Theoretical studies of the motion of the blast wave in an environment with a high density region in the plane of the orbit of the two stars are also worthwhile.  In particular, it would be important to understand how the blast wave is diffracted around the red giant star and how it propagates in an medium with the periodic density enhancements in the plane due to the spiral pattern.  The recurrent nova RS Ophiuchi is a rich system for the study of circumstellar matter under extreme physical conditions.  Continued study will provide important insights into Type Ia supernovae, of which RS Oph may be a progenitor."
    },
    "0801/0801.3898_arXiv.txt": {
        "abstract": "{We analyze the \\ha spectral variability of the rapidly-rotating K1-dwarf LQ\\,Hya using high-resolution \\ha spectra recorded during April-May 2000. Chromospheric parameters were computed from the \\ha profile as a function of rotational phase. We find that all these parameters vary in phase, with a higher chromospheric electron density coinciding with the maximum \\ha emission. We find a clear rotational modulation of the \\ha emission that is better emphasized by subtracting a reference photospheric template built up with a spectrum of a non-active star of the same spectral type. A geometrical plage model applied to the \\ha variation curve allows us to derive the location of the active regions that come out to be close in longitude to the most pronounced photospheric spots found with Doppler imaging applied to the photospheric lines in the same spectra. Our analysis suggests that the \\ha features observed in LQ~Hya in 2000 are a scaled-up version of the solar plages as regards dimensions and/or flux contrast. No clear indication of chromospheric mass motions emerges. ",
        "introduction": "According to the solar-stellar analogy, chromospheric stellar activity can be established by the presence of emission in the  core of the Balmer \\ha line. As in the Sun, \\ha emission intensification has often been observed in surface features (plages) spatially connected with the photospheric starspots \\citep[see, e.g.,][and references therein]{Cata00,Bia06,Bia07}. Thus, the time variability of the \\ha spectral features can be used to estimate the basic properties of the emitting sources, allowing the geometry of the stellar chromosphere to be mapped.  In this paper we analyze the \\ha spectral variability of the young, rapidly-rotating ($P_{\\rm rot}\\approx 1.6$\\,d) single K2-dwarf \\object{LQ\\,Hydrae} (HD\\,82558 = Gl\\,355) using 15 high-resolution spectra recorded during April-May 2000. All \\ha spectra were acquired simultaneously with the mapping lines used for the year-2000 Doppler imaging study presented in our previous paper \\citep[][hereafter Paper~I]{Kova04}. LQ\\,Hya was first recognized as  a chromospherically active star through \\ion{Ca}{ii} H\\&K emission \\citep{bide81,hein81}. The \\ha absorption filled in by chromospheric emission was reported first by \\citet{feke86}, and LQ\\,Hya was classified as a BY\\,Dra-type spotted star. Variable \\ha emission-peak asymmetry was investigated by \\citet[][hereafter Paper~II]{Stra93}, who attributed it to the presence of chromospheric velocity fields in the \\ha forming layer probably surrounding photospheric spots.  In our spectroscopic study in Paper~I, we presented Doppler images using the \\ion{Fe}{i}-6411, \\ion{Fe}{i}-6430, and \\ion{Ca}{i}-6439 lines for both the late April data and the early May 2000 dataset. Doppler imaging was supported by simultaneous photometric measurements in Johnson-Cousins $VI$ bands. Doppler images showed spot activity uniformly at latitudes between $-20$\\degr and $+50$\\degr, sometimes with high-latitude appendages, but without a polar spot. Comparing the respective maps from two weeks apart, rapid spot evolution was detected, which was attributed to strong cross-talks between the neighboring surface features through magnetic reconnections. In Table~\\ref{chart} we give a summary of the stellar parameters as they emerged from Paper~I.  \\begin{table} \\caption{Astrophysical data for LQ\\,Hya \\citep[adopted from][]{Kova04}} \\label{chart} \\centering \\begin{tabular}{ll} \\hline\\hline \\noalign{\\smallskip} Parameter                & Value \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Classification           & K2 V \\\\ Distance (Hipparcos)     & 18.35$\\pm$0.35 pc \\\\ $(B-V)_{\\rm Hipparcos}$  & 0.933$\\pm$ 0.021 mag \\\\ $(V-I)_{\\rm Hipparcos}$  & 1.04$\\pm$ 0.02 mag \\\\ Luminosity, $L$          & 0.270$\\pm$0.009~L$_\\odot$ \\\\ $\\log g$                 & 4.0$\\pm$0.5 \\\\ $T_{\\rm eff}$            & 5070$\\pm$100 K \\\\ $v\\sin i$                & 28.0$\\pm$1.0 km/s \\\\ Inclination, $i$         & 65$^{\\circ}$$\\pm$10$^{\\circ}$ \\\\ Period, $P_{\\rm rot}$    & 1.60066$\\pm$0.00013 days \\\\ Radius, $R$              & 0.97$\\pm$0.07 $R_\\odot$ \\\\ Mass                     & $\\approx$ 0.8 $M_\\odot$ \\\\ Age                      & $\\approx$ ZAMS \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table}   The observations are again briefly presented in Sect.~\\ref{data} and the method used for the \\ha spectral study is described in Sect.~\\ref{sec:analysis}. The results are presented in Sect.~\\ref{results}. Since the \\ha observations of this paper were included in the spectroscopic data used in Paper~I to reconstruct the year-2000 Doppler images, in Sect.~\\ref{disc} we take the opportunity to compare the photospheric features with the contemporaneous H$\\alpha$-emitting regions. ",
        "conclusions": "\\label{disc}  To obtain information about the surface location of the active regions in the chromosphere of LQ~Hya, we applied a simple geometric {\\it plage} model to the rotationally modulated chromospheric emission. This method had been described and successfully applied to the \\ha modulation curves of several other active stars \\citep[e.g.,][]{Fra00, Bia06, Fra07}. On LQ~Hya, two circular bright spots (plages) are normally sufficient for reproducing the observed variations within the data errors \\citep[cf.][]{Fra00,KoBa97}. In our model the flux ratio between plages and the surrounding chromosphere ($F_{\\rm plage}/F_{\\rm chrom}$) is a free parameter. Solar values of $F_{\\rm plage}/F_{\\rm chrom}\\,\\approx\\,2$, as deduced from averaging many plages in H$\\alpha$  \\citep[e.g., ][]{Elli52,LaBo86,Ayres86}, are too low to model the high amplitudes of H$\\alpha$ $\\Delta EW$ curves in very active stars \\citep{Bia06,Fra07}. In fact, extremely large plages, covering a significant fraction of the stellar surface, would be required with such a low flux ratio, and they could not reproduce the observed modulation. In order to achieve a good fit, a flux ratio of 5, which is also typical of the brightest parts of solar plages or of flare regions \\citep[e.g., ][]{Sves76,Ziri88}, was adopted.  The solutions essentially provide the longitude of the plages and give only rough estimates of their latitude and size. The information about the latitude of surface features can be recovered from the analysis of spectral-line profiles broadened by the stellar rotation, as we did for the spots of LQ~Hya with Doppler imaging in Paper I. A similar technique cannot reach the same level of accuracy when applied to the \\ha profile whose broadening is dominated by chromospheric heating effects, which are particularly efficient in the active regions \\citep[e.g.,][]{Lanzafame00}.  We searched for the best solution by varying the longitudes, latitudes, and radii of the active regions. The radii are strongly dependent on the assumed flux contrast $F_{\\rm plage}/F_{\\rm chrom}$. Thus, only the combined information between plage dimensions and flux contrast, i.e. some kind of plage ``luminosity'' in units of the quiet chromosphere ($L_{\\rm plage}/L_{\\rm quiet}$), can be deduced as a meaningful parameter. Note also that we cannot estimate the true quiet chromospheric contribution (network), since the H$\\alpha$ minimum value, $\\Delta EW_{\\rm quiet}=0.84$\\,\\AA, could still be affected by a homogeneous distribution of smaller plages. However, such an approximation seems valid for LQ~Hya because our aim was only to compare the spatial distribution of the main surface inhomogeneities at chromospheric and photospheric levels.  We also searched for solutions with three active regions, finding a small but possibly significant increase in the goodness of the fit, with the $\\chi^2$ passing from 5.77 to 4.45 (see Fig.~\\ref{fig:rotmod}). The model with only two plages requires features at rather high latitudes (Table~\\ref{tab:param_plages}) to explain the nearly flat and long-lasting maximum, while a 3-plage model allows placing plages at lower latitudes, in better agreement with the spot locations from Doppler imaging. However, the longitudes of the two largest plages do not change very much (less than 15 degrees).  The resulting plage longitudes are in good agreement with the Doppler results in Paper\\,I; i.e., photospheric minima at 0.6 and 0.9 with the largest/coolest photospheric spots correspond to the most luminous chromospheric phases, and vice versa, the brightest photometric phase at 0.3 overlap with the lowest chromospheric \\ha emission, again supporting the paradigm that photospheric spots are physically connected with chromospheric plages (Fig.~\\ref{fig:palle}).  The synthetic \\ha EW curve for the model with three plages is plotted in the middle panel of Fig.~\\ref{fig:rotmod} as a continuous line superimposed on the data (dots), while the 2-plage solution is displayed by a dotted line. Both of them reproduce the observed rotational modulation quite well. With the same models we were able to calculate the radial velocity shift between the \\ha synthetic emission profile, resulting from the quiet chromosphere plus the visible plages, and the photosphere. It is clear in the bottom panel of Fig.~\\ref{fig:rotmod} that the amplitude of the theoretical $\\Delta V_{\\rm em}$ curves, both for a 2-plage and a 3-plage model, is very low, consistent with the values derived from the residual profiles ($\\Delta V_{\\rm em}^{\\rm res}$) and in total disagreement with the velocity shifts derived from the raw spectra. This strongly supports the spectral synthesis method we used to evaluate the chromospheric emission in the \\ha line.  We would like to outline the rotational modulation of other features observed in the \\ha profile, such as the peak separation, which is related to the chromospheric electron density (Sect.~\\ref{results}) and the asymmetry of blue/red peak intensity (Fig.\\,\\ref{fig:Ne}). The chromosphere of LQ~Hya displays a higher electron density above active regions. The blue emission peak is stronger than the red one ($\\frac{F_{\\rm R}}{F_{\\rm V}}< 1$) in the ``quiet\" chromosphere of LQ~Hya, while nearly equal peaks ($\\frac{F_{\\rm R}}{F_{\\rm V}}\\simeq 1$) tend to be observed when the most active regions are in the visible hemisphere.  Asymmetry in the peaks of an emission line with a central reversal has been frequently observed for \\ion{Ca}{ii} K line both in the Sun \\citep[e.g.,][and reference therein]{Ziri88,Ding98} and in cool stars \\citep[e.g.,][and reference therein]{Montes00}.\t% The blue asymmetry is more frequently observed in the solar chromosphere and is commonly attributed to the propagation of acoustic waves \\citep[e.g.,][]{Cram76} with an upward velocity on the order of 10 \\kms\\ in the layer in which the K$_2$ emission peaks are formed or a downward motion in the layer producing the central reversal K$_3$ \\citep[e.g.,][]{Durrant76}.  \\citet{Oranje83} has shown that the average \\ion{Ca}{ii}~K emission profile in solar plages has a completely different shape from the surrounding chromosphere, with the red peak slightly stronger than the blue one. This could reflect the different physical conditions and velocity fields in active regions compared to the quiet chromosphere.  A similar analysis cannot be made for the Sun as a star in H$\\alpha$, because this line is an absorption feature in the quiet chromosphere and only a filling in is observed in plages. However, similar changes of the \\ha line profile can be expected in the plages of any stars that are more active than the Sun. Thus, the rotational modulation of $\\frac{F_{\\rm R}}{F_{\\rm V}}$ is consistent with the presence of plages in the chromosphere of LQ~Hya.  These results suggest scaled-up versions (regarding size and brightness) of solar-type plages for the \\ha features observed in LQ~Hya in 2000, without any clear evidence of strong mass motions in its chromosphere.   \\begin{figure}[tbh] \\vspace{0.5cm} \\begin{center} \\includegraphics[width=2.5cm]{9058f4a.eps}\t% \\hspace{0.4cm} \\includegraphics[width=2.5cm]{9058f4b.eps}\t% \\hspace{0.4cm} \\includegraphics[width=2.5cm]{9058f4c.eps}\t% \\includegraphics[width=2.5cm]{9058f4d.eps}\t% \\hspace{0.4cm} \\includegraphics[width=2.5cm]{9058f4e.eps}\t% \\hspace{0.4cm} \\includegraphics[width=2.5cm]{9058f4f.eps}\t% \\includegraphics[width=2.5cm]{9058f4g.eps}\t% \\hspace{0.4cm} \\includegraphics[width=2.5cm]{9058f4h.eps}\t% \\hspace{0.4cm} \\includegraphics[width=2.5cm]{9058f4i.eps}\t% \\end{center} \\caption{Schematic representation of the surface features of LQ~Hya in the chromosphere as reconstructed by the 2-plage {\\it (top)} and by the 3-plage models {\\it (middle)} and in the photosphere {\\it (bottom)}. In these maps, the star is rotating counterclockwise. Dominant features are found at similar longitudes at both chromospheric and photospheric levels. } \\label{fig:palle} \\end{figure}   \\begin{table} \\caption{Plage parameters for LQ~Hya in April-May 2000.}\t% \\label{tab:param_plages} \\begin{center} \\begin{tabular}{crcc} \\hline \\hline \\noalign{\\smallskip} Radius  &  Lon.     &   Lat.    &  $\\frac{F_{\\rm plage}}{F_{\\rm chrom}}$ \\\\ ($\\degr$) & ($\\degr$) & ($\\degr$) &  \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{4}{c}{2 Plages Solution}\\\\ \\hline \\noalign{\\smallskip} 15.5   &  215  &  35  &  5  \\\\ 20.0   &  353  &  66  &  5  \\\\ \\hline \\noalign{\\smallskip} \\multicolumn{4}{c}{3 Plages Solution}\\\\ \\hline \\noalign{\\smallskip} 13.0   &   210  &  20  & 5  \\\\ 13.0   &     8  &  15  & 5  \\\\ 9.5    &   280  &  15  & 5  \\\\ \\hline \\hline \\end{tabular} \\end{center} \\end{table}"
    },
    "0801/0801.1065_arXiv.txt": {
        "abstract": "Strong selection effects are present in observational samples of cataclysmic variables (CVs), complicating comparisons to theoretical predictions.  The selection criteria used to define most CV samples discriminate heavily against the discovery of short-period, intrinsically faint systems.  The situation can be improved by selecting CVs for the presence of emission lines.  For this reason, we have constructed a homogeneous sample of CVs selected on the basis of H$\\alpha$ emission.  We present discovery observations of the 14 CVs and 2 additional CV candidates found in this search.  The orbital periods of 11 of the new CVs were measured; all are above 3~h.  There are two eclipsing systems in the sample, and one in which we observed a quasi-periodic modulation on a $\\sim 1\\,000\\,\\mathrm{s}$ time-scale.  We also detect the secondary star in the spectrum of one system, and measure its spectral type.  Several of the new CVs have the spectroscopic appearance of nova-like variables (NLs), and a few display what may be SW Sex star behaviour.  In a companion paper, we discuss the implications of this new sample for CV evolution. ",
        "introduction": "Cataclysmic variable stars (CVs) are semi-detached binary stars, consisting of a white dwarf primary accreting from an approximately main-sequence companion (see e.g. \\citealt{bible}).  The mass transfer is caused by orbital angular momentum loss, which also drives the secular evolution of CVs.  At long orbital periods ($P_{orb}\\ga3\\,\\mathrm{h}$), angular momentum loss through magnetic braking is thought to dominate over that resulting from gravitational radiation.  According to the `disrupted magnetic braking' model of CV evolution, magnetic braking stops when $P_{orb}$ reaches $\\simeq 3\\,\\mathrm{h}$, the secondary loses contact with its Roche lobe, and gravitational radiation is left as the only angular momentum loss mechanism (e.g. \\citealt{RobinsonBarkerCochran81}; \\citealt{RappaportVerbuntJoss83}; \\citealt{SpruitRitter83}).  The secondary only regains contact with its Roche lobe when gravitational radiation has decreased $P_{orb}$ to $\\simeq 2\\,\\mathrm{h}$.  This model explains the pronounced drop in the number of CVs at $2~\\mathrm{h} \\la P_{orb}\\la3$~h, called the period gap.  Mass loss from the secondary increases its thermal time-scale, so that, for CVs below the period gap, this eventually exceeds the mass-transfer time-scale (even though the mass loss time-scale increases as $M_2$ decreases).  When this happens, the secondary is not able to decrease its radius rapidly enough in response to mass loss, so that the orbital evolution slowly moves back through longer periods (e.g. \\citealt{Paczynski81}; \\citealt{PaczynskiSienkiewicz81}; \\citealt{RappaportJossWebbink82}).  CVs in this final phase of evolution, where $P_{orb}$ is increasing, are referred to as `period bouncers'.  The reversal in the direction of change in $P_{orb}$ causes the period minimum---a sharp cut-off in the $P_{orb}$ distribution of hydrogen-rich CVs at about 76~min.  Magnetic braking is much more efficient than gravitational radiation at removing angular momentum---the resulting mass transfer rates ($\\dot{M}$) above the gap are roughly 10 to 100 times larger than below the gap (e.g. \\citealt{Patterson84}).  This implies that long-period CVs are intrinsically bright, and also that the long-period phase of the evolution of a CV is very short-lived.  The majority of CVs should therefore be short-period, low-$\\dot{M}$ (and therefore intrinsically faint) systems.  In fact, even the time taken to evolve from the bottom of the period gap to the period minimum is short compared to the age of the Galaxy, so that most CVs should be period bouncers.  Population synthesis studies predict the relative sizes of the long-period, short-period, and period bouncer populations (roughly $1:30:70$), as well as the absolute size of the Galactic CV population (e.g. \\citealt{Kolb93}; \\citealt{HowellRappaportPolitano97}).  However, the predictions of relative sizes of sub-populations, as well as the overall space density of CVs, have been disputed on observational grounds (e.g. \\citealt{Patterson84}; \\citealt{Patterson98}; \\citealt{GansickeHagenEngels02}; \\citealt{PretoriusKniggeKolb07}; \\citealt{PretoriusNEP}).  A meaningful quantitative comparison between observations and theory is only possible if observational selection effects can be accounted for.  This requires that the criteria for inclusion in the observational sample are well-defined, and preferably homogeneous.  Furthermore, the typical intrinsic brightness of the lowest-$\\dot{M}$ CV population that can be probed by a given survey clearly depends on the survey flux limit.  If this limit is too bright, the survey has no sensitivity to any but the intrinsically brightest (and rarest) CVs.  Several complete, uniformly selected CV samples already exist, e.g. those that resulted from the Palomar Green Survey, the \\emph{ROSAT} Bright Survey, and the \\emph{ROSAT} North Ecliptic Pole Survey (\\citealt{Ringwald93}; \\citealt{SchwopeBrunnerBuckley02}; \\citealt{PretoriusNEP}).  Ongoing surveys, e.g. the Hamburg Quasar Survey, and the Sloan Digital Sky Survey (\\citealt{GansickeHagenEngels02}; \\citealt{AungwerojwitGansickeRodriguez-Gil06}; \\citealt{sdsscvs1}, 2003, 2004, 2005, 2006, 2007) are in the process of producing similarly well-defined, but much deeper, samples.  All CV samples are affected by a flux-limit\\footnote{This may not be important for the discovery of classical novae, but it is for recovering them and measuring periods.}, which already implies a bias against intrinsically faint systems.  But it is in fact very difficult to construct a sample that is reasonably large and deep, and purely flux-limited.  Of the surveys mentioned above, only the very shallow \\emph{ROSAT} Bright Survey, and the very small \\emph{ROSAT} North Ecliptic Pole Survey do not contain a blue cut, in addition to a flux limit.  The surveys that incorporate a blue selection are biased against low-$\\dot{M}$ CVs, since these systems are not only intrinsically faint, but also relatively red.  The single property through which the most CVs have been discovered is large amplitude variability.  It is not easy to quantify the selection effects acting on the CVs discovered in this way, but they certainly also favour intrinsically brighter systems, since low-$\\dot{M}$ CVs undergo less frequent outbursts.  \\cite{Gansicke05} reviews all surveys that have discovered sizable samples of CVs.  The presence of Balmer emission lines in the spectra of most CVs provides an alternative to the commonly used blue- and variability-based selection techniques.  Selecting CVs for line emission discriminates only against the discovery of the intrinsically rare CVs with very bright, optically thick discs.  In fact, there is a well known (and theoretically expected) anti-correlation between the equivalent widths (EWs) of Balmer emission lines and the luminosity of CVs (\\citealt{Patterson84}; \\citealt{WithamKniggeGansicke06}; we will take EWs of emission lines as positive throughout).  A few surveys have already exploited the promise of discovering CVs via emission lines.  CVs are selected from the Hamburg Quasar Survey based in part on H$\\beta$ emission.  One of the selection criteria for CVs from the Cal\\'{a}n-Tololo Survey is Balmer emission lines \\citep{TappertAugusteijnMaza04}.  The INT Photometric H$\\alpha$ Survey of the Northern Galactic Plane (IPHAS) is currently being used to find CVs (\\citealt{WithamKniggeGansicke06}; \\citealt{WithamKniggeAungwerojwit07}), and \\cite{RogelLuggerCohn06} describe another H$\\alpha$-based search for CVs.  H$\\alpha$ emission has also been used to identify CV candidates in globular clusters (e.g. \\citealt{BailynRubensteinSlavin96}; \\citealt{CoolGrindlayCohn95})  The AAO/UKST SuperCOSMOS H$\\alpha$ Survey (SHS; \\citealt{Parker05}) is currently the best available southern hemisphere resource for identifying H$\\alpha$ emission line point sources.  One CV has previously been discovered in this survey, partly on the basis of H$\\alpha$ emission \\citep{HowellMasonHuber02}.  We have carried out spectroscopic follow-up of sources in this survey with $R<17.5$, selected for H$\\alpha$ emission.  This has allowed us to construct a small, homogeneous sample of CVs.  We present discovery spectra of 16 new CV candidates selected from the SHS, as well as time-resolved observations that have confirmed 14 of these systems as CVs, and yielded orbital period measurements for 11 of the new CVs.  The sample will be used to derive constraints on CV evolution theory in a subsequent paper (\\citealt{HalphaII}; hereafter Paper II). ",
        "conclusions": "\\begin{tabular}{@{}llllllllllllll@{}} \\hline Object/coordinates & $B$  & $R$  &$R_2$ &$R_1$ & $I$  & $J$   & $H$   & $K_S$ & $R-\\mathrm{H}\\alpha$ & $\\mathrm{EW}(\\mathrm{H}\\alpha)/\\mathrm{\\AA}$ & $P_{orb}/\\mathrm{h}$& $d_l/\\mathrm{pc}$ & Notes \\\\ \\hline 073418.56-170626.5 & 17.8 & 16.7 & ---  & 18.3 & 17.4 & 16.23 & 15.93 &$>$14.9& 0.47 & 61(6)         & 3.18542(5)& 320 & 1,2,3\\\\ 074208.23-104932.4 & 16.5 & 15.3 & 15.4 & 16.0 & 14.3 & 14.26 & 13.59 & 13.36 & 1.21 & 63(3), 45(1)  & 5.706(3)  & 490 & 4,5\\\\ 074655.48-093430.5 & 13.8 & 14.0 & 13.6 & 13.4 & 13.3 & 13.47 & 13.29 & 13.19 & 0.46 & 27(1), 23.6(3)& 3.3984(4) & 180 & 6\\\\ 075648.83-124653.5 & 18.3 & 16.0 & 17.0 & 17.0 & 16.2 & 15.39 & 14.73 & 14.61 & 0.61 & 43(7), 47.1(6)& ---       & --- & 7\\\\ 092134.12-593906.7 & 17.3 & 16.8 & 16.7 & 16.9 & 16.7 & 16.49 & 16.20 &$>$15.8& 0.53 & 56(6), 43.7(4)& 3.041(9)  & 530 & 8\\\\ 092751.93-391052.3 & 16.4 & 16.2 & 16.8 & 16.3 & 16.1 & 15.72 & 15.54 &$>$15.7& 0.35 & 15(2), 20.0(3)& 4.1(3)    & 770 & 9\\\\ 094409.36-561711.4 & 17.5 & 16.2 & ---  & 18.4 & 16.7 & 15.72 & 15.40 & 14.96 & 0.99 & 66(4), 67(1)  & 4.506(4)  & 520 & 1,2,3,8\\\\ 102442.03-482642.5 & 17.1 & 16.4 & 17.6 & 16.4 & 16.2 & 15.83 & 15.57 & 15.41 & 0.94 &125(2), 122(1) & 3.673(6)  & 580 & 3\\\\ 103135.00-462639.0 & 17.4 & 16.3 & 17.2 & 16.9 & 16.3 & 16.04 & 15.65 & 15.52 & 0.39 & 35(4), 19.4(3)& 3.76(2)   & 630 & 9\\\\ 103959.96-470126.1 & 17.0 & 16.4 & 16.9 & 16.7 & 16.8 & 15.87 & 15.69 & 15.63 & 0.95 & 31(3), 20.2(5)& 3.785(5)  & 670 & 9\\\\ 112921.67-535543.6 & 16.3 & 15.5 & 15.9 & 16.1 & 15.3 & 15.27 & 15.17 & 15.12 & 0.33 & 23(2), 17.4(5)& 3.6851(4) & 510 & 9\\\\ 115927.06-541556.2 & 18.6 & 17.4 & 17.6 & 17.5 & 17.2 & ---   & ---   & ---   & 0.57 & 57(15)        & ---       & --- & 10\\\\ 122105.52-665048.8 & 17.8 & 17.3 & 17.5 & 16.9 & 17.4 & 16.43 &$>$16.3&$>$16.4& 0.58 & 22(3), 66.0(5)& ---       & --- & 7,9,11\\\\ 130559.50-575459.9 & 17.2 & 16.5 & 16.4 & ---  & 16.2 & 15.99 & 15.55 &$>$14.9& 0.59 & 41(4), 58.3(5)& 3.928(13) & 500 & 8\\\\ 163447.70-345423.1 & 16.8 & 16.7 & 16.7 & 17.0 & 16.3 & 15.88 & 15.42 &$>$15.7& 0.61 & 36(5), 34(3)  & ---       & --- & 3,9\\\\ 190039.83-173205.5 & 17.6 & 16.8 & 17.3 & 16.9 & 17.8 & 15.66 & 15.45 & 15.11 & 0.34 & 26(5)         & ---       & --- & 9,10\\\\ \\hline \\end{tabular} \\noindent Notes: (1) Eclipsing system. (2) Possibly an SW Sex star. (3) Large amplitude variability. (4) Secondary star spectral type M$0\\pm1$. (5) \\emph{ROSAT} source.  (6) Bowen blend detected. (7) Highly variable H$\\alpha$ line profile. (8) Strong He\\,{\\scriptsize II}\\,$\\lambda$4686 emission. (9) Spectroscopic appearance of a nova-like variable (NL). (10) Classification as a CV not certain. (11) Probable quasi-periodic oscillation (QPO) detected in one light curve. \\hfill } \\end{table*}  The SHS is a photographic survey performed with the UK Schmidt Telescope (UKST), and scanned by a digitizing machine called SuperCOSMOS.  The survey imaged the southern Galactic plane in $R$ and H$\\alpha$, down to a limiting magnitude of $\\simeq 20.5$ (see \\citealt{Parker05} for a detailed description).  $I$-band photometry from an older UKST survey is included with the H$\\alpha$ and $R$ data and was used in our target selection.  We selected objects that are clear H$\\alpha$ excess outliers in the $R-\\mathrm{H}\\alpha$ vs $R-I$ colour-colour plane for spectroscopic follow-up.  Most of the target sample was restricted to objects brighter than $R=17.0$, but we also considered objects with $17.0 \\le R < 17.5$ for one of our identification spectroscopy runs.  Two CV candidates were found in the fainter magnitude bin.  The selection criteria and sample completeness will be discussed in more detail in Paper II.  Fig.~\\ref{fig:findingcharts} gives finding charts for the new CVs and CV candidates turned up by the search, and Table~\\ref{tab:summary} lists J2000 coordinates, broad band magnitudes, H$\\alpha$ excesses and equivalent widths, orbital periods, and lower limits on distances.  We will name objects using their right ascension as `H$\\alpha$hhmmss'.  Where two values of $\\mathrm{EW}(\\mathrm{H}\\alpha)$ are given, the first is obtained from the identification spectrum, and the second from the average of higher resolution spectra of that system.  Errors in the equivalent widths were estimated using the method described by \\cite{HowarthPhillips86}.  Near-infrared magnitudes are from the Two Micron All Sky Survey (2MASS; \\citealt{2mass})\\footnote{H$\\alpha$092751 was not detected in $K_S$ by 2MASS; the 95\\% confidence lower limit on $K_S$ is given as 16.65.  This implies either very unusual near-IR colours for the source, or variability between the different IR images (which were taken on the same night).  However, the number-magnitude count of sources detected in a 5 arc min radius around the position of H$\\alpha$092751 indicates that the magnitude limit of the $K_S$-band image is about 15.7.  We will adopt this more conservative limit.}.  The optical magnitudes listed in Table~\\ref{tab:summary} are from SuperCOSMOS scans of the UKST Blue Southern and Equatorial Surveys ($B$), the UKST Red Southern and Equatorial Surveys ($R_2$), the ESO Schmidt Telescope Red Southern Survey ($\\delta < -17.5^\\circ$) or (for $-17.5^\\circ < \\delta < +2.5^\\circ$) the Palomar-I Oschin Schmidt Telescope (POSS-I) Red Southern Extension ($R_1$), and the UKST near Infrared Southern Survey ($I$).  The orbital period measurements are discussed in Section~\\ref{sec:results}, and lower limits on the distances to our sources ($d_l$) are derived in Section~\\ref{sec:distances}.  \\begin{figure*} $\\begin{array}{c@{\\hspace{0.1cm}}c@{\\hspace{0.1cm}}c@{\\hspace{0.1cm}}c} \\includegraphics[width=43mm]{fig01_01.eps} & \\includegraphics[width=43mm]{fig01_02.eps} & \\includegraphics[width=43mm]{fig01_03.eps} & \\includegraphics[width=43mm]{fig01_04.eps} \\\\ \\includegraphics[width=43mm]{fig01_05.eps} & \\includegraphics[width=43mm]{fig01_06.eps} & \\includegraphics[width=43mm]{fig01_07.eps} & \\includegraphics[width=43mm]{fig01_08.eps} \\\\ \\includegraphics[width=43mm]{fig01_09.eps} & \\includegraphics[width=43mm]{fig01_10.eps} & \\includegraphics[width=43mm]{fig01_11.eps} & \\includegraphics[width=43mm]{fig01_12.eps} \\\\ \\includegraphics[width=43mm]{fig01_13.eps} & \\includegraphics[width=43mm]{fig01_14.eps} & \\includegraphics[width=43mm]{fig01_15.eps} & \\includegraphics[width=43mm]{fig01_16.eps} \\end{array}$ \\caption {$4' \\times 4'$ finding charts of the new CVs and CV candidates, made using UKST $R$-band plates digitized by SuperCOSMOS.  The charts are labelled with the right ascension of the sources.  North is at the top and east to the left in all images. } \\label{fig:findingcharts} \\end{figure*}     We have discovered 16 CV candidates by selecting emission line objects from the SHS, and obtained additional observations for 14 of these systems, confirming their CV nature.  Orbital periods were measured for 11 of the new CVs.  All of these are long-period systems, and most have orbital periods in the range 3 to 4 h.  The periods are listed in Table~\\ref{tab:summary}.  Note that 6 of these periods were determined from aliased radial velocity curves, and may therefore only be correct to within $\\sim 10$\\% (the errors given in the table are for the strongest alias only).  The aim of this study is to construct a new CV sample with uniform selection criteria.  Although the bright flux limit implies that the sample is biased against intrinsically faint CVs, there is no explicit second selection cut that compounds this bias (but note that there is an effective $R-I$ cut, which will be discussed in Paper II).  It is also the largest H$\\alpha$-selected CV sample constructed to date.  The blue, weak-lined spectra of H$\\alpha$092751 and H$\\alpha$112921 imply that they are probably both NLs.  Given the orbital periods of our CVs, it is likely that many of them are SW Sex stars.  Specifically, we have reason to suspect that H$\\alpha$073418, H$\\alpha$094409, H$\\alpha$075648, and H$\\alpha$122105 are SW Sex stars, but confirmation is needed for all of them.  H$\\alpha$075648 and H$\\alpha$122105 are unusual in displaying large variations in line profiles, which prevented us from measuring orbital periods.  H$\\alpha$073418, H$\\alpha$094409, H$\\alpha$102442, and H$\\alpha$163447 have been observed to display large amplitude variability.  The secondary star is detected in the spectrum of H$\\alpha$074208.  We measure a spectral type of M$0\\pm1$, and estimate that the secondary contributes between 40 and 80\\% of the flux in the wavelength range from $\\simeq5\\,700$ to $\\simeq7\\,200\\,\\mathrm{\\AA}$."
    },
    "0801/0801.4530_arXiv.txt": {
        "abstract": " ",
        "introduction": "Be stars are well known to be variable on virtually all timescales, reaching from minutes to dozens of years. For the study of the latter, long term data collections as homogeneous as possible are necessary.  The professional astronomer, however, is often hampered in the study of intermediate- to long-term time scale processes as in Be stars. The reasons are the observational practices usually employed at professional observatories, which typically are not suited for observing a bright object with execution times of a few minutes only about every other week for several seasons; as well as the funding timescales, making it hard to start the collection of a long-term database that does not promise a significant number of publications within the first few years.  On the other hand, the interpretation of time-limited observations with professional resources, such as interferometers, polarimeter, or high-resolution spectrographs, in almost all cases can profit from the knowledge of the disc state in the course of the long-term evolution.  The problems in long-term data acquisition for the professional astronomer, however, open a promising field for the dedicated amateur.  Amateur spectrographs at relatively small telescopes of about 20\\,cm diameter, equipped with CCD-detectors meanwhile reach resolution powers well above 10\\,000 and and are sensitive enough to reach many of the brighter Be stars. This work describes a database worth of more than five years of observation of the Be star $\\zeta$ Tau.  The observational data will be made available online together with this communication.  \\IBVSfig{21cm}{5813_f1.ps}{All H$\\alpha$ profiles measure from late 2000 to early 2006. The vertical offset of the profiles is proportional to time and corresponds to 25 days per continuum unit. The lowermost spectra date from Nov.\\ 1, 2000 (left), Sep.\\ 9, 2002 (middle) and Aug.\\ 23, 2004 (right), respectively.} ",
        "conclusions": ""
    },
    "0801/0801.0038_arXiv.txt": {
        "abstract": "$JHK_s$ near-infrared photometry of stars in the Phoenix dwarf galaxy is presented and discussed. Combining these data with the optical photometry of Massey et al. allows a rather clean separation of field stars from Phoenix members. The discovery of a Mira variable (P = 425 days), which is almost certainly a carbon star, leads to an estimate of the distance modulus of $23.10\\pm 0.18$ that is consistent with other estimates and indicates the existence of a significant population of age $\\sim 2$ Gyr. The two carbon stars of Da Costa have $M_{bol} = -3.8$ and are consistent with belonging to a population of similar age; some other possible members of such a population are identified. A Da Costa non-carbon star is $\\Delta K_s \\sim 0.3$ mag brighter than these two carbon stars. It may be an AGB star of the dominant old population.  The nature of other stars lying close to it in the $K_s,(J-K_s)$ diagram needs studying. ",
        "introduction": "The present investigation of the Phoenix dwarf galaxy is part of a programme to study local group galaxies  using the Japanese - South African 1.4m Infrared Survey Facility (IRSF) and {\\sc sirius} three-channel camera (Nagashima et al. 1999, Nagayama et al. 2003) at SAAO Sutherland.  Phoenix is a member of the Local Group and the most distant of the Milky Way's satellite galaxies (e.g. Grebel (1999) fig. 3).  It was discovered by Schuster \\& West (1976) who originally suggested it might be a globular cluster; Canterna \\& Flower (1977) established that it was a galaxy. Though its overall properties are consistent with a classification as a dwarf spheroidal, it also contains a relatively small young component and is thus often referred to as a dIrr/dSph (e.g. Mateo 1998). It is associated with an off-centre H{\\sc i} cloud (Oosterloo, Da Costa \\& Staveley-Smith 1996; Young \\& Lo 1997; St-Germain et al. 1999). The origin of this cloud is not clear, although it may be formed from supernovae winds associated with the most recent epoch of star formation in the galaxy (Young et al. 2007). Though there have been a number of optical studies of Phoenix, this seems to be the first to describe $JHK_s$ observations. ",
        "conclusions": "By combining our own $JHK_s$ observations with the optical photometry of Massey et al. (2007) it has been possible to make a rather clean separation of Phoenix members from field stars.  A clear RGB of an old population is found together with a few highly evolved stars. A Mira variable, almost certainly a carbon star, with a period of 425 days is present in the galaxy and leads to an estimate of $23.10 \\pm 0.18$ for the distance modulus in agreement with other estimates.  The kinematics of carbon Miras in our Galaxy (Feast et al. 2006) suggest an age of $\\sim 2\\;$Gyr for this star. Since Miras are relatively short-lived objects this implies a significant population of this age.  The two Da Costa carbon stars have  $M_{K} = -6.2$ or $M_{bol} = -3.8$ (based on an estimate of the bolometric correction from the work of Frogel et al.(1980)).  These luminosities are consistent with an age $\\sim 1$ to a few Gyr (see e.g. the luminosities of carbon stars in LMC clusters (Frogel et al. 1990)).  Whilst most of the stars fainter than $K_s \\sim 17.2$ are found to be members of an old RGB population, a significant number of them are identified as probably AGB stars of intermediate age. They are likely to belong to the same population as the carbon stars. In this connection, we note that a feature in the colour-magnitude diagram of Holtzman et al. (2000) (their fig. 2), starting at $V$ or $I$ of $\\sim 24.0$, $(V-I) \\sim 0$ and sloping to higher luminosities and redder colours may be a subgiant branch of intermediate age stars. It is reasonably well fitted by a 1 Gyr isochrone (z = 0.002) from Girardi et al. (2002).  The status of the non-carbon star Da Costa C2 which is $\\sim 0.3$ mag brighter than the two carbon stars at $K_s$ is uncertain. It seems most likely to be an AGB star of an old population. Whether other stars of about the same luminosity and colour to C2 are also old AGB stars or carbon stars of an intermediate age population requires further spectroscopic work."
    },
    "0801/0801.2451_arXiv.txt": {
        "abstract": "The rich oscillation spectra determined for the two stars, $\\nu$~Eridani and 12~Lacertae, present an interesting challenge to stellar modelling. The stars are hybrid objects showing a number of modes at frequencies typical for $\\beta$ Cep stars but also one mode at frequency typical for SPB stars. We construct seismic models of these stars considering uncertainties in opacity and element distribution. We also present estimate of the interior rotation rate and address the matter of mode excitation.  We use both the OP and OPAL opacity data and find significant difference in the results. Uncertainty in these data remains a major obstacle in precise modelling of the objects and, in particular, in estimating the overshooting distance. We find evidence for significant rotation rate increase between envelope and core in the two stars.  Instability of low-frequency g-modes was found in seismic models of $\\nu$~Eri built with the OP data, but at frequencies higher than those measured in the star. No such instability was found in models of 12~Lac. We do not have yet a satisfactory explanation for low frequency modes. Some enhancement of opacity in the driving zone is required but we argue that it cannot be achieved by the iron accumulation, as it has been proposed. ",
        "introduction": "}  Transport of chemical elements is still not well understood aspect of stellar interior physics. In the case of the upper main sequence there are uncertainties regarding the extent of mixing of the nuclear reaction products beyond the convective core, known as the overshooting problem, as well as, regarding the survival of element stratification in outer layers caused by  selective radiation pressure and diffusion. Closely related is another difficult problem of the angular momentum transport because there is a role of rotation in element mixing. The upper main sequence pulsators, $\\beta$ Cephei stars, are potential source of constraints on modelling the transport processes massive stars. Recently,  Miglio et al. (2007b) and Montalb\\'an et al. (2007) studied sensitivity of the oscillation frequencies in B stars to the rotationally induced mixing. A comprehensive survey of properties of these variables was published by Stankov \\& Handler (2005).   Pulsation encountered in $\\beta$ Cep stars frequently consists in excitation of a number of modes, which differ in their probing properties. They are found most often in evolved objects, where low order nonradial modes have mixed character: acoustic-type in the envelope and gravity-type in the deep radiative interior. Frequencies of such modes are very sensitive to the extent of overshooting. Their rotational splitting is a source of information about the deep interior rotation rate. Pulsation is found both in very slow ($<10$~km~s$^{-1}$) and very rapid ($>200$~km~s$^{-1}$) rotating stars. Thus, there is a prospect for disentangling the effect of rotation in element mixing. First limits on convective overshooting and some measures of differential rotation have been already derived from data on the $\\beta$ Cep stars HD 129929 (Aerts et al. 2003), $\\nu$~Eri (Pamyatnykh et al. 2004 (PHD), and Ausselooss et al. 2004). Unfortunately, both objects are very slow rotators. To disentangle the role of rotation in element mixing, we need corresponding constraints for more rapidly rotating objects.  It seems that we understand quite well the driving effect responsible for mode excitation in $\\beta$ Cep stars. The effect arises due to the metal (mainly iron) opacity bump at temperature near 200 000 K. Yet, the seismic models of $\\nu$~Eri, which reproduce exactly the measured frequencies of the dominant modes, predict instability in a much narrower frequency range than observed. As a possible solution of this discrepancy, PHD proposed accumulation of iron in the bump zone caused by radiation pressure, following solution of the driving problem for sdB pulsators (Charpinet at al. 1996). PHD have not conducted calculations of the abundance evolution in the outer layers. Instead, they adopted an {\\it ad hoc} factor 4 iron enhancement showing that it leads to a considerable increase of the instability range but also allows to fit the frequency of the troublesome p$_2$ dipole mode. However, recent calculations of the chemical evolution made by Bourge et al. (2007) showed that the iron accumulation in the $\\beta$ Cep proceeds in a different manner than in the sdB stars. Unlike in latter objects, in photospheres of $\\beta$ Cep stars the iron abundance is significantly enhanced for as long as it is enhanced in the bump. In fact, the photospheric enhancement is always greater (P.-O. Bourge private communication). However, neither $\\nu$~Eri nor 12~Lac shows abundance anomaly in atmosphere except perhaps for the nitrogen excess, [N/O]=$0.25\\pm0.3$, in the former object (Morel et al. 2006). So there is no clear evidence for element stratification and we are facing two problems: what is the cause of element mixing in outer layers of this star, and how to explain excitation of high frequency modes ($\\nu>7$~cd$^{-1}$) and the isolated very low frequency mode at $\\nu=0.43$~cd$^{-1}$.  Unlike most of $\\beta$ Cep stars, where frequencies of detected modes are confined to narrow ranges around fundamental or first overtone of radial pulsation, modes in $\\nu$~Eri and 12~Lac are found in wide frequency ranges. Both stars are hybrid objects with simultaneously excited low-order p- and g-modes (typical for $\\beta$~Cep stars) as well high-order g-modes (typical for SPB stars). The greatest challenge is to explain how the latter modes are driven. Our work focuses on these two stars, which are best studied but not the only hybrid pulsators. Other such objects are $\\gamma$~Pegasi, $\\iota$~Her, HD~13745, HD~19374 (De Cat et al. 2007). Very recently Pigulski \\& Pojma\\'nski (2008) reported a likely discovery of five additional objects.  In the next section, we compare oscillation spectra and provide other basic observational data for $\\nu$~Eri and 12~Lac. Seismic models of these two objects are presented in Section 3, after a brief description  our new treatment of the overshooting. In the same section, we give for both objects the estimate of rotation rate gradient in the abundance varying zone outside core. Section 5 is devoted to the problem of mode excitation. Finally, in Section 6, after summarizing the results, we discuss the matter of element mixing in outer layers and what is needed for progress in B star seismology.   \\vspace{4mm} ",
        "conclusions": "We believe that our seismic models of $\\nu$~Eridani yield a good approximation to its internal structure and rotation but there is room for improvement and a need for a full explanation of mode driving. We did not fully succeed in the interpretation of the oscillation spectrum of $\\nu$ Eridani with our standard evolutionary models and our linear nonadiabatic treatment of stellar oscillations. The frequency misfit between the high frequency peak to the $\\ell=1, {\\rm p}_2$ mode is much reduced in models built with the OP opacity data allowing large overshooting but such models are much cooler ($2\\sigma$) than the mean colour of the star implies. Such models also nearly avoid the driving problem. Similar conclusions regarding models allowing large overshooting distance were reached by Ausseloos et al. (2004). As for the consequences of usage of the OP instead of OPAL opacity data for mode instability, our finding agrees with more general observation of Miglio et al. (2007a) that models using the former data predict wider frequency ranges.   An assessment of the inward rise of the angular rate in the $\\mu$-gradient zone was made for both, $\\nu$~Eri and 12~Lac, yielding the values around five for the ratio of the core to envelope rate. The surface equatorial velocities in the two stars are very different. For the former object our seismic estimate gives about 6 km~s$^{-1}$. For 12~Lac, our estimate of about 50 km~s$^{-1}$ is less certain because data on rotational splitting are much poorer but the value agrees with estimate of Desmet (private communication), based on his analysis of line profile changes. There is a large difference in rotation rate between the two stars but unfortunately the accuracy of our modelling is insufficient for addressing the question of the relation between the overshooting distance and rotation.  The surface rotation rate in $\\nu$~Eri is indeed very low and this was the reason why PHD suggested that chemical element stratification in this star may be sustained. Hovever, a closer look at the problem reveals that the effect of rotation on element distribution should not be ignored, even at the equatorial velocity of few km~s$^{-1}$. As Bourge et al. (2007) showed for their $10 M_\\odot$ model, the two-fold excess of the iron abundance in the driving zone is produced in the time scale comparable with main sequence life time. The process is slower than in less massive objects (Seaton 1999) but fast enough so that if it goes unimpeded a significant excess of iron could be created. The effect of rotation could be ignored if the time, $\\tau_{\\rm mc}$, for meridional circulation to travel from the depth where most of iron is moved up ($T\\approx4\\times10^5$K) to the bottom of the convective zone ($T\\approx2\\times10^5$K) around the iron opacity bump is longer than the star age. To estimate $\\tau_{\\rm mc}$ we use the well-known expression (see e.g. Tassoul 2000) for the speed of the meridional flow  \\begin{equation} V_{\\rm mc}={\\epsilon R\\over\\tau_{\\rm KH}}f(x)P_2(\\cos\\theta), \\end{equation} where $\\epsilon={\\Omega^2R^3/GM}$, $\\tau_{\\rm KH}={GM^2\\over RL}$ is the Kelvin-Helmholtz time, and  $$f(x)={4\\over3}\\left({\\partial\\ln\\rho\\over\\partial\\ln T}\\right)_p {x^5\\over1-\\nabla/\\nabla_{\\rm ad}}\\left(1-{\\Omega^2\\over2\\pi G\\rho}\\right).$$  The last term, known as the Gratton-\\\"Opik term, is kept though it is of higher order because it is important near the surface where density $\\rho$ is low, even if $\\epsilon<<1$ and otherwise linearization is valid. In our model of $\\nu$~Eri, the the  Gratton-\\\"Opik term is -3.4 at the depth where most of iron is supposed to be pushed up.  \\begin{equation} \\tau_{\\rm mc}={\\tau_{\\rm KH}\\over\\epsilon}{\\Delta x\\over|f|} \\end{equation}  We evaluate $\\tau_{\\rm mc}$ at the layer where $T=4\\times10^5$K. and for $\\Delta x$ use the distance to the bottom of the convective zone. For the $\\nu$~Eri model, we have $\\epsilon=1.2\\times10^{-4}$, $\\tau_{\\rm KH}=5.9\\times10^4$yr, $x_b=0.907$, $\\Delta x=0.05$, $f=-15$. With this numbers, we get from Eq.\\,(6) $\\tau_{\\rm mc}=1.14\\times10^6$y, which is much less than the age of the star, $1.75\\times10^7$y according to our seismic model. Thus, even in this slowly rotating star the effect of rotation cannot be ignored. For our best model of 12~Lac, we find $\\tau_{\\rm mc}$ by about four orders less than in $\\nu$~Eri. Yet, as we may see in Fig.\\,8, the problem with driving the low frequency mode is even greater than in the case of $\\nu$~Eri.  Meridional flow and/or turbulence developed through instability induced by rotation are expected to prevent accumulation of iron in the driving zone and in photosphere. This is a likely explanation of apparently normal chemical atmospheric composition of $\\beta$~Cephei stars, including such slow rotators as $\\nu$~Eri. This is also the reason to reject iron accumulation in the driving zone as the solution of the driving problem, whose solution must be searched in a different way.  We attach greatest hope for solution of the driving problem with improvements in stellar opacity data. Thus, we would like to encourage further effort in this field. Reliable opacity data are essential for B star seismology. Data on mode frequencies in $\\nu$~Eri and 12~Lac are abundant and accurate enough for probing rotation and structure of the $\\mu$-gradient zone above the convective core. However, to answer open questions regarding macroscopic transport of elements and angular momentum, we need accurate microscopic input data, especially on opacity. Only if we have measurements of  rotational splitting, our inference on differential rotation does not rest on very precise model of stellar structure, but such data are rare. Regarding observational work, we view as most important new spectroscopic observations of the two stars aimed at improving the value of the mean effective temperature and leading to the $\\ell$ and $m$ identifications for a greater number of modes."
    },
    "0801/0801.0454_arXiv.txt": {
        "abstract": "We examine the infrared properties of 43 high redshift (0.1\\,$<$\\,z\\,$<$\\,1.2), infrared-luminous galaxies in the Extended Groth Strip (EGS), selected by a deep 70\\,$\\mu$m survey with the Multiband Imaging Photometer on \\emph{Spitzer} (MIPS). In addition and with reference to starburst-type Spectral Energy Distributions (SEDs), we derive a set of equations for estimating the total infrared luminosity ($L_{IR}$) in the range 8--1000\\,$\\mu$m using photometry from at least one MIPS band. 42 out of 43 of our sources' optical/infrared SEDs ($\\lambda_{observed}$\\,$<$\\,160 $\\mu$m) are starburst-type, with only one object displaying a prominent power-law near-infrared continuum. For a quantitative analysis, models of radiation transfer in dusty media are fit onto the infrared photometry, revealing that the majority of galaxies are represented by high extinction, A$_v$\\,$>$\\,35 and for a large fraction ($\\sim$\\,50 per cent) the SED turns over into the Rayleigh-Jeans regime at wavelengths longward of 90\\,$\\mu$m. For comparison, we also fit semi-empirical templates based on \\emph{local} galaxy data, however, these underestimate the far-infrared SED shape by a factor of at least 2 and in extreme cases up to 10 for the majority ($\\sim$\\,70 per cent) of the sources. Further investigation of SED characteristics reveals that the mid-infrared (70/24\\,$\\mu$m) continuum slope is decoupled from various galaxy properties such as the total infrared luminosity and far-infrared peak, quantified by the L$_{160}$/L$_{70}$ ratio. In view of these results, we propose that these high-redshift galaxies have different properties to their local counterparts, in the sense that large amounts of dust cause heavy obscuration and are responsible for an additional cold emissive component, appearing as a \\emph{far-infrared excess} in their SEDs. ",
        "introduction": "The classes of galaxies known as Luminous, Ultraluminous and Hyperluminous Infrared Galaxies (LIRGs, ULIRGs and HyLIRGs), have been a major subject of focus in observational cosmology since initial, groundbreaking studies such as those of Soifer et al. (1986) and Sanders et al. (1987). These objects, discovered in abundance by the Infrared Astronomical Satellite (\\emph{IRAS}), were found to possess an array of extreme properties, including an excessive energy output, effectively manifested as a large infrared/optical ratio (Sanders et al. 1989). It is now widely accepted that the combination of a powerful UV-photon source and the presence of substantial amounts of dust are responsible for such extreme luminosities of the order of $10^{10}$--$10^{14}$$L_{\\odot}$ (e.g. Sanders $\\&$ Mirabel 1996, Genzel et al. 1998). Determining the nature of the central energy source has been the subject of numerous subsequent studies, very often supplemented by multiwavelength data (e.g. Gregorich et al. 1995, Genzel et al. 1998, Rigopoulou et al. 1999, Klaas et al. 2001, Tacconi et al. 2002, Alonso-Herrero et al. 2005). These have revealed populations dominated by powerful starburst activity, but with a non-negligible Active Galactic Nucleus (AGN) contribution found to increase with bolometric luminosity (e.g. Veilleux et al. 1997, Tran et al. 2001, Brand et al. 2006).  The processes responsible for energy production and radiation transfer in the interstellar medium of such obscured systems are primarily evaluated by studying galaxy Spectral Energy Distributions (SEDs), sometimes replacing spectroscopic diagnostics in determining the nature of the central engine.  For infrared-luminous objects, these processes are greatly influenced by the composition, temperature and distribution of interstellar dust. Quantifying the infrared energy budget is, therefore, key in determining the generic properties of a galaxy and, as it is directly coupled to the central energy source, important in estimating the integrated cosmic star formation history. Cases where photometric data is scarce, benefit from various methods of data extrapolation and interpolation: construction of infrared SEDs has been a popular approach of modelling emission from dusty systems, either directly from first principles or semi-empirically, bridging together contributions from different parts of the interstellar medium (e.g. Guiderdoni et al. 1998, Calzetti et al. 2000, Rowan-Robinson 2000). Such models have enabled the examination of sources with extreme properties and pronounced infrared luminosity, either by solely considering star formation processes or including contribution from AGN. In addition, establishing correlations between monochromatic flux densities or luminosities has proven extremely advantageous and has conveniently been used as a first order approximation of the infrared energy budget, as well as a diagnostic for the physical processes in galaxies, adopted extensively in evolution studies (e.g. Franceschini et al. 2001).  In this paper we investigate the properties of 43 infrared-luminous galaxies from one of the first deep \\emph{Spitzer} far-IR surveys. This sub-sample is part of a population of 178 sources detected at 70\\,$\\mu$m by the Multiband Imaging Photometer on \\emph{Spitzer} (MIPS) (hereafter the 70\\,$\\mu$m population), down to a limiting flux density of 4\\,mJy at 5$\\sigma$. In Symeonidis et al. 2007 (hereafter S07), we gave a brief overview of the sources' properties and derived Star Formation Rate (SFR) estimates. Here, we aim to gain qualitative insight into their physical properties, primarily with respect to dust, comparing them to similar objects in the local universe. For this purpose and to quantitatively characterise the sample, we fit our photometry with three starburst dust models and a set of empirical local galaxy templates.  The paper is set out as follows: Section 2 introduces our selection criteria for the sub-sample, including an overview of infrared colours. Section 3 analyses and compares the model SED templates, associating the quality of fits to the properties of the sources. We investigate the infrared energy budget, with reference to SED characteristics such as the continuum slope and far-IR peak, in section 4. Finally, section 5 focuses on the derivation of a set of equations to calculate the total infrared luminosity in the range 8--1000\\,$\\mu$m, with at least one MIPS band, for starburst-type sources. Our conclusions and discussion are presented in section 6. In all subsequent calculations, we have employed the following values: $H_o=71$ kms$^{-1}$Mpc$^{-1}$, $\\Omega_M=0.3$ and  $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "We have examined a population of 70\\,$\\mu$m selected sources, in the region of the Extended Groth Strip, focusing this study on the infrared properties of 43 objects with available spectroscopic redshifts (0.1\\,$<$\\,z\\,$<$\\,1.2). In the last section, we derived a set of equations in order to estimate the total IR luminosity with at least one MIPS band, applicable to objects with starburst-type SEDs. We would like to emphasise that our selection criteria did not photometrically bias the redshift sample with respect to the full sample or predispose our study to a particular type of galaxy. As this is a flux limited survey (see section 2), the higher redshift objects are also more luminous in the infrared and display lower optical flux. This implies that the sources with no available spectroscopic redshift are also in the LIRG/ULIRG regime. In addition, the MIPS colours (figure \\ref{fig:mipscolours}) and SEDs (figure \\ref{fig:seds}) unanimously show that all objects are photometrically similar in the mid and far-IR and their properties diverge only in the optical and near-IR, not relevant for the work described here.  By fitting 4 libraries of SED templates on our infrared photometry, we have examined the properties of the sources and estimated the infrared energy budget, revealing 12 per cent starbursts, 62 per cent LIRGs and 26 per cent ULIRGs. Evaluation of the fits showed that the SK07 models perform better at reproducing the far-IR SED; their formulation adopts a two-component configuration of dust heated in the immediate locality of massive young stars and cirrus-heated dust, essentially decoupling mid and far infrared emission. As a result, since the total infrared luminosity is not directly related to the position of the SED peak or the steepness of the mid-IR continuum slope, with the SK07 templates it is possible to represent objects with strong mid-IR emission and an additional cold component. CE01 and DH02 employ a scenario with a single radiating central source where temperature varies as a function of distance from the centre. Although this has worked well with local infrared galaxies, it completely underestimates our sources' 160\\,$\\mu$m flux, with a discrepancy of at least a factor of 2 and up to 10 for extreme cases, since it does not allow for types of systems with elevated far-IR emission. Our empirical local galaxy templates, also mostly underestimate the far-IR region, indicative of the fact that an average dust temperature of 30-50\\,K is too high.  Considering all above points, we propose that this deep 70\\,$\\mu$m survey has probed high-z LIRGs and ULIRGs, which are unlike local counterparts: heavy obscuration and large amounts of cold dust potentially lower than 30\\,K, appear as a \\emph{far-infrared excess} component in the SEDs, causing a significant fraction to peak at $\\lambda>$\\,90\\,$\\mu$m and steepening the continuum slope. The existence of such objects has also been put forward by other authors; e.g. Marcillac et al. (2006) have suggested that some high-z infrared-luminous objects display far-IR properties divergent from those of local galaxies because of an additional cold dust component, even if the strength of mid-IR emission is the same. This has further implications: in such systems, the continuum slope cannot be representative of various galaxy characteristics such as dust temperature, total infrared luminosity or the nature of the central energy source. Figures \\ref{fig:lumratios} and \\ref{fig:lratiopeak} demonstrate this: the mid-IR continuum slope (quantified by the $L_{70}/L_{24}$ ratio) satisfies the IRAS warm source criterion of $f_{25}/f_{60}$ $>0.17$ for sources with reduced 60\\,$\\mu$m emission and increased $\\lambda >$\\,90\\,$\\mu$m flux, as opposed to elevated 25\\,$\\mu$m emission and it is completely dissociated from the position of the SED turnover and total infrared luminosity. As this is inconsistent with the representation of far-IR emission as a single temperature black body, it is in line with our earlier suggestion of the presence of an additional cold emissive component.  It is worth noting that this 70\\,$\\mu$m sample is one of the first far-IR-selected samples reaching such low flux-density limits and additional similar surveys will be of great importance in highlighting the differences (if any) between local and high redshift infrared populations. The next infrared observatory, \\emph{Herschel}, will greatly benefit this work, extending to the sub-mm part of the spectrum and, hence, enabling a complete census of high-z infrared galaxy SEDs. Finally, we would like to mention that in order to confirm our conclusions we have planned a follow-up of the work described in this paper with submillimetre observations using the Submillimetre Common-User Bolometer Array (SCUBA 2)."
    },
    "0801/0801.0348_arXiv.txt": {
        "abstract": "Observations of two H$_2$CO ($3_{03}-2_{02}$ and $3_{21}-2_{20}$) lines and continuum emission at 1.3 mm towards Sgr B2(N) and Sgr B2(M) have been carried out with the SMA. The mosaic maps of Sgr B2(N) and Sgr B2(M) in both continuum and lines show a complex distribution of dust and molecular gas in both clumps and filaments surrounding the compact star formation cores. We have observed a decelerating outflow originated from the Sgr B2(M) core, showing that both the red-shifted and blue-shifted outflow components have a common terminal velocity. This terminal velocity is 58$\\pm$2 km s$^{-1}$. It provides an excellent method in determination of the systematic velocity of the molecular cloud. The SMA observations have also shown that a large fraction of absorption against the two continuum cores is red-shifted with respect to the systematic velocities of Sgr B2(N) and Sgr B2(M), respectively, suggesting that the majority of the dense molecular gas is flowing into the two major cores where massive stars have been formed. We have solved the radiative transfer in a multi-level system with LVG approximation. The observed H$_2$CO line intensities and their ratios can be adequately fitted with this model for the most of the gas components. However, the line intensities between the higher energy level transition H$_2$CO ($3_{21}-2_{20}$) and the lower energy level transition H$_2$CO ($3_{03}-2_{02}$) is reversed in the red-shifted outflow region of Sgr B2(M), suggesting the presence of inversion in population between the ground levels in the two K ladders (K$_{-1}$= 0 and 2). The possibility of weak maser processes for the H$_2$CO emission in Sgr B2(M) is discussed. ",
        "introduction": "The giant molecular cloud Sgr B2, located close to the Galactic center ($\\sim$44 arcmin from Sgr A*), is a well-known massive star-forming region in our Galaxy. Sgr B2 consists of an extended envelope, a hot ring and a few compact cores (e.g. Goicoechea, Rodriguze-Fernandez \\& Cernicharo 2004). The radio continuum and recombination line observations of the compact HII regions suggest that Sgr B2(N) and Sgr~B2~(M) are the two most active star forming cores in this region (Gaume \\& Claussen 1990; Gaume et al. 1995, Mehringer et al. 1993; de Pree et al. 1995, 1996, 1998). Masers, outflows and possible rotation of the two dense cores have been revealed from observations of various molecular lines at centimeter and millimeter wavelengths (Reid et al. 1988; Gaume \\& Claussen 1990; Martin-Pintado et al. 1990; Mehringer, Goss \\& Palmer 1994; Lis et al. 1993; Kuan \\& Snyder 1996; Liu \\& Snyder 1999). In addition, previous observations have shown evidence for the two hot cores to be at different evolutionary stages and to have different molecular abundances (e.g. Vogel et al. 1987; Lis et al. 1993; Miao et al. 1995; Kuan, Mehringer \\& Snyder 1996; Liu \\& Snyder 1999).  H$_2$CO pervades the interstellar medium and it has a simple chemical reaction path which has been proven to be a useful probe of physical conditions (e.g. Mangum \\& Wootten 1993). The H$_{2}$CO (1$_{10}-1_{11}$) transition at 6 cm was observed in absorption against discrete continuum sources towards Sgr B2 complex with an angular resolution of $\\sim10^{\\prime\\prime}\\times20^{\\prime\\prime}$, showing nearly the same radial velocity pattern as that of the radio recombination lines (Martin-Pintado et al. 1990; Mehringer, Palmer \\& Goss 1995). These authors suggested that the H$_{2}$CO (1$_{10}-1_{11}$) transition probably arises from the surrounding gas with a relatively low mean H$_{2}$ density of $\\sim 10^{4}$ cm$^{-3}$ (Martin-Pintado et al. 1990; Mehringer, Palmer \\& Goss 1995).  The millimeter H$_{2}$CO lines are an excellent tracer of H$_{2}$ density $> 10^{5}$ cm$^{-3}$ (e.g., Mangum \\& Wootten 1993). In addition, H$_{2}$CO is a planar asymmetric top molecule with very little asymmetry. The symmetry of the spin function of the molecule leads to two transition classes: ortho-H$_{2}$CO levels if the spin wavefunction is symmetric and para-H$_{2}$CO levels if antisymmetric. Since para-H$_{2}$CO is 1-3 times less abundant than ortho-H$_{2}$CO, observations of para-H$_{2}$CO have less opacity effect (Kahane et al 1984; Mangum \\& Wootten 1993). Hence, para-H$_{2}$CO appears to be a better probe to determine the physical conditions of the massive star formation regions.  The millimeter/submillimeter transitions of H$_{2}$CO gas require relatively high excitation temperature and high H$_{2}$ density compared to those in the centimeter wavebands. If the brightness temperature of the continuum emission is higher than the excitation temperature, the absorption against the continuum cores can be observed in millimeter and sub-millimeter wavebands with the high angular resolution of an interferometric array (such as the Submillimeter Array,\\footnote {The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics and is funded by the Smithsonian Institution and the Academia Sinica.} hereafter SMA). Taking advantage of the large bandwidth coverage of the SMA, we have observed multiple H$_{2}$CO lines towards Sgr B2 at 1.3 mm within a bandwidth of 2 GHz. Thus, with the same telescope system and calibration procedure, the uncertainties due to absolute flux density calibration among the different line transitions can be mitigated by measuring the line-intensity ratios, which are needed to determine physical conditions, such as kinetic temperature and H$_{2}$ number density, of the gas. In addition, the SMA is not sensitive to extended larger scale emission ($\\sim$ 50\\arcsec). Thus, the SMA observations are sensitive to the clumps of high density gas rather than the extended diffuse components.  In this paper, we present the results from the SMA observations of Sgr B2 at the H$_{2}$CO lines and continuum at 1.3 mm. The paper is organized as follows: \\S 2 discusses the observations and data reduction. In \\S 3 we present the data analysis and results. In \\S 4, we present the kinematics in Sgr B2(M) by a model incorporating a spherically symmetric inflow along with a decelerating outflow. In \\S 5 we model the physical properties of the H$_{2}$CO gas in Sgr B2 using the large velocity gradient (LVG) approach. In \\S 6, we discuss the important results derived from our observations and analysis. We summarize the results in \\S7. We adopt a distance of 8 kpc to Sgr B2. ",
        "conclusions": "The assessment of molecular cloud mass from molecular lines can be affected by the excitation, opacity, abundance variations and gas dynamics of the molecular lines. The optically thin submillmeter dust continuum emission has been proven to be a good tracer of molecular cloud mass (Pierce-Price et al. 2000, Gordon 1995). If we take an average grain radius of 0.1 $\\mu$m and grain density of 3 g~cm$^{-3}$ and a gas to dust ratio of 100 (Hildebrand 1983, Lis, Carlstrom \\& Keene 1991), the dusty cloud mass and column density are given by the formulae (Lis, Carlstrom \\& Keene 1991)  \\begin{equation} M_{H_{2}}=1.3\\times 10^{4}\\frac{e^{h{\\nu}/{\\kappa}T}-1}{Q({\\nu})} (\\frac {S_{\\nu}}{Jy})(\\frac {D}{kpc})^{2}(\\frac {{\\nu}}{GHz})^{-3}~M_{\\odot}, \\end{equation}  \\begin{equation} N_{H_{2}}=8.1\\times 10^{17}\\frac{e^{h{\\nu}/{\\kappa}T}-1}{Q({\\nu}){\\Omega}}(\\frac { S_{\\nu}}{Jy})(\\frac {{\\nu}}{GHz})^{-3}~ ({\\rm cm}^{-2}), \\end{equation}  \\noindent where $T$ is the mean dust temperature (K), $Q({\\nu})$ is grain emissivity at frequency ${\\nu}$, $S_{\\nu}$ is the flux density corrected for free-free emission, $\\Omega$ is the solid angle subtended by the source. Assuming $Q({\\nu})$ at 1.3 mm is 2$\\times$10$^{-5}$ and the dust temperature is 150 K for Sgr B2 ( Carlstrom \\& Vogel 1989; Lis et al. 1993; Kuan, Mehringer \\& Snyder 1996), we derived the masses, H$_{2}$ column densities and number densities.  Because the continuum at 1.3 mm contains free-free emission, we estimate the physical parameters using the flux densities of the continuum corrected for free-free emission. Assuming 3.6 cm continuum emission with FWHM beam of 3$\\rlap{.}^{\\prime\\prime}8\\times 2\\rlap{.}^{\\prime\\prime}0$ (Mehringer et al. 1993) of Sgr B2(N) (K1-3)  and Sgr B2(M) (F1-4) come from optically thin free-free emission ($S_{\\nu} \\varpropto {\\nu}^{-0.1}$), we estimate the free-free contribution of 4.7 and 8.4 Jy ($\\sim$ 9\\% and  24\\% of the total flux densities at 1.3 mm) at 1.3 mm towards the continuum cores K1-3 and F1-4 in our observations. Our estimates are consistent with the determinations of Lis et al. (1993) (6\\% and 33\\% of the total flux densities at 1.3 mm) and Martin-Pintado et al. (1990) ($\\sim$ 9\\% and 28\\% of the total flux densities at 1.3 mm ) for K1-3 and F1-4. The continuum flux densities at 1.3 mm corrected for free-free emission are 47.4 and 27.2 Jy for K1-3 and F1-4, respectively.  From the flux densities at 1.3 cm (Gaume et al. 1995), the estimated free-free emission contributions at 1.3 mm are 0.02 Jy, 0.04 Jy and 1.26 Jy for K4, Z10.24 and MW, respectively. The corresponding continuum flux densities corrected for free-free emission are 3.5, 9.1 and 11.0 Jy, respectively.  The derived H$_{2}$ masses, column densities and number densities are summarized in Table 4. The estimated H$_{2}$ masses of the Sgr B2(N) core (K1-3) and the Sgr B2(M) core (F1-4) are larger than those given by Lis et al. (1993), while the H$_{2}$ number densities are less than their results. This result is caused by the relatively larger size and higher flux densities of the continuum in our observations.  The H$_{2}$CO ($3_{03}-2_{02}$) spectra show absorption against both Sgr B2(N) and Sgr B2(M) compact cores and multiple absorbing peaks. The absorptions are dominated by red-shifted gas, suggesting that the lower transition H$_{2}$CO ($3_{03}-2_{02}$) traces the cold gas in front of the continuum cores falling into the two compact cores.  Previous observations showed multiple massive young stars in Sgr B2(N) and Sgr B2(M) (e.g., Gaume et al. 1995, de Pree, Goss \\& Gaume 1998). There are multiple absorbing peaks with different optical depths from our H$_{2}$CO spectra in the Sgr B2(N) and Sgr B2(M) cores, which appears to indicate that the gas is falling into the massive stars or massive star forming cores embedded at different depths in the molecular clouds (Mehringer, Palmer \\& Goss 1995). However, the angular resolution of our observations is inadequate for us to determine whether there are multiple regions present or whether the overall gravitation potential dominates the infalling gas. If infalling gas is in simple free-fall, the infalling velocities can be estimated by  \\begin{equation} V_{in}=\\sqrt {\\frac {2MG}{r_{in}}}=0.09\\sqrt {\\frac {M/M_{\\odot}}{r_{in}/pc}} ~ {\\rm km~s^{-1}}, \\end{equation}  \\noindent where $r_{in}$ is the infall radius, $M$ is the sum of the gas and star masses included in the $r_{in}$ and $G$ is the gravitation constant.  The H$_{2}$ masses derived from the continuum are 1.4$\\times$10$^{4}$ and 7.9$\\times$10$^{3}$ M$_{\\odot}$ for the cores of Sgr B2(N) and Sgr B2(M), respectively. The VLA observations of radio continuum at 1.3 cm (Gaume et al. 1995) showed that there are three UCHII regions (K1, K2 and K3) in the core of Sgr B2(N) and four UCHII regions (F1, F2, F3 and F4) in the core of Sgr B2(M). By use of the relationships between stellar spectral type and stellar mass (Vacca, Garmany \\& Shull 1996), a total stellar mass of 68 M$_{\\odot}$ was inferred for the massive stars in the core of Sgr B2(N). The higher resolution observations (0$\\rlap{.}^{\\prime\\prime}$05) at 7 mm (de Pree, Goss \\& Gaume 1998) resolved out F1, F2, F3 and F4 into twenty-one UCHII regions, and a stellar mass of 443 M$_{\\odot}$ was inferred corresponding to the massive stars in the core of Sgr B2(M). The mass of the massive stars in the Sgr B2(M) core is six times larger than that in the Sgr B2(N) core. Taking the major axis sizes of 0.24 and 0.29 pc as the infall radii of the two cores (K1-3 and F1-4), we inferred the infall velocities of 21 and 15 km~s$^{-1}$ for Sgr B2(N) and Sgr B2(M), respectively. Hence, based on our SMA observations, we have shown that high-density molecular gas is continuously feeding onto the active star formation cores in both Sgr B2(M) and Sgr B2(N)."
    },
    "0801/0801.1722_arXiv.txt": {
        "abstract": "{Inflationary scenarios based on simple non-minimal coupling $a \\phi^2 R$ and its generalizations are studied. Generalizing the form of non-minimal coupling to $ K(\\phi) R$ with an arbitrary function $K(\\phi)$, we show that the flat potential still is obtainable when $V(\\phi)/K^2(\\phi)$ is asymptotically constant. Very interestingly, if the ratio of the dimensionless self-coupling constant($\\lambda$) of the inflaton field and the non-minimal coupling constant ($a$) is small, $\\sqrt{\\lambda /a^2} \\sim 10^{-5}$, the cosmological observables for general monomial cases ($K \\sim a \\phi^m$, $V \\sim \\lambda \\phi^{2m}$) are in good agreement with recent observational data.} ",
        "introduction": "An early period of accelerated expansion of the universe, or inflation~\\cite{inf}, can solve many cosmological problems such as flatness problem, homogeneity problem and isotropy problem  and can provide the desired initial conditions for the subsequent hot big bang cosmology~\\cite{books}. In particle physics models, inflation occurs when one or more scalar fields, the inflaton fields, dominate the energy density of the universe with their potential being overwhelming~\\cite{Lyth:1998xn}. Under such a condition, dubbed slow-roll condition, the curvature perturbation is produced which is nearly scale invariant and is heavily constrained by the measurements of the anisotropies of the CMB and the observations of the large scale structure~\\cite{obs}. The slow-roll condition says that the inflaton potential should be very flat, i.e. the effective mass of the inflaton should be very small compared with the inflationary Hubble parameter. The biggest question is the origin of the inflaton field itself.  Recently Bezrukov and Shaposhnikov (BS) reported an interesting possibility that the standard model with an additional non-minimal coupling term of the Higgs field and the Ricci scalar ($\\sim a |\\phi|^2 R$) can give rise to inflation ~\\cite{Bezrukov:2007ep} without introducing any new scalar particle in the theory \\footnote{The models of chaotic inflation with nonzero $a$ were considered in various different contexts \\cite{Spokoiny:1984bd,Salopek:1988qh,Kaiser:1994vs,Komatsu:1999mt,Futamase:1987ua,Fakir:1990eg,Libanov:1998wg}.}. The authors showed that the ``physical Higgs potential'' in Einstein frame is indeed nearly flat at the large field value limit but it is also required from the COBE data $U/\\varepsilon=(0.027 \\mpl)^4$ that the ratio between the quartic coupling of the Higgs field ($\\lambda$) and the non-minimal coupling constant ($a$) should be small $\\sqrt{\\lambda/a^2}\\sim 10^{-5}$ \\footnote{One should note that when one is trying to identify the inflaton field as the Higgs field the self-coupling $\\lambda$ is of the order of unity. In that case the largeness of the non-minimal coupling is required as well.}. Here we would generalize the case of BS by taking more generic form of the nonminimal coupling and read out the required condition for the asymptotically flat potential. It is certainly worthwhile to consider the generalization since we could understand the underlying structure of the theory more closely.  After reviewing the suggestion by BS in the next section, we generalize the suggestion by considering generic form of the gravity-scalar coupling term in a non-minimal way ($K(\\phi)R$) and see the general condition for getting the flat potential or the slow-roll condition in the Sec.\\ref{generalization}. In the Sec.\\ref{monomial}, we work out the monomial case with functions ($K\\sim a \\phi^m$ and  $V\\sim \\lambda \\phi^{2m}$) in detail. Interestingly for any positive integer power ($m$) the scalar spectral index and the tensor-to-scalar perturbation turned out to be in good agreement with the latest cosmological observations once a combination of dimensionless self coupling constant and the coefficient of the nonminimal coupling is fixed as $\\sqrt{\\lambda_0/a_0^2}\\sim 10^{-5}$ (Details of the parameters are given in Sec.\\ref{monomial}). Summary and discussions will be followed. ",
        "conclusions": "In this paper, we examined the inflationary scenarios based on non-minimal coupling of a scalar field with the Ricci scalar ($\\sim K(\\phi)R$). Taking conformal transformation, the resultant scalar potential in the Einstein frame is shown to be flat at the large field limit if the condition in eq.\\ref{condition} is satisfied. This is one of the main result of this paper.  This class of models gets constraints from the recent cosmological observations of the spectral index, tensor-to-scalar perturbation ratio as well as the amplitude of the potential. We explicitly considered the monomial cases $K \\sim \\phi^m$ and found that this class of models are indeed good agreement with the recent observational data: $n_S \\simeq 0.964-0.975$ and $r \\simeq 0.0007- 0.008$ for any value of $m$. In fig.\\ref{rnplot}, the predicted values for $n_S$ and $r$ are depicted. We explicitly read out the condition for fitting the observed anisotropy of the CMBR by which essentially the amplitude of the potential is determined. The condition does not look natural ($\\sqrt{\\lambda/a^2}\\sim 10^{-5}$) at the first sight but we may understand this seemingly unnatural value once we embed the theory in higher dimensional space-time. Details of higher dimensional embedding of the theory and possible solution to the smallness of $\\sqrt{\\lambda/a^2}$ will be given in separate publication \\cite{large volume}."
    },
    "0801/0801.2930_arXiv.txt": {
        "abstract": "Stars at the top of the asymptotic giant branch (AGB) can exhibit maser emission from molecules like SiO, H$_2$O and OH. These masers appear in general stratified in the envelope, with the SiO masers close to the central star and the OH masers farther out in the envelope. As the star evolves to the planetary nebula (PN) phase, mass-loss stops and ionization of the envelope begins, making the masers disappear progressively. The OH masers in PNe can be present in the envelope for periods of $\\sim$1000 years but the H$_2$O masers can survive only hundreds of years. Then, H$_2$O maser emission is not expected in PNe and its detection suggests that these objects are in a very particular moment of its evolution in the transition from AGB to PNe. We discuss the unambiguous detection of H$_2$O maser emission in two planetary nebulae: K~3-35 and IRAS~17347-3139. The water-vapor masers in these PNe are tracing disk-like structures around the core and in the case of K3-35 the  masers were also found at the tip of its bipolar lobes. Kinematic modeling of the H$_2$O masers in both PNe suggest the existence of a rotating and expanding disk. Both PNe exhibit a bipolar morphology and in the particular case of K~3-35 the OH masers are highly polarized close to the core in a disk-like structure. All these observational results are consistent with the models where rotation and magnetic fields have been proposed to explain the asymmetries observed in planetary nebulae. ",
        "introduction": "Planetary nebulae (PNe) represent the final stage of evolution for intermediate mass stars (1-8 M$_\\odot$). A PN consists of an expanding gaseous shell of highly ionized gas that has been ejected by the star during the end of the AGB evolution, and a central star that will end as a white dwarf.  The study of PN is very important not only because our Sun will end as a PN but also because these objects expel heavy elements like carbon and oxygen producing an enrichment of the interstellar medium. The lifetime of a PNe is about 10$^4$ years and it is expected that there are 10$^4$-10$^5$ PNe in our galaxy, from which only $\\sim$1500 have been reported (\\cite{Kohoutek01}). PNe are characterized by exhibiting a variety of shapes (e.g. \\cite{Balick02} ). Several observational studies have revealed that most of the PNe ($\\sim$75$\\%$; \\cite{Manch04}) exhibit asymmetric morphologies that go from elliptical to very collimated outflows. One of the open questions in this field is to understand the origin of these asymmetries.  \\begin{table}\\def~{\\hphantom{0}} \\begin{center} \\caption{OH and H$_2$O maser emission in young PNe} \\label{tab:young} \\begin{tabular}{lcccc}\\hline PNe          & OH   & H$_2$O   & Morphology & References \\\\\\hline% NGC~6302    & YES  & NO        & bipolar & Payne et al. (1988)\\\\ OH~349.36-0.20 &YES&NO&?& Zijlstra et al. (1989)\\\\ IRAS 17207-2855& YES & NO  & ? & Zijlstra et al. (1989) \\\\ PK 356+2$^\\circ$ 1 & YES & NO  & ? & Zijlstra et al. (1989) \\\\ IRAS 17347-3139 & YES & YES  & bipolar & de Gregorio-Monsalvo et al. (2004) \\\\ IRAS 17371-2747 & YES & NO &  ? & Zijlstra et al. (1989) \\\\ IRAS 17375-2759 & YES & NO &  ? & Zijlstra et al. (1989) \\\\ IRAS 17375-3000 & YES & NO &  ? & Zijlstra et al. (1989) \\\\ OH0.9+1.3 & YES & NO & ? & Zijlstra et al. (1989) \\\\ IRAS 17443-2949 & YES & YES &  ? & Zijlstra et al. (1989); Su\\'arez et al. (2007)\\\\ IRAS 17580-3111 & YES & YES$^a$ &  ? &  Zijlstra et al. (1989); Su\\'arez et al. (2007)\\\\ IRAS 18061-2505 & NO  & YES$^b$  & ? & Su\\'arez et al. (2007)\\\\ IC~4997         & YES & NO       & bipolar & \\cite{Tamura89} \\\\ K~3-35          & YES & YES & bipolar & Engels et al. (1989); Miranda et al. (2001)\\\\ Vy2-2  &YES&NO&shell& \\cite{Davis79} \\\\ M1-92 & YES & NO & bipolar & \\cite{lepine74} \\\\ \\hline \\end{tabular} \\begin{tabular}{l} \\noindent \\scriptsize{$^a$ Probably not a PN; $^b$ No OH maser emission (\\cite{Suarez07}).}\\\\ \\end{tabular} \\end{center} \\end{table}  To answer this question we should review the formation process itself in the transition between the AGB and PNe stages. After completion of the H and He core burning, the star evolves to the AGB. At the top of the AGB phase, the star is surrounded by an expanding envelope of gas and dust (\\cite{Casw74} ). The study of the star in this evolutionary phase is better pursued with radio and infrared telescopes and the presence of neutral atomic and molecular gas in the envelope of the star is common (e.g \\cite{Mufson75}; \\cite{Huggins96}; \\cite{Josselin03}; \\cite{Bujarrabal06} ). For AGB stars with oxigen-rich envelopes it is possible to detect maser emission of one or more molecules, such as SiO, H$_2$O and OH that appear stratified in the envelope (e.g. \\cite{reidmoran81}; \\cite{Elitzur92} ). As the slow and massive mass loss of the late AGB phase stops, the gravity contracts the core heating it up, and the star enters its PN phase, from this moment the maser conditions will disappear in short timescales (\\cite{Lewis89} , \\cite{Gomez90}). In particular, the water molecules are expected to disappear in a timescale of decades, and only OH masers seem to persist for a considerable time ($\\sim$1000 yr; \\cite{Kwok93}). In this way H$_2$O masers are not expected to be present in a regular PN. However, two PNe (K3-35 and IRAS 17347-3139; \\cite{Miranda01}; \\cite{deGre04} ) have been found to harbor H$_2$O and OH maser emission, suggesting that these objects can be in an early stage of their evolution as PNe. Recently, two more PNe have been reported to exhibit H$_2$O maser emission and interferometric observations are required to confirm the association (\\cite{Suarez07}). In this talk we will present details about the kinematics and polarization of the maser emission in the two confirmed PNe: K3-35 and IRAS~17347-3139.  \\begin{figure} \\includegraphics[height=5.3in,width=5.3in,angle=0]{ygomez_fig1.eps} \\caption{Left: Contour images of the radio continuum observed with the VLA in K~3-35 (\\cite{Miranda01}) and IRAS~1737-3139 (\\cite{Tafoya07b}). Right: The respective water maser spots with a disk fit (\\cite{Uscanga07}). }\\label{fig:torus} \\end{figure} ",
        "conclusions": "\\label{sec:concl}  The presence of H$_2$O and OH maser emission in PNe is not common, and their detection suggest that these objects are in a very early phase of their evolution. In particular the presence of H$_2$O maser emission has been confirmed only in two PNe (K~3-35 and IRAS~17347-3139), and from a recent H$_2$O survey, made with the Robledo de Chavela Antenna, two more PNe seem to be detected. The H$_2$O maser positions in  K~3-35 and IRAS~17347-3139 show kinematics consistent with masers located in rings. From modeling, it is derived that in K~3-35 there is a rotating and expanding disk perpendicular to the bipolar outflow. There is evidence of magnetic fields in K~3-35 that could help to constrain the models for jets and bipolar morphologies in PNe."
    },
    "0801/0801.2876_arXiv.txt": {
        "abstract": "{Determining the structure of and the velocity field in prestellar cores is essential to understanding protostellar evolution.}  {We have observed the dense prestellar cores L~1544 and L~183 in the $N = 1 \\rightarrow 0$ rotational transition of CN and \\thcn \\ in order to test whether CN is depleted in the high--density nuclei of these cores.}  {We have used the IRAM~30~m telescope to observe along the major and minor axes of these cores. We compare these observations with the 1~mm dust emission, which serves as a proxy for the hydrogen column density.}{ We find that while CN\\jone\\ is optically thick, the distribution of \\thcn\\jone\\ intensity follows the dust emission well, implying that the CN abundance does not vary greatly with density.  We derive an abundance ratio of $\\rm [CN]/[\\hh]=\\dix{-9}$ in L~183 and $1-3$\\tdix{-9} in L~1544, which, in the case of L~183, is similar to previous estimates obtained by sampling lower--density regions of the core.}{We conclude that CN is not depleted towards the high--density peaks of these cores and thus behaves like the N-containing molecules \\nnhp\\ and \\nhhh. CN is, to our knowledge, the first C--containing molecule to exhibit this characteristic.} ",
        "introduction": "Knowledge of the structure and kinematics of prestellar cores is important to our understanding of protostellar evolution.  Cores which are close to the critical point at which collapse sets in are representative of the preliminary phase of evolution which subsequently leads to the formation of a protostar.  In this context, the cores with the highest central density and/or column density are the most interesting; but the fact that CO and many other tracers of the kinematics deplete on to dust grain surfaces at densities above \\dix{5}~\\ccc\\ \\citep{tafalla2002} has been an obstacle to progress in this area.  Fortunately, observations have suggested that some N--containing species, such as \\nhhh\\ and \\nnhp, remain in the gas phase at densities above \\dix{5}~\\ccc. Given that both \\nhhh\\ and \\nnhp\\ form from \\nn\\, and that the volatility of \\nn\\ is similar to that of CO \\citep{oberg2005}, it is puzzling that even these species survive. There is some evidence that they finally freeze out at densities $\\gtrsim\\dix{6}$~\\ccc\\ \\citep{bergin2002,belloche2002,pagani2007}, and, irrespective of the reason for this behaviour, it would be useful to find other tracers of the densest gas in the cores.  Radicals represent a type of ``molecule'' which tends to resist depletion; this comes about because, typically, they are both formed and destroyed by neutral--neutral substitution reactions in which both formation and destruction depend in the same manner on the abundance of an atom, such as C, N, or O. An example is NO, which is believed to be formed and destroyed by reactions with atomic nitrogen. Under these conditions, the fractional abundance of NO is independent of the degree of depletion of atomic N. In a previous article \\citep[][\\, hereafter A07]{akyilmaz2007}, we attempted to test this hypothesis by observing NO in two high density cores. We found, to our discomfiture, that the hypothesis was incorrect: NO becomes depleted in the high density cores of L~1544 and L~183 similarly to many other species.  The explanation of this observational result is not entirely clear, but it probably has to do with the fraction of oxygen which remains in the gas phase at high densities.  The motivation for our NO study arose partly from the issues of the nitrogen chemistry, referred to earlier. NO has been considered to be the main intermediary between atomic and molecular nitrogen, which are probably the main forms of elemental nitrogen in the gas phase in molecular clouds. However, the fact that NO appears to be not present in the densest gas suggested that there may be other intermediaries between N and \\nn\\ and that a possible candidate might be CN. CN can form from hydrocarbons, such as CH, in reactions with atomic N \\citep{pineau1990} and is destroyed, also by N, forming \\nn.  Thus, if elemental carbon is present in the dense gas, CN might mediate in the formation of \\nn\\ and hence in the formation of \\nhhh\\ and \\nnhp. However, the fraction of elemental carbon which is available to form hydrocarbons in the dense regions of prestellar cores is uncertain. The reason for this uncertainty is that the main reservoir of carbon in the gas phase of molecular clouds is CO, and if CO freezes out, it might be expected that there remains little carbon to form hydrocarbons and, by extension, CN.  However, from an observational point of view, one knows only that CO is at least an order of magnitude underabundant in prestellar core nuclei. There are non-thermal processes \\citep[\\eg][]{leger1985} which can maintain gas phase CO at abundances of order $\\rm [CO]/[\\hh] = 10^{-6}$ at densities of order \\dix{6}~\\ccc\\ and which might suffice to produce appreciable amounts of \\cp, a precursor of CH, following the reaction of CO with cosmic--ray produced \\hep.  The present Letter presents observations of CN and \\thcn\\ which demonstrate that CN is relatively abundant in the high density cores of L~183 and L~1544.  In section~2, we summarize the observations and, in section~3, we present the basic observational results. Section~4 contains a brief discussion of the implications of our observational results for chemical models. ",
        "conclusions": "The main conclusion of this work is that CN remains in the gas phase at densities close to $106$~\\ccc\\ in prestellar cores. This result is surprising, given that CO, which is the main repository of gas--phase carbon, is clearly depleted at such densities. It follows that CN can serve as a kinematic tracer of the high--density material. The profiles of \\nnhp\\jone\\ from A07 and of \\thcn\\ from the present observations are in generally good agreement, suggesting that both molecules trace the same density regime. It is worth noting also that a double--Gaussian fit to the \\thcn\\ spectrum in L~1544 at offset $(-20,20)$ yields one component with a velocity of $7.300\\pm0.011$~\\kms\\ and a FWHM of $0.130\\pm0.025$~\\kms. The associated upper limit on the kinetic temperature is $10.0\\pm3.8$~K, which is in good agreement with the value from the best fit temperature profile of \\cite{crapsi2007} at a distance of 30\\arcsec, where $\\tkin=10$~K. In turn, this implies that the non-thermal linewidth is at most 0.65 times the thermal linewidth, and hence that turbulence has almost completely dissipated.  CN is important also as a possible tracer of the magnetic field in dense cores, by means of the Zeeman effect \\citep{crutcher1999}.  There exist already indirect estimates of the magnetic field in L~1544 and L~183, based upon the \\cite{chandrasekhar1953bmag} method, which are of the order of 100~$\\mu$G \\citep{crutcher2004a}. Independent and direct measurements, based on the Zeeman effect, are highly desirable. In this context, we note that OH Zeeman measurements are sensitive only to the outer core, owing to the large beam size and to the depletion of oxygen towards the centre of the core.  The fact that CN is not depleted in the nuclei of L~1544 and L~183 has interesting consequences for the chemistry in these cores. One likely conclusion is that, although CO is depleted by more than an order of magnitude at densities of a few times $105$ to $106$~\\ccc, there remains sufficient CO in the gas phase to supply carbon to other, less abundant species. Quantifying this statement will require model calculations, which are under way. An expectation which is based on the model calculations is that the effective $\\rm [C]/[O]$ gas--phase abundance ratio in dense depleted regions, such as those in L~1544 and L~183, is larger than the value of $\\rm [C]/[O]=0.67$ adopted by \\cite{flower2005}.  The $\\rm [CN]/[NO]$ abundance ratio is sensitive to the amount of free oxygen (i.e. not tied up in CO) in the gas phase. Our observations of L~1544 establish a lower limit to $\\rm [CN]/[NO]$ of order 0.1 in the high density region; this suggests a $\\rm [C]/[O]$ of at least 0.8. There will be consequences for HCN and HNC, and it would be useful to determine limits on the abundances of these species in the dense cores of L~1544 and L~183. We note that our suggestion (A07) that CN acts as an intermediary in the nitrogen chemistry is supported by the present results.  \\begin{figure} \\def\\wa{0.55\\hsize} \\begin{center} \\includegraphics[width=0.8\\hsize,angle=-90]{pl-zero.eps} \\caption{\\thcn$(N = 1\\ra 0, F_2 = 2\\ra 1)$ spectrum at positions $(0,0)$ in L~1544 (top) and L~183 (bottom). The five observed HFS components are indicated.} \\label{fig:zero} \\end{center} \\end{figure}  \\begin{figure} \\def\\wa{0.43\\hsize} \\begin{center} \\includegraphics[width=\\wa,angle=-90]{cut-l1544.eps}\\\\ \\includegraphics[width=\\wa,angle=-90]{cut-l183.eps} \\caption{Cuts through L~1544 (top) and L~183 (bottom). Filled squares: \\thcn total column density (Tables~\\ref{tab:l1544} and ~\\ref{tab:l183}) in \\cc\\ (left hand scale). Open red circles: integrated intensity (in m\\kkms: scale is indicated in top left corner) of the $(J,F_1,F)=1,0,1 \\ra 0,1,2$ transition of \\nnhp\\ (see A07). Light histogram: dust emission in MJy\\,\\unit{sr}{-1} (right hand scale).} \\label{fig:cut} \\end{center} \\end{figure}"
    },
    "0801/0801.0274_arXiv.txt": {
        "abstract": "The steep source counts and negative $K$-corrections of bright submillimetre galaxies (SMGs) suggest that a significant fraction of those observed at high flux densities may be gravitationally lensed, and that the lensing objects may often lie at redshifts above 1, where clusters of galaxies are difficult to detect through other means.  In this case follow-up of bright SMGs may be used to identify dense structures along the line of sight.  Here we investigate the probability for SMGs to experience strong lensing, using the latest $N$-body simulations and observed source flux and redshift distributions.  We find that almost all high redshift sources with a flux density above 100\\,mJy will be lensed, if they are not relatively local galaxies.  We also give estimates of the fraction of sources experiencing strong lensing as a function of observed flux density.  This has implications for planning follow-up observations for bright SMGs discovered in future surveys with SCUBA-2 and other instruments.  The largest uncertainty in these calculations is the maximum allowed lensing amplification, which is dominated by the presently unknown spatial extent of SMGs. ",
        "introduction": "\\label{sec:Introduction}  Sources that are picked up in deep extragalactic millimetre (mm) and submillimetre (submm) wavelength surveys are far from being ordinary galaxies. Due to strong evolution, and also driven by the typically coarse angular resolution of current surveys at these wavelengths, mm/submm sources are rare and luminous compared with well-studied populations of optical galaxies. Submm galaxies (SMGs) are usually interpreted as an early rapidly star-forming phase, perhaps driven by a major merger, in the sequence that will ultimately become a massive elliptical galaxy \\citep[e.g.,][and references therein]{Blain02}.  However, gravitational lensing can also play a role, such that some fraction of SMGs may be more typical, lower-luminosity galaxies at high redshift ($z>1$) that are strongly amplified by lensing. This is already a well known effect for quasars, where the steepening source counts make lensing of increasing importance for the brightest objects \\citep[e.g.,][]{Kochanek04}.  In addition, the prevalence of giant optical arcs is partly explained by the high redshifts of the sources \\citep{Wu06}. There are many anecdotal examples of SMGs that are either known or strongly suspected to be lensed. Examples include sources seen through targeted rich galaxy clusters, most spectacularly in Abell 2218 \\citep{Kneib96, Swinbank03}, and MS0451-03 \\citep{Borys04}, as well as the brightest source in the central Hubble Deep Field North, usually called HDF850.1 \\citep{Dunlop04}, the brightest source in the entire GOODS-North field, referred to as GN20 (Pope et al., submitted), and other bright well-studied sources, such as SMM\\,J14011+0252 \\citep{Ivison01, Smail05}, MIPS\\,J142824.0+352619 \\citep{Borys06}, and SMM\\,J02399-0136 \\citep{Ivison98}.  It has long been recognized \\citep[e.g.,][]{Blain93} that the steep source counts and high redshifts of SMGs make the brightest ones prime lensed candidates.  Despite some efforts to model the lensing contribution to the source counts \\citep[e.g.,][]{Perrotta02, Perrotta03, Negrello07} the predictions for SMG lensing are still very uncertain. The purpose of the present paper is to improve this situation, using the latest numerical models, and to try to quantify the uncertainties. In the process of this investigation we will also be able to address a related question -- namely whether bright SMGs could be used as tracers of galaxy clusters at redshifts that are high enough to be challenging for other techniques.  In order to carry out this study we will need several ingredients, including the unlensed submm source counts and the redshift distribution. However, we start in section~\\ref{sec:ProbDensity} by considering the probability distribution for lensing amplification along random lines of sight. We then look at the observed and modelled source counts for SMGs in section~\\ref{sec:SourceCounts} and use these to calculate the source counts after lensing in section~\\ref{sec:Calculating}, also exploring the dependence on the redshift distribution of sources.  A major simplification of our approach is to assume that the {\\it shape\\/} of the source counts is independent of redshift; in order to test this we calculate in section~\\ref{sec:evolution} lensing expectations from a specific evolutionary model that is consistent with a range of current IR and sub-mm data.  Finally, we consider the uncertainties in the various components of our analysis in section~\\ref{sec:Uncertainties} and interpret our results in the context of upcoming submm surveys in section~\\ref{sec:Predictions}. ",
        "conclusions": "\\label{sec:Predictions}  We can use our study to make approximate predictions for upcoming submm surveys.  There are several relevant instruments, but we focus on the surveys which are planned with the SCUBA-2 instrument \\citep{Holland06} on the James Clerk Maxwell Telescope.  There are 2 relevant $850\\,\\mu$m programmes: the SCUBA-2 Cosmology Legacy Survey, S2CLS; and the SCUBA-2 `All Sky' Survey, SASSy \\citep{Thompson07}.  We have also carried out estimates for surveys at shorter wavelengths, of the sort which might be performed with the $450\\,\\mu$m array of SCUBA-2, or the SPIRE instrument on the {\\it Herschel} satellite (operating at $250$, $350$ and $500\\,\\mu$m).  While it is clear that there may be many examples of strong lenses in such wide surveys, the fraction of bright sources which are lensed is significantly lower than at longer wavelengths.  If one wants to find such `monsters', either as probes of line-of-sight structure or for their own intrinsic value, then one should turn to ground-based surveys in the $850\\,\\mu$m or ${\\sim}\\,1\\,$mm windows.  The S2CLS plans to map approximately $20\\,{\\rm deg}^2$ to an RMS of $0.7\\,$mJy.  From our number counts model we estimate that there will be about 96, 44 and 27 sources detected with $S\\,{>}\\,20$, 25 and $30\\,$mJy, respectively.  The total number of sources above these flux limits which have $\\mu\\,{>}\\,2$ will be about 15 (with most of them at the high flux end) using our best estimate from the Schechter function number counts. The numbers change only slightly over this flux density range, due to the probability of lensing increasing faster than the decrease in source counts. Above $35\\,$mJy the number of expected lensed sources declines steadily. The lensed fractions are considerably smaller using our `maximal' counts model, as can be seen by referring to Fig.~\\ref{fig:mmu-combo}.  There will of course always be bright sources which have negligible lensing -- hence one would like to know how bright to go before the probability of strong lensing is significant.  This can be determined using Fig.~\\ref{fig:mmu-combo}. If one is prepared to accept a 1 in 3 chance of a source being strongly lensed, then one should select sources observed with $S\\,{\\ga}\\,25\\,$mJy.  If one would like the chance to be 1 in 2, then that flux rises to about $30\\,$mJy.  What would one do with such sources in practice?  Since strong lensing is likely to come from either a galaxy cluster, or a massive galaxy, then the existence of strong lensing implies strongly clustered structure along that line of sight.  Hence for each sufficiently bright source one would use follow-up observations at other wavelengths to try to establish whether strong lensing was likely, and then to see if one could find the structure which was responsible for the lensing.  Multiple images or distorted morphologies of optical counterparts would be ways of determining that lensing was taking place -- these would naturally show up as part of the procedure for trying to determine counterparts in deep data at other wavelengths.  For SMGs where lensing was strongly suspected, one would target the area to search for the presence of structure along the line of sight, either with X-ray, Sunyaev-Zel'dovich or `red cluster sequence' observations.  The advantage of this approach is that the high redshift of the SMG sources means that it should be feasible to find cluster (or proto-cluster) lenses at higher redshifts than are easy to achieve with most other techniques.  Of course, the selection effects for clusters found in this way may be complicated to quantify.  Nevertheless, building up samples of $z\\,{>}\\,1$ clusters is sufficiently important for understanding structure formation (as well as constraining dark energy, etc.), that it is worth using every available method.  SASSy is designed to make a shallow $850\\,\\mu$m map over approximately $4{,}000\\,{\\rm deg}^2$, or one tenth of the sky, with the possibility of extension to a larger area later.  The RMS of the maps is planned to be around $30\\,$mJy, so that a robust $5\\sigma$ catalogue will have a limit around $150\\,$mJy.  It may also be possible to reduce this to nearer to $100\\,$mJy using targetted repeat observations for peaks in the maps.  The procedure for identifying strongly lensed sources in SASSy may be a little different than for S2CLS, and because of the brighter flux densities, the level of uncertainty in predictions of the number of lensed sources will be considerably higher.  These predictions depend strongly on the counts model, the normalization of the high amplification tail of the lensing PDF, as well as the amplification cut-off imposed by the finite source size for SMGs.  Since there are still huge uncertainties in all of these factors, the expectations for SASSy cover a wide range of possibilities.  Using the optimistic limit of $100\\,$mJy for SASSy, the unlensed model counts give approximately 1200 sources for the SASSy catalogue. Many of these will be in the `Euclidean counts' regime, and hence should be relatively easy to eliminate as lensing candidates.  These will often be already well-known galaxies, with others typically being in the {\\sl IRAS\\/}, {\\sl Akari\\/} or radio catalogues.  Colours (e.g.,~$850\\,\\mu$m to radio) can be used to distinguish objects which are likely to be at higher redshift and hence have a higher likelihood of being lensed; such methods will also be necessary to eliminate Galactic clouds.  Our best lensing estimate yields a total of 1000 sources in SASSy with lensing amplification $\\mu\\,{>}\\,2$, with most of them being much more strongly lensed than this limit.  We expect that these extremely lensed sources will be fairly easy to distinguish from relatively nearby, intrinsically bright galaxies, and hence the chances of there being structure along the line of sight to such candidates will be very high.  Follow-up of these SMGs at other wavelengths will also be easy, since they should be at least an order of magnitude brighter than the typical SCUBA sources which have been followed up in the past."
    },
    "0801/0801.2271_arXiv.txt": {
        "abstract": "The post-starburst region B in M82 and its massive star cluster component have been the focus of multiple studies, with reports that there is a large population of coeval clusters of age $\\sim$1~Gyr, which were created with a Gaussian initial mass distribution. This is in disagreement with other studies of young star clusters, which invariably find a featureless power-law mass distribution. Here, we present Gemini-North optical spectra of seven star clusters in M82-B and show that their ages are all between 10 and 300~Myr (a factor of 3-100 younger than previous photometric results) and that their extinctions range between near-zero and~4~mag~($A_V$). Using new {\\it HST ACS-HRC U}-band observations we age date an additional $\\sim$30 clusters whose ages/extinctions agree well with those determined from spectroscopy. Completeness tests show that the reported `turn-over' in the luminosity/mass distributions is most likely an artefact, due to the resolved nature of the clusters. We also show that the radial velocities of the clusters are inconsistent with them belonging to a bound region. ",
        "introduction": "% Region B, in the local starburst galaxy M82, has been the object of debate in recent years, due to claims that it hosts a 1~Gyr-old, independent starburst region \\citep{RdG01}. Furthermore, a turn-over was reported for the cluster mass distribution. This would make it unique among young cluster populations, which are generally accepted (and theoretically expected) to have power-law distributions.  As part of a larger project\\footnote{We refer the reader to \\citet[][multi-band photometry]{m82phot07} and \\citet[][spectroscopy]{isk08a} for full details of source selection, data acquisition, reduction and analysis, and full spectroscopic and photometric results.} to study the cluster population of the entire galaxy, we have utilised multi-band ({\\it UBVI}) {\\it HST-ACS} photometry of 35 Young Massive Clusters (YMCs) in the region and additionally, Gemini-North multi-object spectroscopy (GMOS-N) of seven of them, from which we constrain extinctions, ages and radial velocities. We present these results in sections \\S~2, 3 and 4 and a summary and conclusions in \\S~5. ",
        "conclusions": "% We have performed a detailed study of the stellar cluster population of M82 region~B, using spectroscopy and multi-band photometry, and conclude the following:  \\begin{itemize} \\item The clusters are considerably younger than previously reported \\citep{RdG03a} , with ages in the range 10-300~Myr, peaking at 150~Myr. This is in agreement with the 220~Myr timescale for the last encounter with M81, the event that triggered the starburst.  \\item We find significantly higher extinctions than previous studies, ranging up to $A_V\\sim2.5$~mag. However, the region has a lower overall extinction compared to the rest of the disk, hence allowing a view deep into the body of the galaxy.  \\item The radial velocities of the clusters follow the galactic rotation curve, indicating that region B cannot be kinematically distinct from the rest of the disk.  \\item Consideration of the detection limit for resolved clusters shows that the reported turnover in the mass/luminosity distribution of the clusters is caused by incompleteness effects (discussed in Figure.~\\ref{plot:det-limit}).  \\end{itemize}  Overall, our findings contradict previous claims that present region B as having formed $\\sim1$~Gyr ago, during an off-set, independent starburst episode. Our data strongly suggest that region B simply represents a line of sight into the galaxy, allowing a clear view of the cluster population of the disk, which undoubtedly formed during a galaxy-wide starburst.  The results presented here form part of a large spectroscopic study of the star cluster population of M82, which will be presented in \\citet[][spectroscopy of 61 clusters]{isk08b}."
    },
    "0801/0801.1272_arXiv.txt": {
        "abstract": "{Phase self-calibration (or {\\em selfcal}) is an algorithm often used in the calibration of interferometric observations in astronomy. Although a powerful tool, this algorithm presents strong limitations when applied to data with a low signal-to-noise ratio. We analyze the artifacts that the phase selfcal algorithm produces when applied to extremely noisy data. We show how the phase selfcal may generate a spurious source in the sky from a distribution of completely random visibilities. This spurious source is indistinguishable from a real one. We numerically and analytically compute the relationship between the maximum spurious flux density generated by selfcal from noise and the particulars of the interferometric observations. Finally, we present two simple tests that can be applied to interferometric data for checking whether a source detection is real or whether the source is an artifact of the phase self-calibration algorithm.} ",
        "introduction": "Phase self-calibration (or {\\em selfcal}) is an algorithm often used in the calibration of radio astronomical data. It was introduced by Readhead \\& Wilkinson (\\cite{Readhead1978}) and Cotton (\\cite{Cotton1979}), and it has been essential for the success of Very Long Baseline Interferometry (VLBI) imaging. Also, the antenna-based calibrations obtained from the {\\em Global Fringe Fitting} algorithm (Schwab \\& Cotton \\cite{Schwab1983}) are equivalent to a phase self-calibration. The phase selfcal will also be an algorithm widely used with future interferometric instruments, such as the Atacama Large Millimeter Array (ALMA) or the Square Kilometre Array (SKA), now under construction or planned. Optical interferometric observations (like those in the Very Large Telescope Interferometry, VLTI) will also eventually benefit from some form of selfcal, although closure phases and amplitudes are measured in optical interferometry in a very different way than in radio. Thus, the statistical analysis presented here may need some substantial changes to rigorously describe the probability of false detections by optical interferometers.  Given that part of the interferometric observations obtained from all those instruments may come from very faint sources, it is important to take into account the undesired and uncontrollable effects that the instrumentation and/or the calibration and analysis algorithms applied to the data could introduce in the interferometric observations. A deep study of all our analysis tools and their effects on noisy data is essential for discerning the reliability of detections of very faint sources. Some discoveries made by pushing the interferometric instruments to their sensitivity limits could turn out to be the result of artifacts produced by the analysis tools.  The main limitations of the phase self-calibration algorithm have been analyzed in many publications (e.g., Linfield \\cite{Linfield1986}, Wilkinson et al. \\cite{Wilkinson1988}). It is well known that an unwise use of selfcal can lead to imperfect images, even to the generation of spurious source components, elimination of real components, and deformation of the structure of extended sources. In this paper, we focus on the effects that phase self-calibration produces when applied to pure noise. We show that selfcal can generate a spurious source from pure noise, with a relatively high flux density compared to the rms of the visibility amplitudes. We analytically and numerically study how the recoverable flux density of such a spurious source depends on the details of the observations (the sensitivity of the interferometer, the number of antennas, and the averaging time of the selfcal solutions). Finally, we study the effects of phase self-calibration applied to the visibilities resulting from observations of real faint sources, instead of pure noise. We present two simple tests that can be applied to real data in order to check whether a given faint source is real or not, and apply these tests to real data, corresponding to VLBI observations of the radio supernova SN\\,2004et (Mart\\'i-Vidal et al. \\cite{MartiVidal2007}). ",
        "conclusions": "We have analyzed the consequence of the phase-self-calibration algorithm when it is applied to extremely noisy data. We have studied how this algorithm and the statistical fluctuations of the visibility phases can create a spurious source from pure noise. The flux density of the spurious source can be a considerable fraction of the rms of the visibility amplitudes. The application of other other antenna-based calibration algorithms (like the Global Fringe Fitting) to noisy data can have similar consequences to those of selfcal if the SNR cutoff of the gain solutions is set to small values.  We have considered numerical and analytic studies to show how the flux density of a spurious source created by selfcal depends on the number of antennas, the sensitivity of the array, and the averaging time of the selfcal solutions. We have also presented two simple tests that can be applied to real data in order to check if the detection of a faint source could be the result of the application of an antenna-based calibration algorithm to noisy data. These tests basically relate the averaging time of the selfcal solutions and the characteristics of the closure phase distribution to the flux density of a compact source possibly present in the data. To show a practical case, we have applied these tests to a set of real VLBI observations of supernova SN\\,2004et and found good agreement between the flux density recovered by CLEAN from the (phase-referenced) visibilities of this supernova (Mart\\'i-Vidal et al. \\cite{MartiVidal2007}) and the flux density estimate provided by our reliability tests."
    },
    "0801/0801.4661_arXiv.txt": {
        "abstract": "We derive fundamental parameters of the embedded cluster DBSB\\,48 in the southern nebula Hoffleit\\,18 and the very young open cluster Trumpler\\,14, by means of deep JH\\ks\\ infrared photometry. We build colour-magnitude and colour-colour diagrams to derive reddening and age, based on main sequence and pre-main sequence distributions. Radial stellar density profiles are used to study cluster structure and guide photometric diagram extractions. Field-star decontamination is applied to uncover the intrinsic cluster sequences in the diagrams. Ages are inferred from K-excess fractions. A prominent pre-main-sequence population is present in DBSB\\,48, and the K-excess fraction $f_K=55\\pm6\\%$ gives an age of $1.1\\pm0.5$\\,Myr. A mean reddening of $A_{K_s}=0.9\\pm0.03$ was found, corresponding to $A_V=8.2\\pm0.3$. The cluster CMD is consistent with the far kinematic distance of 5\\,kpc for Hoffleit\\,18. For Trumpler\\,14 we derived similar parameters as in previous studies in the optical, in particular an age of $1.7\\pm0.7$\\,Myr. The fraction of stars with infrared excess in Trumpler\\,14 is $f_K=28\\pm4\\%$. Despite the young ages, both clusters are described by a King profile with core radii $\\rc=0.46\\pm0.05$\\,pc and $\\rc=0.35\\pm0.04$\\,pc, respectively for DBSB\\,48 and Trumpler\\,14. Such cores are smaller than those of typical open clusters. Small cores are probably related to the cluster formation and/or parent molecular cloud fragmentation. In DBSB\\,48, the magnitude extent of the upper main sequence is $\\Delta\\,\\ks\\approx2$\\,mag, while in Trumpler\\,14 it is $\\Delta\\,\\ks\\approx5$\\,mag, consistent with the estimated ages. ",
        "introduction": "\\label{Intro}  Infrared clusters represent a new class of objects, virtually undetectable before the 90's (e.g. \\citealt{Dehar97}; \\citealt{Hodapp94}). Until recently, the number of known infrared clusters and stellar groups amounted to 276, as shown in the compilation by \\citet{BiDuBa03}. A systematic survey by \\citet{DuBiSoBa03} and \\citet{BiDuSoBa03} in directions of nebulae using 2MASS\\footnote{The Two Micron All Sky Survey, All Sky data release \\citep{Skru97}, available at {\\em http://www.ipac.caltech.edu/2mass/releases/allsky/}} revealed 346 new infrared embedded clusters and candidates. Thus, the systematic study of embedded clusters is fundamental to understand their nature, to probe the physical conditions associated to the early stages of star clusters, and to derive their fundamental parameters. Just to mention a few efforts in that direction, Serpens \\citep{OT02}, NGC\\,1333 \\citep{Warin96}, NGC\\,3576 \\citep{Persi94}, AFGL\\,5142 \\citep{CNSS93}, and S\\,106 \\citep{HodRay91}.   Embedded clusters may provide the clues  to better understand the formation and evolution processes of star clusters and their interaction with the parent molecular cloud. Colour-colour diagrams (2-CDs) allow identification of Pre-Main-Sequence (PMS) stars and can be used to distinguish them from Main Sequence (MS) and field stars (\\citealt{CNSS93}; \\citealt{LL03}). Together with colour-magnitude diagrams (CMDs) they can be used to derive reddening values, reddening distribution, distance from the Sun and age. K-band infrared excesses originate mostly in dust envelopes and/or discs of PMS stars, and indicate their evolutionary stage up to $\\sim10$\\,Myr (\\citealt{Greaves05}; \\citealt{MB05}), or more \\citep{BoBiOrBa06}. K, and particularly L, excesses are sensitive to the presence of protoplanetary discs (\\citealt{HLL01a}; \\citealt{OJL04}). Thus, these indicators allow dating very young star clusters (e.g. \\citealt{LAL96}; \\citealt{SoBi03}). In this context, it is important to increase the number of embedded clusters studied in detail in order to establish star-formation age spreads and constrain survival time-scales of dust envelopes and circumstellar discs.  Deriving locations of embedded clusters in the Galaxy by means of photometric and spectroscopic methods is useful, since these estimators might provide distances to be compared with the available kinematic ones of the nebulae (\\citealt{GG76}; \\citealt{BFS82}). Embedded clusters are typically observed up to $\\sim4$\\,kpc from the Sun \\citep{LL03}. In turn, these studies can be used to build the Galactic structure around the Sun, in particular to trace spiral arms. Improvement in distance determinations for clusters in all disc directions would contribute to the derivation of more reliable parameters of the rotation curve of the Galaxy. The rotation curve and spiral structure has been studied by \\citet{Rus03}.  In this paper fundamental cluster parameters are derived using infrared photometry of the embedded cluster DBSB\\,48 and the young open cluster Trumpler\\,14.  Hoffleit\\,18 is a southern nebula discovered by \\citet{Hof53}. The cluster embedded in this nebula is Dutra, Bica, Soares, Barbuy\\,48 - DBSB\\,48 - \\citep{DuBiSoBa03}, located at J2000 $\\alpha=10^h31^m29^s$, $\\delta=-58^\\circ02\\arcmin01\\arcsec$ ($\\ell=285.26^{\\circ}$, $b=-0.05^{\\circ}$). The estimated diameter of the cluster is 1.5\\arcmin.  The nebula Hoffleit\\,18 was also identified as a radio H\\,II region designated G285.3+0.0 or G285.253-0.05. The derived radio velocity is $-2\\rm\\,km\\,s^{-1}$, implying in this direction a near distance of 0.3 kpc and a far distance of 5.0 kpc. Given that the object is faint, the far distance should be the correct one, placing it in the Sagittarius-Carina arm \\citep{CaHa87}. In the present paper we check the consistency of the CMD loci with the kinematic distance estimate.  Trumpler\\,14 is usually classified as a young open cluster (\\citealt{VBF96}, and references therein). Besides, it is embedded in the nebula NGC\\,3372 ($\\eta\\,\\rm Car\\ Nebula$) and its optical populous nature and environment properties make it an ideal object to be compared to a bona-fide embedded cluster such as DBSB\\,48. Trumpler\\,14 contains about 13 O stars, and is a relatively massive cluster, with a mass estimated to be around $2000\\,\\ms$ \\citep{VBF96}.  In Sect.~\\ref{ObsDR} the observations are described. In Sect.~\\ref{DBSB48} the structure, CMD and 2-CDs of DBSB\\,48 are discussed, field-star decontamination is applied to the CMDs, and fundamental parameters are derived. Trumpler\\,14 is dealt with similarly in Sect.~\\ref{Tr14}. Discussions and concluding remarks are in Sect.~\\ref{Conclu}. ",
        "conclusions": "\\label{Conclu}  Before reaching the zero-age main sequence, stars are surrounded by optically thick material consisting of an infalling envelope and accretion disc that gradually disperse along the pre-main sequence phase. Because of disc-depleting processes such as irradiation by the central star, viscous accretion and mass loss due to outflow, the median lifetime of optically thick inner accretion discs may be as short as $\\rm2 - 3\\,Myr$, with the final stages of disc accretion lasting as long as $\\rm\\sim10^7\\,yr$ (\\citealt{H05}). In addition, stars in young open clusters appear to form along some period of time (e.g. \\citealt{Sagar95}, and references therein). It is in this context that the analysis of stellar density structure and stellar-mass distribution in young star clusters is important.  In the present paper we studied DBSB\\,48, the cluster embedded in the H\\,II region Hoffleit\\,18, by means of JH\\ks\\ photometry, radial stellar density profiles and fraction of K-excess emission stars. Its properties were compared to those of the young open cluster Trumpler\\,14. Besides an important population of PMS stars of different ages, field-star decontamination shows that the MS has developed only the upper mass range, as expected from PMS contraction time-scales. The relatively short MS extent ($\\Delta\\,\\ks=2$) of DBSB\\,48 reflects its younger age with respect to Trumpler\\,14 ($\\Delta\\,\\ks=5$). The upper MS extent appears to be an age indicator for well populated embedded clusters.  The loci of PMS stars in DBSB\\,48 are described by tracks with ages in the range $0.3-4$\\,Myr set in the CMD with reddening $A_V=8.2\\pm0.3$ and absolute distance modulus $(m-M)_{\\circ}=13.48\\pm0.3$. With a fraction of K-excess emission of $f_K=55\\pm6\\%$ the age of DBSB\\,48 results $1.1\\pm0.5$\\,Myr. Its radial density profile is well represented by a King profile with a core radius $\\rc=0.46\\pm0.05$\\,pc.  For Trumpler\\,14 we derived  $\\rc=0.35\\pm0.04$\\,pc. PMS stars are described by $0.1-2$\\,Myr tracks, consistent with the $1.7\\pm0.7$\\,Myr age implied by $f_K=28\\pm4\\%$.  The PMS age spread suggests that star formation in both clusters did not occur as an instantaneous event, instead it lasted a time equivalent to about the cluster ages.  The core radii of DBSB\\,48 and Trumpler\\,14  are similar to that of the embedded open cluster NGC\\,6611 and significantly smaller than those of classical open clusters (e.g. \\citealt{BoBi05}; \\citealt{OldOCs}). This suggests that core formation is a process partly primordial, probably associated with parent molecular cloud fragmentation, and partly related to subsequent internal dynamical evolution."
    },
    "0801/0801.4382_arXiv.txt": {
        "abstract": "We are conducting a search for supermassive black holes (SMBHs) with masses below $\\sim\\! 10^7\\,\\msun$ by looking for signs of extremely low-level nuclear activity in nearby galaxies that are not known to be AGNs. Our survey has the following characteristics: (a) X-ray selection using the \\cxo\\ X-ray Observatory, since x-rays are a ubiquitous feature of AGNs; (b) Emphasis on late-type spiral and dwarf galaxies, as the galaxies most likely to have low-mass SMBHs; (c) Use of multiwavelength data to verify the source is an AGN; and (d) Use of the highest angular resolution available for observations in x-rays and other bands, to separate nuclear from off-nuclear sources and to minimize contamination by host galaxy light. Here we show the feasibility of this technique to find AGNs by applying it to six nearby, face-on spiral galaxies (NGC 3169, NGC 3184, NGC 4102, NGC 4647, NGC 4713, NGC 5457) for which data already exist in the \\cxo\\ archive. All six show nuclear x-ray sources. The data as they exist at present are ambiguous regarding the nature of the nuclear x-ray sources in NGC 4713 and NGC 4647. We conclude, in accord with previous studies, that NGC 3169 and NGC 4102 are almost certainly AGNs. Most interestingly, a strong argument can be made that NGC 3184 and NGC 5457, both of type Scd, host AGNs. ",
        "introduction": "The past decade has seen extraordinary improvement in our understanding of supermassive black holes (SMBHs) --- their growth and evolution, and their links with their host galaxies. We now realize that galaxies hosting SMBHs at their centers are the rule rather than the exception. The observed correlations of the masses $\\mbh$ of the SMBHs with properties of their host galaxies, for example with the bulge stellar velocity dispersion \\citep{fm00,gea00}, show that there is a close link between the formation and evolution of galaxies and of the SMBHs they host. Furthermore, a comparison of the SMBH mass function required to explain the observed luminosity function of active galactic nuclei (AGNs) with estimates of the local SMBH mass function \\citep[e.g.][]{mea04,sea04} shows that not only must SMBHs be very common in massive galaxies but that most, if not all, of these black holes are relics of AGNs active in previous epochs.  Thus knowledge of the local SMBH mass function enables us to put constraints on theories of galaxy and SMBH formation and growth, and AGN lifetimes.  Estimates of the local SMBH mass function are anchored by resolved stellar or gas dynamical measurements of the masses of $\\sim 30$ SMBHs \\citep[see, e.g., the review by][]{ff05}, and otherwise based on the distribution of host galaxy properties (luminosity of the bulge or bulge stellar velocity dispersion $\\sigma$) and known scaling relationships between the mass of the SMBH and these properties (most prominently $\\mbh - \\sigma$ and $\\mbh - M_\\mathrm{bulge}$ ; but see \\citealt{gh07b} for a more direct estimate ). Most of the measured SMBH masses are $\\sim 10^8 \\msun$ or greater, however, as the sphere of influence of a less massive SMBH is extremely hard to resolve even at moderate distances, even with \\hst. For example, the sphere of influence of a $10^6 \\msun$ SMBH at 15 Mpc is $\\sim 30$ milliarcseconds (mas).  As a result, while different estimates of the mass function agree for $10^8 \\msun \\lesssim \\mbh \\lesssim 10^9 \\msun$ , the low-mass end ($\\mbh \\lesssim 10^6 \\msun$) often has discrepancies \\citep[see, e.g., Fig.~7 of ][ for a comparison of different authors' estimates of the mass function]{gdea07}. A second source of uncertainty at the low-mass end is the fact that it is unknown how the scaling relationships extrapolate to very late-type spiral galaxies, which have little or no bulge component, and to very low mass galaxies (dE and dSph).  Yet SMBHs \\emph{do} exist in very late-type spirals, e.g.\\ NGC 4395, a spiral galaxy of type Sdm, with $\\mbh \\sim 3\\times 10^5 \\msun$ \\citep{pea05}, and in very low mass galaxies, e.g.\\ POX 52, a dwarf galaxy, with $\\mbh \\sim 3\\times 10^5 \\msun$ \\citep{bea04}. Questions that naturally arise at this point are: Do the scaling relationships break down at low masses? What determines the mass of an SMBH: the mass of the bulge or the mass of the dark matter halo? Is there a lower bound to the local SMBH mass function? A well defined sample of low-mass SMBHs is needed to answer these questions.  Since we cannot detect low-mass SMBHs by their dynamical signature, looking for them by signs of their accretion activity may be the only viable way of detecting them. This of course limits detection to the subset of low-mass black holes that are active, but this fraction can be expected to be large, for the following reason. We now understand that the ``quasar era'' is a function of luminosity, with the space density of the most luminous quasars peaking at high redshift and that of lower luminosity quasars peaking at progressively lower redshifts \\citep{fea03,hms05}. This is often referred to as the ``downsizing'' of AGN activity with cosmic epoch.  Moreover, at least some models of black hole growth \\citep{mea04,m04} require anti-hierarchical growth. That is, higher mass SMBHs attain most of their mass at high redshift while lower mass SMBHs grow at progressively lower redshifts. The trends of AGN downsizing and anti-hierarchical growth, extended logically to the smallest mass SMBHs, imply that these objects were active in recent times and may still be accreting at the present epoch.  The Eddington luminosity of a $10^5 \\msun$ SMBH is only $\\sim \\! 10^{43}$ \\es; low-mass SMBHs, even if accreting at high rates, will be low-luminosity AGNs (LLAGNs). The converse is not true; that is, not all LLAGNs have a low-mass SMBH. In addition to a low SMBH mass, the low luminosity of an LLAGN may be caused by a low rate or radiatively inefficient mode of accretion to an SMBH of any mass \\citep[present a study of massive SMBHs accreting at very low rates]{sfea06,sgea06}. Finally, obscuration may further lower the observed luminosity. Low-ionization nuclear emission-line region (LINER) nuclei have been studied at multiple wavelengths \\citep[e.g.][]{escm02,ssd04,dsgs05,fea06} to identify LLAGNs among them, and these studies have demonstrated that these AGNs are not necessarily the same type of object with the same physical characteristics. Since we are interested specifically in the low-mass end of the SMBH mass function, it is necessary to identify among the AGNs those that can be expected to have the smallest black holes. One approach is to use mass estimators based on the luminosity of the AGN and the width of the broad component of emission lines in its spectrum. This is the technique that was used by \\citet{gh07a} in constructing their sample of low-mass SMBHs. A second approach, and the one we use, is to look for AGN activity in galaxies (late-type spirals, dwarf galaxies) where the known host galaxy-SMBH scaling relationships predict the lowest-mass SMBHs would reside. This approach has been used by \\citet{svea07,sea08} to find candidate low-mass SMBHs. However, the scaling relationships are only statistical and cannot be used to estimate the mass of a particular SMBH; thus the second approach requires an independent estimate of the SMBH mass.  A search for active low-mass SMBHs requires confirmation of the presence of an AGN in each candidate nucleus, and measurement of the mass of the SMBHs in the confirmed AGNs. The method used by \\citet{gh07a} has the virtue of effectively combining all of the above into a single step. However, as the luminosity of an AGN decreases, the optical spectrum of the galaxy nucleus becomes more and more dominated by host galaxy light, and the signature of the AGN becomes difficult to detect. Even when the optical spectrum shows no clear evidence of an AGN, however, such evidence may still be present in other wavelengths, such as x-ray and radio \\citep[e.g.][]{fea04}, and infrared \\citep[e.g.][]{dea06,svea07,sea08} \\citep[See][ for a review of nuclear activity in nearby galaxies]{h08}. For the lowest-luminosity AGNs, therefore, it is possible that a system based on optical spectra would not classify the nuclei as AGNs at all. We choose to use x-ray selection to identify candidate AGNs for the following reasons: First, x-rays can penetrate obscuring material which may be hiding the line emitting regions. Second, there are fewer sources of x-rays in a galaxy than there are of optical and UV emission and so dilution of the AGN signature by host galaxy light is less of a problem. Where the luminosity of the AGN is low to begin with, even a modest amount of absorption may result in the signal being below the optical background imposed by the galaxy. Third, even if, as expected in some theories \\citep{elb95,n00,nmm03,l03} AGNs that have luminosities or accretion rates below a cut-off value do not have broad-line regions, they should still be detectable in x-rays. X-ray observations have in fact detected AGNs in what were thought to be ``normal'' galaxies \\citep[e.g.][]{mea02}. The disadvantage, as discussed in \\S\\ref{sec:disc}, is that x-ray observations by themselves cannot always distinguish between AGNs and other x-ray sources, such as x-ray binaries (XRBs) and ultraluminous x-ray sources (ULXs). Multi-wavelength data are needed to determine the type of source.   As the first step towards assembling a sample of low-mass SMBHs, we are conducting an x-ray survey of nearby (within 20 Mpc), \\emph{quiescent} spiral galaxies using the \\cxo\\ X-ray Observatory to look for nuclear x-ray sources, with an emphasis on late-type spirals. The high angular resolution of \\cxo\\ is necessary to ensure that any detected source is really at the center and is not an off-nuclear source. The survey is sensitive to an (unobscured) SMBH of mass $\\mbh = 10^4 \\msun$ radiating at $\\sim\\! 2\\times 10^{-3}\\, \\LEdd$ out to the survey limit. We will present details of the survey sample selection and the \\cxo\\ observations in a future paper. Here, we present the methods used and the feasibility of detecting AGNs using these methods by applying them to six galaxies that meet selection criteria similar to those used for the survey sample, and for which x-ray data already exist in the \\cxo\\ archive.  The paper is organized as follows: \\S\\ref{sec:sampsel} describes the criteria used to select the six galaxies presented here; \\S\\ref{sec:datan} describes the observations and data analysis, with individual targets discussed in \\S\\S\\ref{sec:n3169}--\\ref{sec:n5457}; the results are discussed in \\S\\ref{sec:disc}. ",
        "conclusions": "}  The motivation for this paper was to evaluate the feasibility of detecting low-mass SMBHs in late-type spiral galaxies, which may still be accreting at the current epoch and if so should be detectable in x-rays. The six galaxies studied in this paper are not all late-type, but span the range Sa--Sd. NGC 3169 and NGC 4102 were regarded as low luminosity AGNs. None of the remaining four nuclei, however, was known to have an accreting SMBH, of any mass. NGC 3169 and NGC 4102 are of type Sa and Sb, which have massive bulges and are expected to have massive SMBHs. For galaxies of types Scd and Sd, on the other hand, the lack of a luminous AGN could mean either that there is no SMBH or that the mass of the SMBH is low. As such, the observations studied here examine both aspects of a search for accreting low-mass SMBH: First, are these objects really detectable, given that the accretion rate is expected to be low? Second, are any sources detected in the very latest type spiral galaxies that have small or no bulges?  We first note that of the six galaxies presented here, all six show nuclear x-ray sources.  This implies that it is a very common occurrence. In a survey of late-type spiral galaxies such as our ongoing \\cxo\\ survey, therefore, the predominant concern is not going to be detection efficiency, but rather identification of the AGNs among the detected sources. Given that the sample presented in this paper consists of only six galaxies, we do not draw statistical conclusions here of the prevalence of very low-luminosity AGNs in nearby galaxies. But we note that, as shown in \\S\\S 3.1--3.6, of the six nuclear x-ray sources, NGC 3169 is almost certainly an AGN, and NGC 4102, NGC 3184, and NGC 5457 , have very strong, though not conclusive, arguments in favor of their being AGNs. The two remaining galaxies, NGC 4713 and NGC 4647, are ambiguous but AGNs are not ruled out. We discuss below the issues such surveys will face when attempting to identify the nature of the detected sources. The diagnostic tools traditionally used to distinguish AGNs from non-AGNs \\citep[e.g.\\ optical line ratios,][]{bpt81,vo87} were developed in the course of studying luminous AGN. Dilution of the AGN emission by host galaxy light was not a serious problem and observations with low spatial resolution (several arcseconds) sufficed. In the study of AGN that are either intrinsically less luminous or are heavily obscured, however, host galaxy light becomes increasingly problematic, and surveys relying on optical spectra \\citep[e.g.][]{hfs95,gh04} require careful subtraction of the starlight using galactic spectral templates. In the weakest AGNs, however, signs of AGN emission may not be detected by the usual diagnostics.  This problem will persist until optical observations with angular resolution of 1--10 mas become possible so that host galaxy light can be effectively excluded, though it can be mitigated by using regions of the spectrum where host galaxy emission is negligible, for example very high energy x-rays (tens to hundreds of keV).  The detection of a compact radio source unresolved at milliarcsecond-scales, especially if the source is accompanied by jets, would also unambiguously identify the nucleus as an AGN. This has been the motivation for radio surveys like that of \\citet{nea02}. Some AGNs obscured in the optical and UV may be detectable using infrared emission line strengths and ratios \\citep[e.g.][]{dea06,svea07,sea08}.  While the methods listed above allow the unambiguous identification of AGNs, observations often do not have the angular resolution or sensitivity to distinguish AGN and non-AGN flux.  Other sources of radiation in the vicinity of the AGN are, for example, plasma photoionized by the AGN itself, or a nuclear star cluster. The targeted AGNs have very low luminosity and thus even moderate amounts of obscuration may cause a significant decrement in the observed flux. In most cases, therefore, the AGN contribution should not be expected to dominate the total observed flux. Consequently, identifying these AGNs requires a different approach than what can be used in the case of the more luminous AGNs (Seyferts and QSOs).  AGNs can be identified using x-ray observations (e.g.\\ with \\cxo\\ and \\textit{XMM-Newton}) solely, but only if they are point sources whose inferred luminosities are greater than $\\sim\\! 10^{41}$ \\es. Below that value, AGNs can become indistinguishable from ULXs and XRBs in x-rays. ULXs can have luminosities of a few times $10^{40}$ \\es\\ \\citep[e.g.][]{sea07} and have x-ray spectra that look similar to AGN spectra. It has been suggested \\citep{srw06} that ULX spectra show a break at $\\sim\\! 5$ keV. AGNs are not known to show this break. In high quality spectra with a large number of counts it may be possible to exploit this difference to separate ULXs and AGNs.  XRBs have power law spectra with $\\Gamma \\sim 2$, can show an Fe K$\\alpha$ emission line, and emit hard x-rays, all characteristics of AGNs as well. Additionally, even with \\cxo's angular resolution, the physical space probed ranges from the central $\\sim\\!$ 10 to 100 pc of the galaxy. The existence of one or more XRBs within that region would be unsurprising.  However, the fact that the inferred luminosities are as high as $10^{38}$ \\es\\ severely constrains the expected number of XRBs. For example, for NGC 5457 , one of the four galaxies presented here that are not confirmed AGNs, \\citet{pea01} provide a log N-log S relation as well as the surface density of x-ray point sources as a function of radius (their Figs.~3 and 4). Approximately 12.5\\% of the sources have luminosities exceeding $10^{37}$ \\es. The surface density of sources in the innermost $0.5\\arcmin$ is $\\sim 4.75$ arcmin$^{-2}$. Therefore we may expect $\\sim\\!0.6$ sources arcmin$^{-2}$ above the luminosity cut-off in the central $0.5\\arcmin$, or $\\sim 4\\times 10^{-3}$ such sources within the \\cxo\\ source circle of radius $2.3\\arcsec$ that has been used here.  NGC 5457 is of type Scd, and can be taken to be representative of the other three galaxies, NGC 4647, NGC 3184, and NGC 4713 which are types Sc, Scd, and Sd, respectively. Thus invoking XRBs and ULXs alone to explain the x-ray observations leads to physically implausible conditions, such as requiring that most, or all, quiescent, non-starburst spiral galaxies have a ULX or $\\sim \\! 100$ XRBs in the central $0.5\\arcsec-2\\arcsec$. Nevertheless, x-ray observations by themselves will usually be inadequate to distinguish AGNs from non-AGNs in any particular instance.  Information from other wavebands is thus crucial for the identification process. For the six galaxies in this paper multi-wavelength photometry is summarized in Tables~\\ref{tab:lum} and \\ref{tab:n31spitz}. But here too, traditional methods of identifying AGNs using flux ratios such as $\\alpha_{OX}$, $\\alpha_{KX}$, and $f_X/f_R$, must be used with caution. The low luminosities of the AGNs mean that the emission in the two bands being compared may not be from the same object. For instance, for an obscured AGN surrounded by a nuclear star cluster, the observed x-ray flux may be from the AGN but the optical flux may be dominated by the cluster.  The existence of an AGN must instead be inferred by consistency and plausibility checks considering as much of the spectral energy distribution as possible, and the goal is the rejection of the hypothesis that all of the observed properties can be explained without requiring the presence of an AGN. NGC 3184 provides a good example where different modes of observing, imaging and spectroscopy, in two wavebands, x-ray and IR, together make a compelling argument for the presence of an AGN where each observation individually is ambiguous.  Nuclear star clusters are fairly common in spiral galaxies \\citep[e.g.][]{bea02,wea05,salb08}.  A cluster poses two main problems. First, it makes the presence of XRBs more likely. Second, if the cluster is young and contains many O and B stars, it may overwhelm AGN emission in the UV in addition to the optical \\citep[e.g. NGC 4303,][]{cgml02}. It may be possible in some cases, as for NGC 1042 \\citep{swea08} and NGC 4102 \\citep{gvv99}, to attempt a separation of the cluster and AGN components in the optical spectrum. In addition, stellar spectra, even for late-type stars, fall faster towards the infrared than the power-laws typical of AGNs.  The mid-infrared colors of AGNs, therefore, tend to be redder than those of stellar populations and this color difference can be used to infer the presence of an AGN \\citep{skea05}.  In addition to the spectral energy distribution, source variability can be a discriminant, as AGNs are known to vary in all wavelengths, whereas a star cluster, say, would not. Conversely, if the variation is periodic it would rule out an AGN and argue for an XRB instead.  A survey of the type discussed here finds AGNs and strong AGN candidates, but does not measure the masses of the SMBHs in those AGNs. Measurement of the mass of one of these SMBHs will be a difficult endeavor, since the sphere of influence of the black hole cannot be resolved with current technology and resources. None of the six objects studied in this paper shows broad optical emission lines whose widths could be used to estimate the BH mass, and this is likely to be typical behavior. Spectropolarimetry may uncover broad lines in polarized light in some of the AGNs. Estimates of the SMBH mass may be made by using known scaling relationships, with the caveat that the correlations are all based on SMBHs two or more orders of magnitude more massive than the ones expected to be found by the survey. The least indirect method is an application of the observed correlation between the x-ray power law slope and Eddington ratio \\citep{wmp04}. The Eddington ratio and the luminosity in turn provide an estimate of the mass of the SMBH. Other relationships that can be used are: (a) the $\\mbh$--$\\sigma$ relation \\citep{fm00,gea00}, but this method becomes more and more uncertain as the bulge itself becomes ill-defined in the latest-type spirals; (b) the $\\mbh$--$L_\\mathrm{bulge}$ relationship \\citep{kr95,md01,mh03}, which has more scatter and also faces the problem of the definition of the bulge; (c) the $\\mbh$--$v_\\mathrm{circ}$ relation \\citep{f02,bea03}, which has the advantage that it does not require the presence of a well-defined bulge; (d) the $\\mbh$--$C$ relation \\citep{gea01}, where $C$ is the concentration of light. There is also a reported relationship between black hole mass and core radio power \\citep{sea03,mwea04}, but this relationship is not as well established as the others, and is based on observations of elliptical galaxies, so its applicability to the spiral galaxies here is uncertain. Overall, there is unlikely to be one standard method of measurement that can be applied to these galaxies, but one or more of the above methods may provide useful estimates of or limits to the masses of the SMBHs in these AGNs. Mass measurements independent of the scaling relationships are possible if an object turns out to have broad emission lines, like NGC 4395, in which case line widths or reverberation mapping may be used, or if it has maser emission, like NGC 4258, in which case gas dynamics can be used. Mass measurement in other cases will have to await mas-scale angular resolution in bands other than radio to resolve the sphere of influence of these black holes.  Of the six galaxies here, NGC 3169, NGC 4102, and NGC 5457 have measurements of either the stellar velocity dispersion or the luminosity of the bulge, thus allowing an estimate of their SMBH masses (assuming here that the source in M 101 is an AGN). The scatter in the $\\mbh$--$\\sigma$ and $\\mbh$--$L_\\mathrm{bulge}$ relations, together with the uncertainty in the observed flux and the bolometric correction, result in uncertainties of about an order of magnitude in the inferred Eddington ratio, but all three objects have $L/\\LEdd \\sim 10^{-4}$. This is in the range seen in low-luminosity AGNs \\citep[e.g. $L\\sim 10^{-5}\\LEdd$ for M 81;][]{yea07}, and much higher than the $L\\sim 10^{-9}\\LEdd$ seen in truly quiescent SMBHs. These observations indicate that there is indeed a population of accreting SMBHs in nearby spiral galaxies that do not show optical signs of activity but can be uncovered by looking for their x-ray emission. Such a population will answer the question of whether the bulge is the dominant component that determines the existence, and mass, of a nuclear SMBH. The discoveries of AGNs in the Sd galaxies NGC 4395 \\citep{hfsp97} and NGC 3621 \\citep{svea07}, and the strong evidence for AGNs in the Scd galaxies NGC 3184 and NGC 5457 suggest it is not, at least as far as existence is concerned. Among the SMBHs discovered in the latest-type spirals (with small or no bulges) and in the lowest-mass galaxies should be a population of SMBHs with masses less than $10^6 \\msun$, enabling a systematic study of the low-mass end of the local supermassive black hole mass function.    \\begin{deluxetable}{llcccrcccr} \\tablecaption{Targets and observation parameters\\label{tab:obs}} \\tablewidth{0pt} \\tabletypesize{\\small} \\tablecolumns{10} \\tablehead{ \\colhead{Target} & \\colhead{Morph.} & \\colhead{Nucleus} & \\multicolumn{2}{c}{Coordinates (J2000)} & \\colhead{Dist.\\tablenotemark{b}} & \\colhead{Scale} & \\colhead{Obs. Date} & \\colhead{ObsID} & \\colhead{Exp.} \\\\ \\colhead{} & \\colhead {Type} & \\colhead{Type\\tablenotemark{a}} & \\colhead{RA} & \\colhead{Dec} & \\colhead{(Mpc)} & \\colhead{(pc/\\arcsec)} & \\colhead{} & \\colhead{} & \\colhead{(ks)} } \\startdata NGC 3169 & Sa & L2 & 10 14 15.0 & +03 27 58 & 19.7 & 96 &  2001 May 2 & 1614 & 2.0\\\\ NGC 3184 & Scd & \\ion{H}{2} & 10 18 17.0 & +41 25 28 & 8.7 & 42 & 2000 Jan 8 & 804 & 39.8\\\\ & & & & & & & 2000 Feb 3 & 1520 & 21.3 \\\\ NGC 4102 & Sb & \\ion{H}{2} & 12 06 23.1 & +52 42 39 & 17.0 & 82 & 2003 Apr 30 & 4014 & 4.9\\\\ NGC 4647 & Sc & \\ion{H}{2} & 12 43 32.3 & +11 34 55 & 16.8 & 81 & 2000 Apr 20 & 785 & 36.9\\\\ NGC 4713 & Sd & T2 & 12 49 57.9 & +05 18 41 & 17.9 & 87 & 2003 Jan 28 & 4019 & 4.9\\\\ NGC 5457\\tablenotemark{c} & Scd & \\ion{H}{2} & 14 03 12.6 & +54 20 57 & 7.2 & 35 & \\nodata\\tablenotemark{c} & \\nodata\\tablenotemark{c} & 750\\\\ \\enddata \\tablenotetext{a}{From \\citet{hfs97-3}. L2: Type 2 LINER, T2: Type 2 Transition object.} \\tablenotetext{b}{From \\citet{t88} except from \\citet{ssea98} for NGC 5457.} \\tablenotetext{c}{NGC 5457 (M 101) was observed multiple times. The total observation time analyzed here is given in this table. The individual observations are listed in Table~\\ref{tab:m101}} \\end{deluxetable}     \\begin{deluxetable}{lrrrrrrcc} \\tablewidth{0pt} \\tablecaption{X-ray Measurements\\label{tab:det}} \\tablecolumns{9} \\tablehead{ \\colhead{Target} & \\multicolumn{6}{c}{Counts} & \\colhead{Bkg/Src} & \\colhead{HR\\tablenotemark{a}}\\\\ \\colhead{} & \\multicolumn{2}{c}{Broad} & \\multicolumn{2}{c}{Hard} & \\multicolumn{2}{c}{Soft} & \\colhead{Area Ratio} & \\colhead{}\\\\ \\colhead{} & \\colhead{Src} & \\colhead{Bkg} & \\colhead{Src} & \\colhead{Bkg} & \\colhead{Src} & \\colhead{Bkg} & \\colhead{} & \\colhead{}} \\startdata NGC 3169 & 159 & 23 & 148 & 5 & 11 & 18 & 21.4 & \\strt $+0.86^{+0.05}_{-0.03}$\\\\ NGC 3184 N & 36 & 95 & 4 & 33 & 32 & 62 & 92 & \\strt $-0.75^{+0.08}_{-0.13}$\\\\ \\phm{NGC 3184} S & 117 & 95 & 13 & 33 & 104 & 62 & 60 & \\strt $-0.77^{+0.05}_{-0.07}$\\\\ NGC 4102 all & 354 & 48 & 78 & 8 & 276 & 40 & 6.8 & \\strt $-0.55^{+0.04}_{-0.05}$\\\\ \\phm{NGC 4102} core & 171 & 48 & 68 & 8 & 103 & 40 & 52.5 & \\strt $-0.20^{+0.08}_{-0.07}$\\\\ \\phm{NGC 4102} ext & 115 & 48 & 6 & 8 & 109 & 40 & 26.1 & \\strt $-0.88^{+0.03}_{-0.06}$\\\\ NGC 4647 & 15 & 38 & 1 & 15 & 14 & 23 & 10.1 & \\strt $-0.80^{+0.04}_{-0.20}$\\\\ NGC 4713 & 10 & 4 & 1 & 0 & 9 & 4 & 21.4 & \\strt $-0.69^{+0.09}_{-0.25}$\\\\ NGC 5457 & 314 & 256 & 23 & 45 & 291 & 211 & 20.8 & $-0.86\\pm 0.03$ \\enddata \\tablenotetext{a}{Hardness ratio, HR = (H$-$S)/(H+S), where H and S are the counts in the hard and soft bands respectively, calculated using the method described in \\citet{pea06}. The tool used is available at \\url{http://hea-www.harvard.edu/AstroStat/BEHR/}.} \\end{deluxetable}    \\begin{deluxetable}{rlcrcc|rcc} \\rotate \\tablewidth{0pt} \\tablecaption{NGC 5457 X-ray Measurements\\label{tab:m101}} \\tablecolumns{9} \\tablehead{ \\colhead{ObsID} & \\colhead{Obs Date} & \\colhead{Exp Time} & \\multicolumn{3}{c}{North source} & \\multicolumn{3}{c}{South source}\\\\ \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{Counts} & \\colhead{Rate} & \\colhead{HR} & \\colhead{Counts} & \\colhead{Rate} & \\colhead{HR}\\\\ \\colhead{} & \\colhead{} & \\colhead{(ks)} & \\colhead{} & \\colhead{(ks$^{-1}$)} & \\colhead{} & \\colhead{} & \\colhead{(ks$^{-1}$)} & \\colhead{}} \\startdata 934 & 2000 Mar 26 & $97.8$ & 262 & $2.68$ & $-0.81^{+0.03}_{-0.04}$ & 264 & $2.70$ & \\strt $-0.70 \\pm 0.04$\\\\ 4731\\tablenotemark{a} & 2004 Jan 19 & $56.2$ & 92 & $1.64$ & $-0.64^{+0.07}_{-0.09}$ & 136 & $2.41$ & \\strt $-0.58^{+0.06}_{-0.07}$\\\\ 5300 & 2004 Mar 7 & $52.1$ & 36 & $0.69$ & $-0.75^{+0.08}_{-0.13}$ & 117 & $2.24$& \\strt $-0.60^{+0.06}_{-0.08}$\\\\ 5309 & 2004 Mar 14 & $70.8$ & 38 & $0.54$ & $-0.62^{+0.10}_{-0.14}$ & 158 & $2.24$ & \\strt $-0.64^{+0.05}_{-0.07}$\\\\ 4732 & 2004 Mar 19 & $69.8$ & 56 & $0.80$ & $-0.77^{+0.07}_{-0.09}$ & 151 & $2.17$ & \\strt $-0.58^{+0.06}_{-0.07}$\\\\ 5322 & 2004 May 3 & $64.7$ & 54 & $0.83$ & $-0.79^{+0.06}_{-0.10}$ & 144 & $2.22$ & \\strt $-0.69 \\pm 0.06$\\\\ 5323 & 2004 May 9 & $42.4$ & 42 & $0.99$ & $-0.64^{+0.09}_{-0.14}$ & 120 & $2.83$ &\\strt $-0.54^{+0.06}_{-0.08}$\\\\ 5338\\tablenotemark{a} & 2004 Jul 6 & $28.6$ & 61 & $2.14$ & $-0.56^{+0.09}_{-0.12}$ & 41 & $1.45$ & \\strt $-0.50^{+0.12}_{-0.14}$\\\\ 5339\\tablenotemark{b} & 2004 Jul 7 & $14.0$ & 18 & $1.26$ & $-0.60^{+0.13}_{-0.22}$ & 27 & $1.90$ & \\strt $-0.79^{+0.06}_{-0.15}$\\\\ 5340 & 2004 Jul 9 & $54.4$ & 81 & $1.48$ & $-0.72^{+0.06}_{-0.08}$ & 110 & $2.02$ & \\strt $-0.74^{+0.04}_{-0.07}$\\\\ 4734 & 2004 Jul 11 & $35.5$ & 31 & $0.87$ & $-0.53^{+0.13}_{-0.17}$ & 63 & $1.77$ & \\strt $-0.61^{+0.09}_{-0.11}$\\\\ 6114 & 2004 Sep 5 & $38.2$ & 91 & $2.38$ & $-0.60^{+0.08}_{-0.09}$ & 64 & $1.68$ & \\strt $-0.64^{+0.08}_{-0.11}$\\\\ 6115 & 2004 Sep 8 & $35.4$ & 96 & $2.72$ & $-0.78^{+0.05}_{-0.07}$ & 63 & $1.79$ & \\strt $-0.70^{+0.07}_{-0.10}$\\\\ 4735 & 2004 Sep 12 & $28.8$ & 70 & $2.43$ & $-0.70^{+0.07}_{-0.10}$ & 49 & $1.71$ & \\strt $-0.77^{+0.06}_{-0.11}$\\\\ 4736\\tablenotemark{a}\\tablenotemark{b} & 2004 Nov 1 & $77.4$ & 290 & $3.75$ & $-0.78^{+0.03}_{-0.04}$ & 122 & $1.58$ & \\strt $-0.67^{+0.06}_{-0.07}$\\\\ 6152\\tablenotemark{a} & 2004 Nov 7 & $26.7$ & 133 & $4.99$ & $-0.78^{+0.05}_{-0.06}$ & 51 & $1.92$ & \\strt $-0.63^{+0.09}_{-0.12}$\\\\\\enddata \\tablenotetext{a}{Source extended, or possibly image smeared.} \\tablenotetext{b}{Possible residual contamination from background flare.} \\end{deluxetable}    \\begin{deluxetable}{lccccccccc} \\rotate \\tablewidth{0pt} \\tablecaption{X-Ray Spectral Fits\\label{tab:spec}} \\tablecolumns{10} \\tablehead{ \\colhead{Target} & \\colhead{Galactic} & \\colhead{Model\\tablenotemark{a}} & \\multicolumn{6}{c}{Spectral Fit Parameters} & \\colhead{$\\chi_\\nu^2$(dof)}\\\\ \\colhead{} & \\colhead{N$_{\\mathrm H}$} & \\colhead{} & \\colhead{kT} & \\colhead{N$_{\\mathrm H}$} & \\colhead{$\\Gamma$} & \\colhead{Refl.\\tablenotemark{b}} & \\colhead{Line} & \\colhead{EW} & \\colhead{}\\\\ \\colhead{} & \\colhead{($10^{20}\\, \\mathrm{cm}^{-2}$)} & \\colhead{} & \\colhead{(keV)} & \\colhead{($10^{21}\\, \\mathrm{cm}^{-2}$)} & \\colhead{} & \\colhead{} & \\colhead{(keV)} & \\colhead{(keV)} & \\colhead{} } \\startdata NGC\\,3169 & 2.86 & \\texttt{ab(pl)} & \\nodata & \\strt $99^{+46}_{-36}$ & $2.0^{+1.2}_{-1.1}$ & \\nodata & \\nodata & \\nodata & $1.1 (26)$\\tablenotemark{c}\\\\ \\phm{NGC\\,3169} & & \\texttt{ab(pl)} & \\nodata & \\strt $116^{+90}_{-52}$ & $2.6^{+2.1}_{-1.5}$ & \\nodata & \\nodata & \\nodata & $0.3 (22)$\\\\   NGC\\,4102 core & 1.79 & \\texttt{ab(pl)} & \\nodata & \\strt $0^{+1.4}$  & $2.0\\pm 0.5$ & \\nodata & \\nodata & \\nodata & $7.4(27)$\\tablenotemark{c}\\\\  & & \\texttt{ab(rn)} & \\nodata & \\strt $0^{+0.4}$ & $1.8\\pm 0.4$ & $108^{+149}_{-61}$ & \\nodata & \\nodata & $2.3 (26)$\\tablenotemark{c} \\\\  & & \\texttt{ab(rn+g)} & \\nodata & \\strt $0^{+1.9}$ & $2.3^{+0.6}_{-0.5}$ & $129$\\tablenotemark{d} & $6.4$ & $3.3$ & $0.6 (25)$\\\\  & & \\texttt{ab(rn+g)} & \\nodata & \\strt $0^{+0.6}$ & $2.2^{+0.6}_{-0.5}$ & $129^{+283}_{-85}$ & $6.4$ & $2.5$ & $1.6 (25)$\\tablenotemark{c} \\\\  \\phm{NGC\\,4102} ext & & \\texttt{ab(br)} & \\strt $1.2^{+8.6}_{-0.8}$ & $1.2^{+2.8}_{-1.2}$ & \\nodata & \\nodata & \\nodata & \\nodata & $0.5(18)$\\\\  NGC\\,5457 sep.\\ low & 1.15 & \\texttt{ab(me+pl)} & \\strt $0.3^{+0.2}_{-0.1}$ & $0^{+1.8}$ & $1.7\\pm 0.5$ & \\nodata & \\nodata & \\nodata & $0.5(28)$\\\\  \\phm{NGC\\,5457} sep.\\ high & & \\texttt{ab(pl)} & \\nodata & \\strt $1.5^{+1.1}_{-1.5}$ & $2.2^{+0.4}_{-0.3}$ & \\nodata & \\nodata & \\nodata & $0.5(40)$\\\\  \\phm{NGC\\,5457} sim.\\ low & & \\texttt{ab(me+pl)} & \\strt $0.3^{+0.3}_{-0.1}$ & $0.42^{+1.9}_{-0.42}$ & $2.0^{+0.4}_{-0.2}$ & \\nodata & \\nodata & \\nodata & $0.5 (28)$\\\\ \\phm{NGC\\,5457} sim.\\ high & & \\texttt{ab(pl)} & \\strt \\nodata & $1.1 \\pm 0.6$ & $2.0^{+0.4}_{-0.2}$ & \\nodata & \\nodata & \\nodata & $0.5 (41)$\\\\  \\phm{NGC\\,5457} merged & & \\texttt{ab(pl)} & \\strt \\nodata & $0.6 \\pm 0.4$ & $1.9 \\pm 0.2$ & \\nodata & \\nodata & \\nodata & $0.6 (70)$\\\\ \\enddata  \\tablenotetext{a}{Model labels: \\texttt{ab}=\\texttt{xswabs}, photoelectric absorption; \\texttt{ga}=\\texttt{xswabs} with value frozen at Galactic column density towards this target; \\texttt{br}=\\texttt{xsbremss}, thermal bremsstrahlung; \\texttt{g}=\\texttt{gauss1d}, one-dimensional Gaussian; \\texttt{me}=\\texttt{xsmekal}, thermal plasma; \\texttt{pl}=\\texttt{powlaw1d}, one-dimensional power law; \\texttt{rn}=\\texttt{xspexrav}, power-law reflected by neutral material.} \\tablenotetext{b}{Reflection scaling factor.} \\tablenotetext{c}{Cash statistic (\\texttt{cstat} in \\textit{Sherpa}).} \\tablenotetext{d}{Unconstrained by fit.} \\end{deluxetable}    \\begin{deluxetable}{cccccccc} \\tablewidth{0pt} \\tablecaption{Inferred nuclear luminosities\\label{tab:lum}} \\tablecolumns{8} \\tablehead{ \\colhead{} & \\colhead{Filter} & \\colhead{NGC 3169} & \\colhead{NGC 4102} & \\colhead{NGC 4647} & \\colhead{NGC 3184} & \\colhead{NGC 5457} & \\colhead{NGC 4713}\\\\ \\colhead{} & \\colhead{or Band} & \\colhead{(Sa)} & \\colhead{(Sb)} & \\colhead{(Sc)} & \\colhead{(Scd)} & \\colhead{(Scd)} & \\colhead{(Sd)} } \\startdata X-ray & 0.3--8 keV & 41.7 & 40.2 & 39.0\\tablenotemark{ab} & 37.3\\tablenotemark{c} & 37.5--38.5\\tablenotemark{d} & 38.6\\tablenotemark{c}\\\\ UV \\& Opt. & F300W & \\nodata & \\nodata & \\nodata & $< 39.4$ & \\nodata & \\nodata \\\\ & F336W & \\nodata & \\nodata & \\nodata & \\nodata & 39.0 & \\nodata \\\\ & F547M & \\nodata & \\nodata & \\nodata & \\nodata & 39.5 & \\nodata \\\\ & F606W & \\nodata & \\nodata & \\nodata & 39.5 & \\nodata & 40.1 \\\\ IR & $K_s$ & 42.6 & 42.8 & 38.8 & 40.9 & 41.0 & 41.0 \\\\ Radio & 15 GHz & \\nodata & \\nodata & \\nodata & \\nodata & \\nodata & $< 36.8$ \\\\ & 5 GHz & 37.2 & \\nodata & \\nodata & \\nodata & $< 35.8$ & \\nodata \\\\ & 1.4 GHz & \\nodata & 38.0 & 36.9\\tablenotemark{b} & $< 35.1$ & $< 34.9$ & $< 35.7$ \\enddata \\tablecomments{Values are $\\log(L/\\mathrm{erg\\:s}^{-1})$ in the specified bandpass in the x-ray and $\\log(\\nu L_\\nu/\\mathrm{erg\\:s}^{-1})$ where a single filter, wavelength, or frequency is given. Note that the values are derived from observations with varied instruments, PSFs, apertures, and signal-to-noise ratios, and the uncertainties can be as large as a factor of two.} \\tablenotetext{a}{0.3--12 keV.} \\tablenotetext{b}{{L}arge positional uncertainty makes it unclear whether the detected source is really the nucleus.} \\tablenotetext{c}{Based on a small number of counts. A power-law spectrum was assumed. Intrinsic absorption is unknown.}   \\tablenotetext{d}{Variable source.} \\end{deluxetable}   \\begin{deluxetable}{cccccccc} \\tablewidth{0pt} \\tablecaption{NGC 3184 luminosity in \\textit{Spitzer\\/} bands.\\label{tab:n31spitz}}  \\tablecolumns{3} \\tablehead{ \\colhead{Band} & \\colhead{Aperture} & \\colhead{$\\log(\\nu L_\\nu)$}\\tablenotemark{a} \\\\ \\colhead{$(\\micron)$} & \\colhead{Radius} & \\colhead{(\\es)} } \\startdata 3.6 & $3\\arcsec$ & 40.6\\\\ 4.5 & $3\\arcsec$ & 40.4\\\\ 5.8 & $3\\arcsec$ & 40.7\\\\ 8.0 & $3\\arcsec$ & 40.8\\\\ 24 & $6\\arcsec$ & 41.1\\\\ 70 & $10\\arcsec$ & 41.5\\\\ 160 & $10\\arcsec$ & 41.7\\\\ \\enddata \\tablenotetext{a}{Uncertainties are 10\\% in the $3.6\\micron$, $4.5\\micron$, $5.8\\micron$, $8.0\\micron$, and $160\\micron$ bands; 4\\% at $24\\micron$; 12\\% at $70\\micron$.} \\end{deluxetable}"
    },
    "0801/0801.1875_arXiv.txt": {
        "abstract": "The Einstein radius of a cluster provides a relatively model-independent measure of the mass density of a cluster within a projected radius of $\\sim 150$ kpc, large enough to be relatively unaffected by gas physics. We show that the observed Einstein radii of four well-studied massive clusters, for which reliable virial masses are measured, lie well beyond the predicted distribution of Einstein radii in the standard $\\Lambda$CDM model.  Based on large samples of numerically simulated cluster-sized objects with virial masses $\\sim 10^{15}M_{\\odot}$, the predicted Einstein radii are only $15-25\\arcsec$, a factor of two below the observed Einstein radii of these four clusters. This is because the predicted mass profile is too shallow to exceed the critical surface density for lensing at a sizable projected radius. After carefully accounting for measurement errors as well as the biases inherent in the selection of clusters and the projection of mass measured by lensing, we find that the theoretical predictions are excluded at a 4-$\\sigma$ significance. Since most of the free parameters of the $\\Lambda$CDM model now rest on firm empirical ground, this discrepancy may point to an additional mechanism that promotes the collapse of clusters at an earlier time thereby enhancing their central mass density. ",
        "introduction": "The standard picture of the basic cosmological framework has recently come to rest firmly on detailed empirical evidence regarding the cosmological parameters, the proportions of baryonic and non-baryonic dark matter, together with the overall shape and normalization of the power spectrum \\citep[e.g.,][]{SNV06,Spergel07,BAO07}.  This framework has become the standard $\\Lambda$CDM cosmological model, with the added simple assumptions that the dark matter reacts only to gravity, is initially sub-relativistic, and possesses initial density perturbations which are Gaussian distributed in amplitude. This is a very well defined and relatively simple model, with clear predictions which are amenable to examination with observations. The cooling history of baryons complicates the interpretation of dark matter on galaxy scales, especially for dwarf galaxies that traditionally have been a major focus of studies of halo structure. Clusters have the advantage that the virial temperature of the associated gas is too hot for efficient cooling, so the majority of the baryons must trace the overall gravitational potential and hence we may safely compare lensing-based cluster mass measurements to theoretical predictions that neglect gas physics and feedback.  Lensing-based determinations of the mass profiles of galaxy clusters rely on detailed modeling of the strong lensing region to define the inner mass profile, and also a careful analysis of the outer weak lensing regime. The latter involves substantial corrections for instrumental and atmospheric effects \\citep{KSB}, and a clear definition of the background, free of contamination by the lensing cluster \\citep{Br05b,Medezinski}. In the center we may make use of the Einstein radius of a cluster which is often readily visible from the presence of giant arcs and provides a relatively model-independent determination of the central mass density. In the case of axial symmetry, the projected mass inside the Einstein radius $\\tE$ depends only on fundamental and cosmological constants: $M(<\\tE)=\\tE^2 (c^2/4G) D_{\\rm OL} D_{\\rm OS} / D_{\\rm LS}$, where this combination of angular diameter distances (observer-lens, observer-source, and lens-source) leads to a relatively weak dependence on the lens and source redshifts. More generally, an effective Einstein radius can be defined by axially averaging the projected surface density, which itself is well determined when there are a large number of constraints. Virtually all known massive clusters at intermediate redshifts, $0.15<z<0.8$, show multiple images including obvious arcs in sufficiently deep high-resolution data. The derived Einstein radius of these massive clusters typically falls in the range $10\\arcsec< \\tE< 20\\arcsec$ \\citep{Gioia,Smith}, with the largest known case of $\\sim 50\\arcsec$ for A1689.  Increasingly large simulations have helped to specify the evolution of the halo mass function and the form of the mass profile predicted in the context of the $\\Lambda$CDM model \\citep{NFW,Bull}. These simulations are now becoming sufficiently large and detailed to define the predicted spread of halo structure over a wide range of halo mass, and to quantitatively assess the inherent bias in observing clusters in projection and selecting them by lensing cross-section \\citep{Hennawi,Neto}. For the most massive collapsed objects in these simulations (virial mass $\\Mv \\sim 10^{15} M_{\\odot}$), a mean observed concentration of $\\c2\\sim 6$ is predicted for lenses, where $\\c2$ is defined precisely in the next section. Such profiles are relatively shallow and seem at odds with recent careful lensing studies of massive clusters; although the \\citet{NFW} (hereafter NFW) profile provides acceptable fits to the observations, relatively high concentrations of $\\cv\\sim 10-15$ are derived for several well-studied massive clusters \\citep{Kneib,Gavazzi,Br05b,kling,Limousin,Bradac,halkola08,umetsu}. These values are larger than expected based on simulations of the standard $\\Lambda$CDM model. Given the relatively shallow mass profile predicted for cluster-mass CDM halos, the question arises whether the projected critical surface density for lensing can be exceeded within a substantial radius for this model.  In this paper we compare observations of well-constrained massive clusters with the predictions of $\\Lambda$CDM simulations. The idea is to compare directly the projected 2-D mass distributions in the observations and the simulations. In the observations, the projected surface density is obtained from the lensing analysis, which we emphasize does not assume any symmetry of the lensing mass. In the simulations, the 2-D density is directly measurable from the simulation outputs. We summarize each 2-D density distribution with two defined quantities, the virial mass and the effective Einstein radius. In the simulations, the 2-D density distribution of each halo was already axially averaged by fitting to a projected NFW profile, so we use this profile to obtain the effective $\\tE$ (which, consistently, is also defined through an axial average). The virial mass is measured directly in the simulations (with its usual spherically-averaged definition). In the observations, we obtain $\\Mv$ directly in A1689 (though in projection), and measure it using NFW fits to the final distribution in the other clusters.  This paper is structured as follows. In section~2 we first summarize the theoretical predictions, including a brief review of the NFW profile and its lensing properties, and of the halo concentrations measured by \\citet{Neto} and \\citet{Hennawi} in large numerical simulations. Note that the conflict between high observed concentrations and lower ones determined for the numerical halos was noted in both of these papers (see also \\citet{navarro}). We then present the observational data for the four clusters, followed by a model-independent method for measuring the mass, which we apply to A1689. In section~3 we confront the theoretical predictions with the data, finding a clear discrepancy. We discuss the possible implications in section~4. ",
        "conclusions": "We have presented perhaps the clearest, most robust current conflict between observations and the standard $\\Lambda$CDM model. This model is highly successful in fitting large scale structure measurements, which in turn strongly constrain the free parameters of the model and thus produce precise predictions for comparison with data on smaller scales. Structure on these scales is non-linear and potentially affected by gas physics, but clusters provide perhaps the best opportunity for a robust comparison between the models and the theory. Clusters are so large and massive that their evolution is dominated by gravity, especially since their high virial temperature prevents most of the intracluster gas from cooling. The evolution of clusters including gravitational collapse and virialization can now be accurately numerically simulated, with sufficient resolution for studying cluster structure and for simulating lensing in projection, and in sufficient volumes to produce large samples in a cosmological context. At the same time, observations of clusters combining weak and strong lensing now produce accurate virial mass determinations. Given the virial mass, the cleanest measure of the halo structure is the effective Einstein radius, which is easily obtained observationally from a model constrained by large numbers of arcs and directly measures the central mass density.  We derived the theoretical predictions for cluster lensing in $\\Lambda$CDM using the distribution of 3-D halo profiles measured by \\citet{Neto} in the Millennium simulation, after correcting it for lensing and projection biases based on \\citet{Hennawi}. These analyses of numerical samples expressed halo structure in terms of the NFW concentration parameter. We found two key results (Figure~\\ref{fig:Hennawi}) based on the halo analysis by \\citet{Hennawi}: first, that the distribution of 3-D concentrations of the lens population is the same as that of the general halo population except for a shift upwards by a factor of 1.17; and second, that the concentrations measured in projection are related to the 3-D concentrations, such that the ratio follows a lognormal distribution which corresponds to a factor of 1.14 shift plus a factor of 1.33 spread.  The concentration parameter is higher for relaxed halos than for unrelaxed, and it declines slowly with halo mass, resulting in a predicted Einstein radius that increases roughly linearly with mass for relaxed halos (Figure~\\ref{fig:NetoMean}).  \\begin{figure}% \\includegraphics[width=84mm]{rEinstein2.eps} \\caption{Total number of observable clusters in the universe above redshift $z$, obtained by integrating the halo mass function of \\citet{Sheth} over our past light cone. We consider all cluster halos above virial mass $M=1$, 2, 3, 4, 5, or $6 \\times 10^{15} M_{\\odot}$ (top to bottom). Also listed for each $M$ (top-right corner) is the redshift above which there is a $50\\%$ chance of observing at least one halo of mass greater than $M$. For $M=6 \\times 10^{15} M_{\\odot}$ we instead list the probability of observing at least one halo at any $z>0$.} \\label{fig:Nofz} \\end{figure}  We compared the theoretical predictions with the observed $\\tE$ for four clusters, A1689, Cl0024, A1703, and RXJ1347. For the latter three we used the virial mass as given by NFW fits to the lensing observations, but for A1689 we obtained a model-independent mass directly from the lensing data, only assuming spherical symmetry (section~2.3). For each object, the predicted $\\tE$ values came up short by a factor of two compared with the observations (Figure~\\ref{fig:rEofz}). After including the measurement errors, the full probability distribution functions of the predicted Einstein radii excluded the theoretical model at 2-$\\sigma$ for each object. The total probability of the standard $\\Lambda$CDM model yielding four clusters with such large $\\tE$ is $3 \\times 10^{-5}$, a 4-$\\sigma$ discrepancy.  Lensing work is now being extended to larger samples of clusters, so that in the near future we may examine more fully the relation between the Einstein radius and virial mass and its scatter, over a wider range of cluster masses. The theoretically predicted triaxiality of CDM halos implies a particular scatter in the projected concentration parameter (and thus in the Einstein radius) for a given halo mass. This scatter, which we included in our calculations, is apparently insufficient to explain the observations. If the scatter is observationally determined to be relatively small, then this would further highlight the problem we have discussed and leave the high concentrations unexplained. Determining the scatter observationally will also statistically probe the degree of triaxiality of CDM halos. To ensure the most direct, unbiased comparison, the simulated distributions should be calculated not at a fixed 3-D virial mass, but at a fixed projected, effective virial mass, defined as in section~2.3. We expect this projected virial mass to also be observationally measured in more clusters.  Numerical simulations show a clear correlation between the concentration of a halo and its formation time, i.e., the time at which a significant portion of the halo mass first assembled \\citep[e.g.,][]{Neto}. This agrees with the intuitive notion that a dense halo core must have assembled at high redshift, when the cosmic density was high. Thus, the fact that observed cluster halos are apparently more centrally concentrated than is predicted in $\\Lambda$CDM suggests an additional mechanism that promotes the collapse of cluster cores at an earlier time than expected. Baryons are unlikely to help. Central cD galaxies contribute only a small fraction of the mass within the Einstein radius, which for our four clusters is $\\sim 150$ kpc enclosing a projected mass of $\\sim 2 \\times 10^{14} M_\\odot$, or a 3-D mass of $\\sim 1 \\times 10^{14} M_\\odot$. We can estimate the effect of baryons on the total mass profile using the simple model of adiabatic compression \\citep{blumenthal}. Within this model, conservation of angular momentum implies that the quantity $r M(r)$ (assuming spherical symmetry) is fixed. Assuming that we start out with a halo with the mean expected theoretical concentration (Figure~\\ref{fig:NetoMean}), the observed 3-D mass within the Einstein radius can be obtained through adiabatic compression if the enclosed baryonic mass within this radius is $\\sim 3 \\times 10^{13} M_\\odot$ in the four clusters we considered. Thus, explaining the discrepancy through adiabatic compression requires the baryonic fraction within the Einstein radius to be $\\sim 1/3$, twice the cosmic baryon fraction. This seems highly unlikely, as the observed X-ray emission yields at these radii gas fractions well below the cosmic value (e.g., see Figure~12 of \\citet{Doron} for A1689), and a cD galaxy contains only $\\sim 1 \\times 10^{12} M_\\odot$ in baryons (e.g., see \\citet{M87} for M87).  Modifications in the properties of dark matter or the slope of the primordial power spectrum are generally expected to have a smaller effect on clusters than on smaller-scale objects which are predicted in $\\Lambda$CDM to have earlier formation times and higher concentrations. On the other hand, since clusters are rare objects in the standard model, primordial non-Gaussianity would significantly affect them and allow clusters to form earlier, which may also help explain other observations \\citep{silk,yoel}. Regardless of the mechanism, early collapse of cluster cores may have observable consequences if it is accompanied by star formation.  Finally, we note that the fact that clusters are now being detected with masses $\\sim 10^{15} M_\\odot$ is completely consistent with the $\\Lambda$CDM model. Indeed, Figure~\\ref{fig:Nofz} shows that large numbers of clusters are expected out to significant redshifts, including $M = 2 \\times 10^{15} M_\\odot$ halos out to $z \\sim 1$, as well as more massive halos up to $\\sim 5\\times 10^{15} M_\\odot$ at lower redshift. Clearly, while large samples of halos with precise, profile-independent lensing determinations of both $\\tE$ and $\\Mv$ will make our results completely conclusive, the highly-significant discrepancy we have identified already represents a substantial challenge for $\\Lambda$CDM."
    },
    "0801/0801.1334_arXiv.txt": {
        "abstract": "In the coming year, the Large Hadron Collider will begin colliding protons at energies nearly an order of magnitude beyond the current frontier. The LHC will, of course, provide unprecedented opportunities to discover new particle physics.  Less well-known, however, is that the LHC may also provide insights about gravity and the early universe.  I review some of these connections, focusing on the topics of dark matter and dark energy, and highlight outstanding prospects for breakthroughs at the interface of particle physics and cosmology. ",
        "introduction": "The Large Hadron Collider (LHC) is scheduled to begin running in the summer of 2008.  Conceived around 1984 and approved in 1994, the LHC will provide the first detailed look at the weak energy scale $\\mweak \\sim 100~\\gev - 1~\\tev$ by colliding protons with protons at the center-of-mass energy $E_{\\text{CM}} = 14~\\tev$ and ultimate luminosity ${\\cal L} = 100~\\ifb/\\yr$.  This is far beyond the current energy frontier, where the Tevatron collides protons and anti-protons with $E_{\\text{CM}} = 2~\\tev$ and ${\\cal L} \\sim 1~\\ifb/\\yr$.  As an illustration of the power of the LHC, top quarks, discovered in 1994 with a handful of events and currently produced at the Tevatron at the rate of $\\sim 1000$ per year, will be produced at the rate of 10 Hz when the LHC reaches its design luminosity.  The {\\em raison d'etre} for the LHC is the discovery of the Higgs boson and associated microphysics, including supersymmetric and other postulated particles.  In recent years, however, the LHC's potential for providing insights into gravity and cosmology have taken on increasing importance.  My goal here is to review some recent developments and to highlight a few scenarios in which the implications of the LHC for our understanding of gravity and the early universe may, in fact, be profound. ",
        "conclusions": "In the coming year, the LHC will probe the weak scale $\\mweak \\sim 100~\\gev - 1~\\tev$ in great detail.  This has implications for new particle physics, but may also open up new windows on the early universe, and tests of gravity in rather unusual regimes.  At present, the evidence for particle dark matter is as strong as ever.  The possibility of WIMP dark matter is well-motivated, and there has been a recent proliferation of candidates.  At the same time, there has also been a great deal of progress on the alternative possibility of superWIMP dark matter.  In virtually all cases, the LHC will be able to produce these candidates, and in some cases, precision measurements at the LHC may be able to determine the candidate's relic density.  Comparison with observations may then provide compelling evidence that particles produced at the LHC do, in fact, constitute the dark matter.  Such studies will also be able to determine the microscopic properties of the WIMP particle.  Colliders may also provide an interesting window on gravity in unusual environments.  For example, in the superWIMP scenarios, one may probe gravitational interactions between fundamental particles and provide quantitative evidence for supergravity.  In the WIMP scenarios, the thermal relic density studies simultaneously bound new contributions to dark energy at the time of freezeout, probing variations in the strength of gravity at $\\sim 1$ ns after the Big Bang, and possibly shedding light on the dark energy problem.  It is rather striking that in many of these scenarios, the LHC, along with other observatories and experiments, could solve many old questions, such as the identity and origin of dark matter.  If any of the ideas discussed here is realized in nature, the interplay of collider physics with cosmology and astrophysics in the next few years will likely yield profound insights about the Universe, its contents, and its evolution.  \\ack  I thank the organizers of GRG18/Amaldi7 for the invitation to participate in this stimulating conference and my collaborators for their many insights regarding the work discussed here.  This work was supported in part by NSF Grants PHY--0239817 and PHY--0653656, NASA Grant NNG05GG44G, and the Alfred P.~Sloan Foundation."
    },
    "0801/0801.4511_arXiv.txt": {
        "abstract": "{Blue Stragglers Stars (BSSs) are thought to form in globular clusters by two main formation channels: $i)$ mergers induced by stellar collisions and $ii)$ coalescence or mass-transfer between companions in binary systems. The detailed study of the BSS properties is therefore crucial for understanding the binary evolution mechanisms, and the complex interplay between dynamics and stellar evolution in dense stellar systems.} {We present the first comparison between the BSS specific frequency and the binary fraction in the core of a sample of Galactic globular clusters, with the aim of investigating the relative efficiency of the two proposed formation mechanisms.} {We derived the frequency of BSSs in the core of thirteen low-density Galactic globular clusters by using deep ACS@HST observations and investigated its correlation with the binary fraction and various other cluster parameters.} {We observed a correlation between the BSS specific frequency and the binary fraction. The significance of the correlation increases by including a further dependence on the cluster central velocity dispersion.} {We conclude that the unperturbed evolution of primordial binaries could be the dominant BSS formation process, at least in low-density environments.} ",
        "introduction": "Blue Straggler Stars (BSSs) are objects that, in the color-magnitude diagram (CMD) of evolved stellar populations, lie along an extension of the Main Sequence (MS), in a region which is brighter and bluer than the Turn-off (TO). First discovered by Sandage (1953) in M3, they have been observed in all Galactic globular clusters (GCs; Piotto et al. 2004), in the field population (Carney et al. 2005), and in dwarf galaxies of the local group (Momany et al. 2007).  Their location in the CMD suggests that BSSs have masses of $1.2 \\div 1.5~M_{\\odot}$, significantly larger than those of normal stars in old stellar systems (like GCs). Thus, they are thought to have increased their mass during their evolution. Two mechanisms have been proposed for their formation: $i)$ the merger of two stars induced by stellar collision (COL-BSSs; Hills \\& Day 1976) and $ii)$ coalescence or mass-transfer between two companions in a binary system (MT-BSSs; McCrea 1964). The two formation channels are thought to act with different efficiencies according to the cluster structural parameters (Fusi Pecci et al. 1992) and they can work simultaneously within the same cluster in different radial regions, corresponding to widely different stellar densities (Ferraro et al. 1997; Mapelli et al. 2006). Indeed, collisions are more frequent in the central region of GCs, because of the high stellar density, while MT-BSSs mainly populate the cluster periphery, where binary systems can more easily evolve in isolation without suffering exchange or ionization due to gravitational encounters. The whole scenario is further complicated by the cluster dynamical evolution that leads massive systems (like binaries and BSSs) to sink toward the cluster center in a time-scale comparable to the cluster relaxation time.  A possible tool for distinguishing COL-BSSs from MT-BSSs is based on high-resolution spectroscopic analysis. In fact, anomalous chemical abundances are expected at the surface of BSSs resulting from mass-transfer activity (Sarna \\& de Greve 1996), while they are not predicted for COL-BSSs (Lombardi et al. 1995). However, such studies have just become feasible and they are limited to only a small number of BSSs in just one cluster (47 Tucanae; Ferraro et al. 2006).  As an alternative way for getting insights on the relative efficiency of the two formation mechanisms, here we investigate possible correlations between the BSS population and the host cluster properties. ",
        "conclusions": "We measured the BSS specific frequency in the core of thirteen low-density Galactic GCs and investigated its correlation with different dynamical and general cluster parameters. We found evidences that, at least in this density regime, binary-rich environments are more efficient in producing BSS. No correlations have been found with the cluster central density, concentration, stellar collision rate, and half-mass relaxation time, in agreement with the results of Piotto et al. (2004) and Leigh et al. (2007). These evidences indicate that the collisional channel for the BSS formation has a very small efficiency in low-density GCs, while the mechanisms involving the unperturbed evolution of binary systems are dominant.  The higher significance of the trivariate correlation among the BSS frequency, the binary fractions and the cluster velocity dispersion, indicates that, for a given binary fraction, the BSS specific frequency decreases with increasing velocity dispersion. This finding might be connected with the effect of the cluster velocity dispersion in the dynamical evolution of binary systems. In fact, a small cluster velocity dispersion corresponds to a lower energy limit between soft and hard binaries\\footnote{A binary is defined {\\it soft} ({\\it hard}) if its binding energy ($E = -G~m_{1}~m_{2}/2~a$, with $a$ being the orbital separation of the two components) is smaller (larger) than the mean kinetic energy of normal cluster stars ($K = m~\\sigma_{\\rm v}^2$, with $m$ being the average mass of cluster stars).}, i.e., to a larger fraction of hard binary systems. Since the natural evolution of hard binaries is to increase their binding energy (i.e. decrease their orbital separation; Heggie 1975), this implies that low velocity dispersion GCs should host a larger fraction of hard (and close) binaries, able to both survive possible stellar encounters, and activate mass-transfer and/or merging processes between the companions. A larger fraction of BSSs formed by the evolution of primordial binaries is therefore expected in lower velocity dispersion GCs (see also Davies et al. 2004). Such an effect of $\\sigma_{\\rm v}$ (in terms of both hardening and shrinking the binary systems) might be, in turn, at the origin of the inverse correlation between the BSS frequency and the cluster total luminosity (mass) observed in open clusters (De Marchi et al. 2006), low density GCs (Sandquist 2005), as well as high density GCs (Piotto et al. 2004; Leigh et al. 2007). Indeed, the more massive GCs have larger central velocity dispersions (as a consequence of the virial theorem for systems with similar radii, as GCs; see Fig. 1 of Djorgovski 1995).  For most of these clusters, however, the binary fraction is still unknown and the trivariate correlation between $F$, $\\xi_{bin}$ and $\\sigma_{\\rm v}$ cannot be derived. If its significance and its interpretation in terms of the velocity dispersion effect is confirmed also in high-density clusters, then stellar collisions might play a secondary role in the production of BSSs, and the evolution of primordial binaries should always be the dominant process. For the moment, however, these conclusions remain speculative, since the sample of GCs analyzed here, in spite of being the largest to date with known binary fraction, is still too small for statistically reliable assessments. Enlarging the sample of GCs with known binary and BSS fraction is therefore essential and urgent to verify the findings presented in this paper on a more robust statistical basis."
    },
    "0801/0801.1116_arXiv.txt": {
        "abstract": "We describe new, deep MIPS photometry and new high signal-to-noise optical spectroscopy of the 2.5 Myr-old IC 348 Nebula.  To probe the properties of the IC 348 disk population, we combine these data with previous optical/infrared photometry and spectroscopy to identify stars with gas accretion, to examine their mid-IR colors, and to model their spectral energy distributions.  IC 348 contains many sources in different evolutionary states, including protostars and stars surrounded by primordial disks, two kinds of transitional disks, and debris disks. Most disks surrounding eary/intermediate spectral-type stars ($>$ 1.4 M$_{\\odot}$ at 2.5 Myr) are debris disks; most disks surrounding solar and subsolar-mass stars are primordial disks. At the 1--2 $\\sigma$ level, more massive stars also have a smaller frequency of gas accretion and smaller mid-IR luminosities than lower-mass stars. These trends are suggestive of a stellar mass-dependent evolution of disks, where most disks around high/intermediate-mass stars shed their primordial disks on rapid, 2.5 Myr timescales.  The frequency of MIPS-detected transitional disks is $\\approx$ 15--35\\% for stars plausibly more massive than 0.5 M$_{\\odot}$. The relative frequency of transitional disks in IC 348 compared to that for 1 Myr-old Taurus and 5 Myr-old NGC 2362 is consistent with a transition timescale that is a significant fraction of the total primordial disk lifetime. ",
        "introduction": "Most $\\le$ 1 Myr-old stars are surrounded by optically-thick, accreting {\\it primordial} disks, which produce strong near-to-mid infrared (IR) emission \\citep[L$_{d}$/L$_{\\star}$ $\\ge$ 0.1;][]{Kh95}. The gas and dust in these disks comprise the building blocks of planets.  As stars age, the dust grains grow, settle toward the midplane, and become incorporated into planetesimals. Circumstellar gas depletes by accretion onto the star \\citep{Ha98}. All of these processes reduce the amount and frequency of disk emission on timescales of $\\sim$ 3--7 Myr for most stars \\citep{He07a, Hi08b,Clp09}.  By $\\sim$ 10 Myr, most stars do not have optically-thick primordial disks \\citep{Cu07a,Hi08b}.  Disks around these older stars are typically debris disks \\citep{Cu08a}, which have weaker, optically-thin dust emission (L$_{d}$/L$_{\\star}$ $\\lesssim$ 10$^{-3}$) and lack evidence for circumstellar gas accretion. Because stellar radiation -- radiation pressure and Poynting-Robertson drag -- or stellar wind drag can remove dust in debris disks on timescales much less than the age of the star, their dust requires a replenishment source, which is supplied by active icy or rocky planet formation \\citep[e.g.][]{Bp93, Kb04, Kb08}.  Because debris disks lack copious amounts of circumstellar gas needed to form gas giant planets, identifying when most primordial disks turn into debris disks pinpoints an empirical upper limit for the formation timescale for gas giant planets.  Recent \\textit{Spitzer Space Telescope} studies of 1--5 Myr-old clusters indicate that primordial disks disappear faster around early type, high/intermediate-mass stars than they do around late type, low-mass stars \\citep{Ca06, La06, He07a, Clp09}.  By 5 Myr, debris disks completely dominate the disk population around high/intermediate-mass stars.  At 5 Myr, most disks around lower-mass stars appear to have inner holes and/or extremely low warm dust masses \\citep{Dahm09,Clp09}, which indicates that they may be disappearing and actively making the primordial-to-debris disk transition.  Thus, the epoch at which most high/intermediate-mass stars make the primordial-to-debris disk transition must occur at an age prior to 5 Myr and may be identified by mid-IR observations of younger clusters.  With a median age of $\\sim$ 2--3 Myr, stars in the IC 348 Nebula \\citep{He98} may help to constrain the epoch when most primordial disks around high/intermediate-mass stars turn into debris disks.  Cluster membership is well determined \\citep[e.g.][]{He98,Lu98,Lu03,Mu07}. At only 320 pc distant, the cluster provides a sensitive probe of low-levels of disk emission in spite of its high infrared background.  Recently, \\citet{La06} and \\citet{Muzerolle2006} analyzed Spitzer Cycle 1 IRAC and MIPS IC 348 data and found evidence for substantial disk evolution. Based on the IRAC 3.6 $\\mu m$ to 8 $\\mu m$ flux slope, \\citet{La06} show that only $\\approx$ 30\\% of cluster stars have strong ('thick') IRAC excess emission comparable to stars with primordial disks (e.g. most disk-bearing stars in Taurus). Another $\\approx$ 23\\% have weaker ('anemic') IRAC excess emission, and $\\approx$ 44\\% lack IRAC excess emission ('diskless').  However, inferring the physical state of disks from the observed IRAC slope is not straightforward. Both debris disks and 'evolved' primordial disks ('transitional' disks: those with lower disk masses and/or inner holes) can have weak IRAC-excess emission \\citep{Cu07b, Clp09}.  Furthermore, while SED analysis including longer wavelength data (e.g. 24 $\\mu m$) helps to break this degeneracy \\citep{Clp09}, only $\\approx$ 30\\% of IC 348 stars have Cycle 1 MIPS detections. Therefore, deeper MIPS data coupled with other diagnostics of the disk evolutionary state (e.g. gas accretion signatures) would help to make IC 348 a better laboratory for studying disk evolution.  In this paper, we combine new, high signal-to-noise optical spectra and new, deeper MIPS 24 $\\mu m$ and 70 $\\mu m$ photometry of IC 348 with archival data to study disk evolution in more detail.  We first use the optical spectra to search for evidence of circumstellar gas accretion, which is absent in debris disks. Deeper MIPS data is then used to establish a much larger sample of disk-bearing stars.  From these data, we compare the disks SEDs to models to constrain their evolutionary states, to compare disk properties around high-mass stars to low-mass stars, and to investigate the timescale for and the morphology of the primordial-to-debris disk transition. ",
        "conclusions": "\\subsection{Summary of Results} Using new optical spectra and new, deep MIPS 24 $\\mu m$ and 70 $\\mu m$ photometry we investigated the disk population of the 2.5 Myr-old IC 348 Nebula.  Combining these data with previous work from \\citet{La06}, we analyzed optical spectroscopic data for all stars, performed SED modeling, and examined the mid-IR colors of MIPS-detected members to constrain the disks' evolutionary states and probe how disk properties vary with stellar mass.  Our study yields the following major results: \\begin{itemize} \\item IC 348 sources with MIPS-70 $\\mu m$ detections include flat-spectrum protostars and pre-main sequence stars with optically-thick, luminous far IR disk emission. Some optically-thick disks show evidence for depleted inner disks; others do not.  \\item IC 348 stars with new and previous MIPS detections have disks in evolutionary states ranging from primordial disks to debris disks. In agreement with recent work \\citep{La06, He07a, Clp09}, we find evidence for two separate pathways from primordial disks to debris disks.  In homologously depleted disks, disks deplete their reservoir of small dust grains at all disk locations simultaneously.  In a second sequence ('disks with inner holes'), disks clear their supply of small dust grains from the inside out.  \\item At a $\\sim$ 1--2 $\\sigma$ significance, signatures of circumstellar gas accretion are more frequent for solar and subsolar-mass stars than for high/intermediate-mass stars.  This result is consistent with a gas disk dispersal timescale that is shortest for high/intermediate-mass stars.  \\item The mid-IR disk luminosities of MIPS-detected disks are stellar-mass dependent.  Relative to the stellar photosphere, 24 $\\mu m$ emission from disks is lower for high/intermediate-mass stars than for solar/subsolar-mass stars.  If disks generally decline in luminosity as a function of time, this result implies that disks around high/intermediate-mass stars evolve faster.  \\item The evolutionary states of MIPS-detected disks are also stellar mass dependent.  Most MIPS-detected disks surrounding high/intermediate-mass stars stars appear to be debris disks; primordial disks comprise only $\\sim$ 15\\% of the disk population for M$_{\\star}$ $>$ 1.4 M$_{\\odot}$. In contrast, most MIPS-detected disks around solar and subsolar-mass stars are primordial disks.  For stars of all masses, transitional disks (homologously depleted or disks with inner holes) comprise $\\sim$ 15--35\\% of the MIPS-detected disk population. \\end{itemize}   \\subsection{The Evolutionary State of \\textit{Anemic} Disks} Our SED modeling and investigation of the mid-IR colors of IC 348 stars clarifies the nature of \\textit{anemic} disks identified by \\citet{La06}.  Among late type stars, most anemic disks are homologously depleted disks or are disks with inner holes. Several anemic disks are probably primordial disks. Among earlier-type stars with probable masses $\\gtrsim$ 1.4 M$_{\\odot}$, anemic disks comprise a broader range of evolutionary states. At least two anemic disks (surrounding IDs 6 and 8) are probably debris disks.  In contrast, ID-31 has strong mid-IR emission and gas accretion signatures more consistent with a transitional disk, specifically one with an inner hole.  Thus, as pointed out by \\citet{Clp09}, a single IRAC slope serves as a useful first-order probe of disk properties, but a full analysis of IRAC \\textit{and} MIPS data combined with gas accretion signatures is required for an accurate taxonomy of disks. Our results indicate that a full analysis is especially necessary if the stellar population includes both high/intermediate-mass stars and lower-mass stars.  Given the diversity in disk properties for 2.5 Myr-old IC 348, we suggest that a single flux slope may be best suited as a probe of disk evolution in the youngest regions (e.g. Taurus; NGC 1333), where less diversity in disk states (e.g. few if any debris disks) is expected.  \\subsection{Constraints on the Primordial-to-Debris Disk Transition} Recent Spitzer studies have placed strong constraints on the timescales for the evolution of primordial disks into debris disks.  By 5 Myr, most high/intermediate-mass stars and solar-mass stars either lack evidence for a disk or have disk properties suggestive of a debris disk \\citep{Ca06, He08, Clp09}.  At this age, many (detectable) subsolar-mass stars have weaker levels of mid-IR disk emission than typical primordial disks \\citep{Clp09,Dahm09}.  This work shows that by 2.5 Myr most disks around high/intermediate-mass stars are either actively leaving the primordial disk phase or have already reached the debris disk phase.  Together with previous results \\citep[e.g.][]{Ca06, He07a, Clp09}, our results clearly support a stellar-mass dependent timescale for the disappearance of primordial disks and the emergence of debris disks.   Within the context of planet formation, gas giant planets around high/intermediate-mass stars have much less time to form than around low-mass stars.  If gas giant planets are more frequent around high/intermediate-mass stars \\citep{Jj07}, their formation must be very rapid and efficient.  Only recently have realistic models \\citep{Kb09} that form gas giant planets via core accretion been successful in forming the cores of such planets in the short timescales implied from this paper (2.5 Myr). Some trends in exoplanet properties, such as the semimajor axis distribution, may emerge from the competing effects of core formation timescales and primordial disk dispersal timescales \\citep{Currie2009}.  The fraction of IC 348 disks in a transitional phase serves as a useful contrast to results obtained for younger clusters like Taurus and older clusters like NGC 2362.  Based primarily on IRAS data, the computed frequency of transitional disks in Taurus is small, $<<$ 10\\% \\citep[e.g.][]{Sk90, Sp95, Ww96}. These authors then argued that the lifetime of transitional disks is $\\approx$ 1--10\\% the age of Taurus: $\\approx$ 0.01--0.1 Myr. If this inference were correct, the total disk lifetime may be several Myr, but most of the disk is dissipated rapidly.  However, analysis of Spitzer data for 5 Myr-old NGC 2362 finds a much higher frequency of transitional disks \\citep{Clp09}.  Both types of transitional disks (inner holes, homologously depleted) greatly outnumber primordial disks. This result argues for far longer typical transition timescale ($\\approx$ 1 Myr).  In the intermediate age IC 348, the frequency of transition disks is intermediate between the frequency for 1 Myr-old Taurus and 5 Myr-old NGC 2362 for solar/subsolar-mass stars. Thus, the number of transition disks relative to primordial disks around solar/subsolar-mass stars appears to be an increasing function of stellar age.  This result is expected if the typical transition timescale is an appreciable fraction of the typical primordial disk lifetime.  Thus, the transition disks in Taurus and IC 348 could remain in such a state for an extended period.  However, it is possible that transition disk lifetimes, like primordial disk lifetimes, also have an intrinsic dispersion (i.e., a gaussian distribution of lifetimes), where sources in Taurus represent those that make the primordial-to-debris disk transition fastest and sources in 5--10 Myr-old clusters make the transition the slowest. Quantifying the range of times a disk spends in a transitional phase and whether the typical timescale depends on stellar mass requires observations of many more 1--10 Myr-old clusters.  While debris disks typically only have excess emission longwards of $\\approx$ 20 $\\mu m$, the debris disk population in IC 348 contains at least two sources (ID-6 and ID-8) that have warmer dust more indicative of terrestrial planet formation than icy planet formation. Thus, warm debris disks consistent with the observable signatures of terrestrial planet formation may emerge as early as $\\approx$ 2.5 Myr around high/intermediate-mass stars. Observations of more 2--5 Myr-old clusters are needed to determine if the frequency of warm debris disks is higher than the $\\sim$ 4\\% derived for 10--15 Myr-old clusters \\citep{Cu07a, Cu08}. The existence of both warm debris disks and colder debris disks (lacking IRAC excess emission) suggests that even at 2.5 Myr debris disk populations may exhibit a range of dust temperatures consistent with a range of locations over which the debris-producing stages of planet formation are ongoing."
    },
    "0801/0801.4505_arXiv.txt": {
        "abstract": "We present measurements with the VLBA of the variability in the centroid position of \\SgrA\\ relative to a background quasar at $7\\,\\mm$ wavelength.  We find an average centroid wander of $71\\pm 45\\,\\muas$ for time scales between 50 and  $100\\,\\min$ and $113\\pm50\\,\\muas$ for timescales between 100 and $200\\,\\min$, with no secular trend.  These are sufficient to begin constraining the viability of the accretion hot-spot model for the radio variability of \\SgrA.  It is possible to rule out hot spots with orbital radii above $15\\,G M_{\\rm Sgr A*}/c^2$ that contribute more than 30\\% of the total $7\\,\\mm$ flux.  However, closer or less luminous hot spots remain unconstrained.  Since the fractional variability of \\SgrA\\ during our observations was $\\sim20$\\% on time scales of hours, the hot-spot model for \\SgrA's radio variability remains consistent with these limits. Improved monitoring of \\SgrA's centroid position has the potential to place significant constraints upon the existence and morphology of inhomogeneities in a supermassive black hole accretion flow. ",
        "introduction": "There is now overwhelming evidence that \\SgrA\\ is a supermassive black hole at the center of the Milky Way.  Many stars are observed to orbit about a common focal position, requiring an unseen mass of $\\approx4\\times10^6$~\\Msun\\ contained within a radius of less than 100~AU \\citep{Schoedel:2002,Ghez:2003}, for a distance to the center of 8.0~kpc \\citep{Reid:1993}. Accurate registration of the infrared and radio reference frames \\citep{Menten:1997,Reid:2003} reveal that the common orbital focal position is coincident with \\SgrA\\ to within measurement uncertainty of $\\approx10$~mas.  Finally, the absence of intrinsic motion of \\SgrA\\ at levels near that expected for a $4\\times10^6$~\\Msun\\ object \\citep{Reid-Brun:2004}, coupled with a size less than ~1~AU \\citep{Bower:2004}, provide a lower limit on mass density of $\\sim10^{22}$~\\Msun~pc$^{-3}$, which is only two orders of magnitude less than the density of a $4\\times10^6$~\\Msun\\ non-rotating black hole within its innermost stable orbit.  There can now be little doubt that \\SgrA\\ is a supermassive black hole.  \\citet{Reid-Brun:2004} present measurements of the position of \\SgrA\\ relative to a compact extragalactic radio source (\\EGS, also refered to as J1745-283 in earlier publications). These measurements were conducted with the NRAO \\footnote{The National Radio Astronomy Observatory is operated by Associated Universities, Inc., under a cooperative agreement with the National Science Foundation.} Very Long Baseline Array (VLBA) over a period of $8\\,\\yr$ at a wavelength of $7\\,\\mm$ ($43\\,\\GHz$) and have been used to determine the apparent proper motion of \\SgrA.  Over time scales of months or longer, \\SgrA's {\\it apparent} motion is dominated by the effects of the orbit of the Sun about the center of the Galaxy. The component of the Sun's orbit in the Galactic plane is uncertain at roughly the 10\\% level, and this limits estimation of any intrinsic motion of \\SgrA\\ at about the $\\pm20$~\\kms\\ level.  However, the component of the motion of the Sun out of the Galactic plane is known to high accuracy [$7.16\\pm0.38$~\\kms\\ toward the North Galactic Pole; \\citet{Dehnen-Binney:1998}]. After removing the effects of the Sun's motion, the residual motion of \\SgrA\\ perpendicular to the Galactic plane is very small, $\\lesssim 1\\,\\km\\,\\s^{-1}$, as expected for a supermassive black hole (SMBH) at the dynamical center of the Galaxy. While previous work concentrated on the long-term motion of \\SgrA, here we analyze its short-term position ``wander'' on time scales of hours to weeks.  Short-timescale motion of the centroid position of \\SgrA\\ would be expected if a portion of the emission comes from material orbiting about the SMBH.  The degree of centroid variability would necessarily depend upon the brightness of the orbiting material, the degree to which its emission is nonuniform, and the orbital radius dominating the total flux.  We use a simple hot-spot model to relate the constraint from the observed short-term position wander of \\SgrA\\ (\\S \\ref{OoCM}) to a constraint upon the presence of strong inhomogeneities in the accretion flow onto the SMBH as a function of hot-spot luminosity and orbital period (\\S \\ref{Constraints}). Finally, concluding remarks are contained in \\S \\ref{C}. ",
        "conclusions": "\\label{C} Possible reasons for position wander include intrinsic variations in the position of the emitting plasma (\\eg, variations in the accretion flow or perhaps in a jet) or extrinsic processes such as refractive interstellar scattering. \\SgrA\\ is observed to be diffractively scattered to a size of $\\theta_{sc} \\sim 0.5 (\\lambda/0.7~{\\rm cm})^2$~mas, where $\\lambda$ is the observing wavelength. Flux density fluctuations are modest and decrease in strength with increasing wavelength; thus strong refractive scintillations are not indicated \\citep{Gwinn:1991}. Any refractive position wander should be $\\ll \\theta_{sc}$ and should occur on time scales $>\\theta_{sc}D/v$, where $D$ is the distance and $v$ is the transverse velocity of the scattering ``screen'' relative to the observer \\citep{Romani:1986}. For $D\\approx\\Ro\\approx8$~kpc \\citep{Reid:1993} and $v\\sim100$~\\kms, characteristic of material in the inner $\\sim100$~pc of the Galaxy where large scattering sizes are observed, the refractive time scale is $>10^3$~hours. Thus, we would not expect a significant contribution to the short-term wander of \\SgrA\\ from refractive scattering. For comparison, \\citet{Gwinn:1988}, using VLBI observations of the Sgr~B2(N) \\hho\\ masers near the Galactic center, find a wander limit of $<18~\\muas$ over timescales of months for maser spots, which are diffractively scattered to a comparable size (at 22 GHz) as \\SgrA\\ (at 43 GHz). Of course, our results provide an observation limit to any refractive position wander.  Since extrinsic sources of position wander (scattering) are unlikely to be dominant, we now discuss the implications for intrinsic wander from variations in brightness within an accretion disk given in \\S\\ref{Constraints}. Our observations of the lack of short-term wander of the centroid position of \\SgrA\\ presented in \\S\\ref{section:hours} give an upper limit of $\\approx100~\\muas$ for time scales of $\\approx1~{\\rm to}~4$ hours. This translates to an upper limit on the wander versus orbital period plots in Fig.~\\ref{fig:drs} as indicated by the horizontal red line and  hatched region. (In the very unlikely event that the accretion disk inclination is both near $90^\\circ$ and oriented nearly North-South on the sky, we would need to use our NS limits, which are a factor of three weaker.) Our EW limit is below the dotted blue line in Fig.~\\ref{fig:drs}, which is for hot spot flux density dominating over disk (or possible jet) emission, for orbital periods exceeding 120~min (corresponding to orbital radii larger than $15\\,G M_{\\rm Sgr A*}/c^2$ for $M_{\\rm Sgr A*}= 4\\times10^6$ \\Msun).   For cases in which the hot spot flux density is weaker than that of the disk, somewhat longer periods are allowed.  For example, for $F_{\\rm spot}/F_{\\rm disk} = 0.37$, orbital periods longer than 5~hr are excluded.  In practice, the limits placed by current $7\\,\\mm$ VLBI are significantly weaker.  The limited sensitivity to hot spots on compact orbits is primarily due to two reasons: (1) ``long'' integration times  ($T\\gtrsim 1\\,\\hr$) average much of the short time variability out, and (2) the opacity of the accretion flow itself makes it difficult to view hot-spots on compact orbits at $7\\,\\mm$. The integration time is limited by the sensitivity issues and the small number of antennas yielding interferometer baselines $<1500$~km afforded by the current VLBA; higher bandwidth recording in the future should help alleviate this problem. The optical depth is a property of Sgr A* itself, and can only be addressed by observations at shorter wavelengths. However, even in the absence of an optically thick accretion flow, it is not possible to increase the centroid variability by more than an order of magnitude due to the intrinsically small orbital radii, as seen by comparing the blue and green limits in the lower-left panel of Fig. \\ref{fig:drs}.  Nevertheless, high-resolution astrometry is reaching sensitivities and resolutions sufficient to begin to test the hot-spot model for bright Sgr A* flares.  Unfortunately, the typical fractional variability at $7\\,\\mm$ during our observations was roughly $\\pm20\\%$, implying that significant improvement in positional accuracy will be required to constrain such events.  Since the observed centroid wander is consistent with systematic errors, owing predominantly to centimeter-scale errors in the modeling of the atmospheric path-delays, substantially increasing accuracy will require better calibration techniques. However, for the somewhat rare instances in which the spot is substantially brighter \\citep{Zhao:2001}, the VLBA at 7, or possibly 3,~mm wavelength appears poised to provide significant limits upon the existence and morphology of inhomogeneities in the accretion flow surrounding \\SgrA.  Ultimately, observations at $\\sim1$~mm wavelength with VLBI techniques or at infrared wavelengths with an instrument like GRAVITY \\citep{GRAVITY} may be necessary to image the region within $\\sim3$ Schwarzschild radii on the short time scales needed to test the hot-spot model.  \\vskip 1truecm A.B. is supported by the Priority Programme 1177 of the Deutsche Forschungsgemeinschaft. \\vskip 0.5truecm {\\it Facilities:} \\facility{VLBA}"
    },
    "0801/0801.1320_arXiv.txt": {
        "abstract": "{}{ To determine alpha effect and turbulent magnetic diffusivity for mean magnetic fields with profiles of different length scale from simulations of isotropic turbulence, and to relate these results to nonlocal formulations in which alpha and the turbulent magnetic diffusivity correspond to integral kernels. }{ A set of evolution equations for magnetic fields is solved which gives the response to imposed test fields, that is, mean magnetic fields with various wavenumbers. Both an imposed fully helical steady flow consisting of a pattern of screw-like motions (Roberts flow) and time-dependent statistically steady isotropic turbulence are considered. In the latter case the aforementioned evolution equations are solved simultaneously with the momentum and continuity equations. The corresponding results for the electromotive force are used to calculate alpha and magnetic diffusivity tensors. }{ For both the Roberts flow under the second--order correlation approximation and isotropic turbulence alpha and turbulent magnetic diffusivity are largest at large scales and these values diminish toward smaller scales. In both cases the alpha effect and turbulent diffusion kernels are well approximated by exponentials, corresponding to Lorentzian profiles in Fourier space. For isotropic turbulence the turbulent diffusion kernel is half as wide as the alpha effect kernel. For the Roberts flow beyond the second--order correlation approximation the turbulent diffusion kernel becomes negative at large scales. }{} ",
        "introduction": "\\label{Intro}  Stars and galaxies harbour magnetic fields whose scales are large compared with the scale of the underlying turbulence. This phenomenon is successfully explained in terms of mean--field dynamo theory discussed in detail in a number of text books and reviews (e.g.\\ Moffatt 1978, Krause \\& R\\\"adler 1980, Brandenburg \\& Subramanian 2005a). In this context velocity and magnetic fields are split into large--scale and small--scale components, $\\UU=\\meanUU+\\uu$ and $\\BB=\\meanBB+\\bb$, respectively. The crucial quantity of the theory is the mean electromotive force caused by the small--scale fields, $\\meanEMF=\\overline{\\uu\\times\\bb}$. In many representations it is discussed under strongly simplifying assumptions. Often the relationship between the mean electromotive force and the mean magnetic field is tacitly taken as (almost) local and as instantaneous, that is, $\\meanEMF$ in a given point in space and time is considered as determined by $\\meanBB$ and its first spatial derivatives in this point only, and the possibility of a small--scale dynamo is ignored. Then the mean electromotive force is given by \\EQ \\meanemf_i=\\alpha_{ij}\\meanB_j+\\eta_{ijk} \\partial \\meanB_j / \\partial x_k \\label{meanemfi} \\EN with two tensors $\\alpha_{ij}$ and $\\eta_{ijk}$. If the turbulence is isotropic the two tensors are isotropic, too, that is $\\alpha_{ij}=\\alpha \\delta_{ij}$ and $\\eta_{ijk}=\\eta_{\\rm t} \\epsilon_{ijk}$ with two scalar coefficients $\\alpha$ and $\\eta_{\\rm t}$. Then the expression \\eq{meanemfi} simplifies to \\EQ \\meanEMF=\\alpha\\meanBB-\\eta_{\\rm t}\\meanJJ \\, , \\label{meanEMF} \\EN where we have denoted $\\nab \\times \\BB$ simply by $\\JJ$ (so that $\\JJ$ is $\\mu_0$ times the electric current density, where $\\mu_0$ is the magnetic permeability of free space). The coefficient $\\alpha$ is, unlike $\\eta_{\\rm t}$, only non--zero if the turbulence lacks mirror--symmetry. The coefficient $\\eta_{\\rm t}$ is referred to as the turbulent magnetic diffusivity.  In general, the mean electromotive force has the form \\EQ \\meanEMF=\\meanEMF_0 + \\KK \\circ \\meanBB, \\label{kernel} \\EN where $\\meanEMF_0$ stands for a part of $\\meanEMF$ that is independent of $\\meanBB$, and $\\KK \\circ \\meanBB$ denotes a convolution in space and time of a kernel $\\KK$ with $\\meanBB$ (see, e.g., Krause \\& R\\\"adler 1980, R\\\"adler 2000, R\\\"adler \\& Rheinhardt 2007). Due to this convolution, $\\meanEMF$ in a given point in space and time depends on $\\meanBB$ in a certain neighborhood of this point, with the exception of future times. This corresponds to a modification of \\eq{meanemfi} such that also higher spatial and also time derivatives of $\\meanBB$ occur.  In this paper we ignore the possibility of coherent effects resulting from small--scale dynamo action and put therefore $\\meanEMF_0$ equal to zero. For the sake of simplicity we assume further the connection between $\\meanEMF$ and $\\meanBB$ to be instantaneous so that the convolution $\\KK \\circ \\meanBB$ refers to space coordinates only. The memory effect, which we so ignore, has been studied previously by solving an evolution equation for $\\meanEMF$ (Blackman \\& Field 2002).  For homogeneous isotropic turbulence we may then write, analogously to (\\ref{meanEMF}), \\EQ \\meanEMF=\\hat{\\alpha}\\circ\\meanBB-\\hat{\\eta}_{\\rm t}\\circ\\meanJJ, \\label{meanEMFzz} \\EN or, in more explicit form, \\EQ \\meanEMF(\\xx)=\\int\\left[\\hat{\\alpha}(\\xi)\\meanBB(\\xx-\\xxi) -\\hat{\\eta}_{\\rm t}(\\xi)\\meanJJ(\\xx-\\xxi)\\right]\\,\\dd^3 \\xi \\label{meanEMFzzprime} \\EN with two functions $\\hat{\\alpha}$ and $\\hat{\\eta}_{\\rm t}$ of $\\xi = |\\xxi|$ that vanish for large $\\xi$. The integration is in general over all $\\xi$--space. Although $\\meanEMF$ and $\\meanBB$ as well as $\\hat{\\alpha}$ and $\\hat{\\eta}_{\\rm t}$ may depend on time the argument $t$ is dropped everywhere. For a detailed justification of the relations \\eq{meanEMFzz} and \\eq{meanEMFzzprime} we refer to \\App{justific}. In the limit of a weak dependence of $\\meanBB$ and $\\meanJJ$ on space coordinates, i.e.\\ when the variations of $\\meanBB (\\xx - \\xxi)$ and $\\meanJJ (\\xx - \\xxi)$ with $\\xxi$ are small in the range of $\\xi$ where $\\hat{\\alpha} (\\xi)$ and $\\hat{\\eta}_{\\rm t} (\\xi)$ are markedly different from zero, the relations \\eq{meanEMFzz} or \\eq{meanEMFzzprime} turn into \\eq{meanEMF}, and we see that \\mbox{$\\alpha = \\int \\hat{\\alpha} (\\xi) \\, \\dd^3 \\xi$} and $\\eta_{\\rm t} = \\int \\hat{\\eta}_{\\rm t} (\\xi) \\, \\dd^3 \\xi$.  At first glance the representations \\eq{meanEMFzz} and \\eq{meanEMFzzprime} of $\\meanEMF$ look rather different from \\eq{kernel}. Considering $\\meanJJ = \\nab \\times \\meanBB$ and carrying out an integration by parts we may however easily rewrite \\eq{meanEMFzzprime} into \\EQ \\meanemf_i (\\xx) = \\int K_{ij} (\\xxi) \\meanB_j (\\xx - \\xxi) \\, \\dd^3\\xi \\label{eq1} \\EN with \\EQ K_{ij} (\\xxi) = \\hat{\\alpha} (\\xi) \\delta_{ij} - \\frac{1}{\\xi} \\frac{\\partial \\hat{\\eta}_{\\rm t} (\\xi)}{\\partial \\xi} \\epsilon_{ijk} \\xi_k \\, . \\label{eq3} \\EN We further note that due to the symmetry of $\\hat{\\alpha} (\\xi)$ in $\\xxi$ only the part of $\\meanBB (\\xx - \\xxi)$ that is symmetric in $\\xxi$, i.e.\\ the part that can be described by $\\meanBB (\\xx)$ and its derivatives of even order, contributes to the $\\hat{\\alpha}$ terms in \\eq{meanEMFzzprime} or in \\eq{eq1} and \\eq{eq3}. The symmetry of $\\hat{\\eta}_{\\rm t} (\\xi)$ implies that only the part of $\\meanBB (\\xx - \\xxi)$ antisymmetric in $\\xxi$, which corresponds to the derivatives of $\\meanBB (\\xx)$ of odd order, contributes to the $\\hat{\\eta}_{\\rm t}$ terms.  Finally, referring to a Cartesian coordinate system $(x,y,z)$ we define mean fields by averaging over all $x$ and $y$ so that in particular $\\meanEMF$ and $\\meanBB$ depend only on $z$ and on time. Then \\eq{meanEMFzzprime} turns into \\EQ \\meanEMF(z) =\\int\\left[\\hat{\\alpha}(\\zeta)\\meanBB(z-\\zeta) -\\hat{\\eta}_{\\rm t}(\\zeta)\\meanJJ(z-\\zeta)\\right]\\,\\dd \\zeta \\, . \\label{meanEMFzzprime2} \\EN The functions $\\hat{\\alpha} (\\zeta)$ and $\\hat{\\eta}_{\\rm t} (\\zeta)$ are just averages of $\\hat{\\alpha} (\\xi)$ and $\\hat{\\eta}_{\\rm t} (\\xi)$ over all $\\xi_x$ and $\\xi_y$. They are therefore real and symmetric in $\\xi_z\\equiv\\zeta $. The integration in \\eq{meanEMFzzprime2} is in general over all $\\zeta$. The remark on the limit of weak dependences of $\\meanBB$ and $\\meanJJ$ on space coordinates made above in connection with \\eq{meanEMFzz} and \\eq{meanEMFzzprime} applies analogously to \\eq{meanEMFzzprime2}. We have now $\\alpha = \\int \\hat{\\alpha} (\\zeta) \\, \\dd \\zeta$ and $\\eta_{\\rm t} = \\int \\hat{\\eta}_{\\rm t} (\\zeta) \\, \\dd \\zeta$.  Relation \\eq{meanEMFzzprime2} can also be brought in a form analogous to \\eq{eq1} and \\eq{eq3}, \\EQ \\meanemf_i (z) = \\int K_{ij} (\\zeta) \\meanB_j (z - \\zeta) \\, \\dd \\zeta \\label{eq5} \\EN with \\EQ K_{ij} (\\zeta) = \\hat{\\alpha} (\\zeta) \\delta_{ij} - \\frac{\\partial \\hat{\\eta}_{\\rm t} (\\zeta)}{\\partial \\zeta} \\epsilon_{ij3} \\, . \\label{eq7} \\EN The remarks made under \\eq{eq1} and \\eq{eq3} apply, now due to the symmetries of $\\hat{\\alpha} (\\zeta)$ and $\\hat{\\eta}_{\\rm t} (\\zeta)$ in $\\zeta$, analogously to \\eq{meanEMFzzprime2}, \\eq{eq5} and \\eq{eq7}.  It is useful to consider in addition to \\eq{meanEMFzzprime2} also the corresponding Fourier representation. We define the Fourier transformation in this paper by $Q(z) = \\int \\tilde{Q}(k) \\exp(\\ii k z) \\, \\dd (k/2 \\pi)$. Then this representation reads \\EQ \\tilde{\\meanEMF}(k) = \\tilde{\\alpha}(k) \\tilde{\\meanBB}(k) -\\tilde{\\eta}_{\\rm t}(k) \\tilde{\\meanJJ}(k) \\, . \\label{meanEMFzzprime3} \\EN Both $\\tilde{\\alpha}(k)$ and $\\tilde{\\eta}_{\\rm t}(k)$ are real quantities, and they are symmetric in $k$. The limit of weak dependences of $\\meanBB$ and $\\meanJJ$ on $z$ corresponds here to $k \\to 0$, and we have $\\alpha = \\tilde{\\alpha} (0)$ and $\\eta_{\\rm t} = \\tilde{\\eta}_{\\rm t} (0)$. Detailed analytic expressions for $\\hat{\\alpha}(\\zeta)$ and $\\hat{\\eta}_{\\rm t}(\\zeta)$, or $\\tilde{\\alpha}(k)$ and $\\tilde{\\eta}_{\\rm t}(k)$, can be derived, e.g., from results presented in Krause \\& R\\\"adler (1980). A numerical determination of quantities corresponding to $\\hat{\\alpha} (\\zeta)$ and $\\hat{\\eta}_{\\rm t} (\\zeta)$ has been attempted by Brandenburg \\& Sokoloff (2002) for shear flow turbulence.  In this paper two specifications of the velocity field $\\uu$ will be considered. In the first case $\\uu$ is chosen such that it corresponds to a steady Roberts flow, which is periodic in $x$ and $y$ and independent of $z$. A mean--field theory of a magnetic field in fluid flows of this type, that are of course different from genuine turbulence, has been developed in the context of the Karlsruhe dynamo experiment (R\\\"adler et al.\\ 2002a,b, R\\\"adler \\& Brandenburg 2003). It turned out that the mean electromotive force $\\meanEMF$, except its $z$ component, satisfies relation \\eq{meanEMF} if any nonlocality in the above sense is ignored (see also \\App{robdyn}). Several analytical and numerical results are available for comparison with those of this paper. In the second case $\\uu$ is understood as homogeneous, isotropic, statistically steady turbulence, for which the above explanations apply immediately. Employing the method developed by Schrinner et al.\\ (2005, 2007) we will in both cases numerically calculate the functions $\\tilde{\\alpha}(k)$ and $\\tilde{\\eta}_{\\rm t}(k)$ as well as $\\hat{\\alpha}(\\zeta)$ and $\\hat{\\eta}_{\\rm t}(\\zeta)$. ",
        "conclusions": "The test field procedure turned out to be a robust method for determining turbulent transport coefficients (see Brandenburg 2005, Sur et al.\\ 2008 and Brandenburg et al.\\ 2008). The present paper shows that this also applies to the Fourier transforms of the integral kernels which occur in the nonlocal connection between mean electromotive force and mean magnetic field, in other words to the more general scale--dependent version of those transport coefficients. For isotropic turbulence the kernels $\\hat\\alpha$ and $\\hat\\eta_{\\rm t}$ have a dominant large-scale part and decline monotonously with increasing wavenumbers. This is consistent with earlier findings (cf.\\ Brandenburg \\& Sokoloff 2002), where however the functional form of the decline remained rather uncertain. Our present results suggest exponential kernels, corresponding to Lorentzian profiles in wavenumber space. The kernel for the turbulent magnetic diffusivity is about half as wide as that for the alpha effect. This result is somewhat unexpected and would be worthwhile to confirm before applying it to more refined mean field models. On the other hand, the effects of nonlocality become really important only when the scale of the magnetic field variations is comparable or smaller than the outer scale of the turbulence.  One of the areas where future research of nonlocal turbulent transport coefficients is warranted is thermal convection. Here the vertical length scale of the turbulent plumes is often comparable to the vertical extent of the domain. Earlier studies by Miesch et al.\\ (2000) on turbulent thermal convection confirmed that the transport of passive scalars is nonlocal, but it is also more advective than diffusive. It may therefore be important to also allow for nonlocality in time. This would make the expansion of passive scalar perturbations more wave-like, as was show by Brandenburg et al.\\ (2004) using forced turbulence simulations."
    },
    "0801/0801.4675_arXiv.txt": {
        "abstract": "{} {We  present a 0.2--200 keV broad-band study of absorbed AGN observed with \\textit{INTEGRAL}, \\textit{XMM-Newton}, \\textit{Chandra} and \\textit{ASCA} to investigate the continuum shape and the absorbing/reflecting medium properties.} {The sources are selected in the \\integral\\ AGN sample to have a 20--100 keV flux below 8$\\times$10$^{-11}$ $\\flux$ (5 mCrab), and are characterized by a 2--10 keV flux in the range (0.8--10)$\\times$10$^{-11}$ $\\flux$. The good statistics allow us a detailed study of the intrinsic and reflected continuum components. In particular, the analysis performed on the combined broad-band spectra allow us to investigate the presence of Compton reflection features and high energy cut-off in these objects.} {The column density of the absorbing gas establishes the Compton thin nature for three sources in which a measure of the absorption was still missing. The Compton thin nature of all the sources in this small sample is also confirmed by the diagnostic ratios F$_{\\rm x}$/F$\\left[\\rm O  III\\right] $.  The Compton reflection components we measure, reflection continuum and iron line,  are not immediately compatible with a scenario in which the absorbing and reflecting media are one and the same, i.e. the obscuring torus. A possible solution is that the absorption is more effective than reflection, e.g. under the hypothesis that the absorbing/reflecting medium is not uniform, like a clumpy torus, or that the source is observed through a torus with a very shallow opening angle. The high energy cut-off (a lower limit in two cases) is found in all sources of our sample and the range of values is in good agreement with that found in type 1 Seyfert galaxies. At lower energies there is clear evidence of a soft component (reproduced with a thermal and/or scattering model), in six objects.} {} ",
        "introduction": "The broad-band properties of Seyfert 2 AGN have been extensively studied in the past years mainly using \\sax\\ data (\\cite{risaliti02}). The continuum shape in these objects can be reproduced by a steep power-law with photon index $\\Gamma \\sim$ 2. This model matches well a Comptonization scenario in which seed photons from a cold gas (the accretion disc) are up-scattered in a hot corona of relativistic electrons (\\cite{HM91}, \\cite{HMG97}). A photoelectric cut-off is present at X-ray energies and its value depends on the absorbing column density of the source. For N$_H$ in the $10^{22}-10^{23}$ \\cmM2 range the cut-off is below 2 keV, for N$_H$ in  the $10^{23}-10^{24}$ \\cmM2 range  the cut-off is in 3--10 keV band and for N$_H >\\sigma_{T}^{-1}=1.5\\times 10^{24}$ \\cmM2 (i.e. the Compton thick regime with  $\\sigma_{T}^{-1}$ being the inverse of the Thomson cross section), the continuum can be observed only through its reflected component above 10 keV.  The presence of a high energy cut-off is still an open issue. Risaliti (2002)  found that the high energy cut-off is present only in $\\sim$ 30 per cent of a sample of 20 bright Sy 2 observed with \\sax. Above 10 keV the presence of a Compton reflection component has often been observed, however is its origin still unclear as well as if the reflecting/absorbing medium are one and the same (the putative ''molecular torus''). A cold iron line at 6.4 keV is often associated to this continuum Compton reflection hump.  The IBIS gamma-ray imager on board \\integral\\ is surveying the whole sky above 20 keV with a mCrab ($\\sim $10$^{-11}$ erg cm$^{-2}$ s$^{-1}$) sensitivity in well exposed regions. With its angular resolution of 12 arcmin and the point source location accuracy of 1-2 arcmin for moderately bright sources (Ubertini et al. 2003) \\integral\\ has already detected a number of AGN, a large fraction  of which were previously unknown as high energy emitters (\\cite{bird06}, \\cite{bassani06a}, \\cite{bird07}).  In this paper we present a deep broad-band spectral analysis of seven absorbed AGN detected by \\integral. The sample is extracted from the Bassani et. al (2006a) survey, updated to include a number of optical classifications obtained afterwards (see \\cite{bassani06b}). The sample includes Seyfert 2 galaxies having a 20--100 keV flux below 5 mCrab and with X-ray data available at the time that this work started; from the sample we have excluded sources already studied by \\sax\\ except for one object (\\eso) which was retained in order to provide a direct comparison between this work and previous studies using  \\sax\\ data  and therefore an \\textit{a posteriori} check of our analysis procedure which is affected by the limitation of using non simultaneous X and soft gamma-rays data. Overall we can conclude that the sample used in this study is representative of the population of type 2 AGN detected by \\integral\\ above 10 keV; further broad band X/gamma-ray analysis on a complete sample of type 2 AGN selected in the 20-40 keV band is on going and will be the objective of a future paper\\footnote{6 other type 2 AGN from Bassani et al. survey are below the assumed threshold flux: 3 (IGR J12415+5750, IGR J20286+2544 and IGR J17513-2011) had no X-ray available data when this analysis started while 3 (NGC 1068, Mkn 3 and NGC 6300) had BeppoSAX observational results published but all are peculiar objects  the first two being  Compton thick AGN and the third one the prototype of \"changing look\" Seyfert 2 ( i.e. a source that   goes from thick to thin and viceversa).}. Four sources (IGR J12391-1610=\\leda\\ hereafter, \\075, \\120 and \\j104) have been discovered first in soft gamma-ray band and  immediately after observed with \\chandra\\ (\\leda, \\075 and \\120) and \\xmm\\ (\\j104). Two objects (\\eso\\ and \\ic) are known AGN (although \\ic\\ was never reported at X-ray energies before) and follow up observations with \\xmm\\ have been used in our broad-band analysis. For one object (\\ngc) archival \\asca\\ data have been used in this work. Using X-ray measurements in conjunction with \\integral\\ data allow us to build a spectrum in three decades of energy, and to distinguish between the various spectral components characterizing the broad-band spectra of absorbed AGN. This in turn will allow us, in particular, to study the absorption/reflection medium properties of the sample. These good quality spectra also permit a detailed investigation of the intrinsic continuum slope including the cut-off if present. The broad--band coverage available is a powerful and unique tool to address all these issues, making possible a simultaneous measurement of the intrinsic continuum (both photon index and high energy cut-off) and of the reflected continuum together with the Fe line.  Note that a number of sources in the sample have data already published in the literature both in the X-ray band (Sazonov et al. 2005, Shinozaki et al. 2006) and in soft gamma-ray range (Beckmann et al. 2006, Molina et al. 2006). However, most of these works did not provide a detailed analysis of the broad band spectra of the sources in the sample nor attempted  to put constraints on the high energy cut-off and reflection bump. The only exception is \\eso\\ analysed by our team using \\sax\\ and \\integral\\ data; here we use a recent \\xmm\\ observation of the source in conjunction with an updated \\integral\\ spectrum  in order to have in the sample a well studied object to use as a source of reference.  In Sect. \\ref{data} we present the data of our small sample while the models used to fit the spectra are described in Sect. \\ref{models}. Results and discussion are presented in Sect. \\ref{RD} and draw our conclusions. The single cases are discussed in more details in the appendix. ",
        "conclusions": "\\label{conclusions}  We have presented spectral broad-band analysis of a small sample of seven absorbed Seyfert 2 selected at E$>$ 10 keV by \\integral/IBIS.  The combined 0.2--200 keV spectra allowed us to perform a detailed study of the emission/absorption properties and to test possible correlations between different parameters (e.g. R vs $\\Gamma$ or $E_{c}$ vs $\\Gamma$); in particular, for the first time in these sources a measurement of the reflection components (both Compton bump and iron line) and high energy cut-off has been performed. An \\textit{a posteriori} check of our analysis procedure, which is affected by the limitation of using non simultaneous X-ray soft gamma-ray data, is performed providing for one well known source, \\eso, a direct comparison between our study and previous ones with \\sax\\ data. The values of $\\Gamma$, R and N$_{H}$ we found for this source are completely consistent with those obtained by Akylas et al. (2001) analysing simultaneous broad--band \\sax\\ observations. This evidence strengthens the validity of our fitting procedure and makes  reliable our results. The main conclusions or our analysis are the followings.  \\begin{itemize} \\item The average value of the absorbing column density is $\\bar{N}_H$ = (10.6$^{+0.2}_{-0.2}$) $\\times$ 10$^{22}$ \\cmM2, and the range observed suggests a Compton thin nature for all objects in the sample, that is a new results in the case of \\ngc, \\ic\\ and \\j104. This evidence is also confirmed by the ratio F$_x$/F$\\left [ \\rm OIII\\right ]$ $\\lambda$5007.  \\item The continuum was reproduced with an e-folded power-law having a photon index in the range 1.3-1.9 and a high energy cut-offs well below 300 keV; this finding suggests that a cut-off is a common factor in Seyfert 2 galaxies with values in good agreement with those found in type 1 AGN.  \\item We measure the Compton reflection component in five sources (\\120, \\ngc, \\eso, \\ic\\ and \\j104) and found an upper limit in two others (\\leda\\ and \\075). No evidence for a correlation between R and $\\Gamma$ is found. We also observed the iron line (in the case of \\chandra\\ spectra we found upper limits only on the equivalent width). Both reflection components are not immediately consistent with production in the cold absorbing gas identified in the molecular torus. Both the reflection fraction R and the equivalent width of the line are too high to be produced in the gas with the observed column density. A possible solution is that the absorption is more effective than the reflection, e.g. making the hypothesis that the absorbing/reflecting medium is not uniform, like a clumpy torus, or that the source is observed through a torus with a very shallow opening angle.  \\item A soft excess component emerging below 2--3 keV is found in all sources with the only exception of \\075. This component is well reproduced by a thermal black body model having temperature in the range 0.2-0.9 keV, as typically observed in Sy 1. Alternatively a power-law with photon index equal to that of the illuminating continuum  is also able to reproduce this excess is some cases; the ratio of soft versus primary continuum is of the order of a few per cent, as typically observed in Sy 2. A different case is that of \\120 which has, in the thermal model, a temperature kT=1.7$^{+0.1}_{-0.2}$ keV and, in the scattered scenario, A$_{IC}$/A$_{soft}$=0.12; both value are higher than observed in the other sources and, typically, in Seyfert. This strongly suggests that the source is characterized by a larger soft excess (possibly due to a starburst contribution) that is not present in the other sources.   \\item \\075 is the only source in the sample that does not show any evidence for continuum excess at low energies. A narrow emission line around 2.5 keV is instead found in the \\chandra\\ data.  \\end{itemize}  We acknowledge the Italian Space Agency financial and programmatic support via contracts I/R/046/04 and I/023/05/0. We thank the referee for constructive suggestions."
    },
    "0801/0801.2786_arXiv.txt": {
        "abstract": "The observed 511 keV line from the Galactic Bulge is a real challenge for theoretical astrophysics: despite a lot of suggested mechanisms, there is still no convincing explanation and the origin of the annihilated positrons remains unknown. Here we discuss the possibility that a population of slowly evaporating primordial black holes with the mass around $10^{16}-10^{17}$~g ejects (among other particles) low--energy positrons into the Galaxy. In addition to positrons, we have also calculated the spectrum and number density of photons and neutrinos produced by such black holes and found that the photons are potentially observable in the near future, while the neutrino flux is too weak and below the terrestrial and extra--terrestrial backgrounds. Depending on their mass distribution, such black holes could make a small fraction or the whole cosmological dark matter. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.0260_arXiv.txt": {
        "abstract": "Magnetograms from the Vector SpectroMagnetograph (VSM) of the Synoptic Optical Long-term Investigations of the Sun (SOLIS) project are utilized to study the latitude distribution of magnetic flux elements as a function of latitude in the polar solar caps. We find that the density distribution of the magnetic flux normalized by the surface of the polar cap and averaged over months decreases close to the solar poles. This trend is more pronounced when considering only flux elements with relatively large size. The flux density of the latter is relatively flat from the edge of the polar cap up to latitudes of 70$^\\circ$--75$^\\circ$ and decreases significantly to the solar pole. The density of smaller flux features is more uniformly distributed although the decrease is still present but less pronounced. This result is important in studying meridional flows that bring the magnetic flux from lower to higher solar latitudes resulting in the solar cycle reversal. The results are also of importance in studying polar structures contributing to the fast solar wind, such as polar plumes. ",
        "introduction": "Polar regions of the Sun harbor numerous challenging solar phenomena, such as the source and acceleration process of the fast solar wind. Although it is widely believed that the magnetic field is main driver of most physical processes within the solar polar areas, these areas are not adequately characterized for observational and instrumental limitation reasons.  Several solar phenomena (i.e., solar differential rotation, supergranular diffusion and meridional flows) couple together to transport poleward the mid-latitude magnetic flux of decaying active regions. This process leads to the formation of polar caps and their evolution through the solar cycle (see Babcock \\& Babcock 1955; Leighton 1964; DeVore \\& Sheeley 1987; Sheeley, Nash, \\& Wang 1987; etc.). Solar meridional flows have been shown, both observationally and theoretically, to be the main mechanism of zonal flux transport (Howard 1974; Durrant, Turner, \\& Wilson 2004; Duvall 1979; LaBonte \\& Howard 1982). However, the weak flow, of few times 10~m~s$^{-1}$ at best, is not easily measurable and direct measurements can not be achieved beyond solar mid-latitudes ($\\sim45^\\circ$). Thus, it is important to obtain additional constraints on these flows close to the solar poles, which would be of great use for solar dynamo and flux transport models.  Raouafi, Harvey \\& Solanki (2006a,b; 2007) studied polar plumes EUV spectral emissions using different models. By comparing theoretical results to observations, they found that polar plumes would preferentially be based more than $10^\\circ$ away from the solar pole. Saito (1958) noticed that white-light plumes in eclipse observations were also rooted close to polar hole edges. The study of the polar flux latitude-distribution would confirm Raouafi et al.'s findings and constrain the magnetic flux transport (e.g., meridional flows) that drive such a distribution. ",
        "conclusions": "The high quality of magnetograms from SOLIS allowed us to characterize the flux concentration distribution as a function of latitude over a relatively large period of time. This was not possible earlier mainly because of instrumental limitations. The obtained results are important for studies of magnetic flux transport. They provide additional constraints on solar phenomena such as meridional circulation that is not possible to measure beyond $45^\\circ-50^\\circ$ nor how it functions near solar poles. Meridional flows, for instance, are an important input for solar dynamo and flux transport models. Our results on the density distribution of the magnetic flux concentrations at the polar regions suggests that the mechanisms responsible for the flux transport increasingly lose strength within the last 20 degree latitude before reaching the solar poles."
    },
    "0801/0801.2579_arXiv.txt": {
        "abstract": "For extrasolar planets with orbital periods, P$>$10 days, radial velocity surveys find non-circular orbital eccentricities are common, $\\langle e\\rangle\\sim 0.3$.  Future surveys for extrasolar planets using the transit technique will also have sensitivity to detect these longer period planets.  Orbital eccentricity affects the detection of extrasolar planets using the transit technique in two opposing ways: an enhancement in the probability for the planet to transit near pericenter and a reduction in the detectability of the transit due to a shorter transit duration.  For an eccentricity distribution matching the currently known extrasolar planets with P$>$10 day, the probability for the planet to transit is $\\sim 1.25$ times higher than the equivalent circular orbit and the average transit duration is $\\sim 0.88$ times shorter than the equivalent circular orbit.  These two opposing effects nearly cancel for an idealized field transit survey with independent photometric measurements that are dominated by Poisson noise.  The net effect is a modest $\\sim 4\\%$ increase in the transiting planet yield compared to assuming all planets have circular orbits.  When intrinsic variability of the star or correlated photometric measurements are the dominant source of noise, the transit detectability is independent of the transit duration.  In this case the transit yield is $\\sim$25\\% higher than that predicted under the assumption of circular orbits.  Since the Kepler search for Earth-sized planets in the habitable zone of a Solar-type star is limited by intrinsic variability, the Kepler mission is expected to have a $\\sim$25\\% higher planet yield than that predicted for circular orbits if the Earth-sized planets have an orbital eccentricity distribution similar to the currently known Jupiter-mass planets. ",
        "introduction": "The known extrasolar planets possess a broad distribution of orbital eccentricity \\citep{BUT06} (see Figure~\\ref{fig:ecchist}).  The short period, Hot Jupiter (P$<$10 day) planets predominately have circular orbits.  However, at longer orbital periods circular orbits become a minority and the median eccentricity for extrasolar planets e$\\sim$0.3.  At the extreme eccentricity end, there are three planets, HD 80606b \\citep{NAE01}, HD 20782b \\citep{JON06}, and HD 4113b \\citep{TAM07}, having e$>$0.9.  HD 80606b comes closer to its stellar host (a=0.033 AU) than many of the circular orbit Hot Jupiter planets. \\citet{FOR07} recently reviewed the various mechanisms invoked to explain the distribution of eccentricities for the known extrasolar planets. Interactions with a stellar companion, planetary companion, passing star, gaseous disk, planetesimal disk, and stellar jets have all been proposed to modify the orbital eccentricity of extrasolar planets.  Limited discussions in the literature have been given to the impact orbital eccentricity has on a transit survey for extrasolar planets. \\citet{TIN05} discuss the impact of eccentricity on their $\\eta$ parameter ($\\eta$ is the ratio of the observed transit duration to an estimate of the transit duration).  As expected, they find transits occur near pericenter (apocenter) are shorter (longer) in duration than the circular orbit case, and they show that the transit duration of an eccentric is typically shorter than a circular orbit of the same period.  However, their discussion was focused on the impact of orbital eccentricity on their $\\eta$ parameter. \\citet{MOU06} also discuss how the transit duration is affected by orbital eccentricity, but they do not quantify the impact this will have on transit surveys. Recently, \\citet{BAR07} derives the probability for a planet on an eccentric orbit to transit and conclude that the photometric precision of current surveys and future surveys, such as Kepler, is insufficient to determine the orbital eccentricity solely from the light curve. \\citet{BAR07} concludes that without knowledge of the eccentricity from radial velocity data or independent measurement of the stellar host radius, the habitability of planets detected with Kepler will remain unknown.  Neglecting the impact eccentricity has on transit detections is justified for the current sample of transiting planets given the predominance of circular orbits for the Hot Jupiters and the strong bias of transit surveys against finding long period planets on circular orbits \\citep{GAU05}.  The announced transit for the planet orbiting HD 17156\\footnote{First detected by the radial velocity technique \\citep{FIS07}} with a 21 day period and eccentric orbit \\citep{BARB07} is a precursor for the kinds of planets detectable in transit surveys.  As transit surveys continue, longer period transiting planets may be discovered.  More importantly, the recently launched COROT mission will surpass current surveys for sensitivity to longer period planets \\citep[P$\\sim$ several months; ][]{BORD03}. Also, the Kepler mission, scheduled for launch in 2009, has a goal to find Earth-sized objects at 1 AU from their host star \\citep{BORU04}. The main purpose of this paper is to show that given the distribution of eccentricities for the currently known extrasolar planets, eccentricity should not be ignored in assessing the detectability of transiting giant planets when the transit survey is sensitive to planets with P$>$ 10 day.  In addition to longer period planets, the COROT and Kepler missions also will detect transiting planets with small, Earth-sized radii. The eccentricity distribution for planets less massive than the currently known Jupiter-mass planets is beginning to be explored. \\citet{RIB07} and \\citet{FOR07} provide tentative evidence for the tendency of lower mass planets to have lower eccentricities in the current sample of radial velocity planets.  There is theoretical agreement that a proto-planetary gas disk strongly damps the eccentricity of non-gap opening embedded low-mass planets \\citep{GOL80, CRE07}.  However, the theoretical models show that once the gas disk dissipates, dynamical interactions amongst the planets results in a random-walk diffusion that leads to an increasing eccentricity that takes on a Rayleigh-distribution similar to what is observed \\citep[dotted line in Figure~\\ref{fig:ecchist}; ][]{JUR07,ZHO07}.  Opposing the increases in eccentricity from planet-planet scattering, late stage interactions with planetesimals can preferentially damp the eccentricities of lower mass planets enabling theoretical models to achieve the low eccentricities of the terrestrial planets of the Solar System \\citep{RAY06} and possibly explain the current trend of lower eccentricities for lower mass planets \\citep{FOR07}.  Given the number of potential physical processes that can affect orbital eccentricity, it is premature to assume that the typical Earth-like planet has an eccentricity near zero like the Solar System.  In this study, \\S~\\ref{sec:eccdist} reviews the observed eccentricity distribution of the known radial velocity extrasolar planets.  The broad distribution of orbital eccentricity has two main effects on the sensitivity of a transit survey.  First, the planet-host separation varies along an eccentric orbit, enhancing the probability to transit when the planet is relatively closer to the stellar host. \\S~\\ref{sec:tranprob} quantifies the net affect on the transit probability for a population of planets with non-circular orbits. Second, the planet velocity varies along and eccentric orbit, resulting in a reduction or lengthening of the transit duration. \\S~\\ref{sec:trandur} quantifies the distribution of transit durations resulting from a population of planets on eccentric orbits. \\S~\\ref{sec:disc} describes how the transit duration affects transit detection for transit surveys in the limit of various noise sources. \\S~\\ref{sec:conc} concludes by quantifying the net result of the two aforementioned effects, the enhanced probability to transit and the reduced detectability, on the yield from transit surveys. ",
        "conclusions": "\\label{sec:conc}  Orbital eccentricity results in an enhanced probability for a planet to transit and potentially a reduction in the transit detectability. The overall yield from a transit survey is given by $N_{\\rm det}\\propto{\\rm Prob}_{{\\rm T}}\\times N_{\\rm obj}$.  The results from this study can be used to scale the overall yield from a transit survey based on assuming all planets are on circular orbits for an assumed distribution of orbital eccentricity.  The enhanced probability for a planet to transit, ${\\rm Prob}_{{\\rm T}}$ with a distribution of orbital eccentricity scaled to the circular orbit case is given by Equation~\\ref{eq:probtrane}.  The reduced number of transiting planets detectable $N_{\\rm obj}$ scaled to the circular orbit case for an ideal transit survey where the photometric noise is white and dominated by Poisson error is given in \\S~\\ref{sec:disc}. Multiplying these two factors provides the overall yield of an ideal transit survey scaled to assuming all planets are on circular orbits. Figure~\\ref{fig:tranyield} show that these two opposing effects nearly cancel over the parameters that characterize the eccentricity distribution.  Thus, for an idealized transit survey with the currently observed orbital eccentricity distribution, the overall yield will be 4\\% greater than assuming all planets are on circular orbits.  The result for an idealized transit survey with a Rayleigh distribution of orbital eccentricities gives a similar enhancement (4\\%).   \\begin{figure} \\includegraphics[scale=0.9,viewport=0 100 350 350]{f6.ps} \\caption{The overall yield from an idealized transit survey as function of the observe eccentricity distribution model parameters scaled to the assumption that all planets have circular orbits.  The transit survey is assumed to have independent photometric measurements that are dominated by Poisson noise.  The abscissa indicates $e_{\\rm max}$, and the curves are for selected values of $e_{\\rm crit}$ as labeled.  For the model parameters shown in Figure~\\ref{fig:ecchist}, $N_{\\rm det}\\sim 4\\%$ higher than assuming all planets have circular orbits ({\\it dotted line}).\\label{fig:tranyield}} \\end{figure}   In ground-based transit surveys, correlated measurements limit transit detectability \\citep{PON06} and intrinsic variability of the star will limit the detectability for Earth-sized planets for the Kepler mission \\citep{JEN02}.  In both cases the transit detectability is independent of the transit duration, in which case $N_{\\rm obj}$ is independent of the orbital eccentricity.  For these cases, the transit survey will have higher returns by a factor of $\\langle {\\rm Prob}_{{\\rm T}e}\\rangle$ (Equation~\\ref{eq:probtrane}) than estimated by assuming all planets are on circular orbits.  However, the reduced planet yield in a transit survey due to intrinsic stellar variability or correlated measurements must be properly accounted for.  If every dwarf star has an Earth-sized planet orbiting in the habitable zone, then assuming circular orbits, the Kepler mission expects to detect 100 Earth-sized planets in the habitable zone \\citep{BORU04}.  The work presented here indicates Kepler will have $\\langle {\\rm Prob}_{{\\rm T}e}\\rangle\\sim 25\\%$ higher yield if Earth-sized planets in the habitable zone have a planet eccentricity distribution similar to the currently known sample of giant planets from radial velocity surveys.  However, if Earth-sized planets with e$\\sim$0.9 are as common as e$\\sim$0 then, the yield of Earth-sized planets could be 80\\% higher from the Kepler mission assuming the high-e, short transit duration planets are still detectable.  The dependence of the transit survey yield on the uncertain underlying orbital eccentricity distribution implies an uncertainty in measuring the frequency of terrestrial planets in the habitable zone \\citep[a major goal of the Kepler mission][]{BORU04} .  An analysis of the transit yield from a transit survey that assumes all planets are on circular orbits will overestimate the frequency of habitable planets if high eccentricities are common and not taken into account. In practice a variety of noise regimes affect a transit survey and accurate yields necessitate an accurate understanding of the photometric noise, stellar sample, and underlying eccentricity distribution \\citep{BUR06,GOU06,FRE07}."
    },
    "0801/0801.0110_arXiv.txt": {
        "abstract": "We explore the abundance of light clusters in core-collapse supernovae at post-bounce stage in a quantum statistical approach. Adopting the profile of a supernova core from detailed numerical simulations, we study the distribution of light bound clusters up to alpha particles (2 $\\leq A \\leq$ 4) as well as heavy nuclei ($A > 4$) in dense matter at finite temperature. Within the frame of a cluster-mean field approach, the abundances of light clusters are evaluated accounting for self-energy, Pauli blocking and effects of continuum correlations. We find that deuterons and tritons, in addition to $^3$He and $^4$He, appear abundantly in a wide region from the surface of the proto-neutron star to the position of the shock wave. The appearance of light clusters may modify the neutrino emission in the cooling region and the neutrino absorption in the heating region, and, thereby, influence the supernova mechanism. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.3219_arXiv.txt": {
        "abstract": "{Amino acids are building blocks of proteins and therefore key ingredients for the origin of life. The simplest amino acid, glycine (NH$_2$CH$_2$COOH), has long been searched for in the interstellar medium but has not been unambiguously detected so far. At the same time, more and more complex molecules have been newly found toward the prolific Galactic center source Sagittarius B2.} {Since the search for glycine has turned out to be extremely difficult, we aimed at detecting a chemically related species (possibly a direct precursor), amino acetonitrile (NH$_2$CH$_2$CN).} {With the IRAM 30m telescope we carried out a complete line survey of the hot core regions Sgr~B2(N) and (M) in the  3\\,mm range, plus partial surveys at 2 and 1.3\\,mm. We analyzed our 30m line survey in the LTE approximation and modeled the emission of all known molecules simultaneously. We identified spectral features at the frequencies predicted for amino acetonitrile lines having intensities compatible with a unique rotation temperature. We also used the Very Large Array to look for cold, extended emission from amino acetonitrile.} {We detected amino acetonitrile in Sgr~B2(N) in our 30m telescope line survey and conducted confirmatory observations of selected lines with the IRAM Plateau de Bure and the Australia Telescope Compact Array interferometers. The emission arises from a known hot core, the Large Molecule Heimat, and is compact with a source diameter of 2$\\arcsec$ (0.08 pc). We derived a column density of $2.8 \\times 10^{16}$ cm$^{-2}$, a temperature of 100 K, and a linewidth of 7 km~s$^{-1}$. Based on the simultaneously observed continuum emission, we calculated a density of $1.7 \\times 10^8$ cm$^{-3}$, a mass of 2340 M$_\\odot$, and an amino acetonitrile fractional abundance of $2.2 \\times 10^{-9}$. The high abundance and temperature may indicate that amino acetonitrile is formed by grain surface chemistry. We did not detect any hot, compact amino acetonitrile emission toward Sgr~B2(M) or any cold, extended emission toward Sgr~B2, with column-density upper limits of $6 \\times 10^{15}$ and $3 \\times 10^{12-14}$ cm$^{-2}$, respectively.} {Based on our amino acetonitrile detection toward Sgr~B2(N) and a comparison to the pair methylcyanide/acetic acid both detected in this source, we suggest that the column density of both glycine conformers in Sgr~B2(N) is well below the best upper limits published recently by other authors, and probably below the confusion limit in the 1-3\\,mm range.} ",
        "introduction": "\\label{s:intro}  Among the still growing list of complex molecules found in the interstellar medium, so-called ``bio''molecules garner special attention. In particular, the quest for interstellar amino acids, building blocks of proteins, has engaged radio and millimeter wavelength astronomers for a long time. Numerous published and unpublished searches have been made for interstellar glycine, the simplest amino acid \\citep{Brown79,Hollis80,Berulis85,Combes96,Ceccarelli00,Hollis03,Jones07,Cunningham07}. Its recent ``detection'' claimed by \\citet{Kuan03} has been persuasively rebutted by \\citet{Snyder05}. Since the early days of molecular radio astronomy, Sagittarius B2 has been a favorite target in searches for complex molecules in space.  \\subsection{The target: Sagittarius B2} \\label{ss:sgrb2intro} \\object{Sagittarius B2} (hereafter Sgr~B2 for short) is a very massive (several million solar masses) and extremely active region of high-mass star formation at a projected distance of $\\sim 100$ pc from the Galactic center. Its distance from the Sun is assumed to be the same as the Galactic center distance, $R_0$. \\citet{Reid93}, reviewing various methods to determine $R_0$, arrived at a ``best estimate'' of 8.0 $\\pm$ 0.5 kpc, a value that we adopt in this article. It is supported by recent modeling of trajectories of stars orbiting the central black hole, which yields 7.94 $\\pm$ 0.42 kpc \\citep{Eisenhauer03}.  There are two major centers of activity, \\object{Sgr~B2(M)} and \\object{Sgr~B2(N)} separated by $\\sim 2$ pc. In each of them, recent star formation manifests itself in a multitude of H{\\sc ii} regions of many sizes, from hypercompact to compact \\citep[][]{Gaume95}, and there is abundant material to form new stars evident by massive sources of molecular line and submillimeter continuum emission from dust \\citep{Lis91, Lis93}.  \\subsubsection{Sgr B2 as part of the Central Molecular Zone} \\label{sss:sgrb2cmz} Some of the first detections of interstellar organic molecules (at cm-wavelengths!) were made toward Sgr~B2 \\citep[see][ for a historical perspective]{Menten04}. The low intrinsic line strengths make these cm lines unlikely candidates for detection. However, the situation is helped, first, by the fact that many of the transitions in question may have inverted levels \\citep{Menten04} and amplify background radio continuum emission which is very intense at cm wavelengths \\citep{Hollis07}. Second, the spatial distributions of many species are characterized by spatially extended emission covering areas beyond Sgr~B2 itself, filling single dish telescope beams, thus producing appreciable intensity even when observed with low spatial resolution \\citep[][]{Cummins86,Jones08}. This emission is characterized by low rotation temperatures, favoring lower frequency lines. Recent identifications of  ``new'' species include glycolaldehyde CH$_2$OHCHO \\citep[][]{Hollis00,Hollis01}, ethylene glycol HOCH$_2$CH$_2$OH \\citep[][]{Hollis02}, and vinyl alcohol CH$_2$CHOH \\citep[][]{Turner01}.  Sgr B2 and its surroundings are part of the Central Molecular Zone (CMZ) of our Galaxy, a $\\sim \\pm 0\\fdg3$ latitude wide band stretching around the Galactic center from longitude $l \\sim +1\\fdg6$ to $-1\\fdg1$ \\citep[see, e.g.,][]{Morris96}. The CMZ contains spatially extended emission of many complex organic molecules \\citep{Minh92,Dahmen97,Menten04,RequenaTorres06}.  \\subsubsection{The Large Molecule Heimat} \\label{sss:introlmh} Near Sgr~B2(N), there is a hot, dense compact source  that has a mm-wavelength line density second to no other known object. This source, for which \\citet{Snyder94} coined the name ``Large Molecule Heimat'' (LMH), is characterized by very high densities ($> 10^7$~cm$^{-3}$) and gas temperatures ($> 100$~K). In recent years arcsecond resolution interferometry with the BIMA array  has resulted in the detection and imaging of increasingly complex organic species toward the LMH, such as vinyl cyanide CH$_2$CHCN, methyl formate HCOOCH$_3$, and ethyl cyanide CH$_3$CH$_2$CN \\citep[][]{Miao95,Miao97}, formamide NH$_2$CHO, isocyanic acid HNCO, and methyl formate HCOOCH$_3$ \\citep[][]{Kuan96a}, acetic acid CH$_3$COOH \\citep[][]{Mehringer97b,Remijan02}, formic acid HCOOH \\citep[][]{Liu01}, and acetone (CH$_3$)$_2$CO \\citep[][]{Snyder02}. All the interferometric observations are consistent with a compact ($< $few arcsec diameter) source that had already been identified as the source of high-density-tracing non-metastable ammonia line emission by \\citet{Vogel87} and thermal methanol emission by \\citet[][, their source ``i'']{Mehringer97a}. The LMH also hosts a powerful H$_2$O maser region \\citep[][]{Reid88}, which provides evidence that it is very young (see Sect.~\\ref{ss:sgrb2n}).  \\subsection{The complex spectra of complex molecules} Complex molecules in general have large partition functions, in particular for the elevated temperatures ($> 100$ K) in molecular hot cores, dense and compact cloud condensations internally heated by a deeply embedded, young high-mass (proto)stellar object. Therefore, most individual spectral lines are weak and might easily get hidden in the ``line forest'' found toward these frequently extremely line-rich sources. To a large part, this forest consists of rotational lines, many of them presently unidentifiable, from within relatively low-lying vibrational states of molecules. Most of the candidate molecules from which these lines originate are known to exist in these sources, but laboratory spectroscopy is presently lacking for lines from the states in question. At this point in the game, unequivocally identifying a species in a spectrum of a hot core covering a wide spectral range requires the following steps: as described in detail in Sect.~\\ref{ss:modeling30m}, assuming Local Thermodynamic Equilibrium (LTE) (which applies at the high densities in hot cores) a  model spectrum is calculated for an assumed rotation temperature, column density, line width and other parameters. This predicts lines of a given intensity at all the known frequencies. Then at least two conditions have to be fulfilled: (i) All predicted lines should have a counterpart in the observed spectrum with the right intensity and width -- no single line should be missing. (ii) Follow-up observations with interferometers have to prove whether all lines from the candidate species are emitted from the same spot. Given the chemical variety in hot core regions, this is a powerful constraint. Moreover, interferometer images tend to have less line confusion, since many lines that are blended in larger beam single-dish spectra arise from different locations or are emitted by an extended region that is spatially filtered out. Using an interferometer for aiding molecule identifications was pioneered by L. Snyder and collaborators who (mostly) used the Berkeley-Illinois-Maryland-Array (BIMA) to clearly identify a number of species  in the Sgr~B2(N) Large Molecule Heimat (see Sect.~\\ref{sss:introlmh}).  We carried out a complete line survey of the hot core regions Sgr~B2(N) and (M) with the IRAM 30m telescope at 3\\,mm, along with partial surveys at 2 and 1.3\\,mm. One of the overall goals of our survey was to better characterize the molecular content of both regions. It also allows searches for ``new'' species once we have identified the lines emitted by known molecules (including vibrationally and torsionally excited states). In particular, many complex molecules have enough lines in the covered frequency ranges to apply criterion (i) above. Once a species fulfils this criterion, interferometer measurements of selected lines can be made to check criterion (ii).  \\subsection{Amino acetonitrile} One of our target molecules was amino acetonitrile (NH$_2$CH$_2$CN), a molecule chemically related to glycine. Whether it is a precursor to the latter is under debate (see Sect.~\\ref{ss:precursor}). Not many astronomical searches for amino acetonitrile have been reported in the literature. In his dissertation, \\citet[][]{Storey76} reported searches for the $J_{K_a,K_c} = 2_{11}-2_{12}$ and $1_{01}-0_{00}$ transitions at 1350.5 and 9071.7 MHz, respectively with the Parkes 64\\,m telescope. On afterthought, the only chance of success for their observations would have been if amino acetonitrile existed on large spatial scales, similar to the molecules described in Sect.~\\ref{sss:sgrb2cmz} (see Sect.~\\ref{ss:aanvla} for further limits on extended amino acetonitrile emission). Recently, \\citet{Wirstroem07} reported unsuccessful searches of a number of mm-wavelength transitions of amino acetonitrile toward a number of hot cores.  Here, we report our detection of warm compact emission from amino acetonitrile in Sgr~B2(N) with the IRAM 30m telescope, the Plateau de Bure Interferometer (PdBI) and the Australia Telescope Compact Array (ATCA), and upper limits on cold, spatially extended emission from amino acetonitrile that we obtained with the NRAO Very Large Array (VLA). Section~\\ref{s:obs} summarizes the observational details. We present our results in Sect.~\\ref{s:results}. Implications in terms of interstellar chemistry are discussed in Sect.~\\ref{s:discussion}. Our conclusions are summarized in Sect.~\\ref{s:concl}. ",
        "conclusions": ""
    },
    "0801/0801.3733_arXiv.txt": {
        "abstract": "We aim to understand cloud formation in substellar objects.  We combined the non-equilibrium, stationary cloud model of Helling, Woitke \\& Thi (2008; seed formation, growth, evaporation, gravitational settling, element conservation) with the general-purpose model atmosphere code {{\\sc Phoenix}} (radiative transfer, hydrostatic equilibrium, mixing length theory, chemical equilibrium) in order to consistently calculate cloud formation and radiative transfer with their feedback on convection and gas phase depletion.  We calculate the complete 1D model atmosphere structure and the chemical details of the cloud layers. The {{\\sc Drift-Phoenix}} models enable the first stellar atmosphere simulation that is based on the actual cloud {{\\it formation}} process. The resulting $(T,p)$ profiles differ considerably from the previous limiting {{\\sc Phoenix}} cases {{\\sc Dusty}} and {{\\sc Cond}}. A tentative comparison with observations demonstrates that the determination of effective temperatures based on simple cloud models has to be applied with care.  Based on our new models, we suggest a mean T$_{\\rm eff}=1800$K for the L\\,-\\,dwarf twin-binary system \\objectname{DENIS J0205-1159} which is up to 500K hotter than suggested in the literature. We show transition spectra for gas-giant planets which form dust clouds in their atmospheres and evaluate photometric fluxes for a \\objectname{WASP-1} type system. ",
        "introduction": "Today's most efficient tools to interpret the observed spectra of substellar objects, i.e., brown dwarfs and planets, are 1D static atmosphere simulations. Comparisons with observations, based on the solution of the radiative-transfer problem which is done in great detail with respect to the gas-phase opacities (Tsuji 2002, 2005; Allard et al. 2001, Ackerman \\& Marley 2001, Burrows \\& Sharp 1999, Barman et al. 2005), ideally yield insight into the atmospheric structure and chemistry providing finger prints of its evolutionary state.  Brown dwarf and planetary atmospheres have a far more complex chemistry than stellar objects due to the formation of clouds, which bind chemical elements and, hence, strongly influences the remnant gas phase inside the atmosphere. The presence of such clouds was previously simplified in static model atmosphere codes.  We now move a significant step forward by kinetically treating the chemistry of cloud formation as a phase-non-equilibrium process in the framework of model atmosphere simulations. Our dust model has been studied so far for a given $(T, p, v_{\\rm conv})$ structure ($T$ - gas temperature, $p$ - gas pressure, $v_{\\rm conv}$ - convective velocity), and we present in this letter the first consistent simulation of cloud micro-physics and atmospheric structure which allows for the first time to study the feedback of the dust {\\it formation} onto the atmospheric structure. We present our first consistent {\\sc Drift-Phoenix} results together with a tentative comparison to the observed spectrum of \\objectname{ DENIS J0205--1159} and point out possible uncertainties in present T$_{\\rm eff}$ determinations (Sect.~\\ref{s:appl}).  We, however, leave aside the issue of hydrodynamical cloud formation (e.g. Showman et al. 2006, Knutson et al. 2007, Rauscher et al. 2007) which unavoidably has to deal with the turbulent closure problem (Helling et al. 2004, Helling 2007). We calculate transition spectra for gas-giant planets and a {\\sc WASP}-1 type star, and evaluate the photometric fluxes for the 2MASS system, for the VISIR system, and for the IRAC band systems. ",
        "conclusions": "The modelling of cloud formation in substellar atmospheres has a strong impact on the objects temperature-pressure structure which in turn determines the cloud's chemical composition, the grain size distribution function, and the cloud's location inside the atmosphere. These are results of our consistent simulation of substellar atmospheres and detailed non-equilibrium dust cloud formation.  We demonstrate synthetic transition spectra for gas-giant planets and calculate synthetic photometric fluxes based on our consistent solution of dust cloud formation and radiative transfer problem.  A future goal is to extend our model to solar-system-like planets with much cooler atmosphere.  As we have shown for the field object {\\sc Denis} J0205-1159, stellar parameter determinations based on a comparison with synthetic spectra can vary considerable. As one consequence of this, we have set out to conduct a component-based study \\footnote{http://phoenix.hs.uni-hamburg.de/BrownDwarfsToPlanets1/ } where e.g. cloud compositions, dust-to-gas-ratios, and grain sizes are compared for different cloud models (Helling et al. 2007, 2008)."
    },
    "0801/0801.4325_arXiv.txt": {
        "abstract": "We analyze the hitherto available space-based X-ray data as well as ground-based optical data of the X-ray transient 080109/SN\\,2008D. From the data we suggest that ({\\it i}) The initial transient ($\\lesssim 800$ sec) is attributed to the reverse shock emission of a mildly relativistic ($\\Gamma \\sim$ a few) outflow stalled by the dense stellar wind. ({\\it ii}) The subsequent X-ray afterglow ($\\lesssim 2\\times 10^4$ sec) can be ascribed to the forward shock emission of the outflow, with a kinetic energy $\\sim 10^{46}$ erg, when sweeping up the stellar wind medium. ({\\it iii}) The late X-ray flattening ($\\gtrsim 2\\times 10^4$ sec) is powered by the fastest non-decelerated component of SN\\,2008D's ejecta. % ({\\it iv}) The local event rate of X-ray transient has a lower limit of $\\sim 1.6\\times 10^4~{\\rm yr^{-1}~Gpc^{-3}}$, indicating a vast majority of X-ray transients have a wide opening angle of $\\gtrsim 100^\\circ$. ({\\it v}) Transient 080109/SN\\,2008D indicates a continuum from GRB-SN to under-luminous GRB-/XRF-SN to X-ray transient-SN and to ordinary Ibc SN (if not every Ibc SN has a relativistic jet), as shown in Figure 2 of this {\\it Letter}. ",
        "introduction": "\\label{sec:Into} During the past decade, long-duration ($\\gtrsim 2$sec) $\\gamma-$ray bursts (GRBs), including the subclass of X-ray flashes (XRFs), have been found (1) to be driven by the core-collapse of massive stars \\cite{Woosley93}; thus (2) to be associated with a rare variety ($\\sim1\\%$) of type Ibc supernovae (SNe), the so-called hypernovae (HN) \\cite{Galama98,Hjorth03,Stanek03,Male04,Sollerman06,Camp06} (but also see Fynbo et al. 2006); and (3) in general to be hosted by the star-forming dwarf galaxies with low metallicity \\cite{Fyn03,Fru06,Stanek06}. Though the association of GRB/XRF and Ibc SN has been pinned down, what channels make a dying star to produce a GRB or an XRF, and not just a Ibc SN, is still unclear. The progenitor's mass, metallicity, angular momentum, and the configuration and strength of its internal magnetic field play important roles for the generation of GRBs/XRFs and ordinary Ibc SNe.  The serendipitous discovery of the X-ray transient 080109/SN\\,2008D may shed light on filling in this gap between energetic GRBs/XRFs and ordinary Ibc SNe. We will analyze space- and ground-based data of this transient and SN, focusing on X-ray/radio data because observationally they trace the fastest component of the transient/SN outflow while optical data trace the slower SN ejecta (e.g., Soderberg et al. 2006). ",
        "conclusions": "Transient 080109/SN\\,2008D presents the first evidence for a mild-relativistic outburst, $\\sim 10^{46}\\,{\\rm erg}$, preceding the main SN component, thus confirming previous speculation in SN\\,2005bf (Folatelli et al. 2006).  It sets a lower limit of the local event rate of its kind as $1/3{\\rm{yr}}/({\\rm{0}}{\\rm{.0276 Gpc}})^3 \\sim 1.6 \\times 10^4 {\\rm{yr}}^{{\\rm{ - 1}}} {\\rm{Gpc}}^{{\\rm{ - 3}}}$, comparable with the local rate of Ibc SNe, $\\sim 4.8 \\times 10^4 {\\rm{yr}}^{{\\rm{ - 1}}} {\\rm{Gpc}}^{{\\rm{ - 3}}}$, and thus indicates a vast majority of X-ray transients have a wide opening angle of $\\gtrsim 100^\\circ$. The collimation-corrected energy is of $\\sim 5\\times 10^{45}$ erg. The wide angle budget, together with the self-consistent interpretation of the transient and its early afterglow with the on-beam model, largely rules out the off-axis viewing model for this transient.  The host NGC2770, a spiral galaxy with copious ${\\rm H\\alpha}$ sign, stands out from the star-forming dwarf galaxies typically hosting GRBs/XRFs. Transient 080109 puts itself on the upper border of the nearby GRB/XRF collection in terms of the host metallicity (Berger \\& Soderberg 2008b; Sollerman et al. 2005).  As a result, this event may unveil a continuum from energetic GRB (top-right of Figure 2) to ordinary Ibc SN (bottom-left of Figure 2).  (1) Whether or not every Ibc SN has a quasi-jet outburst proceeding the main SN component is still uncertain even the discovery of Transient 080109. For this reason we mark the ordinary Ibc SN and X-ray transient/SN populations with a dash ellipse in Figure 2.  (2) Soderberg et al. (2006) showed that producing GRBs/XRFs needs a relativistic ejecta carrying at least $10^{48}$ erg. We show in this {\\it Letter} that X-ray transient population couples $\\sim 10^{46}$ erg to relativistic material regarding this found one marks the transition between GRB/XRF and ordinary Ibc.  (3) While under-luminous GRBs/XRFs are likely powered by moderate-relativistic material, X-ray transients are likely powered by mild-relativistic material.  (4) GRBs have an average opening angle of $\\sim \\ 10^\\circ$ while a vast majority (if not all) of X-ray transients have a much wider one of $\\sim 100^\\circ$. There is a negative correlation between radiated energy and opening angle from GRB to XRF to X-ray transient.  (5) Materials with higher bulk Lorentz factor tend to have a shallower energy-velocity distribution leading to spikeful behavior as shown in GRB/XRF prompt lightcurves and hypernova's broad-lined spectra. Materials with lower bulk Lorentz factor tend to have a steeper energy-velocity distribution and thus largely couple with each other leading to spikeless/little-spiked behavior as shown in various optical afterglows. The decay laws in terms of velocity for each event in Figure 2 (from left to right) matches this principle."
    },
    "0801/0801.4864_arXiv.txt": {
        "abstract": "We studied the radio structure of high-redshift ($z>3$) quasars with VSOP at 1.6 and 5~GHz. These sources are the most distant objects ever observed with Space VLBI, at rest-frame frequencies up to $\\sim25$~GHz. Here we give an account of the observations and briefly highlight the most interesting cases and results. These observations allowed us, among other things, to estimate the mass of the central black holes powering these quasars, to identify large misalignments between the milli-arcsecond (mas) and sub-mas scale radio structures, and to detect apparent superluminal motion at sub-mas scale. ",
        "introduction": "Twenty of the most distant ($z>3$) radio quasars have been proposed for Space Very Long Baseline Interferometry (SVLBI) observations at either 1.6 or 5~GHz or both in the VLBI Space Observatory Programme \\citep[VSOP;][]{hira00}. These were the brightest known, with total flux densities higher than $\\sim400$~mJy. Our goal was to study the structural properties of high-redshift sources at the highest possible resolution at emitted frequencies up to $\\sim20-25$~GHz. This allowed us to investigate source compactness, relativistic jet propagation, and eventually to estimate physical parameters of the jets and the active nuclei feeding them. ",
        "conclusions": ""
    },
    "0801/0801.3169_arXiv.txt": {
        "abstract": "We propose a model of allocating galaxies in cosmological N-body simulations. We identify each subhalo with a galaxy, and assign luminosity and morphological type assuming that the galaxy luminosity is a monotonic function of its host subhalo mass. The morphology assignment is made by using two simple relations between subhalo mass and galaxy luminosity of different types. One is using a constant ratio in luminosity of early (E/SO) and late (S/Irr) type galaxies at a fixed subhalo mass. And the other assumes that galaxies of different morphological types but having an equal luminosity have a constant ratio in their subhalo masses. We made a series of comparisons of the properties of these simulated galaxies with those of the SDSS galaxies. The resulting simulated galaxy sample is found to successfully reproduce the observed local number density distribution except for in high density regions. The luminosity function is studied as a function of local density. It was found that the observed luminosity functions in different local density environments are overall well-reproduced by the simulated galaxies. Discrepancy is found at the bright end of the luminosity function of early types in the underdense regions and at the faint end of both morphological types in very high density regions. A significant fraction of the observed early type galaxies in voids seems to have undergone a relatively recent star formation and became brighter. The lack of faint simulated galaxies in dense regions may be due to the strong tidal force of the central halo which destroys less massive satellite subhalos around in the simulation. The mass-to-light ratio is found to depend on the local density in the way similar to that observed in the SDSS sample. We have found an impressive agreement between out simulated galaxies and the SDSS galaxies in the dependence of the central velocity dispersion on the local density and luminosity. ",
        "introduction": "The current galaxy formation paradigm can be characterized by ``hierarchical clustering''.  This means that massive dark matter halos form by merging less massive halos and/or by accreting ambient matter, and also that a dark matter halo governs the evolution of the galaxy residing inside. Most galaxies are believed to be hosted by the dark matter halos becuase the halos can provide a deep potential well for baryonic matter to condense and cool down. This leads to triggering of the star formation; gas sufficiently accumulated in the halo potential center begins to experience many hydrodynamic processes, such as radiative cooling, star-formation, supernova explosion, and chemical enrichments. All of these processes play an important role in making the visible galaxies in the end. Because galaxies as a building block of the large scale structures consequently follows the evolution of their host halos over the cosmic history, understanding gravitational evolution of dark halos is very important for the study of galaxy formation and evolution.  Over the past few decades, cosmological simulations have been proved useful for the study of structure formation. Simulations have boosted up many investigations of the nonlinear structure evolutions and their results have been extensively compared to the observations in various aspects of interest. Many studies have reported successful recovery of observational features of the galaxy distribution like the two-point correlations of galaxies \\citep{conroy06,kravtsov04,berlind02}, topology \\citep{park05a,park05b,gott06}, and the environmental dependence of spin distributions \\citep{cervantes-sodi07}.   Detailed modeling of galaxy evolution and understanding of the galaxy properties have been possible with the recent advent of the huge redshift surveys such as Sloan Digital Sky Survey\\footnote{http://www.sdss.org} (SDSS) and 2dFGRS\\footnote{http://www.aao.gov.au/2df/}.  These larger surveys provide a rich information on the formation and evolution of galaxy. To meet a requirement to establish a relation between simulated structures and observed galaxies many techniques have been introduced. The semi-analytic model(SAM) of galaxy formation \\citep{cole94,kauffmann97,baugh06} is based on the hierarchical clustering of halos whose merging history trees are built by generating a set of Gaussian random numbers. The mass growth history of halos on various mass scales can be generated and traced by this method. Numerous SAM parameters are implemented in the merging history to reflect the hydrodynamical processes of heating, cooling, star formation, and aging the stellar populations. Their parameter values are fine tunned for the best description of collective properties of observed galaxies. But as the number of observational contraints increases, the number of parameters increases and the SAM becomes much complicated.  Hydrosimulations employ direct formulations of the hydrodynamic processes (\\citealt{weinberg06}). Two types of hydro simulations are now widely adopted; the Lagrangian \\citep{monaghan92,hernquist89} and Eulerian methods (\\citealt{tasker06,harten97}, and for comparative study, see \\citealt{heitmann05,thacker00}). They set gas particles or gas grids in the system of interest. Then, gas dynamics are taken into account by calculating the hydrodynamic interactions between neighbor particles or by solving differential hydro equations between adjacent grids. This method has several weak points; it suffers from the lack of resolutions in space and mass as the N-body simulations. And in most cases, the star formation and supernova explosions are far beyond the simulation resolution. Some processes such as radiative transfer, star formation, supernovae feedback, and initial stellar mass function, are not thoroughly understood and are difficult to parameterize in a well established way.  One of the offsprings of the SAM is the Halo Occupation Distribution (HOD; \\citealt{zheng05,seljak00,berlind02}) model. It relates the Friend-of-Friend (FoF) halo mass to the number of contained subhalos (or galaxies) using a conditional probability, $P(N|M)$, where $M$ is the FoF halo mass and $N$ is the number of subhalos inside the FoF halo. This probability function is obtained from numerical simulations \\citep{kravtsov04,berlind02,jing98} or from observations \\citep{zehavi04,abazajian05,zheng07} fitting the model to the observed two-point correlation functions.  Another variant of the SAM is a one that adopts the one-to-one monotonic correspondence between the galaxy luminosity and subhalo mass \\citep{marinoni02,vale04,vale06,shankar06}. This method is often called the ``correspondence model''. It is simpler than other methods because it requires only two prerequisites; the luminosity function of galaxies and mass function of subhalos. Its key assumption that a more massive subhalo hosts a brighter galaxy, is consistent with the hierarchical clustering picture. As the halo mass grows, its luminosity is expected to grow in general because merging of halos may be followed by the merging of galaxies. By matching these two functions, the subhalo mass is mapped to galaxy luminosity.  However, sometimes a galaxy may survive the merger event because baryonic component of a galaxy is usually more concentrated and is more tightly bound than its dark counterpart. These ``orphan'' galaxies \\citep{gao04} which have no separate corresponding halos could exist in the cluster regions. Also halos of mass below a certain characteristic scale are unable to provide baryonic matter with gravitational attraction strong enough to resist against supernova explosions which can tear up the small-mass system. Even though such cases are many, it is sound to expect that the galaxy census is closely related to the population of subhalos because most of the observed galaxies are field galaxies who have their own dark halos. It is worth investigating the hypothesis of the subhalo-to-galaxy correspondence model of galaxy formation.  In this paper, we apply the subhalo-galaxy correspondence model to a large N-body simulation and compare the statistical properties of simulated galaxies with those of galaxies observed by the SDSS. We organize this paper as follows. In section 2 we describe our simulation and the subhalo finding method. In section 3, we implement the subhalo-galaxy correspondence model and show how to assign morphological types to mock galaxies. In section 4, the local density is introduced to quantify local environments and the local density distributions of mock and SDSS galaxies are investigated. We also compare the luminosity distribution of simulated galaxies with those of the SDSS galaxies in section 5. The environmental dependence of central velocity distributions are studied for the mock and SDSS galaxies in section 6. And discussions and conclusions follow in section 7. ",
        "conclusions": "\\label{summary} We have proposed a model to assign galaxy luminosity and morphology to the dark subhalos directly identified in cosmological N-body simulations. It is assumed that the galaxy luminosity is a monotonic function of its host halo mass. In the $\\kappa$ model we assume that the halo masses of early and late type galaxies of equal luminosity have a constant ratio. Another alternative model is the $\\beta$ model which assumes the constant luminosity ratio between equal-mass galaxies of different types. This model has been proposed by \\citet{marinoni02} who found that the observed $B$--band luminosity function of galaxies can be reproduced from the PS function assuming double power-law mass-to-light ratios and derived halo occupation numbers. It has been expanded by \\citet{vale04,vale06} who adopted the satellite halo mass functions and directly link subhalos to the observed galaxies. In this paper, we have introduced the ratio of luminosity of the early and late morphological type galaxies at a given halo mass, and derived type-specific mass-to-light ratios as a function of subhalo mass. They are used to assign luminosity and morphology to subhalos. The mass-to-light ratio of the early type galaxies derived in this way has a minimum value of $\\Upsilon_u\\simeq100$ at the scale of $M_u\\simeq3\\times10^{11}h^{-1}{\\rm M_\\odot}$ in the $\\beta=2$ model. The mass-to-light ratio starts to exponentionaly increase below $M_u$ and increases in a power law above $M_u$.  We use the large-scale background galaxy number density as an environmental parameter. The smooth galaxy density field is obtained by using the adaptive spline kernel which enabled us to resolve crowded regions well. The local density distribution of the SDSS galaxies is well described by the log normal function. The obtained log normal parameter values indicate that brighter galaxies tend to locate in dense regions. The local density distribution of the simulated galaxies is quite similar to that of the SDSS galaixes in voids and moderate-density regions for both morphological types. The underestimation of galaxy population in clusters is thought to be due to the evaporation of subhalos by the dynamical friction and tidal stripping. This can also explain the discrepance in the luminosity function of the early type galaxies in high density regions.  Recently, \\citet{gott06} measured the genus statistic from a large sample of the SDSS galaxies and compared it with those of mock galaxies created by the three distinctive methods such as the semi-analytic galaxy formation model applied to the Millennium run \\citep{springel05}, a hydrodynamic simulation, and the subhalo-galaxy correspondence model adopted in this paper. It was found that the observed topology of large scale structure was best reproduced by the subhalo-galaxy correspondence model even though other models were also consistent with the observation. However, the observed topology was marginally inconsistent with all simulations in the sense that it showed a strong meatball topology at the significance level of $2.5 \\sigma$ at the scale studied. The prominence of the isolated high density regions in the observation seems to be due to the Sloan Great Wall which was a dominant structure in the sample analyzed.  We have found an impressive agreement between our simulated galaxies and the SDSS galaxies in the dependence of the central velocity dispersion on the local density and luminosity. The early type galaxies tend to have higher $\\sigma_v$ or higher mass in high density regions at a given luminosity when they are brighter than about $\\mathcal{M}_*$. In other words, these bright galaxies tend to become relatively fainter in high density regions at a given halo mass. This interesting dependence of the mass-to-light ratio on environment was successfully reproduced by our subhalo-galaxy correspondence model of galaxy formation. A more detailed study of this phenomenon will be presented in a forthcoming paper."
    },
    "0801/0801.4055_arXiv.txt": {
        "abstract": "{ A brief review of various methods to calculate radiative accelerations for stellar evolution and an analysis of their limitations are followed by applications to Pop I and Pop II stars.  Recent applications to Horizontal Branch (HB) star evolution are also described.  It is shown that models including atomic diffusion satisfy Schwarzschild's criterion on the interior side of the core boundary on the HB without the introduction of overshooting.  Using stellar evolution models starting on the Main Sequence and calculated throughout evolution with atomic diffusion, radiative accelerations are shown to lead to abundance anomalies similar to those observed on the HB of M15. ",
        "introduction": "In his book, \\citet{Eddington26} evaluated the equilibrium concentrations to be expected if atomic diffusion in the presence of differential radiation  pressure were efficient.  He concluded that some heavy metals should then completely dominate the spectrum. Since light and heavy elements are present in stellar spectra, he concluded (\\S 193, 196) that some mixing process made atomic diffusion inefficient. He suggested that it be meridional circulation (\\S 199). His argument for the importance of meridional circulation was qualitative.  In a later paper he evaluated quantitatively the meridional circulation velocity without further commenting on its efficiency in mixing stellar interiors \\citep{Eddington29}.  While it was realized in the 1940s that the presence of the giant branch in clusters implied that stars could not be completely mixed, Eddington's argument seems to have had enough weight to prevent the proper calculation of the effect of differential radiation pressure until the late 1960s even if Eddington had partially corrected his argument (see the 1930 correction on page xiii of \\citealt{Eddington26}). Work however continued on gravitational settling in outer solar layers \\citep{Biermann37,Wasiutynski58,AllerCh60}; the most important application was made to white dwarfs \\citep{Schatzman45}.  Instead of assuming that equilibrium concentrations were reached, one of us introduced the differential radiative acceleration term, \\gr, into the diffusion velocity equation of \\citet{AllerCh60} and calculated anomalies to be expected in atmospheric regions \\citep{Michaud70}.  Comparison to observations of ApBp stars suggested that \\gr{} plays a role in at least some stars. As large atomic data bases started becoming available in the 1980s, we realized that it would be possible to calculate \\gr{} throughout stellar interiors with good accuracy at the same time as the evolution proceeds.  All composition changes can then be taken into account self consistently during evolution for the \\gr{} and  the Rosseland averaged opacity.  It was first tried to make those calculations with TOPBase from the Opacity Project \\citep{AlecianMiTu93,LeBlancMi95,GonzalezArMi95,GonzalezLeAretal95} but using those data we could not reproduce the concentration dependence of the Rosseland averaged OPAL opacities around the solar center (see \\S 5.1 of \\citealt{TurcotteRiMietal98}) and we shifted to using OPAL spectra \\citep{RicherMiRoetal98} to calculate \\gr{}.  The original spectra they had used to calculate the OPAL opacity tables   \\citep{IglesiasRo93,IglesiasRo96,RogersIg92a} were kindly made available to us by Iglesias and Rogers.  In this  review of \\gr{} in stellar evolution, we will first briefly describe how the calculations are carried out in stellar evolution codes (\\S\\,\\ref{sec:calculations}) and compare the values of \\gr{} obtained using OP and OPAL data (\\S\\,\\ref{sec:CorrectionFactors}).  A few examples of the effect of \\gr{} in Pop I and Pop II stars are described (\\S\\,\\ref{sec:Role}) and, finally,  results obtained recently for HB stars are presented (\\S\\,\\ref{sec:HBStars}) showing that \\gr{}s play a role in stellar evolution (\\S\\,\\ref{sec:con}).  \\begin{figure*}[t!] \\resizebox{\\hsize}{!}{\\includegraphics[clip=true]{OPcorr_grad.eps}} \\caption{\\footnotesize Correction factors used in evolution calculations except for Fe, see the text. Adapted  from Fig. 6 of \\citealt{RicherMiRoetal98}. }. \\label{fig:correction} \\end{figure*} ",
        "conclusions": "\\label{sec:con} The availability of large atomic data bases has allowed, as described in \\S\\,\\ref{sec:calculations},  to calculate stellar evolution models for Pop I and II stars up and past the giant branch including all effects of atomic diffusion. They cause abundance anomalies not only in Ap stars, as originally suggested by \\citet{Michaud70}, but also in HB stars of clusters (\\S \\ref{sec:SurfaceAbundanceAnomalies}) and possibly in turnoff stars (\\S \\ref{sec:Role}). They play an essential role in driving pulsations in sdB stars \\citep{FontaineBrChetal2003}, the field analogue of HB stars. It is furthermore not only the surface region which is affected but 50\\% of the stellar radius and 10$^{-3}$ of its mass \\citep{RichardMiRi2002,MichaudRiRi2007}.  \\citet{Eddington26} was right however in suggesting that competing processes also have a role to play.  For instance, in M15, rotation plays a role probably through meridional circulation. The arrow in Figure \\ref{fig:Fe} shows the only star with $\\teff> 11500$\\,K which has no abundance anomaly.  It is also the only relatively rapidly rotating star.  \\citet{QuievyChMietal2007} have shown that meridional circulation explained that the stars with $\\teff<11500$\\,K have a normal abundance and not the 5$\\times$ overabundance that Figure \\ref{fig:Fe} would suggest.  While atomic diffusion driven by \\gr{}s plays the main role in creating abundance anomalies on the HB, it has to compete with the effects of rotation just as in HgMn or AmFm stars \\citep{CharbonneauMi91}."
    },
    "0801/0801.1670_arXiv.txt": {
        "abstract": "We present a study of the mass-metallicity ($M-Z$) relation and H~II region physical conditions in a sample of 20 star-forming galaxies at $1.0 < z < 1.5$ drawn from the DEEP2 Galaxy Redshift Survey. We find a correlation between stellar mass and gas-phase oxygen abundance in the sample and compare it with the one observed among UV-selected $z \\sim 2$ star-forming galaxies and local objects from the Sloan Digital Sky Survey (SDSS). This comparison, based on the same empirical abundance indicator, demonstrates that the zero point of the $M-Z$ relationship evolves with redshift, in the sense that galaxies at fixed stellar mass become more metal-rich at lower redshift. Measurements of [O III]/H$\\beta$ and [N II]/H$\\alpha$ emission-line ratios show that, on average, objects in the DEEP2 $1.0 < z < 1.5$ sample are significantly offset from the excitation sequence observed in nearby H~II regions and SDSS emission-line galaxies. In order to fully understand the causes of this offset, which is also observed in $z\\sim2$ star-forming galaxies, we examine in detail the small fraction of SDSS galaxies that have similar diagnostic ratios to those of the DEEP2 sample. Some of these galaxies indicate evidence for AGN and/or shock activity, which may give rise to their unusual line ratios and contribute to Balmer emission lines at the level of $\\sim 20$\\%. Others indicate no evidence for AGN or shock excitation yet are characterized by higher electron densities and temperatures, and therefore interstellar gas pressures, than typical SDSS star-forming galaxies of similar stellar mass. These anomalous objects also have higher concentrations and star formation rate surface densities, which are directly connected to higher interstellar pressure. Higher star formation rate surface densities, interstellar pressures, and H~II region ionization parameters may also be common at high redshift. These effects must be taken into account when using strong-line indicators to understand the evolution of heavy elements in galaxies. When such effects are included, the inferred evolution of the $M-Z$ relation out to $z\\sim 2$ is more significant than previous estimates. ",
        "introduction": "The abundance of heavy elements in the H~II regions of galaxies reflects the past history of star formation and the effects of inflows and outflows of gas. A characterization of the evolution of chemical abundances for galaxies of different masses is therefore essential to a complete model of galaxy formation that includes the physics of baryons \\citep{delucia04,finlator07}. Important observational constraints for such models come from determining the scaling relations at different redshifts among galaxy luminosity, stellar mass, and metallicity, which, for star-forming galaxies, typically consists of the oxygen abundance. However, one of the key challenges is to take the observationally measured quantities, i.e. strong, rest-frame optical emission-line ratios, and connect them with the physical quantity of interest, i.e. oxygen abundance.    \\begin{deluxetable*}{lcccccccc} \\tablewidth{0pt} \\tabletypesize{\\footnotesize} \\tablecaption{Galaxies Observed with Keck~II NIRSPEC\\label{tab:obs}} \\tablehead{ \\colhead{~~~~~~~~~~DEEP ID~~~~~~~~~~} & \\colhead{~~~~~R.A. (J2000)~~~~~} & \\colhead{~~~~~Dec. (J2000)~~~~~} & \\colhead{~~~~~$z_{{\\rm H}\\alpha}$~~~~~} & \\colhead{~~~~~$B$~~~~~} & \\colhead{~~~~~$R$~~~~~} & \\colhead{~~~~~$I$~~~~~} & \\colhead{~~~~~$M_B$~~~~~} & \\colhead{~~~~~$U-B$~~~~~} } \\startdata 42044579 \\dotfill   &    02 30 43.46   &   00 42 43.60 & 1.0180 &   23.22   &    22.97   &    22.40  &  -21.20 &   0.54 \\\\ 22046630 \\dotfill   &    16 50 13.83   &   35 02 01.78 & 1.0225 &   23.64   &    23.02   &    22.31  &  -21.37 &   0.69 \\\\ 22046748 \\dotfill   &    16 50 14.55   &   35 02 04.31 & 1.0241 &   24.43   &    23.76   &    22.90  &  -20.86 &   0.86 \\\\ 42044575 \\dotfill   &    02 30 44.85   &   00 42 51.33 & 1.0490 &   23.08   &    22.94   &    22.56  &  -21.06 &   0.34 \\\\ 42010638 \\dotfill   &    02 29 08.74   &   00 23 28.40 & 1.3877 &   22.93   &    22.85   &    22.54  &  -22.12 &   0.49 \\\\ 42010637 \\dotfill   &    02 29 08.74   &   00 23 32.87 & 1.3882 &   24.20   &    23.98   &    23.72  &  -20.87 &   0.44 \\\\ 42021412 \\dotfill   &    02 30 44.55   &   00 30 50.73 & 1.3962 &   24.07   &    23.74   &    23.12  &  -21.91 &   0.78 \\\\ 42021652 \\dotfill   &    02 30 44.70   &   00 30 46.19 & 1.3984 &   22.97   &    22.24   &    21.32  &  -24.01 &   1.01 \\\\ \\enddata \\tablecomments{Units of right ascension are hours, minutes, and seconds, and units of declination are degrees, arcminutes and arcseconds.} \\end{deluxetable*}   In the local universe, \\citet{tremonti04} have used a sample of $\\sim$53,000 emission-line galaxies from the Sloan Digital Sky Survey (SDSS) to investigate the luminosity-metallicity ($L-Z$) and mass-metallicity ($M-Z$) relationships. For this sample, metallicities were estimated on the observed spectra of several strong emission lines, including [O II] $\\lambda\\lambda$3726, 3729, H$\\beta$, [O III]$ \\lambda\\lambda$5007, 4959, H$\\alpha$, [N II] $\\lambda\\lambda$6548, 6584, and [S II] $\\lambda\\lambda$6717, 6731. At increasing redshifts, as the strong rest-frame optical emission lines shift into the near-IR, metallicities are typically based on smaller subsets of strong emission lines through the use of empirically calibrated abundance indicators \\citep[e.g.][]{pp04,pagel79}. Much progress has been made recently in assembling large samples of star-forming galaxies with abundance measurements at both intermediate redshift \\citep{kewley02,savaglio05} and at $z>2$ \\citep{erb06a}. However, we have only begun to gather chemical abundance measurements for galaxies at $z \\sim 1-2$ \\citep[][hereafter Paper I,]{maier06,shapley05}. In this work, we continue our efforts to fill in the gap of chemical abundance measurements during this important redshift range, which hosts the emergence of the Hubble sequence of disk and elliptical galaxies \\citep{dickinson00}, and the buildup of a significant fraction of the stellar mass in the universe \\citep{drory05,dickinson03} prior to the decline in global star formation rate (SFR) density \\citep{madau96}.      Chemical abundances for high-redshift galaxies are commonly estimated using locally calibrated empirical indicators. Yet it is crucial to recognize the fact that a considerable fraction of the $z \\sim 1$ and $2$ galaxies with measurements of multiple rest-frame optical emission lines do not follow the local excitation sequence described by nearby H~II regions and star-forming galaxies in the diagnostic diagram featuring the [O III] $\\lambda$5007/H$\\beta$ and [N II] $\\lambda$6584/H$\\alpha$ emission-line ratios \\citep[Paper I;][]{erb06a}. On average, the distant galaxies lie offset towards higher [O III] $\\lambda$5007/H$\\beta$ and [N II] $\\lambda$6584/H$\\alpha$, relative to local galaxies. As discussed in Paper I and \\citet{groves06}, several causes may account for this offset, in terms of the prevailing physical conditions in the H~II regions of high-redshift galaxies. The relevant conditions are H~II region electron density, hardness of the ionizing spectrum, ionization parameter, the effects of shock excitation, and contributions from an active galactic nucleus (AGN). It is still unclear which of these are most important for determining the emergent spectra of high- redshift galaxies. Understanding this offset in emission-line ratios is important, not only because it provides evidence that physical conditions in the high redshift universe are different from the local ones, but also because the application of an empirically calibrated abundance indicator to a set of H~II regions or star-forming galaxies rests on the assumption that these objects are similar, on average, to those on which the calibration is based.  In this sense, the current work is also motivated by the interpretation of the offset in emission-line ratios among distant galaxies, and an assessment of the reliability of using local abundance calibrations for high-redshift star-forming galaxies. Instead of focusing on high-redshift objects, another approach is to study the properties in a class of nearby objects, which exhibit similar offset behavior on the emission-line diagnostic diagram. Unravelling the relations between the physical conditions and unusual diagnostic line ratios for such objects aids the understanding of high-redshift galaxies. The SDSS, with its rich set of photometric and spectroscopic information, provides an ideal local comparison sample.   In this paper we expand on the analysis presented in Paper I, with an enlarged sample of DEEP2 star-forming galaxies observed with NIRSPEC on the Keck II telescope. The larger number of DEEP2 objects with near-IR observations enforces the conclusions drawn in the previous work. Furthermore, our detailed study of nearby SDSS objects with similar emission-line diagnostic ratios leads to a clearer physical explanation of the observed properties of the DEEP2 galaxies. The DEEP2 sample, near-IR spectroscopic observations, data reduction, and measurements are described in \\S 2. We present the oxygen abundances derived from measurements of [O III], H$\\beta$, H$\\alpha$, and [N II] emission lines in both individual as well as composite spectra in \\S 3. The mass-metallicity relationship and its evolution through cosmic time are discussed in \\S 4. In \\S 5 we investigate differences in $z \\sim 1.0-1.5$ H II region physical conditions with respect to local samples by examining nearby SDSS galaxies with similar emission-line diagnostic ratios. Finally, we summarize our main conclusions in \\S 6. A cosmology with $\\Omega_m = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $h = 0.7$ is assumed throughout. ",
        "conclusions": "We have compiled a sample of 20 star-forming galaxies at $1.0 < z < 1.5$ drawn from the blue cloud of the color bimodality observed in the DEEP2 survey, to study the correlation between stellar mass and metallicity, across a dynamical range of 2 orders of magnitude in stellar mass, as well as H~II region physical conditions at this redshift range. In order to gain some insights on the causes of the offset in the BPT diagram observed in high-redshift star-forming galaxies, we have examined the H II region diagnostic line ratios and host galaxy properties of the small fraction of SDSS galaxies that have similar diagnostic ratios to those of the DEEP2 sample. Our main results are summarized as follows: \\begin{enumerate}  \\item[1.] There is a correlation between stellar mass and gas-phase oxygen abundance in DEEP2 star-forming galaxies at $z \\sim 1.0$ and at $z \\sim 1.4$. We have found that the zero point of the $M-Z$ relationship evolves with redshift, in the sense that galaxies at fixed stellar mass become more metal-rich at lower redshift, by comparing the $1.0 < z < 1.5$ sample with UV-selected $z \\sim 2$ and SDSS local star-forming galaxies. At the low-mass end ($M_{\\star} \\sim 8\\times10^9 M_{\\odot}$), the relation at $1.0 < z < 1.5$ is offset by $\\sim$0.2 (0.35) dex from the local mass-metallicity relation according to the N2 (O3N2) indicator. The N2-based offset could be larger by as much as $\\sim$0.16 dex, when the systematic bias due to difference in H II region physical conditions between $1.0 < z < 1.5$ and the local universe is taken into account. At the high-mass end ($M_{\\star} \\sim 5\\times10^{10} M_{\\odot}$), the metallicity offset between the DEEP2 $1.0 < z < 1.5$ sample and the local SDSS sample is at least $\\sim 0.2$ dex, according to the O3N2 indicator.    \\item[2.] As observed previously for a very small sample of high-redshift galaxies, on average our new DEEP2 sample at $1.0 < z < 1.5$ is offset from the excitation sequence formed by nearby H~II regions and SDSS emission-line galaxies. By examining the small fraction of SDSS galaxies that have similar diagnostic ratios to those of the DEEP2 sample, we have found two likely causes for the anomalous emission-line ratios. One is the contribution from AGN and/or shock excitation at the level of $\\sim 20$\\%. The other is the difference in H~II region physical conditions, characterized by significantly larger ionization parameters, as a result of higher electron densities and temperatures, and hence higher interstellar ambient pressure, than the typical values of local star-forming galaxies with similar stellar mass. These unusual physical conditions are possibly connected to the host-galaxy properties, most importantly smaller sizes and higher star-formation rate surface densities. Our conclusion drawn from analyzing the SDSS data has been further verified by the fact that the DEEP2 objects more offset from the local excitation sequence in the BPT diagram also have higher electron densities than those closer to the local sequence. We cannot rule out either the contribution from AGN and/or shock excitation, or the difference in H~II region physical conditions, for the unusual emission-line diagnostic ratios of high-redshift star-forming galaxies.  \\item[3.] We have quantified the effects of different H II region physical conditions on the strong-line metallicity calibrations. The direct electron temperature method was used to estimate the ``true'' metallicity difference between offset SDSS objects with anomalous line ratios and more typical objects of similar stellar mass. Strong-line indicators were also used to estimate this difference. A comparison of these results reveals potential biases in the strong-line indicators. According to our test, the metallicities based on strong-line indicators can either be systematically too large by $\\sim$0.16 dex (N2, N2O2), or too low by $\\sim$0.12 dex ($R_{23}$), for objects with similar H~II region physical conditions to those observed in high-redshift galaxies. The bias of O3N2 is much less significant (too large by $\\sim$0.02 dex) for offset SDSS objects with anomalous line ratios, although it may not always be negligible given other ranges of physical conditions in terms of metallicity, ionization parameter, and electron density. \\end{enumerate}    The difference in H~II region physical conditions, which may commonly apply to $z\\sim 1.0-1.5$ star-forming galaxies, must be taken into account when strong-line abundance indicators are used to study the evolution of galaxy metallicity with redshift. There are at least two methods to remove the systematic bias from the effect of significantly different H~II region physical conditions on the strong-line abundance calibrations. One is to gather the abundance information with direct $T_e$ method for a sample of high-redshift H~II regions as the calibration sample, which may be hard to achieve, as auroral lines are difficult to measure except for very metal-poor objects or in deep, composite spectra. The other, which is currently feasible yet relies on photoionization models, is to quantify the biases in strong-line indicators when certain physical conditions are present and then compensate the biases when inferring abundances from strong-line indicators for high-redshift galaxies.  In this study we have presented evidence that high-redshift star-forming galaxies possess distinct H~II region physical properties, as characterized by on average larger ionization parameters, higher electron densities, and temperatures, which are possibly connected to their relatively smaller sizes and higher SFR surface densities. These conditions may be quite common during the epoch at $z \\geq 1$ when at least 50\\% of the local stellar mass density was formed \\citep{bundy06,drory05}. Therefore, they should be characterized in more detail for a full understanding of the star formation history of the universe as well as the buildup of heavy elements in galaxies. The next generation of ground-based near-IR multi-object spectrographs will play a key role in assembling rest-frame optical emission-line measurements for large samples of high-redshift galaxies, enabling the detailed study of star-forming galaxies in the early universe."
    },
    "0801/0801.4749_arXiv.txt": {
        "abstract": "We present a systematic survey of the temporal and spectral properties of all GRB X-ray afterglows observed by Swift-XRT between January 2005 and July 2007.  We have constructed a catalog of all light curves and spectra and investigate the physical origin of each afterglow segment in the framework of the forward shock models by comparing the data with the closure relations.  We search for possible jet-like breaks in the lightcurves and try to explain some of the \"missing\" X-ray jet breaks in the lightcurves. ",
        "introduction": "Studies of the presence or absence of jet breaks in GRB X-ray afterglows have been recently undertaken with several different approaches yielding differing results \\citep{burrows07,liang07b,panaitescu07,kocevski07}.  The importance of this work lies in that the results have vital implications on the energetics, geometry, and frequency of GRBs.  The fact that they do not behave as expected from pre-{\\it Swift} observations is not surprising in the context of how much we do not understand and have only recently learned about all of the aspects of X-ray afterglows.  To understand the jet break phenomena we must understand it in the global context of GRB and afterglow properties.  Therefore, the goal of our study is to do a census of X-ray afterglow properties by fitting a variety of physical models to each component of the afterglows and understand how the jet breaks fit in as one component in the larger coherent picture of this phenomenon. ",
        "conclusions": "Although the jet break phenomena is not fully understood, we are beginning to be able to better explain the paucity of expected observations of them.  We are able to identify a sizeable group of afterglows that likely contain jet breaks or post-jet break data even though they do not present themselves in the classical context of the full canonical model.  Due to observational limitations and additional possible physical model variations, even more afterglows may be consistent with the expectations from those that do contain confident jet breaks, but are currently indistinguishable. While we are beginning to understand or at least be able to explain the majority of our sample, we are also finding an interesting small subset of outliers that confidently do not contain jet breaks during the time interval in which we would expect to see them.  These afterglows require further investigation and perhaps are somehow fundamentally different in their jet and afterglow properties.   \\begin{theacknowledgments} JLR, DNB, DCM, and AF acknowledge support under NASA contract NAS5-00136. \\end{theacknowledgments}"
    },
    "0801/0801.2390_arXiv.txt": {
        "abstract": "We present {\\em GALEX} near-UV ($NUV$) and 2MASS $J$ band photometry for red sequence galaxies in local clusters. We define quiescent samples according to a strict emission threshold, removing galaxies with very recent star formation. We analyse the $NUV$--$J$ colour--magnitude relation (CMR) and find that the intrinsic scatter is an order of magnitude larger than for the analogous optical CMR ($\\sim$0.35 rather than 0.05 mag), in agreement with previous studies. Comparing the $NUV$--$J$ colours with spectroscopically-derived stellar population parameters, we find a strong ($> 5.5 \\sigma$) correlation with metallicity, only a marginal trend with age, and no correlation with the $\\alpha$/Fe ratio. We explore the origin of the large scatter and conclude that neither aperture effects nor the UV upturn phenomenon contribute significantly. We show that the scatter could be attributed to simple `frosting' by either a young or a low metallicity subpopulation. ",
        "introduction": "\\label{sec:intro} The optical colour--magnitude relation (CMR) shows that brighter early-type galaxies are also redder \\citep{vis77}, and is traditionally regarded as arising from the mass--metallicity sequence (cf. \\citealt{dre84}, \\citealt{kod97}). Gas loss, caused by supernova wind, occurs earlier in less massive galaxies. Therefore, a smaller fraction of gas is processed before being expelled from the less massive galaxy \\citep{mat71}, resulting in lower average metallicities \\citep{lar74}. \\citet{bow92} found a very small intrinsic scatter in the $U$--$V$ CMR ($\\sim$0.05 mag) and, due to the sensitivity of the $U$ band to the presence of young stars, interpreted this as a small age dispersion. Age and metallicity are observed to have the same effect on broadband optical colours, whereas spectral line indices can be used to break the degeneracy \\citep{wor94}. \\citet{kun98} claimed that the CMR is driven by metallicity variations with luminosity, although \\citet{nel05} found evidence for a strong age--mass relation in addition to this metallicity--mass trend (see also \\citealt{cal03}, \\citealt{tho05}).  The ultraviolet--optical CMR for non-star-forming galaxies has an intrinsic scatter an order of magnitude larger than its optical counterpart; $\\sim$0.5 mag compared to 0.05 mag (e.g. \\citealt{yi05}). Hot young stellar populations dominate the ultraviolet (UV) flux for $\\sim$100 Myr after an episode of star formation (ten times longer than H$\\alpha$ emission after star formation; \\citealt{lei99}). The large intrinsic scatter in the UV CMR is therefore often interpreted as differing quantities of very recent, albeit low level, star formation \\citep{fer00}.  In intermediate age populations ($\\sim$1--3 Gyr), the near-UV (NUV; 2000--3000\\AA) flux is dominated by hot stars on the main sequence turn-off \\citep[e.g.][]{oco99}. The sensitivity of the turn-off to the epoch of formation emphasises the importance of the UV bands for age determination \\citep{dor03}.  Old ($\\sim$10 Gyr) metal-poor populations have a significant UV flux contribution from very hot ($T_{\\rm eff}$ $\\sim$ 10000K) blue horizontal branch (BHB) stars \\citep{mara00, lee02}. However these tend to reside in globular clusters or galactic haloes (where Fe/H $< -1$), where they are useful age indicators \\citep{kav07}, rather than in relatively metal-rich elliptical galaxies.  The UV picture is further complicated by the presence of the ultraviolet upturn (or UV excess, UVX) phenomenon. First observed by \\citet{cod69}, this unanticipated upturn dominates the far-UV (FUV; $<$ 2000\\AA) in UVX galaxies. In contrast, the NUV can be decomposed into two separate components: the blue end of the main sequence/subgiant branch, and the UVX contribution \\citep{dor97}. \\citet{bur88} found that the UVX can sometimes be appreciable at wavelengths as long as 2700\\AA: for example, in NGC 4649 $\\sim$75 per cent of the NUV flux can be attributed to the UVX component. However, the UVX cannot be explained by the BHB population, as the temperature required to fit the upturn would be $T_{\\rm eff}$ $\\ga$ 20000K, whereas BHBs are usually no hotter than $T_{\\rm eff}$ $\\sim$ 12000K \\citep{oco99}.  \\citeauthor{bur88} further reported that FUV flux (assumed to trace the UVX) is strongly correlated with the Mg$_2$ line strength ($\\sim$metallicity) and also with the velocity dispersion, which is a proxy for galaxy mass. However, more recent studies (e.g. \\citealt{ric05}) have weakened the case for a strong UVX vs metallicity relation. From analysis of internal colour gradients, \\citet{oco92} concluded that the FUV flux in most early-types originates from old stellar components. Drawing on these results, the source of the UVX is tentatively identified as hot, low mass, helium burning stars, such as extreme horizontal branch (EHB) or `failed' AGB (AGB-manqu\\'{e}) stars and their progeny (see \\citealt{yi97}, or the review \\citealt{oco99}).  The {\\em Galaxy Evolution Explorer} ({\\em GALEX}; launched in 2003; \\citealt{mar05}, \\citealt{mor07}) is revolutionising UV astronomy, with high resolution imaging in two bands: near-ultraviolet ($NUV$; $\\lambda_{eff} = 2310$ \\AA) and far-ultraviolet ($FUV$; $\\lambda_{eff} = 1530$ \\AA). Using analysis of both $NUV$--$V$ and $FUV$--$V$ vs $B$--$V$ relations, \\citet{don06} suggest that the $FUV$--$NUV$ colour reflects an extension of the colour--metallicity relation into the UV, as well as deducing that $\\sim$10 per cent of ellipticals have residual star formation. Using the $NUV$--$\\,r$ colour, \\citet{kav06} also find non-negligible young stellar populations in morphologically selected early-type galaxies. \\citet{sal07} derive star formation rates (SFRs) from broadband photometry dominated by the UV, and find good agreement with SFRs deduced from spectroscopic indices (predominantly using H$\\alpha$). However, they also confirm that some galaxies with no H$\\alpha$ emission show signs of star formation in the UV bands and attribute this to post-starburst galaxies.  Here, we build upon these previous studies by exploring the relationship between the $NUV$--$J$ colour and spectroscopic stellar population indicators for a sample of quiescent, red sequence galaxies in nearby clusters. This paper is organised as follows. Section \\ref{sec:data} describes our two red sequence samples and associated 2MASS and {\\em GALEX} datasets. The criteria used to remove galaxies with emission are described. In Section \\ref{sec:results} we show that a large intrinsic scatter is found in the $NUV$--$J$ colours of these quiescent cluster galaxies. Metallicity is shown to be strongly correlated with the $NUV$--$J$ colour, although there remains a large residual scatter. Section \\ref{sec:discuss} discusses possible explanations for this scatter, showing that morphological abnormalities, aperture bias and the UV upturn do not contribute significantly. We investigate simple `frosting' models with a low mass fraction of younger stars (or alternatively a low mass fraction population of low metallicity, blue horizontal branch stars), and show that these could account for the scatter. The uncertainties in the $NUV$ K-correction are also discussed. Our conclusions are given in Section \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion} Using {\\em GALEX} UV and 2MASS $J$ band photometry, we have investigated the relationship between UV--IR colours and spectroscopically derived stellar population parameters (age, metallicity and $\\alpha$-abundance) for red sequence galaxies in local clusters.  We select galaxies using strict emission criteria to avoid contamination from galaxies with very recent star formation. We analyse the $NUV$ colour--magnitude relation (CMR) for our two samples of quiescent galaxies (920 in NFPS; 156 in SSC), and find rms dispersions of 0.37 and 0.30 mag (intrinsic scatter of 0.36 and 0.29 mag) respectively. This is similar to previously reported values of $\\sim$0.5 mag and is an order of magnitude larger than the scatter in the optical CMR ($\\sim$0.05 mag).  We compared the $NUV$--$J$ colour to the spectroscopic stellar population parameters for 87 galaxies in the SSC sample and found the following:  \\begin{itemize} \\item There is a significant $NUV$--$J$ vs metallicity trend, with a slope of $1.27 \\pm 0.23$ and an rms dispersion of 0.32 mag. \\item There is only a weak $NUV$--$J$ vs age trend after the metallicity effect has been removed, and no correlation with $\\alpha$-abundance. \\item There is a large intrinsic scatter ($\\sim$0.25 mag) in the $NUV$--$J$ colour at fixed age and metallicity which cannot be easily accounted for with simple stellar populations. \\end{itemize}  The unexpected blue colours of at least some objects, including an influential outlier, cannot be attributed to large radius contamination from other objects, and aperture bias cannot account for the large scatter. Corrections for the UV upturn (UVX) phenomenon are relatively small ($\\sim$0.2 mag) and are similar galaxy-to-galaxy, so do not reduce the intrinsic scatter.  We find that the large $NUV$--$J$ intrinsic scatter could be attributed to galaxy `frosting' by small ($<$ 5 per cent) populations of either young stars or a low metallicity blue horizontal branch."
    },
    "0801/0801.0489_arXiv.txt": {
        "abstract": "{There is ${}^3P_2$ neutron superfluid region in NS (neutron star) interior. For a rotating NS, the ${}^3P_2$ superfluid region is like a system of rotating magnetic dipoles. It will give out electromagnetic radiation, which may provides a new heating mechanism of NSs. This heating mechanism plus some cooling agent may give sound explanation to NS glitches. ",
        "introduction": "Since the discovery of NSs (neutron star) and glitches in late 1960s, it is generally believed that there is superfluidity in NS interior (Ruderman 1976, Shapiro et al. 1983, Elgar{\\o}y et al. 1996). Cooling mechanism associated with superfluidity was first proposed by Flowers et al. (1976). Not until recently, the importance of superfluidity is considered seriously in the \"Minimal model\" (Gusakov et al. 2004, Page et al. 2004, Kaminker et al. 2006).  Then one may ask: since cooling agent associated with superfluidity must be considered in the \"Minimal model\", what about the heating mechanisms? Heating mechanism accompanied with superfluidity has been discussed by Alpar et al. (1989), and was taken into NS cooling model by Umeda et al (1994), Page et al. (2006), Tsuruta (2006). Here we investigate another possibility: microscope magnetic dipole radiation (MMDR) in NSs, possibly another heating mechanism associated with superfluidity.  The idea is as follows: there is ${}^3P_2$ superfluid region in the interior of a NS. We calculate its paramagnetic properties in the presence of a background magnetic field. Since the neutron superfluid is in a vortex state, the ${}^3P_2$ superfluid region is like a system of rotating magnetic dipoles. This system will give magnetic dipole radiation. The emitted photon can not penetrate the NS matter, thus provides a heating mechanism of NSs. The origin of this heating mechanism is microscopic magnetic dipole radiation, so we call it MMDR heating of NSs.  This heating mechanism plus some cooling agent may give sound explanation to NS glitches (Bai et al. 2006; Peng et al. 2006; Peng 2007).  Before going into detailed calculations of MMDR, we will first look at the superfluidity in NSs. ",
        "conclusions": "Here we have presented another possible heating mechanism of NSs, associated with superfluidity. It can be compared with other heating mechanisms and  cooling agents of NSs (Gusakov et al. 2004, Page et al. 2006). \\begin{figure}[!htbp] \\centering \\includegraphics[width=\\textwidth]{MMDRHeating.eps}\\\\ \\caption{MMDR Heating VS Cooling. The brown one represents MMDR heating. It increase with decreasing temperature, as a result of increasing thermal factor. The red, green, and black one corresponds to photon cooling, MUrca (Modified Urca) process, and PBF (Pair Breaking and Formation) process respectively. (Adapted from Gusakov et al. (2004).)} \\label{MMDR} \\end{figure}  As shown in Fig. \\ref{MMDR}, MMDR heating may be dominate only in the photon cooling stage. So it will not affect the cooling scenario of young and mid age NSs. But for old NSs, e.g. PSR 1055-52, it may serve as a moderate heating (Tsurata 2006, Page 2006).  Using toy model given by Yakovlev et al. (2003), we can make illustrative calculation of NS cooling, including the MMDR heating. \\begin{figure}[!htbp] \\centering \\includegraphics[width=\\textwidth]{CoolingCurves.eps}\\\\ \\caption{Cooling Curves Including MMDR Heating. The lower and upper curve correspond to MUrca and bremsstrahlung process dominated case respectively. Whereas the plateau at latter stage is due to MMDR heating. The physical input are the same as Yakovlev et al. (2003) except that we include the MMDR heating. Remind: these cooling curves are for illustrative use only.} \\label{CoolingCurves} \\end{figure} Fig. \\ref{CoolingCurves} shows that, there is a plateau in the late stage of NS cooling curves. This plateau can be compared with that of Kang's (2007).  This MMDR heating must cease after some time. In our case, the cease of MMDR heating has several possibilities. \\begin{enumerate} \\item As shown by Huang et al. (1982), the energy input of MMDR heating is through Ekman pumping. When this agent is out of work, MMDR heating have to make a pause. \\item When there is a phase transition in the core, e.g. deconfinement of hardrons, all the hadron processes disappear including MMDR heating. \\item Due complexities of superfluidity gap (Elagr{\\o}y et al. 1996), the ${}^3P_2$ superfluidity region becomes slimer when the core becomes more compact. The MMDR heating contribution is ignorable in this case. \\end{enumerate}  We presented another possible NS heating mechanism associated with superfluidity. The exact effect of MMDR heating needs accurate calculation of the cooling curves. This heating mechanism plus some cooling agent may give sound explanation to NS glitches. The detailed investigation is the scope of another work (Peng 2007)."
    },
    "0801/0801.0440_arXiv.txt": {
        "abstract": "We compare the stellar wind torque calculated in a previous work (Paper~II) to the spin-up and spin-down torques expected to arise from the magnetic interaction between a slowly rotating ($\\sim 10$\\% of breakup) pre-main-sequence star and its accretion disk.  This analysis demonstrates that stellar winds can carry off orders of magnitude more angular momentum than can be transferred to the disk, provided that the mass outflow rates are greater than the solar wind.  Thus, the equilibrium spin state is simply characterized by a balance between the angular momentum deposited by accretion and that extracted by a stellar wind.  We derive a semi-analytic formula for predicting the equilibrium spin rate as a function only of the ratio of $\\dot M_{\\rm w} / \\dot M_{\\rm a}$ and a dimensionless magnetization parameter, $\\Psi \\equiv B_*^2 R_*^2 (\\dot M_{\\rm a} v_{\\rm esc})^{-1}$, where $\\dot M_{\\rm w}$ is the stellar wind mass outflow rate, $\\dot M_{\\rm a}$ the accretion rate, $B_*$ the stellar surface magnetic field strength, $R_*$ the stellar radius, and $v_{\\rm esc}$ the surface escape speed.  For parameters typical of accreting pre-main-sequence stars, this explains spin rates of $\\sim 10$\\% of breakup speed for $\\dot M_{\\rm w} / \\dot M_{\\rm a} \\sim 0.1$. Finally, the assumption that the stellar wind is driven by a fraction of the accretion power leads to an upper limit to the mass flow ratio of $\\dot M_{\\rm w} / \\dot M_{\\rm a} \\la 0.6$. ",
        "introduction": "\\label{sec_intro}    The slow rotation rates of low to intermediate mass ($\\la 2 M_\\odot$) pre-main-sequence stars remains one of the most important aspects of star formation that has, so far, resisted a generally accepted explanation.  By the time they become optically visible as T Tauri stars \\citep[TTSs;][]{joy45}, approximately half of them are observed to rotate at approximately 10\\% of breakup speed \\citep[the ``slow rotators''; e.g.,][]{vogelkuhi81, bouvier3ea97, rebull3ea04, herbstea07}.  This is a surprise because many TTSs (the Classical T Tauri stars; CTTSs) are actively accreting material from surrounding Keplerian disks \\citep{lyndenbellpringle74, bertout3ea88, calvetgullbring98, muzerolle3ea01}.  At a typical accretion rate of $\\dot M_{\\rm a} \\sim 10^{-8} M_\\odot$ yr$^{-1}$, the angular momentum deposited by accreting disk material should spin up a CTTS to near breakup speed in $\\sim 10^6$ years \\citep{hartmannstauffer89, mattpudritz07coolstars}.  Since the accretion phase lasts for $10^6$ -- $10^7$ years \\citep{lyolawson05, jayawardhanaea06}, since the stars accrete at much higher rates prior to the TTS phase, and since the stars are still contracting \\citep{rebullea02}, an efficient angular momentum loss mechanism is required to explain the existence of the slow rotators.  A few interesting and important ideas for explaining the TTS slow rotators have been developed over the last two decades.  These have resulted in the star-disk interaction model of \\citet{ghoshlamb78}, applied to CTTSs by \\citet[][and see \\citealp{camenzind90}]{konigl91}, the X-wind model \\citep{shuea94}, and the idea that stellar winds provide strong torques \\citep{hartmannstauffer89, toutpringle92, paatzcamenzind96, ferreira3ea00, mattpudritz05l}.  Although both have advanced our understanding of the magnetic star-disk interaction, neither the Ghosh \\& Lamb nor X-wind models are without problems \\citep{ferreira3ea00, uzdensky04, mattpudritz05}, and the idea that stellar winds are important has not yet been worked out in sufficient detail to compare to the other models.   In \\citet[][hereafter Paper~I]{mattpudritz05l}, we further explored powerful stellar winds as a solution to the angular momentum problem and suggested that a fraction of the accretion power provides the energy necessary to drive the wind.  We showed that stellar winds are capable of carrying off the accreted angular momentum, provided that $\\dot M_{\\rm w} / \\dot M_{\\rm a} \\sim 0.1$, where $\\dot M_{\\rm w}$ is the outflow rate of material that is magnetically connected to the star (the ``stellar wind'').  This analysis included a formulation for the stellar wind torque that contained the Alfv\\'en radius ($r_{\\rm A}$), which is not easily determined a priori in the wind, and the conclusions were based on a one-dimensional scaling estimate of this important physical quantity.  Thus, while it is clear that accretion-powered stellar winds (APSWs) can in principle provide the necessary spin-down torque, this idea requires further development to produce a more detailed model.  Toward this goal, \\citet[][hereafter Paper~II]{mattpudritz08II} used 2-dimensional (axisymmetric) magnetohydrodynamic simulations to solve for $r_{\\rm A}$ and calculate realistic stellar wind torques for a range of parameters.  In the present paper, we use the stellar wind solutions of Paper~II to compare the stellar wind torque to the torques expected to arise from the star-disk interaction. Furthermore, we find new solutions for stellar spins, based upon torque balance between the accretion torque and the APSW spin-down torque.  This paper begins with a brief description of the simulation results of Paper~II (\\S \\ref{sec_simulations}).  We then compare the stellar wind torque to the star-disk spin-down torque in section \\ref{sec_tdsd} and then to the star-disk spin-up torque in section \\ref{sec_equilibrium}, which contains spin-equilibrium solutions. Section \\ref{sec_discussion} contains a summary and discussion. ",
        "conclusions": "\\label{sec_discussion}    In this work, we have further developed the accretion-powered stellar wind model proposed in Paper~I, where the stellar wind magnetic lever arm length, $r_{\\rm A}$, was taken as a parameter to determine the stellar wind torque.  We employed the simulation results of Paper~II (see \\S \\ref{sec_simulations}) to obtain stellar wind torques for several cases representative of T Tauri systems.  We examined the total torque on the star arising from the stellar wind plus the magnetic interaction between the star and its accretion disk.  Our results can be summarized as follows. \\begin{enumerate}  \\item{We found that the spin-down torque from a stellar wind can be orders of magnitude stronger than the spin-down portion of the star-disk interaction torque, for slowly rotating stars with mass loss rates substantially larger than the solar wind outflow rate (see \\S \\ref{sec_tdsd}).  This confirms the assumption of Paper~I that the condition for net zero torque on the star (spin equilibrium) is simply determined by a balance between the stellar wind torque and the accretion torque.}  \\item{Using the computed stellar wind torques for several cases, we looked at the conditions for spin equilibrium (\\S \\ref{sec_balancing}).  We found that a rotation rate of 10\\% of breakup speed typically requires $\\dot M_{\\rm w} / \\dot M_{\\rm a}$ equal a few tens of percent, confirming the original suggestion by \\citet{hartmannstauffer89} that stellar winds may be capable of removing accreted angular momentum.}  \\item{For most cases in spin equilibrium, the disk truncation radius was very close to the corotation radius, though this is not a general requirement of the APSW model (\\S \\ref{sec_state}).}  \\item{Accretion power is generally sufficient to power a stellar wind that is capable of solving the angular momentum problem (\\S \\ref{sec_power}), as suggested in Paper~I.  The energy requirements for most of the cases considered here is relatively large, and more work is needed to further constrain the energy coupling.}  \\item{Under the assumption that the stellar wind is accretion powered, the cases we examined suggested a hard upper limit of $\\dot M_{\\rm w} / \\dot M_{\\rm a} \\la 0.6$.}  \\item{Finally, in section \\ref{sec_semianalytic} we used the results from Paper~II to derive a semi-analytic formulation for the equilibrium spin rate predicted by the APSW model.  We found the that spin rate, expressed as a fraction of breakup speed, generally depends only on the two dimensionless parameters $\\dot M_{\\rm w} / \\dot M_{\\rm a}$ and $\\Psi \\equiv B_*^2 R_*^2 (\\dot M_{\\rm a} v_{\\rm esc})^{-1}$.}  \\end{enumerate}    The APSW model incorporates several previous ideas.  As in all other models that emphasize the role of stellar magnetic fields, the interaction of the magnetized star with the disk leads to the truncation of the disk and accretion of material along field lines onto the star.  The APSW model adopts the finding of the Ghosh \\& Lamb-type models \\citep[e.g.,][]{ghoshlamb78, konigl91, armitageclarke96}, which is also supported by numerical simulations \\citep[][]{romanovaea02, long3ea05}, that the angular momentum of accreting material is transferred to the star.  However, in contrast to the Ghosh \\& Lamb-type models, we found that for slow rotators, any spin-down torque arising from the star-disk interaction is negligible (item 1 above). Instead, the stellar spin-up torque from accretion is counteracted by an accretion-driven stellar wind, which carries a comparable amount of angular momentum out of the system.  Compared to all existing models, APSW is distinct in that it conceptually links the driving of the stellar wind to the energy released by the accretion process (via $\\epsilon$).  In other ways, the general picture of APSW is similar to other angular momentum models that utilize winds.  In particular, the X-wind \\citep{shuea94, ostrikershu95}, the Reconnection X-wind \\citep{ferreira3ea00}, and other works considering stellar winds \\citep{hartmannstauffer89, paatzcamenzind96} all find that, in order to carry away significant angular momentum, the mass outflow rate needs to be of the order of 10\\% of the accretion rate (item 2 above).  Except for the X-wind, in all of the above scenarios, the outflow is magnetically connected to the star, and thus extracts angular momentum directly from the star. By contrast, the X-wind outflow is magnetically connected to the disk. Furthermore, the X-wind is unique in that it assumes that the accretion of material does not deposit angular momentum onto the star.  As mentioned in item 3 above, for most of the specific cases of our simulated winds, we found that in the spin equilibrium state, the truncation radius was very close to the corotation radius.  This is similar to the prediction of the disk locking models, which includes both the Ghosh \\& Lamb-type models and the X-wind.  However, in contrast with the disk locking models, it is not a requirement of APSW that $R_{\\rm t}$ be close to $R_{\\rm co}$.  Measurements of the location of the inner edge of the gas disk \\citep{najita3ea03, carr07} suggest that $R_{\\rm t}$/$R_{\\rm co}$ is typically $\\sim 70$\\%.  We leave a more detailed comparison between models for future work.     There is observational evidence that the outflow rates from accreting young stellar systems are of the order of 10\\% of the accretion rates \\citep[e.g.,][]{hartigan3ea95, calvet97} and are therefore accretion powered \\citep{cabritea90}.  However, it appears that a large fraction of this flow (which is usually probed by forbidden emission coming from large spatial scales) originates in the disk, rather than the star \\citep{ferreira3ea06}.  It is not yet clear how much of the total observed flow may originate in a stellar wind.  There is some evidence specifically for stellar winds from CTTSs \\citep{beristain3ea01, edwardsea03, dupreeea05, edwardsea06, kwan3ea07}, as distinct from disk winds, and that these are accretion powered \\citep[e.g.,][]{edwardsea03, edwardsea06}, but the mass outflow rates are not yet well constrained \\citep{dupreeea05}.  Additional work constraining the value of $\\dot M_{\\rm w} / \\dot M_{\\rm a}$, the stellar wind driving mechanism, and the stellar magnetic field strength and geometry will help to provide stringent and quantitative tests for the APSW model.  The predictions of the spin equilibrium state can also be checked observationally \\citep{johnskrullgafford02}.  This will likely require large samples of stars, due to large uncertainties in measured parameters, and since intrinsic variability in real systems \\citep[e.g.,]{hartmann97} may only allow a spin equilibrium state to be achieved in a time-averaged sense.  The Ghosh \\& Lamb, X-wind, and APSW models all predict an equilibrium spin rate that depends on $\\Psi$, but of these three, only the APSW model contains an additional dependence on the stellar wind mass outflow rate.  For the power law fits to the simulations of Paper~II, the APSW spin equilibrium predicts a slightly different power-law of spin vs.\\ $\\Psi$ than the other models (\\S \\ref{sec_semianalytic})---though the exact dependence of the stellar wind torque has not been determined for all parameters.  In this series of papers, we have focused on the global problem of calculating the magnitude of stellar wind torques and comparing them with other torques acting on accreting stars.  In order to refine the APSW model further and make the predictions more precise, more work is required.  In particular, it is not yet clear how the presence of an accretion disk will influence the stellar wind torque, and conversely, how a stellar wind may influence the accretion process.  Also, although it is clear that there is enough accretion energy to power the stellar wind, it is still not known what actually drives the stellar wind and how the accretion power may transfer to it.  We suspect that a strong flux of hydromagnetic waves can be excited near the base of the accretion shock and can tap the energy released there, which may provide an efficient driver for the APSW.  We defer a rigourous investigation of the APSW driving mechanism to future work."
    },
    "0801/0801.0506_arXiv.txt": {
        "abstract": "{} {Our goal is to demonstrate the potential of the interferometric AMBER instrument linked with the Very Large Telescope Interferometer (VLTI) fringe-tracking facility FINITO to derive high-precision stellar diameters.} {We use commissioning data obtained on the bright single star V3879~Sgr. Locking the interferometric fringes with FINITO allows us to record very low contrast fringes on the AMBER camera. By fitting the amplitude of these fringes, we measure the diameter of the target in three directions simultaneously with an accuracy of 25 micro-arcseconds.} {We showed that V3879~Sgr has a round photosphere down to a sub-percent level. We quickly reached this level of accuracy because the technique used is independent from absolute calibration (at least for baselines that fully span the visibility null). We briefly discuss the potential biases found at this level of precision.} {The proposed AMBER+FINITO instrumental setup opens several perspectives for the VLTI in the field of stellar astrophysics, like measuring with high accuracy the oblateness of fast rotating stars or detecting atmospheric starspots.} ",
        "introduction": "In optical, long-baseline interferometry, the most popular observable is the contrast (also called the visibility) of the interferometric fringes that appear when superposing the beams coming from several distant telescopes. The major limitation is the random optical delay introduced by the atmospheric turbulence, which make these fringes jitter around on the detector by a quantity larger than the fringe spacing. Practically, fringe blurring is avoided, either by reducing the integration time to a few milliseconds, which prevents observations of faint targets; or by using a dedicated fringe-tracking facility, the purpose of which is to stabilize fringes by measuring and correcting the optical delay in real-time. Fringe-tracking is generally used to observe fainter objects or to increase the spectral resolution, but this is not the only application. Another important gain is the possibility of recording very low-contrast fringes that that rise above the noise level after integration times longer than a few seconds.  Such small fringes are produced by astronomical objects spatially resolved by the interferometer. More precisely when observing a single star modeled by a Uniform Disk of diameter \\dUD{}, the fringe visibility \\vis{} drops according to the following law: \\begin{equation} \\label{eq:0} \\visSquare{} = \\tfSquare(\\wave,t) \\;.\\;\\airySquare(\\dUD{}\\,\\blambda{}) \\end{equation} where\\begin{description} \\item[\\blambda{}] is the spatial frequency of the observation, given by the ratio of the distance between the telescopes \\base{} projected on the sky (called the baseline) and the observing wavelength \\wave{}. \\item[\\airy{}] is the Airy function (the Fourier transform of a disk of unitary size). The function \\airySquare{} actually goes to zero (fringes disappear) for a disk of diameter $\\dUD{}\\approx{}1.22/\\blambda{}$. After this first null the fringes reappear, but at lower visibility (second lobe). \\item[\\tfSquare{}] is the interferometric response of the instrumental chain, also called the transfer-function. It changes with the instrumental setup and the atmospheric conditions. Therefore, it has to be frequently calibrated by observing unresolved stars, or stars with known diameters. Statistical errors and fluctuations of this transfer-function generally dominate the error budget when trying to estimate the diameter with high accuracy. \\end{description} Because of the structure of Eq.~\\ref{eq:0}, collecting visibility points in the vicinity of the first null of the \\airySquare{} function makes the diameter estimation much less sensitive to the transfer-function uncertainty. For instance, measuring the exact spatial frequency at which the fringes disappear gives an estimation of \\dUD{}, formally independent from \\tfSquare{}, since one can use the relation: $\\dUD{}\\approx{}1.22/\\blambda{}$.  In this paper, we demonstrate the feasibility of this technique at the Very Large Telescope Interferometer \\citep[VLTI, see][]{Scholler-2006jul} by using an adequate setup of the fringe-tracker FINITO \\citep{Gai-2004oct} and of the scientific instrument AMBER \\citep{Petrov-2007mar}. Section~\\ref{sec:observationsAndDataReduction} describes the instrumental setup, the observations, and the data reduction. Section~\\ref{sec:resultsAndDiscussion} details the results and discusses the obtained accuracy. The paper ends with brief conclusions and perspectives.  \\begin{figure} % \\begin{flushright}  \\includegraphics[scale=1]{LeBouquin_fig1.ps} \\end{flushright} \\centering \\caption{\\label{fig:technique} VLTI array configuration during the observation (top) and corresponding simulated visibility-curves (solid line, bottom), assuming the uniform disk diameter of \\vSgr{} ($7.56$~mas). FINITO tracks the fringes along D0-H0 and G1-D0 in the H-band (star symbols), while AMBER records data of the three baselines across the K-band (circle symbols). The horizontal axes are the spatial frequencies \\blambda{}, marked in meter per micron, the vertical axes are the squared and linear visibilities.} \\end{figure} ",
        "conclusions": "From the instrumental point of view, we have demonstrated that: \\begin{itemize} \\item AMBER is able to record and process very low-contrast fringes when they are decently locked (down to about $1.5$\\% contrast on a bright $\\mathrm{K}=\\-0.5^{\\rm m}$ star, when integrating 60 frames of 1s each); \\item FINITO is able to lock fringes in the second visibility lobe (where fringes have less than $15$\\% contrast) with a sufficiently good efficiency. \\end{itemize} As an on-sky validation, we collected K-band interferometric fringes in the vicinity of the first visibility null of the bright star \\vSgr{}. By fitting the visibility curves, we measured a diameter of $7.56\\pm{}0.025$\\,mas in three directions simultaneously. We reached such sub-percent accuracy without spending additional time on calibrators since the technique is independent from absolute calibration (at least for baselines that fully span the visibility null). We show that, at this level of precision, several systematic error sources have to be taken into account, for example the spectral calibration of the instrument and the pupil lateral position.  From the scientific point of view, this work opens several perspectives for the VLTI in the field of stellar astrophysics:\\begin{itemize} \\item By using the AMBER+FINITO setup presented in this paper, one can measure a stellar diameter in three directions simultaneously with high accuracy. This can be used to constrain quickly the oblate photosphere of fast-rotating stars. \\item By using the same setup, one can precisely measure the value of the visibility minimum between the visibility lobes. Any departure from zero, even small, proves the presence of asymmetric features on (or above) the stellar surface, such as convective or magnetic starspots (in our \\vSgr{} dataset, this minimum is compatible with zero, see Fig.~\\ref{fig:fit}). \\item Finally, by associating fringe-tracking in the second lobe with baseline bootstrapping, one can record visibility points in the third visibility lobe. There, observables become very sensitive to the limb-darkening strength and its spectral dependency, which will allow thorough tests of stellar atmosphere models. \\end{itemize}"
    },
    "0801/0801.0766_arXiv.txt": {
        "abstract": "We study the relationship between different wave phenomena associated with a coronal mass ejection (CME) observed on  05 Mar. 2000. EIT waves were observed in the images recorded by EIT at 195~{\\AA}. The white-light LASCO/C2 images show clear deflection and propagation of a kink along with the CME. Spectroscopic observations recorded by the UVCS reveals excessive line broadening in the two O~{\\sc{vi}} lines (1032 and 1037 {\\AA}). Moreover very hot lines such as Si ~{\\sc{xii}} and Mg ~{\\sc{x}} were observed. Interestingly, the EIT wave, the streamer deflection and the intensity modulation along the slit were all propagating North-East. Spatial and temporal correlations show that the streamer deflection and spectral line broadening are highly likely to be due to a CME-driven shock wave and that the EIT wave is the signature of a CME-driven shock wave in the lower corona. ",
        "introduction": "Coronal mass ejections (CMEs) are one of the most fascinating and intriguing forms of solar activity. They occur when solar plasma threaded with topologically complex magnetic fields are ejected out into the corona and interplanetary medium. Recent technological developments have yielded a steady increase in qualitative as well as quantitative studies of CMEs and related phenomena. However, numerous CME-associated phenomena such as the relationship between different wave features (EIT waves: Thompson et al. 1998; CME-driven shock waves: Hundhausen et al. 1987) are not understood unambiguously. CMEs with speed greater than the Alfv\\'en speed of the local plasma produce shock waves whose effects can be traced through radio type II bursts (Klassen et al. 2000) and/or emission of very hot lines (compared to the ambient coronal temperatures; Raouafi et al. 2004). These waves could also be detected through deflections in remote streamer belts with respect to the ejected CME material (see Hundhausen et al. 1987 and Sime \\& Hundhausen 1987). The high quality data of SOHO/LASCO (Brueckner et al. 1995) yielded numerous cases for streamer deflections associated with super-Alfv\\'enic CME eruptions (see Sheeley et al. 2000). These observations led to the conclusion deduced by the previous authors that these deflections were indeed a consequence of CME-driven shock waves.  Spectroscopic signatures of coronal shock waves due to high speed CMEs ($>600$~km~s$^{-1}$; see Raymond et al. 2000) are line broadening (e.g., \\ion{O}{vi} 1032~{\\AA} \\& 1037~{\\AA} lines) and enhanced emissions in spectral lines that are rarely observed in the corona such as \\ion{Si}{xii} 520~{\\AA} and \\ion{Mg}{x} 625~{\\AA} (e.g., Raouafi et al. 2004) due to their relatively high formation temperatures ($>2$~MK). Up to now only a few CME-driven shock wave events have been reported through SOHO/UVCS (Kohl et al. 1995) observations: Raymond et al. 2000; Mancuso et al. 2002; Raouafi et al. 2004; and Ciaravella et al. 2005 (hereafter RMRC00-05).  Since the discovery of the H${\\alpha}$ Moreton waves (Moreton 1964), it was thought that these waves were due to the intersection of coronal shock waves (due to flares) with the chromosphere (e.g., Uchida 1968). Later when EIT waves (Thompson et al. 1998) were discovered based on the observations recorded by EIT (Delaboudini\\`ere et al. 1995), they were interpreted as the coronal manifestation of the chromospheric Moreton wave (Thompson et al. 1999). However, the difference in propagation patterns and speeds of Moreton and EIT waves questions the similarity between these two phenomena.  Although these waves are being observed more frequently the question remains open as to how these different wave phenomena are related to each other. Based on MHD simulations, (Chen et al. 2005a,b) showed that Moreton waves are the surface counter part of the CME driven shock wave. However, the EIT waves are the slow moving wave fronts traveling behind the Moreton wave due to the opening of the magnetic flux system. This was also suggested by Zhukov \\& Auch\\`ere (2004) and by Delanee (2000). However, this relationship and the nature of EIT waves remains elusive.  A CME event on 5 March 2000 has been observed simultaneously by different instruments on board SOHO (EIT, LASCO and UVCS) with different signatures of wave features. This event provides a unique opportunity to establish a relationship between different phenomena such as the CME driven shock wave, streamer deflection and the EIT waves. Observations and data analysis are described in \\S2 and results and conclusions are presented in \\S3. ",
        "conclusions": "We studied the relationship between different coronal waves (EIT wave, streamer deflection and a CME-driven shock wave) generated by the CME event on 5 March 2000 as observed by different instruments on board SOHO (EIT, LASCO and UVCS).  LASCO-C2 images show a deflection propagating outward at $\\sim260$~km~s$^{-1}$ (projected on the plane of the sky) in a remote streamer located approximately at $10^\\circ-15^\\circ$ CCW from north pole. The streamer kink was first seen in LASCO-C2 FOV ($\\sim2.5$~R$_{\\sun}$) at about 17:30~UT. No evidence for any CME material which could be the origin of such a deflection. The speed of the leading edge of the CME was $\\sim860$~km~s$^{-1}$, which is sufficient to generate a shock wave. Therefore, this propagating kink in the streamer provides strong evidence for a CME-driven shock wave in the corona as suggested by Sheeley et al. (2000).  UVCS spectra show excessive broadening in \\ion{O}{vi} lines and intensity enhancement in the hot lines of \\ion{Si}{xii} 520~{\\AA} and \\ion{Mg}{x} 625~{\\AA}. These are clear evidence for a CME-driven shock wave (see RMRC00-05). The analysis of the intensity modulation along the slit also reveals the propagation direction of the wave, North-East, that is the same as the kinking streamer. UVCS observations show that the shock wave would have reached the UVCS slit around 16:45~UT.  EIT difference images show evidence for an EIT wave front propagating North-East with speed of $\\sim55$~km~s$^{-1}$ (the solar disk shape not taken into account) up to the north polar hole where the propagation has stopped, satisfying the properties of EIT waves. The propagation of the wave in the South-East direction seems to be disabled by the presence of the complex active region. The EIT wave first appeared at 16:10~UT and disappeared gradually after 16:58~UT.  On one hand, temporal and spatial correlations between the events observed by UVCS and LASCO suggest that they are basically different manifestations of the same phenomenon. Measurements of the speed in different part of CME envelopes reveal that noses of CMEs travel faster than their flanks (Sheeley et al. 2000). Therefore we anticipate that the propagating shock would have a similar property and would have speed higher in front of the CMEs and lower in the flanks. Different speed in the different parts would provide evidence for the anisotropic propagation of associated shock waves. On the other hand, the EIT wave observed in the low corona was significantly slower the UVCS and LASCO events. A possible explanation for this is the opening of the magnetic flux system due to the expulsion of the CME rather than the manifestations of the CME driven shock wave as suggested by Chen et al. (2005a,b) based on MHD modeling and by Delanee (2000) and by Zhukov \\& Auch\\`ere (2004). Moreton waves, which are observed in the chromosphere are interpreted as counterpart of CME driven shock waves (e.g., Chen et al. 2005a). Unfortunately, we did not have H${\\alpha}$ high cadence observations for this event, which prevented a comparison with a chromospheric Moreton wave. However, the semi-circular wavefront was moving in the same direction as the coronal shock wave. This supports the hypothesis that all three features are linked to each other. They are generated by the same source, propagate in the same direction and have good time and spatial correlations. The only apparent problem resides in the propagation speed. However, note that the three wave-like structures do not share similar physical conditions. One of these is the space and its characteristics. The speed of Alfv\\'en waves is given by $\\displaystyle{V_A=\\frac{B}{\\sqrt{4\\pi\\rho}}}$, where $B$ is the magnetic field strength and $\\rho$ is the density. Depending on models, the density drops by 3 to 4 orders of magnitude or more between the very sparse corona and 2.0-3.0~R$_{\\sun}$, where the magnetic field drops by an order of magnitude or more (assuming the the coronal field is nearly potential and thus drops as $\\displaystyle{r^{-2}}$; see Altschuler \\& Newkirk 1969). While the high corona is increasingly less dense with increasing altitude the corresponding Alfv\\'en speed tends to increase. This is not the case in the lower corona where the plasma density is higher and thus is characterized by a smaller Alfv\\'en speed. This may explain at least partially the difference in speed. Another parameter is the propagation direction of the waves. The CME-driven shock wave is mainly propagating parallel to the magnetic field lines. However, the wave propagating on the solar disk does not share this property and propagates across. It encounters different magnetic structures that may allow for energy leakage which might consequently damp and slow down the wave (see for instance De Pontieu et al. 2004 on energy leakage of the acoustic oscillation $p$ modes through flux tubes with high inclinations to the chromosphere which contributes to the heating of this layer). A good illustration for that is given by active regions through which EIT waves could not propagate and also across polar holes. We think that energy carried by the wave is progressively leaked to different magnetic structures which may heat and accelerate the plasma in these structures. This interpretation needs to be carried further and deeper through additional analysis of other observational examples and also by numerical simulations."
    },
    "0801/0801.0599_arXiv.txt": {
        "abstract": "% In this contribution, I present a simplified overview of the evolution of the disk galaxy population since $z=1$, and a brief discussion of a few open questions. Galaxy evolution surveys have found that the disk galaxy population forms stars intensely at intermediate redshift.  In particular, they dominate the cosmic star formation rate at $z<1$ --- the factor of ten drop in cosmic average comoving star formation rate in the last 8 Gyr is driven primarily by disk physics, not by a decreasing major merger rate.  Despite this intense star formation, there has been little change in the stellar mass density in disk galaxies since $z=1$; large numbers of disk galaxies are being transformed into non-star-forming spheroid-dominated galaxies by galaxy interactions, AGN feedback, environmental effects, and other physical processes.  Finally, despite this intense activity, the scaling relations of disk galaxies appear to evolve little.  In particular, as individual galaxies grow in mass through the formation of stars, they appear to grow in radius (on average, the population grows inside-out), and they appear to evolve towards somewhat higher rotation velocity (i.e., mass is added at both small and large radii during this inside-out growth). ",
        "introduction": "The properties and evolution of the disk galaxy population are a critical diagnostic of the galaxy formation process in a dark matter-dominated Universe.  Galaxy disks, formed naturally as a result of a collapse in which at least some of the angular momentum is conserved \\citep{fall80}, are relatively fragile beasts. Interactions with low-mass halos are likely to thicken the stellar disk \\citep[e.g.,][]{toth92,benson04} and/or lead to warps and lopsidedness; interactions with galaxies/halos with masses within a factor of a few of the galaxy mass look likely to disrupt the stellar disk entirely \\citep[e.g.,][]{koda08}.  Thus, the evolution and properties of the disk galaxy population are not only a fascinating puzzle in their own right, but also give insight into those environments which are the least affected by the characteristic pummeling that galaxies in $\\Lambda$CDM seem to receive.  In the seven years since the first Vatican disk galaxies meeting, a number of the basic observational features of the evolution of the disk galaxy population have come into place.  In this contribution, I present a brief overview of what I consider to be some of the most important features of the evolution of the disk galaxy population since $z=1$ (roughly the last 7-8 billion years), and pose a few questions that I had after this conference.  These features fall under three broad themes: star formation in disks, stellar mass in disks, and disk galaxy scaling relations.  In what follows, in the spirit of a review, I will deliberately over-simplify the observational results somewhat in order to clarify the three basic messages of this contribution. ",
        "conclusions": "In this contribution, I have presented an over-simplified view of the evolution of the disk galaxy population: \\begin{itemize} \\item Disk galaxies form stars intensely since $z=1$.  They dominate the cosmic average SFR since $z=1$; the factor of 10 decline in cosmic SFR since $z=1$ is primarily driven by disk galaxy physics (gas consumption, decreased accretion from the cosmic web), not changes in the galaxy merging rate. \\item Despite this intense star formation, the disk galaxy population does not grow in mass since $z=1$.  Large numbers of star-forming disks are being transformed into non-star-forming spheroid-dominated galaxies. \\item As disk galaxies gain mass through star formation, they grow in radius (i.e., the population grows inside-out) and rotation velocity (i.e., mass is added at small and large radii during this inside-out growth). \\end{itemize} A final word.  This picture of a dynamic disk galaxy population has important implications for how we interpret the results of our observations.  Statements such as 'massive disk galaxies are in place at $z=1$ and appear to evolve little to the present day' are incorrect.  Some of these massive disk galaxies have become bulge-dominated at the present day (and are therefore not in the $z=0$ control disk galaxy sample), some are in (rare) even more massive disk galaxies.  This has implications for how one should interpret a variety of observations: bar fractions, metallicity--mass relations, mass--radius and mass--rotation velocity relations.  While some studies specifically address such issues (e.g., \\citealp{kassin06}), it is important to bear the (observed!) dramatic evolution of the disk galaxy population in mind when interpreting scaling relations and properties of the evolving disk galaxy population."
    },
    "0801/0801.3848_arXiv.txt": {
        "abstract": "I review recent progresses in the dynamics and the evolution of self-gravitating accretion discs. Accretion discs are a fundamental component of several astrophysical systems on very diverse scales, and can be found around supermassive black holes in Active Galactic Nuclei (AGN), and also in our Galaxy around stellar mass compact objects and around young stars. Notwithstanding the specific differences arising from such diversity in physical extent, all these systems share a common feature where a central object is fed from the accretion disc, due to the effect of turbulence and disc instabilities, which are able to remove the angular momentum from the gas and allow its accretion. In recent years, it has become increasingly apparent that the gravitational field produced by the disc itself (the disc's self-gravity) is an important ingredient in the models, especially in the context of protostellar discs and of AGN discs. Indeed, it appears that in many cases (and especially in the colder outer parts of the disc) the development of gravitational instabilities can be one of the main agents in the redistribution of angular momentum. In some cases, the instability can be strong enough to lead to the formation of gravitationally bound clumps within the disc, and thus to determine the disc fragmentation. As a result, progress in our understanding of the dynamics of self-gravitating discs is essential to understand the processes that lead to the feeding of both young stars and of supermassive black holes in AGN. At the same time, understanding the fragmentation conditions is important to determine under which conditions AGN discs would fragment and form stars and whether protostellar discs might form giant gaseous planets through disc fragmentation. ",
        "introduction": "Disc-like or flattened geometries are very common in astrophysics, from the large scale of spiral galaxies down to the small scales of Saturn's rings. In both these examples, the discs are characterized by a prominent structure either in the form of a spiral structure (in the case of galaxy discs) or in the form of gaps and spirals on small scales (in the case of Saturn's rings). The dynamics underlying the development of such structures is determined by the propagation of density waves, and the important role of the disc's self-gravity in their development has been clearly recognized (see, for example, \\cite{bertinbook}). In the above examples the system is either{\\it  collisionless}, such as in the case of spiral galaxies, where the dominant component is constituted of stars, or {\\it particulate}, such as in the case of Saturn's rings, where the dominant components (mainly rocks and pebbles) do undergo inelastic collisions, but the system cannot be simply described in terms of hydrodynamics. In the last thirty years increasing attention has been given to {\\it fluid} discs, where the dynamically active component is gaseous. Here, dissipative effects, associated with friction or viscosity, can significantly affect the dynamics of the disc. Through dissipation the fluid elements in the disc can lose their energy or, more fundamentally, their angular momentum, as I describe in more detail below, and can fall towards the bottom of the potential well, hence accreting on to a central gravitating body. Such a system, where a disc feeds a central object through accretion, under the effect of viscous forces, is called an {\\bf accretion disc}.  Accretion discs are found in very different contexts and over a wide range of physical scales. On the largest scale, they are one of the main physical ingredient to power the central engine of Active Galactic Nuclei (AGN), through accretion on to a supermassive black hole. The mass of the central black hole ranges from $10^6M_{\\odot}$ to $10^9M_{\\odot}$ and the accretion discs can extend out to large distances, of the order of about 1 pc, where rotating gaseous discs have often been detected through water maser emission \\cite{miyoshi95,greenhill97,kondratko06}. On the galactic scale, accretion discs around compact objects, such as white dwarfs, neutron stars and stellar mass black holes, are often found in binary systems, where the compact object (with a mass of the order of $1M_{\\odot}$) is fed by material outflowing from the companion. Historically, this has been one of the first contexts where accretion disc theory has found widespread application (a detailed review with an emphasis on accretion in galactic binary system and AGN can be found in the textbook by Frank, King and Raine \\cite{franck}). Another class of objects where accretion discs play an important role are Young Stellar Objects (YSO). Here, the central accreting object is a young protostar, which receives most of its mass from a surrounding protostellar discs. In this case, protostellar discs play also another key role, as the site where planet formation takes place. Understanding the dynamics of the disc and the development of the various instabilities that might occur in it, and possibly lead to turbulence, is therefore essential if one wants to understand some properties of our own Solar system, as well as those of extra-solar planets. A thorough description of the current observational state of protostellar discs can be found, for example, in the recent Protostars and Planets V book \\cite{PPV}.  As for the case of spiral galaxies, also for accretion discs the effects of the disc self-gravity can be very important. Indeed, also in this case the development of gravitational instabilities can lead to the formation of grand design spiral structures, which deeply affect the structure and the dynamics of the disc. The analysis of such instabilities naturally reveals several similarities between the case of spiral galaxies and that of accretion discs, but also some important differences. First of all, as mentioned above, accretion discs are fluid and, as such, they are intrinsically dissipative. Indeed, the evolution of gravitationally unstable accretion discs, as I will discuss below, is strongly dependent on the gas cooling rate. On the other hand, gravitationally unstable discs of stars, in the absence of any gas, have a natural tendency to dynamically heat up. Thermal energy, or disordered kinetic energy, plays a key role in stabilizing self-gravitating discs. However, while for stellar discs this energy is mostly in the form of disordered stellar motions, with the possibility of an anisotropic velocity dispersion, for gaseous discs the disordered kinetic energy content is mostly in the form of thermal energy. Furthermore, for gaseous discs it is expected that the role of resonances is much reduced with respect to a collisionless system. Finally, while for galaxy discs the disc mass can be a substantial fraction of the total mass of the system and therefore give a large contribution to the gravitational potential, in the case of accretion discs the disc mass is often much smaller than the mass of the central object (with some important exceptions, described below).  Historically, accretion disc theory has concentrated on the non self-gravitating case, the effects of self-gravity having only been discussed occasionally \\cite{pacinski78,kolikalov79,linpringle87,linpringle90}. In the last ten years, on the other hand, the important role of the disc self-gravity has been clearly recognized, partly due to improved observations, which have shown that in several observed systems (on all scales, from AGN to protostars) the disc mass can be high enough to have a dynamical role, and partly due to the increased computational resources, which have allowed a detailed numerical investigation of the development of gravitational instabilities in the non-linear regime. However, the early papers mentioned above already reveal the main directions around which research in this field would have later developed. Currently, most of the interest revolves around the three issues of self-regulation \\cite{pacinski78}, disc fragmentation \\cite{kolikalov79} and angular momentum transport induced by gravitational instabilities \\cite{linpringle87,linpringle90}. This again reveals a difference with respect to the galaxy discs case, where most of the attention has been traditionally given to the morphology induced by self-gravity.  In this paper, I present an overview of the dynamics of self-gravitating accretion discs. Accretion disc theory is a very vast topic and has been treated extensively elsewhere \\cite{franck,pringle81}, so I will not repeat it here. I will rather summarize the most salient general features, with particular emphasis on those features which are most affected by the disc self-gravity. Similarly, a detailed description of all the different astrophysical systems where self-gravitating accretion discs play a role (protostellar discs, AGN,...) is obviously beyond the scope of this review. The structure of the paper is as follows. In section \\ref{sec:dynamics} I describe the basic equations that determine the evolution of accretion discs and the main effects of the disc self-gravity. In section \\ref{sec:GI} I discuss the development of gravitational instabilities, and present some models which incorporate in a simple way the salient physics behind their development. In section \\ref{sec:simulations} I review the current state of numerical simulations of gravitationally unstable gaseous discs. In section \\ref{sec:transport} I focus on the transport properties induced by gravitational instabilities, while in section \\ref{sec:fragmentation} I describe the related issue of disc fragmentation. Finally, in the last two sections I will present some examples of observed systems where the concepts described in this paper find a natural application, and in particular, in section \\ref{sec:examples} I will discuss the impact of self-gravity in the dynamics of protostellar discs with some application related to planet formation. In section \\ref{sec:AGN} I focus on the impact of self-gravity in AGN discs also showing how self-gravitating accretion discs might have played an important role at high red-shifts, by allowing the formation of the seeds of supermassive black holes. ",
        "conclusions": "\\label{sec:conclusions}  As we have seen, self-gravitating accretion discs are an important aspect of the modeling of several systems, and thus play a key role in modern astrophysics. The influence of the disc self-gravity extends from the very small scales, associated with the formation of planetary systems, up to the large scales, associated with the formation and the feeding of supermassive black holes in the nuclei of galaxies. While clearly the specific properties of systems associated with such hugely different scales are going to differ substantially, there is yet a unifying framework related to the scale-free nature of the gravitational force.  Because of self-gravity, accretion discs can develop gravitational instabilities. These are going to deeply affect the structure and the evolution of the accretion disc. Firstly, we have seen that the instability usually takes the form of a regular spiral structure, that can efficiently redistribute the disc angular momentum and thus allow the accretion process to take place. In this respect the role of the disc self-gravity is complementary to that of other instabilities, such as magneto-hydrodynamic (MHD) instabilities \\cite{balbusreview,balbus03}. Indeed, while MHD instabilities are likely to be the most important source of angular momentum transport in the inner and hotter parts of the disc, it is not clear whether they are able to provide the necessary torques in colder environments, where the ionization level of the gas is low and the coupling to the magnetic field is weaker. This is, for example, the case of the outer parts of protostellar discs \\cite{fromang02}. On the other hand, such cold parts of the discs are naturally subject to gravitational instabilities. It is then possible to envisage a situation where gravitational instabilities are responsible for bringing the accreting material from large distances (which might be of the order of tens of AU for protostellar discs, and a fraction of a parsec for AGN discs) into the inner disc, where MHD instabilities take over and are responsible for the ultimate deposition of the matter onto the central object. In this respect, it would be important to develop models which include both the effects of MHD induced turbulence and that of gravitational instabilities. Semi-analytical models along these lines have been discussed \\cite{kratter07}, but it is still unclear what would be the interplay between magnetic and gravitational instabilities in the general case \\cite{fromang04a,fromang04b}.  A second aspect related to gravitational instabilities is the natural tendency associated with self-gravity to collapse and produce gravitationally bound objects. Disc fragmentation is a possible outcome of the onset of gravitational instabilities. As discussed above, it is now well established that the conditions leading to fragmentation are essentially related to the disc thermodynamics and to the cooling rate. From the modeling point of view, fragmentation has a twofold aspect. On the one hand, if one is interested in producing smaller objects within the disc, it offers a natural way to accomplish the task. On the other hand, there is the danger that most of the disc mass would end up in the fragments, with very little being accreted onto the central object. In this respect, fragmentation is a dangerous outcome, which might inhibit accretion.  We have seen that the biggest difference between protostellar discs and AGN discs with respect to self-gravity lies indeed in their thermal properties, and in particular in the fact that AGN discs are expected to be cold (in the sense that $H/R\\ll 1$) and with a short cooling timescale (in the sense that $\\tau_{\\rm cool}\\Omega\\ll 1$), while protostellar discs are hot ($H/R\\approx 0.1$) and with a long cooling timescale ($\\tau_{\\rm cool}\\Omega\\gg 1$, at least within 100 AU). The behaviour of these two kind of systems in this respect is then significantly different. Most protostellar discs are not expected to fragment, except perhaps in their outermost regions. This is the main problem affecting models of planet formation which rely on direct disc fragmentation. At the other end of the scales that we have considered, we see that AGN discs instead are likely to fragment. This is good news in view of forming young massive stars in the vicinity of the central supermassive black hole, as are observed, for example, close to the black hole in our Galactic Center. On the other hand, we have discussed above the problems that might arise from disc fragmentation in this context, related to the feeding of the black hole.  The last 10-15 years have witnessed a significant improvement of our understanding of the evolution of self-gravitating accretion discs. This has been made possible through the interplay between simple analytical models and the results of detailed numerical simulations of the disc hydrodynamics. The big advances in computing power in the last decade have made it possible to run such numerical simulations with unprecedented resolution, and progressively including the effects of several physical processes in more detail. In this paper, emphasis has been put on the important role of the gas cooling rate in determining the outcome of the instability. This has been often introduced in a simplified way in the simulations, which has led to a valuable investigation of the different regimes, in a somewhat academic way. The theoretical challenge in this respect is to introduce in the numerical codes a more realistic cooling function, in order to understand the behaviour of actual systems. For example, most of the current debate on the applicability of the fragmentation scenario for planet formation is ultimately due to our ignorance of the actual thermodynamics at work in protostellar discs. The difficulty here is to treat numerically the radiative transfer within very optically thick media. Only very recently have numerical codes implemented radiation physics and, although still preliminary, some results are starting to emerge \\cite{boley06,mayer07}. Codes which couple radiative transfer with the disc hydrodynamics are also desired to include the possibly important effect of irradiation from the central object, which has mostly been neglected in the simulations run so far (with some exceptions \\cite{cai06}).  If we look at the problem of star formation within AGN accretion discs, here the issue is ultimately whether accretion can proceed at significant rates even if the disc fragments. In other words, it is still unclear how much gas needs to be turned into stars before fragmentation is quenched. Here, it would be important to include in the models the energy feedback arising from the forming stars. Even in this context then, the key ingredient to be added into numerical codes is a proper treatment of radiative transfer, coupled with hydrodynamics."
    },
    "0801/0801.4373_arXiv.txt": {
        "abstract": "We discuss recent models on the evolution of massive stars at very low metallicity including the effects of rotation, magnetic fields and binarity. Very metal poor stars lose very little mass and angular momentum during the main sequence evolution, and rotation plays a dominant role in their evolution. In rapidly rotating massive stars, the rotationally induced mixing time scale can be even shorter than the nuclear time scale  throughout the main sequence. The consequent quasi-chemically homogeneous evolution greatly differs from the standard massive star evolution that leads to formation of red giants with strong chemical stratification. Interesting outcomes of such a new mode of evolution include the formation of rapidly rotating massive Wolf-Rayet stars that emit large amounts of ionizing photons, the formation of a long gamma-ray bursts and a hypernovae, and the production of large amounts of primary nitrogen. We show that binary interactions can further enhance the effects of rotation, as mass accretion in a close binary spins up the secondary. ",
        "introduction": "Massive stars affect the evolution of the early universe in a number of ways. They are believed to be the main source of reionizing photons at high redshift, and their explosions provide large amounts of energy into the surrounding medium. Most elements heavier than helium begin to exist in the early universe due to nucleosynthesis in massive stars. The history of star formation and the evolution of galaxies in the early universe are therefore closely related to such feedback from massive stars. This has motivated many theoretical studies on the evolution of massive stars of the first and second generations, which are characterized by zero/low metallicity.  Mass loss due to stellar winds from massive stars is mainly driven by metal lines (Castor, Abbott \\& Klein~1975; Pauldrach, Puls \\& Kudritzki R.P.~1986; Vink, de Koter \\& Lamers~2001). As a result, metal poor massive stars are expected to lose only a small amount of mass and angular momentum and to be kept rotating rapidly, during their life times (see Ekstr\\\"om and Meynet, however,  in this volume). Rotation -- one of the primary factors to determine the evolution of massive stars -- may thus play an even more important role in the evolution of massive stars in the early universe, compared to the case of galactic metal-rich massive stars. For instance, the high ratio of nitrogen to carbon abundance observed in extremely metal poor stars may be related to chemical mixing due to rotationally induced instabilities in massive stars in the early universe (Chiappini et al.~2006). Stellar cores could also more easily retain a large amount of angular momentum at lower metallicity, resulting in the abundant production of energetic supernovae and gamma-ray bursts (Yoon \\& Langer~2005; Woosley \\& Heger~2006; Yoon, Langer \\& Norman~2006).  Detailed numerical simulations of the evolution of massive stars including the effects of rotation involve many uncertain physical processes such as the transport of angular momentum and chemical species due to rotationally induced instabilities. Studies by Langer et al. (1999), Heger, Langer \\& Woosley (2000) and Hirschi, Meynet \\& Maeder (2004) indicate that their adopted angular momentum transport mechanisms (Eddington Sweet circulations, shear instability and baroclinic instability) are too inefficient to explain the observed spin rates of white dwarfs and young neutron stars: their models predict one or two orders of magnitude higher spin rates than the observed values. More recent models adopting the prescription by Spruit (2000) for the Tayler-Spruit dynamo in differentially rotating radiative layers are shown to be more consistent with observations in terms of the spin rates of stellar remnants (Heger, Woosley \\& Spruit~2005; Suijs et al. in prep.). The validity of the Talyer-Spruit dynamo has recently been challenged by several authors (Denissenkov, Pavel, \\& Pinsonneault~2007; Zahn, Brun \\& Mathis, S.~2007) but it seems clear that an efficient braking mechanism that is comparable to what Spruit suggests is needed to explain observations.  Here we present recent models of both single and binary massive stars at very low metallicity including the Tayler-Spruit dynamo as well as other effects of rotation such as rotationally induced chemical mixing, enhanced mass loss near the break-up velocity and tidal interactions. We suggest that rotation could lead to very different types of massive star evolution at very low metallicity, compared to the case of metal rich stars. Implications for supernovae and gamma-ray bursts, and massive star feedback in the early universe are briefly discussed. ",
        "conclusions": "We still do not fully understand the details about the effects of rotation such as transport processes of angular momentum and chemical species due to magnetic torques, Eddington Sweet circulations and other possible rotationally induced hydrodynamic instabilities in stars, which are important ingredients in recent stellar evolution models. Improved treatments of the rotation-related physics are thus needed for future studies, in connection with observational tests and multi-dimensional simulations (e.g. Brott et al. in this volume; Talon~2007). However, above discussions clearly indicate that the evolution of metal poor massive stars much depends on rotation. Addressing the role of the massive star feed back in the early universe including the effects of rotation and binarity will be an exciting but challenging subject for the next decade. In particular, future observational studies on massive star populations (e.g. WR/O ratio) in metal poor galaxies and supernovae/GRBs at high redshift may provide strong constraints for theoretical models, as discussed in Yoon, Langer \\& Norman~(2006) in greater detail."
    },
    "0801/0801.3286_arXiv.txt": {
        "abstract": "The observed properties of galaxies vary with inclination; for most applications we would rather have properties that are independent of inclination, \\emph{intrinsic} properties.  One way to determine inclination corrections is to consider a large sample of galaxies, study how the \\emph{observed} properties of these galaxies depend on inclination and then remove this dependence to recover the \\emph{intrinsic} properties.  We perform such an analysis for galaxies selected from the Sloan Digital Sky Survey which have been matched to galaxies from the Two-Micron All Sky Survey.  We determine inclination corrections for these galaxies as a function of galaxy luminosity and Sersic index.  In the $g$-band these corrections reach as as high as 1.2 mag and have a median value of $0.3$ mag for all galaxies in our sample. We find that the corrections show little dependence on galaxy luminosity, except in the $u$ band, but are strongly dependent on galaxy Sersic index.  We find that the ratio of red-to-blue galaxies changes from 1:1 to 1:2 when going from observed to intrinsic colors for galaxies in the range $-22.75 < M_K < -17.75$.  We also discuss how survey completeness and photometric redshifts should be determined when taking into account that observed and intrinsic properties differ.  Finally, we examine whether previous determinations of stellar mass give an intrinsic quantity or one that depends on galaxy inclination. ",
        "introduction": "In our search to understand the formation and evolution of galaxies, some of our primary tools are measurements of the distribution of galaxy properties and relationships among these properties.  From observations of these quantities at different redshifts we can deduce the nature of galactic evolution and by comparing them to the properties of dark matter halos we can constrain models of galaxy formation.  The measurement of galactic distributions includes, but is not limited to: the galaxy luminosity function \\citep{hubb:36}, the galaxy correlation function \\citep[the distribution of galaxies' spatial separations]{ph:74}, the galaxy velocity function \\citep{gonz:00}, and the distributions of galaxy sizes \\citep{chol:85}, surface brightnesses \\citep{free:70}, colors \\citep{baum:59,faber:73}, metalicities \\citep{oste:70} and star formation rates \\citep{td:80}. Relationships between galaxy properties include, the Tully-Fisher relation \\citep{tf:77}, the Faber-Jackson \\citep{fj:76} and fundamental plane \\citep{dd:87,dres:87} relations, the luminosity-size relation \\citep{korm:77}, the luminosity-metalicity relation \\citep{faber:73, lequ:79} and the density-morphology relation \\citep{dres:80}.  However, with the notable exception of the Tully-Fisher relation these distributions and relations are traditionally measured in terms of the {\\it observed} properties of galaxies. That is, the measurements used are K-corrected and corrected for foreground dust extinction, but no correction is attempted to compensate for the viewing angle from which the galaxies are observed.  In contrast, the Tully-Fisher relation is not a relationship between a galaxy's \\emph{observed} luminosity and rotation velocity, but a relation between a galaxy luminosity and rotation velocity \\emph{corrected for inclination}.  The inclination correction attempts to recover the intrinsic properties of a galaxy and not properties that are measured because of the particular angle from which the galaxy is viewed. Spiral galaxies are observed to have redder colors when their disks are more inclined, which is expected if the inclination increases the amount of dust that light traverses when emitted from the galaxy.  Clearly, we would prefer to measure all galactic distributions and relationships in terms of intrinsic galaxy properties instead of observed ones.  The comparison of theory to observations is complicated and often done incorrectly because of confusion between observed and intrinsic galaxy properties.  Early semi-analytic models were unable to match both the galaxy luminosity function and the Tully-Fisher relation in part because they failed to take into account that the first is observed luminosity while the second is intrinsic luminosity \\citep{sp:99}.  Also, when comparing galaxies at different redshifts we would like to be able to distinguish between evolution in their stellar populations and changes in their dust properties.  Furthermore, inclination effects are of great help in understanding the nature of dust in galaxies.  Theoretical modeling of attenuation in galaxies is complicated because it not only depends on the properties of dust, which seem to vary between galaxies, but also on how the dust is distributed and mixed with stars.  By determining the intrinsic properties of galaxies we also learn how those properties change as a function of galaxy inclination and therefore some properties of the dust distribution.  To determine intrinsic properties we need to know how a galaxy's properties change as a function of its inclination.  This is different then removing the effects of dust and dust is still present for a face-on galaxy.  There are a number of approaches for addressing this issue each of which has its own merits and disadvantages.  One approach is to solve for an inclination correction that minimizes the scatter in the Tully-Fisher relation \\citep[e.g.,][]{verh:01}, which assumes that the scatter in this relation should be as small as possible.  Another method is to fit stellar population models to the SED of a galaxy and then assume that any discrepancies are caused by dust \\citep[e.g.,][]{kauf:03}. A third is to observe background objects behind a foreground galaxy to get a direct measure of the extinction through the galaxy \\citep[e.g.,][]{berl:97,holw:05}, but this is difficult to do for more than a handful of cases.  Finally, one can simulate the radiative transfer through a galaxy \\citep[e.g.,][]{rocha:07}, assuming one knows the distribution and scattering properties of the dust.  The approach we explore here is somewhat simpler in that it assumes no knowledge of stellar population or dust properties.  Instead, the main assumption is that a galaxy's properties should be independent of inclination.  Thus any statistical correlation between a galaxy property and inclination can be attributed to dust and the inclination correction is whatever makes the observed correlation go away. This procedure has been applied a number of times \\citep{giov:94,giov:95,tully:98,mgh:03,shao:07}.  In this paper we greatly expand upon this method by applying it to 10,340 galaxies taken from the Sloan Digital Sky Survey \\citep[SDSS,][]{york:00} with accompanying infrared magnitudes from the Two-Micron All Sky Survey \\citep[2MASS,][]{skru:06}.  It is important to have galaxies with near infrared photometry because the effects of attenuation are minimized in these wavebands \\citep{bd:01}.  In a subsequent paper we will extend the analysis performed here to the full SDSS galaxy catalog.  We describe the method for determining inclination corrections in \\S\\ref{sec:method}. In \\S\\ref{sec:data} we describe the sample we will use and discuss some of the properties of galaxies in this sample.  In \\S\\ref{sec:results} we determine inclination corrections using our sample and compare our results to other determinations in \\S\\ref{sec:compare}.  In \\S\\ref{sec:changes} we discuss how consideration of intrinsic properties can change our conclusions about the distribution of galaxy properties focusing on the color-magnitude diagram.  We also comment on the effect on survey completeness, stellar masses and photometric redshifts. \\S\\ref{sec:conc} contains our conclusions and some discussion of future directions. ",
        "conclusions": "\\label{sec:conc}  This paper has focused on the intrinsic properties of galaxies, which are routinely determined for Tully-Fisher studies \\citep[e.g.,][]{piza:07}, but not for other statistical studies of galaxies.  Intrinsic properties are invariant under changes in viewing angle, unlike observed properties which would change if we could view a galaxy from a different vantage point.  Our main goal in this paper has been to clarify the difference between observed and intrinsic properties and to suggest how observed properties can be converted into intrinsic ones.  The method we use in this paper is to identify an observed galaxy property that shows a correlation with axis ratio and then apply the necessary correction to remove this correlation.  We find that both color and size show correlations with axis ratio. We are able to remove these correlations using simple linear formula that depend on $K$-band magnitude and Sersic index.  We therefore can construct a galaxy catalog with intrinsic value for galaxy size and magnitude.  There are many distributions and relationships that should be reconsidered in terms of the intrinsic properties instead of the observed ones.  We can not cover all of these in a paper of this length but we highlight a few points to suggest how things may change.  We focus on the color magnitude diagram as an example.  The observed color magnitude diagram shows a number of galaxies redder than the red sequence and a lack of bright blue galaxies.  When we plot the intrinsic color magnitude diagram these ultra red galaxies are mostly removed and there are many more bright blue galaxies.  We find that the ratio of blue to red galaxies changes from 1:1 to 2:1 for galaxies with absolute luminosities $-23.75 \\ge M_K > -17.75$, a significant change.  In a following paper, we apply the insights we have gained here to producing inclination corrections for the full SDSS catalog.  Having intrinsic properties for this flux limited catalog will then allow us to address how volume densities are effected when going from observed to intrinsic quantities."
    },
    "0801/0801.4529_arXiv.txt": {
        "abstract": "We perform a study of cosmic evolution with an equation of state parameter $\\omega(t)=\\omega_0+\\omega_1(t\\dot H/H)$ by selecting a phenomenological $\\Lambda$ model of the form, $\\dot\\Lambda\\sim H^3$. This simple proposition explains both linearly expanding and inflationary Universes with a single set of equations. We notice that the inflation leads to a scaling in the equation of state parameter, $\\omega(t)$, and hence in equation of state. In this approach, one of its two parameters have been pin pointed and the other have been delineated. It has been possible to show a connection between dark energy and Higgs-Boson. ",
        "introduction": "Cosmological research is mainly concerned with time (and in some cases space as well) evolution of various physical parameters like scale factor, Hubble parameter, matter-energy density etc. Along with these parameters, in recent years a new physical entity $\\Lambda$ has resurrected in the foreground of cosmology. In fact, $\\Lambda$ has become an essential part of the field equations of Einstein after some observational results \\citep{Riess1998,Perlmutter1999} indicated towards an accelerating Universe. It is believed by most of the physicists that the cosmological parameter $\\Lambda$ is responsible for driving the present acceleration because it can exert negative pressure. Moreover, due to some fine-tuning problem (known as cosmological constant problem), $\\Lambda$ is regarded as a variable quantity rather than a constant.  Now, in order to specify exact time-dependence of the unknown physical quantities including $\\Lambda$, one has to take recourse of a relationship between cosmic pressure $p$ and matter-energy density $\\rho$ involving the equation of state parameter $\\omega$. Mathematically speaking, one variable quantity can depend on the product of two other variable quantities. So,  one may construct $\\omega$ as a function of time, red-shift or scale factor \\citep{Chevron2000,Zhuravlev2001,Peebles2003}. In fact, values of $\\omega$ at different stages of cosmic evolution suggest that it may evolve with time. As an instance, for the present pressure-less Universe, the value of $\\omega$ is considered as zero, whereas its value was $1/3$ in the early radiation dominated Universe. However, it is convenient to consider $\\omega$ as a constant quantity because observational data can hardly distinguish between a varying and a constant equation of state \\citep{Kujat2002,Bartelmann2005}. Here some useful limits on $\\omega$ as appeared from SNIa data are $-1.67 <\\omega < -0.62$ \\citep{Knop2003} whereas refined values were indicated by the combined SNIa data (with CMB anisotropy) and galaxy clustering statistics which is $-1.33 < \\omega < -0.79$ \\citep{Tegmark2004}.  As stated above, $\\omega$ may have a functional relationship with scale factor or cosmological redshift. In connection to redshift it may depend linearly, $\\omega(z) = \\omega_o + {\\omega}^{\\prime} z$, where ${\\omega}^{\\prime}= (d\\omega/dz)_{z=0}$ \\citep{Huterer2001,Weller2002} or it may have a non-linear relationship as $\\omega(z) = \\omega_o + {\\omega}_1 z/(1+z)$ \\citep{Polarski2001,Linder2003}.  This suggests for a simple form \\begin{eqnarray} \\label{omegaq} \\omega(t)= \\omega_0 + \\omega_1 (t \\dot H/H), \\end{eqnarray} which  has got an explicit time dependence that disappears with the condition, $t\\dot H =H$.  Using above proposition, we explore the physical features of different stages of cosmic evolution, viz., linearly expanding and inflationary Universes. For this, a phenomenological $\\Lambda$ model is selected to solve the Einstein field equations. There are mathematically motivated works~\\citep{Ray2007,Mukhopadhyay2005,Mukhopadhyay2007a,Mukhopadhyay2007b}, wherein several phenomenological $\\Lambda$ models have been investigated for time-dependent $\\omega$. ",
        "conclusions": "We have discussed two Universes: (i) a linearly expanding Universe from its very beginning, (ii) and also the Universe like ours, which has gone through an inflation at its very early stage followed by a linear expansion later. We notice that these two kind of Universes, which  are direct consequence of our proposition (equation~\\ref{omegaq}), are represented by the same set of equations with a translational shift in the  equation of state parameter in the latter case compared to the former. In both the cases, $a(t)=1$ demands $E=1$, which applies a constraint on the equation of state parameter.  For the inflationary Universe, we have pin pointed $\\omega_0=-1/3$ and have delineated the other parameter with a range $-2/3 <\\omega_1< -0.46$. We observe that former is a special case of the latter with $\\omega_0=-1/3$ and $\\omega_1$=0. Any other value of $\\omega_0$ would invoke a non-linear behaviour in $a(t)$ through $E$. The  effect of the variation of $\\omega_0$ on $p$ is presented in Figure~\\ref{fig2} for a constant $\\omega=\\omega_0+\\omega_1=-0.80$ obtained by adjusting $\\omega_1$ accordingly. The $\\omega_1$ has nothing to do with $E$ and hence has nothing to do with $a(t)$. However, its value is a measure of translation in $\\omega$ due to inflation. The equations for $\\rho$ and $p$ involve $\\omega$ and hence would remain unchanged with its constant value. Thus, variations in curves of Figure~\\ref{fig2} is purely due to the variation in $\\omega_0$. The corresponding variations in $a(t)$ are shown in Figure~\\ref{fig3}.   A negligible value of $A$ is shown to be physically possible from the viewpoint of cosmology and particle physics, which means the absence of $\\Lambda$ in the field equations. So, both from physical and mathematical point of view the nullity of $\\Lambda$ is achieved for the same $\\Lambda$ model. Again, the expression of $q$ in this case has a striking similarity with that of~\\citet{Ray2007}. This work suggests that in the late phase of the Universe, where $t\\dot H=H$, the equation of state parameter behaves as a constant. Perhaps for this reason current data cannot distinguish clearly between a time-dependent $\\omega$ and a constant one as pointed out by some workers \\citep{Kujat2002,Bartelmann2005}.  Separating the entire cosmic history into two phases, it has been possible to derive the time-dependent expressions for the scale factor and the other physical parameters of each phase. It has been found that for inflationary phase, the deceleration parameter $q$ depends on time whereas for the linearly expanding phase it is constant, rather zero. This supports the opinion that $q$ has changed during the course of time~\\citep{Riess2001,Amendola2003,Padmanabhan2003}."
    },
    "0801/0801.3543.txt": {
        "abstract": "{T Tauri stars exhibit variability on all timescales, whose origin is still debated. On WTTS the variability is fairly simple and attributed to long-lived, ubiquitous cool spots.} {We investigate the long term variability of WTTS, extending up to 20 years in some cases, characterize it statistically and discuss its implications for our understanding of these stars.} {We have obtained a unique, homogeneous database of photometric measurements for WTTS extending up to 20 years. It contains more than 9 000 UBV R observations of 48 WTTS. All the data were collected at Mount Maidanak Observatory (Uzbekistan) and they constitute the longest homogeneous record of accurate WTTS photometry ever assembled.} {Definitive rotation periods for 35 of the 48 stars are obtained. Phased light curves over 5 to 20 seasons are now available for analysis. Light curve shapes, amplitudes and colour variations are obtained for this sample and various behaviors exhibited, discussed and interpreted.} {Our main conclusion is that most WTTS have very stable long term variability with relatively small changes of amplitude or mean light level. The long term variability seen reflects modulation in the cold spot distributions. Photometric periods are stable over many years, and the phase of minimum light can be stable as well for several years. On the long term, spot properties do change in subtle ways, leading to secular variations in the shape and amplitudes of the light curves.} ",
        "introduction": "The third catalog of pre-main-sequence emission-line stars assembled by G. Herbig and collaborators includes 742 objects (Herbig \\& Bell 1988). Of these, about fifty objects were classified as weak-line T Tauri stars (WTTS).These include objects from X-ray surveys of several star-forming regions and from a survey of stars with CaII H and K emis\\-sion lines.WTTS exhibit a weak, narrow $H\\alpha$ line [W($H\\alpha$) $<$ 10\\AA], a strong LiI absorption line (6707 \\AA) with W(Li) $>$ 0.1 \\AA, and a substantial CaII H and K emission lines.  The first photoelectric UBVR observations of these stars revealed periodic light variations in most WTTS (Rydgren \\& Vrba 1983; Rydgren et al. 1984; Bouvier et al. 1986, 1993, 1995; Grankin 1992, 1993, 1996; Vrba et al. 1993). Their photometric periods ranged from 1 to 10 days and were in agreement with spectroscopic estimates of their rotational velocities. It was assumed that the periodic light variations were due to the nonuniform distribution of cool photospheric spots over the stellar surface.  Rotation periods are now available for some hundreds of pre-main sequence and recently arrived main sequence stars of solar-like mass in five nearby young clusters: the Orion Nebula Cluster, NGC 2264, $\\alpha$ Per, IC 2602 and the Pleiades (Mandel \\& Herbst 1991; Attridge \\& Herbst 1992; Eaton et al. 1995; Edwards et al. 1993; Choi \\& Herbst 1996; Bouvier et al. 1997a). In combination with estimates of stellar radii these data allow us to construct distributions of surface angular momentum per unit mass at three different epochs: nominally, 1, 2 and 50 Myr (see Herbst \\& Mundt 2005). Recent work is extending the age range to 200 Myr and adding many new clusters (Irwin et al. 2007).  The study of long-term variations in the main param\\-eters of the light curves for young spotted stars is also of considerable interest. One would expect light-curve shapes to evolve as the spots change size, shape, temperature, or location, and a latitudinal migration of spots could cause changes in period, if the surfaces are in differential rota- tion, as in the Sun. It would, of course, be very interest- ing to find any cyclic pattern that could be interpreted as evidence for a magnetic cycle. Unfortunately, most of the photometric programs of observations of WTTS stars have been carried out episodically, in an interval from one to three months. Only one object (V410 Tau) has been in\\-vestigated in detail over several seasons (Vrba et al. 1988; Herbst 1989; Petrov et al. 1994). These authors showed that there were at least two extended and long-lived spot- ted regions on the surface of V410 Tau during several years. The evolution of the shape and amplitude of the light curve from season to season was explained by changes in the spot sizes, temperatures, and relative locations.  The first systematic photoelectric BVR observations for about thirty WTTS in dark clouds in Taurus-Auriga began in 1990 at the Maidanak Observatory (Grankin et al. 1995). During 1990-1993 rotation periods were dis- covered for 12 WTTS and refined for another 9 WTTS. A fundamental result of the analysis of the photometric behavior of these WTTS was the stability of the initial epochs and rotation periods for 17 WTTS on a time scale from 2 to 4 years. In addition, there is a high detection rate of rotation periods among WTTS in the sample con- sidered. It is generally believed that the stability of the initial epochs and rotation periods over several years in- dicates that the active region in each WTTS remains on a definite meridian over this time scale. Variations in the maximum brightness levels, the amplitudes and the shapes of the light curve for many WTTS is evidence for migra- tion of the spots, within the limits of an active zone, and for changes in their size. It is interesting to note, that RS CVn and BY Dra stars show the same kind of long- term photometric behaviour and that active longitudes have long been reported on the Sun (Losh 1939). For RS CVn stars, active longitudes appear to persist for a number of years (Butler 1996). Thus, an activity mechanism with some similarities to that operating on the Sun may be indicated by these common observational characteris- tics with WTTS.  Longer term observations have shown, that the phe- nomenon of stability of the initial epochs is most pro- nounced in seven WTTS: LkCa 4, LkCa 7, V410 Tau, V819 Tau, V827 Tau, V830 Tau and V836 Tau (Grankin 1997, 1998, 1999). Changes in the phase of minimum light for these stars did not exceed $\\pm 7\\%$ of the period on a time scale from 7 to 12 years. Similar results were obtained for 23 WTTS in the young Orion Nebula Cluster (Choi \\& Herbst 1996).  Some interesting results have been received from anal\\-ysis of long-term light curves of 30 WTTS and 80 RS CVn and FK Com stars (Grankin et al. 2002). These authors have found that in 18 of the 110 stars monitored, an in- crease in the amplitude of the periodic variability accom- panies a rise of the maximum brightness level. This result is difficult to understand in the context of a simple model where the star contains only a single spotted region in a high latitude zone that is always visible, because increas- ing the size of the spot (as required to increase the am- plitude) should decrease the maximum brightness of the star. It is possible to explain such behavior by invoking a star covered with mutiple spots at different latitudes. In this case the amplitude will not depend just on the total area of the spots, but will also depend on the distribu- tion of these spots on the stellar surface. Similar results have been obtained recently for 24 WTTS stars in the ex- tremely young cluster IC 348 over five observing seasons (Cohen et al. 2004). It has been shown that the stability of average magnitude from season to season most likely indi- cates that spots on these WTTS stars tend to redistribute themselves and undergo changes in size, temperature, and location rather than simply appear or disappear {\\it en masse}.  The observations to date support to the possible exis- tence of stellar activity cycles on WTTS. In an attempt to find and study such activity cycles we have contin- ued to monitor some tens of WTTS at Mt. Maidanak Observatory over many years. In this paper, we employ the unique and extensive Mt. Maidanak database to in- vestigate WTTS variability on a time scale of many years and, in some cases, multiple decades. In Section 2, we de- scribe the stellar sample and the observations carried out for the last 20 years at Mt. Maidanak. In Section 3, we dis- cuss the phenomenon of rotational modulation of the light curves of 36WTTS. In Section 4, we describe a spot model and some results of modeling. In Section 5, we define the statistical parameters used to characterize the long term variability of WTTS. Further modeling of the light curves will be presented in a later paper in this series. ",
        "conclusions": "Our main conclusion is that most WTTS have very sta- ble long term variability with relatively small changes of amplitude or mean light level. The long term variability seen merely reflects modulation in the cold spot distribu-tions. Photometric periods are stable over many years, as is the phase of minimum light for most objects, within some range. On the long term, the spot properties do change in subtle ways, leading to some variations in the shape and amplitudes of the light curves. The long term variability of WTTS is thus perhaps best understood by assuming that they are covered by a number of spots un- evenly distributed on the stellar surface and located at preferential longitudes and/or latitudes, much like active belts in the Sun and RS CVn systems, rather than result- ing from a single polar spot. Such an uneven spot distri- bution would account for the reported correlation between amplitude and brightness level, as well as for the quasi- sinusoidal shape of the light curves, which indicates that the unspotted photosphere is not seen in these stars at any rotational phase. It also implies that the coverage of the stellar photosphere by cold spots may be much larger than anticipated from simple modelling of the light curves, which only provide lower limits to the spot areal coverage. Doppler imaging of a sample of typical WTTS could be an effective means to obtain further constraints on the distribution of cold spots at the surface of WTTS.  Modelling of the light curves for V819 Tau, V827 Tau, V830 Tau and V836 Tau indicates that the spotted regions cover as much as 30 to 90\\% of the visible stellar hemisphere and that the mean temperature of the spots is 500--1400~K lower than the ambient photosphere. More detailed model calculations of the light curve of V410 Tau showed that:  (i) The decrease in mean brightness is attributable to the increase in total spot area from 47 to 53\\% and that it is essentially independent of the degree of nonuniformity in the spot distribution over the surface.  (ii) The amplitude of the phased light curve depends on the degree of nonuniformity in the spot distribution more strongly than on the star\u0092s total spot area. An in- crease in amplitude was accompanied by an increase in the degree of nonuniformity in the spot distribution over the stellar surface from 21 to 35\\%.  (iii) There is a weak correlation between the star's to- tal spot area and its degree of nonuniformity in the spot distribution with the correlation coefficient k=0.71. The increase in the star\u0092s total spot area from 47 to 53\\% was accompanied by a decrease in the degree of nonuniformity in the spot distribution from 35 to 21\\%.  The most active variable stars in our sample (i.e. those with the largest amplitudes) demonstrate stability of phase of minimum light most strongly over many years. The data presented here should prove useful in fur- ther investigations of the nature and evolution of spots on WTTS."
    },
    "0801/0801.2005_arXiv.txt": {
        "abstract": "We present a possible star formation and chemical evolutionary history for two early-type galaxies NGC~1407 and NGC~1400. They are the two brightest galaxies of the NGC~1407 (or Eridanus-A) group, one of the 60 groups studied as part of the Group Evolution Multi-wavelength Study (GEMS).  Our analysis is based on new high signal-to-noise spatially resolved integrated spectra obtained at the ESO 3.6m telescope, out to $\\sim$~0.6 (NGC~1407) and $\\sim$~1.3 (NGC~1400) effective radii. Using Lick/IDS indices we estimate luminosity-weighted ages, metallicities and $\\alpha$-element abundance ratios. Colour radial distributions from HST/ACS and Subaru Suprime-Cam multi-band wide-field imaging are compared to colours predicted from spectroscopically determinated ages and metallicities using single stellar population models. The galaxies formed over half of their mass in a single short-lived burst of star formation ($\\geq$~100~M$_{\\sun}$/year) at redshift z~$\\geq$~5. This likely involved an outside-in mechanism with supernova-driven galactic winds, as suggested by the flatness of the $\\alpha$-element radial profiles and the strong negative metallicity gradients. Our results support the predictions of the revised version of the monolithic collapse model for galaxy formation and evolution. We speculate that, since formation the galaxies have evolved quiescently and that we are witnessing the first infall of NGC~1400 in the group. ",
        "introduction": "The physical formation and evolution processes taking place at various locations within a galaxy leave their fossil imprint in the stars. These chemodynamical imprints are the regular features observed in the stellar populations, photometry, internal kinematics and structure of early-type galaxies. Furthermore, these features are found to correlate defining scaling relations of remarkable tightness, such as the colour-magnitude relation and the fundamental plane. This class of galaxy, therefore, is one of the best guides to compare observations and theoretical model predictions.  Theoretical scenarios that describe the formation and evolution of early-type galaxies include: a revised version of monolithic collapse, the dissipative (wet/gas-rich) and dissipationless (dry/gas-poor) merger alternatives of hierarchical clustering.  In the classical models of monolithic collapse (\\citealt{eggen62}; \\citealt{larson74a}; \\citealt{larson75}; \\citealt{carlberg84}; \\citealt{arimoto87}), stars form in all regions during a rapid collapse and remain in their orbits, whereas the gas dissipates to the centre of the galaxy. The shape and ellipticity of stellar orbits are expected to depend on the initial angular momentum of the collapsing protogalactic gas cloud and on the effect of gas viscosity in redistributing angular momentum, towards outer radii, during the collapse. The sinking gas is continuously enriched by evolving stars. As a consequence, stars formed in the centre are predicted to be more metal-rich than those born in the outer galaxy regions. The efficiency of this process is proportional to the depth of the galaxy potential well, so that a strong correlation between metallicity gradient and galaxy mass is expected. The evolving stars are also responsible for the $\\alpha$-element enrichment of the interstellar medium (ISM). To reproduce the supersolar abundance of $\\alpha$-elements with respect to the iron-peak elements, observed in early-type galaxies, the models require that the collapse and star forming process occurred on timescales $\\la$~1~Gyr, which will produce null or very small age gradients (\\citealt{arimoto87}; \\citealt{matteucci94}; \\citealt{thomas99}). Recently, feedback processes such as supernova-driven galactic winds have been re-considered in more detailed numerical models of dissipative collapse (\\citealt{larson74b}; \\citealt{arimoto87}; \\citealt{gibson97}; \\citealt{martinelli98}; \\citealt{chiosi02}; \\citealt{kawata03}; \\citealt{pipino06}; \\citealt{pipino07}). It is suggested that they may create and shape the abundance gradients. Galactic winds initiate when the energy injected by supernovae explosion balances the gravitational binding energy of the ISM. The winds evacuate the gas, eliminating the essential fuel for any future star formation. The shallower local potential well of the external galactic regions supports the early development of winds with respect to the central regions. Therefore, in the central regions the star formation and the chemical enrichment last longer. Dissipation of gas towards the galaxy centre and a time-delay in the occurrence of galactic winds cooperate in steepening any metallicity gradient. In summary, dissipative collapse models predict: i) very steep negative metallicity gradients, that strongly correlate with galaxy mass (\\citealt{chiosi02}; \\citealt{kawata03}), ii) small or null age gradients, as a consequence of the short timescales involved in the collapse and star formation processes, iii) $\\alpha$-element enhancement gradients that can be either positive, negative or null (\\citealt{martinelli98}; \\citealt{pipino06}), due to the ``\\textit{interplay between local differences in the star formation timescale and gas flows}\" (\\citealt{pipino07}), iv) isophotal shape and ellipticity proportional to the protogalactic initial angular momentum and its radial redistribution, v) no peculiar internal kinematic structure. The difference between the original and revised dissipative collapse models focuses on the assembly of the initial gaseous material. The revised view accounts for the possibility of either a unique primordial cloud or the coalescence of many gaseous clumps without any preexisting stars. Otherwise, the physical processes in galaxy formation are similar to that described above.  In the hierarchical clustering scenario of galaxy formation (eg. \\citealt{kauffmann93}; \\citealt{cole94}; \\citealt{baugh96}; \\citealt{kauffmann98}; \\citealt{delucia06}), elliptical galaxies are produced by merging events. Early numerical simulations of galaxy mergers predicted contradictory radial variations of stellar properties in merger remnants. \\cite{white80} suggested that dissipationless mergers cause a flattening of metallicity gradients, whereas \\cite{albada82} argued that the position of the stars in the local potential are preserved by violent relaxation and the gradients in the progenitors are only affected moderately. Further simulations of \\cite{barnes91}, considered the hydrodynamical physics of the gas and found that during a dissipative merger a significant fraction of the progenitors' gas tends to migrate toward the central regions of the merger remnant. \\cite{mihos94} showed that the gas accumulated in the central regions may trigger a local secondary burst of star formation. Galaxy merger-induced star formation is predicted to produce metallicity and age gradients. Positive or negative $\\alpha$-element enhancement gradients may also originate, depending on the original enhancement of the gas and the duration of the burst (\\citealt{thomasgre99}; \\citealt{thokauf99}). \\cite{bekki99} simulated dissipative mergers of two gas-rich disc galaxies of varying mass to form an elliptical galaxy remnant. Their models show that metallicity gradients in merger remnants are shallower with respect to those predicted by dissipative collapse models and only weakly dependent by the remnant galaxy mass. \\cite{kobayashi04} simulated the formation and chemodynamical evolution of elliptical galaxies in a cold dark matter cosmology, focusing on internal metallicity gradients. The models ranged from a monolithic collapse scenario, seen as the assembly of tens of gas-rich subunits at high redshift, to dissipative and dissipationless mergers of equal-mass galaxies at low redshift (which she denoted ``major mergers''). The result was that galaxies of a given mass had steep metallicity gradients if formed by collapse and shallower gradients if formed by major merger (dissipative or dissipationless).  Recently, combining the Millennium cosmological N-body simulation (\\citealt{springel05}) with semi-analytic models, \\cite{delucia06} and \\cite{delucia07} focused on the distinction between the formation and assembly time (redshift) of elliptical galaxies and brightest cluster galaxies (BCGs) in the hierarchical clustering scenario. The models make use of the prescription of feedback from a central active galactic nucleus (AGN) and supernova (see \\citealt{croton06}) to prevent excessive galaxy-mass growth from cooling flows. Their model predicts a mass-dependent evolutionary history (``down sizing'' scenario), with more massive galaxies (e.g BCGs, BGGs; brightest group galaxies) forming a large fraction of their stars (50~$-$~80 per cent) in progenitor systems at redshift $\\sim$~2.5. The progenitors then assemble into a single final object at z~$\\sim$~0.8. The formation epoch of less massive galaxies is predicted to be z~$\\sim$~1.9 and the assembly time at redshift $\\sim$~1.5. The peak of the star formation history of massive galaxies is at redshift $\\sim$~5, and it progressively moves towards lower redshifts for less massive systems. Similarly, the duration of the star forming episode is predicted to last longer in low mass galaxies. Mergers are thought to induce bursts of star formation. Therefore, the star formation history of a galaxy presents the ``bursty'' behaviour described in \\cite{delucia06}.  The kinematics and internal structure of a merger remnant are a consequence of the progenitors' mass ratios and the efficiency of the dissipative process (e.g. \\citealt{naab06}). \\cite{naab03} show that a dissipationless binary merger of equal-mass disk galaxies leads to a slowly rotating and anisotropic supported system ($(v/\\sigma)^{*}<$~0.4) with preferentially boxy isophotes and significant minor-axis rotation. Whereas, the merger of unequal-mass disk galaxies produces a rotationally supported elliptical galaxy with a small amount of minor-axis rotation ($(v/\\sigma)^{*}=$~1.2; suggesting isotropic oblate rotators) and disky isophotes. In dissipative mergers the gas is found to settle at the galaxy centre deepening the potential well and making it more axisymmetric. The isophotal shape of an equal-mass merger is then strongly affected by the gas, since most of the stars are gravitationally induced to move from boxy to more elliptical or even disky orbits. Similarly, the disky isophotal shape of an unequal-mass merger remnant is weakly affected by the presence of gas (\\citealt{naab06}). Furthermore, in simulations of an equal-mass merger with induced star formation (\\citealt{bekki97}) gradual gas dissipation forms an elliptical galaxy with disky isophotes, whereas a merger with rapid star formation produces both boxy and disky isophotes depending on the viewing angle of the observer. The merger of unequal-mass gas-rich disk galaxies is found to exhaust a large amount of the gas owing to moderate star formation and ending in the formation of an S0 galaxy (\\citealt{bekki98}). In summary, hierarchical galaxy formation models predict: i) shallow negative metallicity gradients, with little host mass dependence, ii) age gradients and iii) $\\alpha$-element enhancement gradients, that can be either positive or negative, due to the duration and location of the merger-induced star formation and the original abundance pattern of the gas, iv) isophotal shape  and ellipticity depending on the progenitors mass ratio and the efficiency of gas dissipation, v) peculiar internal kinematic structures.  The contrasting predictions of the galaxy formation models underline the lack of complete understanding of the physical processes behind the formation of galaxies. Observationally, we now have the ability to obtain spatially resolved high signal-to-noise integrated spectra out to large galactic radii (and hence larger mass fractions). This allows us to consider the physical mechanisms acting locally and how their properties vary with radius (i.e. investigating beyond the central regions), proving invaluable constraints on the processes of galaxy formation and evolution. Recently, an increasing number of works have focused their attention on stellar population radial profiles of galaxies in different environments (e.g. \\citealt{kobaari99}; \\citealt{mehlert03}; \\citealt{sanchez06}; \\citealt{sanchez07}; \\citealt{brough07}; \\citealt{reda07}). The results underline a common behaviour for the stellar population radial profiles of the studied galaxies: early-type galaxies are often characterised by strong metallicity gradients, shallow age gradients and null or statistically insignificant $\\alpha$-element enhancement gradients.  Following from \\cite{spola08a} (hereafter Paper~I) we probe a possible star formation and chemical evolutionary history of the early-type galaxies NGC~1407 and NGC~1400. We accomplish the task using the spectral Lick/IDS indices (\\citealt{faber85}; \\citealt{worthey94}) to determine the radial profiles in age, metallicity, [Z/H], and $\\alpha$-element abundance ratios, [$\\alpha$/Fe], out to $\\sim$~0.6 (NGC~1407) and $\\sim$~1.30 (NGC 1400) times the galaxies' effective radii ($r_{e}$; the radius within which half the galaxy light is contained). Central values and gradients in age, [Z/H] and [$\\alpha$/Fe] are estimated and used as proxies of the quantity, velocity and duration of gas dissipation, star formation and possible merger events during galaxy formation. Furthermore, we have considered the results from the spatially resolved radial kinematics and surface photometry analysis performed in Paper~I. We compared colour radial profiles with colours predicted from spectroscopically derived ages and metallicities using single stellar population models.  This paper is organised as follows. In Section~2 we describe the two sample galaxies. In Sections~3 and 4 we describe the observations and the relevant data reduction. Section~4 presents the spatially resolved stellar population parameters. In Section~5 observed and predicted colour index radial profiles are examined. In Section~6 we summarise our results, and discuss the possible star formation and evolutionary histories of NGC~1407 and NGC~1400. In Section~7 conclusions are presented. ",
        "conclusions": "We find that the stellar population of NGC~1407 is uniformly old, 12.0~$\\pm$~1.1~Gyr, with a supersolar degree of $\\alpha$-elements enhancement, 0.41~$\\pm$~0.05. We also measure a steep negative metallicity gradient of $-$0.38~$\\pm$~0.04 and a central metallicity of 0.29~$\\pm$0.08. The stellar population of NGC~1400 is found to be constantly old, 13.8~$\\pm$~1.1 Gyr, and largely enriched by $\\alpha$-elements, 0.33~$\\pm$~0.04. The central metallicity is 0.25~$\\pm$~0.06 and it steeply declines towards outer radii with a gradient of $-$0.47~$\\pm$~0.04.  The results of this work (and Paper~I) suggest a similar star formation and evolutionary history for the two galaxies. The outlined scenario is compatible with the revised version of the monolithic collapse model of galaxy formation and evolution. The two galaxies formed the bulk of their stars at z~$\\geq$~5. The star forming episode lasted for no longer than 1~Gyr with a star formation rate $\\geq$~100~M$_{\\sun}$/year. NGC~1407 might have experienced an ancient merger, as inferred by the possible kinematically decoupled core (Paper~I); nevertheless, the detection is uncertain and potentially originated by a misalignment of the slit with respect to the centre of the galaxy, as suggested by the $h_{3}$ radial profile. In Paper~I we have also claimed that NGC~1400 has not interacted with NGC~1407 or the group intergalactic medium. The flatness of the $\\alpha$-element radial profiles and the strong negative metallicity gradients are consistent with an outside-in formation scenario, where supernova-driven galactic winds played a fundamental role and shaped the slope of the gradients. The uniform radial profile of old ages, steep metallicity gradients and lack of signature of recent mergers indicate that the galaxies formed stars quickly some $\\sim$~12~Gry ago and then have evolved quiescently without any significant interaction ever since."
    },
    "0801/0801.0693_arXiv.txt": {
        "abstract": "We present the strategies adopted in the relative and absolute calibration of two different data sets: $U,B,V,I$-band images collected with the Wide Field Imager (WFI) mosaic camera mounted on the 2.2m ESO/MPI Telescope and $u,v,b,y$ \\calastrom images collected with the 1.54m Danish Telescope (ESO, La Silla). In the case of the WFI camera we adopted two methods for the calibration, one for images collected before 2002, with the ESO filters $U/38_{ESO841}$ and $B/99_{ESO842}$, and a different one for data secured after 2002, with the filters $U/50_{ESO877}$ and $B/123_{ESO878}$. The positional and color effects turned out to be stronger for images collected with the old filters. The eight WFI chips of these images were corrected one by one, while in the case of images secured with the new filters, we corrected the entire mosaic in a single step. In the case of the Danish data set, we compared point-spread function (PSF) and aperture photometry for each frame, finding a trend in both the X and Y directions of the chip. The corrections resulted in a set of first and second order polynomials to be applied to the instrumental magnitudes of each individual frame as a function of the star position. ",
        "introduction": "\\subsection{Observations and data reduction} Data have been retrieved from the ESO archive and include 8 $U$, 39 $B$, 51 $V$, and 26 $I$ images of \\calaomc. Data include both shallow and relatively deep images, with exposures times ranging from 1 to 300s for the $B, V, I$ bands, and from 300 to 2400s for the $U$ band, and were collected in several observing runs ranging from 1999 to 2003. During this period two filters were changed: data secured before 2002 were collected with the filters $U/38_{ESO841}$ and $B/99_{ESO842}$, while later ones with the filters $U/50_{ESO877}$ and $B/123_{ESO878}$. These data were obtained in good seeing conditions, and indeed the mean seeing ranges from ~0.6\" for the $I$ band to ~1.1\" for the $U$ band. We accurately selected the best PSF stars uniformly distributed across each chip. A moffat analytical function linearly variable on the chip was assumed for the PSF. The data were reduced with DAOPHOT {\\small IV}/ALLFRAME (Stetson 1994). \\subsection{Relative and absolute calibration} We are interested in providing accurate relative calibration of individual chips of the WFI mosaic camera because the occurrence of positional effects might cause a spurious color broadening of key evolutionary features such as the Red-Giant Branch (RGB) and the Main-Sequence Turn-Off. Moreover, we need to provide an accurate absolute calibration of our data in order to compare theory with observations. The plausibility of this comparison relies on the accuracy of the absolute zero-point of individual bands. This is a fundamental requirement for the distance modulus, and in turn for the absolute ages. Recent findings (Corsi et al. 2003; Koch et al. 2004) based on photometric data collected with the WFI indicate that the eight CCD chips might be affected by positional effects involving zero-point errors of the order of several hundredths of magnitude. This subtle effect could be due to scattered light. It seems that telescopes equipped with mosaic cameras and focal reducers may present this problem (Manfroid et al. 2001, Manfroid \\& Selman 2001). \\begin{floatingfigure}[l]{0.45\\textwidth} \\centering \\includegraphics[height=5.cm,width=5.5cm]{calamidaF1cut.ps} \\caption{Residuals, $\\Delta mag = mag_S - mag_{PSF}$, in the $V$ band, plotted versus the X, Y position on the chip, before (top) and after (bottom) the positional corrections were applied, for an image collected in 2002.} \\end{floatingfigure}  In order to correct the positional effects of the WFI mosaic camera and to perform an accurate absolute calibration of \\calaomc data set, we followed these steps:\\\\ $\\bullet$ obtain a set of local standard stars for \\calaomc;\\\\ $\\bullet$ identify all the trends of PSF magnitudes compared to standard magnitudes versus position on the frame and correct them;\\\\ $\\bullet$ estimate the calibration curves.\\\\ We thus made use of a set of new multi-band ($U,B,V,I$) local standard stars for \\calaomc (Stetson et al. 2007). This star list has been selected in photometric accuracy ($\\sigma \\leq$ 0.03 mag) and in 'separation index' ($sep \\geq$ 2.5). We ended up with a catalog of $\\sim 3\\times 10^4$ local standard stars. The sky area covered by these stars is $\\sim$ 37'$\\times$40', and includes a substantial fraction of our WFI data. We thus applied two methods: in the case of images collected before 2002 ($U/38_{ESO841}$ and $B/99_{ESO842}$ filters), the eight chips were corrected one by one, while for images secured starting from 2002 ($U/50_{ESO877}$ and $B/123_{ESO878}$ filters) we corrected the entire mosaic in a single step. The color terms for these filters have a very steep slope, therefore, we decided to estimate a first color curve and apply it. We then studied the residuals, $\\Delta mag = mag_S - mag_{PSF}$, where $S$ stands for Standard, as a function of the X and Y position on the chip (or mosaic, see Fig. 1). Once corrected for the positional effects, we estimated the color term once again and applied it for the absolute zero-point calibration. This strategy relies on the assumption that the two corrections are independent, and there is no reason why this should not be the case. Fig. 1 shows the magnitude residuals in the $V$ band plotted versus the X and Y positions on the frame, before and after the corrections were applied to an image collected in 2002. The trend with the position is clear, non-linear, and stronger in the four outermost chips. A small residual trend on the position is still present in the external regions, where the lack of standard stars makes it difficult to estimate the corrections.  \\begin{floatingfigure}[l]{0.45\\textwidth} \\centering \\includegraphics[height=5cm,width=5cm]{calamidaF2cut.ps} \\caption{Residuals, $\\Delta mag = mag_{AP} - mag_{PSF}$, in the $u$ band, plotted versus the X, Y position on the chip, before (top) and after (bottom) the positional corrections were applied.} \\end{floatingfigure}  We thus corrected the positional effects by fitting the $\\Delta mag$ vs $X/Y$ with second, third or fourth order polynomials, and we then estimated the calibration curves for each set. We adopted a first order polynomial to calibrate the $V$ and the $I$ bands as a function of the instrumental $V-I$ color. In the case of the $B$ band, we used a first order polynomial to calibrate the new filter, $B/123_{ESO878}$, and a third order one for the old filter, $B/99_{ESO842}$, as a function of the instrumental $B-V$ color. Particular attention has been paid to the $U$ band calibration, either in the case of the old filter, $U/38_{ESO841}$, as well in the case of the new one, $U/50_{ESO877}$. Having applied the positional corrections to each set, we estimated a first color curve, a fourth order polynomial for $U$ vs $U-I$. After applying this calibration curve, the $U$ magnitudes still showed a residual trend as a function of the $U$ magnitude and the $U-I$ color. We thus corrected these trends with first and sixth order polynomials, respectively. The reason of these very complicated trends with colors is probably due to the shape of the old $B$ and $U$ filters, which is quite different from the shape of the standard Johnson filters. ",
        "conclusions": ""
    },
    "0801/0801.2375_arXiv.txt": {
        "abstract": "We present integrated $JHK_s$ 2MASS photometry and a compilation of integrated-light optical photoelectric measurements for 84 star clusters in the Magellanic Clouds.  These clusters range in age from $\\approx200$~Myr to $>10$~Gyr, and have [Fe/H] values from $-2.2$ to $-0.1$ dex.  We find a spread in the intrinsic colours of clusters with similar ages and metallicities, at least some of which is due to stochastic fluctuations in the number of bright stars residing in low-mass clusters.  We use 54 clusters with the most reliable age and metallicity estimates as test particles to evaluate the performance of four widely used SSP models in the optical/NIR colour-colour space. All models reproduce the reddening-corrected colours of the old ($\\ge$ 10 Gyr) globular clusters quite well, but model performance varies at younger ages.  In order to account for the effects of stochastic fluctuations in individual clusters, we provide composite $B-V$, $B-J$, $V-J$, $V-K_s$ and $J-K_s$ colours for Magellanic Cloud clusters in several different age intervals. The accumulated mass for most composite clusters are higher than that needed to keep luminosity variations due to stochastic fluctuations below the 10\\% level. The colours of the composite clusters are clearly distinct in optical-NIR colour-colour space for the following intervals of age: $>10$~Gyr, $2-9$~Gyr, $1-2$~Gyr, and $200$~Myr$-1$~Gyr.  This suggests that a combination of optical plus NIR colours can be used to differentiate clusters of different age and metallicity. ",
        "introduction": "The most efficient method to determine the age and metallicity for unresolved stellar systems (especially at high redshift) is by comparing their observed colours with the predictions of evolution synthesis models \\citep[e.g.][]{bc93, bc03, worthey94, vazdekis99, maraston98, maraston05, af03}. Thus, it is important to test the integrated colours predicted by recent models, based on objects which have accurate ages and metallicities determined independently. In the present paper we focus our attention on the combination of visual and near-infrared (NIR) photometry, which has proven to be important for breaking the age-metallicity degeneracy, particularly in stellar populations older than $\\approx {\\rmn{a\\ few\\ times}} \\times 100$~Myr \\citep[e.g.][]{goudfrooij01, puzia02, hempel04}.  With the advent of the {\\it Spitzer Space Telescope (Spitzer)\\/} and mid-infrared (MIR) instrumentation for some large ground-based telescopes, the NIR spectral region is now accessible at intermediate-to-high redshifts. In a recent paper based on {\\it Spitzer\\/} Infrared Array Camera (IRAC) imaging, \\cite{vanderwel2006} reported significant discrepancies between some model predictions and the observed rest-frame $K$-band properties of early-type galaxies at z $\\approx 1$. Their results show that the interpretation of NIR photometry is hampered by model uncertainties. As a consequence the determination of masses of distant stellar systems based on such data can have uncertainties up to a factor 2.5 \\citep[see][]{bruzual07}.  Unfortunately, providing accurate model predictions in the near-infrared is challenging, since there are limitations imposed by the current lack of understanding of certain stages of stellar evolution (e.g., thermally pulsing asymptotic giant branch, or TP-AGB stars). These objects significantly affect the spectral energy distribution (SED) in the NIR and MIR for stellar populations with ages between $\\approx200$ Myr to 3 Gyr. Another possible complication is that the stellar libraries used by population synthesis models contain mostly stars from the solar neighborhood.  These stars have a star formation history which is not necessarily typical for extragalactic populations (e.g. relatively little variation of [$\\alpha$/Fe] ratios), and there are only a very limited number of AGB spectra available.  The Large and Small Magellanic Clouds (LMC and SMC respectively) provide a unique opportunity to test the accuracy of most current SSP models, since they contain a significant population of intermediate-age massive star clusters which are not easily accessible in our Galaxy. The ages and metallicities of these star clusters can be determined from deep colour-magnitude diagrams (CMDs) reaching below the main sequence turn-off (MSTO)\\footnote{Obtaining photometry with sufficient quality to secure reliable age and metallicity estimates for clusters in galaxies beyond the Magellanic Clouds requires a significant investment of observing time. To date only one such cluster,  SKHB~312 in M31, has a CMD deep enough to probe the MSTO region \\citep{brown2004}. The photometry for this object was obtained as a result of a program utilizing 126 {\\it Hubble Space Telescope (HST)\\/} orbits. }. Medium and high-resolution spectroscopy of individual giants in these clusters also provides independent metallicity estimates. Therefore their integrated-light properties (easily observed with small and moderate-aperture telescopes) can be combined with the accurate age/metallicity measurements and used to test (and calibrate) the SSP models.  In \\cite{pessev2006} (hereafter Paper~I) we used the Two Micron All Sky Survey (2MASS; \\cite{skrutskie2006}) to derive NIR $(JHK_s)$ integrated-light magnitudes and colours for a large sample of Magellanic Cloud star clusters, based on a homogeneous, accurately calibrated dataset. In the present study we use the sample from Paper~I and new photometry for 9 additional objects (forming the largest dataset of integrated NIR magnitudes and colours of LMC/SMC star clusters to date) to test the performance of several SSP models. We combine the 2MASS data with optical photometry originating from the work of \\cite{bica_et_al_96} and the compilation of \\cite{vdb81}. The technique adopted in Paper~I - measuring $JHK_s$ curves of growth to large radii allows us to utilize rather heterogeneous databases of optical photometry, usually performed with a set of fixed apertures. We use 54 clusters from our sample as ``test particles''. These clusters were chosen to have reliable age and metallicity measurements, covering a wide parameter space.  \\setcounter{figure}{0} \\begin{figure} \\centering \\includegraphics[bb=14 14 256 256,width=8.15cm]{fig1.eps} \\caption{ A finding chart of the LMC showing the clusters in our sample. The $R$-band image is centred on $\\alpha_{2000} = 05^{h} 26^{m} 37.7^{s}$ and $\\delta_{2000} = -68^{o} 56' 57.5\"$. The arrow near the centre outlines the position of NGC~1928 ($\\alpha_{2000} = 05^{h} 20^{m} 57.7^{s}$ and $\\delta_{2000} = -69^{\\rmn{o}} 28' 40.2\"$).  This cluster is located close to the geometrical centre of the LMC bar, and was adopted by Bica et al.\\ (1996) as a reference point for the relative coordinates of the LMC cluster system. The labeled arrows show the direction towards the clusters lying outside the boundaries of this chart. The $R$-band image (G. Bothun 1997, private communication) covers $8^{\\rmn{o}}$x$8^{\\rmn{o}}$, while the dimensions of the chart are $16^{\\rmn{o}}$x$16^{\\rmn{o}}$.} \\label{fig:lmc_map} \\end{figure}  \\setcounter{figure}{1} \\begin{figure} \\centering \\includegraphics[bb=14 14 256 256,width=8.15cm]{fig2.eps} \\caption{A finding chart of the SMC showing the clusters in our sample. The $R$-band image (G. Bothun 1997, private communication) is centred on $\\alpha_{2000} = 01^{h} 04^{m} 42.8^{s}$ and $\\delta_{2000} = -72^{\\rmn{o}} 52' 32.4''$. SMC clusters cover a smaller area than LMC objects. There is only one cluster outside the $R$ frame, but for illustrational purposes the dimensions of this finder chart are identical to those of the LMC chart in Figure~\\ref{fig:lmc_map}.} \\label{fig:smc_map} \\end{figure}  This paper is organized as follows: in \\S 2 we define our extended sample and present the new photometry along with the compilation of visual magnitudes and colours. Four sets of SSP models are tested in \\S 3, followed by concluding remarks in \\S 4. Information about the properties of the cluster sample is presented in Appendices A and B. Transformations between the model grids in the \\cite{bb88} system and the photometric system of 2 MASS are provided in Appendix C. \\label{intro} ",
        "conclusions": "We have presented new integrated $JHK_s$ 2MASS photometry for 9 Magellanic Cloud clusters, bringing our total sample (when combined with the results of Paper I) to 54 clusters with reliable ages $\\geq$200~Myr. In addition, we compile integrated-light $B$ and $V$ photometric measurements, extinction estimates, and a database of reliable age and metallicity determinations (mostly recent results) from the literature for our sample clusters.  We divide the clusters into different age (e.g., $\\geq10$~Gyr, $3-9$~Gyr, $3-4$~Gyr, $1-2$~Gyr, and 200~Myr$-$1~Gyr) and metallicity (when possible), and quantify the observed spread in the intrinsic cluster colours in these ranges. Care was taken to account for the spread of the observational data around the model predictions due to the stochastic fluctuations in the stellar populations of the clusters.  The smallest spread in intrinsic colours is found for clusters with ages $\\ga 10$~Gyr, the colours of which are well-reproduced by all four sets of SSP model predictions.  The systematic shift between the model predictions and the observed colours for a sample of old Milky Way globular clusters reported by \\cite{cohen07} is not observed in our Magellanic Cloud cluster analysis.  The largest spread in colour is found for clusters in the age range $2-4$~Gyr.  We believe that much of the spread in the colours for {\\it individual\\/} clusters younger than 10~Gyr results from stochastic fluctuations in the numbers of infrared-luminous stars, since individual clusters tend to have less than ${\\cal{M}}_{\\rm (10\\%)}$\\footnote{stellar mass needed to decrease the the luminosity uncertainty due to stochastic effects in the stellar population to 10\\%)} contributing to the observed colours.  Composite $(B-J)_0$, $(V-J)_0$, and $(J-K_s)_0$ cluster colours are calculated for each age/metallicity interval, and compared with the predictions of four widely used population synthesis models \\citep{maraston05, bc03, af03, vazdekis99}, in order to evaluate their performance.  We interpolate the model grids to calculate the offset or distance in colour-colour space between the model predictions and the age and metallicity for our composite cluster colours.  All four sets of models reproduce the colours of old ($\\geq10$~Gyr) Magellanic Cloud clusters quite well, with the \\citet{maraston05} and \\citet{bc03} models giving slightly better fits than the other two.  In the age range of 2\\,--\\,10 Gyr, the \\citet{maraston05} models have the largest separation in optical-NIR colour-colour space between the 2~Gyr and 10~Gyr model tracks, which best reproduces our observed composite colours in the $2-3$~Gyr and $3-9$~Gyr ranges.  While the composite colour for $2-3$~Gyr-old clusters falls just off the grid for the other three models, actual quantitative distances between the model predictions and composite cluster colours are not significantly different among the four models. In the $1-2$~Gyr and $0.2-1$~Gyr age ranges, the \\cite{bc03} models generally give the best quantitative match to our composite Magellanic Cloud cluster colours. Taking into account the inferred ages and metallicities, there is little difference between the \\cite{bc03} and \\cite{maraston05} model performance. The cluster colours fall off the \\citet{af03} and \\citet{vazdekis99} model predictions in the two youngest age ranges, largely due to their limited coverage at low metallicities.  Based on the comparisons presented in this work, it is found that each model has strong and weak points when used to analyse the optical$+$NIR colours of unresolved stellar populations. There is no model set that clearly outperforms the others in all respects. Overall, the \\cite{bc03}  and  typically yield the best quantitative match to our composite cluster colors. The \\cite{maraston05} models are a close second. The same two models also yield the best match to the composite cluster ages and metallicities."
    },
    "0801/0801.2972_arXiv.txt": {
        "abstract": "We present a sequence of toy models for irradiated planet atmospheres, in which the effects of geometry and energy redistribution are modelled self-consistently. We use separate but coupled grey atmosphere models to treat the ingoing stellar irradiation and outgoing planetary reradiation. We investigate how observed quantities such as full phase secondary eclipses and orbital phase curves depend on various important parameters, such as the depth at which irradiation is absorbed and the depth at which energy is redistributed. We also compare our results to the more detailed radiative transfer models in the literature, in order to understand how those map onto the toy model parameter space. Such an approach can prove complementary to more detailed calculations, in that they demonstrate, in a simple way, how the solutions change depending on where, and how, energy redistribution occurs. As an example of the value of such models, we demonstrate how energy redistribution and temperature equilibration at moderate optical depths can lead to temperature inversions in the planetary atmosphere, which may be of some relevance to recent observational findings. ",
        "introduction": "The past several years have seen a rapid and exciting increase in the amount of information available concerning the physical properties of extrasolar planets. In particular, the class of short orbital period planets, known informally as `Hot Jupiters', has begun to yield information through a variety of observational channels. The discovery of systems in which the orbital plane is almost edge-on to the line of sight has resulted in the detection of planetary transits (Charbonneau et al. 2000; Henry et al. 2000; Bouchy et al. 2005). Optical and ultraviolet photometry and spectroscopy of such systems can yield information about the absorbative properties of the planet atmosphere (Brown et al. 2001; Charbonneau et al. 2002; Vidal-Madjar et al. 2003; Vidal-Madjar et al. 2004). In addition, the unparalleled capabilities of the Spitzer Space Telescope has allowed the detection of thermal infrared emission (Charbonneau et al. 2005; Deming et al. 2005, Deming et al. 2006; Grillmair et al. 2007; Richardson et al. 2007) during secondary eclipses as well as the measurement of phase curves for non-transitting systems (Harrington et al. 2006).  This flurry of observational activity has spurred a similar level of activity on the theoretical front, with several independent models being used to infer the physical properties of the planets from the observations (Seager \\& Sasselov 2000; Barman, Hauschildt \\& Allard 2001; Cho et al. 2003; Menou et al 2003; Sudarsky, Burrows \\& Hubeny 2003; Burrows, Hubeny \\& Sudarsky 2005; Dyudina et al. 2005; Cooper \\& Showman 2005; Seager et al. 2005; Barman, Hauschildt \\& Allard 2005; Fortney et al. 2005; Fortney 2005; Burrows, Sudarsky \\& Hubeny 2006; Cooper \\& Showman 2006; Fortney et al. 2006; Langton \\& Laughlin 2007). The resulting inferences are not always entirely in agreement. As an example, consider each group's response to the first detection of secondary eclipses. Burrows et al. (2005) concluded that the redistribution of energy from the substellar to the antistellar face of the planet was weak. On the other hand, Barman, Hauschildt \\& Allard (2005), as well as Fortney et al. (2005), concluded that the redistribution was strong, while the opinion of Seager et al. (2005) fell somewhere in between. Similarly, predictions for the effects of atmospheric fluid motions on the emergent intensity pattern differ substantially from group to group (Cho et al. 2003; Cooper \\& Showman 2006; Langton \\& Laughlin 2007).  This diversity of opinion is hardly surprising, given the complexity of the problem. The study of Hot Jupiter atmospheres requires that we understand the atmospheric chemistry, radiative transfer and hydrodynamics of the planetary atmosphere. These are affected by the thermal structure and evolution of the planet, by the effects of photochemistry induced by ultraviolet light from the star, the possible existence and behaviour of atmospheric condensates, whether they form clouds, and whether (or if) energy is redistributed across the surface of the planet or reradiated in situ. Furthermore, there are a variety of different observational probes, each of which provides information but in ways that differ, sometimes subtly, from other methods. The complexity of the coupled radiative transfer, chemical equilibrium and atmospheric flow models also restrict somewhat the flexibility of the models in providing the true qualitative or geometric information necessary to develop a real understanding of what the observations are telling us.  The overall intent in the pages to follow is pedagogical. Clearly, detailed model atmospheres are necessary to fully interpret the emerging observations. However, simplified models that focus on key physical elements can be useful to try and develop a qualitative understanding of what elements of the model each of the various extant observations is actually probing. To do this, we construct a sequence of analytic and semi-analytic toy models. These models are chosen to be simple enough to be transparent while still capturing the qualitative features of the more detailed models. The hope is that one can use this approach to tailor our interpretations of the more detailed calculations to reach a proper synthesis between theory and the different observations that we have at our disposal. In \\S~\\ref{Nada} we review our simplest model, based on a variation of the grey atmosphere model. In this first case we will assume no redistribution of energy across the face of the planet. In \\S~\\ref{Apply} we then examine how the various available observations relate to the underlying structure of the model. In \\S~\\ref{Finally} we then introduce the physics of energy redistribution into the model and examine how this affects the observations in \\S~\\ref{MoreApps}. Finally, in \\S~\\ref{Finis}, we place our models within the context of both observations and more detailed models in the literature, and examine possible implications for future generations of models. ",
        "conclusions": "In summary, the intent of the models presented here was to understand the qualitative behaviour of radiative transfer in hot Jupiter atmospheres. The simplified nature of the models is dictated by the difficulties of including both full radiative transfer and detailed hydrodynamics, both of which are important parts of a proper understanding of the phenomenon. In particular we have studied how the models change depending on the parameters $\\gamma$ (which reflects the difference in opacity for incoming optical-wavelength radiation and outgoing infrared-wavelength radiation) and $\\tau_w$ (the depth at which energy is redistributed horizontally across the planets atmosphere).  We have also studied how the various new observational probes of Hot Jupiter atmospheres are influenced by changes in these parameters. We find that models in which redistribution occurs at large infra-red optical depths redistribute little energy to the night side (as expected). Perhaps more surprisingly, we find that significant day/night flux differences can also be found in the case $\\tau_w<1$, as long as $\\gamma < 1$ as well. We find that day-side temperature inversions are a generic feature of atmospheres in which the redistribution occurs at moderate ($\\tau_w \\sim 1$--1000) optical depths. This is of interest because of recent claims that temperature inversions have been found in HD~209458b (Knutson et al. 2007). The initial explanation for such behaviour is high altitude absorption (Burrows et al. 2007; Fortney et al. 2007), although the very low albedo (Rowe et al. 2007) restricts the absorbing species to be a poor reflector. Some level of redistribution is required to explain the HD~209458~b secondary eclipse amplitude and (lack of) phase variations, so that another potential explanation for the temperature inversion may simply be that it is related to the depth at which the energy is removed from the day-side. Certainly, figures~\\ref{BH6} and \\ref{BH7} show that some parts of the planet experience pronounced temperature inversions, even without a significant high altitude absorber. However, it will require detailed radiative transfer models to confirm whether this is a viable model or not.   Our initial exploration has  been far from exhaustive. One important advantage of these toy grey atmosphere models is their flexibility, and we have studied how energy redistribution changes the character of the models. This will hopefully provide a guide as to how one might perform similar calculations using the more detailed models, which presently can only consider changes at either the top or bottom of the atmosphere. Furthermore, we have also not exhausted the parameter space of toy models.  For instance, one can imagine many other redistribution schemes apart from the one we have implemented here."
    },
    "0801/0801.0977_arXiv.txt": {
        "abstract": "An excess over the extrapolation to the extreme ultraviolet and soft X-ray ranges of the thermal emission from the hot intracluster medium has been detected in a number of clusters of galaxies.  We briefly present each of the satellites (EUVE, ROSAT PSPC and BeppoSAX, and presently XMM-Newton, Chandra and Suzaku) and their corresponding instrumental issues, which are responsible for the fact that this soft excess remains controversial in a number of cases.  We then review the evidence for this soft X-ray excess and discuss the possible mechanisms (thermal and non-thermal) which could be responsible for this emission. ",
        "introduction": "\\label{Introduction}  The existence of soft excess emission originating from clusters of galaxies, defined as emission detected below 1~keV as an excess over the usual thermal emission from hot intracluster gas (hereafter the ICM) has been claimed since 1996.  Soft excesses are particularly important to detect because they may (at least partly) be due to thermal emission from the Warm-Hot Intergalactic Medium, where as much as half of the baryons of the Universe could be. They are therefore of fundamental cosmological importance.  Soft excess emission has been observed (and has also given rise to controversy) in a number of clusters, mainly raising the following questions: 1)~Do clusters really show a soft excess? 2)~If so, from what spatial region(s) of the cluster does the soft excess originate? 3)~Is this excess emission thermal, originating from warm-hot intergalactic gas (at temperatures of $\\sim 10^6$~K), or non-thermal, in which case several emission mechanisms have been proposed. Interestingly, some of the non-thermal mechanisms suggested to account for soft excess emission can also explain the hard X-ray emission detected in some clusters, for example by RXTE and BeppoSAX (also see \\citealt{petrosian2008} - Chapter 10, this volume; \\citealt{rephaeli2008} - Chapter 5, this volume).  Several instruments have been used to search for soft excess emission: EUVE, ROSAT and BeppoSAX in the 1990's, and presently XMM-Newton, Chandra and Suzaku. We will briefly present a history of these detections, emphasizing the difficulties to extract such weak signal and the underlying hypotheses, and summarising what is known on each of the observed clusters. For clarity, results will be presented separately for the various satellites. Finally we summarise and discuss the various models proposed to account for Extreme Ultraviolet (hereafter EUV) emission with their pros and cons. ",
        "conclusions": "\\label{conclusions}  While in most cases the origin of the soft excess emission is difficult to prove unambiguously, due to problems with instrument calibration and unknown background/foreground emission level, the independent detection of the phenomenon with many instruments gives confidence in the genuine nature of the phenomenon in a number of clusters. Both thermal and non-thermal emission mechanisms are probably at work in producing the soft excess emission in clusters."
    },
    "0801/0801.3659_arXiv.txt": {
        "abstract": "Since type Ia Supernovae (SNe) explode in galaxies, they can, in principle, be used as the same tracer of the large-scale structure as their hosts to measure baryon acoustic oscillations (BAOs). To realize this, one must obtain a dense integrated sampling of SNe over a large fraction of the sky, which may only be achievable photometrically with future projects such as the Large Synoptic Survey Telescope. The advantage of SN BAOs is that SNe have more uniform luminosities and more accurate photometric redshifts than galaxies, but the disadvantage is that they are transitory and hard to obtain in large number at high redshift. We find that a half-sky photometric SN survey to redshift $z = 0.8$ is able to measure the baryon signature in the SN spatial power spectrum. Although dark energy constraints from SN BAOs are weak, they can significantly improve the results from SN luminosity distances of the same  data, and the combination of the two is no longer sensitive to cosmic microwave background priors. ",
        "introduction": "Type Ia supernovae\\footnote{We consider only type Ia SNe in this \\emph{Letter}.} (SNe) have become a mature tool for studying the cosmic expansion history \\citep[e.g.,][]{phillips93,riess98, perlmutter99a}. A number of SN surveys, such as the Sloan Digital Sky Survey (SDSS) II \\citep{frieman08}, the Supernova Legacy Survey \\citep{astier06}, and the ESSENCE Supernova Survey \\citep{miknaitis07} are being carried out to improve the statistics and our understanding of the systematics. Moreover, the SN technique will be an integral part of almost every proposed dark energy survey including the Large Synoptic Survey Telescope\\footnote{See \\url{http://www.lsst.org}.} \\citep[LSST, see][]{tyson06} and the Joint Dark Energy Mission.  The conventional SN technique, measuring only the relative luminosity distance, $D_{\\rm L}$, is subject to degeneracies between cosmological parameters. For example, the SN constraint on the dark energy equation-of-state (EOS, $w$) parameter $w_{\\rm a}$, as defined by $w(z) = w_0 + w_{\\rm a}z/(1+z)$, is sensitive to the prior on the mean curvature of the universe \\citep*{linder05b, knox06c}. The reason is that the response of the relative distance to a variation in $w_{\\rm a}$ resembles that to a variation in the mean curvature and that the SN technique lacks the calibration of absolute distances \\citep{zhan06e}. Even for a flat universe with $w(z)\\equiv w_0$, the SN constraint on $w_0$ can be tightened considerably if the matter density is known to high precision \\citep{frieman03}. Such priors may come from other techniques, such as the cosmic microwave background (CMB), weak lensing, and baryon acoustic oscillations \\citep*[BAOs,][]{eisenstein98, cooray01b, blake03, hu03b, linder03b, seo03}. It has indeed been demonstrated that the latter three techniques are highly complementary to the SN $D_{\\rm L}$ technique \\citep{frieman03, seo03, knox06c}.  Since SNe explode in galaxies, their distribution bears the BAO imprint as well. To measure the SN spatial power spectrum, one needs the angular position and redshift of each SN, not its luminosity. Hence, the SN BAO technique does not suffer from uncertainties in the SN standard candle, which constitute the largest unknown in the $D_{\\rm L}$ measurements. Nevertheless, the narrow range of the  SN intrinsic luminosity reduces the effect of Malmquist-like biases and luminosity evolution. The SN rate traces the mass and star formation of the host galaxies with a time delay \\citep{sullivan06}. This means that SNe have a different clustering bias than galaxies that are selected by their luminosity or color. Finally, SNe have rich and time-varying spectral features for accurate estimation of photometric redshifts (\\phz{}s) \\citep*{pinto04, wang07a, wang07b}, which is helpful for measuring BAOs from a photometric survey.  There have been discussions of using the SN weak lensing magnification \\citep*{cooray06} and nearby SN peculiar velocities \\citep*{hannestad07} to probe the large-scale structure. We focus on photometric SN surveys for BAOs in this \\emph{Letter} and note in passing that the SN weak lensing technique is more limited by shot noise than the SN BAO technique \\citep{zhan06d, zhan06e} and that the SN peculiar velocity technique requires precise redshift and distance measurements.  For the BAO technique to be useful, one must survey a large volume at a sufficient sampling density as uniformly as possible. Although SN events are rare, the spatial density of SNe accumulated over several years will be comparable to the densities targeted for future spectroscopic galaxy BAO surveys. ",
        "conclusions": "We have demonstrated that \\phz{} SN data can be used to measure BAOs and to constrain cosmological parameters. The BAO constraints on the matter density and the baryon density are sensitive to the priors on the curvature and the scalar spectral index but not to the dark energy parameters. The dark energy constraints from the SN BAO technique alone are not meaningful. However, a combination of the SN BAO and $D_{\\rm L}$ techniques reduces the EP considerably, and the \\emph{Planck} priors are no longer crucial. The SN BAO results in Section~\\ref{sec:cons} are also applicable to \\phz{} galaxy BAOs. We note that \\phz{} errors are a large uncertainty for \\phz{} SN cosmology. Although our assumption about them is conservative compared to the results in \\citet{pinto04}, further studies are needed to make realistic forecasts.  Long and non-uniform cadence may result in uneven sampling of the SN spatial distribution. Fortunately, LSST will be likely to always catch the SNe (especially high-$z$ ones) at their maximum owing to its fast sky coverage and rapid sampling. Furthermore, the effect of the cadence on the SN depth can be simulated and determined, and methods of correcting for uneven depths in galaxy surveys can be applied to the SN data.  A spectroscopic SN BAO survey will be impractical, because one would have to revisit the sky many times spectroscopically over thousands of square degrees to catch the SNe that occur at different times. However, LSST will be able to obtain SNe in the millions over half the sky photometrically. This opens a window for applying the BAO technique to SNe and achieving more robust constraints with \\phz{} SN data. Since this SN BAO analysis requires no additional observations than doing the SN $D_{\\rm L}$ analysis alone, it should be a feature of all large-area SN cosmology analyses. Moreover, the SN data can also help calibrate the host galaxy \\phz{}s and the \\phz{} error distribution of other galaxies through the cross-correlation method \\citep{schneider06,zhan06d,newman08}."
    },
    "0801/0801.4286_arXiv.txt": {
        "abstract": "We present results from {\\it Hubble Space Telescope} ultraviolet spectroscopy of the massive X-ray and black hole binary system, HD~226868 = Cyg X-1.  The spectra were obtained at both orbital conjunction phases in two separate runs in 2002 and 2003 when the system was in the X-ray high/soft state. The UV stellar wind lines suffer large reductions in absorption strength when the black hole is in the foreground due to the X-ray ionization of the wind ions. A comparison of the {\\it HST} spectra with archival, low resolution spectra from the {\\it International Ultraviolet Explorer Satellite} shows that similar photoionization effects occur in both the X-ray high/soft and low/hard states.   We constructed model UV wind line profiles assuming that X-ray ionization occurs everywhere in the wind except the zone where the supergiant blocks the X-ray flux.  The good match between the observed and model profiles indicates that the wind ionization extends to near to the hemisphere of the supergiant facing the X-ray source. We also present contemporaneous spectroscopy of the H$\\alpha$ emission that forms in the high density gas at the base of the supergiant's wind and the \\ion{He}{2} $\\lambda4686$ emission that originates in the dense, focused wind gas between the stars. The H$\\alpha$ emission strength is generally lower in the high/soft state compared to the low/hard state, but the \\ion{He}{2} $\\lambda4686$ emission is relatively constant between X-ray states.  The results suggest that mass transfer in Cyg~X-1 is dominated by the focused wind flow that peaks along the axis joining the stars and that the stellar wind contribution from the remainder of the hemisphere facing the X-ray source is shut down by X-ray photoionization effects (in both X-ray states).  The strong stellar wind from the shadowed side of the supergiant will stall when Coriolis deflection brings the gas into the region of X-ray illumination. This stalled gas component may be overtaken by the orbital motion of the black hole and act to inhibit accretion from the focused wind.  The variations in the strength of the shadow wind component may then lead to accretion rate changes that ultimately determine the X-ray state. ",
        "introduction": "%  The massive X-ray binary Cygnus X-1 is a seminal target in the study of gas dynamics in the vicinity of a stellar mass black hole. Its X-ray luminosity and energetic jets \\citep{gal05} are powered by gas accretion from the nearby companion star HD~226868 (O9.7~Iab; \\citealt{wal73}) in a spectroscopic binary with a 5.6 day orbital period. There are several ways in which mass transfer from the supergiant to the black hole may occur in this system \\citep{kap98}. The O-supergiant, like other massive and luminous stars, has a strong radiatively driven wind that may be partially accreted through the gravitational force of the black hole.  The supergiant is large and is probably close to filling its critical Roche surface \\citep{gie86a,her95}, so a gas stream through the inner L1 point may also be present.  The actual gas flow in the direction of the black hole is probably intermediate between a spherically symmetric wind and a Roche lobe overflow stream, and there is evidence that the flow is best described as a focused wind \\citep{fri82,gie86b,gie03,mil05}.  The gas ions responsible for accelerating the wind may become ionized in the presence of a strong X-ray source, leading to a lower velocity, ``stalled'' wind \\citep{blo90,ste91}.  In situations of very high X-ray flux, photoionization may extend so close to the supergiant's photosphere that the wind never reaches the stellar escape velocity and thus ceases to become an X-ray accretion source \\citep{day93,blo94}. However, such a high X-ray flux may heat the outer gas layers to temperatures where the thermal velocities exceed the escape velocity to create a thermal wind that may fuel black hole accretion.  Important clues about the mass transfer process come from the temporal variations of the observed X-ray flux.  Cyg~X-1 is generally observed in either a low flux/hard spectrum state, with an X-ray spectrum that is relatively flat, or a high flux/soft spectrum state with a steeper power-law spectrum \\citep{sha06}.  The gamma-ray portion of the spectrum is also elevated during the high/soft state \\citep{mcc02}.  The high/soft state usually lasts for periods of days to months, and the fraction of time observed in the high/soft state has increased from $10\\%$ in 1996--2000 to $34\\%$ since early 2000 \\citep{wil06}.   This increase may be related to an overall increase in the supergiant's radius in the period from 1997 to 2003 -- 2004 that is suggested by changes in the long term optical light curve \\citep{kar06a}.   The system sometimes experiences so-called failed-state transitions, when it starts to increase in flux, but then stops at an intermediate state and returns to the low/hard state.  All these transitions probably reflect changes in the inner truncation radius of the accretion disk surrounding the black hole that are caused by a variable accretion rate (largest when the system is in the high/soft state; \\citealt{don02,mcc06}). Thus, the temporal variations in the X-ray state offer us the means to compare the black hole accretion processes with observational signatures related to mass transfer.  The H$\\alpha$ emission formed in the high density gas at the base of the stellar wind is an important diagnostic of the mass loss rate in massive stars \\citep{pul96,mar05}. The H$\\alpha$ emission variations in HD~226868 over the last few years are documented in independent spectroscopic investigations by \\citet{gie03} and \\citet*{tar03}.  Both of these studies concluded that the H$\\alpha$ emission appears strongest when the system is in the low/hard X-ray state, while a range of weak to moderate emission strengths are observed during the high/soft states.  This is a surprising result, since taken at face value, strong emission is associated with a large wind mass loss rate, and the simplest expectation that the X-ray accretion flux increases with mass loss rate is, in fact, not observed.  There are several possible explanations: (1) A denser wind may be more opaque to X-rays.  However, this seems unlikely because the observed inverse relation between H$\\alpha$ emission strength and X-ray flux is observed at all orbital phases, not just when the supergiant and its wind are in the foreground. (2) The X-ray source may photoionize and heat the gas responsible for the H$\\alpha$ emission, so that a larger X-ray flux leads to a decrease in H$\\alpha$ strength.  This clearly occurs at some level, but both \\citet{gie03} and \\citet{tar03} argue that portions of the wind shaded from the ionizing flux also display significant temporal variations.  (3) Changes in the X-ray flux will lead to variations in the ionized volume of gas surrounding the black hole, and consequently, the total acceleration of the wind in the direction towards the black hole will vary with the distance traveled before the atoms responsible for line-driving are ionized.  Thus, a stronger, denser wind might reach a faster speed before ionization,  and since the Bondi-Hoyle accretion rate varies as $\\sim v^{-4}$, the gas captured by the black hole (and the associated X-ray flux) declines.  This last process can be tested through direct study of the degree of wind ionization observed in the ultraviolet P~Cygni lines formed in the supersonic part of the wind outflow.   When the system is observed with the ionization region in the foreground, the absorption cores of these P~Cygni lines will be truncated at a blueshift corresponding to the highest projected speed before encountering the ionization zone, the so-called Hatchett-McCray effect \\citep{hat77}.  The binary is so faint in the ultraviolet that high dispersion spectra were very difficult to obtain with the {\\it International Ultraviolet Explorer (IUE)} satellite \\citep{dav83}, but a good series of observations were obtained with {\\it IUE} at lower spectral resolution that clearly indicate the weakening of the wind lines when the black hole is in front \\citep*{tre80,vlo01}.  Most of these spectra were obtained in the low/hard X-ray state (\\S5), when according to the varying wind strength model the mass loss rate is higher and the wind is less ionized and faster.  We embarked on a new program of high S/N, high dispersion UV spectroscopy to test this hypothesis with the {\\it Hubble Space Telescope} Space Telescope Imaging Spectrograph (STIS). We obtained observations at the two orbital conjunction phases in both 2002 and 2003.  These sets of observations were both made during the rare high/soft X-ray state, and planned observations during the low/hard state were unfortunately scuttled by the STIS electronics failure in 2004.  However, we can rebin the high quality {\\it HST} spectra made during the high/soft state to the lower resolution of the {\\it IUE} archival spectra (mostly low/hard state) in order to test whether or not the wind ionization state does in fact differ significantly between states.   We describe a program of supporting optical spectroscopy we have obtained to check the orbital phase (\\S2) and wind strength (\\S3) at the times of the {\\it HST} observations.  We compare the H$\\alpha$ measurements with the contemporaneous X-ray light curve recorded with the {\\it Rossi X-ray Timing Explorer} All-Sky Monitor instrument \\citep{lev96} and confirm our earlier result showing how H$\\alpha$ tends to strengthen as the X-ray flux declines (\\S3). We then describe the observed variations in the main UV wind lines and present a simple ``shadow wind'' model for the profiles (\\S4). We compare the variations observed in the {\\it HST} and {\\it IUE} spectra in \\S5, and then reassess the question of the mass transfer process in \\S6.  We will discuss the photospheric features in the UV spectra in a forthcoming paper (Caballero Nieves et al., in preparation). ",
        "conclusions": "%  The X-ray accretion flux of Cyg X-1 is fueled by mass transfer from the supergiant.  We argue in this section that the mass transfer process is dominated by a wind focused along the axis joining the stars.  However, the accretion of this gas by the black hole may be influenced by the strength of the radiatively driven shadow wind that is directed away from the black hole.  We begin by reviewing the most pertinent observational results from this investigation, and then we consider the interplay between the dynamics of the wind outflow and the X-ray accretion flux.  First, the {\\it HST} STIS spectra of HD~226868 that we obtained at two epochs when the system was in the high/soft state show dramatic variations in the wind line strength that result from a superionization of the gas atoms illuminated by the X-ray flux. Shadow wind models, in which the wind ions only exist in the region where X-rays are blocked by the supergiant, make a reasonably good match to the observed profile variations, so we suspect that X-ray photoionization dominates much of the zone between the black hole and the facing hemisphere of the supergiant. A similar degree of wind ionization probably also exists in the X-ray low/hard state since similar orbital variations in wind line strength are found in {\\it IUE} low dispersion spectra made during the X-ray low/hard state.  Second, the {\\it HST} spectra suggest that stellar wind gas emanating from parts of the photosphere facing the X-ray source attains only a small velocity before becoming photoionized.  For example, the highest optical depth wind feature, \\ion{C}{4} $\\lambda\\lambda 1548, 1550$, shows only a very modest P~Cygni absorption core at phase $\\phi=0.5$ (see Fig.~11) that extends blueward no more than about $-400$ km~s$^{-1}$ (and it is possible that this small component results from a minor part of shadow wind projected against supergiant at $\\phi=0.5$; see Fig.~8). This very low wind speed is probably less than the stellar escape velocity ($\\sim 700$ km~s$^{-1}$ near the poles).  These results from the UV wind lines indicate that very little mass loss is occurring by a radiatively driven wind for surface regions that are exposed to the X-ray source.  The fact that the wind features appear similarly weak in {\\it IUE} spectra obtained in the low/hard state suggests that a spherical, radiatively driven wind from the hemisphere of the supergiant facing the black hole is probably always weak or absent, and thus, accretion from a spherically symmetric wind must play a minor role in feeding the black hole in Cyg~X-1.  On the other hand, we found that the emission equivalent width of the \\ion{He}{2} $\\lambda 4686$ line is consistently strong between X-ray states and over the available record of observation. The orbital phase variations of this spectral feature \\citep{gie86b,nin87} are successfully matched by models of emission from an enhanced density and slower gas outflow region between the stars that is expected for a focused stellar wind \\citep{fri82}. Thus, while X-ray ionization reduces the wind outflow away from the axis joining the stars, the X-ray flux is apparently insufficient to stop the outflow in the denser gas of the focused wind (the result of both tidal and radiative forces). Consequently, it is this focused wind component that is probably the primary means of mass transfer in Cyg~X-1. We caution that the relative constancy of the \\ion{He}{2} $\\lambda 4686$ emission flux does not necessarily imply that the mass loss rate in the focused wind is also steady.  For example, the increased X-ray photoionization during the high/soft state may lead to an increase in \\ion{He}{2} $\\lambda 4686$ emission (see Fig.~2 in \\citealt{gie86b}), so that a lower mass loss rate but higher ionization fraction might result in the same amount of observed emission.  However, the presence of the \\ion{He}{2} $\\lambda 4686$ emission in both X-ray states indicates that focused wind mass loss always occurs at some level.  Finally, we confirm that the H$\\alpha$ P~Cygni line forms mainly in the base of the stellar wind of the supergiant since we observe that H$\\alpha$ follows the orbital velocity curve of the supergiant (Fig.~1).  The new observations are consistent with earlier results \\citep{gie03,tar03} in demonstrating that the H$\\alpha$ emission strength is generally weaker in the high/soft X-ray state.  Photoionization and heating may extend down to atmospheric levels where the gas densities are sufficient to create H$\\alpha$ emission, so that the reduction in H$\\alpha$ strength in the X-ray high/soft state may partially result from photoionization related processes.  However, \\citet{gie03} showed that H$\\alpha$ emission variability was present in those Doppler shifted parts of the profile corresponding to the X-ray shadow hemisphere of the supergiant (see their Fig.~15), so part of the H$\\alpha$ variations must be related to gas density variations at the base of the stellar wind.  Thus, the observed H$\\alpha$ variations suggest that the high/soft X-ray state occurs when the global, radiatively driven part of the wind is weaker.  Long term, quasi-cyclic variations in wind strength are apparently common among hot supergiants \\citep{mar05}.  \\citet{gie03} and \\citet{tar03} suggested that the variations in X-ray state are caused by changes in wind velocity due to changes in supergiant mass loss rate.  During times when the supergiant's wind is denser and the mass loss rate is higher, the photoionization region would be more restricted to the region closer to the black hole.  Consequently, a radiatively driven wind could accelerate to a higher speed before stalling when the gas enters the ionization zone, and thus, the faster wind would result in a lower black hole accretion rate and X-ray luminosity (creating the low/hard X-ray state).  Conversely, if the wind mass loss rate drops, then the X-ray ionization zone will expand, the maximum wind velocity towards the black hole will decline, and the net accretion rate will increase (perhaps creating the high/soft state; \\citealt{ho87}).  This creates a positive feedback mechanism that may continue until the wind is ionized all the way down to the stellar photosphere facing the X-ray source \\citep{day93,blo94}.  If this scenario is correct, then we expect that the outflow velocities in the direction of the black hole (as measured in blue extent of the P~Cygni lines at phase $\\phi=0.5$) will be larger than the supergiant escape velocity during the X-ray low/hard state.  The superb quality {\\it HST}/STIS spectra indicate that outflow velocities are too low to launch the wind during the X-ray high/soft state. Moreover, the low resolution {\\it IUE} spectra from the low/hard state appear to show a very similar pattern of the loss of the P~Cygni absorption at $\\phi=0.5$ (Fig.~12 and 13), indicating that significant ionization zones still exist in the low/hard state. Taken at face value, these {\\it IUE} results suggest that the spherical component of wind outflow towards the black hole is weak and slow in both X-ray states, so a wind speed modulation is probably not the explanation for the accretion variations associated with the X-ray states.  We will require new, high quality, UV spectroscopy of Cyg~X-1 during the low/hard state in order to make a definitive test of this idea.  The radiatively driven wind of the supergiant leads to effective mass loss only in the X-ray shadowed hemisphere and in the focused wind between the stars in Cyg~X-1.  The outflow in the shadow wind region will experience a Coriolis deflection, so that the trailing regions of the shadow wind will eventually enter the zone of X-ray illumination \\citep{blo94}. Once photoionized, this gas will stall with the loss of the important ions for radiative acceleration, and some of this slower gas may extend around the orbital plane to the vicinity of the black hole.  Although this deflected wind gas is probably not a major accretion source \\citep{blo94}, it may affect the accretion dynamics of the focused wind.  For example, when the shadow wind mass loss rate is high (times of strong H$\\alpha$ emission), the resulting stalled wind component will create a higher ambient gas density on the leading side of the zone surrounding the black hole. The focused wind flow will make a trajectory towards the following side of the black hole, and while gas passing closer to the black hole will merge into an accretion disk, gas further out will tend to move past the black hole before turning into the outskirts of the disk.  The presence of the stalled gas on the leading side may deflect away this outer, lower density part of the flow and effectively inhibit gas accretion from the focused wind. The subsequent reduction in gas accretion by the black hole may correspond to the conditions required to produce the low/hard X-ray state, while conversely a reduction in the stalled gas from the shadow wind may promote mass accretion and produce the high/soft state \\citep{bro99,don02,mcc06}. Clearly, new hydrodynamical simulations are needed to test whether the stalled wind component is sufficient to alter the accretion of gas from the focused wind and create the environments needed for the X-ray transitions."
    },
    "0801/0801.4415_arXiv.txt": {
        "abstract": "{ We show that the entropy created by Ohmic dissipation inside an accreting charged black hole may exceed the Bekenstein-Hawking entropy by a large factor. If the black hole subsequently evaporates, radiating only the Bekenstein-Hawking entropy, then the black hole appears to destroy entropy, violating the second law of thermodynamics. A companion paper discusses the implications of this startling result. Bousso's covariant entropy bound is not violated. } ",
        "introduction": "The purpose of this paper is to show that classical processes of dissipation can generate huge quantities of entropy inside the horizon of a black hole, many orders of magnitude more than the Bekenstein-Hawking \\cite{Bekenstein:1973ur} entropy. The specific black hole model presented is intended to be semi-realistic, albeit over-simplified, with parameters appropriate to a real supermassive black hole. We take charge as a surrogate for angular momentum, and electrical conductivity as a surrogate for angular momentum transport. To see how much entropy might be created, we treat the electrical conductivity as an adjustable free parameter.  If a black hole creates many times the Bekenstein-Hawking entropy and subsequently evaporates, radiating only the Bekenstein-Hawking entropy and leaving no remnant, then entropy is destroyed, violating the second law of thermodynamics. The implications of this startling result are discussed in a companion paper \\cite{Polhemus:2009}.  Throughout this paper we treat entropy in a purely classical fashion. In particular, we assume that locality holds inside the black hole. Locality, the quantum field theory proposition that operators commute at spacelike-separated points, is the assumption that normally makes it legitimate to add entropy over spacelike surfaces. Since the spacetime curvature inside a supermassive black hole is well below Planck, except near the singularity, one might expect classical physics to apply.  It is widely thought that in order to preserve unitarity of black hole evaporation, locality must break down over spacelike surfaces connecting the inside and outside of a black hole \\cite{Susskind:1993if}. In the companion paper \\cite{Polhemus:2009} we argue that the calculations of the present paper point to a profligate breakdown of locality inside black holes.       We work in Planck units, $k_B = c = G = \\hbar = 1$. ",
        "conclusions": "\\label{summary}  We have shown that the dissipation of the free energy of the electric field inside a charged black hole can potentially create many times more than the Bekenstein-Hawking entropy. If the black hole subsequently evaporates, radiating only the Bekenstein-Hawking entropy and leaving no remnant, then entropy is destroyed.  This startling conclusion is premised on the assumption that entropy created inside a black hole accumulates additively on spacelike slices, which in turn derives from the assumption that the Hilbert space of states is multiplicative over spacelike-separated regions, as postulated by locality. This is essentially the same reasoning that originally led Hawking \\cite{Hawking:1976ra} to conclude that black hole evaporation is non-unitary, that black holes must destroy information.  It is widely thought that unitarity should be considered a higher principle than locality. To ensure that black hole evaporation is unitary, locality between the inside and outside of a black hole must break down. In a companion paper \\cite{Polhemus:2009} we argue that the gross violation of the second law found in the present paper points to a wholesale breakdown of locality inside black holes, and provides a compelling argument in favor of the conjecture of ``observer complementarity''.  The black hole respects Bousso's covariant entropy bound \\cite{Bousso:2002ju}, as it should given the theorem of \\cite{Flanagan:1999jp}."
    },
    "0801/0801.4309_arXiv.txt": {
        "abstract": "We present a new method for the reconstruction of the longitudinal profile of extensive air showers induced by ultra-high energy cosmic rays. In contrast to the typically considered shower size profile, this method employs directly the ionization energy deposit of the shower particles in the atmosphere.  Due to universality of the energy spectra of electrons and positrons, both fluorescence and Cherenkov light can be used simultaneously as signal to infer the shower profile from the detected light. The method is based on an analytic least-square solution for the estimation of the shower profile from the observed light signal. Furthermore, the extrapolation of the observed part of the profile with a Gaisser-Hillas function is discussed and the total statistical uncertainty of shower parameters like total energy and shower maximum is calculated. ",
        "introduction": "The particles of an extensive air shower excite nitrogen molecules in the atmosphere, which subsequently radiate ultraviolet fluorescence light isotropically. This fluorescence light signal can be measured with appropriate optical detectors such as the fluorescence telescopes of HiRes~\\cite{fd:HiRes}, the Pierre Auger Observatory~\\cite{fd:Auger} or the Telescope Array~\\cite{fd:TA}.  The number of emitted fluorescence photons is expected to be proportional to the energy deposited by the shower particles. Recent measurements of the fluorescence yield in the laboratory confirm this expectation within the experimental uncertainties~\\cite{Kakimoto:1995pr,Waldenmaier:2007am,airflydEdX}. Non-radiative processes of nitrogen molecule de-excitation lead to a temperature, pressure and humidity dependence of the fluorescence yield (see e.g.~\\cite{airflyAtm}). For atmospheric parameters of relevance to the reconstruction of air showers of ultra-high energy cosmic rays, the pressure dependence of the ionization energy deposit per meter track length of a charged particle is almost perfectly canceled by the pressure dependence of the fluorescence yield (see, for example, \\cite{Keilhauer:2005nk}). Therefore only a weak pressure and temperature dependence has to be taken into account if the number of emitted photons is converted to a number of charged particle times track length, as has been done in the pioneering Fly's Eye experiment \\cite{Baltrusaitis:1985mx}. The reconstructed longitudinal shower profile is then given by the number of charged particles as function of atmospheric depth.  The approximation of assuming a certain number of fluorescence photons per meter of charged particle track and the corresponding expression of the longitudinal shower development in terms of shower size are characterized by a number of conceptual shortcomings. Firstly the energy spectrum of particles in an air shower changes in the course of its development. A different rate of fluorescence photons per charged particle has to be assumed for early and late stages of shower development as the ionization energy deposit depends on the particle energy \\cite{Song:1999wq}. Secondly the tracks of low-energy particles are not parallel to the shower axis leading to another correction that has to be applied \\cite{Alvarez-Muniz:2003my}. Thirdly the quantity ``shower size'' is not suited to a precise comparison of measurements with theoretical predictions. In air shower simulations, shower size is defined as the number of charged particles above a given energy threshold $E_{\\rm cut}$ that cross a plane perpendicular to the shower axis. Setting this threshold very low to calculate the shower size with an accuracy of $\\sim 1$\\% leads to very large simulation times as the number of photons diverges for $E_{\\rm cut}\\rightarrow 0$. Moreover, the shower size reconstructed from data depends on simulations itself since the shower size is not directly related to the fluorescence light signal.  These conceptual problems can be avoided by directly using energy deposit as the primary quantity for shower profile reconstruction as well as comparing experimental data with theoretical predictions. Due to the proportionality of the number of fluorescence photons to the energy deposit, shower simulations are not needed to reconstruct the total energy deposit at a given depth in the atmosphere. Another advantage is that the calorimetric energy of the shower is directly given by the integral of the energy deposit profile \\cite{linsleydEdX}. Furthermore the energy deposit profile is a well-defined quantity that can be calculated straight-forwardly in Monte Carlo simulations and does not depend on the simulation threshold \\cite{Risse:2003fw}.  Most of the charged shower particles travel faster than the speed of light in air, leading to the emission of Cherenkov light. Thus, in general, the optical signal of an air shower consists of both fluorescence and Cherenkov light contributions. In the traditional method \\cite{Baltrusaitis:1985mx} for the reconstruction of the longitudinal shower development, the Cherenkov light is iteratively subtracted from the measured total light. The drawbacks of this method are the lack of convergence for events with a large amount of Cherenkov light and the difficulty of propagating the uncertainty of the subtracted signal to the reconstructed shower profile.  An alternative procedure, used in \\cite{Abbasi:2004nz}, is to assume a functional form for the longitudinal development of the shower, calculate the corresponding light emission and vary the parameters of the shower curve until a satisfactory agreement with the observed light at the detector is obtained. Whereas in this scheme the convergence problems of the aforementioned method are avoided, its major disadvantage is that it can only be used if the showers indeed follow the functional form assumed in the minimization.  It has been noted in~\\cite{cher:giller} that, due to the universality of the energy spectra of the secondary electrons and positrons within an air shower, there exists a non-iterative solution for the reconstruction of a longitudinal shower profile from light detected by fluorescence telescopes.  Here we will present an analytic least-square solution for the estimation of the longitudinal energy deposit profile of air showers from the observed light signal, in which both fluorescence and Cherenkov light contributions are treated as signal. We will also discuss the calculation of the statistical uncertainty of the shower profile, including bin-to-bin correlations. Finally we will introduce a constrained fit to the detected shower profile for extrapolating it to the regions outside the field of view of the fluorescence telescope. This constrained fit allows us to always use the full set of profile function parameters independent of the quality of the detected shower profile.  \\begin{figure*} \\begin{center} \\includegraphics[clip,width=0.95\\textwidth]{sketch.eps} \\caption{Illustration of the isotropic fluorescence light emission (solid circles), Cherenkov beam along the shower axis (dashed arcs) and the direct (dashed lines) and scattered (dotted lines) Cherenkov light contributions.}\\label{fig1} \\end{center} \\end{figure*} ",
        "conclusions": "In this paper a new method for the reconstruction of longitudinal air shower profiles was presented. With the help of simulations we have shown that the least square solution yields robust and unbiased results and that uncertainties of shower parameters can be reliably calculated for each event.\\\\ Events with a large Cherenkov light contribution are currently usually rejected during the data analysis (see for instance \\cite{Abbasi:2004nz,AbuZayyad:2000ay}) However, as we have shown, there is no justification for rejecting such showers, once experimental systematic uncertainties are well understood. Because events with a large Cherenkov contribution have different systematic uncertainties to those dominated by fluorescence light, both event classes can be compared to study their compatibility.\\\\ At energies below 10$^{17.5}$ \\eV{}, where new projects \\cite{heat,tale} are planned to study the transition from galactic to extragalactic cosmic rays, events with a large fraction of direct Cherenkov light will dominate the data samples, because the amount of light, and thus trigger probability, of these events is much larger than that of a fluorescence dominated shower. If at these energies it is still possible to measure an accurate shower geometry, the fluorescence detectors should in fact be used as Cherenkov-Fluorescence telescopes.\\\\"
    },
    "0801/0801.4637_arXiv.txt": {
        "abstract": "This paper addresses the origin of the silicate emission observed in PG QSOs, based on observations with the {\\it Spitzer Space Telescope}. Scenarios based on the unified model suggest that silicate emission in AGN arises mainly from the illuminated faces of the clouds in the torus at temperatures near sublimation. However, detections of silicate emission in Type 2 QSOs, and the estimated cool dust temperatures, argue for a more extended emission region. To investigate this issue we present the mid-infrared spectra of 23 QSOs. These spectra, and especially the silicate emission features at $\\sim 10$ and $\\sim 18$ $\\mu$m, can be fitted using dusty narrow line region (NLR) models and a combination of black bodies. The bolometric luminosities of the QSOs allow us to derive the radial distances and covering factors for the silicate-emitting dust. The inferred radii are 100-200 times larger than the dust sublimation radius, much larger than the expected dimensions of the inner torus. Our QSO mid-IR spectra are consistent with the bulk of the silicate dust emission arising from the dust in the innermost parts of the NLR. ",
        "introduction": "Unified schemes for active galactic nuclei (AGN) postulate an obscuring torus surrounding an accreting super-massive black hole. Models predict that the infrared spectral energy distribution (SED) of the torus depends sensitively on its orientation, geometry and density distribution \\citep[e.g.][]{pier92,granato94,efstathiou95,granato97,nenkova}. In particular, the tori are predicted to exhibit prominent silicate dust features in either absorption or emission, depending on whether an AGN is viewed with the torus edge-on (Type 2) or face-on (Type 1). Previous failures to detect strong 9.7$\\mu$m silicate emission in Type 1 AGN led to several proposed modifications of the unified model. For example modified grain size distributions have been assumed \\citep{laordrain1993,maiolino} or a clumpiness of the torus invoked \\citep{nenkova}.  The {\\it Spitzer Space Telescope} ({\\it Spitzer}), with its good mid-infrared (mid-IR) wavelength coverage and sensitivity, has drastically changed our view of this problem. \\citet{siebenmorgen05a} and \\citet{hao05} reported the first \\spitzer\\ Infrared Spectrograph (IRS) detections of prominent silicate emission features in the mid-infrared spectra of several luminous quasars. Sturm et al. (2005) reported the first detection of 10 and 18 $\\mu$m silicate emission features in a low-luminosity LINER (NGC 3998). Comparison to the 10/18 $\\mu$m feature ratio of optically thin emission from silicate dust at different temperatures suggests a modest temperature ($\\sim$200K) of the emitting dust.  The presence of prominent 10$\\mu$m silicate emission features in AGN covering a broad range in luminosity could be taken as direct evidence for the existence of an obscuring torus. However, it is not at all clear whether this emission actually arises from the inner regions of a face-on torus. Depending on the size and composition of the grains, sublimation occurs between $\\sim$~800 and 1500 K \\citep{kimura} which is also the temperature range expected for silicate dust located near the hot inner torus wall.  The lower temperature indicated for the emitting silicate dust can be interpreted as evidence for dust emission from regions located further away from the central heating source. Several arguments support such a scenario. Silicate emission has also been detected in Type-2 QSOs \\citep{sturm2006,teplitz06}, whereas for an edge-on view of the torus, one would expect to see silicate in absorption only. An extended emitting region, with dimensions much larger than the inner torus dimension, is fully consistent with this result. Broad-band $10\\mu m$ imaging of several nearby AGN suggests extended mid-infrared continuum \\citep{cameron93,tomono,bock00,radomski03,packham05}. \\citet{efstathiou06} has modeled the silicate emission of the Type-2 QSO IRASF10214+4724 \\citep{teplitz06} invoking extended NLR dust in addition to an AGN torus. \\citet{marshall07} conclude that some of the optically thin warm emission in the QSO PG0804+761 may emerge from regions beyond the torus and suggest clouds in the NLR as a possible origin of this emission. These arguments indicate that silicate emission may originate in extended regions ($\\sim 100$ pc dimension).    The nature and location of the extended silicate-emitting region is not yet known. In this paper we explore in a quantitative way one plausible interpretation, namely the association of this cool dust with the NLR. To this end, we present fits of our QSO spectra with a superposition of NLR dust models and spectral components representing the innermost hot dust and the bulk of the inner structure emission (both related to the torus) as well as the large scale host emission. The fitted model and the bolometric luminosity of each source enable us to estimate the cool dust distance and its covering factor. We note that our models cannot exclude the possibility of a torus contribution to the observed silicate emission.  In section \\ref{sec:Obs} we describe our QSO sample. In section \\ref{sec:model}, we introduce the model components and detail our fitting procedure. We also describe how we estimate the silicate dust cloud distances and the related covering factors. In section \\ref{sec:robustness}, we discuss the dependence on model parameters. In section \\ref{sec:results}, we present the results of the fits, which are then discussed in section \\ref{sec:discussion}. Finally section \\ref{sec:conclusions} summarizes our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  Scenarios based on the unified model suggest that silicate emission in AGN arises mainly from the illuminated faces of the clouds in the torus at temperatures near sublimation. However, detections of silicate emission in Type 2 QSOs, and the estimated cool dust temperatures, argue for an origin in a more extended region. To investigate this issue, we have presented the \\spitzer-IRS spectra of 23 QSOs. We have matched physically-based models to the mid-infrared spectra and found that the silicate emission observed in these objects can be reproduced by emission from clouds, outside the central torus. This extended silicate-emitting region is possibly associated with the innermost NLR region or the intermediate dusty region proposed by \\citet{Netzer93}.  The dust cloud distances found here scale with the AGN luminosity as $R_{dust}\\sim 80\\cdot L_{bol46}^{0.5}$ pc, with the bolometric luminosity $L_{bol46} $ given in units of $10^{46}$ $erg s^{-1}$. We have estimated the median distance of the dust cloud responsible for the silicate emission to be 40 pc, while for individual sources distances up to 260 pc are possible. The smallest cloud distance is 9 pc for PG 0050+124. The calculated covering factors for the dust clouds have a median of 0.16 and are in agreement with an NLR origin of the silicate emission. Our models do not exclude the possibility of a torus contribution to the observed silicate emission, but rather emphasize the good agreement and perhaps the necessity of a larger-scale contribution to this emission. Finally, future high spatial resolution infrared observations and further crosschecks including the comparison to Type 1/2 ratios as well as distributions and individual values for the obscuring columns in X-rays and the near infrared are needed to resolve the remaining ambiguities between torus and NLR emission."
    },
    "0801/0801.3684_arXiv.txt": {
        "abstract": "Kant and Laplace suggested the Solar System formed from a rotating gaseous disk in the 18th century, but convincing evidence that young stars are indeed surrounded by such disks was not presented for another 200 years. As we move into the 21st century the emphasis is now on disk formation, the role of disks in star formation, and on how planets form in those disks.  Radio wavelengths play a key role in these studies, currently providing some of the highest spatial resolution images of disks, along with evidence of the growth of dust grains into planetesimals.  The future capabilities of EVLA and ALMA provide extremely exciting prospects for resolving disk structure and kinematics, studying disk chemistry, directly detecting proto-planets, and imaging disks in formation. ",
        "introduction": "The search for forming planetary systems is ultimately a search for our origins.  How do the gas and dust in molecular clouds evolve into rocky and gas giant planets, and how common are planetary systems like the Solar System?  Radial velocity searches for extra solar planets indicate that the Solar System might actually be unusual, but such searches are only just achieving the sensitivities needed to detect the Solar System-like planetary systems.  For several reasons it is much easier to study the disks from which planetary systems are expected to form, and this may shed light on the environment and processes taking place in the early Solar System.  The formation of disks at some stage during the star formation process is made inevitable by one of several ``angular momentum problems'' in star formation.  Specific angular momentum is imparted to molecular clouds by differential Galactic rotation, and many orders of magnitude must be lost before a planetary system can form.  The outcome of the star formation process for solar-type stars is therefore a T Tauri star, descending the Hayashi track towards the main sequence in the HR diagram, surrounded by a disk.  Massive stars reach the main sequence while still deeply embedded and accreting.  A combination of radiative heating by the central star, accretion, and heating by the interstellar radiation field results in temperatures for disks on Solar System size scales of 10 to 20~K for Solar-type stars. The bulk of the disk mass is therefore best traced by millimeter and submillimeter wavelength emission.  Indeed, the first convincing evidence for the ubiquity of disks surrounding T Tauri stars arose from a single-dish survey for 1.3~mm continuum emission by \\citet{chandler:Beckwith1990}.  Although those measurements were not able to resolve the emission the flux densities detected, and corresponding dust column densities, implied that the dust had to be distributed in flattened structures in order to explain the low optical extinctions to the central stars.  The rest of this review focuses on the role of radio interferometry in spatially resolving the emission from circumstellar disks, and in understanding the disk properties. ",
        "conclusions": "Over the last few years, submm, mm, and cm-wave observations of circumstellar disks around low-mass, pre-main sequence stars have demonstrated evidence for grain growth to cm-sized particles, a necessary precursor to planet formation.  Keplerian velocity fields in those disks have enabled independent measurements of central stellar masses, and are able to constrain pre-main sequence evolutionary tracks. Deviations from Keplerian rotation have been observed in some cases, and raise the possibility that low-mass companions may be inferred from such observations.  Evidence for angular momentum transport in disk winds has been demonstrated by the first observations of rotation in stellar jets, in a sense consistent with that of the accompanying disk.  Masers are now able to trace disk and wind interactions for a massive protostar, suggesting that disk-mediated accretion may be a mechanism for forming high-mass stars.  The future of disk studies is very bright; ALMA and EVLA will transform this field by providing images of gaps in disks, detecting dust heated by proto-planets, revealing the early phases of disk formation, and through detailed studies of the gas chemistry."
    },
    "0801/0801.3367_arXiv.txt": {
        "abstract": "In this paper we give a review of the Bowen fluorescence survey, showing that narrow emission lines (mainly N\\,III and C\\,III lines between 4630 and 4660 \\AA) appear to be universally present in the Bowen blend of optically bright low mass X-ray binaries. These narrow lines are attributed to reprocessing in the companion star giving the first estimates of $K_2$, and thereby providing the first constraints on their system parameters.  We will give an overview of the constraints on the masses of the compact objects and briefly highlight the most important results of the survey.  Furthermore, we will point out the most promising systems for future follow-up studies and indicate how we think their estimates of the component masses can be improved. ",
        "introduction": "One of the main aims of optical observations of low-mass X-ray binaries (LMXBs) has always been to find a signature of the donor star, and thereby constrain the masses of both components. Unfortunately, the optical emission of most LMXBs that accrete at high rates is completely dominated by the reprocessing of the X-rays in the outer accretion disk. This is the reason why, in spite of optical counterparts being known for persistent LMXBs for years (sometimes even over 20 years), no strong constraints on their system parameters existed until recently.  \\begin{figure*}[t] \\parbox{8.0cm}{\\psfig{figure=sco.ps,width=8.0cm}} \\parbox{8.0cm}{\\psfig{figure=doppler.eps,angle=-90,width=8.0cm}} \\caption{{\\it Left:} Trailed spectrogram of the Bowen blend and He\\,II $\\lambda$4686 of Sco\\,X-1 showing the presence and the movement of the narrow emission lines. From \\citet{sc02}. {\\it Right:} Example Doppler maps of  He\\,II $\\lambda$4686 (left) and the Bowen region (right) for 4U\\,1636$-$536 (top) and 4U\\,1735$-$44 (bottom). The He\\,II maps trace the accretion disk, while the Bowen maps are dominated by a compact spot at the phasing and velocity where the companion star is expected. The Roche lobe of the secondary and the gas stream trajectory are overplotted for clarity. From \\citet{ccs06}. \\label{scox1}} \\end{figure*}  The discovery of narrow high-excitation emission lines in the bright LMXB Sco\\,X-1 opened new opportunities for system parameter constraints in active X-ray binaries \\citep{sc02}.  Phase-resolved blue spectroscopy showed that these narrow lines moved in anti-phase with the compact object, strongly suggesting that they arise from the irradiated surface of the donor star. This lead to the first estimate of the semi-amplitude of the radial velocity of the secondary, $K_2$, and mass function ($f$($M_1$)=$K_2^3$$P_{\\rm orb}$/4$\\pi$$G$) of Sco\\,X-1. These lines were most prominent in the Bowen region (the region from 4630 to 4660 \\AA), that mainly consists of a blend of N\\,III and C\\,III lines. The N\\,III lines are the result of a UV fluorescence process, while the C\\,III lines are due to photo-ionization and subsequent recombination \\citep{mct75}.  In this paper we review the results of a survey we performed to apply this new technique to other persistent LMXBs or transients during outburst in order to find a donor star signature. ",
        "conclusions": "We have given a short overview of the novel technique of Bowen fluorescence that has proven itself to work for the 9 systems thus far observed (see Table\\,\\ref{limits}), and for most of the systems we have been able to derive the first constraints on their system parameters ever.  There are still a handful of persistent LMXBS left that are bright enough for this technique, and there are plans to also observe these systems. Recently 4U\\,1957$+$11 has been observed with Magellan, and in the near future EXO\\,0748$-$676, Ser\\,X-1 and 4U\\,1556$-$605 will be observed with the VLT. Furthermore, this technique might also be an excellent tool to determine the system parameters of future bright transients.  The next step will be to improve the determined $K_2$ velocities for the most promising candidates by measuring the distribution of the emission across the Roche-lobe using high resolution spectroscopy. This will allow us to more accurately measure the radial velocities of the narrow lines, determine their rotational broadening, and hopefully model the shape of the lines as a function of orbital period in order to get a better constraint on the $K$-correction. Furthermore, we have started using the fact that the Bowen lines are efficiently reprocessed on the donor star to constrain the inclination and the $K$-correction using Echo-tomography with a narrow band filter centered around the Bowen blend (for more information see Mu\\~noz-Darias et~al. in this volume). These next steps will hopefully provide stronger constraints on the masses of LMXBs, and in particular give accurate neutron star masses."
    },
    "0801/0801.1681_arXiv.txt": {
        "abstract": "We report CCD photometry of the dwarf nova BF Ara throughout fifteen consecutive nights in quiescence. Light curve in this interval is dominated by a large amplitude ($\\sim 0.8$ mag) modulation consisting two periods. Higher amplitude signal is characterized by period of 0.082159(4) days, which was increasing at the rate of $\\dot P/P_{\\rm sh} = 3.8(3)\\cdot 10^{-5}$. Weaker and stable signal has period of 0.084176(21) days. Knowing the superhump period of BF Ara determined by Kato et al. (2003) and equal to 0.08797(1) days, the first modulation is interpreted as quiescent negative superhump arising from retrograde precesion of titled accretion disk and the latter one as an orbital period of the binary. The respective period excess and defect are $\\epsilon_+ = 4.51\\% \\pm 0.03\\%$ and $\\epsilon_- = -2.44\\% \\pm 0.02\\%$. Thus BF Ara is yet another in-the-gap nova with mass ratio of around $q\\approx0.21$.  \\noindent {\\bf Key words:} Stars: individual: BF Ara -- binaries: close -- novae, cataclysmic variables ",
        "introduction": "The SU UMa variables are a subclass of dwarf novae characterized by the presence of two types of eruptions - normal outburst having an amplitude of 2-3 mag and repeating typically every few tens of days and superoutburts with amplitude of 3-4 mag, lasting few times longer and occuring every year or so.  These stars are belived to be binary systems consiting of the white dwarf primary and low mass main-sequence secondary filling its Roche lobe and loosing material which forms an accretion disc around the primary. Almost all SU UMa stars are close binaries with orbital period from 1.25 to slightly over 2 hours.  This behaviour in now quite well understood within the frame of the thermal-tidal instability model (see Osaki 1996 for review). Normal outbursts are caused by the thermal instability connected with transition of the material in the disc from neutral to ionized state. Superoutbursts are a result of combined thermal and tidal instability, which, working together, are very effective with sweeping out the matter from the disc.  During the superoutburst, the characteristic tooth-shape light modulations with a period a few percent longer than the orbital period of the binary are observed. They are most probably the result of accretion disc \"precession\" (in fact it is not classical precession but change of the position of the line of apsides) caused by gravitational perturbations from the secondary. These perturbations are most effective when disc particles moving in eccentric orbits enter the 3:1 resonance. Then the superhump period is simply the beat period between orbital and \"precession\" rate periods.  At the beginning of the 1990s, the quite simple class of SU UMas started to be more complicated. The systems showing superhumps were divided into four subgroups:  \\begin{itemize}  \\item WZ Sge stars, characterized by an extremely long quiescent state, going into superoutburst every $\\sim$10 years and showing no or very infrequent ordinary outbursts (Patterson et al. 2002),  \\item ordinary SU UMa stars,  \\item ER UMa stars - systems characterized by an extremely short supercycle (20-60 days), a short interval between normal outbursts (3-4 days) and small amplitude (2--3 mag) of superoutbursts (Kato and Kunjaya 1995, Robertson et al. 1995),  \\item permanent superhumpers - high accretion rate systems being permanently in superoutbursts (Skillman and Patterson 1993).  \\end{itemize} \\smallskip  We now believe that the classes listed above represent increasing mass transfer. Systems with mass transfer rates as low as $10^{15}$ g/s are inactive, quiet and evolved stars of WZ Sge type containing  brown dwarf degenerated secondary. Classical SU UMa stars have mass transfer rates  one magnitude higer and go into the superoutburt every year or so. ER UMa stars are high mass transfer rate systems (few times per $10^{16}$ g/s) showing frequent and relatively long superoutburst lasting for about half of the supercycle. The explanation of the shortest supercycles involves poorly understood mechanisms causing premature quench of the eruption (Osaki 1995). And finally, permanent superhumpers are system with accretion disc which is thermaly stable and tidaly unstable all the time, i.e. being in permanent state of superoutburst. ",
        "conclusions": "Our main conclusions may be summarized as follows:  \\begin{itemize}  \\item Light curve of BF Ara in quiescence is dominated by high amplitude signal with mean period of 0.082159(4) days which during two weeks observing run increased at the rate of $\\dot P/P_{\\rm sh} = 3.8(3)\\cdot 10^{-5}$  \\item We interpret this periodicity as being due to a negative superhump. Such an interpretation is confirmed by the location of BF Ara in the Ritter et al. (2002a) relation,  \\item BF Ara seems to be a twin of V503 Cyg (Harvey et al., 1995) - both stars are very active (supercycles of 80--90 days), both show large amplitude negative superhumps in quiescence characterized by changing value period,  \\item Prewhitening of the original light curve with the main periodicity resulted in the discovery of another modulation with constant period equal to 0.084176(21) days, which is interpreted as the orbital period of the system. This value makes BF Ara another in-the-gap cataclysmic variable,  \\item Knowing the ordinary superhump period measured by Kato et al. (2003) we were able to calculate the period excess as equal to $4.51\\% \\pm 0.03\\%$ which indicates the mass ratio of $q\\approx0.21$.  \\end{itemize}  \\bigskip \\noindent {\\bf Acknowledgments.} ~We acknowledge generous allocation of  the SAAO 1-m telescope time. We would like to thank Prof. J\\'ozef Smak for fruitful discussions. This work was supported for SALT grant number 76/E-60/SPB/MSN/P-03/DWM 35/2005-2007."
    },
    "0801/0801.3198_arXiv.txt": {
        "abstract": "The energy spectra of TeV gamma-rays from blazars, after being corrected for intergalatic absorption in the Extragalactic Background Light (EBL), appear unusually hard, a fact  that poses challenges to the conventional models of particle acceleration in TeV blazars and/or to the EBL models. In this paper we show that the internal absorption of gamma-rays caused by interactions with dense narrow-band radiation fields in the vicinity of compact gamma-ray production regions can lead to the formation of gamma-ray spectra of an almost arbitrary hardness. This allows significant relaxation of the current tight constraints on particle acceleration and radiation models,  although at the expense of enhanced requirements to the available nonthermal energy budget. The latter, however, is not a critical issue, as long as it can be largely compensated by the Doppler boosting, assuming very large ($\\geq 30$) Doppler factors of the relativistically moving gamma-ray production regions. The suggested scenario of formation of hard gamma-ray spectra predicts detectable synchrotron radiation of secondary electron-positron pairs which might require a revision of the  current ``standard paradigm'' of  spectral energy distributions of gamma-ray blazars. If the primary gamma-rays are of hadronic origin related to $pp$ or $p \\gamma$ interactions, the ``internal gamma-ray absorption'' model predicts neutrino fluxes close to the detection threshold of the next  generation high energy neutrino detectors. ",
        "introduction": "The recent reports on detections of very high energy (VHE) gamma-rays from blazars with redshifts $z \\geq 0.1$ (for a review see e.g. Hinton (2007)) initiated renewed debates on the interpretation of TeV gamma-ray spectra of blazars, in particular in the context of the level of the diffuse extragalactic background radiation at optical and infrared wavelengths, often called also as Extragalactic Background Light (EBL). Initially, the tight link  between these two topics - TeV blazars and EBL - became a subject of hot discussions  prompted by multi (up to 20) TeV gamma-rays detected from a nearby BL Lac object,  Mkn~501  \\cite{mkn501_HEGRA},  and by the reports claiming detection of high  fluxes of EBL at far infrared  wavelengths \\cite{Hauser,Schlegel,Lagage,Fink}. However, it was quickly recognised that these two claims hardly could be  compatible within any standard  model of TeV blazars (see, for a review, Aharonian (2001)).  A distinct feature of extragalactic gamma-ray astronomy is that VHE  gamma-rays emitted by distant ($\\geq 100 \\ \\rm Mpc$) objects arrive after significant absorption caused by their interactions with EBL via the process $\\gamma \\gamma \\rightarrow e^+ e^-$ \\cite{Nikishov,Jelley,Gould}. The reconstructed, i.e. the absorption-corrected gamma-ray spectrum from a source at a redshift z, $J_0(E)=J_{\\rm obs}(E) e^{\\tau(E,z)}$ % depends on the flux and energy spectrum of EBL through the optical depth $\\tau(E,z)$. Thus, at energies where $\\tau(E,z) \\geq 1$, the primary gamma-rays suffer strong spectral deformation.  The EBL consists of two emission components produced by stars and partly absorbed/re-emitted by dust throughout the entire history of galaxy evolution. As a result, two distinct bumps are expected in the spectral energy distribution (SED) of EBL at near infrared (NIR) and far infrared (FIR) wavelengths, with a mid-infrared (MIR) ``valley'' between these two bumps (see e.g. Hauser and Dwek (2001)). Generally, for almost all EBL models,  $\\tau(E)$ is a strong function of energy  below 1 TeV and above 10 TeV; between 1 and 10 TeV the energy-dependence of  $\\tau(E)$ is much weaker \\cite{Hamburg01}. Consequently, one  should expect significant distortion of the VHE spectra of blazars at energies below 1 TeV and  above 10 TeV, provided that at these energies $\\tau \\geq 1$. One can re-formulate  this statement in a different way. Namely, for a standard (``decent'') intrinsic gamma-ray spectrum, the observer should detect  very soft (steep) spectra at energies below 1 TeV and above 10 TeV from objects for which $\\tau \\geq 1$  at corresponding energies. This condition is safely satisfied, given the constraints on the minimum  EBL flux imposed by  galaxy counts, for blazars with redshifts $z \\geq 0.15$ like 1ES~1101-232 and for nearby objects with $z \\sim 0.03$ like Mkn~501. Even though the \\textit{detected} gamma-ray spectra from both objects in the corresponding energy intervals are indeed quite steep with a photon index $\\sim 3$ \\cite{mkn501_HEGRA,1101_HESS}, they appear not sufficiently steep to compensate the function $f(E)=e^{\\tau(E)}$, and thus to prevent a robust conclusion that the \\textit{intrinsic} VHE gamma-ray spectra of these blazars are unusually hard.  In the case of Mkn 501, the intrinsic spectrum has a ``non-standard'' shape with a possible pile-up above 10 TeV which has been interpreted as a ``IR background - TeV gamma-ray crisis'' \\cite{IRcrisis} or a need to invoke dramatic assumptions like a violation of the Lorentz invariance (see e.g. Kifune (1999)). However, a more pragmatic view which presently dominates in both  infrared and gamma-ray astronomical communities, treats this ``crisis'' as somewhat exaggerated, especially given the ambiguity of extraction of the  truly diffuse  extragalactic FIR component from the much higher backgrounds of local origin (see e.g. Hauser and Dwek (2001)). Nevertheless, the recently reported low limits on the EBL at mid infrared wavelengths from the Spitzer deep cosmological surveys appeared quite high, for example at 70 $\\mu \\rm m$  the EBL flux should exceed $\\geq 7.1 \\pm 1.0 \\ \\rm nW/m^2 s$ \\cite{Dole}. This implies that the problem is not yet over, and one may still face a challenge with the interpretation of the energy spectra of Mkn 501 and Mkn 421 in the multi-TeV energy domain.   On the other hand, the recent detections of TeV gamma-rays from blazars with redshifts $z \\geq  0.15$  renewed the potential problems and challenges for standard models of TeV blazars. This time the issue has a more solid experimental background, because the gamma-ray spectra corrected for the intergalactic absorption appear very hard (``harder than should be'') even for the minimum possible EBL fluxes at optical and NIR  wavelengths. Namely, the HESS collaboration reported, based on the detection of TeV gamma-rays from the BL Lac object 1ES~1101-232, that any significant deviation from the lower limits of  EBL determined by the integrated light of galaxies resolved by the Hubble telescope \\cite{Madau}, would lead  to very hard intrinsic gamma-ray spectrum with a slope characterized by a photon index $\\Gamma_0 \\leq 1.5$ \\cite{1101_HESS}. The analysis based on a larger sample of TeV blazars leads to the same conclusion \\cite{Mazin}. Recently, the HESS collaboration reported detection of multi-TeV gamma-rays from 1ES 0229+200, a BL Lac object located at a redshift z=0.1396 \\cite{0229_HESS}. It is remarkable that the \\textit{detected} hard gamma-ray spectrum of this source with a photon index  $\\Gamma_{\\rm obs} \\sim 2.5$ extends up to 15 TeV. This, to a certain extent surprising result can be explained by the shape of the energy flux of EBL which between the NIR and MIR  bands is expected to be  proportional to $\\lambda^{-1}$ \\cite{Hamburg01}. Yet, the absolute EBL flux, derived from a rather conservative assumption that  the photon index of the intrinsic spectrum of TeV gamma-rays  does not  exceed $1.5$, appears again close to the EBL lower limit, this time at MIR ($\\approx 2-3 \\ \\rm nW/m^2 sr$ at $10 \\ \\rm \\mu m$), derived from the Spitzer galaxy counts \\cite{Fazio,Dole}. Thus, the gamma-ray observations of 1ES~1101-232 and 1ES 0229+200 can be interpreted  as an argument that the galaxies resolved by the Hubble and Spitzer telescopes provide the bulk of the EBL flux  from optical to mid infrared wavelengths. Given the importance of such a statement, in particular for understanding of contribution of the first stars to the EBL (see e.g. Kashlinsky (2005); Mapelli et al. (2006)), it is essential to explore alternative ways of explanation of  very hard  intrinsic gamma-ray spectra or even sharper spectral features (like pile-ups) in TeV blazars. In this context, recently  some extreme assumptions regarding the distributions of accelerated particles have been proposed. In particular, Katarzynski et al  (2006) argued that a gamma-ray spectrum as hard as $\\Gamma_0 \\sim 0.7$ can be formed in a SSC model assuming a narrow parent electron distribution, e.g. power-law within $E_1$ and $E_2$, with  a low-energy cutoff $E_1$ not much smaller than the high energy cutoff, $E_2$. In similar lines, Stecker et al (2007) argued that electron spectra with power law index $\\leq 1$ can be accommodated within the  models of relativistic shock acceleration . It should be noted, however, that in compact objects relativistic electrons usually suffer very fast synchrotron losses, therefore the assumptions about hard electron \\textit{acceleration } cannot yet guarantee hard gamma-ray spectra. Indeed, the radiatively cooled electron spectrum cannot be harder than $dN/dE \\propto E^{-2}$, independent of the initial (acceleration) spectrum (see e.g. Aharonian 2004). If so, the inverse Compton  scattering  would result in a gamma-ray spectrum steeper than $E^{-1.5}$. In fact, the Klein-Nishina effect makes the spectrum even steeper. In principle, one can avoid the synchrotron cooling of electrons, e.g. in a cold ultrarelativistic wind. However such a hypothesis suggested for Mkn~501\\cite{ColdWind}, in analogy with pulsar winds, needs thorough theoretical studies to clarify  whether such cold ultrarelativistic winds can be formed and survived around supermassive black holes in the cores of AGN.  In this paper we suggest a new scenario which allows formation of  very (in practice, arbitrary) hard gamma-ray spectra in a quite natural way. The model is based on a postulation that gamma-rays before leaving the source suffer significant photon-photon absorption due to interactions with dense radiation fields inside or in the vicinity of compact gamma-ray production region(s). Interestingly,  the presence of high density radiation fields of different origin in the inner parts of blazars generally is treated as a problem for the escape of high energy gamma-radiation from their production region, and, in this regard, the current models of TeV blazars are designed in a way to avoid the internal gamma-ray absorption. Below we show that, in fact,  a moderate internal photon-photon absorption can be a clue to the  very hard intrinsic energy spectra of TeV blazars. ",
        "conclusions": "The energy spectra of VHE  gamma-rays from blazars, after correction for intergalactic absorption,  generally appear very hard, even for the minimum flux level of EBL determined by the integrated light of resolved galaxies at optical (Hubble)  and infrared (Spitzer) wavelengths.  A slight deviation from the  robust lower limits of EBL leads to unusually hard intrinsic gamma-ray spectra which cannot be easily explained within  the standard  particle acceleration and radiation models. In this paper we suggest a scenario which can lead to the formation of instrinsic gamma-ray spectra of arbitrary hardness without introducing modifications in the particle acceleration models. The main idea is that the gamma-rays  before they leave  the source suffer significant internal energy-dependent absorption due to interactions with the ambient low-frequency photons. The existence of dense radiation fields of different origin in blazars (see e.g. Urry and Padovani (1995)) combined with the large photon-photon pair production cross-section, makes this scenario quite natural and effective, in particular in the compact cores of blazars. For the formation of hard VHE gamma-ray spectra, the target  radiation field must have a rather narrow spectral distribution or a sharp low-energy cutoff, with a typical energy of photons of about 1 to 10 eV.  Formally, for very large optical depths, this process can provide an arbitrary hardness of gamma-ray spectra, though at the expense of a significant increase of the required nonthermal energy budget. However, as long as the current blazar models require relativistically-moving gamma-ray production regions with large Doppler factors, $\\delta_j \\geq 30$, and perhaps even more \\cite{PKS_HESS,Fabian}, the available energy budget seems to be not a critical issue.  The unavoidable feature of the proposed model is the radiation of secondary electrons via synchrotron or inverse Compton scattering. If the optical depth inside the gamma-ray production region is small, $\\tau \\ll 1$, e.g. the gamma-ray source  is much smaller than the external source of optical photons, the secondary electrons are produced and radiate mainly outside the gamma-ray production region. Even in the case of heavy absorption of gamma-rays, the secondary radiation of secondary electrons can hardly be detected. Indeed, since the intrinsic gamma-ray luminosity is relatively modest, and the absorbed energy is re-radiated as an isotropic source\\footnote{Unless the electrons are produced in an environment with  a very low magnetic field, and thus are cooled via inverse Compton scattering before any noticeable deflection.}, the lost of the beaming factor dramatically reduces the  signal compared to the primary (Doppler boosted) radiation.  The picure is dramatically changed when the gamma-ray source moves through a very dense photon field, such that the optical depth inside the source becomes larger than 1. In this case the main fraction of the absorbed energy is released in the form of secondary electrons inside the gamma-ray production region, and thus the radiation of the secondary electrons profits, as the primary gamma-radiation does, from the Doppler boosting. The secondary electrons are cooled through synchrotron and/or inverse Compton channels. The latter in fact proceeds via development of pair cascades as long as the typical energies of electrons or gamma-rays and the energy of target photons $\\varepsilon E_{\\rm e, \\gamma} \\gg m^2_e c^4$. The cascade, however,  diminishes the energy-dependent absorption features, thus the  model becomes effective  when the electrons are cooled predominantly via synchrotron radiation, i.e. $B^2/8 \\pi \\geq u_{\\rm r}$. The energy density of the radiation $u_{\\rm r} = \\bar{\\varepsilon} n_{\\rm ph}$ with an average energy of target photons of about $\\bar{\\varepsilon} \\sim 1 \\ \\rm eV$ is estimated from the condition $\\tau_{\\rm max} \\geq 1$, thus for the effective suppression of the cascade  \\begin{equation} B \\geq (40 \\pi \\bar{\\varepsilon}/\\sigma_{\\rm T} l)^{1/2} \\approx 0.5 (l/10^{15} \\rm cm)^{-1/2} \\tau^{1/2}_{\\rm max}\\ \\rm G \\label{B} \\end{equation}  For an optical depth $\\tau_{\\rm max} \\sim 1$ and the size of the gamma-ray source  $l \\sim   10^{16} \\ \\rm cm$, the magnetic field exceeding  0.1 G should be sufficient to prevent the cascade. For smaller gamma-ray production regions, e.g. $l \\sim 10^{14} \\ \\rm cm$, the magnetic field  should be  larger than 1 G. For such magnetic fields the synchrotron radiation of secondary electrons appears in the  optical to hard X-ray energy bands.  Depending on the optical depth, the synchrotron peak can be higher than the gamma-ray peak. Interestingly, unlike the classical ''synchrotron/Inverse-Compton``  models, where the ratio of the synchrotron to IC peak is determined by the ratio $u_{\\rm B}/u_{\\rm r}$, in the ''internal gamma-ray absorption`` scenario the synchrotron peak does not strongly depend on the magnetic field. Whether this scenario can be applied to the the broad-band SEDs of gamma-ray blazars, is an interesting issue which requires special dedicated studies.  Finally, we want to discuss briefly the radiation mechanisms of primary gamma-radiation. Generally, the model  does not give a preference to the leptonic or hadronic origin of radiation, unless the magnetic field exceeds the estimate given by Eq.(\\ref{B}). In this case the synchrotron-to-IC flux ratio produced by directly accelerated electrons would be too high, especially after the internal absorption of gamma-rays, contrary to the detected SEDs of most of the TeV blazars.  Large  magnetic fields in the gamma-ray production region, typically $B \\geq 1 \\ \\rm G$, would favor gamma-ray production by  relativistic protons, with all advantages and disadvantages common for hadronic models. The basic problem of  hadronic models is linked to the low interaction rates which do not allow the most natural explanation of the observed fast gamma-ray variability of blazars in terms of radiative cooling. For example, in the case of interactions of protons with the ambient plasma with number density $n$, the characteristic time of $pp$ interactions with production of $\\pi^0$-mesons is $t_{\\rm pp} \\approx 10^{15} n^{-1} \\ \\rm s$. Thus, in order to explain the variability of gamma-rays as short as several minutes like the TeV flares observed from PKS 21555-301 and Mkn 501, the density of plasma should be as large as $5 \\times 10^{12} \\delta_j^{-1} \\ \\rm cm^{-3}$ which implies a very heavy source  and correspondingly  huge kinetic energy $E_{\\rm kin}= (4/3) \\pi l^3 n m_pc^2 \\gamma_j  \\approx 10^{55} \\ \\rm erg$ (here we assume that $\\delta_ \\approx \\gamma_j$). One may invoke alternative explanations of the variability of blazars, e.g. due to the adiabatic losses or escape of particles from the source, but this assumption leads to dramatic reduction of radiation efficiency, and to an increase the energy requirements to the  accelerated protons.  A similar problem face the photomeson processes at interactions of protons with the ambient radiation fields. Actually in the ''internal gamma-ray absorption`` scenario this mechanisms seems a quite natural choice because  the same  background photons which absorb gamma-rays can play a role of the target for photomeson interactions. However, because of the small cross-section, the efficiency of this process again appears  quite low. The interaction time of protons with energy, $E \\geq 200 \\ \\rm MeV/(\\bar{\\varepsilon} \\gamma_j) \\simeq 2 \\times 10^{16} (\\bar{\\varepsilon}/1 \\ \\rm eV)^{-1} (\\gamma_j/10)^{-1} \\ \\rm eV$ (in the frame of the moving source  with a Lorentz factor $\\gamma_j$) is estimated $t_{\\rm p \\gamma} \\approx 1/(f \\sigma_{\\rm p \\gamma} n_{\\rm ph} c) \\sim (\\sigma_{\\rm \\gamma \\gamma}/\\sigma_{\\rm p \\gamma})f^{-1}  R/c \\tau_{\\rm max}^{-1}$ ($<\\sigma_{\\rm p \\gamma}> \\approx  10^{-28} \\ \\rm cm^2$ is the average cross-section and $f \\sim 0.2$  is the multiplicity of the process). Thus, we can see that during the passage of the source of optical photons of size $R$, the protons transfer only $\\sigma_{\\rm p \\gamma}/\\sigma_{\\gamma \\gamma} \\sim 10^{-3}$ fraction of their energy to gamma-rays. If such a low efficiency can be compensated by very large Doppler boosting (e.g. assuming $\\delta_j \\sim 100$), this channel can provide very large fluxes of neutrinos, which unlike gamma-rays do not suffer internal and extragalactic absorption. In the case of attenuation of VHE gamma-ray fluxes by 2 to 3 orders of magnitude, the expected fluxes of neutrinos from TeV blazars  can be as large as the detection threshold of the km$^3$ volume high energy neutrino telescopes, $F_{\\nu_\\mu} (\\geq 1 \\ \\rm TeV) \\approx 10^{-11} \\ \\rm neutrinos/cm^2 s$.  It is interesting to note that, because of the threshold of photomeson production, the interactions of protons of arbitrary distribution with a narrow band radiation with a characteristic energy $\\bar{\\varepsilon}$, result in a differential gamma-ray spectrum which below the energy $\\approx 10^{16} (\\bar{\\varepsilon}/1 \\ \\rm TeV)^{-1}$ eV is extremely hard, ${\\rm d}N/{\\rm d}E=\\rm const$, thus this process itself can provide very hard gamma-ray spectra independent of the spectrum of parent protons.  Despite certain attractive features, this mechanism faces the same problem  as $pp$ interactions - a low radiation efficiency. Therefore it can work only under conditions of extremely large Doppler boosting of radiation. The efficiency of VHE gamma-ray production can be much higher in the case of synchrotron radiation of protons, provided that the acceleration of protons proceeds at a rate close to the fundamental limit, and  the magnetic field in the proton accelerator well exceeds 10 G. In particular, in the magnetic field of order $100$ G, protons can be accelerated to energies $10^{20} \\ \\rm TeV$ and thus can produce VHE synchrotron gamma-rays on timescales of $10^{4}$ s. Although due to the  self-regulated synchrotron cutoff \\cite{psynch} the spectrum of gamma-rays is limited by sub-TeV energies, an  observer detects  Doppler boosted gamma-radiation extending to multi-TeV energies. The characteristic feature of this mechanism is the very large electromagnetic energy contained in the blob, $1/6 l^3 B^2 \\approx  2 \\times 10^{48} \\ \\rm erg$, hard X-ray emission of the secondary (pair-produced) electrons, and negligible fluxes of  neutrinos.     \\hyphenation{Post-Script Sprin-ger}"
    },
    "0801/0801.4895_arXiv.txt": {
        "abstract": "We present the results of relic density calculations for cold dark matter candidates coming from a model of dark energy and dark matter, which is described by an asymptotically free gauge group $\\SU(2)_Z$ (QZD) with a coupling constant $\\alpha_Z \\sim$ 1 at very low scale of $\\Lambda_Z \\sim 10^{-3}$ eV while $\\alpha_Z \\sim$ weak coupling at high energies. The dark matter candidates of QZD are two fermions in the form of weakly interacting massive particles. Our results show that for masses between 50 and 285 GeV, they can account for either a considerable fraction or the entire dark matter of the Universe. ",
        "introduction": "\\label{sec:intro} It is almost universally accepted that the picture of the Universe made up of approximately $4\\%$ baryonic matter, $23\\%$ dark matter and $73\\%$ dark energy represents a realistic cosmological model. However, it is astounding that almost $96\\%$ of the energy density of the Universe resides in some as-yet-unknown form. What is ``dark matter''? What is ``dark energy''?  In Refs.~\\cite{Hung2005,Hung2006a}, a model of dark energy and dark matter was proposed in which a new unbroken gauge group $\\SU(2)_Z$ -- the shadow sector -- grows strong at a scale $\\sim 10^{-3}\\,$eV. The gauge group $\\SU(2)_Z$ was nicknamed Quantum Zophodynamics, or QZD, in  Refs.~\\cite{Hung2005,Hung2006a}, where the subscript ``\\textit{Z}\" stands for the Greek word \\textit{Zophos}, meaning darkness. The model is described by an $\\SU(2)_Z$ instanton-induced potential of an axion-like particle, $a_Z$, which possesses two degenerate minima. The degeneracy is lifted by a mechanism described in Refs.~\\cite{Hung2006a,Hung2007a}, yielding a false vacuum with energy density $\\sim (10^{-3}\\,\\text{eV})^4$ and a true vacuum with vanishing energy density. The present Universe is assumed to be trapped in the false vacuum~\\cite{Hung2007b}, whose energy density mimics the cosmological constant. This is, in a nutshell, the dark energy model proposed in Ref.~\\cite{Hung2006a}, which also computed various quantities of interest such as the tunneling rate to the true vacuum, etc. A Grand Unified Theory (GUT) involving the SM and $\\SU(2)_Z$ was considered by Ref.~\\cite{Hung2006d} (The models presented in Refs.~\\cite{Hung2006a,Hung2006d} were later revisited by Refs.~\\cite{Das}.).  The particle content of the model includes two shadow fermions, $\\psi _{\\left( {L,R} \\right),i}^{\\left( Z \\right)} $ with $i=1,2$, which transform as $\\left( {1,1,0,3} \\right)$ under $\\SU(3)_c \\otimes \\SU(2)_L \\otimes \\mathrm{U}(1)_Y \\otimes \\SU(2)_Z$, two messenger scalar fields (mediating between the QZD and SM matters; one of which is much heavier than the other \\cite{Hung2006a})  $\\tilde{\\bm{\\varphi}} _i^{\\left( Z \\right)} $ with $i=1,2$ transforming as $\\left( {1,2,Y_{\\tilde \\varphi}   =  - 1,3} \\right)$, and one singlet complex scalar field $\\phi _Z  = \\left( {1,1,0,1} \\right)$ whose imaginary part plays the role of the axion-like particle mentioned above.  As discussed in Ref.~\\cite{Hung2006a}, the masses of the $\\SU(2)_Z$ triplet shadow fermions are found to be of the order of 100 - 200 GeV for the $\\SU(2)_Z$ gauge coupling to grow strong at a scale $\\sim 10^{-3}$ eV, needed for the dark energy scenario. This coupling constant starts out at GUT-scale energy with a value comparable to that of the electroweak couplings, remains relatively flat until an energy comparable to the shadow fermion masses is reached, and then  starts to grow after the shadow fermions drop out of the Renormalization Group (RG) equations. At that dropout point, the $\\SU(2)_Z$ gauge coupling becomes comparable to the weak $\\SU(2)_L$ coupling at the electroweak scale energy. These features have interesting consequences concerning the possibility of the shadow fermions being candidates for cold dark matter (CDM) in the form of weakly interacting massive particles (WIMP's) \\footnote{For a review on various features of CDM and WIMP, see, e.g., Refs. \\cite{CDMRevs}.}. The main reason is the fact that the annihilation cross sections for two shadow fermions into two $\\SU(2)_Z$ ``shadow gluons'' are of the order of the weak cross sections, a typical requirement for WIMP's. The estimates that were made in Ref.~\\cite{Hung2006a} showed that it was possible for shadow fermions to be candidates for CDM with the \\textit{right} relic density.  In this work, we would like to investigate this scenario in more details and by solving shadow fermions' evolution equations to determine the conditions under which they can be considered to be WIMP cold dark matter candidates. It will be seen that the mass range for the shadow fermions obtained by the requirement of having the \\textit{right} density fits in snugly with that used in the RG equations (i.e., the $\\SU(2)_Z$ gauge coupling grows strong at a scale $\\sim 10^{-3}$ eV).  The outline of the paper is as follows. First, we go over the QZD model as far as the issue of dark matter is concerned. Then, we derive the evolution equations for shadow fermions and consequently solve them numerically, to obtain their relic density. Finally, the results of our relic density calculations will be presented and discussed, in comparison with the observational values. The shadow fermions relic density, when computed, would only depend on their masses. Therefore, the parameter space is simply two dimensional. ",
        "conclusions": "We solved evolution equations for number densities of shadow fermions and obtained their total present-day density. The heavier shadow fermion turned out to be long lived if its mass differs from that of the messenger field. In that case, our results revealed an upper bound on the mass of the heavier shadow fermion, i.e., $m_2 \\approx 245$ GeV, above which its \\textit{late} decay can potentially disturb the CMB density of the Universe beyond the measured fluctuation level of $10^{-5}$.  For lighter shadow fermions, the total relic density can account for the entire dark matter of the Universe depending on the mass combination of shadow fermions. When the total density falls short of the observationally suggested density, it still, for most of masses, provides significant fraction of the dark matter of the Universe.  Our results showed that if the heavier shadow fermion's mass is large, considerable mass differences would be needed to comply with experimental bounds. On the other hand, if the heavier shadow fermion's mass is small, little or even no mass differences suffice to give the right relic density. In that sense, degenerate and near-degenerate mass cases become relevant at low mass scales, but not for less than 50 GeV.  A very short lifetime is expected for the heavier shadow fermion if its mass is the same as that of the messenger field. In that case, the calculations reduce to a one-species case. Our results suggest that a sole shadow fermion must have a mass of about 190~--~210 GeV to account for the whole dark matter of the Universe.  Last but not least, possible detections of the shadow fermion CDM candidates are briefly discussed in Ref.~\\cite{Hung2006a}. Needless to say, more work along this line is warranted for this model."
    },
    "0801/0801.0116.txt": {
        "abstract": "Primordial black holes (PBHs) are a profound signature of primordial cosmological structures and provide a theoretical tool to study nontrivial physics of the early Universe. The mechanisms of PBH formation are discussed and observational constraints on the PBH spectrum, or effects of PBH evaporation, are shown to restrict a wide range of particle physics models, predicting an enhancement of the ultraviolet part of the spectrum of density perturbations, early dust-like stages, first order phase transitions and stages of superheavy metastable particle dominance in the early Universe. The mechanism of closed wall contraction can lead, in the inflationary Universe, to a new approach to galaxy formation, involving primordial clouds of massive BHs created around the intermediate mass or supermassive BH and playing the role of galactic seeds.   Primordial black holes (PBHs) are a profound signature of primordial cosmological structures and provide a theoretical tool to study nontrivial physics of the early Universe. The mechanisms of PBH formation are discussed and observational constraints on the PBH spectrum, or effects of PBH evaporation, are shown to restrict a wide range of particle physics models, predicting an enhancement of the ultraviolet part of the spectrum of density perturbations, early dust-like stages, first order phase transitions and stages of superheavy metastable particle dominance in the early Universe. The mechanism of closed wall contraction can lead, in the inflationary Universe, to a new approach to galaxy formation, involving primordial clouds of massive BHs created around the intermediate mass or supermassive BH and playing the role of galactic seeds. ",
        "introduction": "The convergence of the frontiers of our knowledge in micro- and macro- worlds leads to the wrong circle of problems, illustrated by the mystical Uhroboros (self-eating-snake). The Uhroboros puzzle may be formulated as follows: {\\it The theory of the Universe is based on the predictions of particle theory, that need cosmology for their test}. Cosmoparticle physics \\cite{ADS,MKH,book,book3,Khlopov:2004jb,bled} offers the way out of this wrong circle. It studies the fundamental basis and mutual relationship between micro-and macro-worlds in the proper combination of physical, astrophysical and cosmological signatures. Some aspects of this relationship, which arise in the astrophysical problem of Primordial Black Holes (PBH) is the subject of this review.  In particle theory Noether's theorem relates the exact symmetry to conservation of respective charge. Extensions of the standard model imply new symmetries and new particle states. The respective symmetry breaking induces new fundamental physical scales in particle theory. If the symmetry is strict, its existence implies new conserved charge. The lightest particle, bearing this charge, is stable. It gives rise to the deep relationship between dark matter candidates and particle symmetry beyond the Standard model.  The mechanism of spontaneous breaking of particle symmetry also has cosmological impact. Heating of condensed matter leads to restoration of its symmetry. When the heated matter cools down, phase transition to the phase of broken symmetry takes place. In the course of the phase transitions, corresponding to given type of symmetry breaking, topological defects can form. One can directly observe formation of such defects in liquid crystals or in superfluid He. In the same manner the mechanism of spontaneous breaking of particle symmetry implies restoration of the underlying symmetry. When temperature decreases in the course of cosmological expansion, transitions to the phase of broken symmetry  can lead, depending on the symmetry breaking pattern, to formation of topological defects in very early Universe. Defects can represent new forms of stable particles (as it is in the case of magnetic monopoles \\cite{t'Hooft,polyakov,kz,Priroda,preskill,SAOmonop}), or extended structures, such as cosmic strings \\cite{zv1,zv2} or cosmic walls \\cite{okun}.  In the old Big bang scenario  cosmological expansion and its initial conditions were given {\\it a priori} \\cite{Weinberg,ZNSEU}. In the modern cosmology expansion of Universe and its initial conditions are related to inflation \\cite{Star80,Guthinfl,Linde:1981mu,Albrecht,Linde:1983gd}, baryosynthesis and nonbaryonic dark matter (see review in \\cite{Lindebook,Kolbbook}). Physics, underlying inflation, baryosynthesis and dark matter, is referred to extensions of the standard model, and variety of such extensions makes the whole picture in general ambiguous. However, in a framework of each particular physical realization of inflationary model with baryosynthesis and dark matter the corresponding model dependent cosmological scenario can be specified in all details. In such scenario main stages of cosmological evolution, structure and physical content of the Universe reflect structure of the underlying physical model. The latter should include with necessity the standard model, describing properties of baryonic matter, and its extensions, responsible for inflation, baryosynthesis and dark matter. In no case cosmological impact of such extensions is reduced to reproduction of these three phenomena only. A nontrivial path of cosmological evolution, specific for each particular realization of inflational model with baryosynthesis and nonbaryonic dark matter, always contains some additional model dependent cosmologically viable predictions, which can be confronted with astrophysical data. Here we concentrate on Primordial Black Holes as profound signature of such phenomena.  It was probably Pierre-Simon Laplace \\cite{Laplace} in the beginning of XIX century, who noted first that in very massive stars escape velocity can exceed the speed of light and light can not come from such stars. This conclusion made in the framework of Newton mechanics and Newton corpuscular theory of light has further transformed into the notion of \"black hole\" in the framework of general relativity and electromagnetic theory. Any object of mass $M$ can become a black hole, being put within its gravitational radius $r_g=2 G M/c^2.$ At present time black holes (BH) can be created only by a gravitational collapse of compact objects with mass more than about three Solar mass \\cite{1,ZNRA}. It can be a natural end of massive stars or can result from evolution of dense stellar clusters. However in the early Universe there were no limits on the mass of BH. Ya.B. Zeldovich and I.D. Novikov \\cite{ZN} noticed that if cosmological expansion stops in some region, black hole can be formed in this region within the cosmological horizon. It corresponds to strong deviation from general expansion and reflects strong inhomogeneity in the early Universe.  There are several mechanisms for such strong inhomogeneity and we'll trace their links to cosmological consequences of particle theory.  %The simplest one is a collapse of strongly inhomogeneous regions %just after the end of inflation (see e.g. \\cite{2}). Another %possible source of BH could be a collapse of cosmic strings \\cite{3} %that are produced in early phase transitions with symmetry breaking. %The collisions of the bubble walls \\cite{4,5} created at phase %transitions of the first order can lead to a primordial black hole %(PBH) formation.   Primordial Black Holes (PBHs) are a very sensitive cosmological probe for physics phenomena occurring in the early Universe. They could be formed by many different mechanisms, {\\it e.g.}, initial density inhomogeneities \\cite{hawking1,hawkingCarr} and non-linear metric perturbations \\cite{Bullock:1996at,Ivanov:1997ia,Bullock:1998mi}, blue spectra of density fluctuations \\cite{Khlopov:1984wc,polnarev,Lidsey:1995ir,Kotok:1998rp, Dubrovich02,Sendouda:2006nu}, a softening of the equation of state \\cite{canuto,Khlopov:1984wc,polnarev}, development of gravitational instability on early dust-like stages of dominance of supermassive particles and scalar fields \\cite{khlopov0,polnarev0,polnarev1,khlopov1} and evolution of gravitationally bound objects formed at these stages \\cite{Kalashnikov,Kadnikov}, collapse of cosmic strings \\cite{hawking2,Polnarev:1988dh,Hansen:2000jv,Cheng:1996du,Nagasawa:2005hv} and necklaces \\cite{Matsuda:2005ez}, a double inflation scenario \\cite{nas,Kim:1999xg,Yamaguchi:2001zh,Yamaguchi:2002sp}, first order phase transitions \\cite{hawking3,pt1Jedamzik,kkrs,kkrs1,kkrs2}, a step in the power spectrum \\cite{Sakharov0,polarski1}, etc. (see \\cite{polnarev,book,book3,Carr:2003bj,book2} for a review).  Being formed, PBHs should retain in the Universe and, if survive to the present time, represent a specific form of dark matter \\cite{khlopov7,Ivanov:1994pa,book,book3,Blais:2002nd,Chavda:2002cj, Afshordi:2003zb,book2,Chen:2004ft}. Effect of PBH evaporation by S.W.Hawking \\cite{hawking4} makes evaporating PBHs a source of fluxes of products of evaporation, particularly of $\\gamma$ radiation \\cite{Page:1976wx}. MiniPBHs with mass below $10^{14}$~g evaporate completely and do not survive to the present time. However, effect of their evaporation should cause influence on physical processes in the early Universe, thus providing a test for their existence by methods of cosmoarcheology \\cite{Cosmoarcheology}, studying cosmological imprints of new physics in astrophysical data. In a wide range of parameters the predicted effect of PBHs contradicts the data and it puts restrictions on mechanism of PBH formation and the underlying physics of very early Universe. On the other hand, at some fixed values of parameters, PBHs or effects of their evaporation can provide a nontrivial solution for astrophysical problems.  Various aspects of PBH physics, mechanisms of their formation, evolution and effects are discussed in \\cite{carr1,carrMG,LGreen,khlopov6,polnarev,Grillo:1980uj, Chapline:1975tn,Hayward:1989jq,Yokoyama:1995ex,Kim:1996hr, Heckler:1997jv,MacGibbon:2007yq,Page:2007yr,green,Niemeyer:1997mt,Kribs:1999bs, Green:2004wb,Yokoyama:1998pt,Yokoyama:1998xd,Yokoyama:1999xi, Bringmann:2001yp,Dimopoulos:2003ce,Nozari:2007kv,LythMalik,Zaballa:2006kh, Harada:2004pf, Custodio:2005en,Bousso:1995cc,Bousso:1996wy,Elizalde:1999dw, Nojiri:1999vv,Bousso:1999iq,Silk:2000em,Polarski:2001jk, Barrow:1996jk,Paul:2000jb,Paul:2001yt,Paul:2005bk,Polarski:2001yn, Carr:1993aq,Yokoyama:1998qw,Kaloper:2004yj,Pelliccia:2007fh, Stojkovic:2005zh,Soda,Ahn:2006uc,babichev4,babichev5,babichev6,Guariento:2007bs} particularly specifying PBH formation and effects in braneworld cosmology \\cite{Guedens:2002km,Guedens:2002sd,Clancy:2003zd, Tikhomirov:2005bt}, on inflationary preheating \\cite{Bassett:2000ha}, formation of PBHs in QCD phase transition \\cite{Jedamzik:1998hc, Widerin:1998my}, properties of superhorizon BHs \\cite{Harada:2005sc,Harada:2006gn}, role of PBHs in baryosynthesis \\cite{Grillo:1980rt,Barrow:1990he,Turner:1979bt,Upadhyay:1999vk, Bugaev:2001xr}, effects of PBH evaporation in the early Universe and in modern cosmic ray, neutrino and gamma fluxes \\cite{mujana,Fegan:1978zn,Green:2001kw,Frampton:2005fk, MacGibbon:1990zk,MacGibbon:1991tj,Halzen:1991uw,Halzen:1995hu, Bugaev:2000bz,Bugaev:2002yt,Volkova:1994fb,Gibilisco:1996ft, Golubkov:2000qy,He:2002vz,Gibilisco:1996dk,Custodio:2002jv, Sendouda:2003dc,Maki:1995pa,barraupbar,Barrau:2002mc, Wells:1998jv,Cline:1996uk,Xu:1998hn,Cline:1998fx,Sendouda:2006yc, Barrau:1999sk,Derishev:1999xn,Tikhomirov:2004rs,Seto,barrau, barraugamma,barrauprd}, in creation of hypothetical particles \\cite{Bell:1998jk,lemoine,green1,barrau2}, PBH clustering and creation of supermassive BHs \\cite{Bean:2002kx,Duechting:2004dk,Chisholm:2005vm, Dokuchaev:2004kr,Mack:2006gz,Rubin:2005pq}, effects in cosmic rays and colliders from PBHs in low scale gravity models \\cite{barrauADDBH,barrauADDac}. Here we outline the role of PBHs as a link in cosmoarcheoLOGICAL chain, connecting cosmological predictions of particle theory with observational data. We discuss the way, in which spectrum of PBHs reflects properties of superheavy metastable particles and of phase transitions on inflationary and post-inflationary stages. We briefly review possible cosmological reflections of particle physics (section \\ref{Cosmophenomenology}), illustrate in section \\ref{dust} some mechanisms of PBH formation on stage of dominance of superheavy particles and fields (subsection \\ref{particles}) and from second order phase transition on inflationary stage. Effective mechanism of BH formation during bubble nucleation provides a sensitive tool to probe existence of cosmological first order phase transitions by PBHs (section \\ref{phasetransitions}). Existence of stable remnants of PBH evaporation can strongly increase the sensitivity of such probe and we demonstrate this possibility in section \\ref{gravitino} on an example of gravitino production in PBH evaporation. Being formed within cosmological horizon, PBHs seem to have masses much less than the mass of stars, constrained by small size of horizon in very early Universe. However, if phase transition takes place on inflationary stage, closed walls of practically any size can be formed (subsection \\ref{walls}) and their successive collapse can give rise to clouds of massive black holes, which can play the role of seeds for galaxies (section \\ref{MBHwalls}). The impact of constraints and cosmological scenarios, involving primordial black holes, is briefly discussed in section \\ref{Discussion}. ",
        "conclusions": "\\label{Discussion} For long time scenarios with Primordial Black holes belonged dominantly to cosmological {\\it anti-Utopias}, to \"fantasies\", which provided restrictions on physics of very early Universe from contradiction of their predictions with observational data. Even this \"negative\" type of information makes PBHs an important theoretical tool. Being formed in the very early Universe as initially nonrelativistic form of matter, PBHs should have increased their contribution to the total density during RD stage of expansion, while effect of PBH evaporation should have strongly increased the sensitivity of astrophysical data to their presence. It links astrophysical constraints on hypothetical sources of cosmic rays or gamma background, on hypothetical factors, causing influence on light element abundance and spectrum of CMB, to restrictions on superheavy particles in early Universe and on first and second order phase transitions, thus making a sensitive astrophysical probe to particle symmetry structure and pattern of its breaking at superhigh energy scales.  Gravitational mechanism of particle creation in PBH evaporation makes evaporating PBH an unique source of any species of particles, which can exist in our space-time. At least theoretically, PBHs can be treated as source of such particles, which are strongly suppressed in any other astrophysical mechanism of particle production, either due to a very large mass of these species, or owing to their superweak interaction with ordinary matter.  By construction astrophysical constraint excludes effect, predicted to be larger, than observed. At the edge such constraint converts into an alternative mechanism for the observed phenomenon. At some fixed values of parameters, PBH spectrum can play a positive role and shed new light on the old astrophysical problems.  The common sense is to think that PBHs should have small sub-stellar mass. Formation of PBHs within cosmological horizon, which was very small in very early Universe, seem to argue for this viewpoint. However, phase transitions on inflationary stage can provide spikes in spectrum of fluctuations at any scale, or provide formation of closed massive domain walls of any size.  In the latter case primordial clouds of massive black holes around intermediate mass or supermassive black hole is possible. Such clouds have a fractal spatial distribution. A development of this approach gives ground for a principally new scenario of the galaxy formation in the model of the Big Bang Universe. Traditionally, Big Bang model assumes a homogeneous distribution of matter on all scales, whereas the appearance of observed inhomogeneities is related to the growth of small initial density perturbations. However, the analysis of the cosmological consequences of the particle theory indicates the possible existence of strongly inhomogeneous primordial structures in the distribution of both the dark matter and baryons. These primordial structures represent a new factor in galaxy formation theory. Topological defects such as the cosmological walls and filaments, primordial black holes, archioles in the models of axionic CDM, and essentially inhomogeneous baryosynthesis (leading to the formation of antimatter domains in the baryon-asymmetric Universe \\cite{exl1,exl2,crg,kolb,we,khl,CSKZ,zil,sb,dolgmain, khlopgolubkov,bgk,FarKhl,book,book2,book3}) offer by no means a complete list of possible primary inhomogeneities inferred from the existing elementary particle models.  Observational cosmology offers strong evidences favoring the existence of processes, determined by new physics, and the experimental physics approaches to their investigation. Cosmoparticle physics \\cite{ADS,MKH,book,book3}, studying the physical, astrophysical and cosmological impact of new laws of Nature, explores the new forms of matter and their physical properties. Its development offers the great challenge for theoretical and experimental research. Physics of Primordial Black holes can play important role in this process."
    },
    "0801/0801.2157_arXiv.txt": {
        "abstract": "{ Inflation may occur while rolling into the metastable supersymmetry-breaking vacuum of massive supersymmetric QCD. We explore the range of parameters in which slow-roll inflation and long-lived metastable supersymmetry breaking may be simultaneously realized. The end of slow-roll inflation in this context coincides with the spontaneous breaking of a global symmetry, which may give rise to significant curvature perturbations via inhomogenous preheating. Such spontaneous symmetry breaking at the end of inflation may give rise to observable non-gaussianities, distinguishing this scenario from more conventional models of supersymmetric hybrid inflation.}  \\begin{document} ",
        "introduction": "Spontaneously broken supersymmetry (SUSY) has long been one of the most attractive possibilities for physics beyond the Standard Model \\cite{Dimopoulos:1981zb}. Considerable effort has been devoted to elucidating dynamical mechanisms of SUSY-breaking, in which strong-coupling effects give rise to a naturally small SUSY-breaking scale \\cite{Witten:1981nf}. The SUSY-breaking vacua of such theories need not be global minima of the potential, and metastability of phenomenologically-viable vacua appears to be generic when embedding the MSSM into a larger setting. Indeed, the study of metastable dynamical SUSY-breaking has recently undergone something of a renaissance, catalyzed by the observation that massive supersymmetric QCD (SQCD) possesses SUSY-breaking local minima whose lifetimes can be longer than the present age of the universe \\cite{Intriligator:2006dd} (for an excellent review, see \\cite{Intriligator:2007cp}). The simplicity and genericness of such theories makes them particularly well-suited to supersymmetric model-building.  Far from serving only as a possible resolution of the hierarchy problem, spontaneously-broken supersymmetry emerges frequently in inflationary settings as well \\cite{Guth:1980zm, Linde:1981mu, Albrecht:1982wi, Linde:1983gd}. There exist a plethora of possible models aimed at realizing inflation in a natural context \\cite{Lyth:1998xn}, many of which exploit  supersymmetry and supersymmetry breaking \\cite{Kinney:1997hm, Kinney:1998dv, Dvali:1994ms, Dimopoulos:1997fv, Copeland:1994vg}. Based on the moderate success of such theories, it is tempting to suppose that inflation may be built into a supersymmetry-breaking sector. Indeed, it has already been demonstrated that inflation may arise in simple strongly-coupled gauge theories \\cite{Dimopoulos:1997fv}. This suggests that it may be possible to inflate while rolling into the metastable vacuum of massive SQCD from sub-Planckian initial field values. The appeal of successful inflation in this scenario is twofold: first, it provides a generic and technically natural means of simultaneously realizing inflation and supersymmetry breaking; second, it furnishes a robust UV completion for the inflationary sector. Moreover, the spontaneous global symmetry-breaking that accompanies the end of inflation in these models may give rise to additional curvature perturbations exhibiting observable non-gaussianities. In this note we address the prospects for, and signatures of, slow-roll inflation in models with metastable dynamical supersymmetry breaking.  The organization of the paper is as follows.  In Sec.~\\ref{sec:hybrid} we review the traditional scenario of supersymmetric hybrid inflation proposed in \\cite{Dvali:1994ms}, wherein one-loop corrections drive inflation along a flat direction of a theory with a relatively simple superpotential. In Sec.~\\ref{sec:hybridSQCD} we review the means by which supersymmetric hybrid inflation may be realized in strongly-coupled gauge theories \\cite{Dimopoulos:1997fv}. Subsequently, in Sec.~\\ref{sec:ISS} we turn to the Intriligator, Seiberg, and Shih (ISS) model \\cite{Intriligator:2006dd}  of metastable supersymmetry breaking in supersymmetric QCD and explore its implications. In the vicinity of the supersymmetry breaking metastable vacuum, features of the ISS model are reminiscent of supersymmetric hybrid inflation. With this motivation, in Sec.~\\ref{sec:metastable} we explore the prospects for, and constraints on, inflation while rolling into the metastable supersymmetry breaking vacuum of the ISS model; in Sec.~\\ref{sec:predictions} we discuss the subsequent inflationary predictions assuming the inflaton is primarily responsible for the primordial curvature perturbation. However, the spontaneous symmetry breaking that accompanies the end of slow-roll inflation may give rise to additional curvature perturbations through inhomogenous preheating \\cite{Kolb:2004jm}. In Sec.~\\ref{sec:broken} we explore this possibility and its consequences, including the prospects for observable non-gaussianities. In Sec.~\\ref{sec:conclusions} we conclude the analysis and consider directions for further work. ",
        "conclusions": "\\label{sec:conclusions}  We have seen that inflation may naturally occur while rolling into the supersymmetry-breaking metastable vacuum of massive supersymmetric QCD. Although the combination of supersymmetry-breaking and inflation is more a matter of novelty than profundity, this scenario is particularly attractive in that it  contains a concrete UV completion of the inflationary sector and the potential for distinctive observational signatures. Successful slow-roll inflation in the ISS model requires a large number of flavors and relatively small magnetic gauge symmetry, as well as a natural hierarchy $m \\ll \\mu \\ll \\Lambda \\lsim M_P.$ Although quadratic corrections to the inflationary potential arise in the presence of a non-canonical K\\\"{a}hler potential, the ensuing contribution to slow-roll parameters may be small enough to forestall the conventional SUSY $\\eta$ problem. Moreover, the spontaneous global symmetry breaking that accompanies the end of slow-roll inflation in the ISS model may give rise to dominant curvature perturbations through inhomogenous preheating. Such perturbations may possess observable non-gaussianities, further distinguishing ISS-flation from its more conventional cousins.  It is difficult to simultaneously obtain weak-scale SUSY breaking and the observed inflationary spectrum strictly from the dynamics of the ISS model; standalone ISS-flation would seem to favor split SUSY or other high-scale mediation. However, weak-scale SUSY-breaking using conventional gauge- or gravity-mediation is feasible if the primary contribution to primordial curvature perturbations arises from coupling to a separate preheating sector.  It would be interesting to consider concrete realizations of split supersymmetry using ISS SUSY-breaking dynamics, given the high scale of SUSY-breaking required to match inflationary observables in the absence of inhomogenous preheating. Likewise, a more careful analysis of non-gaussianities arising from inhomogenous preheating may be useful in light of recent evidence for significant non-gaussianities in the WMAP 3-year data \\cite{Yadav:2007yy}."
    },
    "0801/0801.2966_arXiv.txt": {
        "abstract": "{ Known classes of radio wavelength transients range from the nearby---stellar flares and radio pulsars---to the distant Universe---$\\gamma$-ray burst afterglows.  Hypothesized classes of radio transients include analogs of known objects, e.g., extrasolar planets emitting Jovian-like radio bursts and giant-pulse emitting pulsars in other galaxies, to the exotic, prompt emission from $\\gamma$-ray bursts, evaporating black holes, and transmitters from other civilizations.  A number of instruments and facilities are either under construction or in early observational stages and are slated to become available in the next few years.  With a combination of wide fields of view and wavelength agility, the detection and study of radio transients will improve immensely.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.2820_arXiv.txt": {
        "abstract": "High spatial resolution spectro-polarimetric observation of a decaying spot was observed with the Diffraction Limited Spectro-Polarimeter. The spatial resolution achieved was close to the diffraction limit (0.\"18) of the Dunn Solar Telescope. The fine scales present inside the decaying active region as well as surrounding areas were studied. Two interesting phenomenon observed are: (i) Canopy like structures are likely to be present in the umbral dots as well as in the light bridges providing evidence for field-free intrusion, (ii) There are opposite polarity loops present outside of the spot and some of them connects to the main spot and the surrounding magnetic features. ",
        "introduction": "\\label{sec:intro} Over the past few years, high spatial resolution solar observations became feasible with the development of versatile Adaptive Optics (AO; Rimmele, 2004) systems. The success in obtaining consistent high spatial resolution images, from the ground, revived the development of new instrumentation for observations close to the telescopes' diffraction limit. The versatility of the solar instrumentation made it possible to simultaneously observe a field-of-view (FOV) of interest at different wavelengths. Simultaneous imaging and spectroscopic observations are feasible for quantitative study of physical parameters.  The success of the AO system is well appreciated during spectroscopic and spectro-polarimetric observations. The long integration time and the time required to scan the FOV of interest imposes the consistent image quality requirement from the AO for longer durations (as long as an hour). To supplement the spectroscopic observations, a set of imageries are also essential in order to track the features of interest. One such combination was developed for the Dunn Solar Telescope (DST) of the National Solar Observatory (NSO) at Sacramento Peak, Sunspot situated in New Mexico, USA. An exit port was assigned as a dedicated port for fixed and well calibrated instrumentation where several set of instruments were deployed and facilitate observations with very minimal setup time. This setup also facilitates standardised data reduction procedures. A set of instrumentation, G-band \\& Ca-K imagery, Diffraction Limited Spectro-polarimeter (DLSP), and a light feed for either a Universal Birefringent Filter (UBF) setup or a dual Fabry-Perot setup, were deployed in this port. More details on this instrumentation can be found in Sankarasubramanian et al. (2004; 2006) and Rimmele \\& Sankarasubramanian (2004).  The diffraction limited imaging capability of the current day telescopes produced several recent papers on the fine scale structures of active as well as quiet regions. Langhans et al. (2005), Lites \\& Socas-Navarro (2004), Sankarasubramanian \\& Rimmele (2003), Sankarasubramanian, Rimmele, \\& Lites (2004), Rimmele \\& Sankarasubramanian (2004), and Rimmele \\& Marino (2006) are few examples from a long list. The positive aspects of some of these high resolution observations are the availability of vector magnetic field data close to the diffraction limit and the provision to observe Doppler velocities and other physical parameters near simultaneously. With several available space programs, it is now even possible to co-ordinate ground based observations with space borne instruments.  In this paper, we report on an observation of a decaying active region using the DST and the fixed port instrumentation. We used the DLSP along with G-band and Ca-K imagers to study the vector magnetic field and its morphology of the decaying active. H-alpha images were also simultaneously obtained using the UBF. These observations were also co-ordinated with TRACE observations, however the TRACE observations will not be discussed in this paper. Section 2 discusses the observing setup and details of the observations. Preliminary results from these observations are spelled out in section 3. ",
        "conclusions": "High spatial resolution observation of a decaying active region NOAA 0781 were studied. A set of instruments, DLSP, G-band \\& CaK imageries, and UBF for H-alpha images were used to obtain the observations. The data were then analysed for the morphological structures of the small scale fields present in and around this active region. The variations of the Stokes profiles observed in and around the UD and LB of this active region support the field free intrusion model suggested by Parker (1979). It is also observed that the edges of the LBs show strong linear polarisation signals compared to the center. The small scale fields surrounding this active region show opposite polarity profiles connected to the parent spot as well as to the surrounding magnetic field regions.  This active region was also monitored using the Michelson Doppler Imager (MDI) on-board SoHO. The spot decayed into plages few days after our observations. The MDI movie also showed bunch of opposite polarity moving magnetic features surging out of the parent spot. The opposite polarity small scale fields observed in the DLSP may belong to one such field region during the surge. A more detailed analysis combining the timing of the MDI observations with the DLSP observation is required to confirm this. The data from the co-ordinated TRACE observations may provide some more evidence on the morphological structures of these small scale loops at higher atmospheric layers."
    },
    "0801/0801.4017_arXiv.txt": {
        "abstract": "We present multi-waveband optical imaging data obtained from observations of the Subaru/XMM-Newton Deep Survey (SXDS). The survey field, centered at R.A. = $02^{h}18^{m}00^{s}$, decl. = $-05^{\\circ}00'00''$, has been the focus of a wide range of multi-wavelength observing programs spanning from X-ray to radio wavelengths. A large part of the optical imaging observations are carried out with Suprime-Cam on Subaru Telescope at Mauna Kea in the course of Subaru Telescope ``Observatory Projects''. This paper describes our optical observations, data reduction and analysis procedures employed, and the characteristics of the data products. A total area of 1.22 deg$^{2}$ is covered in five contiguous sub-fields, each of which corresponds to a single Suprime-Cam field of view ($\\sim34'\\times 27'$ ), in five broad-band filters $B, V, R_c, i', z'$ to the depths of $B=28.4, V=27.8, R_c=27.7, i'=27.7$ and $z'=26.6$ ($AB, 3\\,\\sigma, \\phi=2''$). The data are reduced and compiled into five multi-waveband photometric catalogs, separately for each Suprime-Cam pointing. The $i'$-band catalogs contain about 900,000 objects, making the SXDS catalogs one of the largest multi-waveband catalogs in corresponding depth and area coverage. The SXDS catalogs can be used for an extensive range of astronomical applications such as the number density of the Galactic halo stars to the large scale structures at the distant universe. The number counts of galaxies are derived and compared with those of existing deep extragalactic surveys. The optical data, the source catalogs, and configuration files used to create the catalogs are publicly available via the SXDS web page ({\\tt http://www.naoj.org/Science/SubaruProject/SXDS/index.html}). ",
        "introduction": "\\label{sec:intro} % Understanding the formation and evolution processes of the individual galaxy and the growth of the large scale structures (LSSs) in the universe is one of the major goals in extragalactic astronomy today. In the scheme of a typical $\\Lambda$-CDM model, the growth of structures is governed by the gravitational growth of initial fluctuations of dark matter. The baryonic material cools in these dark matter structures and grows through hierarchical clustering to galaxies and clusters of galaxies. The subsequent evolution of galaxies is closely connected to their environments, which, of course, relate to the LSSs, where those galaxies are located. The ultimate goal of extragalactic surveys is to trace these evolutionary processes by well-defined statistical galaxy samples.  Optical imaging is arguably the cornerstone of any extragalactic survey, since it provides identifications and positions of celestial objects for follow-up spectroscopy. In recent years, there have been a number of deep imaging survey projects which devote significant amounts of telescope time, such as Great Observatories Origins Deep Survey (GOODS: Giavalisco et al. 2004) and Cosmic Evolution Survey (COSMOS: HST treasury project: Feldmann et al. 2006; Scoville et al. 2007). These surveys provide multi-waveband galaxy samples at faint magnitude. Then, the photometric redshift techniques (Furusawa et al. 2000; Bolzonella et al. 2000; Feldmann et al. 2006) are frequently used to pre-select candidates for spectroscopy.  Although a great deal of information can be obtained from the optical data alone, the value of the data set grows significantly as data at other wavelengths from other facilities are added. We therefore elected to use a multi-wavelength approach for the Subaru/XMM-Newton Deep Survey (SXDS; Sekiguchi et al. 2007, in preparation) from the very start, ensuring that our chosen field would be accessible and suitable for observations at all wavelengths. An equatorial field, centered at R.A. = $02^{h}18^{m}00^{s}$, decl. = $-05^{\\circ}00'00''$ is chosen to tie up with the deep X-ray observations with XMM-Newton observatory (Ueda et al. 2007), and also due to the accessibility by all major observatories, both existing and planned. The major multi-wavelength observation programs on the SXDS field (hereafter SXDF) include deep radio imaging with the VLA (Simpson et al. 2006), sub-mm mapping with SCUBA (Coppin et al. 2006; Mortier et al., 2005), mid-infrared observations with Spitzer Space Telescope (Lonsdale et al. 2003), deep near-infrared imaging with the UKIRT/WFCAM (Foucaud et al. 2006; Dye et al. 2006; Lawrence et al. 2007), and the X-ray observations with XMM-Newton observatory. Importantly, the survey field of an infrared ultra deep survey (UDS; Foucaud et al. 2006) covering 0.77 square degrees as part of the UKIDSS project (Lawrence et al. 2007) is centered on the SXDF. It is expected that our extensive multi-wavelength data set will provide photometric redshifts accurate to $\\Delta z\\lesssim 0.1$ over a wide range of redshift, as well as detailed spectral energy distributions for the vast majority of the objects in the field.  The SXDS has been undertaken as a part of the Subaru Telescope ``Observatory Projects'', in which a large amount of observing times are devoted to carry out intensive survey programs by combining observing times rewarded to builders and observatory staff of the Subaru Telescope. Note that Subaru Deep Field (SDF; Kashikawa et al. 2004) is another observatory project, which targets on a field different from the SXDF. The data set of the SDF is slightly deeper than the SXDF, but concentrates on one-fifth narrower field than the SXDF. Thus, these surveys complement each other and they can be used for a wide variety of studies.  In this paper, we describe the observation details and data reduction and analysis of our optical imaging survey with the prime-focus camera on Subaru Telescope, as well as the resultant data products. A detailed description of the survey strategies, scientific objectives, and multi-wavelength survey plans is given in a companion paper (Sekiguchi et al., 2007 - Paper I, in preparation). In Section~2, we describe optical imaging observations by Subaru Telescope, and in Section~3 we explain the data reduction procedure employed.  We present creation of multi-waveband photometric catalogs including object detection and aperture photometry in Section~4, and calibrations of photometry and astrometry in Section~5. In Section~6, we discuss characteristics of the catalogs such as limiting magnitude, detection completeness, and number counts of galaxies.  The summary is given in Section~7.  Throughout this paper, all the magnitudes are expressed in the AB system unless otherwise mentioned. ",
        "conclusions": "% The optical imaging observations of the SXDS project are carried out using the Suprime-Cam on Subaru Telescope in $B, V, R_c, i'$ and $z'$ bands. The SXDF has a contiguous area coverage of $\\sim$1.2 deg$^{2}$, which consists of five Suprime-Cam pointings.  The photometric zeropoints are determined with an absolute accuracy of no larger than 0.05 mag r.m.s. in the photometry. The r.m.s. of the astrometric accuracies across the field is on the order of 0.2 arcsec in both RA and Dec.  The systematic differences in the adopted world coordinates between different pointings are within about 1 arcsec at the outer edge of the field of view. Thus, our SXDS catalogs have a good enough positional accuracies to perform follow-up spectroscopy.  Multi-waveband photometric catalogs of detected objects are created for each band and for each pointing. Each of the catalogs contains more than a hundred sixty thousand of objects.  The catalogs have quite homogeneous S/Ns across the field. The achieved limiting magnitudes in each band are $B=28.4, V=27.8, R_{c}=27.7, i'=27.7$ and $z'=26.6$ ($AB,  3\\,\\sigma, \\phi = 2''$).  The detection completeness as a function of magnitude is estimated by a Monte-Carlo simulation assuming a Gaussian profile. The number counts of galaxies in the SXDF in each band are computed by correcting the detection completeness, which are consistent with each other among the five pointings with a slight difference at faint magnitudes. The mean number counts of galaxies averaged over the five pointings show a good agreement with results from previous surveys down to the faint-end magnitude.  With the aid of the wide coverage area of SXDS data, the uncertainty of the number counts of galaxies due to the Poisson error is greatly decreased. It is emphasized that the SXDS data is extremely useful for pursuing studies on celestial objects spreading in a wide field without suffering from the Poisson error and the field-to-field variation. The SXDS catalogs can be applied to studies, scientific objectives of which range from the Galactic objects to the large scale structures of the universe.  The optical data, the compiled photometric catalogs, and configuration files used to create the catalogs have been released to the public and can be retrieved from the public data archives server of National Astronomical Observatory of Japan, {\\tt http://www.naoj.org/Science/SubaruProject/SXDS/index.html}.   \\smallskip  Acknowledgments  We thank Dr. Masafumi Yagi for his kind suggestions in analyzing the data based on the software package developed by himself. We thank Dr. Nobunari Kashikawa for fruitful discussions. The anonymous referee would also be appreciated for his/her constructive suggestions. We would like to convey our gratitude to the Subaru Telescope builders and staff for their invaluable help to complete the observing run with Suprime-Cam. Data reduction/analysis in this work was in part carried out on \"sb\" computer system operated by the Astronomical Data Center (ADC) and Subaru Telescope of the National Astronomical Observatory of Japan. Use of the UH 2.2-m telescope for the observations is supported by NAOJ."
    },
    "0801/0801.1404_arXiv.txt": {
        "abstract": "8-Gsps 1-bit Analog-to-Digital Converters (ADCs) were newly developed toward the realization of the wideband observation. The development of the wideband ADCs is one of the most essential developments for the radio interferometer. To evaluate its performance in interferometric observations, the time (phase) stability and frequency response were measured with a noise source and a signal generator. The results of these measurements show that the developed ADCs can achieve the jitter time less than 0.05 psec at the time interval of 1 sec, the passband frequency response with the slope less than 0.73 dB/GHz and the ripple less than 1.8 dB, and the aperture time less than 20 psec. The details of the developed ADC design, the measurement methods, and the results of these measurements are presented in this paper. ",
        "introduction": "A high-speed Analog-to-Digital Converter (ADC) can greatly contribute to the frontier of radio astronomy. The performance of ADC is very important because it decides the sensitivity of radio astronomical interferometric observations and the accuracy of astrometry observations and geodetic experiments. To realize the wideband observation for the Atacama Large Millimeter/Submillimeter Array (ALMA), the 2-bit ADC operating at the 4-GHz sampling rate have been developed \\citep{oki02}. A fringe spectrum that contains 20 line features as well as continuum emission at $\\sim$86.2 GHz was obtained though an interferometric observation of the Orion-KL regions using these ADCs connected to the Nobeyama Millimeter Array (NMA). This result shows that the realization of high frequency-resolution over a wide bandwidth is technically feasible \\citep{oku02}. This is the first time to obtain such a wideband astronomical fringe from a \\emph{single} correlation product of digitized signals; 20 line features can be found by a careful inspection of the spectrum.  In this paper, the development of newly-developed 8-Gsps 1-bit ADC is described in Section 2, the measurement methods and the performance of the ADC in Section 3, and the conclusion in Section 4. ",
        "conclusions": "This paper has described the details of the developed 8-Gsps 1-bit ADCs and presented the measurement methods and the performances of the ADCs. Summary of this paper is as follows. \\begin{itemize} \\item The ADCs were developed with the commercially available DEC and DMUX at a bit rate of 10 Gbps. Input signals with power of $-20$ dBm are quantized with a sampling frequency of 8192 MHz or 4096 MHz. \\item The Allan standard deviation of the ADC boards indicates that their performance is stable in the timescale of $\\sim10^4$ sec. The jitter time in one second is estimated to be less than 0.05 psec. \\item There are slopes of $\\leq0.73$ dB/GHz in the passband responses of the developed ADC boards. The peak-to-peak values of the ripple are less than $1.8$ dB. \\item The aperture time of the developed ADCs was derived from power spectra when a strong CW signal is input into the developed ADC boards. The aperture time of the ADCs is estimated to be less than 20 psec. \\end{itemize} The performance of the developed ADCs is summarized in Table \\ref{tab:performance}."
    },
    "0801/0801.4826_arXiv.txt": {
        "abstract": "It has long been known how to  analytically relate the clustering properties of the collapsed structures (halos)  to those of the underlying dark matter distribution for Gaussian initial conditions. Here we apply the same approach to physically motivated non-Gaussian models. The techniques we use were developed in the 1980s to deal with the clustering of peaks of non-Gaussian density fields. The description of the clustering of halos for non-Gaussian initial conditions has  recently received renewed interest, motivated by the forthcoming  large galaxy and cluster surveys. For inflationary-motivated non-Gaussianites, we find an analytic expression for  the halo bias  as a function of scale, mass and redshift, employing only the approximations of high-peaks and large separations. ",
        "introduction": "Constraining primordial non-Gaussianity offers a powerful test  of the generation mechanism of cosmological perturbations in the early universe. While standard single-field models of slow-roll inflation lead to small departures from Gaussianity, non-standard scenarios  allow for a larger level of non-Gaussianity (\\citet{BKMR04} and references therein). The standard observables to constrain non-Gaussianity are the cosmic microwave background and large-scale structure. A powerful technique is based on the abundance \\citep{MVJ00, VJKM01, Loverdeetal07, RB00, RGS00} and clustering \\citep{GW86, MLB86, LMV88} of rare events such as  dark matter density peaks  as they trace the tail of the underlying distribution.  These theoretical predictions have been tested against numerical N-body simulations \\citep{KNS07,Grossietal07,DDHS07}.  \\citet{DDHS07} showed that primordial non-Gaussianity affects the clustering of dark matter halos inducing a scale-dependent bias.  This effect will be useful for constraining non-Gaussianity from  future surveys which will provide a large sample of galaxy clusters over a volume comparable to the horizon size (e.g., DES, PanSTARRS, PAU, LSST, DUNE, ADEPT, SPACE, DUO) or mass-selected  large clusters  samples via the Sunyaev-Zel'dovich effect (e.g., ACT, SPT),  considered alone or via cross-correlation techniques (e.g., ISW, lensing).  Here, we resort to results and techniques developed in the 1980s \\citep{GW86, MLB86, LMV88} to  extend this work and derive an accurate analytical expression for halo bias,  in the presence of general non-Gaussian initial conditions, accounting for  its scale, mass and redshift dependence. ",
        "conclusions": "We have obtained an  analytic expression for the bias of dark matter halos for non-Gaussian  initial conditions. The only approximations used in our approach are: {\\it i)} high peaks, {\\it i.e.} large values of $\\nu=\\delta_c(z)/\\sigma_R$ (as in the original Kaiser's formula), which essentially amounts to a limitation on the mass range over which one can apply this formula, and {\\it ii)} large separation among the halos, which is the standard assumption allowing to use linear bias.  While it is true that on large scales ($k\\rightarrow 0$)  the form factor, the transfer function and the window function  go to unity, on the scales of interest  neglecting these terms may lead to errors on $\\Delta b_h$ and therefore on $f_{\\rm NL}$ of the order of 100\\%. Comparison of these analytical findings with simulations  will be presented elsewhere (Grossi {\\it et al.}, in preparation).   An advantage of our approach is that it  can be easily generalized to non-local and scale-dependent non-Gaussian models in which  $B_{\\phi} (k_1, k_2,k_3)$ is the dominant higher-order correlation and  has a general form, obtaining \\begin{eqnarray} \\frac{\\Delta b_h}{b_h}&=& \\frac{\\Delta_c(z)}{D(z)}\\frac{1}{8\\pi^2\\sigma_R^2}\\int dk_1 k_1^2 {\\cal M}_R(k_1) \\times \\nonumber \\\\ &&\\!\\int_{-1}^1\\!d\\mu {\\cal M}_R\\left(\\sqrt{\\alpha}\\right)\\frac{B_{\\phi}(k_1,\\sqrt{\\alpha},k)}{P_{\\phi}(k)}. \\label{eq: fnlk} \\end{eqnarray} Modeling the clustering of hot and cold CMB spots for non-Gaussian initial conditions, is a straightforward extension of this calculation (Heavens {\\it et al.}, in preparation).  We envision that this calculation will be useful for constraining non-Gaussianity from  future surveys which will provide a large sample of galaxy clusters over a volume comparable to the horizon size (e.g., DES, PanSTARRS, PAU, LSST, DUNE, ADEPT, SPACE, DUO) or mass-selected (via the Sunyaev-Zel'dovich effect)  large clusters  samples (e.g., ACT, SPT)."
    },
    "0801/0801.3895_arXiv.txt": {
        "abstract": "\\ifdraft \\par\\hm\\par \\vspace{-13.5ex} \\fbox{\\Large\\sf Draft of \\isodate } \\par\\hm\\par \\vspace{4ex} \\fi  This paper presents the sixth part to the Very Long Baseline Array (VLBA) Calibrator Survey. It contains the positions and maps of 264 sources of which 169 were not previously observed with very long baseline interferometry (VLBI). This survey, based on two 24 hour VLBA observing sessions, was focused on 1)~improving positions of 95 sources from previous VLBA Calibrator surveys that were observed either with very large a~priori position errors or were observed not long enough to get reliable positions and 2)~observing remaining new flat-spectrum sources with predicted correlated flux density in the range 100--200~mJy that were not observed in previous surveys. Source positions were derived from astrometric analysis of group delays determined at the 2.3 and 8.6~GHz frequency bands using the Calc/Solve software package. The VCS6 catalogue of source positions, plots of correlated flux density versus projected baseline length, contour plots and fits files of naturally weighted CLEAN images, as well as calibrated visibility function files are available on the Web at \\url{http://vlbi.gsfc.nasa.gov/vcs6}. ",
        "introduction": "\\label{s:introduction}  This work is a continuation of the project of surveying the sky for bright compact radio sources. These sources can be used as phase referencing calibrators for imaging of weak objects with very long baseline interferometry (VLBI) and as targets for space navigation, monitoring the Earth's rotation, differential astrometry, and space geodesy. Precise positions of radio sources are needed for these applications. Since 1979 more than 4400 24~hour VLBI experiments have been scheduled in geodetic mode. These observations also determined parameters associated with the Earth orientation and rotation, antenna position and motions and other astrometric/geodetic parameters. By November 01, 2007, 1137 sources were observed in these experiments and 1045 of them were detected. The observations are described by \\citet{icrf98} and the latest catalog of 776 sources is given as the ICRF-Ext2 catalog \\citep{icrf-ext2-2004}.  Because the astrometric and image quality of the radio sources (targets) improve with decreasing angular separation of the calibrator and the target, a higher density of sources than that from the ICRF catalogue was needed. Since 1994, 22 dedicated 24-hr experiments with the VLBA, called the VLBA calibrator survey (VCS), were made in order to increase the density of suitable VLBI phase reference sources \\citep{vcs1,vcs2,vcs3,vcs4,vcs5}. In these experiments, 3601 sources with declinations $> -40^\\circ$ were observed, including 764 objects previously observed under geodetic programs, and 3301 of them were detected. The catalogue of source positions derived from a single least square (LSQ) solution using all geodetic and VCS observations forms the pool of sources with positions at a milliarcsec level of accuracy that is widely used for phase referencing observations, for statistical analysis, and for other applications. Improving the precision of this catalogue and an increase in the number of objects is an important task.  In this paper we present an extension of the VCS catalogue, called the VCS6 catalogue. The objectives of this campaign were to improve the catalogue of compact radio sources. Our approach to this problem was 1)~to improve positions of sources observed in the previous VCS campaigns that had a)~poor a~priori positions; b)~were not observed long enough in the first two VCS1 24-hour experiments recorded at the 64~Mbit/sec rate; and 2)~to get positions of new sources. The new sources were taken from the lists of a)~intra-day variable sources observed in the framework of the MASIV survey \\citep{masiv-2003}, not observed before with the VCS, and b)~leftovers from the list of candidate sources prepared for the VCS5 campaign \\citep{vcs5} that were not observed due to lack of resources.  Since the observations, calibrations, astrometric solutions and imaging are similar to that of VCS1--5, most of the details are described by \\cite{vcs1}, \\cite{vcs3}, and \\cite{vcs4}. In \\S\\ref{s:selection} we describe the strategy for selecting 285 candidate sources observed in two 24~hour sessions with the VLBA. In \\S\\ref{s:obs_anal} we briefly outline the observations and data processing. We present the VCS6 catalogue in \\S\\ref{s:catalogue}, and summarize our results in \\S\\ref{s:summary}. ",
        "conclusions": "\\label{s:summary}  The VCS6 Survey has added 169 new compact radio sources not previously observed with VLBI and significantly improved coordinates of 95 other objects. Among the new sources, 103 objects turned out to be suitable as phase referencing calibrators and as target sources for geodetic applications. Their coordinates have position accuracy better than 5~mas."
    },
    "0801/0801.1068_arXiv.txt": {
        "abstract": "We present CCD photometric and mass function study of 9 young Large Magellanic Cloud star clusters namely \\CLN. The \\ubvri data reaching down to $V \\sim$ 21 mag, are collected from 3.5-meter NTT/EFOSC2 in sub-arcsec seeing conditions. For NGC 1767, NGC 1994, NGC 2002, NGC 2003, NGC 2011 and NGC 2136, broad band photometric CCD data are presented for the first time. Seven of the 9 clusters have ages  between 16 to 25 Myr while remaining two clusters have ages $32\\pm4$ Myr (NGC 2098) and $90\\pm10$ Myr (NGC 2136). For 7 younger clusters, the age estimates based on a recent model and the integrated spectra are found to be systematically lower ($\\sim$ 10 Myr) from the present estimate. In the mass range of $\\sim 2 - 12$ $M_{\\odot}$, the MF slopes for 8 out of nine clusters were found to be similar with the value of $\\gamma$ ranging from $-1.90\\pm0.16$ to $-2.28\\pm0.21$. For NGC 1767 it is flatter with  $\\gamma = -1.23\\pm0.27$. Mass segregation effects are observed for NGC 2002, NGC 2006, NGC 2136 and NGC 2098. This is consistent with the findings of \\citet{kontizas98} for NGC 2098. Presence of mass segregation in these clusters could be an imprint of star formation process as their ages are significantly smaller than their dynamical evolution time. Mean MF slope of $\\gamma = -2.22\\pm0.16$ derived for a sample of 25 young ($\\le 100$ Myr) dynamically unevolved LMC stellar systems provide support for the universality of IMF in the intermediate mass range $\\sim 2-12\\ M_{\\odot}$. ",
        "introduction": "\\label{sec:intro}  The distribution of stellar masses that form in one star-formation event in a given volume of space is called initial mass function (IMF). Some theoretical studies consider that the IMF should vary with the pressure and temperature of the star-forming cloud in such a way that higher-temperature regions ought to produce higher average stellar masses while others have exactly opposite views (see \\citealt{larson98, elmegreen00} and references therein). Detailed knowledge of the IMF shape in different star forming environments is therefore essential. One would like to know whether it is universal in time and space or not. For this, in a galaxy, its young (age $\\le$ 100 Myr) star clusters of different ages, abundance etc. need to be observed, as they contain dynamically unevolved, (almost) coeval sets of stars at the same distance with the same metallicity. For a number of such reasons, populous young star clusters of the Large Magellanic Clouds (LMC) are the most suitable objects for investigating the IMF. They provide physical conditions not present in our Galaxy e.g. stellar richness, metallicity and mass ranges (see \\citealt{sagar93, sagar95} and references therein). Unlike the galactic counterparts, where corrections for interstellar absorption are not always trivial since it could be large as well as variable \\citep{sagar87, yadav01, kumar04}, for LMC star clusters it is relatively small. Its treatment is therefore not a problem. Furthermore, choosing young (age $\\le 100$ Myr) clusters reduces the effects of dynamical evolution on their MF. The present day mass function of these stellar systems can therefore be considered as the IMF. The study of young LMC star clusters is thus important for providing the answer to the question of universality of the IMF. Both ground and Hubble Space Telescope (HST) observations have therefore been obtained (see \\citealt{sagar91a, brocato01, matteucci02} and references therein) for a few of the large number of young LMC star clusters \\citep{bica99}. The potential offered by them  has not been fully utilised as still a large fraction of them are unobserved.  \\begin{figure} \\centering \\includegraphics[width=9cm]{radec.eps} \\caption{Small dots show the location of identified LMC star clusters from the catalog of \\citet{bica99}. A sky area of about 8\\degr $\\times$ 8\\degr is shown centered around the optical center ($\\alpha_{\\rm J2000} =5^h 20^{m} 56^{s}$, $\\delta_{\\rm J2000} = -69^{\\degr} 28^{\\arcmin} 41^{\\arcsec}$) of the LMC. The bar region is clearly seen. The target clusters are shown with filled triangles.} \\label{fig:radec} \\end{figure}   In this paper we derive mass function (MF) slopes using new broad band $BVRI$ CCD photometric observations of the stars in 9 young LMC star clusters namely \\CLN. Their integrated photometric colours indicate that all of them belong to SWB \\citep{searle80} class of 0 or I and hence are very young with ages $\\le 30$ Myr \\citep{elson95, bica96} except NGC 2136 which belongs to SWB class of III indicating and age between 70\\,-\\,200 Myr. Table \\ref{tab:sample} lists the relevant information available prior to this study. All the clusters are rich indicating higher stellar density \\citep{bica95} and thus most suitable for the MF study. Except NGC 2011, all are elliptical in size with major axes diameters ranging from 1\\farcm3 to 2\\farcm8. Except for NGC 2002 and NGC 2098, all others are candidate members of either a pair or a multiple system (see \\citealt{dieball02}). Locations of the target clusters are shown in Fig. \\ref{fig:radec}. Most of them lie towards north-east side of the LMC bar harboring young star forming regions in contrast to the intermediate age (1 - 3 Gyr) cluster field of the bar. NGC 1767 lies south-west region of the LMC bar. Being spread over a wide region ($\\sim$ 5\\degr $\\times$ 10\\degr), the sample may reflect different star forming environments. It is therefore suitable for testing the universality of the IMF. When CCD observations were carried out in 1990, no detailed photometric observations and MF studies had been published. However, in the mean time some CCD photometric observations have been published for a few of the clusters under study. A brief description of the previous work on the clusters under study is given below.  \\input{./sample.tab}   \\subsection{Previous work}  \\begin{itemize}  \\item  {\\bf NGC 1767.} This, a member of triple star cluster system, is located in the OB association LH 8. Integrated ($U-B$) and ($B-V$) colors indicate that the cluster is young with an age of $\\sim$ 10 Myr.  \\item {\\bf NGC 1994.} This, located in LMC DEM 210 region, is a member of a 5-cluster system. Its irregular size is largest amongst them. An age of about 5 - 30 Myr has been derived for the cluster from its integrated photometric colour observations.  \\item {\\bf NGC 2002.} This single cluster is  located in the OB association LH 77 in the supergiant shell LMC 4 region. The cluster center is condensed,  but the outer part is resolved. Integrated light observations indicate an age of $\\sim$ 10 - 30 Myr along with the presence of a few red supergiants \\citep{bica96}.  \\item {\\bf NGC 2003.} Integrated photometric observations indicate an age of 10 - 30 Myr for this cluster located in the Shapley III region of the LMC. Its shape on the photographic image is elongated with resolved outer parts.  \\item {\\bf NGC 2006 and SL 538.} This binary star cluster is located in the northwestern  part of the OB association LH 77 in supergiant shell LMC 4. The clusters are separated by $\\sim 55''$ on the sky corresponding to a linear separation of 13.3 pc at the distance of LMC. Integrated photometric observations obtained by \\citet{bhatia92} and \\citet{bica96} indicate a similar age for both the clusters. Using low-resolution objective prism spectra and integrated IUE spectra, \\citet{kontizas98} suggested that this binary cluster may merge in $\\sim$ 10 Myr. Broad band and H$_{\\alpha}$ CCD photometric observations were obtained  by \\citet{dieball98}. Based on the colour-magnitude diagrams (CMDs) of the clusters they derived an age of $18\\pm2$ Myr for SL 538 and of $22.5\\pm2.5$ Myr for NGC 2006. The MF slopes obtained for both the clusters were consistent with that of \\citet{salpeter55} and indicated similar total masses. These studies thus indicates near-simultaneous formation of the cluster pair in the same giant molecular cloud.  \\item  {\\bf NGC 2011.} This is located in the OB association region LH 75. Its age estimated from the integrated photometric observations is between 10 to few tens of Myr. Its photographic image indicates that it is elongated, fairly condensed and partly resolved cluster. A recent analysis of its stellar content using HST observations reveal that it has two parallel main sequence branches, and may be a binary system \\citep{gouliermis06}. However, the analysis also indicate that both populations might have formed in a single star forming event as the redder stars are situated in the central half arcmin region and are thought to be embedded in the dust and gas, while the blue stars are spread in the outer region up to 1 arcmin.  \\item {\\bf NGC 2098.} This is another single cluster amongst the objects under study. The first $BR$ broad band CCD photometric observations have been presented by \\citet{kontizas98}. They derived an age of 63 - 79 Myr and found strong evidence for mass segregation in agreement with their earlier studies based on the photographic observations. However, poor quality of their CCD data was indicated by the authors.  \\item  {\\bf NGC 2136.} This is the brighter component of the young binary globular cluster NGC 2136/ NGC 2137 in the LMC. The angular separation between the components is about 1\\farcm3. \\citet{hilker95} using Stromgren CCD photometry of the clusters indicates their common origin. They indicate an age of 80 Myr and metallicity [Fe/H] = $-0.55\\pm0.06$ dex for the cluster while \\citet{dirsch00} derive an age of $100\\pm20$ Myr but the same metallicity. The cluster contains a number of Cepheids as well as red giants.  \\end{itemize}  The present CCD observations, in combination with earlier observations, have been used to estimate and/or interpret the interstellar reddening to the cluster regions, ages and mass functions of the clusters. Section 2 deals with the observational data, reduction procedures and comparisons with the published photometric data. In section 3, we analyse the stellar surface density profiles, CMDs and MFs of the sample clusters. Last section is presented with the results and discussions.  \\input{./obslog.tab} ",
        "conclusions": "We present \\ubvri CCD data obtained from 3.5-meter ESO NTT/EFOSC2 observations for 9 young Large Magellanic Cloud star clusters namely \\CLN\\ and their nearby field regions reaching down to $V \\sim 20$ mag for $\\sim$ 6400 stars altogether. They are the first accurate broad band CCD photometric data for all the clusters except for the binary cluster NGC 2006 and SL 538. The observations are made in a region of $\\sim 2\\arcmin \\times 2\\arcmin$ around the cluster center. The data were collected during Jan 10 to Jan 13, 1990 in good seeing conditions ranging from 0\\farcs7 to 1\\farcs0 and reduced using DAOPHOT and MIDAS softwares. Photometric calibrations are done using Landolt (1992) stars and the zero point accuracy is better than $0.02$ mag. Photometric errors become large (\\ga 0.1 mag) for stars fainter than $V = 20$ mag.  We examine radial density profiles, general features of the main sequence and estimate age and reddening for individual clusters using Padova isochrones. The various CMDs of the clusters under study were used to estimate their MF, age and reddening. In order to study radial variation in MF, the LFs are derived for inner, outer, and entire cluster regions. Due to compactness of the clusters, such study could not be carried out for the core regions of the clusters. The LFs are corrected for both data incompleteness and field contamination. The main conclusions of the present study are as follows.  \\begin{enumerate}  \\item Seven of the nine clusters have ages $\\le 25$ Myr, while the remaining two clusters have ages of $32\\pm4$ Myr (NGC 2098) and $90\\pm10$ Myr (NGC 2136). Our age estimates for NGC 2006 and SL 538 were found to be consistent with the previous $BVR$ photometric estimate by \\citet{dieball98}. For NGC 2098, our estimates are lower by about 30 Myr than \\citet{kontizas98}. Thus, the ages of all the clusters in our sample are significantly lower than their typical dynamical ages of a few 100 Myrs.  \\item For younger ($\\le$ 25 Myr) clusters, the age estimates based on a recent population synthesis models by \\citet{wolf07} and integrated spectra are systematically lower by about 10 Myr than the present age estimates based on CMDs.  \\item Assuming an LMC distance modulus of 18.5 mag, the derived reddening for the  clusters in our sample was found to be consistent with that derived from HI emission and 100 $\\mu m$ all sky dust maps.  \\item In the mass range of $2 - 12\\ M_{\\odot}$, the MF slopes for 8 out of 9 sample clusters were found to be similar with values of $\\gamma$ ranging from  $-1.90\\pm0.16$ to  $-2.28\\pm0.21$. For NGC 1767 the slope was found to be significantly shallower with  $\\gamma = -1.23\\pm0.25$. The present MF values are consistent with those derived by \\citet{dieball98} for NGC 2006 and SL 538. \\citet{selman05} studied the star formation history and IMF of the field population of 30 Doradus super association and found that it has a Salpeter slope in the mass range of 7 to 40 $M_{\\odot}$.   \\item Mass function slopes of the inner and outer cluster regions indicates the presence of mass segregations in NGC 2002, NGC 2006, NGC 2136 and NGC 2098. For NGC 2098, \\citet{kontizas98} derive the dynamical relaxation time, $T_{\\rm e}$ between 640 to 1050 Myr. This may indicate that the value of $T_{\\rm e}$ for LMC star clusters could be few hundreds of Myr. The ages of LMC star clusters under study are therefore significantly smaller than their dynamical relaxation time. Consequently, observed mass segregation in these clusters is probably  primordial in nature. A compilation of both ground and space based observations of extremely young galactic and MC star clusters (cf.  \\citealt{hunter95}; \\citealt{sagar88}; \\citealt{hillenbrand98}; \\citealt{chen07} and references therein) indicates presence of mass segregation in most of them, although to varying degrees. All these indicate that in most of the young star clusters located in different galaxies, mass segregation effects are observed and most likely they are imprint of star formation processes.  \\item A mean MF slope of $\\gamma = -2.22\\pm0.16$ derived for a sample of 25 young ($< 100$ Myr) stellar systems in LMC provide support for the universality of IMF in the intermediate mass range $\\sim 2-10\\ M_{\\odot}$. An IMF study of the 30 Doradus star forming region of LMC by \\citet{selman05} also support this conclusion.  \\end{enumerate}"
    },
    "0801/0801.0402_arXiv.txt": {
        "abstract": "We present the Spitzer Infrared Spectrograph (IRS) spectrum of SR20, a 5--10 AU binary T Tauri system in the $\\rho$ Ophiuchi star forming region.  The spectrum has features consistent with the presence of a disk; however, the continuum slope is steeper than the $\\lambda^{-4/3}$ slope of an infinite geometrically thin, optically thick disk, indicating that the disk is outwardly truncated.  Comparison with photometry from the literature shows a large increase in the mid-infrared flux from 1993 to 1996.  We model the spectral energy distribution and IRS spectrum with a wall $+$ optically thick irradiated disk, yielding an outer radius of 0.39$_{+0.03}^{-0.01}$ AU, much smaller than predicted by models of binary orbits.  Using a two temperature $\\chi^2$ minimization model to fit the dust composition of the IRS spectrum, we find the disk has experienced significant grain growth: its spectrum is well-fit using opacities of grains larger than 1 $\\mu$m.  We conclude that the system experienced a significant gravitational perturbation in the 1990s. ",
        "introduction": "Classical T Tauri stars are typically surrounded by optically thick accretion disks, indicated by strong and broad H$\\alpha$ emission lines and by a characteristic infrared excess.  Over the last twenty years, observations have revealed that many of these young stars are components in binary, and sometimes higher order, systems \\citep[e.g.][]{simon92, ghez93}. A binary system may produce up to three distinct disks: a circumprimary, circumsecondary, and circumbinary disk \\citep[][(AL94; JM97)]{artymowicz94, jensen97}.  Gravitational interactions between binaries affect the structure of their disks; in particular, the maximum outer radius for a circumstellar disk in a binary system is approximately 0.18 to 0.4 times the semi-major axis (AL94), depending on the mass ratio and orbital eccentricity.  SR20 is one such close binary in the $\\rho$ Ophiuchi star forming region.  Its components have an angular separation ranging from 0.038 to 0.071\\arcsec \\citep{ghez93, ghez95} which, at a distance to Ophiuchus of 140 pc \\citep{dz99} corresponds to a projected linear separation of 5.3--9.9 AU.  The secondary is 2.2 magnitudes fainter than the primary at 2.2 $\\mu$m \\citep{ghez93} and therefore is less massive than the primary, diskless, or both.  H$\\alpha$ emission has been detected from SR20 with equivalent widths ranging from 15 to 21 \\AA \\citep[][(W05)]{r80, ba92, wilking05}.  Further observations by \\citet{tbc03} did not find a spectro-astrometric signature in the H$\\alpha$ line, indicating that there is little to no accretion in the secondary relative to the primary. Based on the lack of evidence for a circumsecondary disk, we assume that the infrared excess of SR20 is dominated by a circumprimary accretion disk.  Here we present Spitzer Space Telescope Infrared Spectrograph (IRS) \\citep{houck04} observations of the SR20 system. We model the structure of the system and the dust it contains and analyze the nature of the SR20 system. While instances of sculpture of proto-planetary accretion disks by planetary or stellar companions \\emph{within} the disk are becoming widely known \\citep[e.g. JM97,][]{furlan07}, SR20 is an apparently rarer example of sculpture from \\emph{without}. ",
        "conclusions": "\\label{discussion}  The SED, system geometry, and dust composition shed light on the nature of the SR20 system, while simultaneously raising further questions.  The IRS spectrum is well-fit by our composite wall $+$ circumprimary disk truncated beyond 0.39 AU.  Outward truncation could be caused by the orbit of a companion or lack of replenishment from a circumbinary disk.  Using the tidal truncation models of AL94, the range of likely eccentricities (0.2-0.3) and semi-major axes (8-12 AU) of the A-B orbit are not extreme enough to cause truncation.  Even with the lower limit on the semi-major axis, 5 AU, and the highest order resonances and eccentricities tested in AL94, the minimum stable outer truncation radius caused by the known binary is 0.9 AU, considerably larger than the 0.39 AU we derive (AL94).  However, the AL94 models may not be appropriate for a disk with such a small ratio between the outer and inner radii, so we cannot {\\it absolutely} rule out truncation by the known binary, at least on this basis.  Alternatively, it may be that SR20 is a hierarchical triple.  A companion object with an eccentricity $<$ 0.8 could orbit in the region between 0.98 AU and 2.17 AU from the primary, causing a disk truncation at 0.39 AU.  Such a close, low mass companion would be difficult to detect.  Based on the 0.39 AU truncation radius, we speculate that replenishment of the circumprimary disk by a circumbinary disk occurred in the past. A disk between 0.1 and 10 AU would accrete on a timescale of 1,000--100,000 years \\citep{quillen04} and the median age of the Ophiuchus cluster is 2.1 Myr (W05). SR20 has been observed at 800, 850, and 1300 $\\mu$m with only 3 $\\sigma$ upper limit results, implying less than $5.4$ $\\times$ $10^{4}$ $M_{lunar}$ (540 $M_{lunar}$ of dust) in disk material \\citep{aw07}.  Low circumbinary disk masses are typical of binaries with semi-major axes between 1 and 50 AU \\citep{jensen94}, so submillimeter non-detection does not indicate a depleted circumbinary disk; however, it is unlikely that lack of replenishment is responsible for the 0.39 AU truncation radius.  A nearby companion truncating the circumprimary disk is perhaps also the best explanation for the photometric variability.  Before 1993, the infrared excess in SR20 started in the N band, indicating that there were few small dust grains close to the central star. However, the H$\\alpha$ emission indicates that a gaseous accretion disk, perhaps mixed with much larger planetesimals, was present. Between 1993 and 1996, the infrared excess increased sharply, indicating an increase in the amount of dust in the disk, which could be explained by a gravitational interaction between the known binary and a third, unseen, body, causing the unseen companion to orbit closer to the primary, perturbing the circumprimary disk.  Coupling between the gas and dust could confine the newly created dust to the accretion disk. The appeareance of the observed excess over 3 years, preceded by 20 years without an excess and followed by 10 years of little variation in the observed excess, seems consistent with truncation by a nearby source with a shorter period than is plausible for the known companion.  The silicate dust composition of the system is not typical of a Class II object; our fit indicates that the optically thin portion of the disk is comprised primarily of large grains ($>$ 1 $\\mu$m).  Although having a large fraction of large grains is consistent with debris disks or advanced dust processing in the inner regions of a protoplanetary disk \\citep{chen06}, this is the first time we have seen a Class II disk with negligible submicron grains.  Taken with the aforementioned explanations of the 1993--1996 variability, the large grains could support a collision scenario:  gravitional interactions between the known secondary and the unseen companion cause collisions between planetesimals, producing dust grains that are detected as the current disk. If there were a population of planetesimals -- invisible in our spectra -- comprising an extension of the disk beyond the outer truncation radius of the dust, truncation by the known companion would be more probable.  Near infrared interferometry and orbital monitoring would be necessary to determine the nature of this complex system more precisely."
    },
    "0801/0801.2631_arXiv.txt": {
        "abstract": "We present theoretical evolutionary sequences of intermediate mass stars (M=3-6.5$\\msun$) with metallicity Z=0.004. Our goal is to test whether the self-enrichment scenario by massive Asymptotic Giant Branch stars may work for the high metallicity Globular Clusters, after previous works by the same group showed that the theoretical yields by this class of objects can reproduce the observed trends among the abundances of some elements, namely the O-Al and O-Na anticorrelations, at intermediate metallicities, i.e [Fe/H]=-1.3. We find that the increase in the metallicity favours only a modest decrease of the luminosity and the temperature at the bottom of the envelope for the same core mass, and also the efficiency of the third dredge-up is scarcely altered. On the contrary, differences are found in the yields, due to the different impact that processes with the same efficiency have on the overall abundance of envelopes with different metallicities. We expect the same qualitative patterns as in the intermediate metallicity case, but the slopes of some of the relationships among the abundances of some elements are different. We compare the sodium--oxygen anticorrelation for clusters of intermediate metallicity (Z$\\approx 10^{-3}$) and clusters of metallicity large as in these new models. Although the observational data are still too scarce, the models are consistent with the observed trends, provided that only stars of M$\\simgt$5\\msun\\ contribute to self--enrichment. ",
        "introduction": "Spectroscopic investigations of Globular Clusters (GC) show star to star differences in their surface chemistries (Kraft 1994). These inhomogeneities trace clear abundance patterns involving all the light elements: sodium is correlated to aluminium and anticorrelated to oxygen, whereas magnesium is anticorrelated with aluminium (Carretta 2006), though in some clusters the existence of this latter relationship is still under debate (Cohen \\& Mel\\'endez 2005). A common result of these analysis is the approximate constancy of the overall CNO abundances (Ivans et al. 1999).  The detection of such anomalies in scarcely evolved stars, like Turn-Off (TO) and Sub-Giant Branch (SGB) stars (Gratton et al 2001), ruled out the possibility of any in situ mechanism as a possible unique explanation (Denissenkov \\& Weiss 2004), and pointed in favour of a self-enrichment scenario, i.e. that the stars with the anomalous chemistry formed from the ashes of a previous stellar generation, which polluted the interstellar medium with matter which had been processed via the CNO cycle. Ventura et al. (2001) indicated the massive Asymptotic Giant Branch (AGB) stars as likely candidates, because with appropriate hypothesis concerning convection modelling they can easily achieve Hot Bottom Burning (HBB), i.e. the bottom of their envelope becomes so hot to ignite strong nucleosynthesis, whose products, convected to the surface of the star, are ejected into the interstellar medium via the strong winds suffered by these sources (Ventura \\& D'Antona 2005); also, strong HBB favours high mass loss, reduces the number of thermal pulses (TP) experienced by the stars during their life, thus diminishing the number of third dregde-up (TDU) episodes, and keeping approximately constant the C+N+O mass fraction in the envelope, for a reasonably large range of initial masses. This conclusion is at odds with the results by Fenner et al. (2004), who, based on AGB models calculated with the traditional, low efficiency, mixing length theory (MLT, Vitense 1953) convective model, found that oxygen can be hardly depleted at the surface of these stars, and the CNO sum is expected to largely exceed unity.  Recently, an alternative self-enrichment scenario for GCs was proposed by Maeder \\& Meynet (2006), Prantzos \\& Charbonnel (2006), and described in details by Decressin et al. (2007): in this case the enrichment of the interstellar medium is produced by the envelopes of fast rotating massive stars. In order to choose between the AGB and the massive stars self-enrichment, it is necessary to explore in detail both models. In particular, the variation of predictions of yields with metal abundance must be examined. Ventura \\& D'Antona (2008, hereinafter VD08) showed that the enrichment by AGBs can work in the case of intermediate metallicity GCs, since the ejecta of their theoretical models, calculated with a metallicity Z=0.001, with initial masses in the range 5-6$\\msun$, are in qualitative and quantitative agreement with the abundance patterns observed in scarcely evolved stars of M3, M13, M5, NGC6218, NGC6752.  The goal of the present paper is to investigate whether the self-enrichment scenario hypothesis may work also in the case of metal rich GCs, having metallicities [Fe/H]$\\sim -0.7$. We present a new set of AGB models with metallicity Z=4$\\times 10^{-3}$, typical of metal rich GCs like 47Tuc and M71. We discuss the possibility of reproducing the O-Al and O-Na anticorrelations, and determine the expected slope of these relationships compared to the intermediate metallicity case.  \\begin{figure} \\includegraphics[width=.48\\textwidth]{fig1.eps} \\caption{The variation of the temperature at the bottom of the convective envelope as a function of the core mass at the beginning of the AGB phase (Top) and of the initial mass (Bottom) of AGB models with two different metallicities. The temperatures refer to the phase of maximum luminosity during the AGB evolution.} \\label{phys} \\end{figure} ",
        "conclusions": "We present new AGB models with metallicity Z=0.004, to test the hypothesis of AGB self-enrichment in high metallicity GCs and compare the results with previous models having Z=0.001 (VD08). We find that the main physical features of the AGB evolution of the intermediate mass stars do not change by increasing the metallicity, the main effects being only a shift towards higher masses of the mass-luminosity relationship; models with the same core mass and different Z follow very similar evolution, in terms of the temperature reached at the bottom of their convective envelope.  The theoretical yields of the most massive AGB models are in agreement with the aluminum, sodium and oxygen abundances of the most anomalous stars (those showing the strongest depletion of oxygen ---and excluding the giant in which the very low O should be attributed to deep extramixing) in the few high metallicity GCs so far examined; a reasonable dilution scheme gives results consistent with the O--Na and C--N anticorrelations. Yet, other confirmation is needed, since the data currently available involve only few stars per cluster.  One robust prediction of this investigation, which is independent of all the uncertainties associated to the proton capture cross sections by Ne-Na-Al nuclei, is that with increasing metallicity we expect a higher slope of the O-Na anticorrelation, and a similar slope (though less extended) of the O-Al trend."
    },
    "0801/0801.0634_arXiv.txt": {
        "abstract": "The size distribution of mini-filaments in voids has been derived from the Millennium Run halo catalogs at redshifts $z=0,0.5,1$ and $2$. It is assumed that the primordial tidal field originated the presence of filamentary substructures in voids and that the void filaments have evolved only little, keeping the initial memory of the primordial tidal field. Applying the filament-finding algorithm based on the minimal spanning tree (MST) technique to the Millennium voids, we identify the mini-filaments running through voids and measure their sizes at each redshift. Then, we calculate the comoving number density of void filaments as a function of their sizes in the logarithmic interval and determine an analytic fitting function for it. It is found that the size distribution of void mini-filaments in the logarithmic interval, $dN/d\\log S$, has an almost universal shape, insensitive to the redshift: In the short-size section it is well approximated as a power-law, $dN/d\\log S \\approx S$, while in the long-size section it decreases exponentially as $dN/d\\log S \\approx \\exp(-S^{\\alpha})$. We expect that the universal size distribution of void filaments may provide a useful cosmological probe without resorting to the rms density fluctuations. ",
        "introduction": "\\label{intro}  In the classical theory of structure formation, it was generally believed that the gravity is fully responsible for the formation and evolution of the large scale structure in the universe. Recently, however, it has been realized that the overall characteristics of the large scale structures cannot be understood only in terms of the gravitational influence. It has been pointed out that the tidal field provides a driving force in establishing the observed large scale filamentary structures in the universe \\citep{bon-etal96,lee-evr07, par-lee07b,hah-etal07,ara-etal07}. The cosmic voids are most vulnerable to the tidal influence from the surrounding matter distribution due to their extreme low-density \\citep{sah-sha96,sha-etal04,sha-etal06,lee-par06,par-lee07b}. The tidal squeezing and distortion effect tends to deviate the void shapes from spherical symmetry and could sometimes lead even to the collapse and disappearance of the voids \\citep{sah-etal94,sah-sha96,sha-etal04}.  The presence of the filamentary structures in voids marks the most striking evidence for the strong tidal influence on voids. The anisotropy in the spatial distribution of void halos is induced by their alignments with the principal axes of the tidal tensors \\citep{pee01}. Since the void halos evolve very little and most particles remain primordial in voids \\citep{ein06}, the void mini-filaments are pristine, keeping well the initial memory of the spatial coherence of the primordial tidal effect, unlike the large scale filaments which also originated from the large-scale coherence of the primordial tidal field but have undergone highly nonlinear processes in the subsequent evolution.  Hence, the void filaments are the most idealistic probe of the primordial tidal field and its effect on voids.  The size distribution of void filaments should be a function of the strength of the tidal effect and the host void sizes. The maximum size of a void filament cannot exceed the size of its host void.  But, even when a host void has a large size, it would not have long-size filaments if there is no tidal influence. Here, we attempt to derive the size distributions of the void filaments at various redshifts and to explore their statistical properties. Before presenting our results, however, we would like to caution the readers for the difficulty in using the void filaments as a probe of the primordial tidal field. Unlike the galaxy clusters, there is no unique way to define voids and their filaments. Recently, the seminal work of \\citet{col-etal08} compared thoroughly various void-finding algorithms and showed clearly that different void finders different void characteristics, under-density profiles, void galaxies etc. Thus, our results are likely to be dependent on our specific choice of the void-finder. The organization of this paper is as follows. In \\S \\ref{void}, we briefly describe the void catalog was obtained from the Millennium Run simulations. In \\S \\ref{filament}, we explain how the void filaments are identified from the void catalog. In \\S \\ref{result}, we derive numerically the size distributions of void filaments and determine an analytic fitting formula for it. In \\S \\ref{end}, we discuss the final results and assess the future work. ",
        "conclusions": "\\label{end}  It is generally thought that the large scale structure of the universe provides a window on the initial condition of the early universe. In previous cosmological studies,  it was the halo mass function that has been most highlighted as a statistical tool for probing cosmology with the large scale structure \\citep{pre-sch74,she-tor99,she-etal01,jen-etal01,ree-etal03}. The most advantageous aspect of the halo mass function is that it can be written in a simple universal form, independent of redshift \\citep{she-tor99}. Yet, the halo mass function has a generic weakness as a cosmological probe: its dependence on the cosmological parameters comes indirectly from the dependence of the halo mass on the rms fluctuation, $\\sigma(R)$, of the linear density field smoothed on a scale radius of $R$.  It is desirable to develop another statistical tool which can overcome the weakness of the halo mass function. Recently, several authors have noted the possibility of probing cosmology with cosmic voids \\citep[e.g.,][]{she-van04,par-lee07a,pla-etal08}. Yet, no new statistical tool that can really compete with the halo mass function has been suggested so far. For instance, in our previous work \\citep{par-lee07a}, we have proposed that the cosmic voids provide another window on cosmology. Given that the void ellipticities are induced by the primordial tidal effect which depends on the background cosmology, we have shown analytically that the void ellipticity distribution can be used to constrain the cosmological parameters.  However, it has turned out that the void ellipticity distribution suffers from the same weakness that the halo mass function has. Its dependence on cosmology is secondary, resulting from the dependence of the void scale on the linear rms density fluctuation, $\\sigma(R)$. Furthermore, the void ellipticity distribution cannot be written in a universal form unlike the halo mass function. Therefore, true as it is that the void ellipticity distribution provides another way to determine the cosmological parameters with the large scale structure, it is not competitively compared with the halo mass function.  In this work, we have for the first time quantified the effect of the tidal field on cosmic voids by the sizes of the void filaments. It is shown that the size distribution of void filaments is almost independent of redshift, having a simple universal form. A crucial implication of our result is that the number density of long filaments in voids should depend very sensitively on the background cosmology since the sizes of the void filaments reflect the primordial tidal field.  Our result that the size distribution of void filaments is almost independent of redshift in spite of the well known fact that the sizes of voids grow very strongly with redshift can be understood as follows. It is true that the comoving sizes of voids keep growing as the voids expand faster than the rest of universe due to their extreme low density. However, the comoving sizes of void filaments should not necessarily keep growing with redshift.  But for the effect of the tidal field, the spatial distribution of void halos would be isotropic due to the gravitational rarefaction effect caused by the fast expansion of voids. In other words, while the gravitational rarefaction effect tends to make the spatial distribution of void halos isotropic, the tidal effect on voids tends to increase the anisotropy in the distribution of void halos, resulting in growth of void filaments. Therefore, as far as the tidal effect exists and counteracts the expansion effect, the sizes of void filaments will also grow as the sizes of their host voids grow with redshift.  However, when the void size grows large enough to finally reach its threshold at which the expansion effect overcomes the tidal effect, the degree of anisotropy in the spatial distribution of void halos will diminish and accordingly the sizes of void filaments will stop growing. Since this critical threshold of void size is determined by the single condition that the expansion effect overcomes the tidal effect, its value should be redshift-independent.  At high-redshift when the sizes of voids were small, the strong dominant tidal effect increases the anisotropy in the spatial distribution of void halos, increasing the sizes of void filaments. At low-redshift when the sizes of voids are large, the strong expansion effect prevents the sizes of void filaments from keep growing. Hence, the universal size distribution of void filaments reflects the fact that the sizes of void filaments are regulated by the counter-balance between the expansion effect and the tidal effect.  The fitting formula, eqs. (\\ref{eqn:fit1})-(\\ref{eqn:fit2}), indicate that the number density of void filaments, $dN/d\\log\\nu$, behaves like a power-law with power index of $n=1$ in the short filament section while it decreases exponentially in the long filament section. A crucial implication of our result is that the exponential decrease of the abundance of long-size filaments in voids should be very sensitive to the key cosmological parameters. Especially it is expected to depend sensitively on the amplitude of the linear power spectrum, $\\sigma_{8}$. In our previous work, we have already shown that the void ellipticity caused by the tidal effect increases with the value of $\\sigma_{8}$ \\citep{par-lee07a}. Accordingly, we expect that the number density of long filaments in voids would increase as the value of $\\sigma_{8}$ increases. Unlike the void ellipticity distribution whose dependence on $\\sigma_{8}$ is weak and indirect through its dependence on the rms density fluctuation, equation (\\ref{eqn:fit1}) suggests that the size distribution of void filaments should depend strongly on the power spectrum amplitude.  To use the size distribution of void filaments as a cosmological probe, however, there should be some additional tasks to be done in the future. First, the mass-to-light bias and the redshift distortion effects have to be account for. What one can observe is not halos in real space but galaxies in redshift space. The size distribution of void filaments measured from the galaxy catalogs in redshift space could differ from the current result. Thus, it will be very necessary to investigate how the bias and the redshift distortion effect change the size distribution of void filaments. It is expected that the sizes of void filaments may increase in redshift space since the redshift space distortion effect tend to increase the anisotropy in the spatial distribution of void galaxies. Meanwhile, the matter-to-light bias might decrease the sizes of void filaments, compensating for the redshift distortion effect, since the massive halos are found to be distributed less anisotropically (Park \\& Lee 2009 in preparation). Henceforth, the overall size distribution of void galaxy filaments in redshift space might be similar to that in real space, due to the competition between the two effects (in private communication with van de Weygaert).  Second, a general analytic formula for the size distribution of void filaments in an arbitrary cosmology is desirable to derive from physical principles. If an analytic model is found, it would allow us to understand the true physical meaning of eqs. (\\ref{eqn:fit1})-(\\ref{eqn:fit2}). Third, it should be worth examining whether the results depend on the choice of the filament-finding and the void-finding algorithms. As mentioned in \\S 1, it has been found that different void-finders provide different void properties and the HV02 algorithm that we used here has a tendency to find larger voids than most of the other void-finders \\citep{col-etal08}. Furthermore, our results are not general ones but valid only for the specific cosmology of the Millennium Run simulation. Therefore, it is required to test the universality of the size distribution of void filaments with other void-finders and with other cosmologies as well. We plan to address these issues and report the results elsewhere in the near future."
    },
    "0801/0801.4420_arXiv.txt": {
        "abstract": "In cyclic cosmology based on phantom dark energy the requirement that our universe satisfy a CBE-condition ({\\it Comes Back Empty}) imposes a lower bound on the number $N_{\\rm cp}$ of causal patches which separate just prior to turnaround. This bound depends on the dark energy equation of state $w = p/\\rho = -1 - \\phi$ with $\\phi > 0$. More accurate measurement of $\\phi$ will constrain $N_{\\rm cp}$. The critical density $\\rho_c$ in the model has a lower bound $\\rho_c \\ge (10^9 {\\rm GeV})^4$ or $\\rho_c \\ge (10^{18} {\\rm GeV})^4$ when the smallest bound state has size $10^{-15}$m, or $10^{-35}$m, respectively. ",
        "introduction": "\\bigskip  Recently two of the authors have proposed~\\cite{BF,F,BF2} a scenario for a cyclic universe based on a dark energy component with constant equation of state satisfying $w = p/\\rho = -1-\\phi$ where $w < -1$ and hence $\\phi > 0$. The model involves two key ideas: (i) that the universe deflate just prior to the turnaround from expansion to contraction by disintegrating into a very large number $N_{\\rm cp}$ of causal patches. In the notation of \\cite{BF}, note that $N_{\\rm cp}=1/f^3$; (ii) that the contracting universe be empty, meaning that one causal patch at turnaround must contain no matter or black holes, only dark energy. This is called the CBE condition ({\\it Comes Back Empty}). Implementation of CBE requires, as we shall explain, a lower bound on $N_{cp}$ which depends on the length scale $L$ characterising the smallest bound system. We shall consider both $L = 10^{-15}$m for a nucleon then $L \\ge 10^{-35}$m for a PPP ({\\it Presently Point Particle}) meaning a particle which at present is considered to be pointlike, like a quark or a lepton, but which actually has a characteristic size at least a few orders of magnitude smaller than a nucleon but greater than or equal to the Planck scale.  \\bigskip  In the foreseeable future, it is expected that the equation of state of the dark energy $w$, and hence $\\phi$, will be measured with higher accuracy by, for example, the Planck Surveyor satellite~\\cite{Planck}. What we shall show is that this measurement can, within this model, constrain for a given $w$ the number $N_{\\rm cp}$ of causal patches at turnaround by imposing a lower bound thereon.  \\bigskip  The plan of the paper is that in Section 2 we discuss the times at which unbinding, causal disconnection and turnaround occur. In Section 3, the constraints on $N_{\\rm cp}$ from measurement of $\\phi$ are derived. Finally, Section 4 is a discussion. In the Appendices is technical material to supplement the main text.       \\newpage ",
        "conclusions": "\\bigskip  What we have deduced is that the parameters $\\rho_c, w$ and $N_{\\rm cp}$ in cyclic cosmology are already constrained by existing data. For example one requires $w \\gtrsim -2$ for the CBE aspect to work.  \\bigskip  This constraint is already known to be respected in Nature but as better and more accurate cosmological data become available it will shed further light on the viability of the theory.  \\bigskip  In particular, the accurate measurement of the equation of state $w = -1 -\\phi$ is of special interest. Fortunately the Planck Surveyor~\\cite{Planck} is anticipated to acquire improved accuracy on $w$ in the near future. As we have discussed, this will provide a lower bound on the number $N_{\\rm cp}$ of causal patches necessary to dissociate the smallest bound systems at turnaround and hence to solve the entropy problem and, via CBE, enable the possibility of infinite cyclicity.  \\bigskip  It is amusing that the physical conditions at the approach of deflation are so extraordinary that it is natural to ask whether the systems presently regarded as point particles may be composite because the phantom dark energy density grows to unimaginably large values and can disintegrate bound systems down to arbitrarily small scales. We have conservatively limited our attention to systems bigger than the Planck length. However, although this requirement seems dictated by considerations of quantum gravity, it is possible that the dark energy will dissociate even smaller systems if they exist.  \\bigskip  The advantage of cyclic cosmology is that it removes the initial singularity associated with the Big Bang, about 13.7 billion years ago, and allows that time never began. The previous attempts to create a consistent infinite cyclicity were stymied between about 1934~\\cite{Tolman} and 2002~\\cite{ST} primarily because of the entropy problem and the second law of thermodynamics. The discovery of the accelerated expansion rate of the universe and the concomitant necessity of dark energy has permitted more optimism that the cyclic cosmology is, after all, on the right track.  \\newpage  \\begin{center}"
    },
    "0801/0801.3036_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{intro}  As we move towards the era of ELTs, it is timely to think about the future role of the 8-m class telescopes.  Under the OPTICON programme novel technologies have been developed that are intended for use in multi-object and integral-field spectrographs.  To date, these have been targeted at instrument concepts for the European ELT, but there are also significant possibilities for their inclusion in new VLT instruments, ensuring the continued success and productivity of these unique telescopes. ",
        "conclusions": "In summary, there are considerable technical challenges in developing a wide-field spectrograph for the VLT, particularly when one considers the available resources within ESO and of its partners; a wide-field prime focus instrument would also lead to operational problems for VLTI.  Future smart focal plane and photonics technologies may enable a lighter, more compact, prime focus (or foward-Cassegrain) instrument, but current technology readiness levels suggest that a compromise would be to exploit the maximum field available at the Nasmyth foci, e.g. Super-Giraffe \\cite{giraffe}.  Meanwhile, technology development for ELT instrument concepts such as EAGLE, and through the next Framework Programme, should be supported to help break the existing paradigm.  \\vspace{-0.1in}"
    },
    "0801/0801.1205_arXiv.txt": {
        "abstract": "{} {We revisit the vicinity of the microquasar Cygnus X-3 at radio wavelengths. We aim to improve our previous search for possible associated extended radio features/hot spots in the position angle of the Cygnus X-3 relativistic jets focusing on shorter angular scales than previously explored. } {Our work is mostly based on analyzing modern survey and archive radio data, mainly including observations carried out with the Very Large Array and the Ryle Telescopes. We also used deep near-infrared images that we obtained in 2005. } {We present new radio maps of the Cygnus X-3 field computed after combining multi-configuration Very Large Array archive data at 6 cm and different observing runs at 2 cm with the Ryle Telescope. These are probably among the deepest radio images of Cygnus X-3 reported to date at cm wavelengths. Both interferometers reveal an extended radio feature within a few arc-minutes of the microquasar position, thus making our detection more credible. Moreover, this extended emission is possibly non-thermal, although this point still needs confirmation. Its physical connection with the microquasar is tentatively considered under different physical scenarios. We also report on the serendipitous discovery of a likely Fanaroff-Riley type II radio galaxy only $3^{\\prime}$ away from Cygnus X-3. } {} ",
        "introduction": "Cygnus X-3 is one of the X-ray binaries considered as the prototype of the family of galactic sources with relativistic jets, i.e., microquasars. It was originally discovered in X-rays by \\cite{gia67}. This system made headlines more than three decades ago when its first strong radio outbursts were detected, as described in \\cite{g72} and subsequent papers. Such flaring events have been repeating since then, typically one or two times per year, with flux density increments of almost two orders of magnitude above the normal quiescent level of $\\sim0.1$ Jy at cm wavelengths (\\cite{w1994}). At present, Cygnus X-3 is considered to be a high-mass X-ray binary (HMXB) with a WN Wolf-Rayet (WR) companion star of WN8 type (see e.g. \\cite{ke1996} and \\cite{ko2002}). The nature of the compact object is not well constrained, and neither a neutron star nor a black hole accreting from the WR star can be ruled out (see e.g. \\cite{er1998} and \\cite{mi1998}). The observed X-ray (\\cite{p1972}) and infrared (\\cite{bk1973}) modulation of 4.8 h is believed to be connected with the system orbital period, which is rather short for an HMXB. Orbital-phase-resolved spectra of Cygnus X-3 in the near infrared are consistent with it (\\cite{han2000}). The distance has been estimated to be approximately 9 kpc (\\cite{pred00}), which agrees with a strong interstellar absorption ($A_V \\sim 20$ mag, \\cite{ke1996}) that renders the optical counterpart undetectable in the visual domain.  \\begin{table*} \\begin{center} \\caption[]{\\label{vlaobs} VLA archive observations used in this paper \\label{obslog}} \\begin{tabular}{cccccccc} \\hline \\hline Date  &  wavelength  & VLA  & Number of  & Bandwidth  & Number of    &  Project & Time on    \\\\ &  (cm)   & configuration. &  IFs  &  (MHz)      &  visibilities &  id. & source (hrs) \\\\ \\hline 1992 Jun 8 &   6 &  DnC & 2 & 50 & 432229  & UT002 &  5   \\\\ 1997 May 4 &   6 &  B & 2  & 50 & 718939 & AM551   &  5.7   \\\\ 2000 Aug 29 &  6 &  D & 2 & 50 & 29167 & AH669     &  0.33    \\\\ \\hline \\hline \\end{tabular} \\end{center} \\end{table*}   Galactic microquasars such as Cygnus X-3 are known to release a significant amount of energy in the form of collimated relativistic jets into their surrounding inter-stellar medium (ISM). For Cygnus X-3, the averaged energy injection rate of its relativistic jets is estimated to be at least $10^{37}$ erg s$^{-1}$ (\\cite{m2005}), i.e., enough to supply $\\sim 10^{50}$ erg (10\\% of a supernova explosion) during the expected lifetime of the WR star. The Cygnus X-3 jets have been repeatedly resolved as sub-parsec transient radio features, propagating at a significant fraction of the speed of light in the north-south direction, thanks to interferometric radio techniques from arc-second (\\cite{m2001}) to milli arc-second angular scales (e.g. \\cite{g1983,miller04,tu2007a}). However, up to now there is no robust evidence of interaction between the relativistic ejecta of the system and the surrounding ISM on larger, few pc scales. This is in contrast with other microquasars, such as Cygnus X-1 or Circinus X-1, where a clear signature of interaction between their relativistic jets and the ISM does exist with a ring-like or lobe morphology several pc wide (\\cite{ga2005,tu2007b}). Additional examples of relativistic jet/ISM interaction come from the also pc-scale lobes of the Galactic Centre Annihilator 1E1740.7$-$2942 (\\cite{mira92}) or GRS 1758$-$258 (\\cite{m2002}). The detection of such ring/lobe features is a useful tool to better constrain the system's energetics.   The search for large scales features associated with the Cygnus X-3 radio jets has been a concern of the authors during recent years. Two H\\ion{II}\\ regions located $\\sim40^{\\prime}$ from Cygnus X-3 along the jet position angle were first tentatively proposed as possible large scale lobes (\\cite{m2000}). However, no further evidence for a physical connection could be found beyond the mere geometric alignment. A closer search revealed later the existence of two possible hot spot candidates (HSCs) associated with Cygnus X-3, thus suggesting an analogy with Fanaroff-Riley type II (FR II) radio galaxies (\\cite{m2005}). The apparent hot spots were two faint radio sources with non-thermal spectra at angular distances of 7\\prp 07 and 4\\prp 36 from Cygnus X-3. The line joining them was also within one degre of the almost North-South position angle of the inner arc-second radio jets (\\cite{m2001}). Unfortunately, follow up radio and near infrared observations of both the HSCs and the Cygnus X-3 nearby environment did not confirm the proposed hot spot nature and indicated that they were most likely background or foreground objects (\\cite{m2006}). This fact left open again for Cygnus X-3 the issue of searching for signatures of energy deposition from its relativistic jets into the ISM.  In this context, the main purpose of this paper is to present new very deep radio images of Cygnus X-3 and its environment obtained after combining multi-epoch archive and survey data, together with our own observations. The resulting maps enable us to put very strong limits for any extended or compact radio features that could be associated to Cygnus X-3 on a scale of a few pc. We also report on several extended radio features in the field with apparent non-thermal spectra that were previously unknown. Their possible connection with the microquasar is discussed from a skeptical point of view. ",
        "conclusions": "\\begin{enumerate}  \\item We revisited the issue of large scale radio features associated to the microquasar Cygnus X-3. By combining different archive VLA observations, we have been able to create a very deep (rms noise 9.5 $\\mu$Jy beam$^{-1}$) radio map with sensitivity to arc-minute angular scales. Cygnus X-3 appears superposed onto a diffuse radio emission with apparent non-thermal index with an angular size of a few arc-minutes extending South and South-West from it. The reality of such feature is independently confirmed by Ryle Telescope observations and when inspecting in detail maps from the CGPS and other survey data.  \\item With all cautions in mind, we tentatively suggest the possibility that such an extended emission could be physically associated to Cygnus X-3. It could be either a distorted lobe, plume or partial shell-like structure resulting from the accumulated flaring history of the microquasar interacting with the ISM. If connected with Cygnus X-3, the feature linear size would be $\\sim5$ pc. This is of the same order as extended lobes or bubbles seen around other microquasars. Based on simple equipartition assumptions, the observed size is roughly in agreement with expectations from simple theory of radio lobe dynamics.  \\item A triple radio source with FR II morphology has been also discovered at few arc-minute from Cygnus X-3 in our deep VLA map. In addition to provide a curious `family picture' of two accreting sources in the same shot, the angular proximity of the FR II compact core could enable future proper motion studies of this microquasar using future generations of highly sensitive interferometers.  \\end{enumerate}"
    },
    "0801/0801.3941_arXiv.txt": {
        "abstract": "The dynamical mass of a star cluster can be derived from the virial theorem, using the measured half-mass radius and line-of-sight velocity dispersion of the cluster. However, this dynamical mass may be a significant overestimation of the cluster mass if the contribution of the binary orbital motion is not taken into account. In these proceedings we describe the mass overestimation as a function of cluster properties and binary population properties, and briefly touch the issue of selection effects. We find that for clusters with a {\\em measured} velocity dispersion of $\\sigmalos\\ga 10$~\\kms{} the presence of binaries does not affect the dynamical mass significantly. For clusters with  $\\sigmalos\\la 1$~\\kms{} (i.e., low-density clusters), the contribution of binaries to $\\sigmalos$ is significant, and may result in a major dynamical mass overestimation. The presence of binaries may introduce a downward shift of $\\Delta \\log (L_V/\\mdyn) = 0.05-0.4$ in the $\\log (L_V/\\mdyn)$ vs. age diagram. ",
        "introduction": "%  An estimate for the mass of star clusters can be obtained from the virial theorem, using the projected half-mass radius $\\halfmass$ and line-of-sight velocity dispersion $\\sigmalos$. This dynamical mass estimate, $\\mdyn$, is obtained using the equation derived by \\cite{spitzer1987}: \\begin{equation} \\label{equation:spitzer} \\mdyn = \\eta \\, \\frac{\\halfmass  \\sigmalos^2}{G} \\,, \\end{equation} where $\\eta \\approx 9.75$ is a dimensionless proportionality constant. The derivation of $\\mdyn$ using the expression above is valid under the following assumptions: (1) the cluster dynamics are described by the Plummer model, (2) all stars are single and of equal mass, (3) the cluster is in virial equilibrium, and (4) no selection effects are present. Dynamical mass estimates for numerous clusters have been obtained this way \\citep[e.g.,][]{mandushev1991,smith2001,maraston2004,bastian2006,larsen2007,moll2007}. When binary stars are present, however, Eq.~(\\ref{equation:spitzer}) results in an overestimation of the cluster mass.  Observations have shown that the majority of the field stars are part of a binary or multiple system \\citep{duquennoy1991,fischermarcy1992}. It is believed that most (if not all) binary stars are formed in binary systems, which is supported by both observations \\citep{mathieu1994,mason1998,kobulnicky2007,kouwenhoven_adonis,kouwenhoven_recovery} and theory \\citep{goodwinkroupa2005}. Stars in binary systems exhibit not only motion in the gravitational potential of the cluster ({\\em particle motion} $\\sigmapart$), but also {\\em orbital motion} $\\sigmaorb$ in the binary system. Eq.~(\\ref{equation:spitzer}) is applicable for $\\sigmapart$ (the centre-of-mass motion of the binaries), but results in an overestimation if $\\sigmalos$, the superposition of $\\sigmapart$ and $\\sigmaorb$, is measured. In this paper we therefore address the question: ``How do binaries affect the dynamical mass of a star cluster?'' ",
        "conclusions": "%  The total mass of a star cluster is often inferred from its line-of-sight velocity dispersion $\\sigmalos$ and half-mass radius $\\halfmass$, assuming virial equilibrium. The latter approach includes the assumption that no binaries are present. However, most stars are known to be in binary systems, and their orbital motions provide an additional contribution to the measured $\\sigmalos$. The latter value is now no longer representative for the (centre-of-mass) motion of the stars and binaries in the cluster potential. This effect may result in a significant dynamical mass overestimation. Depending on the magnitude of the dynamical mass overestimation, we distinguish between three types of clusters: particle-dominated ($\\sigmaorb \\ll \\sigmapart$), intermediate-case ($\\sigmaorb \\approx \\sigmapart$), and binary-dominated clusters ($\\sigmaorb \\gg \\sigmapart$).  The orbital velocity of binaries is independent of the cluster properties (size, mass, etc.). Whether or not binary motion affects $\\sigmalos$ is thus depends on the cluster properties. For clusters with a high stellar density (i.e, large $\\mcl$ or small $\\halfmass$), $\\mdyn$ is generally unaffected. In the latter case $\\sigmapart > 10$~\\kms{}. The dynamical mass overestimation increases strongly with (i) a higher binary fraction and (ii) a smaller average orbital size. The dependence on the mass ratio distribution and eccentricity distribution is small: $\\Delta(\\mdyn/\\mcl) \\la 5\\%$. We additionally show that observing a certain subset of the cluster introduces a selection effects, which may result in a further mass overestimation by up to 40\\%. Analysis of the brightest (generally massive) stars in the cluster may further overestimate the dynamical mass. The dynamical mass overestimation can be negligible for the most massive clusters, while it may overestimate the true mass by up to an order of magnitude for sparse OB associations. The full results of this study have been presented in \\cite{kouwenhoven_sigma}. A follow-up paper, which treats the selection effects properly, is in preparation."
    },
    "0801/0801.3988_arXiv.txt": {
        "abstract": "V5116 Sgr (Nova Sgr 2005 No. 2), discovered on 2005 July 4, was observed with XMM-Newton in March 2007, 20 months after the optical outburst. The X-ray spectrum shows that the nova had evolved to a pure supersoft X-ray source, with no significant emission at energies above 1~keV. The X-ray light-curve shows abrupt decreases and increases of the flux by a factor $\\sim8$. It is consistent with a periodicity of 2.97~h, the orbital period suggested by \\cite{dob07}, although the observation lasted just a little more than a whole period. We estimate the distance to V5116~Sgr to be $11\\pm3\\,kpc$. A simple blackbody model does not fit correctly the EPIC spectra, with $\\chi^2_{\\nu}>4$. In contrast, ONe rich white dwarf atmosphere models provide a good fit, with $N_{\\rm H}= 1.3(\\pm0.1)\\times 10^{21}\\,cm^{-2}$, $T=6.1(\\pm0.1)\\times10^5\\,K$, and $L=3.9(\\pm0.8)\\times10^{37}(D/10~\\rm{kpc})^2{\\rm erg}\\,{\\rm s}^{-1}$ (during the high-flux periods). This is consistent with residual hydrogen burning in the white dwarf envelope. The white dwarf atmosphere temperature is the same both in the low and the high flux periods, ruling out an intrinsic variation of the X-ray source as the origin of the flux changes. We speculate that the X-ray light-curve may result from a partial coverage by an asymmetric accretion disk in a high inclination system. ",
        "introduction": "Classical novae occur in close binary systems of the cataclysmic variable type, when a thermonuclear runaway results in explosive hydrogen burning on the accreting white dwarf (WD). It is theoretically predicted that novae return to hydrostatic equilibrium after the ejection of a fraction of the accreted envelope. Soft X-ray emission arises in some post-outburst novae as a consequence of residual hydrogen burning on top of the WD. As the envelope mass is depleted, the photospheric radius decreases at constant bolometric luminosity (close to the Eddington value) with an increasing effective temperature. This leads to a hardening of the spectrum from optical through UV, extreme UV and finally soft X-rays, with the post-outburst nova emitting as a Supersoft Source (SSS) with a hot WD atmosphere spectrum. The duration of this SSS is related to the nuclear burning timescale of the remaining H-rich envelope and depends, among other factors, on the WD mass \\citep{TT98,SH05a,SH05b}.  All post-outburst novae are expected to undergo SSS emitting phase, provided a H-rich envelope is left. Nevertheless, the number and duration of SSS states observed in Galactic novae is small. In a systematic search for the X-ray emission from novae in the ROSAT archive, \\cite{Ori01a} found only three SSS novae from a total of 39 novae observed less than ten years after the outburst (GQ Mus, V1974~Cyg, N~LMC~1995). After the ROSAT era, observations by Beppo-SAX and Chandra revealed SSS X-ray emission for some more post-outburst novae: V382~Vel \\citep{Ori02,Bur02}, V1494~Aql, with a puzzling light curve showing a short burst and oscillations \\citep{Dra03}, and V4743~Sgr, also with a very variable light-curve \\citep{Nes03,Pet05}. The SSS phase of the recurrent nova RS~Oph in the 2006 outburst was monitored with Swift, XMM-Newton and Chandra, and the end of the SSS state occurred less than 100 days after outburst \\citep{Bode06,Hachi07}. XMM-Newton has contributed with some more observations of novae. Nova~LMC 1995 showed the SSS emission 5 years after outburst \\citep{Ori03}. Other post-outburst novae were detected with a harder spectrum associated to the expanding ejecta, as Nova LMC 2000 \\citep{Gre03} and V4633 Sgr \\citep{HS07}, or reestablished accretion flow \\cite[V2487~Oph,][]{HS02,fer07}. More recently, the SSS XMMSL1~J070542.7-381442 was identified as a nova \\citep{rea07a,rea07b,tor07}.  X-ray observations of the central area of M31, with its moderate absorption, offer a good chance to monitor the SSS phase of novae: \\cite{pie05} reported 21 X-ray counterparts for novae in M31 -- mostly identified as SSS by their hardness ratios -- and two in M33. Following that work, XMM-Newton and Chandra monitoring of M31 between July 2004 and February 2005 provided the detection of eleven out of 34 novae within a year after optical outburst \\citep{pie07}. This suggests that the low fraction of SSS found by \\cite{Ori01a} was due to selection effects and/or too poor sampling.  X-ray observations of post-outburst novae provide crucial information: soft X-rays yield a unique insight into the remaining burning envelope on top of the WD, while hard X-rays reveal shocks in the nova ejecta. The properties of  ``quiescent novae'', once they have turned-off and once accretion is reestablished (which can occur before or after the turn-off), are revealed both in hard and soft X-rays; they emit then as ``standard'' cataclysmic variables. In view of the scarcity of objects observed and the diversity of behaviors detected, only the monitoring of as many novae as possible, with large sensitivity and spectral resolution (as those offered by XMM-Newton and Chandra) can help to understand post-outburst novae.  V5116 Sgr (Nova Sgr 2005 No. 2) was one of the targets included in our X-ray monitoring programme of post-outburst Galactic novae with XMM-Newton. It was discovered by \\cite{lil05} on 2005 July 4.049~UT, with magnitude $\\sim$8.0, rising to mag 7.2 on July 5.085. The expansion velocity derived from a sharp P-Cyg profile detected in a spectrum taken on July 5.099 was $\\sim$1300~km/s. IR spectroscopy on July 15 showed emission lines with FWHM~$\\sim$2200~km/s \\citep{rus05}. Photometric observations obtained during 13 nights in the period August-October 2006 show the optical light curve modulated with a period of $2.9712\\pm0.0024$~h \\citep{dob07}, which the authors interpret as the orbital period. They propose that the light-curve indicates a high inclination system with an irradiation effect on the secondary star. A first X-ray observation with Swift/XRT (0.3--10~keV) in August 2005 yielded a marginal detection with 1.2($\\pm$1.0)$\\times10^{-3}$~cts/s \\citep{nes07a}. Two years later, on 2007 August 7, Swift/XRT showed the nova as a SSS, with 0.56$\\pm$0.1~cts/s \\citep{nes07b}. A first fit with a blackbody indicated $T\\sim4.5\\times10^5K$. More recently, a 35~ks Chandra spectrum obtained on 2007 August 28 was fit with a WD atmospheric model with N$_{\\rm H}=4.3\\times10^{21}$~cm$^2$ and $T=4.65\\times10^5K$ \\citep{NelOri07}.   Here we report on the X-ray light-curve and broad-band spectra of V5116 Sgr as observed by XMM-Newton in March 2007, 609 days after outburst. ",
        "conclusions": "V5116 Sgr was detected as a bright SSS by XMM-Newton 20 months after the nova outburst. The temperature of the WD atmosphere indicates that residual hydrogen burning is occurring in the WD envelope. To further confirm the origin of the emission and determine the size of the emitting surface, we need an estimation of the distance. Fortunately enough, the time and magnitude at maximum are quite well determined for V5116~Sgr, with a pre-maximum observation only 24~hours before the maximum. \\cite{lil05} reported $m_V=7.15$ at maximum, and $t_{2}$ (time required to decline two magnitudes from maximum) was 6.5$\\pm$1.0~days \\citep{dob07}. Using the empirical relation for novae between $M_V$ and $t_2$ \\citep{del95} the observed $t_2$ implies an absolute magnitude at maximum $M_V=-8.8(\\pm0.4)$ (we add a 5\\% error to account for the scatter in the $M_V-t_2$ relation). The photometry at maximum indicates $B-V=+0.48$ \\citep{gil05}, and for novae at maximum, intrinsic $B-V=0.23\\pm0.06$ \\citep{ber87}. This implies $A_{V}=3.1\\,E_{B-V}=0.8\\pm0.2$. With all this, we estimate the distance to V5116~Sgr to be $d=11\\pm3\\,kpc$. For simplicity, we scale the distance dependent parameters to 10~kpc. It is worth noticing that the $N_{\\rm H}$ obtained from our fits indicates $A_{V}=0.7$ \\cite[$N_{\\rm H}=5.9\\times10^{21}E_{B-V}\\,\\mbox{cm}^{-2}$,][]{zom07}, consistent with the value obtained from the observed colors.  The luminosity during the high flux periods, together with the atmosphere temperature, indicate a radius of the emitting object $R=6.2(\\pm0.9)\\times10^8(D/10kpc)\\,cm$. This corresponds to the whole WD surface emitting during the high flux periods and supports residual H-burning as the origin of the SSS emission, rather than any feature related to the accretion stream (boundary layer, hot spot, or hot accretion poles).  The most remarkable feature in V5116~Sgr is the SSS light-curve. Although the observation lasted just a bit longer than an orbital period, we obtain a period compatible with the orbital period of 2.97~hours found by \\cite{dob07}. The fact that the temperature is the same for both the high and the low flux phases indicates that the change in flux cannot be caused by an intrinsic change in the H-burning envelope. The luminosity during the high flux periods also indicates that the whole WD is visible in this state. The decrease by a factor $\\sim8$ in the flux could be caused by a partial eclipse of the WD. The long duration of the low flux phase and short duration of the high flux phase cannot be reconciled with a partial eclipse by the secondary star. Alternatively, some asymmetric accretion disk could do the job, being responsible for a partial eclipse during most of the orbital period, with a 3~ks window during which the whole WD would be visible. Since no change is observed in the hydrogen absorption column and it is compatible with the average interstellar value both in the low and the high flux spectra, the portion of accretion disk producing the partial eclipse should be optically thick to the soft X-rays. In the particular case of V5116 Sgr, we have checked that an absorbing column of $N_{H}>2\\times10^{22}\\,cm^{-2}$ would make it completely opaque to the observed SSS emission. Longer X-ray observations covering at least two periods together with a good timing with the OM camera would be required to disentangle the origin of the X-ray light curve behaviour."
    },
    "0801/0801.1496_arXiv.txt": {
        "abstract": "We report high-cadence time-series photometry of the recently-discovered transiting exoplanet system HD~17156, spanning the time of transit on UT 2007 October 1, from three separate observatories.  We present a joint analysis of our photometry, previously published radial velocity measurements, and times of transit center for 3 additional events. Adopting the spectroscopically-determined values and uncertainties for the stellar mass and radius, we estimate a planet radius of $R_{p} = 1.01 \\pm 0.09 \\, R_{\\rm Jup}$ and an inclination of $i = 86.5^{+1.1}_{-0.7}$~degrees.  We find a time of transit center of $T_{c} = 2454374.8338 \\pm 0.0020$~HJD and an orbital period of $P = 21.21691 \\pm 0.00071$~days, and note that the 4 transits reported to date show no sign of timing variations that would indicate the presence of a third body in the system.  Our results do not preclude the existence of a secondary eclipse, but imply there is only a 9.2\\% chance for this to be present, and an even lower probability (6.9\\%) that the secondary eclipse would be a non-grazing event.  Due to its eccentric orbit and long period, HD~17156b is a fascinating object for the study of the dynamics of exoplanet atmospheres.  To aid such future studies, we present theoretical light curves for the variable infrared emission from the visible hemisphere of the planet throughout its orbit. ",
        "introduction": "\\label{intro_sect}  Twenty-eight transiting planets are now reported in the literature\\footnote{See {\\tt http://www.inscience.ch/transits/}}, and the doubling time scale for new detections is now roughly one year.  As reviewed by \\citet{char2007}, it is these objects that have allowed the study of the physical structure and atmospheric chemistry and dynamics of gas giant exoplanets, and opened the field of comparative exoplanetology. However, the majority have orbital periods of a few days, due both to the lower geometric probability of transits occurring as the orbital period increases, and due to the limitations of ground-based photometric transit surveys that strongly favor detection of short-period systems.  The latter problem can be circumvented by searching for transits of known radial velocity planet-bearing stars, putting to use the radial velocity information to constrain the possible time of transit.  This is the primary aim of the Transitsearch.org network \\citep{sea2003}\\footnote{See {\\tt http://www.transitsearch.org}}, which, starting in 2002, has been conducting photometric follow-up observations of known radial velocity detected planets.  This network is a prime example of successful collaboration between amateur and professional astronomers.  Although longer period planets possess a lower geometric probability to transit, there exists a loop-hole for some planets on highly eccentric orbits.  The transit probability for a planet on an eccentric orbit, with periastron near inferior conjunction, is amplified by a factor \\begin{equation} A = \\left[ {{1+e\\cos(\\frac{\\pi}{2} - \\omega)} \\over{1-e^2}}\\right] \\end{equation} where $e$ is the eccentricity, and $\\omega$ is the argument of pericenter \\citep{sea2003}.  An amplification factor near 2.9 for the planet orbiting HD~17156 \\citep{fischer2007}, along with an extremely short time window for possible transits, brought the system to the attention of the Transitsearch.org network.  This was rewarded by the detection of transits as recently announced by \\citet{bar2007}.  This unique system contains a $3.1$-$M_{\\rm Jup}$ transiting planet in a $21.2$ day highly-eccentric orbit ($e = 0.67$) about a bright ($V = 8.2$) G0V star.  \\citet{gil2007} have recently presented refined estimates of the system parameters based on new photometric data. Even among planets with periods beyond 5 days, where the timescale for tidal circularization quickly grows to exceed the age of most systems, HD~17156b's eccentricity is unusually high, thus making it an interesting test case for models of planetary migration.  Preliminary measurements of the Rossiter-McLaughlin effect \\citep{rossiter24,mclaughlin24,gw2007,winnasp2007} by \\cite{nar2007} show evidence for a misalignment between the stellar spin axis and the planetary orbit axis, which may indicate a migration mechanism involving planet-planet scattering \\citep[see e.g.][]{chat2007}, or Kozai migration under perturbation by a yet undetected stellar companion or second planet \\citep{fab2007,wu2007}.  The large eccentricity and long period also make HD~17156 a particularly attractive target for several additional follow-up studies.  First, it presents a unique opportunity for the study of the structure and dynamics of a gas-giant atmosphere under strongly varying illumination, through infrared monitoring of the planetary emission (e.g. \\citealt{har2006}; \\citealt{cow2007}; \\citealt{knut2007a,knut2007b}).  Second, by monitoring successive times of transit, the presence of additional bodies in the system can be detected or constrained from the resulting perturbations on the orbit of HD~17156b (\\citealt{hm2005}; \\citealt{agol2005}; \\citealt{sa2005}).  The purpose of our paper is three-fold: First, we seek to refine the estimates of the system parameters that were only poorly constrained by the light curves of \\citet{bar2007}.  Second, we document the time of center of transit and search for transit timing variations.  Third, we evaluate the likelihood that a secondary eclipse will be observable for HD~17156, and present theoretical predictions of the infrared phase variations as might be detected with the {\\it Spitzer Space Telescope}.  We elected to perform an independent analysis from that of \\citet{gil2007}, and therefore do not use their orbital parameters as constraints, since at the time of writing these are still subject to change as their paper has not yet been accepted for publication. We do however make use of the time of mid-transit presented in their Table 1.  We begin by presenting a description of our observations of the transit event on UT 2007 October 1 in \\S \\ref{obs_sect}.  We then present our combined analysis of the extant radial velocity measurements (\\S \\ref{rv_sect}) and photometric data (\\S \\ref{phot_sect}), and our resulting estimates of the system parameters and likelihood of a secondary eclipse.  In \\S \\ref{disc_sect}, we compare our constraints on the planet mass and radius with planetary structural models, summarize the current constraints on the presence of transit timing variations in the system, and present model predictions for the planetary infrared light curve.  We conclude in \\S \\ref{summ_sect} with a discussion of compelling avenues for future research. ",
        "conclusions": "\\label{summ_sect}  HD~17156 thus represents a highly interesting system for follow-up studies in the mid infra-red (e.g. using the {\\it Spitzer Space Telescope}), even if it does not undergo secondary eclipses. Ground-based observations should also be pursued to continue the search for transit timing variations, and hence to constrain the presence of additional bodies in the system, which, given the high orbital eccentricity of the planet HD~17156b, may show interesting dynamical properties and place constraints on planetary formation and evolution scenarios.  This discovery represents an important success for the Transitsearch.org project, and highlights the potential rewards of collaboration between distributed networks of amateur astronomers and the professional community."
    },
    "0801/0801.4085_arXiv.txt": {
        "abstract": "We use several main-sequence models to derive distances (and extinctions), with statistically meaningful uncertainties for 11 star-forming-regions and young clusters. The model dependency is shown to be small, allowing us to adopt the distances derived using one model. Using these distances we have revised the age order for some of the clusters of \\cite{2007MNRAS.375.1220M}. The new nominal ages are: $\\approx2$ Myrs for NGC6530 and the ONC, $\\approx3$ Myrs for $\\lambda$ Orionis, NGC2264 and $\\sigma$ Orionis, $\\approx4-5$ Myrs for NGC2362, $\\approx13$ Myrs for h and $\\chi$ Per, $\\approx20$ Myrs for NGC1960 and $\\approx40$ Myrs for NGC2547. In cases of significantly variable extinction we have derived individual extinctions using a revised Q-method \\citep{1953ApJ...117..313J}.  These new data show that the largest remaining uncertainty in deriving an age ordering (and necessarily ages) is metallicity. We also discuss the use of a feature we term the R-C gap overlap to provide a diagnostic of \\textbf{isochronal} age spreads or varying accretion histories within a given star-formation-region. Finally, recent derivations of the distance to the ONC lie in two groups. Our new more precise distance of $391^{+12}_{-9}$ pc allows us to decisively reject the further distance, we adopt 400 pc as a convenient value. ",
        "introduction": "\\label{intro}  Colour-magnitude diagrams (CMDs) of star-formation regions (SFRs) provide, in combination with model isochrones, an excellent tool with which to determine distances, ages and individual stellar masses. These parameters are critical for determining initial mass functions (IMFs) for stellar populations and discovering the possible impacts of local environment (such as the effect of ionising winds from massive stars) on disc lifetimes and on star and planet formation and evolution. Many calculations of IMFs, disc fractions etc are available but they are derived in heterogeneous ways, thus hints of the effects of environment are only recently beginning to emerge \\citep[e.g.][]{2007MNRAS.375.1220M,2004AJ....128..765S}.  Very precise photometry ($\\approx1\\%$) is routinely available, along with sophisticated stellar models. However, current parameter derivations from CMDs still have relatively large uncertainties and are model dependent \\citep[see the discussions in][]{2004A&A...415..571B, 2004ApJ...600..946P,2002MNRAS.335..291N,2007MNRAS.375.1220M}. Thus current age (and distance) uncertainties all but `wash-out' any environmental effects. Clearly more robust constraints would be available for current stellar theories if more precise parameters could be extracted from the CMDs of SFRs.  In \\cite{2007MNRAS.375.1220M} we created an age ladder for a range of pre-MS populations. The first stage was to create empirical isochrones by fitting splines to the pre-MS locus. Overlaying them in absolute magnitude and intrinsic colour results in an age ladder, with the youngest SFRs at the brightest absolute magnitudes. SFRs with almost indistinguishable positions in the CMD were grouped, and nominal ages assigned to each group. Thus the age sequence (though not the nominal ages) are free from the problems associated with pre-main-sequence (pre-MS) models. In \\cite{2007MNRAS.375.1220M} we had to adopt literature distances for the studied SFRs.  These distances were derived using a range of different methods and their uncertainties proved to be the largest remaining contributor to the uncertainties in our age ladder placements.  Previous distances have chiefly been derived using main-sequence (MS) isochrone fitting, pre-MS isochrone fitting or from \\textit{HIPPARCOS} parallax measurements. MS isochrone fitting provides distances based on the positions of MS stars in a CMD, which are independent of uncertainties in age. Pre-MS isochrone fitting also uses the positions of stars in a CMD, but in this method the derived distances are degenerate with age \\citep[see e.g.][]{2006MNRAS.373.1251N}. Finally, distances derived from \\textit{HIPPARCOS} parallax measurements are only available for a few SFRs included in this paper, $\\sigma$ Ori, NGC2547 and $\\lambda$ Ori, with all except $\\lambda$ Ori having large uncertainties.  Of those methods used to derive distances, the most suitable for the derivation of an age ladder is clearly MS isochrone fitting.  Fully convective pre-MS stars (in young SFRs) are separated in a CMD from those stars on the MS which have radiative cores. The transistion region or gap in the CMD (measurable in colour), we term the radiative-convective gap \\citep[R-C gap, see][for introduction of term and discussion]{2007MNRAS.375.1220M}. Once stars have crossed the R-C gap their position in a CMD is almost independent of age until they reach the turn off. However MS fitting has not yielded the precision one would expect in distance estimates.  This is due to two significant problems. Firstly the position of MS isochrones, although temporally static, is model dependent, with different studies adopting different models. Secondly distances are most often derived using `by eye' fitting of models to the data, yielding ill-defined uncertainties. A full discussion of previous fitting methods can be found in \\cite{2006MNRAS.373.1251N}.  In this paper we solve both of these problems. Firstly we show that the model dependency is small for the model isochrones studied.  We then adopt the distances from a single MS model. This allows us to derive a set of precise distances,  accurate relative to each other, which we term \"relative distances\". Second we use the $\\tau^2$ fitting technique \\citep{2006MNRAS.373.1251N}, a new rigorous and self-consistent method of fitting stars to isochrones, which yields statistically meaningful uncertainties. This presents us with the opportunity to achieve more precise distances from the fitting of high-mass (HM) or MS stars.  The rest of this paper is laid out as follows. In Section \\ref{data} we detail the literature sources, the nature of the data used and any sequence selection carried out on the stars.  Section \\ref{calibration} details the different model isochrones and photometric calibrations used. Section \\ref{fit_method} describes the fitting process. This is done primarily by way of an example in Section \\ref{chiper_eg}.  Section \\ref{ind_ext} describes the derivation of individual extinctions, in particular Section \\ref{Q} describes an revised Q-method for calculating approximate individual extinctions. The results for all the isochrone calibrations and methods are presented in Section \\ref{results_app}. Section \\ref{results} outlines our results for one adopted model with our best-fitting distances given in Table \\ref{dist_comp}. In section \\ref{implications} we discuss the implications of the individual distances to several key SFRs (Section \\ref{imp_ind}) and those of the entire dataset. The implications of the dataset on metallicity (Section \\ref{imp_met}), age spreads and the R-C gap overlap (Section \\ref{imp_cuts}) and secular evolution within the SFRs with particular reference to disc fractions (Section \\ref{imp_sec}) are discussed. The reader interested in distance and reddening values should skip to Section \\ref{results_app} for the values derived from all the models used. ",
        "conclusions": "\\label{imp_sum}  \\begin{enumerate} \\item We have derived a self-consistent set of distances of generally higher precision than previously available for a set of SFRs in the age range 1-40 Myrs. We have also derived distances using several other models and  calibrations (see Section \\ref{results_app}). \\item In addition to these new distances and reddenings (or extinctions) we have reconstructed the age ladder of \\cite{2007MNRAS.375.1220M} assigning new nominal ages, as shown in Table \\ref{ladder}. To enable the reader to add other SFRs to this ladder the pre-MS splines are freely available from the cluster collaboration home page\\footnote{http://www.astro.ex.ac.uk/people/timn/Catalogues/description.html}, and the CDS archive. \\item We have shown that metallicity information is now vital for accurate relative distances to SFRs. This is especially true if one is attempting to characterise evolutionary indicators such as disc fractions as a function of age, or trying to uncover environmental effects (such as the effect of ionising winds from massive stars on planet formation from discs). \\item We have discussed that the overlapping region of the R-C gap could be used to derive spreads in isochronal age in a CMD,  i.e. the apparent age difference between the maximum mass star still on the convective pre-MS and the minimum mass star which has reached the MS. If star formation is slow and isochronal ages of individual stars are reliable this would provide a direct measurement of the age spreads present in SFRs. If star formation is rapid the R-C gap overlap region reveals the underlying spread in accretion histories within an SFR. This is important as for rapid star formation, if an accretion history is unknown isochronal ages derived from a position in a CMD do not represent the true age of a star. Indeed it is therefore likely that if a rapid star formation model is accurate median or mean ages drawn from a population are also invalid. A more useful approach may be to compare SFRs using age ladder arguments or perhaps to use the age of the youngest stars which have the lowest accretion history. \\item We have shown further evidence for non-uniform decay of discs in SFRs, although new comparisons must be made using consistent disc fraction indicators and mass ranges. \\end{enumerate}"
    },
    "0801/0801.3346_arXiv.txt": {
        "abstract": "{% We combine extinction maps from the Two Micron All Sky Survey (2MASS) with Hipparcos and Tycho parallaxes to obtain reliable and high-precision estimates of the distance to the Ophiuchus and Lupus dark complexes.  Our analysis, based on a rigorous maximum-likelihood approach, shows that the $\\rho$-Ophiuchi cloud is located at $(119 \\pm 6) \\mbox{ pc}$ and the Lupus complex is located at $(155 \\pm 8) \\mbox{ pc}$; in addition, we are able to put constraints on the thickness of the clouds and on their orientation on the sky (both these effects are not included in the error estimate quoted above).  For Ophiuchus, we find some evidence that the streamers are closer to us than the core.  The method applied in this paper is currently limited to nearby molecular clouds, but it will find many natural applications in the GAIA-era, when it will be possible to pin down the distance and three-dimensional structure of virtually every molecular cloud in the Galaxy. ",
        "introduction": "\\label{sec:introduction}  In recent years, much effort has been dedicated to the study of dark molecular clouds and their dense cores.  One of the main motivations for these investigations is the study of the process of star and planet formation in its entirety, and a deeper understanding of the effects of the local environment.  A key aspect of the scientific analysis of a dark molecular cloud is its distance, which is related to many physically relevant properties (in particular, the mass scales as the square of the distance).  Unfortunately, the distance estimates for many clouds are often plagued by very large uncertainties and it is not rare to see in the literature measurements that differ by large factors.  Often, this quantity is inferred by associating the cloud to other astronomical objects whose distance is well known.  For example, the Lupus complex is located near the Sco OB2 association \\citep{1978ApJS...38..309H, 1999AJ....117..354D}, whose distance is estimated to be $\\sim 150 \\mbox{ pc}$ based on the photometry of OB stars \\citep{1991ESOSR..11....1R}.  However, this method is based on some degree of arbitrariness when making the link between the cloud and the other objects, and thus the deduced distance can be completely unreliable.  The cloud complexes considered in this paper, the $\\rho$ Ophiuchi and the Lupus dark clouds, are good examples of how different distance estimates made by different authors can be.  Quoted values for the $\\rho$-Ophiuchi cloud are in the range $120$--$165 \\mbox{ pc}$, as recently summarized by \\citet{2004AJ....127.1029R}. \\citet{1981A&A....99..346C} estimated the distance to be $\\sim 160 \\mbox{ pc}$ from multi-color photometry of heavily absorbed stars in the $\\rho$ Oph core, while \\citet{1998A&A...338..897K} suggested a significantly smaller figure, $\\sim 120 \\mbox{ pc}$.  For Lupus, the situation is even more extreme, with estimates from $100 \\mbox{ pc}$ \\citep{1998A&A...338..897K} to $190 \\mbox{ pc}$ \\citep{1998MNRAS.301L..39W}, both based on Hipparcos data; the most widely accepted distance is in the middle of these two extremes, at approximately $140 \\mbox{ pc}$ \\citep{1993AJ....105..571H, 1999AJ....117..354D}.  These uncertainties severely hamper our understanding of the physical properties of molecular clouds, and thus our knowledge of star formation.  For example, an error by a factor two on the distance of a cloud translates into an error by a factor four in the mass, and has a thus a huge impact on the estimate of the density, size, stability, and star formation efficiency of their cloud cores (see \\citealp{2007A&A...462L..17A} for further discussions on this point).  In this paper we present accurate distance measurements of the Ophiuchus and Lupus complexes based upon Hipparcos and Tycho parallaxes.  Our technique is statistically sound and, when applied to nearby giant molecular clouds, is able to provide accurate distance measurements and related errors.  The paper is organized as follows. In Sect.~\\ref{sec:basic-analysis} we present the main datasets used and discuss a simple approach to the problem.  A more quantitative technique is developed in Sect.~\\ref{sec:likelihood-analysis}, where we also presents the results obtained for the two clouds considered in this paper.  The technique is further developed in Sect.~\\ref{sec:second-order-geom} to include the effects of the cloud thickness and orientation.  Finally, in Sect.~\\ref{sec:discussion} we briefly discuss the results obtained. ",
        "conclusions": "\\label{sec:discussion}  Our analysis indicates that the Ophiuchus cloud is very likely to be closer than the we think: the ``standard'' distance, $160 \\mbox{ pc}$, is indeed excluded at $5 \\sigma$.  Not surprisingly, our estimate is in excellent agreement with the one obtained by \\citet{1998A&A...338..897K} using a similar technique (but note that the two results are based on a different datasets and use different statistical methods).  Recently, \\citet{WilsonPhD} and \\citet{2007arXiv0709.0505M} reported an Hipparcos estimated distance of Ophiuchus of $d = 136^{+8}_{-7} \\mbox{ pc}$ and $d = 135^{+8}_{-7} \\mbox{ pc}$, respectively, results which are marginally in agreement with our distance estimate of the Ophiuchus core.  \\citeauthor{2007arXiv0709.0505M} mainly focused on the Lynds~168 cloud, and inferred the cloud distance from the parallaxes of seven stars presumably associated with illuminated nebulae.  Although his method is plagued by a potential uncertainty, the real association between the Hipparcos stars and the Ophiuchus cloud complex, and suffers from the small number statistics (7 parallaxes), the agreement obtained is reassuring and suggests that we finally have in hand a reliable estimate of the Ophiuchus cloud complex.  \\citet{WilsonPhD}, instead, used a more rigorous maximum-likelihood approach, but his analysis is substantially different from ours.  In particular, his method is based on a preliminary classification of stars as foreground and background based on their observed reddening; the distance of the cloud is then inferred from the constraints imposed by the parallaxes of the stars considered.  In contrast, in our method we do not need to explicitly classify stars, but on the contrary can use directly all data in the likelihood (an advantage of this is that we are able to use in our method also stars that cannot be uniquely classified as foreground or background because, e.g., of their large photometric errors).  Very recently, \\citet{Loinard} used phase-referenced multi-epoch Very Long Baseline Array (VLBA) observations to measure the trigonometric parallax of the Ophiucus star-forming region.  Their method is based on accurate (to better than a tenth of a milli-arcsecond), absolute measurements of the position of individual magnetically active young stars though to be associated with the star-forming region.  If multi-epoch data are available, one can infer the parallax of the star, and thus of the star-forming regions, with exquisite accuracy (below $1 \\mbox{ pc}$ for objects within $\\sim 200 \\mbox{ pc}$).  This technique was applied to four stars in the $\\rho$-Ophiuchi region, and interestingly it lead to somewhat contradictory results: two objects associated with the sub-condensation Oph~A (S1 and DoAr21) have a measured parallax close to $\\sim 120 \\mbox{ pc}$, and thus in excellent agreement with our estimate; two others, associated with the condensation Oph~B (VSSG14 and VL5), have a much higher distance of $\\sim 165 \\mbox{ pc}$.  A possible explanation of this is provided by a closer examination of the two sources in the Oph~B condensation. From their spectral energy distributions (SEDs), they appear to be Class~III A-type stars, with little or no circumstellar emitting dust, and an extinction as large as $50 \\mbox{ mag}$ in the $V$ band (see \\citealp{1989ApJ...340..823W}, and especially \\citealp{1994ApJ...434..614G} for a discussion on the individual SEDs of VSSG14 and VL5).  These results indicate that VSSG14 and VL5 are probably background stars, not directly associated with the Ophiucus complex, and likely to be part of a smaller OB association at $\\sim 165 \\mbox{ pc}$.  The distance of the Lupus complex is substantially in agreement with the ones reported in the literature, but both the thickness analysis and the orientation analysis suggest noticeable differences in the various Lupus subclouds."
    },
    "0801/0801.2389_arXiv.txt": {
        "abstract": " ",
        "introduction": "About one third of all known asteroids belong to families (\\cite{zappala}), which are clusters of asteroids believed to result from collisional disruption of parent bodies (\\cite{brien}).  The hypothesis of collisional origin is based entirely on the observed similarity of dynamical and spectral properties, and not on the understanding of the collisional physics itself (\\cite{michel2001}). Members of asteroid families share similar orbital elements (semimajor axis,  eccentricity and orbital inclination) (\\cite{nesvorny+}) and chemical composition (\\cite{juric}). They often have low density { which suggests a high porosity -- the asteroids are often made of} gravitationally bound fragments or ``rubble piles'' (\\cite{richardson}). { This structure is consistent with a collisional origin, and} explains  both the lack of fast rotators among asteroids larger than a few hundred meters (\\cite{pravec}) and the presence of large craters revealed by close-up surface images (\\cite{chapman}). The rotational rates of the asteroids display a Maxwellian distribution, indicating angular momentum redistribution by collisions (\\cite{binzel,fulchignoni}). As such, prominent asteroid families are therefore important tracers of high-velocity impacts, one of the principal geologic processes affecting small bodies in the Solar System.  Fig 1. comes here   Collisions modify the shape and size of an asteroid. { The most energetic collisions disrupt the bodies and generate a number of smaller, gravitationally unbound fragments which will evolve further as new individual asteroids. The less energetic collisions can only fragment the body and/or form craters, without disruption. Collisions are rare, and therefore cannot be observed directly, and their complete understanding requires numerical  experiments that have to reproduce the observable constraints of the impacts. Currently these constraints  are: asteroid size distribution, the number of asteroid families, the age distribution of meteorites, and surface cratering of 5 asteroids which were encountered by spacecraft. The last of these cannot be easily interpreted, because craters can overlap or be completely saturated, and crater size depends on the impactor/crater scaling law model. The other constraints are only sensitive to the most energetic, disruptive collisions. Significant information on the less energetic impacts may be extracted from asteroid shapes: it has been proposed  (\\cite{richardson,K+A}) that the shape of asteroids can evolve with age due to several microimpacts. However, there are only a few asteroids with known shape and hence the evolutionary effects could not have not been fully exploited yet.}  While sizes can be constrained relatively easily from absolute photometry (\\cite{eddington}), determining shapes is much more difficult. Space missions have visited only a few asteroids and shapes for these have been determined accurately (\\cite{fujiwara,abe,okada,saito}). Radar observations have revealed shapes of  Near-Earth Objects encountering the Earth ( e.g. \\cite{ostro}), while high-resolution imaging is restricted to the largest asteroids (\\cite{tanga,marchis}). For the rest of the asteroid belt, time-resolved photometry is the only method to determine the shape. The key effect is the rotation: it causes periodic brightness variations which, in principle, reveal the shape of a given asteroid if there are enough data for full inversion of the lightcurve, or at least for fitting a triaxial ellipsoid.  Despite the relatively simple principles of the method, there has been slow progress because of the time-consuming nature of the observations.  Recent advances in all-sky surveys offer new approaches to the problem. These surveys often observe thousands of asteroids in a night,  and with the appropriate data processing the shape distribution can be  constrained. Here we deduce the shape distribution of 11,735 main-belt asteroids belonging to eight prominent asteroid families, using data from the Sloan Digital Sky Survey (SDSS).  SDSS was a deep imaging survey covering over a quarter of the Celestial Sphere in the Northen Galactic Cap, resulting in high-quality photometric measurements of 50 million stars and a similar number of galaxies. Although its main purpose was cosmology, SDSS has also collected  204,305 measurements of asteroids, listed in the SDSS Moving Object Catalog 3 (SDSS MOC) \\footnote{Available at http://www.sdss.org} (\\cite{ivezic}).  Of these, 67,637 data points have been linked to 26,847 known asteroids (\\cite{juric}). This means that all the identified asteroids have been observed an average of 2.5 times. The calibration of the SDSS is very accurate because all measurements were acquired with the same instrument and from the same observing site. Therefore, systematic effects are minimized and the individual brightness measurements are accurate to 1--2\\%{} down to magnitude 20 in $r$ band. ",
        "conclusions": ""
    },
    "0801/0801.3420_arXiv.txt": {
        "abstract": "The present standard model of cosmology states that the known particles carry only a tiny fraction of total mass and energy of the Universe. Rather, unknown dark matter and dark energy are the dominant contributions to the cosmic energy budget. We review the logic that leads to the postulated dark energy and present an alternative point of view, in which the puzzle may be solved by properly taking into account the influence of cosmic structures on global observables. We illustrate the effect of averaging on the measurement of the Hubble constant. ",
        "introduction": "The concordance model of cosmology states that only $4\\%$ of the mass/energy content of the Universe is carried by identified particles, and the dominant contributions are unknown dark matter ($22\\%$) and mysterious dark energy ($74\\%$) \\cite{Spergel}. It is said that we are living in the age of ``precision cosmology'', which is justified as a statement on the precision of observations, but certainly not on their understandings. The purpose of this talk is to recapitulate the genesis of the dark energy hypothesis and to show that it poses a serious crisis to theoretical cosmology. We argue that a possible resolution for this crisis may come from the proper treatment of effects from averaging. ",
        "conclusions": ""
    },
    "0801/0801.1339_arXiv.txt": {
        "abstract": "A short discussion of theoretical results on cosmic ray first-order Fermi acceleration at relativistic shock waves is presented. We point out that the recent results by Niemiec with collaborators change the knowledge about these processes in a substantial way. In particular one can not expect such shocks to form particle distributions extending to very high energies. Instead, distributions with the shock compressed injected component followed by a more or less extended high energy tail are usually created. Increasing the shock Lorentz factor leads to steepening and decreasing of the energetic tail. Also, even if a given section of the spectrum preserves the power-law form, the fitted spectral index may be larger or smaller than the claimed `universal index' $\\sigma \\approx 2.2$~.  A reported simple check of real shapes of electron spectra in the Cyg A hot spots provides results clearly deviating from the standard expectations for such shocks met in the literature. The spectrum consist of the very flat low energy part ($\\sigma \\approx 1.5$), up to electron energies $\\sim 1$ GeV, and much steeper part ($\\sigma > 3$) at higher energies. We conclude with presentation of a short qualitative discussion of the Fermi second-order processes acting in relativistic plasmas. We suggest that such processes can be the main accelerating agent for very high energy particles. In particular its can accelerate electrons to energies in the range of $1$ - $10^3$ TeV in  relativistic jets, shocks and radio-source lobes. ",
        "introduction": "Relativistic plasma flows are observed in a number of astrophysical objects, ranging from a mildly relativistic jets of  the sources like SS433, through the-Lorentz-factor-of-a-few jets in AGNs and galactic  `micro-quasars', up to highly ultra-relativistic outflows in sources of gamma  ray bursts or pulsar winds. As nearly all such objects are efficient emitters of nonthermal radiation, what requires existence of energetic particles, our attempts to  understand the processes generating cosmic rays are essential for understanding the fascinating phenomena observed. Below I will discuss the work carried out in order to understand the cosmic ray acceleration processes acting at relativistic shocks and within highly turbulent regions accompanying such shocks and shear layers. I will not include here the interesting work involving collisionless shocks modelling with {\\it particle in cell} simulations. This approach uses quite different modelling method in comparison to the other work discussed here,  relating in most cases to the characteristic energy range of shock compressed thermal plasma particles.  The present paper is a modified and updated version of some of my previous reviews of the subject. Also, essentially the same slightly shortened text is provided as my contribution to the HEPRO Workshop in Dublin (September 2007).  Below we will append the index `1' (`2') to quantities measured in the plasma rest frame upstream (downstream) of the shock. ",
        "conclusions": ""
    },
    "0801/0801.4200_arXiv.txt": {
        "abstract": "Fluorescence light is induced by extensive air showers while developing in the Earth's atmosphere. The number of emitted fluorescence photons depends on the conditions of the air and on the energy deposited by the shower particles at every stage of the development. In a previous model calculation, the pressure and temperature dependences of the fluorescence yield have been studied on the basis of kinetic gas theory, assuming temperature-independent molecular collision cross-sections. In this work we investigate the importance of temperature-dependent collision cross-sections and of water vapour quenching on the expected fluorescence yield. The calculations will be applied to simulated air showers while using actual atmospheric profiles to estimate the influence on the reconstructed energy of extensive air showers. ",
        "introduction": "\\label{intro}  Several air shower experiments like HiRes~\\cite{hires}, the Pierre Auger Observatory~\\cite{auger}, and Telescope Array~\\cite{ta}, are using the fluorescence technique for detecting extensive air showers (EAS) induced by ultra-high energy cosmic rays. Measuring the fluorescence light that nitrogen molecules emit after being excited by charged particles of EAS is currently the most direct method for determining the energy of EAS in a model-independent way. A thorough understanding of the light emission process is necessary to obtain the primary energy of EAS with high precision.  In this paper, we extend our previous model calculation for the fluorescence light emission~\\cite{BK2006} by including the latest results on input parameters and their temperature dependence as obtained in laboratory measurements. For the reconstruction of air shower events, the light emission has to be known in dependence on altitude in the Earth's atmosphere at which the shower is observed. Up to now, the altitude dependence has been considered by including air density profiles and collisional quenching of nitrogen-nitrogen and nitrogen-oxygen molecules as described by kinetic gas theory. The cross-sections for collisional quenching were taken to be temperature independent. However, the cross-sections describing collisional quenching are known to be temperature-dependent~\\cite{gs}. Gr\\\"un and Schopper~\\cite{gs} found a decreasing collisional quenching cross-section with increasing temperature. Recently, the AirFly experiment has studied collisional quenching cross-sections in dependence on temperature~\\cite{airfly_icrc}. These data are also included in the model calculations presented in this article. In addition we investigate the influence of water vapour on the fluorescence yield, using relative humidity measurements performed at the site of the Auger detector in Argentina. ",
        "conclusions": "\\label{concl}  The effects of temperature-dependent collisional quenching cross-sections and of quenching due to water vapour have been studied. Both effects lead to a significant reduction of the fluorescence yield in the lower part of the atmosphere. Applying these calculations to simulated EAS, a distortion of the longitudinal shower development is found. A reduction of the emitted ligth is expected, which varies from about 7\\% to 11\\% depending on seasonal atmospheric model and on zenith angle of the EAS. Hence, accounting for these effects in the reconstruction of the primary energy of EAS, the primary energy will be increased by this amount as compared with the former model calculations. The position of the shower maximum is hardly shifted, in all atmospheric models the shift is less than 50~m.  \\subsection*{Acknowledgement} On of the authors (BK) is supported by the German Research Foundation (DFG) under contract KE 1151/1-2."
    },
    "0801/0801.4493_arXiv.txt": {
        "abstract": "APEX single-dish observations at sub-millimeter wavelengths toward a sample of massive star-forming regions reveal that C$_2$H is almost omni-present toward all covered evolutionary stages from Infrared Dark Clouds via High-Mass Protostellar Objects to Ultracompact H{\\sc ii} regions. High-resolution data from the Submillimeter Array toward one hot-core like High-Mass Protostellar Object show a shell-like distribution of C$_2$H with a radius of $\\sim$9000\\,AU around the central submm peak position.  These observed features are well reproduced by a 1D cloud model with power-law density and temperature distributions and a gas-grain chemical network. The reactive C$_2$H radical (ethynyl) is abundant from the onset of massive star formation, but later it is rapidly transformed to other molecules in the core center. In the outer cloud regions the abundance of C$_2$H remains high due to constant replenishment of elemental carbon from CO being dissociated by the interstellar UV photons. We suggest that C$_2$H may be a molecule well suited to study the initial conditions of massive star formation. ",
        "introduction": "Spectral line surveys have revealed that high-mass star-forming regions are rich reservoirs of molecules from simple diatomic species to complex and larger molecules (e.g., \\citealt{schilke1997b,hatchell1998b,comito2005,bisschop2007}). However, there have been rarely studies undertaken to investigate the chemical evolution during massive star formation from the earliest evolutionary stages, i.e., from High-Mass Starless Cores (HMSCs) and High-Mass Cores with embedded low- to intermediate-mass protostars destined to become massive stars, via High-Mass Protostellar Objects (HMPOs) to the final stars that are able to produce Ultracompact H{\\sc ii} regions (UCH{\\sc ii}s, see \\citealt{beuther2006b} for a recent description of the evolutionary sequence). The first two evolutionary stages are found within so-called Infrared Dark Clouds (IRDCs).  While for low-mass stars the chemical evolution from early molecular freeze-out to more evolved protostellar cores is well studied (e.g., \\citealt{bergin1997,dutrey1997,pavlyuchenkov2006,joergensen2007}), it is far from clear whether similar evolutionary patterns are present during massive star formation.  To better understand the chemical evolution of high-mass star-forming regions we initiated a program to investigate the chemical properties from IRDCs to UCH{\\sc ii}s from an observational and theoretical perspective. We start with single-dish line surveys toward a large sample obtaining their basic characteristics, and then perform detailed studies of selected sources using interferometers on smaller scales. These observations are accompanied by theoretical modeling of the chemical processes.  Long-term goals are the chemical characterization of the evolutionary sequence in massive star formation, the development of chemical clocks, and the identification of molecules as astrophysical tools to study the physical processes during different evolutionary stages. Here, we present an initial study of the reactive radical ethynyl (C$_2$H) combining single-dish and interferometer observations with chemical modeling.  Although C$_2$H was previously observed in low-mass cores and Photon Dominated Regions (e.g., \\citealt{millar1984,jansen1995}), so far it was not systematically investigated in the framework of high-mass star formation. ",
        "conclusions": "To understand the observations, we conducted a simple chemical modeling of massive star-forming regions. A 1D cloud model with a mass of 1200\\,M$_\\sun$, an outer radius of 0.36\\,pc and a power-law density profile ($\\rho\\propto r^p$ with $p=-1.5$) is the initially assumed configuration. Three cases are studied: (1) a cold isothermal cloud with $T=10$\\,K, (2) $T=50$\\,K, and (3) a warm model with a temperature profile $T\\propto r^q$ with $q=-0.4$ and a temperature at the outer radius of 44\\,K. The cloud is illuminated by the interstellar UV radiation field (IRSF, \\citealt{draine1978}) and by cosmic ray particles (CRP). The ISRF attenuation by single-sized $0.1\\mu$m silicate grains at a given radius is calculated in a plane-parallel geometry following \\citet{vandishoeck1988}. The CRP ionization rate is assumed to be $1.3\\times 10^{-17}$~s$^{-1}$ \\citep{spitzer1968}. The gas-grain chemical model by \\citet{vasyunin2008} with the desorption energies and surface reactions from \\citet{garrod2006} is used. Gas-phase reaction rates are taken from RATE\\,06 \\citep{woodall2007}, initial abundances, were adopted from the ``low metal'' set of \\citet{lee1998}.  Figure \\ref{model} presents the C$_2$H abundances for the three models at two different time steps: (a) 100\\,yr, and (b) in a more evolved stage after $5\\times10^4$\\,yr. The C$_2$H abundance is high toward the core center right from the beginning of the evolution, similar to previous models (e.g., \\citealt{millar1985,herbst1986,turner1999}). During the evolution, the C$_2$H abundance stays approximately constant at the outer core edges, whereas it decreases by more than three orders of magnitude in the center, except for the cold $T=10$~K model.  The C$_2$H abundance profiles for all three models show similar behavior.  The chemical evolution of ethynyl is determined by relative removal rates of carbon and oxygen atoms or ions into molecules like CO, OH, H$_2$O. Light ionized hydrocarbons CH$^+_{\\rm n}$ (n=2..5) are quickly formed by radiative association of C$^+$ with H$_2$ and hydrogen addition reactions: C$^+$ $\\rightarrow$ CH$_2^+$ $\\rightarrow$ CH$_3^+$ $\\rightarrow$ CH$_5^+$.  The protonated methane reacts with electrons, CO, C, OH, and more complex species at later stage and forms methane.  The CH$_4$ molecules undergo reactive collisions with C$^+$, producing C$_2$H$_2^+$ and C$_2$H$_3^+$. An alternative way to produce C$_2$H$_2^+$ is the dissociative recombination of CH$_5^+$ into CH$_3$ followed by reactions with C$^+$.  Finally, C$_2$H$_2^+$ and C$_2$H$_3^+$ dissociatively recombine into CH, C$_2$H, and C$_2$H$_2$. The major removal for C$_2$H is either the direct neutral-neutral reaction with O that forms CO, or the same reaction but with heavier carbon chain ions that are formed from C$_2$H by subsequent insertion of carbon. At later times, depletion and gas-phase reactions with more complex species may enter into this cycle.  At the cloud edge the interstellar UV radiation instantaneously dissociates CO despite its self-shielding, re-enriching the gas with elemental carbon.  The transformation of C$_2$H into CO and other species proceeds efficiently in dense regions, in particular in the ``warm'' model where endothermic reactions result in rich molecular complexity of the gas (see Fig.~\\ref{model}).  In contrast, in the ``cold'' 10\\,K model gas-grain interactions and surface reactions become important.  As a result, a large fraction of oxygen is locked in water ice that is hard to desorb ($E_{\\rm des} \\sim 5500$~K), while half of the elemental carbon goes to volatile methane ice ($E_{\\rm des} \\sim 1300$~K). Upon CRP heating of dust grains, this leads to much higher gas-phase abundance of C$_2$H in the cloud core for the cold model compared to the warm model. The effect is not that strong for less dense regions at larger radii from the center.  Since the C$_2$H emission is anti-correlated with the dust continuum emission in the case of IRAS\\,18089-1732 (Fig.\\,\\ref{18089}), we do not have the H$_2$ column densities to quantitatively compare the abundance profiles of IRAS\\,18089-1732 with our model. However, data and model allow a qualitative comparison of the spatial structures. Estimating an exact evolutionary time for IRAS\\,18089-1732 is hardly possible, but based on the strong molecular line emission, its high central gas temperatures and the observed outflow-disk system \\citep{beuther2004a,beuther2004b,beuther2005c}, an approximate age of $5\\times10^4$\\,yr appears reasonable.  Although dynamical and chemical times are not necessarily exactly the same, in high-mass star formation they should not differ to much: Following the models by \\citet{mckee2003} or \\citet{krumholz2006b}, the luminosity rises strongly right from the onset of collapse which can be considered as a starting point for the chemical evolution. At the same time disks and outflows evolve, which should hence have similar time-scales.  The diameter of the shell-like C$_2$H structure in IRAS\\,18089-1732 is $\\sim 5''$ (Fig.\\,\\ref{18089}), or $\\sim$9000\\,AU in radius at the given distance of 3.6\\,kpc.  This value is well matched by the modeled region with decreased C$_2$H abundance (Fig.\\,\\ref{model}).  Although in principle optical depths and/or excitation effects could mimic the C$_2$H morphology, we consider this as unlikely because the other observed molecules with many different transitions all peak toward the central submm continuum emission in IRAS\\,18089-1732 \\citep{beuther2005c}. Since C$_2$H is the only exception in that rich dataset, chemical effects appear the more plausible explanation.  The fact that we see C$_2$H at the earliest and the later evolutionary stages can be explained by the reactive nature of C$_2$H: it is produced quickly early on and gets replenished at the core edges by the UV photodissociation of CO. The inner ``chemical'' hole observed toward IRAS\\,18089-1732 can be explained by C$_2$H being consumed in the chemical network forming CO and more complex molecules like larger carbon-hydrogen complexes and/or depletion.  The data show that C$_2$H is not suited to investigate the central gas cores in more evolved sources, however, our analysis indicates that C$_2$H may be a suitable tracer of the earliest stages of (massive) star formation, like N$_2$H$^+$ or NH$_3$ (e.g., \\citealt{bergin2002,tafalla2004,beuther2005a,pillai2006}). While a spatial analysis of the line emission will give insights into the kinematics of the gas and also the evolutionary stage from chemical models, multiple C$_2$H lines will even allow a temperature characterization. With its lowest $J=1-0$ transitions around 87\\,GHz, C$_2$H has easily accessible spectral lines in several bands between the 3\\,mm and 850\\,$\\mu$m.  Furthermore, even the 349\\,GHz lines presented here have still relatively low upper level excitation energies ($E_u/k\\sim42$\\,K), hence allowing to study cold cores even at sub-millimeter wavelengths.  This prediction can further be proved via high spectral and spatial resolution observations of different C$_2$H lines toward young IRDCs."
    },
    "0801/0801.1080_arXiv.txt": {
        "abstract": "The line-of-sight velocities and [OIII] 5007 \\AA\\ expansion velocities are measured for 11 planetary nebulae (PNs) in the Virgo cluster core, at 15 Mpc distance, with the FLAMES spectrograph on the ESO VLT.  These PNs are located about halfway between the two giant ellipticals M87 and M86. From the [OIII] 5007 \\AA\\ line profile widths, the average half-width at half maximum expansion velocity for this sample of 11 PNs is $ \\bar{v}_{HWHM} = 16.5$ kms$^{-1}$ (RMS = 2.6 kms$^{-1}$). We use the PN subsample bound to M87 to remove the distance uncertainties, and the resulting [OIII] 5007 \\AA\\ luminosities to derive the central star masses. We find these masses to be at least $0.6$ M$_\\odot$ and obtain PN observable life times $t_{PN} < 2000 $ yrs, which imply that the bright PNs detected in the Virgo cluster core are compact, high density nebulae. We finally discuss several scenarios for explaining the high central star masses in these bright M87 halo PNs. ",
        "introduction": "Since the discovery of free-floating intracluster planetary nebulae (ICPNs) in the Virgo cluster \\citep{Arnaboldi+96}, extensive imaging and spectroscopic observations were carried out to determine their projected phase space distribution and the fraction of diffuse cluster light not bound to Virgo galaxies. To enlarge the sample of more than 40 ICPN line-of-sight (LOS) velocities available from \\citet{Arnaboldi+03,Arnaboldi+04}, we have obtained new PN spectra with FLAMES at the ESO VLT \\citep{Doherty+08}. The high spectral resolution of the new data allows us for the first time to measure expansion velocities for the [OIII] nebula shells of 11 PNs in the Virgo cluster core, nearly half way between M87 and M86.  The expansion velocity of a planetary nebula (PN) is one of the most important parameters determining its evolution, but currently it is known only for a few hundred Galactic PNs, mostly bright objects in the Milky Way Disk \\citep{gesickizijlstra00}, and for a few tens in the Magellanic Clouds and M31. When interpreted using dynamically evolving nebular models \\citep[e.g.][]{Schoenberner+05}, PN shell expansion velocities provide reliable extimates of PN dynamical ages and distance estimates for Galactic PNs.  In this letter, we give a brief summary of our Observations in Section~\\ref{obser}. In Section~\\ref{PNexp}, we present the PN [OIII] 5007 \\AA\\ half-width half maximum velocity measurements $v_{HWHM}$, and $m(5007)$ magnitudes as defined by \\citet{Jacoby89}. In Section~\\ref{M87PNLF} we build the PN luminosity function (PNLF) for the spectroscopically confirmed PN sample in Virgo and for the subsample bound to the M87 halo.  In Section~\\ref{outerradius}, we then estimate the outer radii of the M87 halo PNs from the observed PN expansion velocities and their visibility lifetimes.  We finally discuss in Section~\\ref{fdiscuss} the remaining distance uncertainties, the shape of the PNLF in the M87 halo, and the mechanisms that may be responsible for the large core masses of the brightest PNs in the halo of M87. ",
        "conclusions": "We have measured the nebular [OIII] 5007\\AA\\ spectroscopic expansion velocities for 11 PNs in the Virgo cluster core, at 15 Mpc distance, with the FLAMES spectrograph on the ESO VLT. Based on the [OIII] line profile width, the average spectroscopic expansion velocity for this sample is $ \\bar{v}_{HWHM} = 16.5$ kms$^{-1}$ (RMS = 2.6 kms$^{-1}$), which is in agreement with the predictions of dynamically evolving nebular models for high density nebulae close to their maximal $m(5007)$ emission. Large central star masses M$_{CS} > 0.6$~M$_\\odot$ are inferred from the bright measured [OIII] luminosities and the known distance for the PNs bound to the M87 halo. From the large central star masses and the measured $v_{HWHM}$, we derive short PN visibility times and small nebular outer radii, $ \\sim 0.07$ pc.  The PNs in the M87 halo have large central star masses, are compact and their nebula shells may be similar to those observed for the brightest Galactic Bulge PNs.  Three mechanisms are reviewed as possible explanation for the large core masses of the M87 halo PNs: intermediate age population, metallicity effects, and blue stragglers, with the latter being the most likely, given the old age and the low metallicities of the IRGB stars in the Virgo core \\citep{Williams+07}."
    },
    "0801/0801.1755_arXiv.txt": {
        "abstract": "We study the dynamical behaviour of the interacting holographic dark energy model whose interaction term is $Q=3H(\\lam_d\\rho_d + \\lam_c\\rho_c)$, where $\\rho_d$ and $\\rho_c$ are the energy density of dark energy and CDM respectively. To satisfy the observational constraints from SNIa, CMB shift parameter and BAO measurement, if $\\lam_c = \\lam_d$ or $\\lam_d, \\lam_c >0$, the cosmic evolution will only reach the attractor in the future and the ratio $\\rho_c/\\rho_d$ cannot be slowly varying at present. Since the cosmic attractor can be reached in the future even when the present values of the cosmological parameters do not satisfy the observational constraints, the coincidence problem is not really alleviated in this case. However, if $\\lam_c \\neq \\lam_d$ and they are allowed to be negative, the ratio $\\rho_c/\\rho_d$ can be slowly varying at present and the cosmic attractor can be reached near the present epoch. Hence, the alleviation of the coincidence problem is attainable in this case. The alleviation of coincidence problem in this case is still attainable when confronting this model to SDSS data.   \\vspace{3mm} \\begin{flushleft} \\textbf{Keywords}: Dark Energy. \\end{flushleft} ",
        "introduction": "Observations suggest that the expansion of the universe is accelerating \\cite{Riess:98,Perl:99}. The acceleration of the universe may be explained by supposing that the present universe is dominated by a mysterious form of energy whose pressure is negative, known as dark energy. One problem of the dark energy model is the coincidence problem, which is the problem why the dark energy density and matter density are of the same order of magnitude in the present epoch although they differently evolve during the expansion of the universe. A possible way to alleviate the coincidence problem is to suppose that there is an interaction between matter and dark energy. The cosmic coincidence can then be alleviated by appropriate choice of the form of the interaction between matter and dark energy leading to a nearly constant ratio $r=\\rho_c/\\rho_d$ during the present epoch \\cite{Zimdahl:03, Campo:06, Sadjadi:06} or giving rise to attractor of the cosmic evolution at late time \\cite{Amen:99, Chimento:03}. Since the existence of the cosmic attractor implies constant $r$ but the attractor does not always occur at the present epoch, we first find a range of dark energy parameters for which the attractor occurs and then check the evolution of $r$ during the present epoch.  Based on holographic ideas \\cite{Cohen:98, Li:04}, one can determine the dark energy density in terms of the horizon radius of the universe. This type of dark energy is holographic dark energy \\cite{Setare:07} - \\cite{Elizalde:05}. By choosing Hubble radius as the cosmological horizon, the present amount of dark energy density agrees with observations. Nevertheless, dark energy evolves like matter at present, so it cannot lead to an accelerated expansion. However, if the particle horizon is chosen to be the cosmological horizon, the equation of state parameter of dark energy can become negative but not negative enough to drive an accelerating universe. The situation is better when one uses the event horizon as the cosmological horizon. In this case, dark energy can drive the present accelerated expansion, and the coincidence problem can be resolved by assuming an appropriate number of e-foldings of inflation. Roughly speaking, the coincidence problem can be resolved because the size of the cosmological horizon during the present epoch depends on the amount of e-folds of inflation, and the amount of holographic dark energy depends on the horizon size. Nevertheless, the second law of thermodynamics will be violated if $w_d < -1$ \\cite{Li:04, Gong:06}. Hence, $w_d$ should not cross the boundary $w_d = -1$. The boundary $w_d = -1$ can be crossed if dark energy interacts with matter. Since now the horizon size has a dependence on the interaction terms,the alleviation of cosmic coincidence should also depend on the interaction term.  In this work we suppose that the holographic dark energy interact only with cold dark matter (CDM) and treat baryons as non-interacting matter component. Our objective is to compare the region of dark energy parameters for which the cosmic evolution has an attractor within the parameter region that satifies the observational constraints from combined analysis of SNIa data \\cite{Riess:06}, CMB shift parameter \\cite{Wang:06} and BAO measurement \\cite{Eisenstein:05}. The results of the comparison can tell us about the range of parameters that alleviate the cosmic coincidence. ",
        "conclusions": "For the interacting holographic dark energy model, we study the fixed points and their stabiliby, and compare a range of model parameters for which attractor exists with the $99.7\\%$ confidence levels from the combined analysis of SNIa data, CMB shift parameter and BAO measurement. Neglecting baryons, the observational constraints require that the value of $\\Om_c$ at the attractor point must be small if $\\lam_d = \\lam_c$ or $\\lam_d, \\lam_c >0$. This implies that the cosmic evolution will reach the attractor point in the future when $\\Om_c$ becomes small. In this case, $r$ cannot be slowly varying during the present epoch and the cosmic attractor cannot be reached near the present. Hence, the coincidence problem is not really alleviated for this case. However, if $\\lam_d$ and $\\lam_c$ are allowed to be negative, the cosmic evolution can reach the attractor near the present epoch for a narrow range of $\\lam_d$ and $\\lam_c$. Therefore, the coincidence problem is  possible  to  alleviate in this case. Including baryons in our consideration, the attractor of the cosmic evolution cannot occur at present due to the non-vanishing baryon fraction. According to observations, the fixed point in this case is possible only when $Q$ is small and positive. Hence, the fixed point will be slowly reached in the future. These results indicate that for the interacting holographic dark energy model with the interaction terms considered here, the cosmic coincidence problem cannot be alleviated very well. We also briefly considered the constraint from SDSS matter power spectrum on the dark energy parameters. We have found that the parameters ranges that lead to the alleviation of cosmic coincidence are allowed by SDSS data."
    },
    "0801/0801.0279_arXiv.txt": {
        "abstract": "Polarimetry is extensively used as a tool to trace the interstellar magnetic field projected on the plane of sky. Moreover, it is also possible to estimate the magnetic field intensity from polarimetric maps based on the Chandrasekhar-Fermi method. In this work, we present results for turbulent, isothermal, 3-D simulations of sub/supersonic and sub/super-Alfvenic cases. With the cubes, assuming perfect grain alignment, we created synthetic polarimetric maps for different orientations of the mean magnetic field with respect to the line of sight (LOS). We show that the dispersion of the polarization angle depends on the angle of the mean magnetic field regarding the LOS and on the Alfvenic Mach number. However, the second order structure function of the polarization angle follows the relation $SF \\propto l^{\\alpha}$, $\\alpha$ being dependent exclusively on the Alfvenic Mach number. The results show an anti-correlation between the polarization degree and the column density, with exponent $\\gamma \\sim -0.5$, in agreement with observations, which is explained by the increase in the dispersion of the polarization angle along the LOS within denser regions. However, this effect was observed exclusively on supersonic, but sub-Alfvenic, simulations. For the super-Alfvenic, and the subsonic model, the polarization degree showed to be intependent on the column density. Our major quantitative result is a generalized equation for the CF method, which allowed us to determine the magnetic field strength from the polarization maps with errors $< 20\\%$. We also account for the role of observational resolution on the CF method. ",
        "introduction": "It is believed that giant molecular clouds in the interstellar medium (ISM) are threaded by large scale magnetic fields \\citep{sch98, cru99}. However, it is still not completely clear what is the role of the magnetic field in the dynamics of the ISM and what is its effect on the star formation process. Also, the ratio of the magnetic and turbulent energy in these environments is a subject of controversy \\citep{padoan02, girart06}. Magnetic fields can influence the injection and evolution of turbulence bringing more complexity to this issue (see Lazarian \\& Cho [2004] for review). As an example, simulations have shown that strongly magnetized turbulent media develop structures with lower density contrasts when compared to pure hydrodynamic turbulence \\citep{kowal07, kritsuk07}.  Observationally, different techniques can be used to measure the ISM magnetic field and determine its intensity and topology. Zeeman splitting of spectral lines provides a direct and precise derivation of the magnetic field component along the line of sight (LOS), mainly for clouds presenting strong spectral lines \\citep{heiles05}. However, it cannot be applied to the clouds where the line intensities are too weak. Typically, when observed, Zeeman measurements of molecular clouds give $B_{\\rm LOS} \\sim 10^{1-3} \\mu$G, and suggest the correlation with density $B_{\\rm LOS} \\propto \\rho^{0.5}$, which is consistent to the expected relation for compressions of magnetic fields frozen into plasma. Spectral line broadening show that molecular clouds present supersonic, but critically Alfvenic motions \\citep{cru99b}. This fact shows that the turbulent motions may be excited by MHD modes instead of being purely hydrodynamical.  One of the most readily available methods of studying the perpendicular component of the magnetic field is based on the polarization of dust thermal emissions at infrared and submillimetric wavelengths \\citep{hil00}. The alignment of grains in respect to the magnetic field is a hot research topic (see Lazarian [2007] for review). Radiative torques (RATs) can promote alignment of irregular dust particles, resulting in different intensities for polarized radiation parallel and perpendicular to the local magnetic field \\citep{dolg76, draine96, draine97, lazho07a}. Grains with long axis aligned perpendicular to the magnetic field induce polarization parallel to the magnetic field for transmitted star light, and perpendicular to the field lines for the dust emission. \\cite{cho05} showed that RATs are very efficient on the grain alignment process in molecular clouds, even for the very dense regions (up to $A_V<10$). They also showed that the alignment efficiency strongly depends on the grain size, being practically perfect for large grains ($a>0.1$ $\\mu$m). More detailed studies of the RATs efficiency by \\cite{lazho07a} confirmed this claim. Therefore, for a range of $A_V$ it is acceptable to assume that the grains are well-aligned.  For a given polarization map of an observed region, the mean polarization angle indicates the orientation of the large scale magnetic field. On the other hand the polarization dispersion gives clues on the value of the turbulent energy. This, as a consequence, can be used to determine the magnetic field component along the plane of sky. \\cite{chandra53} introduced a method (CF method hereafter) for estimating the ISM magnetic fields based on the dispersions of the polarization angle and gas velocity. Simply, it is assumed that the magnetic field perturbations are Alfvenic and that the rms velocity is isotropic.  A promising approach to test this method is to create two-dimensional (plane of sky) synthetic maps from numerically simulated cubes. \\cite{ostriker01} performed 3D-MHD simulations, with $256^3$ resolution, in order to obtain polarization maps and study the validity of the CF method on the estimation of the magnetic field component along the plane of sky. They showed that the CF method gives reasonable results for highly magnetized media, in which the dispersion of the polarization angle is $< 25^{\\circ}$. However, they did not present any other statistical analysis or predictions that could be useful for the determination of the ISM magnetic field from observations.  \\cite{heitsch01} presented a complementary work, with a more detailed analysis regarding the limited observational resolution on the CF method, and presented a modified equation to account for the differences obtained previously. They concluded that lower observational resolution leads to an overestimation of the magnetic field from the CF equation. They also showed a good agreement between the CF technique and the expected magnetic field of their simulations, except for the weak field models.  Polarization maps from numerical simulations can also be used in the study of the correlation between the polarization degree and the total emission intensity (or dust column density). Observationally, the polarization degree in dense molecular clouds decreases with the total intensity as $P \\propto I^{-\\alpha}$, with $\\alpha = 0.5 - 1.2$ \\citep{goncalves05}. \\cite{padoan01} studied the role of turbulent cells in the $P \\ versus \\ I$ relation using supersonic and super-Alfvenic self-graviting MHD simulations. They found a decrease of polarization degree with total dust emission within gravitational cores, in agreement with observations, if grains are assumed to be unaligned for $A_V > 3$. When the alignment was assumed to be independent on $A_V$, the anti-correlation was not observed. Recently, \\cite{pelkonen07} extended this work and refined the calculation of polarization degree introducing the radiative transfer properly. In that work, the decrease in the alignment efficiency arises without any {\\it ad hoc} assumption. The alignment efficiency decreases as the radiative torques become less important in the denser regions. However, it is still not clear the role of the magnetic field topology and the presence of multiple cores intercepted by the line of sight on the decrease of polarization degree.  In this work we attempt to extend the previously cited studies improving and applying the CF method for different situations. For that, we studied both sub and super-Alfvenic models, to study the role of the magnetic field topology in the observed polarization maps. We simulate different observational resolutions in the calculations of polarization maps and provide combined statistical analysis for both dust absorption and emission maps. We also present statistics based methods to characterize the turbulence and magnetic properties from polarization maps. We performed numerical simulations of magnetized turbulent plasma with higher resolution, which are described in Sec.\\ 2. From the data, we computed ``observable\" polarization maps, as shown in Sec.\\ 3. We then present the statistics and spatial distributions of angle and polarization degree for different models in Sec.\\ 4. In Sec.\\ 5 we propose the generalized equation for the CF method and compare it with the expected values to study its validity. In Sec.\\ 6, we discuss the improved procedures of polarization vector statistics, which allow observers to characterize the mean and fluctuating magnetic field of the cloud. We also discuss the applicability of our approach for polarized molecular and atomic lines, and compare the results with previous works. Ou summary is provided in Sec.\\ 7. ",
        "conclusions": "Emission and extinction polarimetric measurements provide an unique technique for the study of the magnetic field, projected into the plane of sky, in molecular clouds. Synthetic extinction maps depend on additional assumptions about the stellar population and may be more explored in a future work. In this work we focused on providing synthetic emission polarimetric maps, as well as different statistical analysis that could be used in the future by observers to infer the physical properties of the studied region. The physical interpretations of our results, as well as the comparisons with previous theoretical works, are given as follows.  \\subsection{Our models}  In this work we presented four different models: (1) $\\beta = 1.0$, $M_{\\rm S}=0.7$ and $M_{\\rm A}=0.7$, (2) $\\beta = 0.1$, $M_{\\rm S}=2.0$ and $M_{\\rm A}=0.7$, (3) $\\beta = 0.01$, $M_{\\rm S}=7.0$ and $M_{\\rm A}=0.7$ and (4) $\\beta = 0.1$, $M_{\\rm S}=7.0$ and $M_{\\rm A}=2.0$. Similar studies provided by \\cite{ostriker01} characterized different models by their pressure ratio. However, our results show completely different polarization maps for the two coincident $\\beta$-value models. This because the super-Alfvenic flows tend to tangle the magnetic field lines, what is not seen in the sub-Alfvenic models (Model 2), even with similar pressure ratio.  The super-Alfvenic case shows a randomly distributed column density maps, with high constrast between the denser and rarefied regions. On the other hand, sub-Alfvenic cases are more filamentary, with contrasts increasing with the sonic Mach number. This general picture is independent on the angle between the external mean magnetic field and the LOS $\\theta$. However, for $B_{ext}$ nearly parallel to the LOS, the observed polarization will mostly be due to the random fluctuation component $\\delta B$. This effect is noticeable comparing Figs. 1 and 2.  We see that for sub-Alfvenic turbulence the large scale density enhancements are mostly parallel to the mean magnetic fields, with exception to the very dense cores, which can easily change the orientation of the magnetic field. As a consequence, polarization maps will present dense structures mostly aligned with the mean magnetic field. This effect also play a role on the generation of the polarization maps. Since we integrate the polarization vectors along the LOS, the low density cells will systematically increase the homogeneous contribution, as well as the resulting polarization degree. To avoid this effect, we disregarded the contribution from low density cells using a threshold, which depends on the model used. For the models where the magnetic field is oriented parallel to the LOS, the polarization maps will show polarization vectors randomly oriented in respect to the density structures. It reveals the degeneracy on the polarimetric maps between the super-Alfvenic models with those with $B$ nearly parallel to the LOS.   \\subsection{Polarization degree versus emission intensity}  The polarization maps showed that the polarization degree is anti-correlated to the column density, in exception to the subsonic case. This result is in agreement with the observations, which revealed ``polarization holes\" associated to the dense cores for most of the regions observed.  Observationally, \\cite{wolf03} showed that the polarization maps of molecular clouds follow the relation $P \\propto I^{-\\gamma}$, with $\\gamma \\sim 0.5 - 1.2$. The same trend is observed from polarized extinction of background stars \\citep{arce98}. It was proposed that this behavior occurs due to changes on dust properties inside denser cores, or even by an increase in thermal pressure, causing depolarization. \\cite{cho05} studied the role of the radiative torques on the grain alignment at dense cores and obtained $\\gamma \\sim 0.5 - 1.5$, depending on the dust size distribution.  \\cite{padoan01} obtained a similar behavior from their numerical simulations of protostellar cores, though for only three dense cores of one single simulation. They assumed a cut-off on grain alignment efficiency for $A_V > 3$mag. For their model, if the alignment efficiency is independent on $A_V$, the polarization degree was shown to be independent on the column density. \\cite{pelkonen07} extended this work, improving the radiative transfer. They naturally obtained a decrease in the grain alignment at denser regions, explaining the lower degree of polarization.  In our models, we assumed perfect grain alignment (independent on $A_V$). Therefore, the depolarization is exclusively due to the dispersion increase of the polarization angles in denser regions. We obtained $\\gamma \\sim 0.5$ for Models (2) and (3), but no correlation (i.e.\\ $\\gamma \\sim 0$) was found for Models (1) and (4). For Model (1), even the densest cores are unable to tangle the magnetic field lines and the polarization degree is homogeneously large. For Model (4), we have the opposite situation. The super-Alfvenic turbulence causes a strong dispersion of the magnetic field even at the less dense regions. For this case, the polarization degree is low everywhere. For Models (2) and (3), the turbulence is unable to destroy the magnetic field structure, but is able to create the dense cores by shocks. The cores are dense enough to drag the magnetic field lines and to increase the local $M_A$.  Possibly, our results differ from the obtained by \\cite{padoan01} because of: i - numerical resolution, ii - $M_A$, and iii - self-gravity. We used $512^3$ simulations (instead of a $128^3$) and, as a result, our magnetic field structure is less homogeneous and the density constrast is higher. The larger complexities present in our cubes increase the effect described in the previous paragraph. \\cite{padoan01} and \\cite{pelkonen07} used a single, super-Alfvenic, model. We showed that the polarization maps for $M_A$ present a flat $P \\times I$ correlation. Finally, self-gravity causes the colapse of the denser regions compressing the magnetic field within these cores. As a consequence, if no strong diffusion takes place, the polarization degree tends to grow. As a future work, we plan to study properly the depolarization at dense cores considering the grain alignment process and self-gravity.  \\subsection{Statistics of polarization angles}  We found that the distributions of polarization angles of sub-Alfvenic models are similar, even for different magnetic to gas pressure ratios. However, the dispersion of angles increases with $M_{\\rm A}$ and with the inclination of the external magnetic field regarding the line of sight ($\\theta$). Actually, we noticed that the critical parameter is the Alfvenic Mach number considering the magnetic field component projected into the plane of sky (i.e. $M_{\\rm A}^{\\rm sky} = \\delta v \\sqrt{4 \\pi \\rho}/B_{\\rm sky}$). We can compare these results with \\cite{padoan01} and \\cite{ostriker01}. The first used one single super-Alfvenic model, and obtained irregular (flat) distributions of polarization angle. The latest analysed supersonic models for $\\beta = 0.01, 0.1$ and 1.0. They obtained clearly gaussian distributions for $\\beta = 0.01$, with increasing dispersion for larger inclinations. Also, they obtained flatter distributions as $\\beta$ increases (i.e. as the Alfvenic Mach number increases). These are all in agreement with our results.  On the other hand, the power spectra analysis was showed to depend on the sonic Mach number. The spectra of the polarization angles show an increase in the power of small scales for increasing $M_{\\rm S}$. This occurs due to the amplification on the perturbations of the smallest scales for stronger turbulence. However, the same behavior is seen varying the inclination of the mean magnetic field. In this sense, there is a degeneracy between the Alfvenic Mach number and the orientation of the magnetic field. In this sense, structure functions of the polarization angle showed to be useful to avoid this degeneracy.  The sub-Alfvenic models presented SFs with slope $\\alpha \\sim 0.5$ ($SF \\propto l^{\\alpha}$), independent on the magnetic field orientation. On the other hand, the super-Alfvenic model presented flatter SFs ($\\alpha \\sim 0.3$). Therefore, we could conjecture that it could be possible to obtain the mean magnetic field intensity, independently on its orientation regarding the LOS, from the SF slopes. Needless to say that more models are needed to confirm this possibility.  Besides, SFs can potentially be used to study the turbulence eddies. Its flat profile at small scales may provide informations about the amplitude and size of the smallest turbulence cells. However, we showed that the results are affected by the observational resolution.   \\subsection{Improved CF technique}  The CF method has been proven an useful tool for the determination of the magnetic field in the ISM. We studied its validity using the obtained polarization maps from our models. \\cite{ostriker01} proposed that the CF method would only be applicable in restricted cases, in which the dispersion of the polarization angle is small ($< 25^{\\circ}$). It is consistent with the original approximations involving the derivation of the CF equation. We derived a generalized formula for the CF method, based on the same assumptions of the original work \\citep{chandra53}, but that accounts for larger dispersion models. Basically, we assume that the perturbations are Alfvenic and that we have an isotropic distribution of $\\delta V$. We, for the first time, successfully applied this equation to the super-Alfvenic and large inclination models.  We also studied dependency of the CF method and the observational resolution. As also shown by \\cite{ostriker01} and \\cite{heitsch01}, the CF method overestimates the magnetic field for coarser resolutions. Therefore, we propose a general equation to fit the observational data considering maps with different resolution. The asymptotic value of the given procedure provides the ``infinite resolution\" measurement from the CF method and is consistent with the expected values from the simulations.  As stated before, a possible limitation in the presented model is the absence of self-gravity effects. At the denser regions, the magnetic field configurations may possibly be different as the cloud collapses and drags the field lines. This process is responsible for the hour-glass structures observed in several gravitationally unstable clouds (e.g. Vall\\'ee \\& Fiege 2007). As a consequence, self-gravity increases the magnetic field locally and reduces the dispersion of the polarization angles within the dense clumps. To test the stability of the dense clumps in our simulations we may estimate the Jeans length ($\\lambda_J = c_s \\sqrt{\\pi/G\\rho}$). For the parameters chosen in Section 5.3, the denser structures, considering all models, are characterized by $l_{\\rm core} \\sim 0.1 - 0.5$pc, $n_{\\rm core} \\sim 10^{5-6}$cm$^{-3}$ and $T=10 - 100$K, resulting in $\\lambda_J \\sim 0.1 - 1$pc. Therefore, since $l_{\\rm core} \\sim \\lambda_J$, the denser structures may be unstable, at least for the given parameters. On the other hand, self-gravity plays a role at small regions and may not be statistically important for the previous results (except for the $P \\times I$ correlation) as we studied regions much larger than the very dense clumps. We plan to study the effects of self-gravity on the obtained results in a further work.  \\subsection{Sub-Alfvenic versus super-Alfvenic turbulence}  The ratio of thermal gas to magnetic pressures ($\\beta$) is typically used as the dominant parameter on the characterization of the degree of magnetization of a cloud. In this sense, systems with similar $\\beta$ values should present similar distribution of structures and dynamics. However, we showed that the sonic ($M_{\\rm S}$) and Alfvenic ($M_{\\rm A}$) Mach numbers, which quantify the ratio of the kinetic to the thermal and magnetic pressures, respectively, divide the models in different regimes. For the case of polarization vectors, our simulations showed that $M_{\\rm A}$ is decisive.  For clouds with $M_{\\rm A} < 1$, the gas motions excited by turbulence are confined by the magnetic field and are not able to change its configuration. Actually, the perturbations in the magnetic field occur, but are small compared to the mean field ($\\delta B \\ll B_{\\rm m}$). In this case, the polarization vectors are uniform, as shown in Fig.\\ 4. For $M_{\\rm A} > 1$, the magnetic pressure is small compared to the kinetic energy of the turbulent gas and the mean magnetic field can be easily distorted. As a consequence, the polarization maps would show large dispersion of $\\phi$.  Obviously, a large dispersion of $\\phi$ can also be related to a projection effect. If the mean magnetic field is projected along the line of sight ($\\theta \\sim 90\\degr$), only the small perturbations $\\delta B$ will be seen as the polarization vectors. However, considering a large number of clouds, there is a very low probability for all to present $\\theta \\sim 90\\degr$. In this case, if observations systematically show very large dispersions of $\\phi$, it means that the turbulence in the ISM may be typically super-Alfvenic. Otherwise, the ISM then presents sub(quasi)-Alfvenic turbulence. Observations of a given cloud could then be compared to our Fig.\\ 4 to determine under which regime the turbulence is operating. It is particularly interesting since the ratio of magnetic to turbulent energy in the ISM is still subject of controversy \\citep{padoan02, girart06}.  \\subsection{Procedure for observational data analysis}  The number of simulations presented in this work, as well as the numerical resolution, must be increased in future works. In any case, from the models we have, we intend to provide observers with a straightforward procedure to characterize the magnetic field and turbulence properties of molecular clouds.  Firstly, from the polarization maps of a given region one should obtain the second order structure function of the polarization angle. From the SF, it is possible to characterize the turbulent cascade and the magnetic field. The extension of the flat profile at small scales give the turbulence cut-off scale. On the other extreme, the flat profile at large scales indicate the energy injection lengths. From the maximum slope of $SF \\propto l^{\\alpha}$, it is possible to determine the averaged Alfvenic Mach number and, as a consequence, the magnetic field intensity.  Another method to obtain the magnetic field intensity is based on the CF technique. From the observed velocity dispersion it is possible to estimate the amplitude of the random component of the magnetic field from Eq.\\ (7). The total magnetic field is then obtained from Eq.\\ (9), using the dispersion of the polarization angle. To avoid the dependence on the observational resolution, it is suggested to evaluate the dispersion of the polarization angle for different resolutions (which may be simulated by averaging neighboring vectors of the polarization maps) and determine the asymptotic total magnetic field from Eq.\\ (10). Finally, subtracting the total field by the random component, it is possible to determine the mean magnetic field projected in the plane of sky.  Also, combining the mean magnetic field obtained from both methods, it is possible to estimate the angle between the mean magnetic field and the LOS.  \\subsection{Grain alignment}  Although it was not considered in the present calculations, a correct treatment of grain alignment is needed for a full understanding of the polarization in molecular clouds. For instance, we showed that the polarization degree is anti-correlated with the column density with slope $\\gamma \\sim -0.5$, while observations sometimes give $\\gamma < -1.0$ \\citep{goncalves05}. This difference is related with the alignment efficiency at different regions of the cloud \\citep{cho05}.  The theory of grain alignment has developed fastly during the past decade (see Lazarian [2007]). It is currently believed that radiative torques play a major role on the alignment process and it strongly depends on $A_V$ (see Lazarian \\& Hoang [2007]). With increasing extinction ($A_V > 2$), the radiative torques are less effective and only large grains are aligned. All in all, both observations \\citep{arce98, whittet08} and theory \\citep{hoang08}, suggest that there is a range of $A_V$ for which our assumptions are correct. It might happen that subsonic mechanical alignment of irregular grains, introduced in \\cite{lazho07b}, extends the range of $A_V$ over which grains are aligned when compared to the estimates based on radiative torques only.  The observed band is also selective regarding the dust sizes and different bands reveal the polarization of different dust components. All these effects will be included in a future work, and a more realistic study of the polarization intensity distribution will be obtained.   \\subsection{Polarization from molecular and atomic lines}  In the present work we focused on calculating synthetic polarization maps of FIR emission from dust particles, which were assumed to be perfectly aligned with the magnetic field. Unfortunately, due to inefficient grain alignment at the dense cloud cores, the dust polarization degree may decrease and different methods have to be used.  The polarization of molecular lines have been shown to be an additional tool for the study of the magnetic fields in the ISM \\citep{girart99,greaves02, girart04,cortes05}. Molecules are present in the dense and cold cores of the clouds and may be detected by thermal line emissions. Polarimetric maps of molecular emission can be used on the study of regions with $A_V > 10$, and can be directly associated with the Zeeman measurements. Based on the Goldreich-Kylafis effect \\citep{gold81,gold82}, the molecular sublevel populations will present imbalances due to the magnetic field, generating polarized rotational transitions. However, the survival of molecules depend on restrict conditions, as for $A_V > 10$ molecules may be frozen into dust particles. Another difficulty regarding this method is the fact that the GK effect generates polarization either parallel or perpendicular to the magnetic field lines, making the polarization maps.  On the other hand, polarized scattering and absorption from atoms and ions provide information about the magnetic field in warm and rarefied regions, like the diffuse ISM and the intergalactic medium.  Polarization arising from aligned atoms and ions is a new method \\citep{yan07a}. Unlike molecular lines that live in the excited state long enough to be imprinted by the magnetic field, the atomic exited states are short lived and tend to decay in timescales shorter than the Larmor precession of the atom. However, species with fine and/or hyperfine structure of ground or metastable states can be aligned. This fact opens new horizons for polarimetric studies of magnetic fields \\citep{yan06, yan07a, yan07b}.  It is useful to comment here that the results shown in this work are also valid for the observed polarization maps of molecular and atomic emission lines. This because the assumptions made for the calculations disregard any special consideration about the emitting species, which could be atoms, molecules or dust particles.  Also, it is worth mentioning that these techniques are either a substitute, for regions where no FIR dust emission is detectable, or complementary to the dust polarized emission but at different wavelengths (e.g.\\ optical and UV radiation). Depending on the $A_V$ range considered, polarization of atoms and molecules may complement the dust emission and absorption maps, which are usually much more detailed.   \\subsection{Comparison with previous works}  In this work we studied of polarization maps and its applicability on the determination of magnetic fields in molecular clouds based on numerical simulations. Here we compare the obtained results with the previous theoretical works.  \\cite{ostriker01} performed numerical simulations, with $256^3$ resolution, considering plasma $\\beta$ values 0.01, 0.1 and 1.0. Their results showed homogeneous polarization maps for $\\beta = 0.01$ (strongly magnetized turbulence), and a complex distribution of polarization vector for $\\beta = 0.1$ and 1.0 (weakly magnetized turbulence). They also obtained an increase in the dispersion of polarization angles with the increase of the external magnetic field inclination regarding the line of sight. These results are in agreement with our models, except for the fact that two of our models with equal $\\beta$ presented completely different polarization maps. This because the $\\beta$ value does not reveal how the magnetic field lines respond to the turbulence. The Alfvenic and sonic Mach numbers reflect how strong is the turbulent pressure compared to the magnetic and thermal pressures, respectively.  They obtained a higher polarization degree for $\\beta$ = 0.01, obviously because of the magnetic field intensity, but larger values of $P$ for larger column densities (i.e.\\ for larger $I$), in disagreement with observations. We believe that this was caused by their method for obtaining the polarization degree. They obtained the integrated Stokes parameters, weighted by local density, for all cells along the LOS. We, on the other hand, used a threshold on density to avoid the contribution of very rarefied regions (where the magnetic field structure is systematically more uniform). \\cite{padoan01} focused their work on the polarization of dust emission from dense cores. They implemented a more realistic calculation of the polarization degree based on the efficiency of the alignment for different values of $A_{\\rm V}$. In this case, they were able to obtain a decreasing polarization degree with the total intensity related to the grain properties, and not to the statistics of polarization vectors along the line of sight. We also believe that the numerical resolution may be playing a role on the polarization degree. More refined simulations systematically result in more complex structures for density and magnetic field lines. As a consequence, the alignment vectors along the line of sight present larger dispersion resulting in a lower polarization degree.  We found no previous theoretical work presenting an extended statistical anaylsis considering all the PDF, Spectra and Structure Function of polarization angle, and therefore no comparison can be made.  \\cite{ostriker01} and \\cite{padoan01} also tested the CF technique using their simulations, considering Eq.\\ (8). They obtained good agreement, with a correction factor of $\\sim 0.5$, between the calculated estimations and the expected values only for the models with $\\delta \\phi < 25\\degr$. \\cite{heitsch01} presented an extended analysis using a larger number of models, with different physical parameters and numerical resolutions (including 1 model with $512^3$ resolution). They also studied the effects of observational resolution on the obtained maps. They concluded that coarser resolutions result in more uniform polarization vectors. As a consequence, the CF method overestimates the magnetic field intensity. This is in full agreement with our results. They also tested the reliability of the CF technique in weakly magnetized clouds. They proposed the modified equation $B_{CF}B_{CF}^{mod}=4\\pi \\rho [\\delta v_{\\rm LOS}/\\delta(\\tan \\phi)][1+3\\delta(\\tan \\phi)^2]^{1/2}$, \\footnote{Here, $B_{CF}$ is the value obtained using the standard CF equation, and $B_{CF}^{mod}$ is the corrected value proposed by \\cite{heitsch01}.} which gave good results compared with the expected values for their models, with discrepancies of a factor $<2$.  We tested their equation to our models 3 and 4 with $\\theta = 0$, representing a strong and weakly magnetized cloud, respectively. For Model 3, the obtained value is in agreement with that shown in Table 2. The ratio between the two measurements is $B_{CF}^{mod}/B_{\\rm CF}^0 = 0.9$. For Model 4, the proposed equation underestimates the magnetic field, and compared to with our method it gives $B_{CF}^{mod}/B_{\\rm CF}^0 \\sim 0.3$. Despite of the few simulations available for compariron, their method seems to systematically underestimate the magnetic field intensity. More tests are needed to determine which method may give the best results.  Furthermore, even though not addressed by \\cite{heitsch01}, we studied the dependence of the polarization angles and the CF technique with the inclination of the magnetic field regarding the LOS. We showed that there is a degeneracy between the results of weakly magnetized clouds and strongly magnetized clouds with high $\\theta$. The modified CF formula presented in this work gave good results for all cases.  As discussed before, even though not taking into account the self-gravity in our simulations it mostly induces changes at small scales, as noted by \\cite{heitsch01}. As a result, they showed the CF technique to be insensitive to self-gravity."
    },
    "0801/0801.2560_arXiv.txt": {
        "abstract": "Terrestrial planets form in a series of dynamical steps from the solid component of circumstellar disks. First, km-sized planetesimals form likely via a combination of sticky collisions, turbulent concentration of solids, and gravitational collapse from micron-sized dust grains in the thin disk midplane. Second, planetesimals coalesce to form Moon- to Mars-sized protoplanets, also called ``planetary embryos''. Finally, full-sized terrestrial planets accrete from protoplanets and planetesimals. This final stage of accretion lasts about 10-100 Myr and is strongly affected by gravitational perturbations from any gas giant planets, which are constrained to form more quickly, during the 1-10 Myr lifetime of the gaseous component of the disk. It is during this final stage that the bulk compositions and volatile (e.g., water) contents of terrestrial planets are set, depending on their feeding zones and the amount of radial mixing that occurs.  The main factors that influence terrestrial planet formation are the mass and surface density profile of the disk, and the perturbations from giant planets and binary companions if they exist.  Simple accretion models predicts that low-mass stars should form small, dry planets in their habitable zones.  The migration of a giant planet through a disk of rocky bodies does not completely impede terrestrial planet growth.  Rather, \"hot Jupiter\" systems are likely to also contain exterior, very water-rich Earth-like planets, and also \"hot Earths\", very close-in rocky planets.  Roughly one third of the known systems of extra-solar (giant) planets could allow a terrestrial planet to form in the habitable zone. ",
        "introduction": "Recent research has developed a model for the growth of the terrestrial planets in the Solar System via collisional accumulation of smaller bodies (e.g., Wetherill 1990).  Several distinct dynamical stages have been identified in this process, from the accumulation of micron-sized dust grains to the late stage of giant impacts.  In this article, I first review the current state of knowledge of the stages of terrestrial planet formation ($\\S$ 2).  Second, I explore the effects of external parameters on the accretion process, including new work showing that simple accretion models predict that low-mass stars should preferentially harbor low-mass terrestrial planets in their habitable zones ($\\S$ 3).  These include the orbit of a giant planet and the surface density profile of the protoplanetary disk.  Third, I examine the effects of the migration of a Jupiter-mass giant planet on the accretion of terrestrial planets ($\\S$ 4), and also apply accretion models to the known set of extra-solar planets to show that roughly one third of the known systems could have formed an Earth-like planet in the habitable zone ($\\S$ 5).  Finally, conclusions and avenues for future study are presented ($\\S$ 6). ",
        "conclusions": "In this proceedings, I have summarized the current state of knowledge about terrestrial planet formation in extra-solar planetary systems.  In $\\S$ 2, I reviewed the stages of terrestrial planet formation, from micron-sized grains to planetesimals, protoplanets, and full-sized planets.  In $\\S$ 3 I described the effect of external parameters such as the disk and giant planet properties on the terrestrial accretion process as well as radial mixing and water delivery.  In $\\S$ 4, I presented recent results showing that terrestrial planets can form in systems with close-in giant planets, assuming those to have migrated to their final locations.  In fact, close-in giant planets should be accompanied by both very close-in \"hot Earths\" and exterior ocean planets (Raymond \\etal 2006b). In $\\S$ 5, I used results from previous work to derive limits designed to predict, based on the observed orbits of giant planets, where to search for other Earths among the known extra-solar systems. Approximately one third of the known systems of giant planets are good candidates for harboring an Earth-like planet (Mandell \\etal 2007).  Many issues remain to be resolved in each topic I described.  The details of planetesimal formation are still poorly known, although recent results suggest that several process including turbulence and migration of meter-sized bodies acting in tandem might be the solution (Johansen \\etal 2007).  The details and consequences of giant impacts are not well understood, in terms of the fate of collisional debris and compositional changes induced by the impacts (Genda \\& Abe 2005; Asphaug \\etal 2006; Canup \\& Pierazzo 2006).   In addition, some of the effects of external parameters are perhaps less well understood than we would like to think.  For example the giant planet eccentricity - terrestrial planet water content correlation (Chambers \\& Cassen 2002; Raymond \\etal 2004) inherently assumes that eccentric giant planets acquire their eccentricities early, before terrestrial accretion.  If this assumption is wrong, then one can imagine a scenario in which  terrestrial planets tend to form with relatively circular giant planets; this is a beneficial situation in terms of water delivery.  However, a late, impulsive eccentricity increase can destabilize the orbits of terrestrial planets and remove them from the system entirely (Veras \\& Armitage 2006).  Indeed, there are many known extra-solar planets that would be favorable for terrestrial planet formation if their orbits were more circular (see Fig.~\\ref{fig:xsp}).  Thus, if giant planet eccentricity is acquired late, many systems may undergo ``planetary system suicide'', forming an Earth-like planet and subsequently destroying it.  A more complete, holistic view of planet formation and evolution is needed to distentangle these effects."
    },
    "0801/0801.0423_arXiv.txt": {
        "abstract": "{} {A recent catalogue by Flesch \\& Hardcastle presents two major anomalies in the spatial distribution of QSO candidates: $i/$ an apparent excess of such objects near bright galaxies, and $ii/$ an excess of very bright QSO candidates compared to random background expectations in several regions of the sky. Because anyone of these anomalies would be relevant in a cosmological context, we carried out an extensive analysis of the probabilities quoted in that catalogue.} {We determine the nature and redshift of a subsample of 30 sources in that catalogue by analysing their optical spectra (another 11 candidates were identified from existing public databases). These have allowed us to statistically check the reliability of the probabilities QSO status quoted by Flesch \\& Hardcastle for their candidates.} {Only 12 of the 41 candidates turned out QSOs (7 of which have been identified here for the first time).} {The probabilities of the QSOs' being the candidates given by Flesch \\& Hardcastle are overestimated for $m_B\\le 17$  and for objects projected near ($\\le 1$ arcmin) bright galaxies. This is the cause of the anomalies mentioned above.} ",
        "introduction": "At present, statistics on QSO distributions can be carried out using the recent public compilations of QSOs identified in the SDSS (Adelman-McCarthy et al.\\ 2007) and 2dF (Croom et al.\\ 2004) surveys. Furthermore, Flesch \\& Hardcastle (2004; hereafter FH04) present an all-sky catalogue with 86\\,009 optical counterparts of radio/X-ray sources as QSO candidates (with a probability greater than 40\\%, according to FH04) that were not identified previously as such. This could be a useful catalogue to look for QSOs in regions that were not covered by SDSS or 2dF, or for deeper magnitudes than their completeness limit. Nevertheless, since the FH04 catalogue does not give the identification of each source but just a statistical estimate of the probability that it be a QSO (or a star, a galaxy, or a wrongly identified optical counterpart of the X-ray/radio source), it must be used with care.  Flesch\\footnote{See his web-page http://quasars.org/ .} claims that his catalogue `reveals a few galaxy-centered fields which  are populated by QSOs/candidates more thickly than expected  by the background density of such QSOs,' which would argue in favour of the hypothesis of QSOs with anomalous redshift and at the same distance as the galaxies with which they are associated (Arp 1998, Burbidge 2001). Because this is against one of the most basic assumptions in standard modern cosmology, we consider it interesting to check it on the basis of rigorous statistical computation. This is the main objective of this paper. In \\S \\ref{.excesses}, we  analyse in which cases (on the basis of the probabilities quoted by FH04)  there are anomalies in the distribution of QSOs. By conducting optical spectroscopy (\\S \\ref{.select}), we directly check the nature of a subsample of FH04 sources. On the basis of these observations, in  \\S \\ref{.probab} we check the reliability of the probabilities quoted by FH04. In \\S \\ref{.conclusions} we  reanalyse the distribution of QSOs and discuss the implications of our results. ",
        "conclusions": "\\label{.conclusions}  In FH04 there are 88 different objects (106 cases with repetitions) with $m_B\\le 18$ within a radius of 3 arcminutes of RC3 galaxies. Of these 88 objects, there are spectroscopic observations of the optical counterparts in 26 cases (36 cases including repetitions), and 10 of these (the same number with repetitions) are in fact QSOs (while the 16 other objects (26 with repetitions) are stars, galaxies, or HII regions). The average expected number of background cases around all RC3 cases should be $\\approx 70$ (\\S \\ref{.explorc3}) and, in our subsample, should be $\\approx \\frac{36}{106}\\times 70=24$ cases, which is higher than the value derived from the observations. This stems from the incompleteness of the FH04 catalogue.  Within a radius of 10 arcminutes of RC3-galaxies, there are 14 objects (no repetitions) with $m_B\\le 15.5$, of which five (36\\% of the sample) were observed and only one (\\# 5 with $m_B=15.4$) turned out a QSO. The average expected number of cases around all RC3 galaxies is 7.6 (\\S \\ref{.explorc3}), and in our subsample should be $\\frac{5}{14}\\times 7.6=2.7$ cases. This result agrees with expectations from a background distribution of QSOs.  Limiting the study to RC3 galaxies by $m_{B,gal}\\le 14.5$, $m_{B,gal}\\le 13.5$, and $m_{B,gal}\\le 12.5$, the number of galaxies is reduced to 7,675/2,566/883. For the subsample with $m_b\\le 18$, $d\\le 3'$ we obtained 5/2/2 QSOs to be compared with the expected 7.8/2.7/0.92. For the subsample with $m_b\\le 15.5$, $d\\le 10'$, we obtained 1/0/0 to be compared with 0.50/0.17/0.06. There is no statistical inconsistency with the claim that all QSOs are background QSOs rather than  being associated with the galaxies.  As for the distribution of the QSOs with magnitude or distance from RC3 galaxies, as pointed out in \\S \\ref{.probab}, there is no longer any inconsistency with background predictions: almost all the confirmed QSOs in our subsample are faint (all of them with magnitudes $m_B\\ge 17.0$, except \\#5)  and distant enough from the centres of RC3 galaxies to be compatible with background expectations. See Fig. \\ref{Fig:QSOdist2} for the distance distribution and note how the observed QSOs are fewer than the background expectations for all distances. Some of the most troublesome candidates, such as those shown in Fig. \\ref{Fig:examples}, were in the end confirmed to be objects different from QSOs.  The only QSO with $m_B\\le 15.5$ (\\# 5) within 10$'$ from an RC3 galaxy is associated with NGC 1136, which was previously reported as a galaxy with peculiar associations (Arp 1981). Objects claimed by Arp to be anomalous are much fewer than 23,011 RC3 galaxies. They are $\\sim 10^2$, so the probability of finding a QSO with $m_B\\le 15.5$ within 10$'$ from some of them is $P\\sim 0.06$. It is not few enough to claim it separately from other facts as a possible anomalous case, but a fact to be added in the discussion of the case of NGC 1136.  For the object \\# 6, now confirmed to be a QSO with $m_B=16.5$ at only 30$\"$ from PHL 1459 (a galaxy with $m_B=17.2$ and X-ray emission), we are not within the statistics of RC3 galaxies. There are $\\approx 4.3\\times 10^5$ galaxies in the whole sky up to the magnitude $m_B=17.2$ (Metcalfe et al.\\ 1991), so, using Eq.\\ (\\ref{densQSO}), the expected number of cases like this within a 30-arcsecond distance of these galaxies is 2.2. The discovery of only one case is not enough to claim any statistical anomaly. If we take into account the peculiarity that PHL 1459 is an X-ray source and has $\\approx 120$ counts/hour in ROSAT (Voges et al.\\ 1999, 2000), the probability is somewhat lower. There are $\\sim 0.4$ galaxies/deg$^2$ up to this X-ray flux (Voges et al. 1999; Zickgraf et al. 2003), i.e. around 16,000 galaxies like this in the whole sky. The probability of finding a QSO with $m_B=16.5$ within 30$\"$ of one of them is 0.09; not low enough to discard a background projection coincidence.  \\begin{figure} \\vspace{1cm} \\resizebox*{7cm}{7cm}{\\includegraphics{QSOdist2.eps}}\\par \\caption{Cumulative counts of QSOs of this paper up to magnitude $m_B=18$ [34\\% (36/106) of the QSO/RC3--galaxy pairs] as a function of the maximum distance to the galaxy. QSOs associated with two galaxies count twice with each corresponding distance.} \\label{Fig:QSOdist2} \\end{figure}  Summing up, within the subsample of 41 objects from FH04, we have not found any statistical anomaly with respect to the expectation of QSOs as background sources. Since we observed a representative subsample of the most controversial cases that could in principle present some excesses of QSOs, we think that there should be in principle no reason to expect statistical anomalies in the FH04 catalogue, which is anyway a useful database for looking for QSOs in regions not covered by other surveys, such as SDSS or 2dF, or for deeper magnitudes. Nevertheless, there are still anomalies in the QSO distributions presented with some other catalogues (e.g.\\ Arp 1998; Burbidge 2001) that require further attention. Further rigorous statistical analysis like those presented here and in L\\'opez-Corredoira \\& Guti\\'errez (2006$a,b$) are necessary.   \\  {\\bf Acknowledgments:}  Thanks are given to Rub\\'en J. D\\'\\i az (Gemini Observatory) for providing telescope time at 2.15 m CASLEO, to Eric Flesch for his comments and criticism on a draft of this paper, to the referee Mira Veron for her helpful comments for improving it, to T. J. Mahoney (IAC) and Joli Adams (A\\&A) for proofreading the paper.  Based on observations made with the telescopes: 2.15 m CASLEO in the observatory of ``El Leoncito'', San Juan, Argentina; 1.93 m OHP (these observations were funded by the Optical Infrared Coordination network, OPTICON, a major international collaboration supported by the Research Infrastructures Programme of the European Commission Sixth Framework Programme), Haute-Provence, France; and 2 m IUCAA (Pune, India).  This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France.  The Digitized Sky Surveys were produced at the Space Telescope Science Institute under US 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.  Funding for the creation and distribution of the SDSS Archive has  been provided by the Alfred P. Sloan Foundation, the Participating  Institutions, the National Aeronautics and Space Administration, the  National Science Foundation, the U.S. Department of Energy, the Japanese  Monbukagakusho, and the Max Planck Society. The SDSS Web site is  http://www.sdss.org/. The SDSS is managed by the Astrophysical Research  Consortium (ARC) for the Participating Institutions. The Participating  Institutions are The University of Chicago, Fermilab, the Institute for  Advanced Study, the Japan Participation Group, The Johns Hopkins University,  the Korean Scientist Group, Los Alamos National Laboratory,  the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington.  The first author (MLC) was supported by the {\\it Ram\\'on y Cajal} Programme of the Spanish Science Ministry."
    },
    "0801/0801.2432_arXiv.txt": {
        "abstract": "We have performed a joint analysis of prompt emission from four bright short gamma-ray bursts (GRBs) with the Suzaku-WAM and the Konus-Wind experiments. This joint analysis allows us to investigate the spectral properties of short-duration bursts over a wider energy band with a higher accuracy. We find that these bursts have a high E$_{\\rm peak}$, around 1 MeV and have a harder power-law component than that of long GRBs. However, we can not determine whether these spectra follow the cut-off power-law model or the Band model.  We also investigated the spectral lag, hardness ratio, inferred isotropic radiation energy and existence of a soft emission hump, in order to classify them into short or long GRBs  using several criteria, in addition to the burst duration. We find that all criteria, except for the existence of the soft hump, support the fact that our four GRB samples are correctly classified as belonging to the short class. In addition, our broad-band analysis revealed that there is no evidence of GRBs with a very large hardness ratio, as seen in the BATSE short GRB sample, and that the spectral lag of our four short GRBs is consistent with zero, even in the MeV energy band, unlike long GRBs. Although our short GRB samples are still limited, these results suggest that the spectral hardness of short GRBs might not differ significantly from that of long GRBs, and also that the spectral lag at high energies could be a strong criterion for burst classification. ",
        "introduction": "The bimodal distribution of the duration of gamma-ray bursts (GRBs) indicates that there are two distinct classes of events. The long-duration bursts have typical durations of around 20 s, while the short-duration bursts (about 1/4 of the total) have durations of around 0.3 s (\\cite{Mazets 1981}; \\cite{Norris 1984}; \\cite{Dezalay 1992}; \\cite{Hurley 1992}; \\cite{Kouveliotou 1993}; \\cite{Norris 2000}). This suggests that the short and long GRBs may have different progenitors. Core collapse of massive stars and mergers of  compact binaries are considered likely models for the long and short GRBs, respectively (\\cite{Katz_Canel}; \\cite{Ruffert 1999}). Recently, thanks to the rapid position information provided by HETE-2 (\\cite{Rikker 2003}) and Swift (\\cite{Gehrels 2004}),  afterglow observations have progressed dramatically and X-ray and optical afterglow emissions were discovered from some short GRBs (\\cite{Hjorth 2005}; \\cite{Fox 2005}), as well as long ones. Some short GRB afterglows have been found to be  associated with galaxies not undergoing star formation, while some long GRBs have been found to be associated with energetic supernovae. These results support the hypothesis that long and short GRBs indeed have different progenitors.  From the point of view of prompt gamma-ray emission, although the spectral characteristics of long GRBs have been well studied (\\cite{Frontera 2000}: BeppoSAX; \\cite{Kaneko 2006}: BATSE; \\cite{Sakamoto 2005b}: HETE-2), our understanding of that of short GRBs is still incomplete, in part due to their very short durations. The spectral characteristics of BATSE short GRBs are often characterized by the hardness ratio, that is, the ratio of 100--300 keV to 25--100 keV counts or fluence (\\cite{Kouveliotou 1993}; \\cite{Cline 1999}). These studies suggest that, although there is considerable overlap,  short GRBs tend to be harder than long GRBs. Paciesas et al. (2001) and Ghirlanda et al. (2004) compared the spectral parameters of bright BATSE short GRBs with those of long GRBs by spectral fitting. They pointed out that the spectra of short GRBs are well described by a cut-off power-law model and that the Band model (\\cite{Band 1993}) did not improve the fit. They also confirmed that the spectra of short GRBs are harder than those of long GRBs. This was found to be a consequence of a flat low energy photon index, rather than a difference in the peak energies.  There is another complicating issue in the classification of short and long GRBs. Since their duration distributions overlap, other distinguishing criteria have been proposed, such as differences in the host galaxy (\\cite{Hjorth 2003}; \\cite{Hjorth 2005}), spectral hardness (\\cite{Cline 1999}), spectral lag (\\cite{Norris 2002}; \\cite{Norris 2006}), isotropic radiation energy (\\cite{Amati 2002}; \\cite{Amati 2006}), and the existence of a soft hump (\\cite{Norris 2006}), or a combination of these and other characteristics (\\cite{Donaghy 2006}). Indeed, some GRBs cannot be classified as short or long using the burst duration alone. For example, GRB 040429 was classified as short using the burst duration (2.39 s), but other properties such as spectral lag and hardness resembled those of long GRBs (\\cite{Wiersema 2004}; \\cite{Fox_Moon 2004}; \\cite{Donaghy 2006}). GRB 051227 (\\cite{Hullinger 2005}; \\cite{Barthelmy 2005a}; \\cite{Sakamoto 2005}) had a long enough duration of 8.0 s to be classified as long. However, this burst had no significant spectral lag in the initial spike, which is seen in many long GRBs (\\cite{Norris 1996}), and it also had a long soft hump in the light curve. These two properties are similar to those of short GRBs. GRB 060614 (\\cite{Gehrels 2006}) exhibited a long duration (102 s). However, it did not have any significant spectral lag in the initial spike, and also there were no signs of any associated supernova, despite its distance of z = 0.125, which should have been close enough to detect it. These examples show why the burst duration should no longer be considered a sole indicator for distinguishing between the short and long classes of bursts. Indeed Donaghy et al. (2006) have suggested the terms ``short population bursts'' and ``long population bursts'' to distinguish the two classes. In this paper, we will use simply ``short'' and ``long'' to describe these classes.  Here, we report on the joint spectral analysis of four bright short GRBs simultaneously observed by the Suzaku-WAM (Suzaku Wide-band All-sky Monitor) and Konus-Wind. The WAM and  Konus have wide energy ranges, from 50--5000 keV and 10--10000 keV, respectively, and the WAM has the largest effective area from 300 keV to 5000 keV of any current experiment with spectral capabilities. Thus, we can investigate the spectral characteristics of short GRBs up to MeV energies. We also classify the four GRBs as short or long using our data with high statistics up to the MeV energy region using a number of criteria: spectral lag, hardness ratio, isotropic radiation energy and existence of a soft hump. ",
        "conclusions": "We performed a joint analysis of four bright, hard-spectrum, short-duration GRBs, localized by the IPN, with the Suzaku-WAM and Konus-Wind, in order to investigate the spectral properties of short GRBs into the MeV energy band. From the spectral analysis, we found that these bursts have a high E$_{\\rm peak}$ around 1 MeV; the spectral parameters can be constrained tightly by joint fitting. However, we could not determine whether these spectra follow the CPL model or the Band model, because there is no  significant improvement between these models. The power-law photon index is slightly flatter than that of long BATSE GRBs. In particular, GRB 051127 has a very flat spectrum with a photon index, $\\alpha$, that is larger than the synchrotron limit of $-$2/3.  We also examined several other spectral and temporal properties to confirm that the bursts belong to the short class. Most properties, such as the spectral lag, the spectral hardness, and the isotropic radiation energy, indicated that the four bursts were indeed in this class. In addition, a broad-band spectral analysis with the WAM and Konus revealed some interesting properties of short GRBs: i) There is no evidence for very hard GRBs that have the large hardness ratios seen in BATSE short GRBs. ii) The spectral lag of our short GRBs does not show any clear evolution, even in the MeV energy region, unlike that of long hard GRBs. Although our short GRB sample is still limited, these results might suggest one possibility, that the hardness ratios of the short GRBs are not so different from those of long GRBs, and also that the spectral lag at high energies can be a strong criterion for classifying short and long GRBs.   \\bigskip M. O. is supported by the Research Fellowships of the Japan Society for the Promotion of Science for Young Scientists (2006). K. H. is grateful for IPN support under the NASA LTSA program, grant FDNAG5-11451, the INTEGRAL Guest Investigator program, NAG5-12706 and NNG06GE69G, and the Suzaku Guest Investigator program NNX06AI36G. The Konus-Wind experiment is supported by a Russian Space Agency contract and RFBR grant 06-02-16070. We are grateful to Richard Starr for providing the Mars Odyssey data, and to Giselher Lichti, Arne Rau, and Andreas von Kienlin for providing the INTEGRAL SPI-ACS data. We would like to thank Jay Norris for useful advice concerning the spectral lag analysis."
    },
    "0801/0801.3744_arXiv.txt": {
        "abstract": "Observations show that small-amplitude prominence oscillations are usually damped after a few periods. This phenomenon has been theoretically investigated in terms of non-ideal magnetoacoustic waves, non-adiabatic effects being the best candidates to explain the damping in the case of slow modes. We study the attenuation of non-adiabatic magnetoacoustic waves in a slab prominence embedded in the coronal medium. We assume an equilibrium configuration with a transverse magnetic field to the slab axis and investigate wave damping by thermal conduction and radiative losses. The magnetohydrodynamic equations are considered in their linearised form and terms representing thermal conduction, radiation and heating are included in the energy equation. The differential equations that govern linear slow and fast modes are numerically solved to obtain the complex oscillatory frequency and the corresponding eigenfunctions. We find that coronal thermal conduction and radiative losses from the prominence plasma reveal as the most relevant damping mechanisms. Both mechanisms govern together the attenuation of hybrid modes, whereas prominence radiation is responsible for the damping of internal modes and coronal conduction essentially dominates the attenuation of external modes. In addition, the energy transfer between the prominence and the corona caused by thermal conduction has a noticeable effect on the wave stability, radiative losses from the prominence plasma being of paramount importance for the thermal stability of fast modes. We conclude that slow modes are efficiently damped, with damping times compatible with observations. On the contrary, fast modes are less attenuated by non-adiabatic effects and their damping times are several orders of magnitude larger than those observed. The presence of the corona causes a decrease of the damping times with respect to those of an isolated prominence slab, but its effect is still insufficient to obtain damping times of the order of the period in the case of fast modes. ",
        "introduction": "\\label{sec:intro}  Solar prominences are large-scale coronal magnetic structures whose material, cooler and denser than the typical coronal medium, is in plasma state. Prominences are supported against gravity by the coronal magnetic field, which also maintains the prominence material thermally isolated from the corona. Small-amplitude oscillations in solar prominences were detected almost 40 years ago \\citep{harvey}. These oscillatory motions seem to be of local nature and their velocity amplitude is typically less than 2--3~km~s$^{-1}$. Observations have also allowed to measure a wide range of periods between 30~s \\citep{balthasar} and 12~h \\citep{foullon}. More recently, some high-resolution observations of prominence oscillations by the Hinode/SOT instrument have been reported \\citep{okamoto,berger,ofman}. From the theoretical point of view, the oscillations have been interpreted by means of the magnetoacoustic eigenmodes supported by the prominence body. A recent example is the work by \\citet{hinode} in which the observations of \\citet{okamoto} are interpreted as fast kink waves.  The reader is referred to \\citet{oliverballester02,ballester,banerjee} for extensive reviews of both observational and theoretical studies.  Evidence of the attenuation of small-amplitude prominence oscillations has been reported in some works \\citep{molowny99,terradasrad,lin2004}. A typical feature of these observations is that the oscillatory motions disappear after a few periods, hence they are quickly damped by one or several mechanisms. The theoretical investigation of this phenomenon in terms of magnetohydrodynamic (MHD) waves has been broached by some authors by removing the ideal assumption and by including dissipative terms in the basic equations. Non-adiabatic effects appear to be very efficient damping mechanisms and have been investigated with the help of simple prominence models \\citep{ballai,carbonell,spatial,terradas}. Nevertheless, other damping mechanisms have been also proposed, like wave leakage \\citep{schutgensA,schutgensB,schutgensToth}, dissipation by ion-neutral collisions \\citep{forteza} and resonant absorption \\citep{arregui}.  In a previous work \\citep[][hereafter Paper~I]{soler}, we have studied for the first time the wave attenuation by non-adiabatic effects of a prominence slab embedded in the corona. In that work the magnetic field is parallel to the slab axis and it is found that the corona has no influence on the internal slow modes, but it is of paramount importance to explain the damping of fast modes, which are more attenuated than in simple models that do not consider the coronal medium. Following the path initiated in Paper~I, here we investigate the wave damping due to non-adiabatic mechanisms (radiative losses and thermal conduction) in an equilibrium made of a prominence slab embedded in a coronal medium, but now we consider a magnetic field transverse to the slab axis. This configuration and that studied in Paper~I correspond to limit cases, since measurements with Zeeman and Hanle effects indicate that the magnetic field lines are skewed to the long axis of prominences. On average, the prominence axis and the magnetic field form an angle of about 20~deg. Thus, the skewed case is relegated to a future investigation.  The equilibrium configuration assumed here was analysed in detail by \\citet{JR92} and \\citet{oliver} in the case of ideal, adiabatic perturbations. The main difference between both works is in the treatment of gravity. \\citet{JR92} neglected the effect of gravity and so straight field lines were considered. On the other hand, \\citet{oliver} took gravity into account and assumed curved field lines according to the \\citet{kippen} model modified to include the surrounding coronal plasma \\citep{poland}. Despite this difference, both studies agree in establishing a distinction between different normal modes depending on the dominant medium supporting the oscillation. Hence, internal modes are essentially supported by the prominence slab whereas external modes arise from the presence of the corona. In addition, hybrid (or string) modes appear due to the combined effect of both media.  The investigation of the thermal attenuation of oscillations supported by such equilibrium is unsettled to date and, indeed, this is the main motivation for the present study. However, two works \\citep{terradasrad01,terradas} studied the wave damping in an isolated prominence slab. \\citet{terradasrad01} considered radiative losses given by the Newtonian law of cooling as damping mechanism and studied the attenuation in the \\citet{kippen} and \\citet{menzel} prominence models. Subsequently, \\citet{terradas} considered a more complete energy equation including optically thin radiation, plasma heating and parallel thermal conduction, and assumed straight field lines since gravity was neglected. The main conclusion of both works is that non-adiabatic mechanisms are only efficient in damping slow modes whereas fast modes remain almost undamped. Nevertheless, in the light of the results of Paper~I, the presence of the coronal medium can have an important repercussion on the wave damping. The investigation of this effect is the main aim of the present work. Therefore, we extend here the work of \\citet{terradas} by considering the presence of the corona and neglect the effect of gravity as in \\citet{JR92} for simplicity.  This paper is organised as follows. Section~\\ref{sec:equilibrium} contains a description of the equilibrium configuration and the basic equations which govern non-adiabatic magnetoacoustic waves. Then, the results of this work are extensively discussed in Sect.~\\ref{sec:results}. Finally, our conclusions are given in  Sect.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we have studied the wave attenuation in a system representing a quiescent solar prominence embedded in the coronal medium. The prominence has been modelled as a homogeneous plasma slab surrounded by a homogeneous medium with coronal conditions. Magnetic field lines have been assumed transverse to the prominence slab axis and the whole system has been bounded in the field direction by two photospheric rigid walls, in order to establish a realistic length for the field lines. The attenuation of the normal modes of such equilibrium has been investigated by considering parallel thermal conduction, radiative losses and plasma heating as non-adiabatic mechanisms, and focusing our study on the fundamental oscillatory modes. The main conclusions of this work are summarised next. \\begin{enumerate} \\item Slow modes are strongly attenuated by non-adiabatic mechanisms, their damping times being of the order of the corresponding periods. Fast modes are less affected and present greater damping times. \\item The most relevant damping mechanisms are prominence radiation and coronal thermal conduction. The first one dominates the damping of internal modes, while the second one is responsible for the attenuation of external modes. The combined effect of both mechanisms governs the damping of hybrid modes. Neither prominence conduction nor coronal radiation become of importance for realistic values of the length of magnetic field lines and the prominence width. \\item The attenuation of slow modes is not affected by the value of the free component of the wavenumber, $k_z$. On the contrary, the behaviour of fast modes is strongly dependent on $k_z$. \\item Thermal conduction allows energy transfer between the prominence slab and the coronal medium. Prominence radiation has an essential role in dissipating the extra heat injected from the corona and stabilises the oscillations. Thermal instabilities appear if the radiative losses from the prominence plasma are omitted or significantly reduced (e.g. caused by an increase of the optical thickness) since the plasma cannot dissipate the extra injected heat in an efficient way. \\item The damping time of fast modes is strongly sensitive to the equilibrium physical parameters while slow waves are less affected by the variation of the equilibrium conditions. \\item The presence of the corona produces a decrement of the damping time of internal modes with respect to the solutions supported by an isolated prominence slab. Nevertheless, this effect is not enough to obtain damping times of the order of the period in the case of fast modes. \\end{enumerate}   Considering the equilibrium parameters of Paper~I, the efficiency of non-adiabatic mechanisms on the damping of fast modes is smaller in the present case. This fact suggests that the orientation of magnetic field lines with respect to the slab axis has a relevant influence on the attenuation of fast modes, the configuration of Paper~I and the present one being limit cases. Moreover, fast modes are strongly sensitive to the equilibrium physical conditions, and it is possible to obtain small values of the damping time by considering extreme equilibrium parameters, such as very weak magnetic fields and very large prominence densities. In this way, fast modes show a wide range of theoretical damping times. On the other hand slow modes are always efficiently attenuated, with damping times of the order of their periods. This result suggests that the attenuation of prominence fast waves may be caused by other damping mechanisms not considered here. Some candidates could be resonant absorption \\citep{arregui} and ion-neutral collisions \\citep{forteza}. Among these mechanisms, resonant absortion may be a very efficient damping mechanism if non-uniform equilibria are considered, e.g. models with a transition region between the prominence and the corona. Other effects, as wave leakage, might only play a minor role in the damping of disturbances. Finally, future studies should take into account the prominence fine structure on the basis that small-amplitude oscillations are of local nature. Therefore, the investigation of the damping of fibril oscillations should be the next step.  R.~O. and J.~L.~B. want to acknowledge the International Space Science Institute teams ``Coronal waves and Oscillations'' and ``Spectroscopy and Imaging of quiescent and eruptive solar prominences from space'' for useful discussions. The authors acknowledge the financial support received from the Spanish MCyT and the Conselleria d'Economia, Hisenda i Innovaci\\'o of the CAIB under Grants No. AYA2006-07637 and PCTIB-2005GC3-03, respectively. Finally, R.~S. thanks the Conselleria d'Economia, Hisenda i Innovaci\\'o for a fellowship."
    },
    "0801/0801.4955_arXiv.txt": {
        "abstract": "We present the results of a systematic study of mid-IR spectra of Galactic regions, Magellanic \\hii\\ regions, and galaxies of various types (dwarf, spiral, starburst), observed by the satellites \\iso\\ and \\spitz. We study the relative variations of the 6.2, 7.7, 8.6 and $11.3\\mic$ features inside spatially resolved objects (such as \\M{82}, \\M{51}, \\xxxdor, \\M{17} and the \\orb), as well as among 90 integrated spectra of 50 objects. Our main results are that the 6.2, 7.7 and $8.6\\mic$ bands are essentially tied together, while the ratios between these bands and the $11.3\\mic$ band varies by one order of magnitude. This implies that the properties of the PAHs are remarkably universal throughout our sample, and that the relative variations of the band ratios are mainly controled by the fraction of ionized PAHs. In particular, we show that we can rule out both the modification of the PAH size distribution, and the mid-infrared extinction, as an explanation of these variations. Using a few well-studied Galactic regions (including the spectral image of the \\orb), we give an empirical relation between the $\\ipah{6.2}/\\ipah{11.3}$ ratio and the ionization/recombination ratio $G_0/n_e\\times\\sqrt{T_\\sms{gas}}$, therefore providing a useful quantitative diagnostic tool of the physical conditions in the regions where the PAH emission originates. Finally, we discuss the physical interpretation of the $\\ipah{6.2}/\\ipah{11.3}$ ratio, on galactic size scales. ",
        "introduction": "\\label{sec:intro}  The reprocessing of stellar light by dust in the infrared (IR) is widely used to probe embedded star formation. However, in the absence of other constraints, the physical properties which can usually be derived from an almost-featureless grain continuum emission is limited to global quantities, such as the dust mass and its average temperature. The ubiquity of numerous mid-IR aromatic features, in a wide variety of astrophysical objects and environments, potentially provides more articulate diagnostics of the physical conditions. Indeed, these features dominate the mid-IR spectra of evolved stars \\citep[{e.g.}][]{blommaert05,kraemer06}, the cool ISM \\citep[{e.g.}][]{abergel05,flagey06,povich07}, as well as whole galaxies \\citep[{e.g.}][]{verma05,smith07}, that have been extensively observed by the \\isoST\\ (\\iso), and are currently investigated with a higher sensitivity by the \\spitzST. In our Galaxy, one third of the stellar light is reprocessed by dust, while this fraction can go up to $99\\%$ and higher, in starburst galaxies. At solar metallicity, roughly $15\\%$ of the cooling is radiated through the most powerful mid-IR bands, centered at 3.3, 6.2, 7.7, 11.3 and 12.7$\\mic$.  Historically, these emission features were attributed to very small grains ($\\simeq10\\,$\\AA), transiently heated by single photon absorption, in order to account for the independence of the color temperature with the distance from the illuminating star, in several reflection nebulae \\citep{sellgren84}. In parallel, the central wavelengths of these bands were recognized to coincide with the vibrational modes of aromatic material \\citep{duley81}. These features are now commonly attributed to the molecular modes of Polycyclic Aromatic Hydrocarbons \\citep[hereafter PAHs;][]{leger84,allamandola85,allamandola89}, which are planar molecules made of $\\simeq10$ to 1000 carbon atoms, excited primarily by ultraviolet (UV) photons. With silicate and carbon grains, PAHs are a main component of dust models \\citep{desert90,dwek97,draine01,zubko04}. Their absorption efficiency has been modeled using astrophysical observations, laboratory measurements and quantum theory \\citep[in particular][]{desert90,joblin92,verstraete01,li01,mattioda05,mattioda05b,draine07,malloci07}. In addition to being major radiative coolants of the interstellar medium (ISM), PAHs are responsible for most of the photoelectric heating of the gas in photodissociation regions (hereafter PDRs) and the neutral interstellar medium, due to their high cumulative surface area \\citep[{e.g.}][]{bakes94,hollenbach97}. For the same reason, they probably play an important role in grain surface chemistry \\citep[{e.g.}][]{tielens87}. In our Galaxy, they contain $\\simeq15-20\\%$ of the depleted carbon \\citep[with solar abundance constraints]{zubko04}. As a consequence, they are believed to be part of the interstellar carbon condensation chain \\citep{cherchneff00,dartois05}.  From an extragalactic point of view, the luminosity of the $6.2\\mic$ feature can be used as a tracer of star formation \\citep{peeters04}. However, this tracer is biased by global parameters such as the ISM metallicity. Indeed, PAHs are underabundant in low-metallicity galaxies \\citep[{e.g.}][]{galliano03,galliano05,galliano08a,draine07b}. There is a general correlation between the PAH-to-continuum intensity ratio and the ISM metallicity \\citep{madden06,wu06,ohalloran06,smith07}, and consequently between the \\IRACiv/\\MIPSi\\ broadband ratio and the metallicity \\citep{engelbracht05}. The origin of this trend has been attributed to radiative and mechanical destruction mechanisms by \\citet{madden06} and \\citet{ohalloran06} respectively. Conversely, from the detailed modeling of the spectral energy distribution (SED) of nearby galaxies, \\citet{galliano08a} showed that the PAH-to-gas mass ratio at different metallicities coincides with the relative amount of carbon dust condensed in the envelopes of low-mass stars, during the Asymptotic Giant Branch phase (AGB). This study suggests that PAHs are injected into the ISM by their progenitors, the AGB stars, several hundreds of Myr after the beginning of the star formation, when the gas has already been enriched by more massive stars. This delay corresponds to the time needed for AGB stars to evolve off the main sequence. Therefore, the delayed injection of AGB-condensed carbon dust into the ISM offers a natural explanation for the paucity of PAHs in low-metallicity environments.  From a cosmological point of view, the large luminosity in these IR emission features coupled with the high sensitivity of \\spitz\\ has allowed the detection of PAHs in distant luminous infrared galaxies out to redshift $z\\simeq2$ \\cite[{e.g.}][]{elbaz05,yan05,houck05}. Hence, understanding what controls the properties of the aromatic features on large scales is required, in order to properly interpret broadband surveys.  The detailed characteristics of the mid-IR features, such as their shape, their central wavelength or the intensity ratio between the different bands are known to vary \\citep[see][]{peeters04b}. Such variations are essentially due to modifications of the molecular structure of the PAHs, in different astrophysical environments. In particular, since each feature is attributed to a given vibrational mode, the ratio between these features will vary with quantities such as the charge, the hydrogenation or the size and shape of the molecule. The $3.3\\mic$ PAH band arises from the radiative relaxation of CH stretching modes, while the $11.3\\mic$ feature originates in the CH out-of-plane bending modes; CC stretching modes are responsible for the features between 6 and $9\\mic$; CH in plane bending excitation produces part of the $8.6\\mic$ band. Now, laboratory studies and quantum calculations shows that the CC modes are instrinsically weak in neutral PAHs, and become stronger when the PAHs are ionized \\citep{langhoff96,allamandola99,bauschlicher02,kim02}. Therefore, the 6 to $9\\mic$ bands will be much more intense for a PAH$^+$ than for a PAH$^0$, while it will be the opposite for the 3.3 and $11.3\\mic$. Consequently, the ratios between the CC and the CH feature intensities depend on the charge of the PAHs, which is directly related to the physical conditions (e.g.\\ intensity of the ionizing radiation field, electron density, etc.) in the environment where the emission is originating.  Evidence of variations between features in different astrophysical environments have been reported by many authors. For example, \\citet{joblin96} showed that the $\\ipah{8.6}/\\ipah{11.3}$ ratio decreases with increasing distance from the exciting star, in the reflection nebulae \\ngc{1333} --~where $\\ipah{\\lambda}$ is the integrated intensity of the feature centered at $\\lambda\\mic$. \\citet{hony01} found a good correlation between the 3.3 and $11.3\\mic$ CH bands, in a sample of Galactic \\hii\\ regions, YSOs, and evolved stars, while they reported variations of $\\ipah{6.2}/\\ipah{11.3}$ by a factor of 5. The observations of Galactic and Magellanic \\hii\\ regions, presented by \\citet{vermeij02}, indicate that the ratios $\\ipah{6.2}/\\ipah{11.3}$, $\\ipah{7.7}/\\ipah{11.3}$ and $\\ipah{8.6}/\\ipah{11.3}$ are correlated. Furthermore, they suggest a segregation between the values of these ratios in the Milky Way and those in the Magellanic Clouds. \\citet{bregman05} studied the variation of $\\ipah{7.7}/\\ipah{11.3}$ in three reflection nebulae. Assuming that this variation is controlled by the charge of the PAHs, they could relate this band ratio to the ratio $G_0/n_e$ between the integrated intensity of the UV field, $G_0$, and the electron density, $n_e$. Similarly, \\citet{compiegne07}, studying the detailed variations of the mid-IR spectrum in the Horsehead nebula, attributed the high relative strength of the $\\ipah{11.3}$ feature to a high fraction of neutral PAHs, due to the high ambient electron density. On the contrary, \\citet{smith07} studied the variation of $\\ipah{7.7}/\\ipah{11.3}$ coming from the nuclear regions of the SINGS legacy program galaxies. They find that this ratio is relatively constant among pure starbursts, but varies by a factor of 5 among galaxies having a weak AGN. They interpret this effect as a selective destruction of the smallest PAHs by the hard radiation arising from the accretion disk, ruling out the explanation in terms of ionization of the molecules, in these particular environments. This interpretation is also supported by the high $11.3\\mic$ and the weakness of the $3.3\\mic$ band in the \\akari\\ spectrum of the giant elliptical galaxy \\ngc{1316} \\citep{kaneda07}.  The previous considerations stress the diversity of the possible interpretation of the mid-IR feature variations in galaxies. We need to identify the main physical processes controling the PAH bands, if we are to use them as diagnostic tools. This is the aim of this paper. It presents a quantitative analysis of mid-IR spectra (\\iso\\ and \\spitz) of Galactic regions, low-metallicity dwarf galaxies, quiescent spirals, starbursts and AGNs. To achieve our goal, we focus on identifying the main trends between the various PAH features, at different spatial scales and in different environments. We then link these variations to the physical conditions inside the studied region. Preliminary results of this study were published in \\citet{galliano04,galliano07c}.  The paper is organized as follows. In \\refS{sec:obs}, we present our sample and the data reduction. We discuss the spectral decomposition of the mid-IR spectrum, in \\refS{sec:decomp}. The results of this decomposition are presented in \\refS{sec:correl}; we study the various trends between the band ratios within galaxies, and among different types of environment. Then, in \\refS{sec:explanation}, we provide a physical interpretation of these trends, when the structure of the ISM is resolved, and when it is not. Finally, we summarize our conclusions in \\refS{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  We presented the results of a systematic study aimed at understanding the main properties of the mid-IR features in different astrophysical environments. It is based on observations of Galactic regions and galaxies of various types, observed by the satellites \\iso\\ and \\spitz. We have developed two distinct methods of spectral decomposition with different hypotheses, in order to test the robustness of our trends, and overcome eventual biases. These two methods have shown similar trends between the physical quantities that we have studied. We explored the variations of the different mid-IR features between the wavelengths 5 and $16\\mic$, among integrated spectra of galaxies, and inside galaxies and Galactic regions. Our main results are the followings. \\begin{enumerate} \\item We find that the 6.2, 7.7, $8.6\\mic$ features are essentially tied together, while the ratio of these bands to the $11.3\\mic$ feature can vary by one order of magnitude, in our sample. These variations are seen both inside individual sources (like \\M{82}, \\M{17}, the \\orb, \\xxxdor, etc.), as well as among integrated spectra. In general, the $\\ipah{6.2}/\\ipah{11.3}$ ratio is spatially correlated with the power radiated by the PAHs. It indicates that the ratio $\\ipah{6.2}/\\ipah{11.3}$ is higher in regions of intense star formation. \\item With the help of a stochastic heating model and realistic absorption efficiencies, we show that the variations of the mid-IR features are essentially due to the variation of the fraction of ionized PAHs. We conclude that the properties of the PAHs, throughout our sample, are remarkably universal. In particular, we rule out both the modification of the grain size distribution, and the extinction by the $9.7\\mic$ silicate feature, as an explanation of these variations. Indeed, we show that a modification of the size distribution could explain the observed variation of the $\\ipah{6.2}/\\ipah{11.3}$ ratio. However, it would cause the $\\ipah{6.2}/\\ipah{7.7}$ and $\\ipah{6.2}/\\ipah{8.6}$ (and $\\ipah{3.3}/\\ipah{11.3}$) ratios to vary significantly and to be correlated with the $\\ipah{6.2}/\\ipah{11.3}$, contradicting our observations. Similarly, a deep absorption by the silicate feature would vary the $\\ipah{6.2}/\\ipah{11.3}$ ratio, but it would decouple the $8.6\\mic$ from the 6.2 and $7.7\\mic$ features at the same time. \\item The universality of the properties of the PAHs and the fact that the band ratios are mainly sensitive to the charge of the molecules allow us to use these band ratios as a tracer of the physical conditions inside the emitting region. Using a few well-studied Galactic regions (including the spectral image of the \\orb), we give an empirical relation between $\\ipah{6.2}/\\ipah{11.3}$ and the ratio $G_0/n_e\\times\\sqrt{T_\\sms{gas}}$. \\item In the case where several regions with different physical conditions are integrated, the band ratios are dependent on the morphology of the ISM, as well as on the evolutionary properties of the illuminating stellar cluster. Thus, we find that the $\\ipah{6.2}/\\ipah{11.3}$ band ratios in the star forming nuclei of \\M{51} and \\M{82} are very similar to those in the galactic compact \\hii\\ regions, \\IR{15384} and \\IR{18317}, while the halo of \\M{82} and the spiral arms of \\M{51} are similar to those in the Galactic reflection nebula \\ngc{2023}. These differences in band ratio reflect differences in the ionization over recombination rate and hence trace back to variations in the ratio of the ionizing radiation field to the electron density. \\end{enumerate}     \\appendix"
    },
    "0801/0801.1431_arXiv.txt": {
        "abstract": "Frequencies, powers and damping rates of the solar p modes are all observed to vary over the 11-yr solar activity cycle. Here, we show that simultaneous variations of these parameters give rise to a subtle cross-talk effect, which we call the ``devil in the detail'', that biases p-mode frequencies estimated from analysis of long power frequency spectra. We also show that the resonant peaks observed in the power frequency spectra show small distortions due to the effect. Most of our paper is devoted to a study of the effect for Sun-as-a-star observations of the low-$l$ p modes. We show that for these data the significance of the effect is marginal. We also touch briefly on the likely $l$ dependence of the effect, and discuss the implications of these results for solar structure inversions. ",
        "introduction": "\\label{sec:intro}  Long timebase observations of the ``Sun as a star'' made, for example, by the ground-based BiSON network (Chaplin et al. 2007) and the GOLF (Gabriel et al. 1995) and VIRGO/SPM (Fr\\\"ohlich et al. 1997) instruments on the \\emph{ESA/NASA} SOHO spacecraft, provide key data on the low-degree (low-$l$) p modes for probing the deep solar interior and core. This paper is concerned chiefly with aspects of analysis of these Sun-as-a-star data.  Accurate and precise p-mode parameter data are a vital prerequisite for making robust inference on the interior structure of the Sun, and improved precision, and usually also improved accuracy, accrue from use of long helioseismic datasets (e.g., see discussion in Chaplin \\& Basu 2007). However, long helioseismic datasets must by necessity span a sizable fraction, or more, of an 11-yr solar activity cycle. Mode frequencies, damping rates, powers and peak asymmetries (Jim\\'enez-Reyes et al 2007) all vary systematically with the solar cycle, and these variations can give rise to complications for the analysis and interpretation, with implications for parameter accuracies, when long datasets are analyzed in one piece.  There is the concern that changes in frequency over a long dataset may distort the underlying shapes of the mode peaks to such an extent that the shapes no longer match the Lorentzian-like functions that are used to model them. This `distortion' effect turns out not to be a significant cause for concern, at least where the Sun is concerned (Chaplin et al., 2008), although it may be an issue for stars that have shorter, and more pronounced, activity cycles than the Sun. There is, however, another potential problem for the solar observations, a `cross-talk' effect that has its origins in the simultaneous variations of the mode frequencies, powers and linewidths over the solar cycle. This cross-talk effect can bias the estimated frequencies of the modes. The effect, which we call the ``devil-in-the-detail'', is the subject of this paper.  To introduce the effect, let us begin with a fictitious scenario, in which the p-mode frequencies are the only mode parameters that are observed to vary over the solar cycle. A mode peak observed in the power frequency spectrum of a long dataset will then have a centroid frequency that corresponds to the unweighted average of its time-varying frequency, over the period of observation. This simple rule of thumb holds because the frequency shifts of the low-$l$ solar p modes are small, both in fractional terms and compared to the peak linewidths. In sum: the unweighted time-averaged frequency may be directly recovered.  Now consider what happens if the power of the mode is also observed to vary in time. This will mean that timeseries data at different epochs will have varying contributions to the final peak profile observed in the power frequency spectrum. Data samples from times when the mode is more prominent (i.e., at high power) will carry proportionately larger weights in determining the final peak profile than samples from times when the mode is less prominent (i.e., at low power). Because the frequency also varies in time, one would expect the mode peak to be pulled, or biased, to a lower or higher frequency than the unweighted average frequency, with the size and sign of the effect dependent on the comparative time variations of the frequency and power. This is our devil in the detail.  Why might the devil in the detail matter? First, it raises uncertainty over what it is we are actually measuring when we determine mode frequencies from long datasets. If we have two sets of data that are not contemporaneous, it creates uncertainty if we wish to correct their frequencies to a common level of activity. Correction procedures, which attempt to `remove' the solar-cycle effects from the fitted frequencies, can be quite sophisticated (e.g., the BiSON correction procedure; see Basu et al. 2007); but they assume that the observed mode frequencies correspond to unweighted time averages of the frequencies. The devil-in-the-detail effect means this is not true in practice. Furthermore, because the frequency bias that is introduced depends on the relative sizes of the parameter variations -- the larger the fractional power variation, for a given frequency variation, the more pronounced the effect -- we might also expect there to be some dependence of the bias on the angular degree, $l$. Both issues above have potential implications for structure inversions made with the estimated p-mode frequencies.  The layout of the rest of the paper is as follows. In Section~\\ref{sec:prob}, we spell out why cross-talk from simultaneous variations of the mode parameters introduces the devil-in-the-detail bias in estimation of frequencies from Sun-as-a-star data. We then describe how artificial data were used to illustrate, and quantify, the problem.  We give some important background on the Sun-as-a-star observations in Section~\\ref{sec:sas}, and details on the artificial Sun-as-a-star data in Section~\\ref{sec:data}. We then discuss results from analysis of long (9.5-yr) artificial datasets in Section~\\ref{sec:full}, where we attempt to quantify the likely size of the devil-in-the-detail effect for real observations.  Finally, we pull together the main points of the paper in Section~\\ref{sec:conc}, where we discuss the implications for analysis of real data. We also touch on the issue of the $l$ dependence of the effect, and the possible impact on results of structure inversions. ",
        "conclusions": "\\label{sec:conc}  Over the course of the solar cycle, variations are observed not only in the p-mode frequencies, but also in the p-mode powers. Variation of the powers over time gives rise to a cross-talk effect with the varying frequencies, which we call the devil-in-the-detail effect.  The effect, and its impact, may be summarized as follows. In the absence of any power variation, the location in the power frequency spectrum of the centroid of a mode peak would in principle give an accurate measure of the unweighted average mode frequency over the period of observation. However, when significant time variation of the power is present, data samples from times when the mode is more prominent (i.e., at high power) will carry proportionately larger weights in determining the final peak profile than samples from times when the mode is less prominent (i.e., at low power). As the level of solar activity rises, the frequencies of the low-$l$ p modes under study here increase, but corresponding mode powers decrease. This means that, for a long timeseries, the frequency of a mode peak will be biased so it is \\emph{lower} than the unweighted time-averaged frequency. This is our so-called devil in the detail.  \\subsection{Impact on Sun-as-a-star data} \\label{sec:concsas}  We have used extensive Monte Carlo simulations to estimate the magnitude of the resulting devil-in-the-detail bias, in BiSON- and GOLF-like Sun-as-a-star helioseismic data spanning almost one full activity cycle. Our simulations imply the bias should rise monotonically in frequency up to $\\approx 3000\\,\\rm \\mu Hz$, and then level off, or even decrease slightly in magnitude, at higher frequencies.  It is at approximately the same frequency that the magnitude of the bias reaches a size that is comparable to the expected frequency uncertainties. Should we be worried by this devil-in-the-detail effect? Let us consider briefly, in this and the next section, two issues where the effect might be a potential cause for concern.  There is the impact of the effect on attempts to correct the frequencies for solar activity. Since our in-depth analysis has been made for Sun-as-a-star data on the low-$l$ p modes, we are in a position to comment in detail on correction procedures for these data. Procedures which attempt to `remove' the solar-cycle effects from the fitted frequencies (e.g., the BiSON correction procedure; see Basu et al. 2007) assume that the observed mode frequencies correspond to unweighted time averages of the frequencies. The devil-in-the-detail effect therefore introduces a bias in the correction, which will be comparable in size to the frequency uncertainties for modes at frequencies above $\\approx 3000\\,\\rm \\mu Hz$. Will this also be true for analysis of resolved-Sun data on modes of higher $l$?   \\subsection{Brief discussion on $l$ dependence of effect} \\label{sec:concl}  The issue of the $l$ dependence of the effect is of course important where inversions for the internal structure are concerned, since those inversions must make use of data on modes covering a wide range in $l$. It is a potential complication not just for analysis of non-contemporaneous data from two instruments (or more); but also for analysis of data from a single instrument, since significant variation of the bias with $l$ could affect the accuracy of inversions. The impact of the bias on analysis of medium and high-$l$ data, from resolved-Sun observations, will be covered in detail in a future paper. Here, in order to get a feel for the likely size of the bias, its variation with $l$, and the possible impact on structure inversions, we conducted the following experiment.  First, we estimated the size of the devil-in-the-detail bias for different $l$ in the range $0 \\le l \\le 150$, but only at fixed frequency (here, a notional mode frequency of $3000\\,\\rm \\mu Hz$). We assumed the frequency shifts should scale inversely with the mode inertias (and used the model `S' inertias; see Christensen-Dalsgaard et al. 1996); and we estimated the $l$-dependent mode height and linewidth changes from GONG data analyzed in Komm, Howe \\& Hill (2000) [their Figs.~10 and 11). With estimates of the solar-cycle frequency, height and linewidth changes to hand for each $l$, we were in a position to use Equation~\\ref{eq:int} to construct the expected final peak profiles -- assuming a timeseries of length 9.5\\,yr (as in Sections~\\ref{sec:data} and~\\ref{sec:full}) -- and to thereby estimate the devil-in-the-detail bias as a function of $l$. Our results showed little variation of the estimated bias with $l$ below $l \\la 60$ (it being $\\approx 0.02\\,\\rm \\mu Hz$ in size). However, at higher degrees the bias increased in a linear manner with increasing $l$, reaching a size of $\\approx 0.05\\,\\rm \\mu Hz$ at $l \\sim 150$.  Next, we scaled the estimated bias for each $l$ by the inertia ratio of Christensen-Dalsgaard \\& Berthomieu (1991), and compared these scaled differences with sizes of scaled frequency differences that are typically large enough to affect inversion results for the internal structure at the $1\\sigma$ level (see also Eff~Darwich et al. 2002). This assumed that the observed p-mode frequencies came from the MDI frequency set of Korzennik (2005), which includes frequencies on modes from $0 \\le l \\le 125$. The results of this comparison suggest that the devil-in-the-detail bias is not quite large enough to affect the inversions at the $1\\sigma$ level, although the amounts by which it fell short were marginal.  So, our preliminary conclusion, regarding the impact of the $l$ dependence, is that the devil in the detail may not be a cause for concern where inversions for structure using data from a single instrument are concerned. That said, we should stress that this comparative analysis was an approximate one. At no point did we account for any possible $m$ dependence, which might have a significant influence on the results. Without having done a full Monte Carlo simulation of the data, we were not in a position to provide an accurate estimate of the frequency dependence for different $l$. These factors will be tested in upcoming work.   \\subsection*"
    },
    "0801/0801.0810_arXiv.txt": {
        "abstract": "{We investigate early time inflationary scenarios in an Universe filled with a dilute noncommutative bosonic gas at high temperature. A noncommutative bosonic gas is a gas composed of bosonic scalar field with noncommutative field space on a commutative spacetime. Such noncommutative field theories was recently introduced as a generalization of quantum mechanics on a noncommutative spacetime. As key features of these theories are Lorentz invariance violation and CPT violation. In the present study we use a noncommutative bosonic field theory that besides the noncommutative parameter $\\theta$ shows up a further parameter $\\sigma$. This parameter $\\sigma$ controls the range of the noncommutativity and acts as a regulator for the theory. Both parameters play a key role in the modified dispersion relations of the noncommutative bosonic field, leading to possible striking consequences for phenomenology. In this work we obtain an equation of state $p=\\omega(\\sigma,\\theta;\\beta)\\rho$ for the noncommutative bosonic gas relating pressure $p$ and energy density $\\rho$, in the limit of high temperature. We analyse possible behaviours for this gas parameters $\\sigma$, $\\theta$ and $\\beta$, so that $-1\\leq\\omega<-1/3$, which is the region where the Universe enters an accelerated phase. \\\\  \\vspace{1cm}  Keywords: Noncommutative Fields, Lorentz Invariance Violation, CPT Violations, Inflationary Cosmology} ",
        "introduction": "Quantum Field Theories (QFT) with Lorentz invariance violation are subject of growing attention among theoretical physicists in the last few years \\cite{Kostelecky:1988zi,Kostelecky:1995qk,Colladay:1998fq,Kostelecky:1999mr,Jackiw:1999yp,Kostelecky:2000mm}. Indeed, there is no reason a priori to believe that the cosmological principle\\footnote{The cosmological principle states that spacetime is locally Lorentz invariant.} is true at all energy scales. Further analysis on the problem of matter-antimatter asymmetry \\cite{Dine:2003ax}, ultra high energy cosmic rays \\cite{Anchordoqui:2002hs,Takeda:2002at,Abraham:2006ar}, primordial magnetic field \\cite{Grasso:2000wj}, neutrino physics \\cite{Athanassopoulos:1997pv} and some other cosmological measurements demands further criticisms on the cosmology principle at ultra high energy scale. Therefore the proposal of validity/bound tests for the cosmology principle, i.e., cosmological phenomenology based upon QFT with Lorentz invariance violation is a subject of major relevance.  There are distinct approaches to QFT with Lorentz invariance violation. A first approach is based on the noncommutativity of spacetime itself. The most studied situation was that of QFTs on the Groenenwold-Moyal \\cite{Szabo:2001kg,Douglas:2001ba}. Some of these models may be obtained from a fundamental theory such as string/M-theory \\cite{Douglas:2001ba}. Recently \\cite{Chaichian:2004za,Aschieri:2005yw,Aschieri:2005zs}, it was shown that one could restore Lorentz invariance of these QFT by a proper twisting of the action of the symmetry group on the fields. Cosmology applications of this twisted Lorentz invariant QFT is recently being under study \\cite{Akofor:2007fv}. There is also an interesting approach based on using a noncommutative version of Wheeler-deWitt equation, where the noncommutativity is introduced as a Moyal deformation \\cite{GarciaCompean:2001wy}.   A second approach is known as extended standard model \\cite{Colladay:1998fq}. It is based on a famous work of S. Carroll, R. Jackiw and G. Field \\cite{Carroll:1989vb}. In this approach the breaking of Lorentz invariance is performed by the introduction of a Chern-Simons term on the action of some gauge theory\\footnote{It is worth mentioning that there is some controversy on the correctness of this approach \\cite{Bonneau:2006ma}}. One should also mention the approach known as doubly special relativity \\cite{Magueijo:2001cr}.  Another approach, the one we are going to consider in this paper, is QFT based on noncommutative fields \\cite{Carmona:2002iv,Carmona:2003kh,Balachandran:2007ua}. On such theories one breaks Lorentz invariance by modifying the canonical commutation relations of the fields (therefore the name noncommutative fields) by the inclusion of an ultraviolet or/and infrared scale factor. These noncommutative fields arise as a generalization of the quantum mechanics on the Groenenwold-Moyal plane \\cite{Nair:2000ii}, where one considers the noncommutativity in the degrees of freedom of the fields. One interesting feature of such theories is that there is a connection of this approach with Carroll et al proposal \\cite{Carroll:1989vb} for gauge theories. See for instance \\cite{Carmona:2002iv,Gamboa:2005bf,Falomir:2005it,Gamboa:2005pd,Falomir:2006hp}.  One common feature of all these theories with Lorentz invariance broken is the deformation of the dispersion relations of the plane waves. It is this deformation that allows one to build phenomenological models that could be in principle be tested in future cosmology or particle physics experiments. The kind of deformation to be seen or ruled out by these experiments will thus decide which theory will survive.  The most promising arena for these phenomenological models is the inflationary scenario of the Universe. Models based on the Universe inflation has successfully explained several observational data, such as the origin of density fluctuations on large scales. The general belief is that inflation that supports large and fast growth of scales can connect the region where the Lorentz invariance might break (i.e., at very high energy) to cosmological scales. On the other hand, it is known that so far there is no explanation of the essential mechanisms of inflation based on fundamental physics such as string/M-theory.   In the context of inflationary scenarios based on QFT with Lorentz invariance violation, based on the noncommutativity of space-time, several developments have been put forward recently \\cite{Alexander:2001dr,Lizzi:2002ib,Hassan:2002qk,Koh:2007rx,Koh:2007wa,Chu:2000ww, Bertolami:2002eq}. In these works several efforts have been made in the direction of finding imprints of primordial fluctuations that can survive the inflationary phase. Such imprints have been investigated both in theories with \\cite{Chu:2000ww} and without \\cite{Koh:2007rx} inflaton fields. See \\cite{Brandenberger:2007rg} and references therein for a recent review on this subject.  In the present paper, we consider the approach based on noncommutative fields as studied by A. P. Balachandran et al. in \\cite{Balachandran:2007ua}. In that work a noncommutative free massless bosonic field was studied. For that theory, they considered a regularization of the delta function appearing in one of the commutators of the fields. The equal time commutation relation for the fields then reads \\begin{equation} \\label{eq:comm-rel-modified} \\left[ \\hat{\\varphi}^a(x;t),\\hat{\\varphi}^b(y;t)\\right]=i\\epsilon^{ab}\\theta(\\sigma;x-y), \\end{equation} where $\\theta(\\sigma;x-y)$ is a Gaussian distribution with a parameter $\\sigma$ related with the standard deviation of the distribution. Besides introducing one further parameter $\\sigma$, the resulting theory is finite. That fact allowed the study of the deformed black body radiation spectrum. One interesting feature of this deformed spectrum is that the energy density for boson with higher frequency is greater than that for the usual black body radiation spectrum.  We investigate here an inflationary model based on an Universe filled with a similar gas composed of a noncommutative boson in the relativistic limit. In this scenario we study the thermodynamics of such gas. We aim at computing and analyzing the equation of state $p=\\omega\\rho$, with $p$ being the pressure, $\\rho$ being the energy density. In the present situation this equation of state is far from trivial. Therefore we have to make some careful approximations, such as considering a power series of the parameters $\\theta$ and temperature $1/\\beta$. Now, for suitable values of the parameters $\\theta$, $\\sigma$ and the temperature $1/\\beta$ of the gas, one obtain $-1\\leq\\omega< -1/3$. For these values of $\\omega$ the Universe presents an accelerated phase. Furthermore, the present case should be contrasted with the tachyon and Chaplygin cosmology that is governed by an equation of state given by the general form $p=-A/\\rho$ \\cite{Gibbons:2003gb}, since here the inflationary regime is followed by a radiation dominated epoch for sufficiently small density $\\rho$.  The present paper is organized as follows: in Sec.~\\ref{prel} we review the formulation of the free noncommutative boson field as presented in Balachandran et al. The goal is to obtain the Hamiltonian for such noncommutative field so that we may formulate a statistical quantum mechanics problem for the gas at high temperature; in Sec.~\\ref{stme} we study the gas of noncommutative bosons in the grand canonical formalism. Then by some suitable approximations, we obtain an equation of state $p=\\omega(\\sigma,\\theta;\\beta)\\rho$ for this gas; in Sec.~\\ref{cosmonq} we consider a cosmology model with the noncommutative gas by presenting the proper Friedman equation and then the behavior of the energy density of this gas with the scale parameter. Next we consider, by fine tuning the parameters, possible scenarios yielding inflation; finally in Sec.~\\ref{conclu} we finish the paper with some concluding remarks. ",
        "conclusions": "\\label{conclu}  In this paper we have used noncommutative fields to find cosmological scenarios that can develop inflation at early Universe. A gas composed with a noncommutative bosonic field shows up several interesting properties. In particular, in the dilute regime at high temperature, these properties can be read off from Eq. (\\ref{omega}).  The aim of the paper was essentially to find the equation of state of the gas by tackling approximately the statistical quantum mechanics problem for the noncommutative field gas at high temperature. Since we are dealing with highly relativistic bosons not necessarily massless, the approximate relation between energy and momenta $E\\simeq p$ used along the paper is valid. This in turn is consistent with the formulation of free massless boson with noncommutative target space adopted previously. The grand canonical ensemble used in the quantum statistical formulation is also valid because, on the other hand, the masslessness of the bosons is just an approximation.  The noncommutative field gas we are dealing with here can develop inflation for a period of time that can be adjusted to observational data. We showed some scenarios that exit inflation phase to a radiation dominated epoch, as it should be according to the perspective of the inflationary cosmology. This feature is certainly a motivation for further studies of these scenarios. In particular, the scenarios with $\\theta$ fixed.  In scenarios with $\\sigma$ and $\\beta$ fixed, we are left with only the parameter $\\theta$ free to be determined for a given temperature in Eq. (\\ref{omega_min}). Let us estimate the value of $\\theta$ for $T=10^{14}GeV\\simeq10^{27} K$, the approximate temperature at beginning of the inflation, for $\\omega_{min}\\geq-1$ and $\\omega_{min}<-1/3$, respectively \\ben \\theta_{-1}\\geq 0.4\\times 10^{-25}cm, \\qquad \\theta_{-1/3}<0.5\\times 10^{-25}cm. \\een This result should be contrasted with other astrophysical bounds on the parameter $\\theta.$ For instance, in \\cite{Carmona:2002iv}, it was estimated a value for $\\theta$ to be of the order of $(10^{20} eV)^{-1}\\simeq10^{-25}cm$. They have used the fact that the highest cosmic ray energy be no greater than $\\sim 10^{20} eV$. Recall that the GZK cutoff is of the order of $10^{19} eV$. Furthermore, as we can easily see above, there is a short range for $\\theta$ that allows for inflation in these scenarios.  A proper understanding of the nature of the parameter $\\sigma$, for instance, a dynamical theory for it, could shed more light on these developments. Further possibilities may arise if one consider other kinds of regularization for the delta function in (\\ref{eqNCTA:69}).  Several other interesting issues such as the implications of the sound speed and the density perturbations can also be investigated to know in what extent the model is in accordance with CMB data. One can estimate the adiabatic sound speed by using \\ben c_s^2=\\left(\\frac{\\partial P}{\\partial \\rho}\\right)_{S,N}=-\\frac{V^2}{m_0N}\\left(\\frac{\\partial P}{\\partial V}\\right)_{S,N}, \\een with $\\rho=m/V=m_0N/V$. Once we have previously assumed a dilute gas $(z\\ll1)$ at high temperature, then by using eqs.(\\ref{therm_quantity2})-(\\ref{log_grand_approx}) one can easily show that the gas satisfies $PV=NT$. The sound speed is then given by \\ben c_s=\\sqrt{\\frac{T}{m_0}}, \\een where $m_0$ is the mass of ultra relativistic bosons satisfying $E^2=m_0^2+p^2\\simeq p^2$. So, at first order in $z$, the approximation maintained throughout this work, the sound speed does not depend on non-commutative parameters. The first correction comes from higher orders in $z$ expansion which can be computed quite easily to be: \\begin{equation} \\frac{T\\,\\left( 1 - \\frac{7\\,z}{64} \\right) }{m_0} + \\frac{91\\,z\\,{\\theta }^2}{16384\\,m_0\\,{\\pi }^2\\,T^3\\,{\\sigma }^6} + \\frac{z\\,\\log (z)}{64\\,m_0\\,T}. \\end{equation}  Since in this work we consider the high temperature limit, in which we can disregard the mass, further analysis on the speed of sound cannot be pursued. We let further studies related to this issue for future investigations, together with the low temperature limit."
    },
    "0801/0801.1816.txt": {
        "abstract": "We have measured the bias of QSOs as a function of QSO luminosity at fixed redshift ($z<1$) by cross-correlating them with Luminous Red Galaxies (LRGs) in the same spatial volume, hence breaking the degeneracy between QSO luminosity and redshift. We use three QSO samples from 2SLAQ, 2QZ and SDSS covering a QSO absolute magnitude range, $-24.5<M_{b_J}<-21.5$, and cross-correlate them with 2SLAQ ($z\\approx0.5$) and AAOmega ($z\\approx0.7$) photometric and spectroscopic LRGs in the same redshift ranges. The spectroscopic QSO samples, in the spectroscopic LRG areas, contain 300-700 QSOs and in photometric LRG areas up to 7000 QSOs. The 2-D and 3-D cross-clustering measurements are generally in good agreement. \u00a0Our (2SLAQ) QSO-LRG clustering amplitude ($r_0=6.8_{-0.3}^{+0.1}$h$^{-1}$Mpc) as measured from the semi-projected cross-correlation function appears similar to the (2SLAQ) LRG-LRG auto-correlation amplitude ($r_0=7.45\\pm0.35$h$^{-1}$Mpc) and both are higher than the (2QZ+2SLAQ) QSO-QSO amplitude ($r_0\\simeq5.0$h$^{-1}$Mpc). Our measurements show remarkably little QSO-LRG cross-clustering dependence on QSO luminosity. If anything, the results imply that brighter QSOs may be less highly biased than faint QSOs, the opposite direction expected from simple high peaks biasing models, where \u00a0more luminous QSOs are assumed to occupy rarer high mass peaks. Assuming a standard $\\Lambda$CDM model and values for $b_{LRG}$ measured from LRG autocorrelation analyses, we find $b_Q=1.45\\pm0.11$ at $M_{b_J}\\approx-24$ and $b_Q=1.90\\pm0.16$ at $M_{b_J}\\approx-22$.  We also find consistent results for the QSO bias from a $z-$space distortion analysis of the QSO-LRG cross-clustering at $z\\approx0.55$. The velocity dispersions fitted to QSO-LRG cross-correlation, $\\xi(\\sigma,\\pi)$, at $\\pm680$ kms$^{-1}$ are intermediate between those for QSO-QSO and LRG-LRG clustering, as expected given the larger QSO redshift errors. The dynamical infall results give $\\beta _Q=0.55\\pm0.10$, implying $b_Q=1.4\\pm0.2$. Thus both the $z-$space distortion and the amplitude analyses yield $b_Q\\approx1.5$ at $M_{b_J}\\approx-23$. The implied dark matter halo mass inhabited by QSOs at $z\\approx0.55$ is $\\sim10^{13}h^{-1}M_{\\sun}$, again approximately independent of QSO luminosity. ",
        "introduction": "Useful constraints on the nature of QSOs can be drawn from even the simplest measure of QSO clustering, the amplitude of the real-space 2-point correlation function. For example, measuring QSO clustering on small scales can constrain models of galaxy mergers and quasar formation (Myers et al. 2007). Furthermore, the redshift evolution and luminosity dependence of the QSO bias can be studied. Recent work on the clustering dependence of the QSO bias on redshift suggested evolution of the 2QZ QSO bias (Croom et al. 2005). However, the fact that the most luminous QSOs lie at high redshifts, while the faintest are at low redshifts, made it difficult to study how QSO properties depend on luminosity. Now, surveys such as 2SLAQ (Cannon et al. 2006) that span a wide range of QSO luminosities, have broken that redshift-luminosity degeneracy and revealed little QSO clustering dependence on luminosity, at fixed redshift (da $\\hat{A}$ngela et al. 2006). Moreover, the amplitude of the QSO clustering is correlated with the average mass of the halos associated with the QSOs, providing indications of QSO lifetimes and making it possible to constrain QSO evolutionary models (Croom et al. 2005, da $\\hat{A}$ngela et al. 2008).  Redshift space distortions of the clustering pattern also contain dynamical information on the QSO bias that are independent of any assumption about underlying mass clustering. The clustering is affected at small scales by the rms velocity dispersion of QSOs along the line-of-sight (Fingers of god) and by dynamical infall of matter into higher density regions, which causes a flattening of the clustering pattern in the redshift direction. In addition to  these dynamical distortions, geometrical distortions are introduced if an incorrect cosmological model is used in order to convert the observed redshifts into comoving distances. They therefore also constrain, more weakly, the value of the cosmological density parameter, $\\Omega_m$.  In the linear regime of clustering, dynamical infall is governed by the parameter $\\beta =\\Omega_m^{0.6}/b$. (Peebles, 1980, Kaiser 1987, Loveday et al. 1996, Matsubara \\& Suto 1996, Matsubara \\& Szalay 2001, Peacock et al. 2001, Hoyle et al. 2002, Coil et al. 2005, Myers et al. 2006, 2007, Porciani \\& Norberg 2006, da $\\hat{A}$ngela et al. 2008, Ross et al. 2007). In recent years, measurements of QSO clustering (da $\\hat{A}$ngela et al. 2008) yielded a $\\beta_{QSO}(z=1.4)=0.60_{-0.11}^{+0.14}$ and $b_{QSO}(z=1.4)=1.5\\pm0.2$ for a combined sample of QSOs from the 2dF QSO Redshift Survey (2QZ, Croom et al. 2004) and the 2dF-SDSS LRG and QSO Survey (2SLAQ, Cannon et al. 2006). Ross et al. (2007) performed similar measurements on the 2SLAQ Luminous Red Galaxies (LRGs) clustering and found $\\beta_{LRG}(z=0.55)=0.45_{-0.05}^{+0.05}$ and $b_{LRG}(z=0.55)=1.66\\pm0.35$.  In this paper we use QSOs from the 2QZ, 2SLAQ and SDSS (York et al. 2000) Data Release 5 (DR5; Adelman-McCarthy 2007) surveys and LRGs from 2SLAQ and AAOmega first to  study the dependence of QSO-LRG clustering amplitude on QSO luminosity. We also measure QSO-LRG redshift distortions to estimate the  dynamical infall parameter, $\\beta$.  These surveys provide large numbers of QSOs and LRGs with a range of luminosities at fixed redshifts. So our results, for example on QSO bias, should be statistically improved over those from QSO-QSO clustering. Moreover, we use photometric LRG samples, which are significantly larger than the spectroscopic ones and measure QSO 2-D correlation function amplitudes, using Limber's formula to convert to 3-D real-space measurements.  In section 2 we describe the QSO and LRG samples we use in our measurements. In section 3 we present results from the 2-point cross-correlation function $\\omega (\\theta)$, in Sections 4 and 5 we show our results for the redshift-space cross-correlation function, $\\xi _s$, and the semi-projected cross-correlation function, $w_p(\\sigma)/\\sigma$, respectively. Section 6 has our results for the real-space cross-correlation function and in Section 7 we present our $\\xi (\\sigma, \\pi )$ results, model the redshift-space distortions and estimate the QSO infall mass and bias. We also study, in Section 8, the dependence of the redshift-space cross-correlation function, the QSO bias and the mass of the QSO-LRG Dark Matter Haloes ($M_{DMH}$) on QSO luminosities at fixed redshifts. Finally, in Section 9 we present our conclusions. ",
        "conclusions": "In this paper we have performed an analysis of the clustering of QSOs with LRGs. For this purpose, we first used the 2-point angular cross-correlation function, $w(\\theta)$, and measured the cross-correlation between 2SLAQ and AAOmega LRGs and different luminosity QSOs. The results show that there is little cross-correlation dependence on QSO luminosity.   Next, we measured the redshift-space cross-correlation function. We again see no QSO-LRG clustering dependence on QSO luminosity, as all the QSO-spectroscopic LRG samples gave similar results. We used Limber's formula to fit $r_0$ to 2-D results. The fits for $r_0$ from 3-D $\\xi (s)$ are in very good agreement with the fits to the 2-D $w(\\theta)$ results. Then, we compared our QSO-LRG clustering with 2SLAQ LRG-LRG (Ross et al. 2007) and 2QZ+2SLAQ QSO-LRG (da $\\hat{A}$ngela et al. 2008) clustering results. On small scales, the QSO-QSO and QSO-LRG results appear flatter than the LRG-LRG results. As confirmed later by the $w_p(\\sigma )/\\sigma$ measurements, this appears to be due to the larger QSO redshift errors (broad-lines). On larger scales ($\\geq 5$h$^{-1}$Mpc) the QSO-QSO correlation amplitude appears lower than the QSO-LRG and LRG-LRG amplitudes, suggesting that $b_Q\\lesssim b_L$. The fractional errors on $\\xi _{QL}$ are $\\sim 50\\%$ smaller than those on $\\xi _{QQ}$ in same redshift range.   The results from the semi-projected cross-correlation function, $w_p(\\sigma )/\\sigma$, are in agreement with our $\\xi (s)$ observations, yielding consistent $r_0$ and $\\gamma$ values. The comparison with the 2SLAQ LRG-LRG (Ross et al. 2007) and 2QZ+2SLAQ QSO-LRG (da $\\hat{A}$ngela et al. 2008) confirms our results from redshift-space, i.e. that the small-scale amplitude difference in $\\xi (s)$ is due to the larger QSO redshift errors and that QSO-QSO clustering amplitude is lower than QSO-LRG and LRG-LRG amplitude. The real-space cross-correlation function, $\\xi (r)$, also seems to agree with the $\\xi (s)$ and $w_p(\\sigma )/\\sigma$ results, although it is noisier.   Then, we measured the $\\xi (\\sigma, \\pi)$ cross-correlation function for all our (spectroscopic) samples and used the results to model the redshift-space distortions. For that, we used two models which gave consistent results with each other and between the different samples. The redshift-space distortions yielded an average of $\\beta _Q=0.55\\pm0.10$, $b_Q=1.4\\pm0.2$ which is slightly higher than $b_Q=1.1\\pm0.2$ at $z\\simeq0.6$, from da $\\hat{A}$ngela et al. (2008). Note that this latter result does not come from analysing the amplitude of $\\xi _{QQ}$; there being too few QSO pairs to make redshift-distortion auto-correlation analysis viable at $z\\sim 0.6$.   After calculating the average absolute magnitude of each QSO sample we measured the integrated cross-correlation function for each one of our QSO-LRG samples. No evidence was found for QSO-LRG clustering dependence on QSO luminosity. Then, using LRG biases as estimated by previous studies, we calculated the QSO biases. There were indications of luminosity dependence, in the sense that $b_Q$ may reduce as we move to brighter QSO samples. Our analysis yielded a $b_Q=1.5\\pm0.1$ (at $M_{b_J}\\approx-23$) which is in very good agreement with our result from redshift-space distortions.  Finally, using the relation between bias and $M_{DMH}$ suggested by Sheth et al. 2001, we calculated the corresponding masses of the QSO hosts. QSO halo masses were estimated to be $\\sim10^{13}h^{-1}M_{\\sun}$ at $z\\approx0.55$. Our $M_{DMH}$ estimations are higher than those from Croom et al. (2005) and da $\\hat{A}$ngela et al. (2008) (at $z=1.4$) and may be explained by long-lived QSO population models. Since the bias values are independent of QSO luminosity at fixed redshift, the halo masses are also independent of luminosity and this represents the main result of this paper."
    },
    "0801/0801.2738_arXiv.txt": {
        "abstract": "We present a discussion of the partial Paschen-Back (PB) effect in magnetic Ap stars. An overview of the theory is illustrated with examples of how splittings deviate non-linearly from the simple Zeeman picture; normally forbidden ``ghost lines'' appear in strong fields. Resulting asymmetric stellar Stokes profiles for a dipolar magnetic geometry are shown for the Fe\\,{\\sc ii} $\\lambda\\,6149$ line and it is established that PB lines may be subject to wavelength shifts. Modelling of Stokes profiles in the PB regime opens exciting new diagnostics. ",
        "introduction": "\\label{intr} The discovery in 1897 by P.\\,Zeeman of the splitting of spectral lines in magnetic fields has provided astrophysicists with a valuable diagnostic tool. Following the pioneering work of G.E.\\,Hale (1908) on sunspots, Zeeman observations have been carried out on all kinds of magnetic structures found throughout the solar atmosphere, leading to a deeper understanding of the role of magnetic fields in atmospheric dynamics. The later discovery by Friedrich Paschen and Ernst Back (1921) of the transition in very strong magnetic fields from the anomalous Zeeman effect to the normal Zeeman effect has had much less impact. Indeed, even in fields of 0.4\\,T as encountered in sunspots, very few spectral lines of special astrophysical interest are affected by the Paschen-Back (PB) effect and so only a handful of atomic lines have been analysed for PB signatures (see e.g. Engvold {\\it et al.}, 1970; Socas Navarro {\\it et al.}, 2004).  When H.W.\\,Babcock ((1960) discovered a 3.4\\,T field in the upper main sequence chemically peculiar (Ap) star HD\\,215441, it should have been obvious that a lot of atomic transitions would find themselves somewhere between the Zeeman and the full PB regime (the partial or incomplete PB effect). However, not until much later was any thought spent on this problem (Kemic, 1975; Stift, 1977). A major discussion of the partial PB effect and a line profile modelling attempt is due to Mathys (1990); Landolfi {\\it et al.} (2001) had a look at the PB effect on hyperfine splitting and predicted how this would affect the interpretation of observations.  In strongly magnetic Ap stars, it turns out that high-resolution spectra of the Zeeman doublet of Fe\\,{\\sc ii} at 6149\\,{\\AA} cannot be fitted with synthetic line profiles calculated in the Zeeman approximation. Mathys (1990) attributed the observed non-symmetric relative intensities of the two components and their shift in wavelength (relative to the magnetic null lines of iron) to the partial PB effect: at 2\\,T the magnetic splitting attains about 35\\% of the distance between the nearest fine structure levels of the lower term. We decided to have a close look at this particular iron line which is heavily used in magnetic field measurements, to model it in detail for realistic stellar atmospheres and various magnetic geometries, and to search for other spectral lines affected by the partial PB effect. ",
        "conclusions": "\\label{out} A number of atomic lines observed in strongly magnetic Ap stars manifest asymmetries and line shifts that cannot be modelled in the Zeeman regime and have to be ascribed to the partial Paschen-Back effect. We can show that these are not just some exotic few because no less than 846 lines originate from a single Cr\\,{\\sc i} term with fine structure splittings $0.25 - 2.77$\\,cm$^{-1}$. Whenever magnetic splittings become of comparable size, a detailed PB treatment of the whole multiplet containing this term is warranted and we have illustrated this with the first ever realistic Stokes profile modelling of the Fe\\,{\\sc ii} $\\lambda\\,6149$ line. Expecting exciting new diagnostic possibilities thanks to such improved modelling, we have started further explorations in this fascinating new field."
    },
    "0801/0801.2162_arXiv.txt": {
        "abstract": "{ The technique of gravitational microlensing is currently unique in its ability to provide a sample of terrestrial exoplanets around both Galactic disk and bulge stars, allowing to measure their abundance and determine their distribution with respect to mass and orbital separation. Thus, valuable information for testing models of planet formation and orbital migration is gathered, constituting an important piece in the puzzle for the existence of life forms throughout the Universe. In order to achieve these goals in reasonable time, a well-coordinated effort involving a network of either 2m or 4 $\\times$ 1m telescopes at each site is required. It could lead to the first detection of an Earth-mass planet outside the Solar system, and even planets less massive than Earth could be discovered. From April 2008, ARTEMiS (Automated Robotic Terrestrial Exoplanet Microlensing Search) is planned to provide a platform for a three-step strategy of survey, follow-up, and anomaly monitoring. As an expert system embedded in eSTAR (e-Science Telescopes for Astronomical Research), ARTEMiS will give advice for follow-up based on a priority algorithm that selects targets to be observed in order to maximize the expected number of planet detections, and will also alert on deviations from ordinary microlensing light curves by means of the {\\sc SIGNALMEN} anomaly detector. While the use of the VOEvent (Virtual Observatory Event) protocol allows a direct interaction with the telescopes that are part of the HTN (Heterogeneous Telescope Networks) consortium, additional interfaces provide means of communication with all existing microlensing campaigns that rely on human observers. The success of discovering a planet by microlensing critically depends on the availability of a telescope in a suitable location at the right time, which can mean within 10~min. To encourage follow-up observations, microlensing campaigns are therefore releasing photometric data in real time. On ongoing planetary anomalies, world-wide efforts are being undertaken to make sure that sufficient data are obtained, since there is no second chance. Real-time modelling offers the opportunity of live discovery of extra-solar planets, thereby providing ``Science live to your home''.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.0956_arXiv.txt": {
        "abstract": "Ultrarelativistic electron-positron plasmas can be produced in high-intensity laser fields and play a role in various astrophysical situations. Their properties can be calculated using QED at finite temperature. Here we will use perturbative QED at finite temperature for calculating various important properties, such as the equation of state, dispersion relations of collective plasma modes of photons and electrons, Debye screening, damping rates, mean free paths, collision times, transport coefficients, and particle production rates, of ultrarelativistic electron-positron plasmas. In particular, we will focus on electron-positron plasmas produced with ultra-strong lasers. ",
        "introduction": "Ultrarelativistic plasmas, in which the thermal energy of the particles is much larger than their rest mass energy, were discussed first in the fifties in the context of astrophysics. They occur in the early Universe as well as stellar and galactic high-energy processes (Zel'dovich and Novikov, 1971; Zel'dovich and Novikov, 1983; Raffelt, 1996). For example, ultrarelativistic electron-positron plasmas (EEP) can be created either by high temperatures in supernovae explosions (see e.g. Hardy and Thoma (2001)) or strong magnetic fields in so-called magnetars (Beskin {\\it et al.}, 1993). The theoretical description of ultrarelativistic plasmas is based on transport theory (Silin, 1960) or thermal field theory (Tsytovich, 1961). The latter is an extension of quantum field theory, such as QED or QCD, in vacuum to finite temperature (or chemical potential) describing high-energy particle interactions in matter. For this purpose, two different approaches were developed, the imaginary time formalism (Matsubara, 1955; Kapusta, 1989) and the real time formalism (Landsmann and van Weert, 1987).  In the eighties a new type of ultrarelativistic plasmas became of interest, the quark-gluon plasma (QGP) (M\\\"uller, 1985). It corresponds to a new state of matter of deconfined quarks and gluons at extremely high temperatures above $k_BT=150$ MeV. The early Universe should have been in this state for the first few microseconds after the Big Bang. Furthermore, the QGP is expected to be created in relativistic heavy-ion collisions, in particular within the accelerator experiment RHIC (Gyulassy and McLerran, 2005). There a hot and dense fireball (``little bang'') of the size of an atomic nucleus is produced, which could be in the QGP phase for less than $10^{-22}$ s. Therefore the QGP in those experiments cannot be observed directly but its discovery relies on the comparison of theoretically predicted signatures with experimental data (``circumstantial evidence''). For understanding the properties of the QGP and for calculating signatures for its formation, QCD at finite temperature has been applied. Besides a non-perturbative method based on a numerical solution of the QCD equations (lattice QCD), perturbation theory at finite temperature has been used widely (Thoma, 1995a; Le Bellac, 1996). Lattice QCD allows only the calculation of static quantities, and hence not of most of the proposed signatures following from particle production. Perturbative QCD, on the other hand, has been used for computing static as well as dynamic quantities. However, its predictions are limited  by the fact that the QGP is a strongly-coupled plasma at temperatures achievable in the experiments. (At extremely high temperatures far above the transition temperature perturbative QCD should be reliable due to asymptotic freedom of QCD.) Nonetheless, a wide variety of properties and quantities of the QGP have been considered in this way. It started with the calculation of self-energies and dispersion relations in the high-temperature limit (Klimov, 1982; Weldon, 1982a; Weldon, 1982b). In the high-temperature limit the results differ from the QED results only by trivial factors containing the QCD degrees of freedom (color, flavor). In addition, it can be shown that the gluon self-energy or polarization tensor, which is directly related to the dielectric function (see e.g. Elze and Heinz (1989)), can be derived also within the transport (Vlasov) approach (Silin, 1960), since the high-temperature limit corresponds to the classical limit. Other quantities, however, such as damping or production rates and transport coefficients, require a quantum field theoretic treatment. In particular, a resummation of certain diagrams, the hard thermal loops (HTL), is needed for a consistent description, i.e. to obtain gauge invariant, infrared finite (or improved) results complete to leading order in the coupling constant (Braaten and Pisarski, 1990). Details about quantum field theoretic methods and their application to the physics of the QGP can be found in the review article by Thoma (1995a).  Recently a third possibility to study ultrarelativistic plasmas has been suggested: in extremely strong laser fields the creation of an ultrarelativistic EPP with temperatures around 10 MeV could be realized soon (Liang {\\it et al.}, 1998). For example two opposite, circular polarized laser pulses with a duration of 330 fs and an intensity of $7\\times 10^{21}$ W/cm$^2$ could hit a thin gold foil. In this way the target electrons can be heated up to about 10 MeV producing an ultrarelativistic EPP by pair creation. The positron density could be about $5\\times 10^{22}$ cm$^{-3}$ (Shen and Meyer-ter-Vehn, 2001). This will allow for the first time to study an ultrarelativistic EPP in the laboratory. Therefore predictions of the physical properties of such a system are needed.  The interaction of relativistic electrons and positrons is described by QED. In the presence of a thermal plasma background (heat bath) QED at finite temperature has to be considered. As discussed above, perturbative methods based on the imaginary (Matsubara) or real time formalism have been developed and applied to the physics of a QGP. Here we will transfer the methods and results for properties of a QGP to the case of an EPP, where a perturbative treatment is more reliable than in the case of a strongly coupled QGP. Some applications of thermal field theory to astrophysical plasmas have been discussed before (see e.g. Thoma (2002) and references therein and Altherr and Kraemmer (1992)).  An important difference to non-relativistic ion-electron plasmas (Lifshitz and Pitaevskii, 1981) are the relevant scales. In the non-relativistic case case they are given by the masses of the plasma particles and the temperature $T$. For example, the electron plasma frequency reads \\begin{equation} \\omega_{pl}=\\sqrt{\\frac{4\\pi e^2\\rho_e}{m_e}} \\label{e1} \\end{equation} and the Debye screening length due to the electrons in the plasma \\begin{equation} \\lambda_D=\\sqrt{\\frac{k_BT_e}{4\\pi e^2\\rho_e}}, \\label{e2} \\end{equation} where $\\rho_e$ is the electron number density, $T_e$ the temperature of the electron component, and $m_e$ the electron mass.  In an ultrarelativistic plasma with $T \\gg m$ the masses can be neglected and the important scales are the temperature $T$, called the hard scale, and the soft scale $eT$, which determines the collective phenomena as we will see below. Here we use natural units, i.e. $\\hbar =c=k_B=1$, as usual in quantum field theory. In these units $e=0.3$ corresponding to a fine structure constant $\\alpha =e^2/(4\\pi)=1/137$. For converting to SI-units we use $1\\> {\\rm MeV}\\> \\hat =\\> 1.60\\times 10^{-13}\\> {\\rm J} \\> \\hat =\\> 5.08\\times 10^{12}\\> {\\rm m}^{-1}\\> \\hat =\\> 1.52\\times 10^{21}\\> {\\rm s}^{-1}$.  In the next section we will discuss the equation of state of an equilibrated ultrarelativistic EPP. Then collective phenomena will be considered. Afterwards transport properties and particle production will be discussed. Finally we will describe properties of an EPP which is not in chemical equilibrium as in the case of laser induced plasmas. We will not consider here the formation process and equilibration of an EPP.  Many results presented here can be found in the literature for the case of a QGP (Thoma, 1995a) differing only by numerical factors, e.g. number of degrees of freedom. The purpose of this colloquium is to summarize these results and to extend them to laser induced EPPs as a reference for future experiments. As an example we consider a temperature of 10 MeV as it can be typically realized in laser produced and supernovae EPPs.  Laser produced QED plasmas have also been discussed recently in two review articles (Mourou {\\it et al.}, 2006; Marklund and Shukla, 2006) with emphasis on the production mechanism and non-linear effects. Here, however, we want to focus on the properties of an equilibrated EPP as they can be calculated from perturbative QED. Such an EPP in thermal and chemical equlibrium might be the outcome of future laser experiments if the intensity can be increased further on. As an example, we have chosen the predictions of Shen and Meyer-ter-Vehn (2001) based on a numerical simulation (PIC) and cross sections for electron-positron productions, in which two opposite laser beams are focussed on a thin gold foil leading to a chemically non-equiliibrated plasma (see section VI). However, this interesting proposal still waits for its experimental confirmation. Other production mechanisms, based on the Schwinger pair production effect in strong fields (Schwinger, 1951), have shown that pair production can be efficient already far below the Schwinger critical field strength, requiring laser intensities of $5 \\times 10^{29}$ W/cm$^2$, in the case of time dependent and inhomogeneous fields, e.g. two oppositely directed pulsed laser beams in vacuum (see e.g. Avetissian {\\it et al.}, 2002; Narozhny {\\it et al.}, 2004; Di Piazza, 2004; Dunne and Schubert, 2005; Gies and Klingmuller, 2005; Blaschke {\\it et al.} 2006a, Blaschke {\\it et al.}, 2006b, Sch\\\"utzhold {\\it et al.}, 2008). Furthermore the pair production in an X-ray free electron laser (XFEL) has been discussed (Ringwald, 2001; Alkofer {\\it et al.}, 2001; Roberts {\\it et al.}, 2002). QED plasmas in strong magnetic fields have been considered by Danielsson and Grasso (1995) and more general in strong electromagnetic fields by Morozov {\\it et al.} (2002). QED plasmas have also been studied in great detail in the book by Melrose (2008). ",
        "conclusions": ""
    },
    "0801/0801.4572_arXiv.txt": {
        "abstract": "Circumstellar disks are an integral part of the star formation process and the sites where planets are formed. Understanding the physical processes that drive their evolution, as disks evolve from optically thick  to optically thin, is crucial for our understanding of planet formation. Disks evolve through various processes including accretion onto the star, dust settling and coagulation, dynamical interactions with forming planets, and photo-evaporation. However, the  relative importance and timescales of these processes are still  poorly understood. In this review, I summarize current models of the different processes that control the evolution of primordial circumstellar disks around low-mass stars (mass $<$ 2 M$_{\\odot}$). I also  discuss recent observational developments on circumstellar disk evolution with a focus on new \\emph{Spitzer} results on transition objects. ",
        "introduction": "It is currently believed that virtually all stars form surrounded by a circumstellar disk, even if this disk is sometimes very short lived. This conclusion follows from simple conservation of angular momentum arguments and is supported by mounting evidence. This evidence ranges from the excess emission, extending from the near-IR (Strom et al. 1989) to the sub-millimeter (Osterloh $\\&$ Beckwith 1995),  that is observed in most young (age $<$ 1 Myr) pre-main-sequence (PMS) stars  to direct \\emph{Hubble} images of disks (McCaughrean \\& O'Dell 1996) seen as silhouettes in front of the Orion Nebula. Also, even though there is not \\emph{direct} evidence that planets actually grow from circumstellar material, it has become increasingly clear that they are in fact the birthplaces of planets since their masses, sizes, and compositions are consistent with the theoretical minimum-mass solar nebula (Hayashi 1981). Recently, the discovery of exo-planets orbiting nearby main sequence stars has confirmed that the formation of planets is a common process and not a rare phenomenon exclusive to our Solar System. Thus, any theory of planet formation should be robust enough to account for the high incidence  of planets and cannot rely on special conditions or on unlikely processes to convert circumstellar dust and gas into planets.  Star and planet formation are intimately related. Standard low-mass star formation models (e.g. Shu et al. 1987) describe the free fall collapse of a slowly rotating molecular cloud core followed by the development of a hydrostatic proto-star surrounded by an envelope and a disk of material supported by its residual angular momentum. This early phase is expected to occur on a timescale of about $10^5$ years (Beckwith 1999) and results in an optically revealed classical T Tauri star (CTTS, low-mass PMS star that shows clear evidence for accretion of circumstellar material). This stage is characterized by intense accretion onto the star, strong winds, and bipolar outflows. As the system evolves, presumably into a weak-lined T Tauri star (WTTS, PMS star mostly coeval with CTTSs but that do not show evidence for accretion), accretion ends, and the dust settles into the mid-plane of the disk where the solid particles are believed to stick together and to grow into planetesimals as they collide. Once the objects reach the kilometer scale, gravity increases the collision cross-section of the most massive planetesimals, and runaway accretion occurs (Lissauer 1993). In the standard core accretion model (e.g., Pollack et al. 1996), massive enough proto-planets still embedded in the disk can accrete the remaining gas and become giant planets. The early stages are the most uncertain, and many people suspect that grains will not grow into planetesimals by collisions alone at the rate necessary to go through all stages of giant planet formation before the gas nebula has dissipated. Also, even assuming that grains can grow into macroscopic bodies at the necessary rate, it is still doubtful whether they can form large planetesimals. Once objects reach the meter-size scale, they are expected to rapidly spiral inward due to the dynamical interactions with the gas in the disk, which rotates at a slightly sub-Keplerian velocity because it is supported  against gravity by pressure in addition to the centrifugal force.  However, since current statistics of extra-solar planets indicate that giant planets are common, the difficulties of the standard core accretion model have led some researchers (e.g., Boss 2000) to revisit an  alternative planet formation mechanism that had been put aside for several decades, namely, gravitational instability. In the gravitational instability scenario,  giant planets form through the direct gravitational collapse of a massive unstable disk over timescales $<$ 10$^{4}$ yrs. In some hybrid models (e.g.,  Youdin $\\&$ Shu 2002), solid particles settle to the mid-plane of the circumstellar disk and form a dense gravitationally unstable sub-disk. In this manner, planetesimals form from the gravitational collapse of the material in the mid-plane.  The collision of planetesimals leads to runaway accretion, and the process continues in a way  analogous to the core accretion model.  Thus, while the existence of planets around a significant fraction of all the stars is considered verified, the precise mechanisms through which planets form still remain largely unknown. So far, none of the proposed theories have proven satisfactory. On one hand, the standard model of continuous accretion of solid particles relies on doubtful sticking properties of rocks and on unknown processes to prevent the migration of meter-size objects. On the other hand, gravitational instability relies on unproven mechanisms to enhance the surface density of the disk's mid-plane to trigger the process of  planet formation. Clearly, more observational constraints are necessary to help the theoretical work on planet formation to proceed forward.  \\begin{figure}[!ht] \\plotone{f1.eps} \\caption{ Inner accretion disk fraction vs. stellar age inferred from H-K excess measurements, binned by cluster or association, for $\\sim$3500 stars from the literature. (Figure taken from Hillenbrand 2006). } \\end{figure} ",
        "conclusions": "The many transition disks discovered by \\emph{Spitzer} are likely to be the subjects of many follow-up observations and studies for years to come. These are objects that are undergoing rapid transformations and have properties that place them  somewhere in between two much better defined evolutionary states: those of regular CTTSs and debris disks. Understanding the physical processes responsible for the diversity of transition disks and its implications for planet formation will be one of main challenges for the field.  Another key outstanding disk evolution issue is the question of when the transition from the primordial to the debris disk stage actually occurs. As mentioned in Section 3, some of the WTTS disks have very small fractional luminosities and thus seem to be optically thin. These objects  \\emph{could} be younger analogs of the $\\beta$ Pic and AU Mic debris disks, and thus some of the youngest debris disks ever observed. However, this interpretation depends on the assumption that these young WTTS disks are gas poor. It is also possible that some of the WTTS disks have very low fractional disk luminosities as a result of most of their grains growing to sizes $\\gg$10-20 $\\mu$m, in which case they would still be primordial disks. Since a real debris disks requires the presence of second generation of dust produced by the collision of much larger objects, the ages of the youngest debris disks can constrain the time it takes for a disk to form planetesimals.  It has even been proposed that detectable second generation dust is not produced until Pluto-sized objects form and trigger a collision cascade (Dominik $\\&$ Decin, 2003). In that case, the presence of debris disks around $\\sim$1-3 Myrs old stars would have even stronger implications for planet formation theories. The fact that, as discussed in Section 3, the incidence of 24 $\\mu$m excesses around low-mass PMS stars with ages $\\sim$10 Myr is $\\sim$0 $\\%$, is consistent with the idea that a quiescent period exists between the dissipation/coagulation of the primordial dust and the onset of the debris phenomenon. However, the confirmation of such a scenario will require far-IR observations sensitive enough to detect, at the distance of the nearest star-forming regions (150-200 pc),  debris disks as faint as those observed in the solar neighborhood (10-30 pc). In the near future, \\emph{Herschel} will provide such sensitivity. It will also provide  the information on the gas content and grain size distribution of optically thin disks around WTTS required to establish whether they represent the end of the primordial disk phase or the beginning of the debris disk stage.  Even though much still remains to be learned about planet formation from the study of circumstellar disks, ultimately, we would like to be able to directly observe planets as they form. ALMA could achieve this milestone in the next decade. With its most extended baseline, ALMA will have a resolution of 0.007$''$ at 625 GHz. This corresponds to an angular resolution of less than 1 AU at the distance of nearby  star forming regions like Ophiuchus and Taurus.  Narayanan et al. (2006) model the molecular emission line from a gravitationally unstable protoplanetary disk (Figure 5, left panel) and conclude that forming giant planets could be directly imageable using dense gas tracers such as  HCO+ (Figure 5, right panel). Such observations could provide the ultimate tests needed to distinguish among the competing theories of planet formation."
    },
    "0801/0801.4091_arXiv.txt": {
        "abstract": "We have conducted an optical and infrared imaging in the neighbourhoods of 4 triplets of quasars. R, z$^\\prime$, J and K$_s$ images were obtained with MOSAIC II and ISPI at Cerro Tololo Interamerican Observatory. Accurate relative photometry and astrometry were obtained from these images for subsequent use in deriving photometric redshifts. We analyzed the homogeneity and depth of the photometric catalog by comparing with results coming from the literature. The good agreement shows that our magnitudes are reliable to study large scale structure reaching limiting magnitudes of  R = 24.5, z$^\\prime$ = 22.5, J = 20.5 and K$_s$ = 19.0.  With this catalog we can study the neighbourhoods of the triplets of quasars searching for galaxy overdensities such as groups and galaxy clusters. ",
        "introduction": "Our present understanding of galaxy formation and evolution can be improved significantly through studies of quasars. In the unified scenario, merger-driven processes include the origin of a quasar by self-regulating, supermassive, black-hole growth \\citep{hop}. Thus, most massive spheroids should host an optically-bright quasar during a brief period of its evolutionary history.  There is well-established evidence indicating that several galaxy properties  depend on the environment where the galaxies formed and evolved. Several processes - such as stimulated or truncated star formation, tidal stripping, merging, etc. - can determine the nature of galaxy populations and dynamics. Also to be considered is the possibility that AGN feedback processes could induce significant changes in the evolution of galaxies near to the quasar (see for instance \\citet{croton}). By studying quasar environments we may deepen our understanding of the relation of quasar fueling and galaxy interactions and merging; and as a consequence, the formation of bright galaxies.  The availability of large quasar samples has increased significantly in recent years, mainly due to the advent of large surveys such as the 2dF Galaxy Redshift Survey \\citep{colles} and the Sloan Digital Sky Survey \\citep{stou}.  These samples provide an opportunity for more robust statistical studies of the quasar phenomenon and its link with galaxy formation.  At low redshifts (z $\\le $ 0.3), existing cross-correlation analysis of quasar environments show that quasars inhabit environments similar to those of normal galaxies (\\citet{smitha}, \\citet{cold02}. However, on small scales (projected distance, $r_p \\leq 1 \\mpc$; and radial velocity difference, $\\Delta V = 500 \\kms$), the quasar environment is overpopulated by blue disk galaxies having a strong star-formation rate with respect to typical galaxy neighbourhoods \\citep{cold03,cold06}. This result is in agreement with those of \\citet{ilona02,ilona04} in the sense that low-redshift quasars follow the Large Scale Structure traced by galaxy clusters but they are not located near the center of clusters.  These quasars are mainly found in the periphery of such structures or between two, possibly merging, clusters.  At higher redshifts, quasars are often associated with rich environments \\citep{hall,djor} where the comoving density of galaxies  is higher than that expected for the general field.  This should be interpreted as the core of future rich clusters.  Some results at $z \\sim 1$ suggest that quasars reside on the peripheries of clusters or in cluster mergers \\citep{haines01,haines04,tanaka00,tanaka01}. \\citet{tanaka01}, for example, investigated a group of 5 quasars tracing a $\\sim$ 10 Mpc structure of 4--5 clusters, with only a single radio-loud quasar appearing to be directly associated with one of the clusters.  In general, the existing studies of quasar environments suggest their association with forming structures, marking merging clusters and filaments. Consequently, groupings of quasars can be expected to trace regions of extraordinary activity. Pairs of quasars have been strongly studied and they are used to identify cluster of galaxies at high redshifts. \\citet{zhdanov} suggested that quasar pairs can be used as tracers of high redshift large-scale structures. Moreover, \\citet{djor03} found a quasar pair at  $z \\sim 5$ associated with large scale structure, possibly a protocluster at that redshift. More recently, \\citet{boris} found that quasar pairs at $z \\sim 1$ are excellent tracers of high density environments suggesting that they can be used to find galaxy clusters.  Triplets of quasars in a relative small volume are extremely rare and until now, no cluster-scale triplets have been published or investigated. A very compact quasar association ($\\sim$ 50 Kpc separation) was recently reported by \\citet{djor07} suggesting that this could be a compact galaxy group in a merging process. We propose to use triplets of quasars as ``lighthouses'' to find and study clusters of galaxies that favour such rare events. With this aim in mind, we performed an optical and infrared photometric study of the environments of triplets of quasars at different redshifts on the scale of galaxy clusters. The images allow us to observe galaxies in the early Universe and they are used to investigate the environments of quasar triplets We also study the group and cluster morphology, and the photometric properties of individual galaxies in the neighbourhood of each triplet. We have already found that low-z triplets of Seyfert 1 galaxies are associated with extremely rich clusters of galaxies \\citep{ilona07}.  If the quasar triggering mechanism does not evolve strongly over redshift we can expect to find some of the richest galaxy clusters ever discovered at $z \\sim 1$. A negative result will be clear evidence of the evolution of the quasar formation mechanism.   This paper is organized as follows: In section 2 we describe the sample selection and section 3 shows the optical and infrared observations and data reduction. Section 4 describes the photometric galaxy catalog, and in section 5 we provide a summary of the main results. We adopt the latest cosmology with $\\Omega_M = 0.3$, $\\Omega_{\\lambda} = 0.7$ and a Hubble constant, $H_0=70 ~ {\\rm{km ~ s^{-1} ~ Mpc^{-1}}}$. ",
        "conclusions": "We selected a sample of triplets of quasars as those triple systems with percolation longitudes smaller than 2 $\\mpc$ and 2000 $\\kms$.  The goal is to use them as ``lighthouses'' to find richest clusters. For this purpose we conducted a photometric study of the neighbourhoods of 4 triplets of quasars in the redshift range of 0.9 to 1.7. Images in R, z$^\\prime$, J and K$_s$ bands were obtained with MOSAIC II and ISPI at CTIO.  Magnitudes and coordinates were found to be in excellent agreement with the USNO-A2.0 and 2MASS catalogs. We analyzed the homogeneity and depth of the photometric catalog by comparing the number counts with results coming from the literature. The excellent agreement shows that our magnitudes are reliable, down to limiting magnitudes of R = 24.5, z$^\\prime$ = 22.5, J = 20.5 and K$_s$ = 19.0.  These results allow us to study large- scale structure and in a forthcoming paper we will deal with the statistical analysis of this data."
    },
    "0801/0801.1211_arXiv.txt": {
        "abstract": "We report results based on the monitoring of the BL Lac object Mrk 501 in the optical (B, V and R) passbands from March to May 2000. Observations spread over 12 nights were carried out using 1.2 meter Mount Abu Telescope, India and 61 cm Telescope at Sobaeksan Astronomy Observatory, South Korea. The aim is to study the intra-day variability (IDV), short term variability and color variability in the low state of the source. We have detected flux variation of 0.05 mag in the R-band in time scale of 15 min in one night. In the B and V passbands, we have less data points and it is difficult to infer any IDVs. Short term flux variations are also observed in the V and R bands during the observing run. No significant variation in color (B$-$R) has been detected but (V$-$R) shows variation during the present observing run.  Assuming the shortest observed time scale of variability (15 min) to represent the disk instability or pulsation at a distance of 5 Schwarschild radii from the black hole (BH), mass of the central BH is estimated $\\sim$ 1.20 $\\times$ 10$^{8} M_{\\odot}$.  {\\bf PACS:}  98.54.Cm, 95.85.Kr, 95.75.De, 95.75.Wx ",
        "introduction": "Blazars constitute a class of radio-loud active galactic nuclei (AGNs) consisting of BL Lacertae objects (BL Lacs) and flat spectrum radio quasars (FSRQs). In the unified schemes of AGNs, blazars are believed to be the beamed counterparts of the FR I/FR II radio galaxies and many of their properties can be well understood in such a scenario (e.g. Browne 1983; Antonucci 1993; Urry \\& Padovani 1995). According to orientation based unification scheme for radio-loud AGNs, blazar's jet make angles $\\leq$ 20$^{\\circ}$ or so to the line of sight (Urry \\& Padovani 1995). Significant flux variations in the whole range of electromagnetic spectrum on time scale of less than a day to several years are common in blazars. Their radiation at all wavelengths is predominantly non thermal and strongly polarized ($> 3\\%$) in radio to optical bands.  Blazar variability can be broadly divided into 3 classes viz. micro variability or intra-night variability or intra-day variability, short term outbursts and long term trends. Variations in flux of a few tenth of a magnitude in a time scale of tens of minutes to a few hours (less than a days) is often known as intra-day variability (IDV) (Wagner \\& Witzel 1995). Short term outbursts can range from weeks to months and long term trends can have time scales of months to several years.  Study of optical and near-IR variability of blazars on diverse time scales was carried out by several groups (Miller \\& McGimsey 1978; Xie et al. 1988; Webb et al. 1988; Kidger \\& de Diego 1990; Carini 1990; Carini et al. 1992; Heidt \\& Wagner 1996; Bai et al. 1998, 1999; Fan \\& Lin 1999; Ghosh et al. 2000; Gupta et al. 2002, 2004; Stalin et al. 2005; and references therein). The first convincing evidence of optical IDV was found in BL Lacertae by Miller et al. (1989). The typical optical IDV found in most of the blazars is $\\sim$ 0.01 mag hr$^{-1}$. Carini (1990) observed a sample of 20 blazars and found most of them showing detectable inter-night variations. He also noted that in observing runs of more than 8 hours, the probability of detecting significant IDV was $\\sim$ 80\\%. Subsequently, Heidt \\& Wagner (1996) observed a sample of 34 radio selected BL Lacs and detected IDV in 28 (i.e. 82\\%) sources out of which 75\\%  sources showed significant variation with in $<$ 6 hours which further supports the finding of high frequency of occurrence of IDV in blazars. In our recent paper (Gupta \\& Joshi, 2005), we studied the frequency of occurrence of optical IDV in different classes of AGNs and found about 10\\% (18/174) radio-quiet AGNs, 35$-$40\\% (41/115) radio-loud AGNs (excluding blazars) and 70\\% (79/113) blazars to show IDV (in fact $\\sim$ 80$-$85\\% blazars show IDV, if observed for more than 6 hours).  Mrk 501 is one of the most interesting and nearby BL Lac object at z = 0.034 which makes it the second closest known BL Lac object after Mrk 421 (z = 0.031).  Mrk 501 has been studied for variability in all regions of the electromagnetic  spectrum (e.g. Fan \\& Lin 1999, Kataoka et al. 1999, Sambruna et al. 2000, Ghosh et al.  2000, Xue \\& Cui 2005, Gliozzi, et al. 2006, and references therein). In optical region  this source has been studied by several groups e.g Stickel et al. (1993) have reported  variation of 1.3 magnitude in luminosity; Heidt \\& Wagner (1996) have reported $\\sim$ 32\\% variation in flux in less than two weeks time; Ghosh et al. (2000) have made 10 nights  observations on Mrk 501 during March to June, 1997 to search for rapid variability  and reported rapid variability in 7 nights; and recently, Fan et al. (2004) observed  Mrk 501 for two hours in one night in October 2000 and found no significant variation.  The variability in blazars is mainly attributed to either due to the disc instability or the activity in the jet. IDV reported in the flaring and high state is generally attributed to the shock moving down the inhomogeneous medium in the jet. But small amplitude variation in the low state of blazar could be due to instability in the accretion disk (Witta et al. 1991, 1992; Chakrabarti \\& Wiita 1993; Mangalam \\& Wiita 1993). Recently, it is noticed that, in the low luminosity AGNs, accretion disk is radiatively inefficient (Chiaberge et al. 2006; Chiaberge \\& Macchetto 2006; Macchetto \\& Chiaberge 2007; Capetti et al. 2007). So, there will be an alternative way to explain the IDV in the low-state of blazars, in which a weak jet emission will be responsible for the IDV. So far there is no clear answer to this. There is a need to study IDV and short term variability in the optical bands with a good sampling in time when the source is in low state. In view of this we pursued the present study to address both the IDV and short term variability on Mrk 501.  This paper is structured as follows: section 2 describes photometric observations and data reductions, section 3 present the results, and discussions and conclusions are presented in section 4. ",
        "conclusions": ""
    },
    "0801/0801.3214_arXiv.txt": {
        "abstract": "Famous for the extraordinary richness of its molecular content, the Sgr B2 molecular cloud complex is the prime target in the long-standing search for ever more complex species. We have completed a molecular line survey of the hot dense cores Sgr B2(N) and SgrB2(M) in the 3~mm wavelength range with the IRAM 30\\,m telescope. We performed the analysis of this huge data set by modeling the whole spectrum at once in the LTE approximation. Ongoing analyses yield an average line density of about 100 features/GHz above 3$\\sigma$ for Sgr B2(N), emitted and/or absorbed by a total of 51 molecular species. We find lines from 60 rare isotopologues and from 41 vibrationally excited states in addition to the main species, vibrational ground state lines. For Sgr B2(M), we find about 25 features/GHz above 3$\\sigma$, from 41 molecular species plus 50 isotopologues and 20 vibrationally excited states. Thanks to the constant updates to the Cologne Database for Molecular Spectroscopy, we are working our way through the assignment of the unidentified features, currently 40$\\%$ and 50$\\%$ above 3$\\sigma$ for Sgr B2(N) and SgrB2(M), respectively. ",
        "introduction": "\\label{s:intro}  With several active regions and a total mass of more than $10^6$ M$_\\odot$,  Sagittarius B2 (hereafter SgrB2) is one of the most complex and massive sites of star formation in the Galaxy. It is located close to the galactic center and harbors two hot dense cores (M and N), which fragment into several sub cores. SgrB2(N) has a very rich chemistry and was called the Large Molecule Heimat (LMH) by \\citet{Snyder94}. Many complex molecules were discovered there like, e.g., acetic acid \\citep[CH$_3$COOH,][]{Mehringer97}, glycolaldehyde \\citep[CH$_2$(OH)CHO,][]{Hollis00}, and acetamide \\citep[CH$_3$CONH$_2$,][]{Hollis06}. It is therefore a prime target to look for new complex molecules.  Line surveys at (sub)mm wavelengths are needed to search for large complex molecules since these molecules emit hundreds of (weak) lines. Regions with rich chemistry produce very crowded spectra with many blended lines and the confusion limit is easily reached. In this context, the secure identification of a new molecule requires the identification of many lines and no single line predicted by a temperature and column density derived from multiple lines should be missing \\citep[see, e.g., the rebuttal of glycine by][]{Snyder05}. Since the spectra are complex, modeling the emission of all known molecules is necessary to prevent mis-assignments and point out blended lines before claiming the detection of a new molecule.  To search for new complex molecules, we carried out a complete line survey at 3~mm toward the two hot cores SgrB2(N) and SgrB2(M) and we present here some preliminary results.   \\begin{figure}[t] \\begin{center} \\includegraphics[width=9.0cm,angle=270]{belloche_fig1.eps} \\caption{Spectra obtained toward SgrB2(N) (\\textit{top}) and SgrB2(M) (\\textit{bottom}) with the IRAM 30\\,m telescope between 80 and 116 GHz. The blow-ups reveal the full complexity of the SgrB2(N) spectrum and the prominent multi-component absorption lines in the SgrB2(M) spectrum .} \\label{f:survey} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0801/0801.4328_arXiv.txt": {
        "abstract": "{} {We searched for massive stars with Balmer-emission consistent with magnetically confined circumstellar material.} {Archival spectroscopic and photometric data were investigated.} {\\object{HR\\,7355} is a formerly unknown He-strong star showing Balmer emission.  At $V=6.02$\\,mag, it is one of the brightest objects simultaneously showing anomalous helium absorption and hydrogen emission. Among similar objects, only $\\sigma$\\,Ori~E has so far been subjected to any systematic analysis of the circumstellar material responsible for the emission.  We argue that the double-wave photometric period of 0.52\\,d corresponds to the rotation period.  In tandem with the high projected equatorial velocity, $v \\sin i=320$\\,km\\,s$^{-1}$, this short period suggests that HR\\,7355 is the most rapidly rotating He-strong star known to date; {a class that was hitherto expected to host stars with slow to moderate rotation only.}} {} ",
        "introduction": "In the early B-type spectral range a subclass of He-strong stars is found, i.e.\\ stars showing Helium lines with abnormally large equivalent widths.  The chemical surface abundances of these stars are influenced by the presence of a strong magnetic field, resulting in a He overabundance that typically varies in strength over the stellar surface.  Because He-strong stars are sufficiently luminous to harbour radiatively driven winds \\citep[as diagnosed by ultraviolet absorption line diagnostics; see][]{1990ApJ...365..665S}, they represent ideal laboratories for understanding the process of magnetic wind confinement \\citep[][]{1997A&A...323..121B}.  Typically, the fields of these stars are too strong for them to be amenable to the magnetohydrodynamical (MHD) simulations \\citep[e,g,][]{2002ApJ...576..413U}. However, an alternative Rigidly Rotating Magnetosphere (RRM) model for the circumstellar distribution of magnetocentrifugally confined wind plasma by \\citet{2005MNRAS.357..251T} has shown much promise in reproducing the detailed optical variability of the archetype emission-line He-strong star \\object{$\\sigma$\\,Ori~E}.  {To date, our knowledge on He-strong stars is limited to slow to moderate rotators, as no rapid rotators had been found. This goes as far as to the conclusion that slow rotation is an intrinsic property of He-strong stars \\citep{1983ApJ...268..195W,1999A&A...345..244Z}, that has to be taken into account by the search for the origin of the magnetic field.  This work not only reports the discovery of one more bright massive star to host a magnetosphere for application of the above model, but also extends the parameter range in which He-strong stars are found by a factor of about two in rotational velocity space.}    \\object{HR\\,7355} (\\object{HD\\,182\\,180}, \\object{HIP\\,95\\,408}) is a little-observed B2Vn star of $6^{\\rm th}$ magnitude ($V=6.02, B=5.91$), located toward the galactic center. It was listed as a MK-standard by \\citet{1969ApJ...157..313H}, but never examined in detail. Other investigators have instead classified it as B5IV \\citep[see][]{1964PLPla..28....1J}. From studies of larger samples of stars that included HR\\,7355, we know that the star is very rapidly rotating: \\citet{2002ApJ...573..359A} measured $v\\sin i = 320$\\,km\\,s$^{-1}$, while \\citet{2000AcA....50..509G} report $v\\sin i = 270\\pm30$\\,km\\,s$^{-1}$. Hipparcos photometric data indicate that the star is a periodic variable, with $P=0.26\\,$d \\citep{2002MNRAS.331...45K}.   \\begin{figure*}[t] \\parbox{18cm}{\\centerline{\\psfig{figure=spectra.ps,width=16cm,clip=t,angle=0}}}% \\caption[xx]{\\label{spectra}Changes in several representative lines between the spectra taken in 1999 (solid line) and 2004 (dotted). The 1999 profile has H$\\alpha$ emission extending out to several times beyond $v \\sin i$ (the latter indicated by the vertical dotted lines, upper left) Note that for the Balmer lines (left column) a wider range in velocity is shown than in the other panels.} \\end{figure*} ",
        "conclusions": "HR\\,7355 is a previously unknown spectroscopically variable star, and as such it should no longer be used as a spectral standard star. In its capacity as the newest member of the He-strong class, it is not only one of the brighter stars in this class, but is also the most rapidly rotating.  In addition to its spectral variability, HR\\,7355 is periodically variable in photometry, with either a single-wave sinusoidal lightcurve of $P_{\\rm sin}=0.260714\\pm0.000003$\\,d or a double-wave pattern with $P_{\\rm dw}=0.521428\\pm0.000006$\\,d.  At this point we cannot firmly exclude the possibility that the photometric variations originate in some mechanism other than the surface abundance inhomogeneities. However, the spectra do not show the typical signature of pulsation, and moreover we do not find any other signal in the photometric data consistent with the rotation period. Thus, if the spectral and photometric variations repeat on the same period, then that period is the double-wave period, which also is the rotational period. {It is the first rapidly rotating He-strong star, and as such may pose a challenge to field origin hypotheses that would have led to strong magnetic braking.}  We intend to begin a monitoring campaign on the star, to obtain a spectroscopic time series {for further analysis of the fundamental parameters of HR\\,7355, as well as for application of} the framework of the RRM model of \\citet{2005MNRAS.357..251T}. This model has proven extremely successful in explaining the emission-line variations of $\\sigma$\\,Ori~E \\citep{2005ApJ...630L..81T}, and we are optimistic that it can explain the behaviour of HR\\,7355 too."
    },
    "0801/0801.3028_arXiv.txt": {
        "abstract": "% Puzzles often give birth to the great discoveries, the false discoveries sometimes stimulate the exiting ideas in theoretical physics. The  historical examples of both are described in Introduction and in section ``Cosmological Puzzles''. From existing puzzles most attention is given to Ultra High Energy Cosmic Ray (UHECR) puzzle and to cosmological constant problem. The 40-years old UHECR problem consisted in absence of the sharp steepening in spectrum of extragalactic cosmic rays caused by interaction with CMB radiation. This steepening is known as Greisen-Zatsepin-Kuzmin (GZK) cutoff. It is demonstrated here that the features of interaction of cosmic ray protons with CMB are seen now in the spectrum in the form of the dip and beginning of the GZK cutoff. The most serious cosmological problem is caused by large vacuum energy of the known elementary-particle fields which exceeds at least by 45 orders of magnitude the cosmological vacuum energy. The various ideas put forward to solve this problem during last 40 years, have weaknesses and cannot be accepted as the final solution of this puzzle. The anthropic approach is discussed. ",
        "introduction": "\\begin{center} ALL GREAT DISCOVERIES IN ASTROPHYSICS\\\\ APPEARED UNPREDICTABLY AS PUZZLES.\\\\ WHAT WAS PREDICTED WAS NOT FOUND.\\\\ \\end{center} \\vspace{5mm} Not many good things fall down on us from the sky, but discoveries do. They arrive as the puzzles, and usually disappear as the errors. But sometimes they become real, leaving behind the great discoveries. I will give below a short list of astrophysical discoveries of the last four decades, separating intuitively astrophysics from cosmology.\\\\*[3mm] {\\it Quasars} were discovered in early 1960s as compact radio sources. G. Mathews and A. Sandage in 1960 identified radio source 3C48 with a stellar-like object. M. Schmidt in 1963 deciphered the optical spectrum of quasar 3C273 assuming its redshift, $z=0.158$. Surmounting resistance of sceptics, this explanation moved the source to the distance of 630~Mpc and made its luminosity uncomfortably large, $L \\sim 10^{46}$~erg/s. It was a puzzle, and many respectable astrophysicists spent years trying to squeeze the source back into Galaxy. All of them failed, and in the end the puzzling energy release resulted in the discovery of a {\\it black hole}, an object of general relativity.  \\\\ % *[3mm] {\\it Pulsars} were discovered first in 1967 by a student of A. Hewish, Jocelyn Bell. She observed a puzzling periodicity of radio-pulses from an unknown source. After short but intense discussion of different possible sources, including extraterrestrial civilizations and ``little green men'', the magnetized rotating neutron stars, the pulsars, were found to be responsible. It opened a new field of cosmic physics: {\\em relativistic electrodynamics}.\\\\ % *[3mm] {\\it The atmospheric neutrino anomaly and the solar neutrino problem} went along most difficult road  to the status of discovery. The puzzling phenomenon in both cases was a neutrino deficit as compared with calculations. The solar neutrino experiments have been started by Ray Davis in 1960 in Brookhaven. With time the solar-neutrino deficit raised, but scepticism of the community raised too. Pushed first by Ray Davis and John Bahcall, the solar neutrino problem moved like  a slow coach along a road of three decades long. Fortunately, physics differs from democracy: opinion of majority means usually less than that of ONE. In the end the two obscure puzzles (solar and atmospheric neutrino deficits) have turned  into discovery of the most fascinating phenomenon, {\\it neutrino oscillations}, theoretically predicted by Bruno Pontecorvo. \\\\*[3mm] {\\it Supernova SN 1987a} became an elementary-particle laboratory in the sky for a study of properties of neutrinos, axions, majorons {\\it etc}. Detection of neutrinos became a triumph of the theory: the number of detected neutrinos, duration of the neutrino pulse and estimated neutrino luminosity have appeared in agreement with theoretical prediction. Gravitational collapse as a phenomenon providing the SN explosion has been established ... and a puzzle appeared. It was the  triumph of the incomplete theory, without rotation of collapsing star. The asymmetric ring around SN 1987a implies that the presupernova was a rotating star (it would be a surprise if not!). Rotation should diminish the neutrino luminosity or even change the collapsing scenario. Why then there is a beautiful agreement between theory and observation? This problem still expects its solution. \\subsection{Greatness of false discoveries} \\begin{center} WHEN FIRST APPEARED THE PUZZLES LOOK WEAK.\\\\ SAVE YOUR TIME SAYING: IT'S RUBBISH.\\\\ IN 90$\\%$ OF CASES YOU WILL BE RIGHT,\\\\ BUT YOU MAY MISS THE GREAT PHYSICS. \\end{center} \\vspace{3mm} False discoveries often have great impact on physics, and {\\it Cyg X-3 } is a famous example.\\\\*[1mm] Cyg X-3 is a galactic binary system well studied in all types of radiations, most notably in X-rays. In 80s many EAS arrays detected from it the 4.8 hour periodic ``gamma-ray'' signal in VHE (Very High Energy, $E\\geq 1$~TeV) and UHE (Ultra High Energy, $E\\geq 0.1-1$~PeV) ranges. The list of these arrays included Kiel, Haverah Park, Fly's Eye, Akeno, Carpet-Baksan, Tien-Shan, Plateau Rosa, Durham, Ooty, Ohya, Gulmarg, Crimea, Dugway, Whipple and others. Probably it is easy to say that there was no single EAS array which claimed no-signal observation. Additionally, some underground detectors (NUSEX, Soudan, MUTRON) marginally observed high energy muon signal from the direction of this source. Apart from the Kiel array, which claimed $6\\sigma$ signal, the confidence level of detection was not high: $3-4 \\sigma$.  In 1990 - 1991 two new-generation detectors, CASA-MIA and CYGNUS, put the stringent upper limit to the signal from Cyg X-3, which excluded early observations.  Apart from two lessons: \\begin{enumerate} \\item good detectors are better than bad ones, \\item ``$3\\sigma$'' discoveries should not be trusted, even if many detectors confirm them, \\end{enumerate} experience of Cyg X-3 has taught us how to evaluate statistical significance searching for periodic signals.  The false discovery of high energy radiation from Cyg X-3 had great impact on theoretical high energy astrophysics, stimulating study of acceleration in binary systems, production of high energy gamma and neutrino radiation  and creation of high energy astrophysics with new particles, such as light neutralinos, gluinos {\\it etc}. ",
        "conclusions": ""
    },
    "0801/0801.4578_arXiv.txt": {
        "abstract": "We present mid-infrared spectroscopy of a sample of 16 optically faint infrared luminous galaxies obtained with the Infrared Spectrograph (IRS) on the $Spitzer~Space~Telescope$. These sources were jointly selected from $Spitzer$ and $Chandra$ imaging surveys in the NDWFS Bo\\\"otes field and were selected from their bright X-ray fluxes to host luminous AGN. None of the spectra show significant emission from polycyclic aromatic hydrocarbons (PAHs; 6.2$\\rm \\mu m$ equivalent widths $<$0.2$\\rm \\mu m$), consistent with their infrared emission being dominated by AGN. Nine of the X-ray sources show 9.7$\\rm \\mu m$ silicate absorption features. Their redshifts are in the range $0.9<z<2.6$, implying infrared luminosities of log(L$_{IR}$)=12.5$-$13.6 $L_\\odot$. The average silicate absorption strength is not as strong as that of previously targeted optically faint infrared luminous galaxies with similar mid-infrared luminosities implying that the X-ray selection favors sources behind a smaller column of Si-rich dust than non-X-ray selection. Seven of the X-ray sources have featureless power-law mid-IR spectra. We argue that the featureless spectra likely result from the sources having weak or absent silicate and PAH features rather than the sources lying at higher redshifts where these features are shifted out of the IRS spectral window. We investigate whether there are any correlations between X-ray and infrared properties and find that sources with silicate absorption features tend to have fainter X-ray fluxes and harder X-ray spectra, indicating a weak relation between the amount of silicate absorption and column density of X-ray-absorbing gas. ",
        "introduction": "Mid-infrared spectroscopy is a useful tool in investigating the extent to which starbursts and AGN contribute to the luminosity of luminous infrared galaxies (LIRGs; \\citealt{san88}). The infrared spectra of local LIRGs have been studied extensively with ISOPHOT, ISOCAM and SWS on the $Infrared$ $Space$ $Telescope$ (ISO; e.g., \\citealt{lut98}; \\citealt{gen98}; \\citealt{rig99}; \\citealt{tra01}) and more recently with the infrared spectrograph (IRS; \\citealt{hou04}) on the $Spitzer~ Space~ Telescope$ (\\citealt{wee05}; \\citealt{brl06}; \\citealt{arm07}; \\citealt{des07}). $Spitzer$ IRS has also extended these studies to fainter sources at higher redshifts. A particularly interesting population of high redshift optically faint infrared luminous sources (\\citealt{yan04}; \\citealt{bra07}; \\citealt{fio08}; \\citealt{dey08}) has been discovered using the Multiband Imaging Photometer (MIPS; \\citealt{rie04}). Various observing programs with $Spitzer$ IRS have found that MIPS sources at flux density levels of f$_{\\nu}$(24$\\rm \\mu m$) $\\ge$ 0.8 mJy with optical magnitudes $R$ $\\ga$ 24 are typically at z $\\sim$ 2$-$3 (\\citealt{hou05}; \\citealt{wee06} \\citealt{yan07}; \\citealt{saj07}). The infrared spectra exhibit a wide range of properties from deep 9.7$\\rm \\mu m$ silicate absorption features and/or strong polycyclic aromatic hydrocarbon (PAH) emission features to featureless power-law continua. This obscured population is important in determining how much light originates from starbursts and AGN at epochs in the Universe when their luminosity density was at its maximum. Given that their optical faintness makes optical follow-up difficult, understanding how to characterize these sources using their infrared spectra is important in understanding their energetics.  One approach in understanding the contribution of starbursts and AGN to the infrared emission of optically faint infrared sources is to investigate the infrared spectra of sources also selected in other wavebands. For example, at one extreme we can study sources selected on the basis of large far-infrared flux densities, which we might expect to contain cool dust and show signatures of starbursts. In \\citet{bra08}, we presented the IRS spectra of 11 70$\\rm \\mu m$-selected optically faint infrared luminous galaxies and found that seven had large 6.2$\\rm \\mu m$ PAH equivalent widths, consistent with them being powerful starburst galaxies. Four of the galaxies showed deep silicate absorption features and their properties suggested that AGN are likely to be the dominant origin of their large mid-infrared luminosities.  At the other extreme, we can consider sources selected on the basis of X-ray flux, which we might expect to show spectral signatures of AGN. \\citet{wee06b} presented IRS spectra of 9 X-ray-loud optically faint infrared luminous galaxies selected from the $Spitzer$ Wide-Area Infrared Extragalactic Survey (SWIRE; \\citealt{lon04}). Eight sources have silicate absorption features and one source has a featureless power-law mid-IR spectrum.  In this paper, we present $Spitzer$ IRS spectra of 16 optically faint infrared galaxies in the Bo\\\"otes field of the NOAO Deep Wide-Field Survey (NDWFS; \\citealt{jan99}) which are selected to be X-ray loud. This significantly increases the number of X-ray loud optically faint infrared luminous galaxies thus far studied, and extends the sample to higher X-ray and 24$\\rm \\mu m$ luminosities.  A cosmology of $H_0 = 70 {\\rm ~km ~s^{-1} ~Mpc^{-1}}$, $\\Omega_M$=0.3, and $\\Omega_\\Lambda$=0.7 is assumed throughout. ",
        "conclusions": "We have presented mid-infrared spectra obtained using $Spitzer$ IRS of a sample of 16 X-ray loud 24 $\\rm \\mu m$ sources. The spectra of these sources are characterized by either silicate absorption or a featureless power-law, with no evidence of PAH emission. This suggests that the mid-infrared emission from these sources is dominated by warm dust emission surrounding the active nucleus with little contribution from star formation. The absorption-dominated sources lie in the range 0.9$<z<$2.6 and have very high X-ray and infrared luminosities (log L$_{IR}$=12.5-13.6 L$_\\odot$). The average absorption depth in our absorption-dominated X-ray sources is similar to that of local Seyfert 2 AGN and much smaller than that of typical ULIRGs. This suggests a difference in the distribution of the dusty material between our sample and typical ULIRGs. The weak silicate absorption feature in our sample could indicate a more clumpy dust distribution or a lower line of sight obscuration. There is evidence in the average spectrum for weak [Ne II], [Ne V], and [NeIII] lines. The [Ne V] is relatively stronger than is seen in local Seyferts, suggesting a higher ionization of the narrow-line region in these more powerful AGN. Given their higher X-ray fluxes, it is unlikely that the power-law sources lie at $z>$2.6. Instead, they must have only weak silicate absorption or emission features. Comparison with other samples suggests that the fraction of power-law sources appear to increase with X-ray luminosity and decrease with X-ray obscuration. Both the infrared and infrared to optical colors suggest that the sources are relatively unobscured compared with $R$-[24]$>$15 sources. Our X-ray selection is likely to be preferentially biased to less obscured and/or more powerful AGN. We find that sources with silicate absorption features tend to have fainter X-ray fluxes and larger hardness ratios. This suggests that more absorbed sources also have higher column densities of absorbing gas along the same line of sight."
    },
    "0801/0801.4863_arXiv.txt": {
        "abstract": "We show that as many as $\\sim$50 quasars with at least mJy-level expected flux density can be pre-selected as potential in-beam phase-reference targets for ASTRO-G. Most of them have never been imaged with VLBI. These sources are located around strong, compact calibrator sources that have correlated flux density $>100$~mJy on the longest VLBA baselines at 8.4~GHz. All the targets lie within $12\\arcmin$ from the respective reference source. The basis of this selection is an efficient method to identify potential weak VLBI target quasars simply using optical and low-resolution radio catalogue data. The sample of these dominantly weak sources offers a good opportunity for a statistical study of their radio structure with unprecedented angular resolution at 8.4~GHz. ",
        "introduction": "Phase-referencing is a way to increase the sensitivity of the Space Very Long Baseline Interferometry (SVLBI) observations that provide extremely high angular resolution due to the baselines exceeding the Earth diameter. Phase-reference imaging in ground-based VLBI is usually done in cycles of interleaving observations between a weak target source and a nearby strong reference source. Delay, delay-rate and phase solutions obtained for the phase-reference calibrator are interpolated and applied for the target source within the atmospheric coherence time, thus increasing the coherent integration time on the weak target source.  Unlike the first dedicated SVLBI satellite HALCA \\citep{hira00}, the next-generation satellite ASTRO-G will be capable of rapid attitude changing maneuvers. This, and the accurate orbit determination will allow us to observe suitable nearby reference--target source pairs in the traditional ``nodding'' style \\citep{asak07}. There is another, technically less demanding method which does not require rapid changes in the space antenna pointing if the reference--target separation is so small that both sources are within the primary beam of the 9.3-m ASTRO-G paraboloid antenna ($\\sim12\\arcmin$ at 8.4~GHz). In this scenario, the ground-based part of the SVLBI network performs the usual reference--target switching cycles, while the space antenna remains pointed to the same celestial position. (The diameters of the ground-based VLBI antennas are at least a factor of $\\sim3$ larger, and thus their primary beam sizes are considerably smaller than that of the orbiting antenna.)  Successful in-beam phase-referencing experiments have already been conducted with HALCA which could not quickly change its antenna pointing \\citep{pori00,bart00,porc00,guir01}. However, the use of in-beam phase-referencing is severely limited by the small number of sufficiently close source pairs known in the sky. Generally speaking, for any given target source of interest, it is very unlikely to find a suitable phase-reference calibrator within the primary beam of even the relatively small-diameter space antenna.  One may reverse the usual logic and select the phase-reference calibrator sources first. {\\it Then} it becomes possible to look for potential weak target sources that are located so close to one of the reference sources that in-beam phase-referencing observations with the orbiting antenna are feasible. Here we show that as many as $\\sim$50 quasars with at least mJy-level expected flux density can be pre-selected as potential in-beam phase-reference targets for ASTRO-G. This prospective sample is large enough for a statistical study of sub-milliarcsecond radio structures of weaker sources in comparison with bright ones. Such a sample, most of which have never been studied with VLBI, would certainly contain individually interesting radio quasars as well. The suitability of at least some of the candidate sources could be verified with ground-based VLBI observations prior to the launch of ASTRO-G. ",
        "conclusions": ""
    },
    "0801/0801.1488_arXiv.txt": {
        "abstract": "We present two-dimensional (2D) stellar and gaseous kinematics of the inner $\\sim$\\,130$\\times$180 pc$^2$ of the Narrow Line Seyfert 1 galaxy NGC\\,4051 at a sampling of 4.5\\,pc, from near-infrared $K$-band spectroscopic  observations obtained with the Gemini's Near-infrared Integral Field Spectrograph (NIFS) operating with the ALTAIR  adaptive optics module. We have used the CO absorption bandheads around 2.3\\,$\\mu$m to obtain the stellar kinematics which show the turnover of the rotation curve at only $\\approx$55\\,pc from the nucleus, revealing a  highly concentrated gravitational potential. The stellar velocity dispersion of the bulge is $\\approx$\\,60\\,km\\,s$^{-1}$ -- implying on a nuclear black hole mass of $\\approx\\,10^6$\\,M$_\\odot$ -- within which patches of lower velocity dispersion suggest the presence of regions of more recent star formation. From measurements of the emission-line profiles we have constructed two-dimensional maps for the flux distributitions, line ratios, radial velocities and gas velocity dispersions for the H$_2$, H\\,{\\sc i} and [Ca\\,{\\sc viii}]  emitting gas. Each emission line samples a distinct kinematics. The Br$\\gamma$ emission-line shows no rotation as well as no blueshifts or redshifts in excess of 30\\,km\\,s$^{-1}$, and is thus not restricted to the galaxy plane. The [Ca\\,{\\sc viii}] coronal region is compact but resolved, extending over the inner 75\\,pc. It shows the highest blueshifts -- of up to $-250$\\,km\\,s$^{-1}$, and the highest velocity dispersions, interpreted as due to outflows from the active nucleus, supporting an origin close to the nucleus. Subtraction of the stellar velocity field from the gaseous velocity field has allowed us to isolate  non-circular motions observed in the H$_2$ emitting gas. The most conspicuous kinematic structures are two nuclear spiral arms -- one observed in blueshift in the far side of the galaxy (to the NE), and the other observed in redshift in the near side of the galaxy (to the SW). We interpret these structures as inflows towards the nucleus, a result similar to those of previous studies in which we have found streaming motions along nuclear spirals in ionized gas using optical IFU observations.  We have calculated the mass inflow rate along the nuclear spiral arms, obtaining $\\dot{M}_{H_2} \\approx 4\\times10^{-5}\\,{\\rm M_\\odot\\,yr^{-1}}$, a value $\\sim$\\,100 times smaller than the accretion rate necessary to power the active nucleus. This can be understood as due to the fact that we are only seeing the hot ``skin'' (the H$_2$ emitting gas) of the total mass inflow rate, which is probably dominated by cold molecular gas. From the H$_2$ emission-line ratios we conclude that X-ray heating can account for the observed emission, but the H$_2\\,\\lambda2.1218\\,\\mu$m/Br$\\gamma$ line ratio suggests some contribution from shocks in localized regions close to the compact radio jet. ",
        "introduction": "The presence of supermassive black holes (SMBHs) at the centres of all galaxies which have stellar bulges is nowadays widely accepted by the astronomy comunity \\citep{gebhardt00,ferrarese00}. According to this scenario the energy emitted by an Active Galactic Nucleus (AGN) is due to the accretion of material onto the SMBH and implies the presence of a gas reservoir close to the AGN. \\citet{lopes07} using archival Hubble Space Telescope (HST) optical images for a large sample of  early-type galaxies  with and without AGNs, found that all AGN hosts have circumnuclear gas and dust, while this is observed in only 26\\% of a pair-matched sample of inactive galaxies.   A gas reservoir close to the AGN is also supported by the presence of recent star formation  in the circumnuclear region of active galaxies \\citep{schmitt99,boisson00,storchi-bergmann00,cid-fernandes01,cid-fernandes05,storchi-bergmann05}. However, the strongest signatures that feeding to the SMBH is occurring include the observation of streaming motions in ionized gas along nuclear spirals towards the nucleus of nearby active galaxies using two-dimensional (2D) optical spectroscopy \\citep{fathi06,storchi-bergmann07}.         Streaming motions as feeding signatures to the nuclear region have been previously observed in radio wavelengths. \\citet{adler96}, for example, have found streaming motions towards the center along the spiral arms of M\\,81 in H\\,{\\sc i}, while \\citet{mundell99} have found similar streaming motions towards the nucleus along the weak bar of NGC\\,4151. Closer to the center, most of the gas is in the molecular phase, and CO observations have been used to map the gas kinematics and inflows  \\citep[e.g.][]{garcia-burillo03,krips05,boone07}.  Molecular hydrogen emission lines are also relatively strong in the near-IR $K$-band spectra of active galaxies, and previous studies suggest that its distribution and kinematics is distinct from that observed in the other emission lines, which are usually dominated by outflows \\citep[e.g.][]{crenshaw00,das05}. In the Seyfert galaxies NGC\\,2110 and Circinus, for example, \\citet{storchi-bergmann99} have found broader emission-line profiles for [Fe\\,{\\sc ii}] and Pa$\\beta$ than for the H$_2\\lambda$2.1218$\\mu$m emission-line in a study using near-IR long-slit observations, suggesting a stronger influence of nuclear outflows on the former emission lines and a different origin for the H$_2$ emitting gas, consistent with the colder kinematics of the galaxy disk. More recently, using two-dimensional (hereafter 2D) near-IR spectroscopy of the Seyfert galaxy ESO\\,428-G14, we \\citep{riffel06} have found that the H$_2$ emission distribution was mostly restricted to the plane of the galaxy and was less affected by the AGN outflow than the [Fe\\,{\\sc ii}] and Pa$\\beta$ emission lines.  In this paper we use adaptive optics IFU spectroscopic data obtained with the Gemini's Near-infrared  Integral Field Spectrograph  \\citep[NIFS - ][]{mcgregor03}, in the near-IR $K$-band at a sampling of  0$\\farcs$1$\\times$0$\\farcs$1, to investigate the stellar and gaseous kinematics of the inner $\\sim 3 \\times 4$\\,arcsec$^2$ of the nearby active galaxy NGC\\,4051. The stellar kinematics will be used to constrain the galaxy potential and mass of the SMBH, but our main goal is to look for signatures of feeding mechanisms at parsec scales through the H$_2$ kinematics.    NGC\\,4051 is a SABbc galaxy at a distance of only $\\sim$9.3\\,Mpc\\,\\citep{barbosa06}, such that 1$\\farcs$0 corresponds to 45\\,pc at the galaxy. It harbors one of the closest AGN, classified as a Narrow-Line Seyfert 1 (NLS1). We have selected NGC\\,4051 for this study partially on the basis of the recent work of \\citet{barbosa06}, who obtained 2D stellar kinematics using the IFU of the Gemini Multi-Object Spectrograph (GMOS) to observe the stellar absorption lines of the Calcium triplet around 8500\\AA. These authors have  shown that the turnover of the stellar rotation curve occurs at only $R\\sim$50\\,pc from the nucleus, indicating that the stellar motions are dominated by a highly concentrated gravitational potential. As NIFS with the ALTAIR adaptive optics module provides a much better image sampling and resolution than the GMOS-IFU, we decided to further investigate the kinematics of the nuclear region of NGC\\,4051 in order to better constrain the gravitational potential and the mass of the SMBH. In adition, this galaxy shows strong H$_2$ emission \\citep{rogerio06}, making it a good candidate for a study of its kinematics.   Previous studies of NGC\\,4051 include HST narrow-band [O\\,{\\sc iii}] images which show an unresolved nuclear source and faint extended emission (by 1\\farcs2) along the position angle PA=100$^\\circ$ \\citep{schmitt96}, the approximate direction of the alignment of two radio components at 6\\,cm separated by 0$\\farcs$4 \\citep{ulvestad84}. \\citet{veilleux91} reported that the profiles of optical forbidden emission lines present blue wings reaching velocities of up to 800\\,km\\,s$^{-1}$, and proposed a model for the narrow-line region (hereafter NLR) with outflows and obscuring dust. Evidence for outflows have also been observed by \\citet{christopoulou97} for the [O\\,{\\sc iii}] emission to 1\\farcs5 NE of the nucleus. \\citet{nagao00} using the [Fe\\,{\\sc X}]$\\lambda$6374 emission line report a high ionization region extending to 3$\\farcs$0\\ SE of the nucleus.  \\citet{lawrence85} found a X-ray variability on time scales of the order of 1 hour and \\citet{salvati93} reported flux changes by a factor of 2 in 6 months observed at 2.2\\,$\\mu$m and suggest that emission by dust reprocessing of the UV radiation from the nucleus is an acceptable explanation. \\citet{ponti06} modelled the nuclear emission by two power law components, one due to the AGN and other due to reflection by the accretion disk.    This paper is organized as follows: in section 2 we describe the observations and data reduction. In section 3 we present the results for the stellar kinematics. In section 4 we present the emission-line flux distributions and in section 5 we present the results for the gas kinematics. In section 6 we discuss the results and in section 7 we present the conclusions of this work. ",
        "conclusions": "We have analysed two-dimensional near-IR $K-$band spectra from the inner $\\approx$\\,150\\,pc of the Narrow Line Seyfert 1 galaxy NGC\\,4051 obtained with the Gemini NIFS integral field spectrograph at a sampling of 4.5$\\times$4.5\\,pc$^2$ at the galaxy and spectral resolution of $\\approx$3\\,\\AA. We have mapped the stellar and gaseous kinematics, and the emission-line flux distributions and ratios of the molecular and ionized hydrogen. The main conclusions of this work are:  \\begin{itemize}   \\item The turnover of the stellar rotation curve is at only $\\approx$55\\,pc from the nucleus, suggesting that the stellar motions are dominated by a highly concentrated gravitational potential, a result confirmed by modeling using a Plummer gravitational potential, as we obtain a small scale length, $A\\approx$39\\,pc. This result supports the findings of \\citet{barbosa06} based of optical data. The mean velocity dispersion of the bulge ($\\sim$60\\,km\\,s$^{-1}$) implies a SMBH mass of $\\sim 10^6$\\,M$_\\odot$. Within the bulge, we find patches of lower velocity dispersions, which we attribute to recent star formation.  \\item The gas kinematics is distinct for each emission line: the Br$\\gamma$ emission line shows velocities restricted to within $\\sim$30\\,km\\,s$^{-1}$ from systemic, without evidence of bultk motions, suggesting that the ionized gas is not restricted to the galactic plane, while  much larger blueshifts and redshifts are observed for both the [Ca\\,{\\sc viii}] and H$_2$ emitting gas. The Br$\\gamma$ velocity dispersion is small ($\\sim$40\\,km\\,s$^{-1}$) over most of the field, except around the nucleus at a location close to the SW tip of the radio jet, suggesting some energy injection in the gas by interaction with the radio jet.  \\item The [Ca\\,{\\sc viii}] coronal emission line is compact but resolved, extending over a circular region of radius  $\\approx$35\\,pc around the nucleus. It also present the largest velocity blueshifts ($-250\\,$km\\,s$^{-1}$) and velocity dispersion (150\\,km\\,s$^{-1}$) of the emission lines studied here, supporting an origin close to the active nucleus, possibly in the transition region between the BLR and the NLR.  \\item The H$_2$ emitting gas seems to be mostly restricted to the galaxy plane. The most conspicuous kinematic features are a curved elongated blueshifted structure to the NE, interpreted as gas inflow along a nuclear spiral arm in the far side of the galaxy, and a curved, redshifted structure to the SW, interpreted as gas inflow along a nuclear spiral arm in the near side of the galaxy. Estimates of the mass inflow rate in hot H$_2$ gives $\\dot{M}_{H_2} \\approx 4\\times10^{-5}\\,{\\rm M_\\odot\\,yr^{-1}}$, which is $\\sim\\,100$ times smaller than the nuclear accretion rate necessary to power the active nucleus of NGC\\,4051, supporting the presence of additional inflow of cold non-emitting gas.  This is not the first time we have been able to map inflows towards an active nucleus along nuclear spiral arms. In two previous studies we \\citep{fathi06,storchi-bergmann07} have mapped streaming motions towards the active nucleus along nuclear spirals in the galaxies NGC\\,1097 and NGC\\,6951 although in ionized gas. As nuclear spirals are ubiquitous around active galactic nuclei \\citep{lopes07,prieto05}, such inflows seem to be an ``universal'' mechanism to bring gas inwards to feed the SMBH within the inner few hundred parsecs of the galaxy.   \\item The total mass of hot H$_2$ is estimated to be of the order of 66\\,M$_\\odot$, while that of H\\,{\\sc ii} is estimated to be 1.4$\\times\\,10^5$ M$_\\odot$.  \\item From the H$_2$ line ratios we conclude that H$_2$ is excited by thermal processes -- heating by X-rays from the AGN and shocks produced by the radio jet. We conclude, based in X-ray excitation models of \\citet{maloney96} that X-ray heating can account for the observed emission, but the H$_2\\,\\lambda2.1218\\,\\mu$m/Br$\\gamma$ line ratio supports some contribution from shocks in the regions where the radio jet interacts with the H$_2$ emitting gas.   \\end{itemize}"
    },
    "0801/0801.2674_arXiv.txt": {
        "abstract": "We first summarize work that has been done on the effects of binaries on theoretical population synthesis of stars and stellar phenomena. Next, we highlight the influence of stellar dynamics in young clusters by discussing a few candidate UFOs (unconventionally formed objects) like intermediate mass black holes,  $\\eta$ Car, $\\zeta$ Pup, $\\gamma$$^{2}$ Velorum and WR 140. ",
        "introduction": "\\medskip  Many (most?) of the massive stars are formed in clusters and the observations reveal that many cluster members are binary components (to illustrate, see Sana et al., 2007). The effect of binaries on population studies has been discussed many times by different research groups (the Brussels group was one of them) but the following points may be worth repeating:  \\medskip  \\begin{itemize} \\item binaries with initial primary mass larger than 40 $M_{\\odot}$ and with a period that is large enough so that LBV mass loss happens before the Roche lobe overflow (RLOF) would start, avoid RLOF (the LBV scenario as it was introduced in Vanbeveren, 1991); this may not be true for case A binaries \\item most of the primaries of binaries with orbital period as large as 10 yrs and with initial mass smaller than 40 $M_{\\odot}$ lose most of their hydrogen rich layers by RLOF; the evolution of these stars is therefore quite different from the evolution of single stars with the same mass \\item the RLOF in late case B and case C binaries may be accompanied by the common envelope process that may result in the merger of the two components; it can be expected that this merger will be a rapid rotator with peculiar chemical abundances \\item the RLOF in Case A and early case B binaries implies mass and angular momentum transfer towards the gainer; together with the merger process discussed above, this process is a natural way to make rapid rotators and therefore massive close binaries may be natural progenitors of long gamma ray bursts (Cantiello et al., 2007) \\item the favored model for short gamma ray bursts is the binary merger of two relativistic compact stars and it cannot be excluded that such mergers are important sites of galactic r-process enrichment (De Donder and Vanbeveren, 2003a) \\item\tthe observed massive binary frequency is not necessarily the binary frequency at birth. This is due to the fact that a significant fraction of all massive binaries become single stars during their evolution (many binaries with small mass ratio merge, late case B and case C binaries evolve through a common envelope phase and may merge as well, the supernova explosion of one of the components disrupts the binary), e.g. many single stars may have a binary past with an evolution which is distinctly different from stars that are born as single stars. \\end{itemize}  \\medskip  The proceedings of the conference on 'Massive Stars in Interacting Binaries' (eds. A. Moffat and N. St.-Louis) were published a few months ago. They can be considered as an update of the monograph 'The Brightest Binaries' published in 1998 (Vanbeveren et al. 1998c) and we recommend careful reading to anyone with an open mind willing to admit that many theoretical population studies, where the effects of binaries are ignored, have mainly an academic value. In section 2 we list a few extra references that may be interesting in order to learn more about the effects of massive binaries on population studies.  Many massive objects (single or binary) are born in dense (embedded) clusters, containing a few 10s up to a few 1000s massive stars and even more (e.g., the clusters in the Carina association, the Arches and Quintuplet clusters in the galactic bulge, the Orion nebula cluster, R136 in the LMC, MGG-11 in M82, etc.). In these clusters the evolution of each object may be affected by the presence of all the others, or at least by the closest neighbors. In sections 3 and 4 we discuss a few peculiar massive objects that may be the result of the dynamical evolution of stars in dense clusters, we like to call them Unconventionally Formed Objects (UFOs). ",
        "conclusions": ""
    },
    "0801/0801.0671_arXiv.txt": {
        "abstract": "We provide a new interpretation of ultraviolet transition region emission line widths observed by the SUMER instrument on the Solar and Heliospheric Observatory ({\\em SOHO}). This investigation is prompted by observations of the chromosphere at unprecedented spatial and temporal resolution from the Solar Optical Telescope (SOT) on {\\em Hinode} revealing that all chromospheric structures above the limb display significant transverse (\\alfvenic{}) perturbations. We demonstrate that the magnitude, network sensitivity and apparent center-to-limb isotropy of the measured line widths (formed below 250,000K) can be explained by an observationally constrained forward-model in which the line width is caused by the line-of-sight superposition of longitudinal and \\alfvenic{} motions on the small-scale (spicular) structures that dominate the chromosphere and low transition region. ",
        "introduction": "\\citet{Alfven1947} speculated that the broadening of coronal emission lines was the signature of \\alfven{} waves of sufficient strength to be a potential heating source for Edl\\'{e}n's recently observed hot solar corona \\citep[][]{Edlen1943}. The physical origin of the enhanced UV (and EUV) transition region (TR) and coronal line widths, and the amount of ``hidden'' energy that they represent, has been an open topic of debate, conflicting interpretation and conjecture in the community since. Interest in \\alfven{} waves in this regard is still strong \\citep[e.g.,][]{Chae1998, Erdelyi1998, Banerjee1998, Moran2003, Peter2003, OShea2005} owing to the recent numerical models of quiet coronal heating and fast solar wind acceleration depend critically on the presence of a significant energy carried by them \\citep[e.g.,][]{Suzuki2005, Cranmer2005, Verdini2007}. Only one thing has prevented the validation of the observational and theoretical inferences; the direct observation of \\alfven{} waves in the solar atmosphere. Recent unequivocal observations of low-frequency ($<$5mHz) propagating \\alfvenic{} motions in the solar corona \\citep[][]{Tomczyk2007} and chromosphere \\citep[][]{DePontieu2007b}, plus the relationship of the latter with finely structured spicules observed at the solar limb \\citep[][]{DePontieu2007a} with the Solar Optical Telescope \\citep[SOT;][]{SOT} on {\\em Hinode} \\citep[][]{Hinode}, have inspired us to look again at the relationship between non-thermal emission line widths and spatially unresolved \\alfvenic{} motions\\footnote{For a discussion on why we call these motions \\alfvenic{} and not MHD kink-mode waves, see the supporting online material of \\citet{DePontieu2007b}}. A reappraisal of TR line widths is clearly necessary. Much of the previous scientific effort in this area, prior to the launch of the Solar and Heliospheric Observatory \\citep[{\\em SOHO};][]{Fleck+others1995}, is expertly summarized in Ch.~5.2 of \\cite{Mariska1992}, but we will focus our discussion on the measurements of the Solar Ultraviolet Measurement of Emitted Radiation \\citep[SUMER;][]{Wilhelm1995} instrument.  We present a new analysis of SUMER full-disk spectroheliograms taken in early 1996 \\citep[e.g.,][]{Peter1999a,Peter1999b} paying close attention to the spatial variation of the line intensities and non-thermal widths. We compare the results of the data analysis with those of an observationally constrained forward model based on the relationship of chromospheric flows and \\alfven{} waves that occur on spatial structures that are considerably smaller than the SUMER spatial resolution (1\\arcsec). The synthesis of data and model analysis demonstrates that the superposition of transverse and longitudinal motions on sub-resolution structure can explain the observed magnitude, network sensitivity and center-to-limb isotropy of the measured non-thermal line broadening in cooler ($<$250,000K) TR lines. The last finding leads us to the conclusion that previous arguments made against the presence of  \\alfven{} waves in the TR, based on the apparent isotropy of the line width variation \\citep[e.g.,][]{Chae1998}, are invalid. ",
        "conclusions": "Our results indicate that the SUMER \\ion{C}{4} observations of line widths are indeed fully compatible with the LOS superposition of a large number of finely scaled spicules that carry significant longitudinal mass flows as well as vigorous \\alfven{} waves. Since the longitudinal and transverse velocity components are of roughly the same order, the overall center-to-limb variation is relatively limited, naturally producing an isotropy of line widths which had been used as an argument against \\alfven{} waves in the past \\citep[e.g.][]{Chae1998}. There is a slight but significant increase of line widths in a region close to and off the limb, which seems to be caused by \\alfvenic{} motions along the much increased number of structures (the first few 2\\arcsec{} pixels around the limb contain spicules from three neighboring network elements!). The steep drop-off of the line width as we go above a certain height above the limb is directly related to the dwindling number of spicular structures: spicules have typical heights of 5,000 km with very few above 10,000 km. This straightforward interpretation contradicts earlier work that explains this decrease as a result of ion-cyclotron damping of high-frequency \\alfven{} waves \\citep{Peter2003}. Our model also suggests that the increase of line widths in and around the magnetic network is directly related to the concentration of spicule flows along the field {\\em and} \\alfvenic{} motions.  We expect that similar results will hold for UV emission lines formed in the low TR (T$<$250,000K) although we note that opacity plays a significant role in the broadening of cooler lines \\citep[e.g.,][]{Doschek2004}. It is possible that \\ion{C}{4} has some opacity at the limb as well, explaining why the forward model predicts intensity increases at the limb that are twice those observed. % It is interesting that, for emission lines formed at temperatures above 500,000K (e.g., like those of \\ion{Ne}{8}), the bump at the limb does not appear, instead we observe a linear increase of line width with the decreasing density above the limb, indicative of the undamped growth of \\alfven{} waves in the extended corona \\citep[e.g.,][]{Banerjee1998}. Our model does not apply to these hotter lines, since the spatial distribution of emission from these lines is no longer dominated by spicules, but rather by thermal conduction and the overlying coronal structure.  Further study will be necessary to explain the larger (by $\\sim$5km/s) line widths in polar and equatorial coronal holes. These enhancements may well be caused by the lower densities and the resulting higher amplitude of the \\alfven{} waves (for an equal energy flux). Detailed studies with SOT are necessary to determine whether different amplitudes for the spicular flows perhaps also play a role, especially since the mix of spicule types changes considerably in coronal holes \\citep{DePontieu2007b}. We note that the absolute magnitude of the line width is of significant importance in constraining the amplitude of the \\alfven{} waves that are needed for solar wind models. Unfortunately, the literature reveals a wide scatter of values for the same TR UV emission lines \\cite [see, e.g.,][]{Mariska1992,Chae1998} from different instruments, epochs, locations, etc. This is perhaps not surprising since our model shows that we can expect significantly different values of line widths for the same line depending on proximity to the limb, the small-scale distribution of magnetic flux (which varies considerably over the solar cycle) and the orientation of the magnetic field (since inclined spicules will show varying contributions of longitudinal and transverse velocities). In addition to these physical causes, we have found that the fitting algorithm and form of functional fit (Gaussian with constant, linear or quadratic background) can have a significant impact ($\\sim 10$ km/s) on the magnitude of the line width; e.g., the more complex the background form of the fit the smaller the resulting line width.  Our forward model is only a first, primitive, step towards a comprehensive understanding of what dominates the emission of the TR. There has been extensive discusson on the detailed physical nature of the TR in the past \\citep[e.g.,][]{Judge1999,Peter2000a}, much of which has ignored the intrinsically dynamic nature of the TR \\citep[see, e.g.,][]{Wikstol1998}, which is inevitable given its connection to spicules. It is clear from our results that more sophisticated models of the TR are needed that take into account the intrinsically dynamic and finely structure nature of its dominant components."
    },
    "0801/0801.2168_arXiv.txt": {
        "abstract": "We summarize three recent publications which suggest that the Galactic center region Sagittarius B (Sgr B) may contain non-thermal radio components (Crocker et al. 2007, Hollis et al. 2007 and Yusef-Zadeh et al. 2007a). Based on new VLA matched-resolution continuum data at 327 MHz and 1.4 GHz, we find no evidence for large scale non-thermal radio emission at these frequencies; the spectral behavior is likely determined by the complex summation of multiple HII region components with a wide range of emission measures and hence radio turn-over frequencies.  Also, we discuss a possible additional interpretation of the radio continuum spectrum of individual component Sgr B2-F carried out by Yusef-Zadeh et al; confusion from nearby HII components with widely different turn-over frequencies may contribute to the the change in slope of the radio continuum in this direction at low frequencies. Finally, we discuss the uncertainties in the determination of the spectral index of the GBT continuum data of Sgr B carried out by Hollis et al. We find that the apparent spectral index determined by their procedure is also likely due to a summation over the many diverse thermal components in this direction. ",
        "introduction": "The radio spectrum of the Galactic Center (GC) has been a longstanding and complicated problem for astronomers.  The GC radio source was slowly recognized to be a complex of individual sources embedded in diffuse Galactic background emission. By 1965, when Bernard Burke's review article, entitled \"Radio Radiation from the Galactic Nuclear Environment\", appeared in the {\\it Annual Review of Astronomy and Astrophysics}, images of the GC region had been made at a range of frequencies as high as 14 GHz. One of the images in this review article was from Cooper \\& Price (1964), made with the Parkes, Australia, 210-foot telescope at 3 GHz. This image had a resolution of 6.7\\arcmin, making it possible to resolve the Sagittarius B (Sgr B) complex from the other major GC sources. Cooper \\& Price (1964) were the first to suggest that the Sgr B source is an optically-thick thermal source based on a spectral index measurement between 3 and 8 GHz (Drake 1959a,b) that had large uncertainties. During the last 40 years, the Sgr B region has been the subject of numerous detailed studies carried out with the increasing resolution and sensitivity of new instrumentation.  \\subsection{The SgrB Region: Sgr B2 and Sgr B1}  The Sgr B region is now understood to be one of the most complex star-forming regions in the Galaxy.  Located at $\\sim$105 pc in projection from Sgr A$^*$, the Sgr B giant molecular cloud has a total mass of 8-20 x 10$^6$ \\msun~(Tsuboi et al. 1999). Therefore, many interstellar molecules have been detected for the first time in the Sgr B cloud. Most studies of Sgr B have been carried out in the far-infrared and radio regimes because this region of the Galaxy is highly obscured (by up to 50 magnitudes) at visible wavelengths.  The Sgr B region is comprised of two distinct complexes:  Sgr B2 (G0.7-0.0; to the North) and Sgr B1 (G0.5-0.0; to the South). Both sources are strong radio continuum sources and both have been well-studied on fine scales ($\\sim$1\\arcsec~or better) with interferometric observations.  In higher frequency radio observations ($\\nu$ $>$ 1.4 GHz), Sgr B2 dominates due to the numerous ultra-compact, optically-thick HII regions (Benson \\& Johnston 1984; Gaume \\& Claussen 1990).  Sgr B2 is well-known for its incredibly high densities ($>$ 10$^5$ cm$^{-3}$) in both ionized and molecular gas (Huttemeister et al. 1993; DePree, Goss \\& Gaume 1998).  The Sgr B2 complex has three bright cores - Sgr B2 ``Main'', ``North'' and ``South''. Sgr B2 shows radio emission over a factor of 6000 in size scale, from $<$1\\arcsec~(ultra-compact HII regions) to the 10\\arcmin~extent of diffuse emission along the complex. The highest resolution radio study of Sgr B2 was made at 43 GHz with the Very Large Array (VLA)\\footnotemark\\footnotetext{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} radio telescope with an angular resolution of $\\sim$64 milli-arcseconds (600 AU at the GC distance of 8 kpc) in SgrB2 Main. This study revealed over 20 ultra-compact HII regions and stellar wind sources in Sgr B2 Main alone (De Pree et al. 1998). In addition, the range of observed emission measures in the ionized gas are from 1$\\times$10$^4$ to 1$\\times$10$^9$ pc cm$^{-6}$.  Radio recombination line studies at numerous frequencies have helped cement the thermal nature of many of the individual radio components in the Sgr B complex (e.g., Mehringer et al. 1993). The radio emission arising from the Sgr B1 complex indicates that this region is a more highly evolved star-forming region with numerous filamentary and shell-like ionized structures (Mehringer et al. 1992). Lying between Sgr B1 and Sgr B2 is a complex of radio sources known as G0.6-0.0. Gas velocities and morphology suggest that G0.6-0.0 physically connects Sgr B1 and Sgr B2 (Mehringer et al. 1992).  \\subsection{Recent Papers: Non-thermal Emission in Sgr B?}  Three recent papers on Sgr B complex have addressed the nature of the radio emission arising from this region, and have argued that there is a non-thermal component of the radio emission (Crocker et al. 2007; Hollis et al. 2007; and Yusef-Zadeh et al. 2007a). Because Sgr B is one of the most active regions of star formation in the Galaxy, containing more than 60 HII regions, it would not be surprising to find an embedded supernova remnant or two.  However, because the preponderance of evidence for the last four decades has indicated that the Sgr B complex is a thermal radio source, claims of non-thermal emission deserve close scrutiny.  One of the main motivations for looking for non-thermal emission has come from high-energy investigations of this region. Using HESS, Aharonian et al. (1996) have made a detection of diffuse $\\gamma$-ray flux between 0.2-20 TeV (1 TeV = 10$^{12}$ eV) from the Galactic center (GC) region, distributed along the Galactic plane. In addition, a separate HESS measurement of the $\\gamma$-ray flux and spectrum in a 0.5\\arcdeg$\\times$0.5\\arcdeg~region was centered on the SgrB molecular cloud. After removal of several bright point-like sources, the diffuse flux is believed to be correlated with the density of molecular gas (H$_2$) in the GC region, which has been imaged in the CS (J=1-0) line transition (Tsuboi et al. 1999). One possible origin of the high-energy radiation is the collision between cosmic ray protons and the ambient, dense molecular gas in the GC. Crocker et al. (2007) predict the radio synchrotron spectrum from their broadband (radio to $\\gamma$-ray) emission models and compare their results to the measured radio spectrum of this region.  Crocker et al. (2007) base their radio spectrum of this region primarily on low-frequency data at 327 and 843 MHz (see below for a description). They determine that the radio spectrum of Sgr B is non-thermal in nature, with a measured excess of non-thermal emission.  In addition, Yusef-Zadeh et al. (2007a) consider the global heating of molecular clouds by cosmic rays, based on their earlier observations of correlations between molecular gas and 6.4 keV K$\\alpha$ line emission (Yusef-Zadeh et al. 2007b). In order to illustrate that SgrB2 has a non-thermal component, Yusef-Zadeh et al. (2007a) use low frequency observations between 255 MHz and 1.4 GHz to show that a part of the SgrB2 complex (a well-studied part, SgrB2 ``F'', a cluster of ultra-compact HII regions) has a component of non-thermal emission with a flux density of $\\sim$82 mJy. They suggest the presence of enhanced heating by cosmic ray particles (which increases the ionization fraction) in the molecular gas.  Finally, Hollis et al. (2007) present a study of the continuum temperature of Sgr B2 based on GBT spectral line observations. They argue that there is a non-thermal component in Sgr B2 on size scales of $\\sim$143\\arcsec~at 1.4 GHz with an optically-thin spectral index of $\\alpha$=$-$0.7.  In this paper, a discussion of the measurements used to derive the spectral index of radio emission in the above papers follows. In addition, a high-resolution 1.4 GHz image of the GC region is presented and used for comparison to the three recent results on the non-thermal emission in Sgr B. ",
        "conclusions": "In this paper, we have summarized three recent papers which point out possible non-thermal radio emission arising from the Sgr B region in the GC. We also present a high-resolution and sensitive VLA image of the Sgr B region at 1.4 GHz. Using this image and a matched-array 327 MHz VLA image, we derive a thermal spectrum for the Sgr B complex and suggest that the radio emission is a mixture of optically thin and optically thick emission over the frequency range discussed here (255 MHz to 1.4 GHz). In addition, we show that the the apparent non-thermal power law slope for the Sgr B2 continuum temperature observed by the GBT is likely determined by source structure and provides limited information about the physical processes in the Sgr B region. While the structure Sgr B region is complex and furthermore confused by the Galactic background, there does not appear to be substantial evidence for a non-thermal component in the Sgr B complex.  The authors would like to thank Roland Crocker, Mike Hollis, and Farhad Yusef-Zadeh for their comments on this article."
    },
    "0801/0801.0995_arXiv.txt": {
        "abstract": "Large-scale structure formation, accretion and merging processes, AGN activity produce cosmological gas shocks. The shocks convert a fraction of the energy of gravitationally accelerated flows to internal energy of the gas. Being the main gas-heating agent, cosmological shocks could amplify magnetic fields and accelerate energetic particles via the multi-fluid plasma relaxation processes. We first discuss the basic properties of standard single-fluid shocks. Cosmological plasma shocks are expected to be collisionless. We then review the plasma processes responsible for the microscopic structure of collisionless shocks.  A tiny fraction of the particles crossing the shock is injected into the non-thermal energetic component that could get a substantial part of the ram pressure power dissipated at the shock. The energetic particles penetrate deep into the shock upstream producing an extended shock precursor. Scaling relations for postshock ion temperature and entropy as functions of shock velocity in strong collisionless multi-fluid shocks are discussed. We show that the multi-fluid nature of collisionless shocks results in excessive gas compression, energetic particle acceleration, precursor gas heating, magnetic field amplification and non-thermal emission. Multi-fluid shocks provide a reduced gas entropy production and could also modify the observable thermodynamic scaling relations for clusters of galaxies. ",
        "introduction": "\\label{Introduction}  The observed large scale structure of the Universe is thought to be due to the gravitational growth of density fluctuations in the post-inflation era. In this model, the evolving cosmic web is governed by non-linear gravitational growth of the initially weak density fluctuations in the dark energy dominated cosmology. The web is traced by a tiny fraction of luminous baryonic matter. Cosmological shock waves are an essential and often the only way to power the luminous matter by converting a fraction of gravitational power to thermal and non-thermal emissions of baryonic/leptonic matter.    At high redshifts ($z> 1100$) the pre-galactic medium was hot, relatively dense, ionised, with a substantial pressure of radiation. The cosmic microwave background (CMB) observations constrain the amplitudes of density inhomogeneities to be very small at the last scattering redshift $z \\sim$ 1000. Strong non-linear shocks are therefore unlikely at that stage. The universe expands, the matter cools, and eventually recombines, being mostly in neutral phase during the \"dark ages\" of the universe. At some redshift, $6 < z < 14$, hydrogen in the universe is reionised, likely due to UV radiation from the first luminous objects, leaving the intergalactic medium (IGM) highly reionised (see e.g. \\citet{Fan_ea06} for a recent review). The reionisation  indicates the formation of the first luminous objects at the end of the \"dark ages\", either star-forming galaxies or Active Galactic Nuclei (AGN). The compact luminous objects with an enormous energy release would have launched strong (in some cases, relativistic) shock waves in the local vicinity of the energetic sources. At the same evolution stage, formation of strong density inhomogeneities in the cosmic structure occurs. Since then the non-linear dynamical flows in the vicinity of density inhomogeneities  would have created large scale cosmic structure shocks of modest strength, thus heating the baryonic matter and simultaneously producing highly non-equilibrium energetic particle distributions, magnetic fields and electromagnetic emission.  Most of the diffuse X-ray emitting matter was likely heated by cosmological shocks of different scales. Accretion and merging processes produce large-scale gas shocks. Simulations of structure formation in the Universe predict that in the present epoch about 40~\\% of the normal baryonic matter is in the Warm-Hot Intergalactic Medium (WHIM) at overdensities $\\delta \\sim 5-10$ (e.g. \\citealt{CenO99,Dave_ea01}). The WHIM is likely shock-heated to temperatures of 10$^5 -- 10^7$~K during the continuous non-linear structure evolution and star-formation processes.   The statistics of cosmological shocks in the large-scale structure of the Universe were simulated in the context of the $\\Lambda$CDM-cosmology using PM / Eulerian adiabatic hydrodynamic codes (e.g. \\citealt{Miniati_ea00,Ryu_ea03,Kang_ea07}) and more recently with a smoothed particle hydrodynamic code by \\citet{Pfrommer_ea06}. They identified two main populations of cosmological shocks: $(i)$ high Mach number \"external\" shocks due to accretion of cold gas on gravitationally attracting nodes, and $(ii)$ moderate Mach number ($2 \\leq {\\cal M}_{\\rm s} \\leq 4$) \"internal\" shocks. The shocks are due to supersonic flows induced by relaxing dark matter substructures in relatively hot, already shocked, gas. The internal shocks were found by \\citet{Ryu_ea03} to be most important in energy dissipation providing intercluster medium (ICM) heating, and they were suggested by \\citet{Bykov_ea00cl} to be the likely sources of non-thermal emission in clusters of galaxies.  Hydrodynamical codes deal with N-body CDM and single-fluid gas dynamics. However, if a strong accretion shock is multi-fluid, providing reduced post-shock ion temperature and entropy, then the internal shocks could have systematically higher Mach numbers.  Space plasma shocks are expected to be {\\sl collisionless}. Cosmological shocks, being the main gas-heating agent, generate turbulent magnetic fields and accelerate energetic particles via collisionless multi-fluid plasma relaxation processes thus producing non-thermal components. The presence of these non-thermal components may affect the global dynamics of clusters of galaxies \\citet{Ostriker_ea05} and the $\\sigma_{\\rm v}$-$T$, $M$-$T$, $L_{\\rm X}$-$T$ scaling relations \\citet{Bykov05}. Detailed discussion of the cosmological simulations of the scaling relations with account of only thermal components can be found in \\citealt{13_borgani2008} - Chapter 13, this volume.  In Sect.~\\ref{colls} we discuss the basic features of the standard collisional shocks. The main part of the review is devoted to physical properties of cosmological shocks with an accent on collisionless shocks and associated non-thermal components. In Sect.~\\ref{CRS} we discuss the most important features of multi-fluid collisionless shocks in the cosmological context including the effects of reduced entropy production, energetic particle acceleration and magnetic field amplification in the shocks. ",
        "conclusions": "Cosmological shocks convert a fraction of the energy of gravitationally accelerated flows to internal energy of the gas. They heat and compress the gas and can also accelerate energetic non-thermal particles and amplify magnetic fields. We discussed some specific features of cosmological shocks.  $\\bullet$ The standard Rankine-Hugoniot relations based on the conservation laws for a steady single-fluid MHD shock allow to calculate the state of the fluid behind the shock once the upstream state and the shock strength are known. The coplanarity theorem for a plane ideal MHD shock states that the upstream and downstream bulk velocities, magnetic fields and the shock normal all lie in the same plane.  $\\bullet$ Cosmological plasma shocks are likely to be collisionless as many other astrophysical shocks observed in the heliosphere and in supernova remnants.  We review the basic plasma processes responsible for the microscopic structure of collisionless shocks.   $\\bullet$ Collisionless shock heating of ions results in a non-equilibrium  state just behind a very thin magnetic ramp region with a strongly anisotropic quasi-Maxwellian ion distributions. The possibility of collisionless heating of electrons by electromagnetic fluctuations in the magnetic ramp region depends on the extension of the fluctuation spectra to the electron gyro-scales, and could depend on the shock Mach number. Then the Coulomb equilibration processes are operating on the scales much larger than the collisionless shock width.   $\\bullet$ Extended MHD shock waves propagating in turbulent media could accelerate energetic particles  both by Fermi type acceleration in converging plasma flows and by DC electric field in quasi-perpendicular shocks. If the acceleration is efficient, then the strong shock could convert a substantial fraction (more than 10~\\%) of the power dissipated by the upstream bulk flow to energetic particles (cosmic rays). The compression ratio $r_{\\rm tot}$ at such a shock can be much higher, while the ion temperature behind the shock $\\propto r_{\\rm tot}^{-2}$ and the post-shock entropy are lower, than that in a standard single fluid shock. The shock structure consists of an extended precursor and a viscous velocity jump (subshock) indicated in Fig.~\\ref{sketch}.    $\\bullet$ Strong collisionless plasma shocks with an efficient Fermi acceleration of energetic particles could generate strong MHD waves in the upstream and downstream regions and strongly  amplify the upstream magnetic fields. A distinctive feature of the shock is a predicted possibility of gas pre-heating in the far upstream region due to MHD wave dissipation, that can produce an extended filament of temperature $\\gtrsim$ 0.1 keV.   $\\bullet$ Shock waves both from the cosmic web formation processes and those due to cluster merging activity can play an important role in clusters of galaxies. Direct evidences for such shocks, as traced by radio relics and the temperature jumps in X-ray observations havebeen found only in a small number of clusters, and thus we need more observations."
    },
    "0801/0801.0447_arXiv.txt": {
        "abstract": "We present equivalent width measurements and limits of six diffuse interstellar bands (DIBs, $\\lambda$4428, $\\lambda$5705, $\\lambda$5780, $\\lambda$5797, $\\lambda$6284, and $\\lambda$6613) in seven damped {\\Lya} absorbers (DLAs) over the redshift range $0.091 \\leq z \\leq 0.524$, sampling $20.3 \\leq \\log N({\\HI}) \\leq 21.7$. DIBs were detected in only one of the seven DLAs, that which has the highest reddening and metallicity.  Based upon the Galactic DIB--$N({\\HI})$ relation, the $\\lambda$6284 DIB equivalent width upper limits in four of the seven DLAs are a factor of 4-10 times below the $\\lambda$6284 DIB equivalent widths observed in the Milky Way, but are not inconsistent with those present in the Magellanic Clouds. Assuming the Galactic DIB--$E(B-V)$ relation, we determine reddening upper limits for the DLAs in our sample.  Based upon the $E(B-V)$ limits, the gas-to-dust ratios, $N({\\HI})/E(B-V)$, of the four aforementioned DLAs are at least $\\sim 5$ times higher than that of the Milky Way ISM. The ratios of two other DLAs are at least a factor of a few times higher.  The best constraints on reddening derive from the upper limits for the $\\lambda$5780 and $\\lambda$6284 DIBs, which yield $E(B-V) \\leq 0.08$ for four of the seven DLAs.  Our results suggest that, in DLAs, quantities related to dust, such as reddening and metallicity, appear to have a greater impact on DIB strengths than does {\\HI} gas abundance; the organic molecules likely responsible for DIBs in DLA selected sightlines are underabundant relative to sightlines in the Galaxy of similarly high $N({\\HI})$.  With regards to the study of astrobiology, this could have implications for the abundance of organic molecules in redshifted galaxies.  However, since DLAs are observed to have low reddening, selection bias likely plays a role in the apparent underabundance of DIBs in DLAs. ",
        "introduction": "Since their discovery by \\citet{hege22}, several hundred diffuse interstellar bands (DIBs) have been studied \\citep{gala00,jenn94,tuai00,wese00,hobbs07}, and yet no positive identifications of the carriers have been made.  The DIBs span the visible spectrum between 4000 and 13,000~{\\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[e.g.,][]{herb95, snow01, cox06a, hudg05}.  The organic-molecular origin of the DIBs may give them an importance to astrobiology; PAHs are now considered an important early constituent to the inventory of organic compounds on Earth \\citep{bada02}.  Via their infrared emitting vibrational bands, PAHs have been observed in high redshift dusty ultra-luminous infrared galaxies \\citep[ULIRGs, e.g.,][]{yan05}.  Searching for DIBs using the technique of quasar absorption lines provides a different approach for charting the presence of possible organics to high redshift.  As such, observing DIBs in high redshift galaxies may offer an independent method for constraining the environmental conditions in early-epoch galaxies governing the abundances of organic molecules, determining the cosmic epoch at which organic molecules first formed, and ultimately charting their evolution with redshift.  Aside from the hundreds of detections within the Galaxy \\citep[e.g.][]{gala00,jenn94,tuai00,wese00,hobbs07}, DIBs have been detected in the Magellanic Clouds \\citep{welt06,cox06b,cox07}, M31 \\citep{cord08}, seven starburst galaxies \\citep{heck00}, active galaxy Centaurus A via supernova 1986A \\citep{rich87}, spiral galaxy NGC 1448 via supernovae 2001el and 2003hn \\citep{soll05}, one damped {\\Lya} absorber (DLA) at $z=0.524$ toward the quasar (QSO) AO~0235+164 \\citep{junk04, lawt06}, and one $z=0.157$ {\\CaII} selected absorber toward QSO J0013--0024 \\citep{elli07}.  There are several environmental factors, such as {\\HI} column density \\citep{herb95,welt06}, reddening \\citep{welt06}, and metallicity and ionizing radiation \\citep{cox07}, that are related to DIB strengths. In the Galaxy, DIB absorption strengths correlate strongly with $N({\\HI})$ \\citep{herb95,welt06}.  However, in the Magellanic Clouds, DIBs are weaker by factors of 7-9 (LMC) and $\\sim20$ (SMC) compared to those observed in the Galaxy with similar $N({\\HI})$ \\citep{welt06}. It is not known whether other galaxies in the Local Group and beyond obey the Galactic DIB--$N({\\HI})$ relation, or, like the Magellanic Clouds, show departures from this relation.  Galaxies with high $N({\\HI})$ observed in absorption (i.e., DLAs) that reside at low to intermediate redshifts (where the prominent DIBs fall in the optical region) provide excellent astrophysical laboratories with which to investigate this issue.  In this paper we search for $\\lambda$4428, $\\lambda$5780, $\\lambda$5797, $\\lambda$6284, and $\\lambda$6613 DIB absorption in seven low to intermediate redshift DLAs.  In \\S~\\ref{sec:sample}, we give a brief summary of each intervening DLA in our sample.  In \\S~\\ref{sec:obs}, we discuss the spectroscopic observations and data reduction of the background QSOs.  In \\S~\\ref{sec:analysis}, we explain the procedure of our analysis, and the resulting spectra.  In \\S~\\ref{sec:results}, we present our results and compare our data to the Galactic DIB--$N({\\HI})$ relation, deduce upper limits for the reddening, $E(B-V)$, determine lower limits on the gas-to-dust ratios, and discuss the role of metallicity for our sample of DLAs.  We conclude in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we employ a generalized method of the \\citet{schn93} technique to find lines and determine conservative equivalent width limits for DIBs in seven DLAs along with an assessment of uncertainties in these limits.  We find:  (1) The $\\lambda$6284 DIB in four of the DLAs in our sample have equivalent width upper limits that are 4--10 times lower then expected for similar $N({\\HI})$ relative to Galactic sight lines. These limits are not inconsistent with the $\\lambda$6284 DIB strengths found in the LMC and SMC.  (2) Assuming the $\\lambda$5780 and $\\lambda$6284 DIB--$E(B-V)$ relations hold for DLAs, as it does for Galactic and Magellanic Cloud sightlines, we estimated upper limits on the reddening for our sample.  In four of our DLAs we estimate $E(B-V) \\leq 0.08$ with two DLAs having an $E(B-V) \\leq 0.05$.  These results are consistent with high redshift DLA samples, which have $E(B-V) < 0.04$ \\citep{elli05} and $E(B-V) < 0.02$ \\citep{murp04}.  (3) Applying the $E(B-V)$ limits, one of our DLAs is consistent with having the same or larger gas-to-dust ratio as the SMC. Two of our DLAs are consistent with having gas-to-dust ratios at least as large as sightlines in the LMC.  Three of our DLAs have less stringent limits that give lower limit gas-to-dust ratios consistent with Galactic, LMC, or SMC sightlines.  The AO~0235+164 DLA has a measured $E(B-V)$ and $N({\\HI})$ that puts its gas-to-dust fraction on the high end of the LMC sightlines as stated in other work \\citep{junk04}.  The $E(B-V)_{lim} \\leq 0.06$ constrained from the $\\lambda$6284 DIB is inconsistent with the known $E(B-V) = 0.23$ measured in \\citet{junk04}.  We should have been able to detect the $\\lambda$6284 DIB in the AO~0235+164 DLA given the limit adopted by \\citet{lawt06}.  Three possibilities are that our method of determining the $E(B-V)_{\\rm lim}$ does not apply to DLAs, our limits do not adequately take into account the large atmospheric absorption band (see Fig.~\\ref{s:0235}$b$), or that conditions are not favorable to the formation or survival of this DIB carrier.  It is interesting to speculate whether ionization conditions may be an important factor in inhibiting or enhancing DIB strengths. \\citet{welt06} test the ionization effects of DIBs along Galactic, LMC, and SMC lines of sight.  However, they do not find any significant trends.  On the contrary, \\citet{cox07} measure the UV radiation field along lines of sight toward the SMC and find UV radiation is an important environmental factor in DIB strengths. Laboratory spectroscopists claim that the DIBs may be due to partially ionized PAHs because they produce a wealth of absorption features in the optical spectrum \\citep{snow01}.  If this is the case then a significant UV radiation may be required.  However, UV radiation that is too high will dissociate the molecules.  In this work, we are unable to explore the affects of ionization conditions; knowledge of ionization conditions in specific DLAs is difficult to obtain.  It is also interesting to speculate whether metallicity may be an important factor in inhibiting or enhancing DIB strengths.  However, since we have robust metallicity measurements for only two of the DLAs in our sample, we cannot directly address the affects of metallicity. We do point out that obtaining deep spectra of DIBs in DLAs with known metallicity could be a fruitful future research direction, especially if reddening were also known. A direct comparison between the low metallicity Q1229--020 DLA and the high metallicity AO~0235+164 DLA might be fruitful as a first examination of the affects of metallicity in determining what governs DIB strengths in DLAs.  \\subsection{Implications of our Results}  Our results imply that reddening is a more crucial indicator of DIB strengths than is {\\HI} content for DLAs; the low reddening in DLA selected galaxies inhibits the presence of DIBs.  However, our sample is small and additional observations are required to ascertain the strength of this statement (metallicity and ionization conditions likely play a role as well).  The weakness of DIB strengths in our sample hints that the environments of high $N({\\HI})$ DLA-selected galaxies may be less suitable to create and/or sustain the organic molecules than those of the Galaxy.  Not only is the immediate solar environment beneficial for sustaining life, but it may be that the Galaxy is a more hospitable location for the survival of organic molecules that may have been the precursors to biology on Earth.  If these molecules are important as precursors to life in the universe, charting their presence to high redshift places constraints on how long ago and in which environments life could have potentially formed.  The presence of DIBs in the AO~0235+164 DLA demonstrates that organic molecules existed in at least one DLA selected environment at a redshift of $z \\sim 0.5$, or some 5 Gyrs ago \\citep{lawt06}.  Due to their weakness relative to DIBs in the Galaxy, observing DIBs in the general population of high redshift galaxies (as opposed to ULIRGs and star bursting galaxies) remains a challenge.  Selecting galaxies by DLA absorption may yet prove to be a lucrative method for detecting DIBs at high redshifts, provided the detection sensitivity can be increased.  However, most known DLAs reside at high redshift where the DIBs move into the near-IR, where high sensitivity spectroscopy is time intensive."
    },
    "0801/0801.2504_arXiv.txt": {
        "abstract": "In the framework of nonextensive statistical mechanics, the equilibrium structures of astrophysical self-gravitating systems are stellar polytropes, parameterized by the polytropic index $n$. By careful comparison to the structures of simulated dark-matter halos we find that the density profiles, as well as other fundamental properties, of stellar polytropes are inconsistent with simulations \\textit{for any value of $n$}. This result suggests the need to reconsider the applicability of nonextensive statistical mechanics (in its simplest form) to equilibrium self-gravitating systems. ",
        "introduction": "Nonextensive statistical mechanics is a generalization of thermodynamics and statistical mechanics proposed by C.\\ Tsallis in 1988 \\cite{tsallis88}, aiming to overcome the limitations of the Boltzmann entropy in its conventional applications \\cite{karlin02,taruya03b,chavanis05,gross05}. The theory has considerably widened the fields of application of statistical physics, allowing the description of systems affected by nonlocal effects, such as long-range forces and memory effects.  The study of astrophysical self-gravitating systems was one of the first applications of the theory \\cite{plastino93}. This approach has recently witnessed a renewed interest in the hope that it may provide a theoretical basis for the description of the universal structure of dark-matter (DM) halos \\cite{hansen05,leubner05,kronberger06,zavala06}. Nonextensive statistical mechanics predicts their equilibrium states to be stellar polytropes (SP) \\cite{plastino93,taruya03}, which have been claimed  to fit DM halos as well as the usual Navarro, Frenk \\& White model \\cite{kronberger06,zavala06} (hereafter NFW \\cite{NFW96}). Moreover, SP have the distinct advantage of being analytically derived from nonextensive statistical mechanics, while most models describing DM halos are empirical fits to N-body simulations. In this context, nonextensive statistical mechanics appears as an attractive framework for providing a theoretical understanding of the structure of DM halos and self-gravitating systems in general.  In this paper we compare, in a parameter independent way, the equilibrium configuration of astrophysical self-gravitating systems predicted by nonextensive statistical mechanics to simulated DM halos. We clarify the issue of which central boundary conditions to use for comparing SP to simulated halos.  On this basis, we establish that simulated DM halos do not corroborate the predictions of nonextensive statistical mechanics. This result calls into question the direct applicability of nonextensive statistical mechanics to equilibrium astrophysical self-gravitating systems. ",
        "conclusions": "We have established here that the predictions of nonextensive statistical mechanics are not corroborated by simulations of DM halos. Our results are based on readily observable quantities, i.e., the matter density itself, and profiles derived from it. These findings are in direct contrast to previous works which found reasonable agreement because they either considered only the outer parts of simulated halos \\cite{leubner05}, or used a non-zero density slope as initial condition near the center \\cite{kronberger06}, or used values of $n<5$ that lead to too steep density profiles in the outer parts \\cite{zavala06}. Support for our conclusions, however, comes from Barnes et al.\\ \\cite{barnes07} who, even though fixing a non-zero density slope near the center to match the NFW density profile, found inconsistency between SP and DM halo velocity dispersion profiles. While we are cautious about the physical relevance of such profiles, it is of interest to emphasize that nor cored neither cuspy polytropes can describe properly simulated DM halos.  These results imply a fundamental difference in the matter distribution of astrophysical self-gravitating systems predicted by nonextensive statistical mechanics and by N-body simulations. Such a discrepancy raises three main questions. Is our comparison with  simulated DM halos valid? Is nonextensive statistical mechanics the proper theory to describe collisionless long-range interaction systems at equilibrium? And how does this study relate to observed DM halos? We briefly address these issues below.  1a. We study idealized systems which do not take into account complex effects in DM halos such as velocity anisotropy \\cite{hansen06c} or triaxiality \\cite{Hayashi07}. However, the error we make by using isotropic and spherically averaged models to represent DM halos is small compared to the discrepancy we find between SP and simulated halos. Moreover, we do not address here general problems of statistical mechanics of self-gravitating systems like, e.g., infinite mass \\cite{hjorth91}, so as to focus on the issues specific to nonextensive statistical mechanics.  1b. N-body simulations depend on various parameters, such as the choice of the softening length and the number of particles, which can introduce numerical effects in the resulting halos. The softening of the gravitational force on small scale, introduced to avoid two-body interaction, creates a core at the center of the halo, approximately the size of the softening length. The number of particles fixes the maximum phase-space density resolved at redshift zero, and if not large enough, a core appears due to the lack of particles in the center. Both effects lead to a shallower density profile in the center of halos if the simulation is of low resolution \\cite{moore98}. Therefore, the discrepancy we observe between simulated profiles and SP is genuine and increases with the resolution of numerical simulations.  1c. We compare theoretical statistical predictions to cosmological simulations, but the latter depend on initial conditions and cosmological parameters. Initial conditions are fixed by the shape of the power spectrum of initial fluctuations $P(k) \\propto k^{n_s}$, i.e., by the choice of the index $n_s$. From CMB observations, we have an indication that  $n_s=0.958 \\pm 0.016$ \\cite{spergel07}, but tests on cosmological simulations proved that they depend very little on the choice of $n_s$ and on the cosmological model used \\cite{navarro97}. Therefore, the universal profiles of simulated halos can be compared to statistical theories as a general prediction.  1d. The choice of a scaling is necessary to compare DM halos models, but can bring a visualization bias. We checked if the use of $r_{-2}$ leads  us to overestimate the disagreement between nonextensive statistical mechanics and N-body simulations. However, adjusting the fit in the outer parts of the halo, e.g., scaling to the virial radius, increases even more the discrepancy in the inner parts, while scaling profiles to fit well at smaller radii makes the discrepancy appear in the outer parts of the halo.  2. Nonextensive statistical mechanics relies on three assumptions: \\textit{a)} the generalized entropy $S_q$ is the right form to describe long-range interaction systems; \\textit{b)} the system is at equilibrium, and its most probable state is given by the maximum entropy principle; \\textit{c)} $S_q$ is maximized at fixed energy, leading to a system whose distribution function depends only on the energy per unit mass, and has isotropic velocity dispersion. Though simulated halos show evidence of velocity anisotropy, the relation between density slope and velocity anisotropy \\cite{hansen06c} shows that it is isotropic in the center, where violent relaxation \\cite{lyndenbell67} (the process by which strong potential fluctuations efficiently drives the system towards equilibrium) is most effective. Therefore, the assumptions \\textit{b)} and \\textit{c)} hold true in the center.   However, it is in the inner parts that the disagreement between SP and simulated halos is the strongest. Hence, we suggest that the assumption \\textit{a)}, i.e., the choice of the generalized entropy $S_q$ motivated by nonextensive statistical mechanics, cannot be used to predict the equilibrium structures of simulated DM halos.  3. Simulated DM halos are in good agreement with observations \\cite{pointecouteau05}, except for the inner core found from spiral galaxy rotation curves \\cite{gentile04}. Stellar polytropes, interestingly,  have an inner core too, but the discrepancy we observe between SP and simulated DM halos is too large to explain the core of observed DM halos. Adding adiabatic contraction in simulations results in DM halos with steeper inner parts, which would not change our conclusions.  In summary, we have established that nonextensive statistical mechanics \\cite{tsallis88, plastino93, taruya03}, a theory generalizing classical statistical mechanics and thermodynamics, does not describe the equilibrium state of astrophysical self-gravitating systems, as represented by DM halos formed in N-body simulations."
    },
    "0801/0801.1111_arXiv.txt": {
        "abstract": "As first Paper of a series devoted to study the old stellar population in clusters and fields in the Small Magellanic Cloud, we present deep observations of NGC\\,121 in the F555W and F814W filters, obtained with the Advanced Camera for Surveys on the {\\it Hubble Space Telescope}. The resulting color-magnitude diagram reaches $\\sim 3.5$ mag below the main-sequence turn-off; deeper than any previous data.  We derive the age of NGC\\,121 using both absolute and relative age-dating methods. Fitting isochrones in the ACS photometric system to the observed ridge line of NGC\\,121, gives ages of $11.8 \\pm 0.5$~Gyr (Teramo), $11.2 \\pm 0.5$~Gyr (Padova) and $10.5 \\pm 0.5$~Gyr (Dartmouth). The cluster ridge line is best approximated by the $\\alpha$-enhanced Dartmouth isochrones.  Placing our relative ages on an absolute age scale, we find ages of $10.9 \\pm 0.5$~Gyr (from the magnitude difference between the main-sequence turn-off and the horizontal branch) and $11.5 \\pm 0.5$~Gyr (from the absolute magnitude of the horizontal branch), respectively. These five different age determinations are all lower by 2 -- 3 Gyr than the ages of the oldest Galactic globular clusters of comparable metallicity. Therefore we confirm the earlier finding that the oldest globular cluster in the Small Magellanic Cloud, NGC\\,121, is a few Gyr younger than its oldest counterparts in the Milky Way and in other nearby dwarf galaxies such as the Large Magellanic Cloud, Fornax, and Sagittarius.  If it were accreted into the Galactic halo, NGC\\,121 would resemble the ``young halo globulars'', although it is not as young as the youngest globular clusters associated with the Sagittarius dwarf.  The young age of NGC\\,121 could result from delayed cluster formation in the Small Magellanic Cloud or result from the random survival of only one example of an initially small number star clusters. ",
        "introduction": "Characterizing old stellar populations provide important constraints on the early star formation histories of galaxies. Only the satellite galaxies of the Milky Way (MW) are sufficiently close to resolve individual stars well below the oldest main-sequence turn-offs, which is a pre-condition for accurate photometric age dating of old stellar populations. All Local Group galaxies, for which adequate data exist, appear to contain stars older than 10~Gyr \\citep{grebel04}. This result is based on main-sequence turn-off photometry of globular clusters and field populations in Galactic satellites and a few more distant Local Group galaxies \\citep[e.g., ][]{brown07, cole07}, as well as the detection of horizontal branch stars (including RR Lyrae variables) in the Local Group and beyond \\citep[e.g., ][] {held00,harb01, saraj02, clementini03, pritzl04}.  Globular clusters are preferred as the basis for old stellar population age tracers since they are usually single-age, single-metallicity objects facilitating comparative studies. Moreover, while globular cluster systems exhibit a range of ages \\citep[e.g., ][]{deang05}, the oldest ones may belong to the most ancient surviving stellar systems to have completed their formation in the youthful Universe \\citep[e.g., ][]{moore06}.  In those nearby galaxies where relative age dating based on main-sequence photometry was carried out in comparison to the oldest globular clusters in the Milky Way, no age difference within the measurement accuracy was found \\citep[e.g., ][ and references therein]{grebel04, brown07, cole07}. The relative age dating of the oldest identifiable Population~II objects thus indicates a common epoch of substantial early star formation in the Milky Way and its companions, although information about a putative, even older Population III remains to be uncovered in these objects.  A galaxy that may {\\em not}\\ share this common epoch of early star formation -- at least not with respect to its globular clusters \\citep[e.g., ][]{saraj98} -- is the Small Magellanic Cloud (SMC)\\footnote{There may be additional exceptions in more distant dwarf irregular galaxies regarding the common epoch of earliest Population~II star formation, although also these galaxies evidently contain old populations \\citep[e.g., ][]{grebel01, maka02}}. The SMC is one of the closest and therefore best studied dwarf galaxies orbiting our Galaxy.  While the SMC hosts a large number of intermediate-age and young star clusters, it only contains one ''old'' globular cluster, NGC\\,121, which is also the most massive star cluster. NGC\\,121 is located $\\sim$ 2.4\\degr ($\\sim$3~kpc) west of the SMC bar at ($\\alpha_{J2000.0}$, $\\delta_{J2000.0}$) = ($0^h26^m47.0^s$, $-71\\arcdeg32'12.0''$).  NGC\\,121 is the only cluster in the SMC that is sufficiently old to have developed an extended red horizontal branch \\citep{stry85} and to contain RR Lyrae stars. Indeed, whether or not to call a star cluster a globular cluster is a matter of definition. In this case we refer to \\citet{salgir02} who consider Lindsay\\,1 as having a stumpy red clump and not a red horizontal branch. Three RR Lyrae stars were discovered in NGC\\,121 by \\citet{thack58}. \\citet{grah75} found a fourth RR Lyrae variable in the cluster and an additional 75 in a 1 $\\times$ 1.3 square degree field centered on NGC\\,121. Studies of various clusters in the LMC and in the MW showed that the presence of RR Lyrae variables indicates that the parent population is as old as or older than $\\sim$ 10~Gyr.  An important question is whether NGC\\,121 is as old as the typical old globular clusters in the Large Magellanic Cloud (LMC) and in the MW. Previous studies found ages ranging from 8 to 14~Gyr for NGC\\,121 \\citep{stry85, Walker91, migh98, udal98, shara98, dol01} using a variety of different techniques.  Studies based on the deepest available color-magnitude diagrams from Hubble Space Telescope (HST) observations with the Wide Field and Planetary Camera 2 (WFPC2) indicate an age of 10 to 10.6~Gyr for NGC\\,121, suggesting that this globular cluster is several Gyr younger than the oldest globulars in other nearby galaxies and in the MW \\citep{shara98, dol01}.  The capabilities of the Advanced Camera for Surveys (ACS) provide an improvement in both sensitivity (depth) as well as angular resolution, which is essential for a reliable photometric age determination in this dense star cluster. Here we present deep photometry of NGC\\,121 obtained with ACS aboard the HST.  We determine the age of NGC\\,121 utilizing both absolute and relative methods \\citep[e.g., ][]{chab96Sci}. The current study is the first in a series of papers based on HST studies of rich intermediate-age and old star clusters in the SMC.  In addition to NGC\\,121, six intermediate-age SMC star clusters have been observed as part of our program: Lindsay\\,1, Kron\\,3, NGC\\,339, NGC\\,416, Lindsay\\,38 and NGC\\,419. We will derive fiducial ridgelines and fit isochrones to obtain accurate ages for each cluster using the same reduction techniques and isochrone models as described here (see $\\S$~\\ref{sec:obs}-\\ref{sec:age}), and will present our results in future papers. In Table~\\ref{tab:journalobs} we list the cluster identification, date of observation, passband, exposure times and location of all clusters in our HST program (GO-10396; principal investigator: J.~S.~Gallagher).  In the next Section we describe the data reduction procedure. In $\\S$~\\ref{sec:CMD} we present the color-magnitude diagram (CMD) of NGC\\,121 and discuss its main features.  In $\\S$~\\ref{sec:age} we describe our age derivation methods and present our results. ",
        "conclusions": "We derived ages for the old SMC globular cluster NGC\\,121 based on our high dynamic range HST/ACS photometry that extends at least three magnitudes below its MSTO.  In order to obtain absolute ages, we applied three different isochrone models. These isochrone models yielded ages of $11.2 \\pm 0.7$~Gyr (Padova), $11.8 \\pm 0.5$~Gyr (Teramo), and $10.5 \\pm 0.5$~Gyr (Dartmouth).  We find the $\\alpha$-enhanced Dartmouth isochrones provide the closest approximation to the MS, SGB, and RGB, whereas the other models cannot reproduce the slope of the upper RGB. High-resolution spectroscopy indicates that NGC\\,121 is indeed $\\alpha$-enhanced \\citep{john04}, a property that it shares with many of the old outer Galactic halo globulars. Given the proximity of NGC\\,121 to the SMC on the sky and its distance, its physical association with the SMC seems well-established.  Our determinations of relative ages for NGC\\,121 are consistent with the results of our absolute age determination. Relative age estimates, when converted to an absolute age scale, are $10.9 \\pm 0.5$~Gyr ($\\Delta V_{TO}^{HB}$), $10.8 \\pm 1.0$ ($M_V(HB)$) and $11.5 \\pm 0.5$~Gyr ($M_{V(BTO)}$). These numbers agree well with the absolute age derivations. Our results confirm that NGC\\,121 is 2--3~Gyr younger than the oldest MW and LMC clusters (as also found in earlier WFPC2 studies).  NGC\\,121 is similar in age to the youngest globular cluster in the Fornax dSph \\citep{buon99}, and to several of the young Galactic halo clusters. On the other hand, NGC\\,121 is not as young as some of the Sgr dwarf galaxy's globular clusters or the youngest Galactic globular clusters.  It is intriguing that the SMC -- in contrast to other Galactic companion dwarf galaxies with globulars -- does not contain any old classical globular clusters. But given the existence of only one cluster and the question of star cluster survival, this could be a result of the one survivor from the SMC's epoch of globular cluster formation randomly sampling an initial distribution of star cluster ages.  On the other hand, in low-mass galaxies without  bulges, spiral density waves, and shear it is much more difficult to destroy globular clusters through external effects. That this cluster is both younger than the Galactic mean and enhanced in $\\alpha$-elements may have interesting implications for the early development of the SMC.  It also is intriguing that the only globular cluster in the SMC is not very metal-poor.  The SMC must have experienced substantial enrichment prior to the formation of NGC\\,121.  In the LMC, where two main epochs of the formation of populous compact star clusters have been found \\citep[e.g., ][]{Bertelli92}, a few globular clusters are found that are old enough to exhibit blue HBs. Interestingly, these globular clusters, which are similarly old as the oldest Galactic globulars \\citep{olsen98}, have a similar metallicity to NGC 121 \\citep{john04} (e.g., NGC\\,1898, NGC\\,2019), indicating very early chemical enrichment.  The MW also contains old classical globular clusters (with blue HBs) that have similarly high metallicities as the somewhat younger NGC 121. Evidently, the conditions for and the efficiency of star formation varied in these three galaxies at early epochs.  After NGC 121 formed there was a hiatus in surviving stars clusters and thus possibly in cluster formation activity in the SMC: The second oldest SMC cluster is Lindsay 1 with an age of $\\sim$ 8~Gyr (Glatt et al. 2007, in preparation). Since then compact populous star clusters formed fairly continuously until the present day in the SMC (e.g., Da Costa 2002) -- in contrast to both the LMC and the MW. In forthcoming papers on our ACS photometry of SMC clusters and field populations we will explore the evolutionary history of the SMC in more detail.  Clearly, clues about the early star formation history of the SMC will have to come from its old field populations."
    },
    "0801/0801.4734_arXiv.txt": {
        "abstract": "The observed  velocities of pulsars suggest the possibility that sterile neutrinos with mass of several keV are emitted from a cooling neutron star. The same sterile neutrinos could constitute all or part of cosmological dark matter.  The neutrino-driven kicks can exhibit delays depending on the mass and the mixing angle, which can be compared with the pulsar data.  We discuss the allowed ranges of sterile neutrino parameters, consistent with the latest cosmological and X-ray bounds, which can explain the pulsar kicks for different delay times. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.3863_arXiv.txt": {
        "abstract": "Supergiant Fast X-ray Transients are obviously related to persistent Supergiant X-ray Binaries. Any convincing explanation for their behaviour must consistently take into account all types of X-ray sources powered by wind accretion. Here we present a common framework for wind accreting sources, within the context of clumpy wind models, that allows a coherent interpretation of their different behaviours as an immediate consequence of diverse orbital geometries. ",
        "introduction": "Traditionally, High Mass X-ray binaries (HMXBs) have been divided in three classes: ``classical'' bright ($L_{{\\rm X}}\\sim10^{38}\\:{\\rm erg}\\,{\\rm s}^{-1}$) sources fed by incipient localised Roche lobe overflow, Be/X-ray binaries, presenting different behaviours but always associated to the presence of a circumstellar decretion disk, and wind-fed Supergiant X-ray Binaries (SGXBs). In the last group, the compact object is fed by accretion from the strong radiative wind of an OB supergiant, leading to a persistent X-ray source with $L_{{\\rm X}}\\sim10^{36}\\:{\\rm erg}\\,{\\rm s}^{-1}$, displaying large variations on short timescales, but a rather stable flux in the long run. The compact object is a relatively close ($P_{{\\rm orb}}\\la15\\:{{\\rm d}}$) orbit with low or negligible eccentricity.    Over the last few years, there has been a huge increase in the number of HMXBs known \\citep[e.g.,][]{wal06}. Many new sources have been found to display very brief outbursts, with a rise timescale of tens of minutes and lasting only a few hours \\citep{sgue05}. Several have been identified with OB supergiants: \\citep[e.g., the prototype XTE~J1739$-$302;][]{smith06,neg06b}. The distances to their counterparts imply typical $L_{{\\rm X}}\\sim10^{36}\\:{\\rm erg}\\,{\\rm s}^{-1}$ at the peak of the outbursts. A number of other systems have been shown to have similar X-ray behaviours leading to the definition of a class of Supergiant Fast X-ray Transients (SFXTs) \\citep{esa,smith06}.  Around 12 such systems are known \\citep{wzh07}. It is perhaps possible to divide them in two groups, the first one characterised by very low quiescence $L_{{\\rm X}}$ and high variability and the second one with higher average $L_{{\\rm X}}$ and smaller variability factors. The second group, then, could be thought of as persistent SGXBs with an average $L_{{\\rm X}}$ below the canonical $\\sim10^{36}\\:{\\rm erg}\\,{\\rm s}^{-1}$ over which flares are superimposed \\citep{wzh07}. To some degree, this distinction may be due to observational biases (cf. the case of XTE~J1739$-$302; Blay et al., submitted). In any case, it is clear that the separation between SFXTs and SGXBs is not well defined.  Two other systems have shown flaring only during episodes separated by regular intervals: IGR J00370+6122 \\citep{zand07} and IGR J11215-5952 \\citep{sidoli}. This modulation is most likely related to an orbital period. Because of their very different properties, we do not count these two systems as SFXTs, but they are also wind-fed accretors and any theory attempting an explanation to the behaviour of SFXTs must take into account the existence of other wind accretors.  At least two SFXTs have been observed to increase their $L_{{\\rm X}}$ by a factor $>100$ (going from deep quiescence to outburst) in only a few minutes: XTE~J1739$-$302 \\citep{sak02} and IGR~J17544$-$2619 \\citep{zand05}.  AX~J1841.0$-$0535 showed a $>10$ increase in count-rate over $\\la 1\\:$h \\citep{bam01}. Such sharp rises are incompatible with any explanation of their behaviour based solely on orbital motion through a smooth medium. ",
        "conclusions": ""
    },
    "0801/0801.3264_arXiv.txt": {
        "abstract": "We present an abundance analysis based on high dispersion and high signal-to-noise ratio Keck spectra of a very highly microlensed Galactic bulge dwarf, \\mystar, with \\teff $\\sim 5400$~K.  The amplification at the time the spectra were taken ranged from 350 to 450. This bulge  star is highly enhanced in metallicity with [Fe/H]\\footnote{We adopt the usual spectroscopic notations that [A/B] ~ $\\equiv ~ log_{10} (N_A/N_B)_* - log_{10} (N_A/N_B)_{\\odot}$, and that log$[\\epsilon(A)] ~  \\equiv ~ log_{10} (N_A/N_H) + 12.00$, for elements $A$ and $B$.}  = \\fehogle ~dex.  The abundance ratios for the 28 species of 26 elements for which features could be detected in the spectra are almost all solar.  In particular, there is no evidence for  enhancement of any of the $\\alpha$-elements including O and Mg. We conclude   that the high [Fe/H] seen in this star, when combined with the equally high [Fe/H] derived in previous detailed abundance analysis of two other Galactic bulge dwarfs, both also highly magnified by microlensing, implies that the median metallicity in the Galactic bulge is very high. We thus infer that many previous estimates of the metallicity  distribution in the Galactic bulge have substantially underestimated the mean Fe-metallicity there due to sample bias, and suggest a  candidate mechanism for such. If our conjecture proves valid, it may be necessary to update the calibrations for the algorithms used by many groups to interpret   spectra and broad band photometry of the integrated light of very metal-rich old  stellar populations, including luminous elliptical galaxies. ",
        "introduction": "Microlensing occurs when a ``lens'' (star, planet, black hole, etc) becomes closely aligned with a more distant ``source'' star, whose image it both magnifies and distorts.  Normally the observer is most interested to learn about the lens, but microlensing data can simultaenously serve as a powerful probe of the source.  If the source transits the lens (or the ``caustics'' generated by the lens, in the case of binary lenses), then it resolves the source, allowing detailed limb-darkening profiles from a photometric times series \\citep{fields03} or even spectral resolution of surface features like the chromosphere \\citep{cassan04} from judiciously taken spectra.  The lens can also serve simply to amplify the role of the telescope as a ``light bucket''. \\citet{minniti98} provocatively titled their report on observations of a bulge dwarf that was magnified by a factor $A=2.25$ ``Using Keck I as a 15m Diameter Telescope'' and \\cite{cavallo} presented a preliminary abundance analysis for six stars with small magnification ($2.5 < A < 30$), including this star and two other bulge dwarfs. Obviously, this trick could in principle be improved to arbitrary ``diameters'' simply by observing the events at higher magnification.  The problem is that it is extremely difficult to recognize high-magnification events in advance, and harder still to activate large-telescope observations in time to take advantage of them.  The process has been facilitated by microlensing planet hunters, who prize high-magnification events because of their extreme sensitivity to planets \\citep{ob05071,ob05169,ob06109}.  \\citet{johnson07} were the first to piggy-back on the microlensing planet hunters, obtaining a 15 minute Keck spectrum of the Galactic bulge dwarf \\johnstar\\ at magnification $A=135$.  This star, the first bulge dwarf with a high-quality spectrum, proved to be extremely metal-rich, a fact that might be telling us that the bulge is much more metal-rich than seems indicated by available spectra of giants, but could also just be the ``luck of the draw''.  The \\citet{johnson07} results therefore substantially raised the premium on obtaining highly-magnified spectra of bulge dwarfs.  On 2 July 2007, the OGLE collaboration\\footnote{http://www.astrouw.edu.pl/$\\sim$ogle/ogle3/ews/ews.html} \\citep{ews} announced OGLE--2007--BLG--349\\ (RA=18:05:24.43; DEC = $-26$:25:19.0) at Galactic coordinates $(l,b)=(4.4,-2.5)$, i.e., 5.1$^{\\circ}$ from the Galactic center, as a probable microlensing event.  The Microlensing Follow Up Network\\footnote{http://www.astronomy.ohio-state.edu/$\\sim$microfun/} ($\\mu$FUN) began monitoring the event on 18 August to determine whether the event would be high-magnification and on 3 September issued a general alert that it would reach at least $A>200$ two nights hence.  On this basis, $\\mu$FUN organized world-wide photometric observations, whose outcome will be reported elsewhere \\citep{dong08} and also contacted JGC at the Keck telescope to recommend intensive observations.   The ability to obtain high resolution, high quality spectra of Galactic bulge stars and to carry out a detailed abundance analysis offers an unbiased way to determine the metallicity distribution of stars in the Galactic bulge, as well as their detailed chemical inventory.  Our abundance analysis of \\mystar\\ is described  in the first few sections of the paper, with the key results presented in \\S\\ref{section_abund_results}. In an effort to explain our rather surprising results, we indicate in \\S\\ref{section_discussion} how past studies of the brightest giants, which are the only bulge stars for which detailed abundance analyses can be derived from spectra obtained under normal conditions, may be subject to previously ignored selection effects, and how this might impact studies of the integrated light for metal-rich old simple stellar populations. Abundance ratios in the Galactic bulge giants and in the microlensed dwarfs are discussed in \\S\\ref{section_abund_ratios}. A brief summary concludes the paper, while an  appendix discusses the behavior of selected diffuse interstellar bands in the spectrum of \\mystar. ",
        "conclusions": "}    An early attempt to determine the metallicity distribution in the Galactic bulge was that of \\cite{sadler96}, who used  low resolution spectroscopy for 268 bulge giants and red clump stars to derive a mean [Fe/H] of $-0.11\\pm0.04$~dex.  \\cite{ramirez00} studied a sample of M giants in the near-IR; they found a very similar mean [Fe/H] of $-0.21$~dex with a dispersion of 0.30~dex. \\cite{zocalli} used extensive optical and near-IR photometry with CMD fitting; they derived a somewhat lower mean [Fe/H]. In all these cases, the calibration of the metallicity scale relied on Galactic globular clusters. \\cite{zocalli} suggest that the differences between several of these studies depend crucially on the abundances adopted for the two highest metallicity GCs with high resolution abundance analyses, NGC~6553 and NGC~6558.  [Fe/H] values for these two GCs have ranged over more than 0.5~dex in the literature, but since the work of \\cite{cohen99} and \\cite{carretta01}, who suggested values higher  than most previous studies, more recent analyses have settled toward the higher values, see, e.g., \\cite{zoccali04} and \\cite{vlt_6553}.  An early high dispersion spectroscopic study of Galactic bulge K giants is that of \\cite{mcwilliam94}, who found a mean [Fe/H] in their sample of $-0.25$~dex. \\cite{fulbright06} updates and expands upon this earlier work.  They then use their detailed abundance analyses of 27 K giants in Baade's window to recalibrate the metallicities for the much larger samples of \\cite{sadler96} and of \\cite{rich88}.  The mean [Fe/H] they thus deduce is $-0.10\\pm0.04$~dex.  The median [Fe/H] of their sample is also sub-solar.  As infrared echelle spectrographs have become available on 8-m class telescopes, high dispersion studies of Galactic bulge giants in the near-IR have become possible, see, e.g., \\cite{rich05}, \\cite{cunha06} and \\cite{rich07}.  The mean [Fe/H] from the  sample of M giants studied by \\cite{rich07} in the Galactic bulge is well below solar, with [Fe/H] $-0.22\\pm0.01$~dex.  The median Fe-metallicity of the small sample of giants with near-IR spectra analyzed by \\cite{cunha06} is  also slightly  below solar metallicity.  Since this sample largely overlaps that of \\cite{fulbright06} it is not included in the figures.  Very recently, \\cite{zoccali08} have presented initial results of a survey of Fe-metallicity in the Galactic bulge from spectra with $\\lambda/\\Delta\\lambda = 20,000$ of about  800 stars. They find a radial gradient in [Fe/H] within the bulge with the mean value going from +0.03~dex at $b = -4^{\\circ}$ to $-0.12$~dex at $b = -6^{\\circ}$, and a sharp cutoff towards higher metallicities.  All of these samples of Galactic bulge stars are of luminous giants and/or of red clump stars.  They all have mean and median [Fe/H] values that are slightly sub-solar and are similar to the mean [Fe/H] of $-0.1$~dex found for local disk stars by \\cite{localdisk}.  Yet \\mystar, analyzed here, and \\johnstar\\ \\citep{johnson07}, another Galactic bulge microlensed dwarf with a high quality detailed abundance analysis, both have [Fe/H] $\\sim +0.5$~dex; the third such star found to date, \\moa\\, with a lower signal-to-noise ratio spectrum, analyzed by \\cite{johnson08}, has a somewhat lower Fe-metallicity, [Fe/H] = $+0.36\\pm0.18$~dex. The comparison between the [Fe/H] distribution for the recalibrated sample of \\cite{sadler96} by \\cite{fulbright06} and the [Fe/H] values deduced for these three Galactic bulge dwarfs is shown in Figure~\\ref{figure_feh_hist}.  While three stars is an uncomfortably small sample, it is difficult to believe that this is consistent with the published metallicity distributions for the Galactic bulge giants and/or red clump stars. The probability that the sample of three stars would have such high [Fe/H] values  by chance is less than 1\\%  given the  [Fe/H] distributions of the larger of the relevant  studies in the Galactic bulge claiming to have unbiased samples, including \\cite{ramirez00} with 110 M giants, which must be considered unbiased at the high metallicity end, \\cite{sadler96} (268 bulge stars),  and the recalibrated verion of the latter by \\cite{fulbright06}.  A KS test  shows that the probability of finding the three carefully studied microlensed bulge dwarfs at their very high  [Fe/H] values if the underlying metallicity distribution is that found by \\cite{zoccali08} from high dispersion spectra of a very large sample of giants in three bulge fields is less than 10$^{-4}$.  This calculation takes into account the radial gradient in the metallicity distribution function found by \\cite{zoccali08}.     We next assemble the evidence that \\mystar\\ is actually a dwarf in the Galactic bulge. We have already demonstrated that the microlensed bulge dwarf studied here has stellar parameters that, when combined with the unlensed magnitude from the mirolensing light curve, yield an unlensed luminosity consistent with that expected for a dwarf of its \\teff\\ at the distance of the Galactic center.  This  is true  to within 0.2~mag for the other two such stars as well \\johnstar\\ and \\moa\\ \\citep{johnson07,johnson08}. The three microlensed bulge dwarfs studied to date come from a kinematically hot population; they have radial velocities of +99, $-$154 and +113 \\kms\\ with a typical uncertainty of less than 2~\\kms, thus showing a dispersion in $v_r$ consistent with that inferred from large samples of bulge giants \\citep[see, e.g.][]{sadler96,zoccali08}.  About 4\\% of the dwarfs in the solar neighborhood from the proper motion sample of \\cite{grenon89} appear to be old and very metal-rich, with [M/H] $> +0.30$~dex.  These stars are almost all on eccentric orbits with small pericenters ($\\lesssim 3$~kpc). \\cite{castro97} and \\cite{pompeia03} have carried out detailed abundance analyses of some of these stars to find that the most metal rich of them reach [Fe/H] +0.55~dex, and have solar abundance ratios in general. These local dwarfs are thus similar in their abundances to \\mystar.  The references cited suggest that they are possibly stars on chaotic orbits ejected from the Galactic bar or older central regions of the Galactic disk. Little is known about the inner regions of the Galactic disk, whether it exists at all in the inner kpc, what its scale height might be should it exist, etc. The model of \\cite{dirbe_phot}, constructed to reproduce the surface brightness within 5~kpc of the Galactic center seen by COBE/DIRBE at 3.5 and at 4.9~$\\mu$ after correction for extinction, suggests a double exponential disk with a scale heights of 42 and of 210~pc (1.5$^{\\circ}$ at a distance of 8~kpc to the Galactic center) combined with a truncated power-law bulge.  However, the beam size of COBE/DIRBE is $0.7^{\\circ} \\times 0.7^{\\circ}$, and this was smoothed to a 1.5$^{\\circ}$ angular resolution for their analysis. The scale height of 210~pc thus corresponds to their minimum angular resolution, and may well be smaller, which would reduce any potential disk contamination of bulge samples.  In any case, using their model as a guide, we find that the probability of contamination by disk stars of a sample at $b = 2.5^{\\circ}$ in this maximal case is only $\\sim$30\\% larger than it is at $b = 4^{\\circ}$.  Some disk models such as model 2 of \\cite{2mass_bulge}, constructed to fit 2MASS star counts, contain an inner hole, which would lower potential disk contamination of bulge samples substantially.  Rather than speculate further on whether or not there is a disk within the central kpc and what its properties might be, we refer back to the large samples of ``bulge'' giants studied in the many references cited above. The innermost field included in most of these is Baade's Window, with $b$ ranging from $-3.9$ to $-4.1^{\\circ}$. The galactic latitudes for the three microlensed dwarfs are $-2.5$, $-3.6$ and $-4.8^{\\circ}$, so that two of these are slightly smaller than that of Baade's window.  However, recently \\cite{rich07} presented a detailed abundance analysis from near-IR spectra of M giants in a field at ($l,b$) = ($0^{\\circ},-1^{\\circ}$), significantly closer to the Galactic center than any of the three microlensed dwarfs.  Their sample of 17 M giants has the usual properties seen in the Baade's Window samples, a mean [Fe/H] of $-0.22$~dex, with the most metal rich at [Fe/H] $+0.02\\pm0.11$~dex.  These  giants show the usual $\\alpha$-enhancement.  Surely the microlensed dwarfs, each located more than twice as high above the Galactic plane than the giants in this inner bulge field, then cannot be from a different population than the \\cite{rich07} M giant sample;  we thus do not believe that the sample of microlensed bulge stars is contaminated by any disk stars that might be co-located near the Galactic center.    We suggest that sample bias is responsible for the difference between the median Fe-metallicity seen in the bulge giant samples and that of the microlensed bulge dwarfs. A possible mechanism is the very high mass loss rates predicted to occur  at the very high metallicities and moderately high luminosities being discussed here.  These are  high enough that in an old stellar population, stars are predicted to lose enough mass to peel off the red giant branch (RGB) before reaching the He-flash, which they never go through.  Red clump stars, which are burning He in their cores on the horizontal branch, are even more evolved than giants at the RGB tip, and so their evolutionary tracks, once mass loss is properly taken into account, may indicate that their expected numbers in an old very-metal rich population may be even more depleted  than are luminous first ascent RGB stars. This effect may have been detected through CMD studies of the extremely metal-rich open cluster NGC~6791 with [Fe/H] +0.45~dex \\citep{carretta07}. In this cluster, \\cite{kalirai07} found a strong relative absence of luminous RGB stars. We propose that this relative paucity of giants on the upper RGB for very metal-rich old populations produces a bias against the highest metallicity stars by preferentially eliminating them from the samples being studied by all previous investigations in the Galactic bulge, which samples consist of one or more of the luminous K giants, luminous M giants, and red clump stars found there.  If this is correct, then the mean bulge metallicity may be comparable to that expected from the radial gradients prevailing within the Galactic disk at a time $\\sim$5~Gyr ago extrapolated to the Galactic center.  It thus may be that at the present time there is a gradient within the solar circle of stellar metallicity with Galactic radius that is roughly comparable to that measured for the interstellar medium (ISM) at present from HII regions by \\cite{esteban05} and from planetary nebulae (PN) by \\cite{pn_grad}.  The latter present an estimate of the change of the radial abundance gradient in the Galaxy as a function of time from PN of varying ages.  They suggest that the radial abundance gradient in the ISM in the Galactic disk was twice as large $\\sim$5~Gyr ago when the Sun was formed than it is now.  There are important consequences for the chemical evolution of extragalactic objects as well if our conjecture regarding the metallicity distribution of the Galactic bulge is correct.  The interpretation of spectra and broad band photometry of the integrated light of simple old metal-rich stellar populations such as are believed to exist in luminous elliptical galaxies may be affected.  We expect an underestimate of the true Fe-metallicity of such systems to occur with commonly used tools such as Lick indices, calculated with stellar evolutionary tracks and isochrones that ignore mass loss, as is normally the case.  We expect an overestimate of the mass-to-light ratio to occur if the luminous RGB stars we expect to be present  near the RGB tip in a metal-rich population are not present in their expected numbers as they contribute only a small fraction of the mass, but a larger fraction of the total luminosity. Changes in  the strength of specific absorption features can be expected as well, particularly in the IR, where luminous red giants dominate the integrated light assuming they are present.  }   We have analyzed a  high dispersion spectrum of a microlensed dwarf, \\mystar, in the Galactic bulge.  The magnification of this event was very high, HIRES on the 10-m Keck~I Telescope was used, the weather was clear with good seeing, and the exposure time was long compared to any previous such data, so the resulting spectrum has a relatively  high signal-to-noise ratio.  We stress that in principle the abundance analysis of a upper main sequence dwarf is much easier and less prone to error for spectra of a fixed signal-to-noise ratio than that of a much cooler but much brighter bulge giant with a very complex spectrum full of blends and of strong molecular bands. The advantages of analyzing microlensed bulge stars, for which the required high signal-to-noise ratio can sometimes be achieved, to improving our understanding of the [Fe/H] distribution and chemical evolution of the Galactic bulge are large.  We have derived  for \\mystar, which we believe to be a dwarf below the main sequence turnoff with \\teff $\\sim 5400$~K, a very high Fe-metallicity, [Fe/H] = \\fehogle ~dex.  This is very peculiar given that many previous surveys of the metallicity distribution carried out with large samples of K or M giants in the Galactic bulge find both the mean and median [Fe/H] to be sub-solar.  The two other highly magnified Galactic bulge dwarfs studied in detail, \\johnstar\\ by \\cite{johnson07} and \\moa\\ by \\cite{johnson08}, also have very high  Fe-metallicities.    In order to produce consistency, we suggest that there is a sampling bias in the bulge giant samples such that very metal-rich giants are strongly depleted.  We suggest a physical mechanism for this, the very high mass loss rates expected for such metal-rich old giants can exhaust their envelopes prior to the normal He-flash.   We also find that \\mystar\\ does not show enhancements of the $\\alpha$ elements; neither does \\johnstar, analyzed by \\cite{johnson07}.  However, most bulge giants from samples with high dispersion spectroscopy, e.g. the work of \\cite{fulbright07} and particularly those from \\cite{lecureur07}, do show large (and varying from star to star) enhancements of [Na/Fe],  [Mg/Fe] and [Al/Fe] even at super-solar metallicities. We suggest that it is the abundances deduced for the microlensed dwarfs that best represent the initial chemical inventory of the interstellar medium at the time these stars formed, while those derived for the bulge giants may not.   We recognize that three stars is a very small sample, but the implications of our results and inferences for the chemical evolution of the Galactic bulge and for the interpretation of integrated light spectra and broad-band photometry of old simple stellar populations such as  luminous ellitical galaxies, are so important that we offer these hypotheses at this time.   The study of additional highly microlensed Galactic bulge dwarfs to increase the sample from just three such stars is clearly urgent. Now that ongoing microlensing surveys make such observations of Galactic bulge dwarfs feasible,   we expect substantial improvements in the sample size of high quality spectra for Galactic bulge dwarfs. Suitable high magnification events are  rare and lining up the necessary instruments/telescopes/clear weather at just the right time is difficult. Several years may be required to accumulate a suitable sample of spectra of highly microlensed dwarf stars in the Galactic bulge. We are  now appropriately positioned to carry out such a time critical program at the Keck Observatory over the next few years, and eagerly look forward to confirmation of our perhaps premature and provocative hypotheses in the not too distant future."
    },
    "0801/0801.0321_arXiv.txt": {
        "abstract": "We present detailed radiative transfer spectral synthesis models for the Iron Low Ionization Broad Absorption Line (FeLoBAL) active galactic nuclei (AGN) \\FQBS\\ and \\ISO.  Detailed NLTE spectral synthesis with a spherically symmetric outflow reproduces the observed spectra very well across a large wavelength range.  While exact spherical symmetry is probably not required, our model fits are of high quality and thus very large covering fractions are strongly implied by our results.  We constrain the kinetic energy and mass in the ejecta and discuss their implications on the accretion rate. Our results support the idea that FeLoBALs may be an evolutionary stage in the development of more ``ordinary'' QSOs. ",
        "introduction": "Spectroscopic observations of quasars show that about 10--20\\% have broad absorption troughs in their rest-frame UV spectra \\citep[see][for example]{trump06}. These absorption lines are almost exclusively blueshifted from the rest wavelength of the associated atomic transition, indicating the presence of an outflowing wind in our line of sight to the nucleus. The line-of-sight velocities range from zero to up to tens of thousands of kilometers per second \\citep[e.g.,][]{narayanan04}.  While understanding these outflows is of fundamental interest for understanding the quasar central engine, it is also potentially important for understanding the role of quasars in the Universe.  The observation that the black hole mass is correlated with the velocity dispersion of stars in the host galaxy bulge \\citep[e.g.,][]{magorrian98,fm00,gebhardt00} indicates a co-evolution of the galaxy and its central black hole.  The close co-evolution implies there must be feedback between the quasar and the host galaxy, even though the sphere of gravitational influence of the black hole is much smaller than the galaxy.  Energy arguments, however, show that is quite feasible that the black hole can influence the galaxy; as discussed by \\citet{begelman03},  the accretion energy of the black hole easily exceeds the binding energy of the host galaxy's bulge.  The nature of the feedback mechanism that carries the accretion energy to the galaxy is not known.  Since AGNs are observed to release matter and kinetic energy into their environment via outflows, it is plausible that these outflows contribute to the feedback in an important way. One of the difficulties in using quasar outflows in this context is that they are sufficiently poorly understood that there are significant uncertainties in such basic properties as the total mass outflow rate and the total kinetic energy.   What is the kinetic luminosity of the broad absorption line quasar winds?  That turns out to be very difficult to constrain.  While the presence of the blueshifted absorption lines unequivocally indicates the presence of high-velocity outflowing gas, the other fundamentally important properties of the gas, including the density, column density, and covering fraction are very difficult to constrain.  The density is difficult to constrain because the absorption lines are predominately resonance transitions, and their strengths are not very sensitive to density.  Without knowing the density, the distance of the gas from the central engine cannot be constrained; the same ionization state can be attained by dense gas close to the central engine, or rare gas far from the central engine.  Density estimates are possible when absorption lines are seen from non-resonance transitions, but even then, they can differ enormously.  For example, \\citet{dekool01} analyzed metastable \\ion{Fe}{2} absorption lines in FBQS~0840$+$3633, and inferred a electron density $<1000-3000\\rm\\, cm^{-3}$ and a distance from the central engine of several hundred pc. In contrast, \\citet{eracleous03} analyze the metastable \\ion{Fe}{2} absorption in Arp~102B with photoionization models and infer a density of at least $10^{11}\\rm \\, cm^{-3}$ and a distance of less than $7 \\times 10^{16}\\rm \\, cm$.   The global covering fraction is also difficult to constrain directly from the quasar spectrum; we know the gas, at least, partially covers our line of sight, but we have little information about other lines of sight.  Covering fraction constraints are generally made based on population statistics.  In a seminal paper, \\citet{weymann91} showed that for most BALQSOs, the emission line properties are remarkably similar to non-BALQSOs.  Thus, the fact that 10--20\\% of quasar spectra contain broad absorption lines is interpreted as evidence that there is a wind that covers 10--20\\% of sight lines to all similar quasars, and whether or not we see absorption lines depends on our orientation.   Alternatively, some BAL quasars have notably different line emission than the average quasar; examples are the low-ionization BALQSOs studied by e.g., \\citet{boroson92b}.  These objects may instead represent an evolutionary stage of quasars, as the quasar emerges from the cloud of gas and dust in which it formed \\citep{becker97}.  While it seems that the column density should be easy to constrain, more recent work has shown that it can be very difficult to measure. It was originally thought that non-black absorption troughs indicated a relatively low column density for the absorbing gas \\cite[equivalent hydrogen column densities of $10^{19-20}\\rm \\,cm^{-2}$, e.g.,][]{hamann98}.  But it has now been found that the non-black troughs indicate velocity-dependent partial covering, where the absorption covers part of the emission region, and the uncovered part fills in the trough partially \\citep[e.g.,][]{arav99}.  Thus, the column density appears to be high, but it is very difficult to constrain directly from the data except in a few very specialized cases \\citep[see for example][]{gabel06,arav05}.  How can we make progress on this problem?  It is becoming clear that because of the difficulties described above, the traditional techniques for analysis of troughs (e.g., curve of growth) and modeling (e.g., photoionization modeling to produce absorption line ratios and equivalent widths) are limited.  An approach that may be profitable is to construct a physical model for the outflow, and constrain the parameters of the model using the data.  Our first foray into constructing physical models for quasar winds was performed by \\citet{branch02}.  In that paper, the FeLoBAL\\footnote{FeLoBALs are distinguished by the presence of absorption in low-ionization lines such as \\ion{Al}{3} and \\ion{Mg}{2} as well as absorption by excited states of \\ion{Fe}{2} and \\ion{Fe}{3}.} FIRST~J121442$+$280329 was modeled using SYNOW, a parameterized, spherically-symmetric, resonant-scattering, synthetic spectrum code more typically used to model supernovae \\citep{fisher00}.  The difference between this treatment and a more typical one applied to the same data by \\citet{dekool02b} is that SYNOW assumes that emission and absorption are produced in the same outflowing gas.  In contrast, the approach taken by \\citet{dekool02b} assumes that absorption is imprinted upon a typical continuum+emission line quasar spectrum; that is, the absorbing gas is separated from the emission-line region. In fact, based on the analysis of the \\ion{Fe}{2} metastable absorption lines, they find that the  absorber is 1--30 parsecs from the central engine, much farther than the quasar broad emission-line region. Note that FIRST~J121442$+$280329 is not the only object that can be  modeled using SYNOW; \\citet{casebeer04} present a SYNOW model of another FeLoBAL, ISO~J005645.1$-$273816.  The SYNOW model is attractive because it is simple; only one component is needed to model both the emission and absorption lines.  However, this model is limited.  The primary purpose of the SYNOW program is to identify lines in complicated supernova spectra.  Thus, individual ions can be added to a SYNOW run at will in order to see if features from emission and absorption from those ions is present.  It does not solve the physics of the gas, so physical parameters beyond the existence of a particular species and its velocity extent cannot be extracted from the results.  We test the ideas of \\citet{branch02} and \\citet{casebeer04} by using the generalized stellar atmosphere code \\texttt{PHOENIX} to model the spectra of the two FeLoBALs that were successfully modeled using SYNOW, and including spectra that extend to rest-frame optical wavelengths for FIRST~J121442$+$280329.  {\\tt PHOENIX} is a much different code than \\texttt{SYNOW} in that it contains all the relevant physics to determine the spectrum of outflowing gas.  It solves the fully relativistic NLTE radiative transfer problem including the effects of both lines and continua in moving flows.  For a discussion of the use of both \\texttt{SYNOW} and \\texttt{PHOENIX} in the context of modeling supernovae spectra, see \\citet*{bbj03}.  We find that \\texttt{PHOENIX} is able to model the spectra from these objects surprisingly well, and we are able to derive several important physical parameters from the model.  In \\S  \\ref{Models} we describe the \\texttt{PHOENIX} model in detail. In \\S  \\ref{Modelfits} we describe our determination of the best-fitting model.  In \\S  \\ref{Results} we describe the results of our model fitting.  In \\S  \\ref{Discussion} we discuss the physical implications of the model, how it   relates to other BAL spectra and where it fits in the BAL picture.  An appendix includes a flowchart of a \\phx\\ calculation. ",
        "conclusions": "Using the spectral synthesis code  \\texttt{PHOENIX} we compute synthetic spectra that provide very good fits to the observed restframe UV and optical spectra of two FeLoBALs. While our models are limited to exact spherical symmetry, they provide excellent fits. In order to reconcile our results with the polarization data on these objects, we require some asymmetry, but still a high global covering fraction, which would only modestly affect the flux spectrum.  We are able to determine a luminosity distance estimate which is direct and is accurate to around 50\\%. The question arises: could these objects be used as distance indicators at high z, even if only as a sanity check on the really high-z Hubble diagram from GRBs?  Our results lend support to the inference that FeLoBALs are an evolutionary stage of the QSO as opposed to a pure orientation effect. Our model with a smaller covering factor may be able to explain other BAL QSO such as overlapping trough QSOs.  In addition our model column densities, which are Compton thick, match those that are  expected from X-ray non-detections of these objects.  For future work we plan to explore metallicity effects on the model spectrum. This is not trivial, as it is not just a matter of changing the metallicity and comparing the model. Completely different sets of dynamical and luminosity parameters may be required to  achieve the best fitting model spectrum. We plan to continue our analysis with the \\citet{hall07} spectrum, which shows $H\\beta$ absorption. We have compared our observed spectrum with the one from that paper and think that it is an excellent candidate for this type of modeling."
    },
    "0801/0801.2923_arXiv.txt": {
        "abstract": "The sources of nuclear uncertainties in nova nucleosynthesis have been identified using hydrodynamical nova models. Experimental efforts have followed and significantly reduced those uncertainties. This is important for the evaluation of nova contribution to galactic chemical evolution, gamma--ray astronomy and possibly presolar grain studies. In particular, estimations of expected gamma--ray fluxes are essential for the planning of observations with existing or future satellites. ",
        "introduction": "Novae are thermonuclear runaways occurring at the surface of a white dwarf accreting hydrogen rich matter from its companion in a close binary system\\cite{Sta72,Geh98,JH98,JH07}. Material from the white dwarf $^{12}$C and $^{16}$O (CO nova) or $^{16}$O, $^{20}$Ne plus some Na, Mg and Al isotopes (ONe nova) provide the seeds for the operation of the CNO cycle and further nucleosynthesis. Novae are supposed to be at the origin of galactic $^{15}$N and $^{17}$O and contribute to the galactic chemical evolution of $^7$Li and $^{13}$C. In addition they produce radioactive isotopes that could be detected by their gamma--ray emission: $^7$Be (478~keV), $^{18}$F ($\\le$511~keV), $^{22}$Na (1.275~MeV) and $^{26}$Al (1.809~MeV). The yields of these isotopes depend strongly on the hydrodynamics of the explosion but also on nuclear reaction rates involving stable and radioactive nuclei. Tests of sensitivity to the reaction rates uncertainties have been done using parametrized\\cite{Wor94}, semi--analytic\\cite{Coc95}, post--processed\\cite{Ili02} nova models but also with a 1-D hydrodynamical model. Indeed, in a series of papers the impact of nuclear uncertainties in the hot-pp chain\\cite{Be96}, the hot-CNO cycle\\cite{F00}, the Na--Mg--Al region\\cite{NaAl99} and Si--Ar region\\cite{SiAr01} have been investigated with the Barcelona (SHIVA) hydrocode. In this way, the temperature and density profiles, their time evolution, and the effect of convection time scale were fully taken into account. The nuclear reaction rates whose uncertainties affected most nova nucleosynthesis having been identified, many nuclear physics experiments were conducted to reduce these uncertainties. In this review, we will shortly summarize the experimental progress made in this domain. ",
        "conclusions": "Detailed calculations performed with the SHIVA hydrodinamical code have enable the identification of nuclear uncertainties affecting nova nucleosynthesis. We see that less than ten years after, great progress have been made thanks to experimental efforts, in particular for the $^{17}$O(p,$\\gamma)^{18}$F, $^{17}$O(p,$\\alpha)^{14}$N, $^{18}$F(p,$\\gamma)^{19}$Ne, $^{18}$F(p,$\\alpha)^{15}$O, $^{21}$Na(p,$\\gamma)^{22}$Mg, $^{22}$Na(p,$\\gamma)^{23}$Mg $^{25}$Al(p,$\\gamma)^{26}$Si, $^{26g.s.}$Al(p,$\\gamma)^{27}$Si and $^{30}$P(p,$\\gamma)^{31}$S reactions that were identified as the most influential. However, further efforts are required for the $^{22}$Na(p,$\\gamma)^{23}$Mg, $^{25}$Al(p,$\\gamma)^{26}$Si, $^{30}$P(p,$\\gamma)^{31}$S reactions and especially for the $^{18}$F(p,$\\alpha)^{15}$O reaction. For this last reaction where contribution of interfering broad resonance tails are essential, progress should come from direct measurements with intense $^{18}$F beam (TRIUMF) or from indirect (THM\\cite{THM}) measurement planned at the CRIB of the Center for Nuclear Studies (Wako).   \\begin{theacknowledgments} I am indebted to Margarita Hernanz and Jordi Jos\\'e for a now more than twelve years collaboration on nova nucleosynthesis and to Nicolas~de~S\\'er\\'eville for frequent discussions. Many thanks also to Carmen Angulo, Christian Iliadis, Fa\\\"{\\i}rouz Hammache, Fran\\c{c}ois de Oliveira Santos and Claudio Spitaleri for long time collaborations. \\end{theacknowledgments}"
    },
    "0801/0801.2937_arXiv.txt": {
        "abstract": "OH megamasers (OHMs) emit primarily in the main lines at 1667 and 1665 MHz, and differ from their Galactic counterparts due to their immense luminosities, large linewidths and 1667/1665 MHz flux ratios, which are always greater than one. We find that these maser properties result from strong 53 \\mic\\ radiative pumping combined with line overlap effects caused by turbulent linewidths $\\sim$ 20 \\kms; pumping calculations that do not include line overlap are unreliable. A minimum dust temperature of \\about\\ 45 K is needed for inversion, and maximum maser efficiency occurs for dust temperatures \\about\\ 80 -- 140 K. We find that warmer dust can support inversion at lower IR luminosities, in agreement with observations. Our results are in good agreement with a clumpy model of OHMs, with clouds sizes \\la\\ 1 pc and OH column densities \\about\\ 5\\x\\E{16}\\cmsq, that is able to explain both the diffuse and compact emission observed for OHMs. We suggest that {\\em all} OH main line masers may be pumped by far-IR radiation, with the major differences between OHMs and Galactic OH masers caused by differences in linewidth produced by line overlap. Small Galactic maser linewidths tend to produce stronger 1665 MHz emission. The large OHM linewidths lead to inverted ground state transitions having approximately the same excitation temperature, producing 1667/1665 MHz flux ratios greater than one and weak satellite line emission. Finally, the small observed ratio of pumping radiation to dense molecular gas, as traced by HCN and HCO$^+$, is a possible reason for the lack of OH megamaser emission in NGC 6240. ",
        "introduction": "OH megamasers are extremely luminous extragalactic OH maser sources that have isotropic luminosities a million or more times greater than their Galactic counterparts. OHMs are found only in the nuclear regions of luminous (LIRGs) and ultraluminous (ULIRGs) infrared galaxies, where intense star formation is occurring. These masers emit primarily in the main lines at 1667 and 1665 MHz, although weaker emission in the satellite lines at 1612 and 1720 MHz has also been detected \\citep{baan87}. In addition to their immense luminosities, OHMs differ from their Galactic counterparts in their very large linewidths and their 1667/1665 MHz flux ratios. Galactic OH mainline masers in star forming regions tend to have linewidths \\la\\ 1 \\kms\\ and 1667/1665 MHz flux ratios less than one. In contrast, OHMs have overall linewidths \\ga\\ 100 \\kms\\ and their 1667/1665 MHz flux ratios are always greater than one. Although this ratio varies, the typical ratio is \\about\\ 5 \\citep{lonsdale02}. \\citet{baan89} found a clear separation between OH emitters and absorbers based on the IR properties of the host galaxies. The OHMs have larger IR luminosity (\\LIR) and tend to be warmer. OHMs may occur for lower \\LIR\\ as long as the dust is warm enough. The Arecibo OHM survey conducted by \\citet{darling02} targeted LIRGs and increased the number of OHMs to \\about\\ 100. That survey substantiated the conclusions of \\citet{baan89} and showed that the fraction of OHMs in LIRGs is an increasing function of far-IR (FIR) luminosity and color, with most of the detections found in warm ULIRGs. These observations strongly suggest that the masers are pumped by the intense FIR radiation in the OH pump lines at 53 \\mic\\ and 35 \\mic. Support for FIR pumping is provided by observations using the Infrared Space Observatory (ISO). \\citet{skinner} observed the 35 \\mic\\ pumping transition in Arp 220 and concluded that absorption in this line alone could power its maser emission. Significantly, OH and water megamasers have not been observed in the same galaxy \\citep{lo05} even though both are often found in the same Galactic star forming regions.   The standard OHM model was introduced by \\citet{baan85}, who proposed that the maser emission is produced by low gain unsaturated amplification of background radio continuum. This proposal was confirmed by the comprehensive study of OHMs performed by \\citet{henkelwilson90}. They found good agreement with observations if the maser transitions have approximately equal excitation temperatures, giving a 1667/1665 MHz optical depth ratio of \\about \\ 1.8. However, subsequent VLBI observations revealed compact maser emission on parsec scales with amplification factors \\ga \\ 800 and very large 1667/1665 MHz line ratios \\citep{lonsdale98}. A surprise of these observations was that linewidths remained large (tens of \\kms) even on the smallest observed angular scales \\citep{lonsdale02}. The VLBI observations led to the suggestion that there are two different modes of maser operation --- low gain, unsaturated amplification responsible for the diffuse maser emission and high gain, saturated emission producing the compact sources. It was also suggested that the two classes may have different pumping mechanisms, with the diffuse emission pumped by IR radiation and the compact masers by a combination of collisions and radiation \\citep{lonsdale02}. However, an extension of the standard model is able to explain both the diffuse and compact emission from IIIZw35 by assuming a clumpy maser medium \\citep{parra05}. Each cloud generates low gain, unsaturated emission, and strong compact emission occurs when the line of sight intersects more than one cloud. A similar model has been used to explain the megamaser emission from IRAS 17208-0014 \\citep{momjian}. The clumpy maser model also explains the observation that compact masers are always found embedded in the diffuse emission and do not occur in isolation \\citep{lonsdale02}.  An observational difficulty is that the nearest OHMs are located at distances of \\about \\ 100 Mpc and thus cannot be probed nearly as accurately as Galactic masers. The study of OHMs thus relies heavily on analyzing their global properties. It is important to first explain the global properties of OHMs and then construct more comprehensive models for those sources which have been observed in more detail using VLBI techniques. Previous OHM models \\citep{henkelwilson90,randell} were performed before the discovery of the compact maser sources. Although the clumpy maser model seems to provide a phenomenological explanation to the compact and diffuse maser emissions, it is not yet backed by pumping calculations. The time is right to develop a pumping model that can explain both the compact and diffuse emission using physical conditions that are consistent with those expected in the maser regions. This is the aim of the present paper. In section 2 we summarize what is known concerning the physical conditions existing in the maser region. In section 3 we explain our model and present the results of our calculations. Section 4 contains a summary and discussion. ",
        "conclusions": "OHMs are characterized by their extreme luminosity, 1667/1665 MHz flux ratio greater than one, and large linewidths. Our detailed pumping analysis of OHMs shows that the above properties are a natural outcome of an intense FIR pump combined with line overlap effects due to large turbulent linewidths. The 53 \\mic\\ lines provide the primary pump, and a minimum dust temperature of \\about \\ 45 K is needed for main line maser production. In agreement with the observations of \\citet{baan89} and \\citet{darling02}, we find that a smaller \\LIR \\ can support inversion, as long as the temperature is warm enough. We find no conditions where collisions can produce the observed maser emission, and when combined with a strong FIR radiation field, collisions tend to weaken and then thermalize the masers. These results are consistent with the detection of the 183 GHz water maser combined with the lack of detection of the 22 GHz water maser in Arp 220 \\citep{cernicharo}. The 22 GHz maser is collisionally pumped and requires relatively high densities and temperatures for its production, while the 183 GHz maser is produced at much lower densities and temperature. The large OHM linewidths produce significant overlap among the the 53 \\mic\\ pump lines resulting in nearly equal negative excitation temperatures for all ground state transitions, in agreement with observations \\citep[figure \\ref{fig:color-color}; see also][]{henkelwilson90}. The 1667/1665 MHz flux ratios are thus greater than one and only weak satellite line emission is produced. Our pumping results are in agreement with those required by clumpy maser models that explain both the diffuse and compact maser emission from OHMs such as IIIZw35. Finally, the small observed ratios of \\LIR /\\LHCN \\ and \\LIR /\\LHCO+ \\ relative to those found in OHMs is a possible reason for the lack of OH megamaser emission in the ULIRG NGC6240. A minimum ratio of pumping radiation to dense molecular gas may be a prerequisite for becoming a OHM.  Excited state transitions at 5 cm and 6 cm have also been observed in OHMs. However, these lines have only been detected in absorption. \\citet{henkel87} observed absorption in the three 6 cm $^2\\Pi_{1/2}(J = 1/2)$ transitions in five OHMs, and absorption has been detected also in one 5 cm $^2\\Pi_{3/2}(J = 5/2)$ transition in Arp 220 \\citep{henkel86}. These observations are surprising since masing in the 5 cm and 6 cm lines has often been observed in conjunction with Galactic ground state OH maser sources \\citep{cesaroni}. Our pumping models of ground state maser emission in OHMs usually predict absorption in the 5 cm lines. However, the 6 cm lines are weakly inverted with opacities a hundred times smaller than the 1667 MHz line. Our calculations find that absorption in the 6 cm lines requires smaller dust temperatures than exist in the maser regions. We expect the absorption to occur outside the maser regions where the dust temperature is lower. Understanding the absorption at 5 cm and 6 cm is an important problem that we plan to treat in a future paper.  \\subsection{OH main-line maser pump}  Our results suggest that {\\em all} OH main-line masers could be pumped by the same mechanism: far-IR radiation. The major differences between Galactic star forming regions on one hand and evolved stars and OHMs on the other can be attributed to differences in line overlap effects. As shown in \\S\\ref{sec:overlap}, linewidth is the most important factor in determining the 1667/1665 MHz flux ratio and thus is the dominant factor causing the extreme differences between the flux ratios of the two classes. In Galactic star forming regions the observed linewidths are \\la \\ 1 \\kms\\ and the 1665 MHz maser is usually strongest, with the 1667 MHz line often comparable in strength to the satellite lines. In OHMs the linewidths are large (\\ga \\ 10 \\kms) and the 1667 MHz is always stronger than the 1665 MHz line, with the satellite lines much weaker. Our results show that such large linewidths produce approximately equal negative excitation temperatures for the four ground state transitions, resulting in the 1667 MHz line being stronger than the 1665 MHz line and much weaker satellite lines. Smaller linewidths tend to produce a stronger inversion in the 1665 MHz line. In evolved stars, the main-line maser velocities are larger than those found in Galactic star forming regions and they also tend to have 1667/1665 MHz flux ratios \\ga\\ 1. Thus the linewidth could be the main reason why OH main-line masers display differences in spite of sharing a common pump mechanism."
    },
    "0801/0801.3270_arXiv.txt": {
        "abstract": "We present an investigation into the potential effect of systematics inherent in multi-band wide field surveys on the dark energy equation of state determination for two 3D weak lensing methods. The weak lensing methods are a geometric shear-ratio method and 3D cosmic shear. The analysis here uses an extension of the Fisher matrix framework to jointly include photometric redshift systematics, shear distortion systematics and intrinsic alignments. Using analytic parameterisations of these three primary systematic effects allows an isolation of systematic parameters of particular importance.  We show that assuming systematic parameters are fixed, but possibly biased, results in potentially large biases in dark energy parameters. We quantify any potential bias by defining a Bias Figure of Merit. By marginalising over extra systematic parameters such biases are negated at the expense of an increase in the cosmological parameter errors. We show the effect on the dark energy Figure of Merit of marginalising over each systematic parameter individually. We also show the overall reduction in the Figure of Merit due to all three types of systematic effects.  Based on some assumption of the likely level of systematic errors, we find that the largest effect on the Figure of Merit comes from uncertainty in the photometric redshift systematic parameters. These can reduce the Figure of Merit by up to a factor of $2$ to $4$ in both 3D weak lensing methods, if no informative prior on the systematic parameters is applied. Shear distortion systematics have a smaller overall effect. Intrinsic alignment effects can reduce the Figure of Merit by up to a further factor of $2$. This, however, is a worst case scenario. By including prior information on systematic parameters the Figure of Merit can be recovered to a large extent, and combined constraints from 3D cosmic shear and shear ratio are robust to systematics. We conclude that, as a rule of thumb, given a realistic current understanding of intrinsic alignments and photometric redshifts, then including all three primary systematic effects reduces the Figure of Merit by \\emph{at most} a factor of $2$. ",
        "introduction": "\\label{Introduction} It has recently been shown that the equation of state of dark energy could be constrained to a high degree of accuracy using wide and deep imaging surveys (see Albrecht et al., 2006; Peacock et al., 2006; for recent and extensive reviews). 3D weak lensing has been shown to be a particularly powerful way to use the information from such surveys in the determination of dark energy parameters (see Munshi et al., 2006 for a recent review). 3D weak lensing, in which the shear and redshift information of every galaxy is used, has the potential to constrain the dark energy equation of state, $w(z)=\\rho_{\\rm de}(z)/p_{\\rm de}(z)$, to $\\Delta w(z)\\sim 0.01$ using surveys such as Pan-STARRS (Kaiser et al., 2002) and DUNE (Refregier et al., 2006). However, the predictions made thus far (Heavens et al., 2006; Taylor et al., 2007) have only included statistical errors and have not included systematic effects. Since the scientific goal of many future surveys is to constrain $\\Delta w(z)\\sim 0.01$ such systematic effects have the potential to render any cosmological constraints impotent.  In this paper we will address astrophysical, instrumental and theoretical systematic effects relevant to multi-band weak lensing surveys in an analytic way. We specifically study the systematic effects of photometric redshifts, intrinsic alignments and shear distortion. We consider these three primary systematics to have potentially the largest effect on the ability of weak lensing surveys to constrain cosmological parameters. Secondary effects, such as source clustering, non-Gaussian effects and theoretical approximations such as the Born approximation have been shown to have a smaller effect on shear statistics (Shapiro \\& Cooray, 2006; Semboloni et al., 2007; Schneider et al., 2002). Note that the effect of  non-Gaussianity (Semboloni et al., 2007) has only been studied via simulations, a full analytic investigation could reveal non-Gaussianity to be a larger systematic effect (Takada, private communication). We will examine the degradation that these primary effects may produce in the determination of the dark energy equation of state constraints for two 3D weak lensing methods: 3D cosmic shear (Heavens, 2003; Castro et al., 2005; Heavens et al., 2006) and the shear-ratio method (Jain \\& Taylor, 2003; Taylor et al., 2007).  The spirit of the approach taken here is to use simple, analytic, descriptions of systematic effects. By distilling complex effects into simple components the change in cosmological parameter determination due to any particular aspect of a systematic effect can be analysed independently. For example, is the bias or the fraction of outliers in photometric redshifts a more important factor? The obvious penalty in taking such an approach is that the analytic approximations made may not be fully representative of real systematic effects.  The two 3D weak lensing methods are introduced in Section \\ref{Methodology}, however we urge the reader to refer to Heavens et al. (2006) and Taylor et al. (2007) for a complete and in-depth introduction to the methods. In Section \\ref{Methodology} we also discuss dark energy parameter prediction. In Section \\ref{Primary Systematic Effects} we introduce the primary systematic effects considered, and the parameterisations used. The potential bias in dark energy parameters due to each systematic effect is presented in Section \\ref{Bias in Dark Energy Parameters}. A marginalisation over systematic parameters in presented in Section \\ref{Marginalising over systematic parameters}. We conclude and recap with a discussion in Section \\ref{Discussion}. For any technical details concerning the 3D weak lensing methods and systematic parameters see Appendix A and B. ",
        "conclusions": "\\label{Conclusion}  We have shown using simple analytic approximations that systematic effects can have a substantial impact on the ability of 3D weak lensing methods, the shear-ratio method and the 3D cosmic shear method, to constrain the dark energy equation of state. We used the Figure of Merit (FoM) to gauge the ability of a next generation experiment to constrain the dark energy equation of state. The systematic effects we considered are those associated with photometric redshifts, an overall distortion in the image plane and intrinsic alignments (both the GI and II terms).  The dark energy FoM can be degraded a factor of $2$ due to the photometric redshift systematics alone, for both 3D weak lensing methods. Shear distortion systematics have a small effect for both methods. Intrinsic alignment effects alone can degrade the FoM by a further factor of $2$ for the 3D cosmic shear method, but have a small effect for the shear-ratio method. This difference is due to the way in which the methods use shear information.   When an extra systematic parameter is encountered it can either be marginalised over using the available data (equivalent to self-calibration), thereby increasing the marginal error on any cosmological parameter of interest, or it can be assumed to be fixed. If a parameter is fixed then any deviation away from the assumed value will bias the most likely value of any measured cosmological parameter. This bias is a function of both the cosmological parameter error and the sensitivity of a method to any systematic parameter. From this analysis it has been shown that assuming some parameters to be fixed can lead to large biases in $w(z_p)$ and $w_a$, this is complimentary to the analysis done by Amara \\& Refregier (2007) who investiagted the bias in cosmological parameters due to shear measurement systematics using weak lensing tomography.  The methods are remarkably insensitive to many parameters, in particular most of the photometric redshift parameters, including the fraction of outliers, and an overall distortion of the shear field. By adopting a parameterisation of the photometric-spectroscopic plane we have shown that a bias in any photometric redshift redshift technique needs to be known to within $\\pm 10^{-3}$ for the dark energy FoM to remain unaffected. We have shown that to calibrate the photometric redshift parameters approximately $10^5$ spectroscopic redshifts are required for the shear-ratio method and approximately $10^2$ to $10^3$ for the 3D cosmic shear method. This difference can be attributed to the binning in redshift required by the shear-ratio method, which may also explain the agreement between the predicted spectroscopic requirements of the shear-ratio and shear tomography methods.  The intrinsic alignment terms were modelled using the Heymans et al. (2006) analytic approximations. The GI and II terms had a small effect on the FoM from the shear-ratio method, we found a drop in the FoM of approximately $10\\%$ when the extra covariances were included. For 3D cosmic shear the FoM is reduced by approximately $50\\%$, but there is a very small sensitivity to the extra intrinsic alignment systematic parameters. The 3D cosmic shear result is in agreement to what has been found using shear tomography in Bridle \\& King (2007). The caveat to these comparisons is that we use a fully 3D cosmic shear method, with no binning, and only investigate the Heymans et al. (2006) parameterisation; Bridle \\& King (2007) consider a variety of parameterisations and investigate a binning tomographic method.  We have shown that a good prior on systematic parameters can improve the FoM. The relative reduction in the FoM can be limited to $\\leq 30\\%$ if the prior on all systematic parameters has a Gaussian error of $\\sigma_P=0.001$. In particular we have shown that a prior on $z_{\\rm bias}$ can improve the FoM. Good priors on $z_{\\rm bias}$, $c_{\\rm calibration}$ would be particularly helpful in limiting the effect on the FoM.  By combining the 3D weak lensing methods the FoM can be increased by a up to a factor of $6$ relative to the methods individually. Furthermore the photometric redshift systematic parameter degeneracies are complementary leading to less systematic degradation in the combined constraints.    The bottom line is that the most important systematics to control are those concerning the photometric redshift distribution.  If these can be controlled to 1\\% accuracy, then the FoM for proposed future surveys such as DUNE and Pan-STARRS may be reduced by at most a factor of order two. In order to reduce these systematics to a negligible level needs $0.001$ accuracy in median redshifts, requiring ${\\mathcal O} (10^4)$ redshifts. We make a number of recommedations and observations with which to guide future systematic investigations \\begin{itemize} \\item Photometric redshift systematics play a dominant role in the systematics that affect 3D weak lensing. The individual systematic parameter which can have the largest effect on the FoM is the bias in photometric redshifts. However approximately ${\\mathcal O} (10^4)$ spectroscopic redshifts should be enough to calibrate photometric redshifts to the required accuracy. These would need to be representative of the photometric galaxies and complete. \\item If shear calibration bias is assumed to be fixed then an uncertainty in the bias of $\\sim 0.008$ can bias dark energy parameters by $>0.01$. However marginalising over shear bias has a smaller effect on the FoM. \\item Intrinsic alignments play a major role in 3D weak lensing systematic effects, and can reduce the maximum achievable FoM by up to $4$. The broad agreement between the parameterisations investigated here and in Bridle \\& King (2007) suggest that the general trends are robust. \\end{itemize}  Despite degrading systematic effects 3D weak lensing retains the potential to be the most powerful cosmological probe of dark energy. Weak lensing is entering a formative period in its development, given that the statistical ability of the method to constrain cosmology is accepted attention must now be focussed on understanding and reducing systematic effects."
    },
    "0801/0801.3931_arXiv.txt": {
        "abstract": "\\label{abst}  We study the dynamics of 2D and 3D barred galaxy analytical models, focusing on the distinction between regular and chaotic orbits with the help of the Smaller ALigment Index (SALI), a very powerful tool for this kind of problems. We present briefly the method and we calculate the fraction of chaotic and regular orbits in several cases. In the 2D model, taking initial conditions on a Poincar\\'{e} $(y,p_y)$ surface of section, we determine the fraction of regular and chaotic orbits. In the 3D model, choosing initial conditions on a cartesian grid in a region of the $(x, z, p_y)$ space, which in coordinate space covers the inner disc, we find how the fraction of regular orbits changes as a function of the Jacobi constant. Finally, we outline that regions near the $(x,y)$ plane are populated mainly by regular orbits. The same is true for regions that lie either near to the galactic center, or at larger relatively distances from it. ",
        "introduction": "\\label{Intro}  The dynamical evolution of galactic systems depends crucially on their orbital structure and in particular on what fraction of their orbits is regular, or chaotic. Thus to permit further studies, it is essential to be able to distinguish between these two types of orbits in a manner that is both safe and efficient. This is not trivial and becomes yet more complicated in systems of many degrees of freedom. A summary of the methods that have been developed over the years can be found in \\cite{Con_spr}.  In the present paper we use a method based on the properties of two deviation vectors of an orbit, the ``Smaller ALingment Index\" (SALI) \\cite{sk:1}. It has been applied successfully in different dynamical systems \\cite{sk:1,sk:3,sk:5,Pan:1,Ant:2,Ant:3,Bou:1,Man:1,Man:2,Man:3,Man:4}, frequently also under the name Alignment Index (AI) \\cite{VKS1,VKS2,VHC,KVC,KEV} and has been shown to be a fast and easy to compute indicator of the chaotic or ordered nature of orbits. We first recall its definition and we then show its effectiveness in distinguishing between ordered and chaotic motion by applying it to a barred potential of 2 and 3 degrees of freedom. Recently, a generalization of the SALI, the ``Generalized ALignment Index\" (GALI) was introduced by Skokos et al. (2007) \\cite{sk:6}, which includes the full set of the  $k$ initially linearly independent deviation vectors of the system to determine if an orbit is chaotic or not.  \\vspace{-0.5 cm} ",
        "conclusions": "\\label{Conclusions}  We used SALI to study the dynamical behavior of Hamiltonian models of 2D and 3D barred galaxies. We found that in both cases there is a significant amount of chaotic orbits. In the 2D model, we were able to chart a subspace of the phase space, to identify rapidly even tiny regions of regular motion and measure their percentages. In the 3D model, apart from computing the global percentages of regular and chaotic orbits, we calculated these percentages as a function of the energy and found that low values of the energy are mainly dominated by `regular' orbital motion. We also followed the distribution of the chaotic and regular orbits in the configuration space and we found that orbits which lie near the $(x,y)$--plane with relatively small mean deviations in $z$--direction are generally regular. Finally, we monitored the variation of their percentages as a function of their initial spherical radius and their mean spherical radius. We find that the fraction of regular orbits is dominant in regions near the center, as well as at relatively larger distances from it.  \\vspace{-0.5 cm}"
    },
    "0801/0801.1105_arXiv.txt": {
        "abstract": "We investigate the properties of one--dimensional flux ``voids'' (connected regions in the flux distribution above the mean flux level) by comparing hydrodynamical simulations of large cosmological volumes with a set of observed high--resolution spectra at $z\\sim 2$. After addressing the effects of box size and resolution, we study how the void distribution changes when the most significant cosmological and astrophysical parameters are varied. We find that the void distribution in the flux is in excellent agreement with predictions of the standard $\\Lambda$CDM cosmology, which also fits other flux statistics remarkably well. We then model the relation between flux voids and the corresponding one--dimensional gas density field along the line--of--sight and make a preliminary attempt to connect the one--dimensional properties of the gas density field to the three--dimensional dark matter distribution at the same redshift. This provides a framework that allows statistical interpretations of the void population at high redshift using observed quasar spectra, and eventually it will enable linking the void properties of the high--redshift universe with those at lower redshifts, which are better known. ",
        "introduction": "Over the past decade, properties of voids have been more widely investigated, using different observational probes and tracers (mainly in the local universe, see for example \\cite{hoyle}, \\cite{rojas}, \\cite{goldberg}, \\cite{rojas05}, \\cite{hoyle05}, \\cite{ceccarelli}, \\cite{patiri06a}, \\cite{patiri06b}, \\cite{tikhonov}), and theoretical -- analytical or numerical -- models in the framework of the $\\Lambda$ Cold Dark Matter ($\\Lambda$CDM) concordance cosmology (e.g. \\cite{peebles}, \\cite{arbabi}, \\cite{mathis}, \\cite{benson}, \\cite{gottlober}, \\cite{sheth}, \\cite{goldberg04}, \\cite{bolejko}, \\cite{colberg}, \\cite{padilla}, \\cite{furlanettopiran}, \\cite{hoeft}, \\cite{leepark}, \\cite{patiri06c}, \\cite{shandarin}, \\cite{aloisio}, \\cite{tully}, \\cite{parklee}, \\cite{brunino}, \\cite{vdw07}, \\cite{neyrinck}, \\cite{peebles07}).   The emerging picture is encouraging. On the one hand, some results appear to be somewhat hard to understand, such as the fact that galaxies observed at the edges of and in voids appear to be a fair sample of the whole galaxy population (the so-called 'void phenomenon', \\cite{peebles}), the fact that dwarf galaxies are not found in voids contrary to expectations from numerical simulations (e.g. \\cite{peebles07}, or the cold spot observed in the Cosmic Microwave Background (CMB) data (\\cite{rudnick}, \\cite{naselsky}). On the other hand, other observational results, such as those from the 2dF or SDSS galaxy redshift surveys, from studies of voids in quasar spectra, or the CMB spectrum (\\cite{caldwell}), are in reasonably good agreement with theoretical predictions from numerical simulations of the standard $\\Lambda$CDM cosmology. Linking the low--redshift properties of voids to those at higher redshifts offers the opportunity to constrain the void population over a significant fraction of cosmic time and to possibly find out more about voids, to ultimately understand their role in the context of cosmic structure formation.  Unfortunately, there are very few observables that constrain the population of voids at high redshift (see, however, \\cite{aloisio}). Here, we will use the \\lya forest as a tracer of voids in the high--redshift universe. The \\lya forest (e.g. \\cite{bi}) has been shown to arise from the neutral hydrogen embedded in the mildly non--linear density fluctuations around mean density in the Intergalactic Medium (IGM), which faithfully traces the underlying dark--matter distribution at scales above the Jeans length. It is a powerful cosmological tool in the sense that it probes the dark matter power spectrum in a range of scales ($1$ to $80\\,h^{-1}$\\,Mpc comoving) and redshifts ($z = $2--5.5) not probed by other observables. However, the observable is not the matter distribution itself but the flux: a one--dimensional quantity sensitive not only to cosmological but also to the astrophysical parameters that describe the IGM.  The so--called fluctuating Gunn--Peterson approximation \\citep{gunnpeterson} relates the observed flux to the IGM overdensity in a simple manner once the effects of peculiar velocities, thermal broadening, and noise properties are neglected: \\begin{equation} F({\\bf{x}},z)=\\exp[-A(z)(1+\\delta_{IGM}({\\bf{x}},z))^{\\beta}]\\, , \\end{equation} with $\\beta(z) \\sim 2-0.7(\\gamma(z)-1)$, $\\gamma$ the parameter describing the power--law ('equation of state') relation of the IGM, $T=T_0(z)((1+\\delta_{IGM})^{\\gamma-1}$, and $A(z)$ a constant of order unity which depends on the mean flux level, on the Ultra Violet Background (UVB) and on the IGM temperature at mean density. In practice, linking the flux to the density of the dark (or gaseous) matter requires the use of accurate hydrodynamical simulations that incorporate the relevant physical and dynamical processes. However, for approximate calculations equation\\,(1), in conjunction with simple physical prescriptions of \\lya clouds (e.g. \\cite{schaye}), can offer precious insights.  Provided the flux--density relation is accurately modelled, using the \\lya forest to check the properties of the population of voids in the matter distribution would avoid having to know the bias between galaxies and dark matter and would allow to sample over larger volumes the values of cosmological parameters that affect the void population. On the other hand, the fact that the flux information is one--dimensional requires non--trivial efforts to relate it to the three--dimensional (dark) matter density field.  Motivated by huge amounts of new data of exquisite quality, we here present a first attempt in this direction. Note that in the early 1990s, before the new paradigm of the \\lya forest absorption was proposed, voids from the transmitted \\lya flux of QSO spectra were analysed, using discrete statistics on the lines, by several groups (e.g. \\cite{crotts,duncan,ostriker,dobrzycki,rauch92}). These attempts, however, did not reach conclusive results due to small--number statistics. We defer to a future paper for a link between the observational properties outlined here and those at low redshift (e.g. \\cite{mathis,hoyle}) and for a careful investigation of the absorber--galaxy relation (e.g. \\cite{stocke,mclin}).  This paper is organized as follows. In Section \\ref{data}, we briefly describe the data set used. In Section \\ref{simulations} we present the hydrodynamical simulations. Section \\ref{results} contains the bulk of our analysis and the results, which are summarized in the conclusions in Section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions} We used \\lya forest QSO spectra to constrain the void population at $z\\sim 2$ at a range of scales and redshifts, which cannot be probed by other observables.  The main conclusions can be summarized as follows: \\begin{itemize} \\item the properties (sizes) of voids identified in simulations from  the flux distribution as connected regions above the mean flux level are in good agreement with observed ones; \\item using a set of hydrodynamical simulations and varying all the astrophysical and cosmological parameters, along with checking box size and resolution effects, we find that the flux void size distribution is a robust statistics that depends primarily on the flux threshold chosen to define the voids; \\item the flux seems to be a reliable tracer of the underlying dark matter distribution, altering the statistical properties of the dark matter density field (changing amplitude, shape, adding or suppressing power in a wavenumber dependent manner) changes the properties of the void population although at present is difficult to give constraints on cosmological parameters using these statistics; \\item flux voids, i.e. connected regions above the mean flux levels, correspond to gas densities around the mean density, regardless of the role of peculiar velocities, which contribute to a scatter in this relation, but do not alter significantly the mean values; \\item linking 1D voids to the corresponding 3D dark matter voids is more difficult, and we used a void--finding algorithm for that: we find that 3D DM voids with a mean overdensity of $\\delta_{t}=-0.5$ correspond to the flux voids that we have defined from the QSO spectra. However, the correspondence is good only for voids with sizes larger than about $7-10$\\,Mpc$/h$. \\end{itemize}   In this paper, we made a preliminary attempt to link the population of voids in the transmitted \\lya flux to the underlying 1D gas density and temperature and 3D dark matter density. The use of \\lya high--resolution spectra is important in the sense that it explores a new regime in scales, redshifts and densities which is currently not probed by other observables: the scales are of order few to tens of Mpc, the redshift range is between $z=2-4$, while the densities are around the mean density.  Further studies, that could possibly rely on wider data sets such as the low--resolution SDSS data (\\cite{shang}), on other future observables at higher redshifts (\\cite{aloisio}) or on tomographic studies involving QSO pairs or multiple line--of--sights (\\cite{dodorico}, \\cite{saitta07}) would be important in understanding the dynamical, thermal and chemical evolution of the void population and the interplay between galaxy and the IGM over a large fraction of the Hubble time."
    },
    "0801/0801.4516_arXiv.txt": {
        "abstract": "We study the anisotropy signature which is expected if the sources of ultra high energy, $>10^{19}$~eV, cosmic-rays (UHECRs) are extragalactic and trace the large scale distribution of luminous matter. Using the PSCz galaxy catalog as a tracer of the large scale structure (LSS) we derive the expected all sky angular distribution of the UHECR intensity. We define a statistic, that measures the correlation between the predicted and observed UHECR arrival direction distributions, and show that it is more sensitive to the expected anisotropy signature than the power spectrum and the two point correlation function. The distribution of the correlation statistic is not sensitive to the unknown redshift evolution of UHECR source density and to the unknown strength and structure of inter-galactic magnetic fields. We show, using this statistic, that recently published $>5.7\\times10^{19}$~eV Auger data are inconsistent with isotropy at $\\simeq98\\%$ CL, and consistent with a source distribution that traces LSS, with some preference to a source distribution that is biased with respect to the galaxy distribution. The anisotropy signature should be detectable also at lower energy, $>4\\times10^{19}$~eV. A few fold increase of the Auger exposure is likely to increase the significance to $>99\\%$ CL, but not to $>99.9\\%$ CL (unless the UHECR source density is comparable or larger than that of galaxies). In order to distinguish between different bias models, the systematic uncertainty in the absolute energy calibration of the experiments should be reduced to well below the current $\\simeq25\\%$. ",
        "introduction": "The origin of UHECRs is still  unknown: Sources have not been identified by cosmic-ray observations and all models of particle acceleration in known astrophysical objects are challenged by the extension of the spectrum to energies exceeding $10^{20}$~eV \\cite{Bhattacharjee:1998qc,Waxman_CR_rev}. The spectrum flattens (becomes harder) \\cite{Nagano:2000ve} and there is evidence for a composition change from heavy to light nuclei \\cite{Bird93,Abbasi05} at $\\sim10^{19}$~eV. This suggests that while the cosmic-ray flux below $\\sim 10^{19}$~eV is dominated by Galactic sources of heavy nuclei, it is dominated at higher energy by different sources of lighter nuclei. The isotropy of the UHECR arrival direction distribution suggests that these sources are extra-Galactic.  If UHECRs are produced by a population of extra-Galactic astrophysical objects, that traces the distribution of luminous matter, the inhomogeneous distribution of luminous matter is expected to imprint a characteristic anisotropy signature on the UHECR arrival direction distribution. If UHECRs are indeed light nuclei, this anisotropy is expected to be significant above $\\sim5\\times10^{19}$~eV,  where the propagation distance of light cosmic-ray nuclei is limited by interaction with the background radiation field to a distance of the order of 10's of Mpc (\\cite{Greisen:1966jv,Zatsepin:1966jv,Stecker:1999}, see \\cite{Sigl01} for a detailed review). This distance is comparable to that over which order unity variations in the density distribution of luminous matter are observed. Identification of the expected anisotropy signature would provide strong support to models where UHECRs are accelerated in known astrophysical objects, and would be inconsistent with most new physics, \"top-down,\" models, in which UHECRs are produced by the decay of heavy relic particles or topological defects \\cite{Bhattacharjee:1998qc}.  The expected anisotropy signal was analyzed in \\cite{Waxman:1996hp}. It was shown there that the signal should be detectable with high statistical significance once the number of detected UHECR events is increased beyond that available at that time, which was $\\approx20$ events above $4\\times10^{19}$~eV, by a factor $\\gtrsim10$. In this paper we revisit this subject, as the commissioning of the Auger detector may provide us within a few years with the required exposure \\cite{Yamamoto:2007xj}. While the analysis of the current paper largely follows that of \\cite{Waxman:1996hp}, it is improved in several respects. First, for the derivation of the large-scale structure of luminous matter we use a larger, as well as more complete and uniform, galaxy redshift survey, the PSCz catalog \\cite{Saunders:2000af} instead of the IRAS 1.2~Jy catalog \\cite{Fisher:1995kd}. The PSCz catalog contains 3 times more galaxies, including many more galaxies at distances of $150-300$ Mpc. Second, we use an updated cosmological model, $\\{\\Omega_m=0.27,\\ \\Omega_\\Lambda=0.73,\\ H_0=75\\ {\\rm km/s/Mpc}\\}$ instead of $\\{\\Omega_m=1,\\ \\Omega_\\Lambda=0,\\ H_0=100\\ {\\rm km/s/Mpc}\\}$, and analyze the effects of a possible redshift evolution of the UHECR source density. Third, we study the sensitivity of the correlation statistic to systematic errors in the determination of UHECR energies. Finally, we compare our proposed method for identifying the expected anisotropy signature to the more commonly used methods based on the power spectrum (e.g.~\\cite{Sommers:2000us,Deligny:2004dj,Anchordoqui:2003bx,Sigl}) and on the two point correlation function (e.g.~\\cite{Takeda:1999sg,DeMarco:2006a,Cuoco:2007id,Takami:2007vv} and references therein) of the angular distribution of UHECR arrival directions. In order to facilitate the use of the proposed statistic for the detection of anisotropy in UHECR data, we have made available at \\verb\"http://www.weizmann.ac.il/\"$\\sim$\\verb\"waxman/criso\" a numerical description of the intensity maps (see \\S~\\ref{sec:summary}).  An analysis of the expected anisotropy signal based on the PSCz catalog has recently been carried out in \\cite{Cuoco:2006}. The most important effects taken into account in the present analysis, and neglected in ref.~\\cite{Cuoco:2006}, are the dependence of the anisotropy signal on the finite number density of sources and on the systematic errors in the determination of UHECR energies. As shown in \\S~\\ref{sec:results} and discussed in \\S~\\ref{sec:summary}, both the finite number density of sources and the systematic uncertainties in energy determination have important implications to the predicted signal and its interpretation.  We assume in our analysis that the  UHERCs are protons. This assumption is motivated by two arguments. First, the observed spectrum of $>10^{19}$~eV cosmic-rays is consistent with a cosmological distribution of proton accelerators producing (intrinsically) a power-law spectrum of high energy protons, $d\\log n/d\\log E\\approx -2$, \\cite{Waxman:1995dg,Bahcall:2002wi} (see also fig.~\\ref{fig:spectrum}). This intrinsic power-law spectrum is consistent with that expected in most models of particle acceleration (\\cite{Waxman_CR_rev} and references therein). Second, the leading candidate extra-Galactic sources (GRBs and AGN, \\cite{Waxman_CR_rev} and references therein) are expected to accelerate primarily protons.  The paper is organized as follows. In sec.~\\ref{sec:Method} we describe the formalism used to analyze the anisotropy signal and present maps of the predicted angular distribution of UHECR intensity. In sec.~\\ref{sec:results} we study the sensitivity of several statistics to the predicted anisotropy signature. In sec.~\\ref{sec:Auger} we analyze Auger data, which was published during the preparation of this manuscript \\cite{AugerSc:2007}, using our correlation statistic. Sec.~\\ref{sec:summary} contains a brief summary of the results and a discussion of their implications.  The following point should be made here regarding the isotropy analysis of ref.~\\cite{AugerSc:2007}. In ref.~\\cite{AugerSc:2007} a correlation is found between the arrival direction distribution of $>5.7\\times10^{19}$~eV cosmic-rays (detected by the Auger experiment) and between the angular distribution of low-luminosity AGNs included in the V-C AGN catalog \\cite{V-C}. According to the analysis of ref.~\\cite{AugerSc:2007}, the probability that the detected correlation would arise by chance for an isotropic UHECR arrival direction distribution is $\\simeq1$\\%. Since low luminosity AGNs trace the distribution of luminous matter, one may argue that the detected correlation provides the first evidence for a correlation of UHECR sources and LSS. Unfortunately, the V-C catalog is merely a compilation of AGN data available in the literature, and is therefore incomplete both in its sky coverage and in its luminosity coverage. It does not, therefore, provide a correct description of the local LSS, as clearly pointed out in the introduction of ref.~\\cite{V-C}: \"The V-C catalog should not be used for any statistical analysis as it is not complete in any sense, except that it is, we hope, a complete survey of the literature\". The interpretation of the correlation reported in ref.~\\cite{AugerSc:2007} is thus unclear. ",
        "conclusions": "\\label{sec:summary}  We have derived the expected angular dependence of the UHECR intensity, under the assumption that the UHECRs are protons produced by extra-Galactic sources that trace the large scale distribution of luminous matter. All sky maps of the UHECR intensity, $I(E,\\hat\\Omega)$ with several energy thresholds $E$, are presented in figure~\\ref{fig:CRmap1} for an unbiased UHECR source distribution, where the UHECR source density is proportional to the LSS density, and in fig.~\\ref{fig:CRmap2} for a biased model, where the UHECR source density is proportional to the LSS density in over-dense region and vanishes elsewhere. A numerical representation of the intensity maps may be downloaded from \\verb\"http://www.weizmann.ac.il/\"$\\sim$\\verb\"waxman/criso\". The angular structure of the UHECR intensity map reflects the local large scale structure of the galaxy distribution, as can be seen by comparing figures~\\ref{fig:CRmap1} and~\\ref{fig:CRmap2} to figure~\\ref{fig:densityMap75}, which presents the integrated galaxy density out to a distance of 75~Mpc. The anisotropy is larger for higher energy threshold, since the propagation distance increases at lower energy (see fig.~\\ref{fig:DiffFlux}) and the contribution to the UHECR flux of sources beyond $\\sim100$~Mpc constitutes a roughly isotropic \"background\". We have used simple analytic arguments (\\S~\\ref{sec:B}) to show that inter-galactic magnetic fields may modify the intensity map on scales of a few degrees, but not on larger scales (a result which is consistent with detailed semi-analytic and numeric analyses \\cite{Dolag:2004kp,Lemoine08}).  We have defined a statistic, $X_C$ (eq.~\\ref{eq:X_C}), that measures the correlation between the predicted and observed UHECR arrival direction distributions, and showed that it is more sensitive to the expected anisotropy signature than the power spectrum and the two point correlation function (see table~\\ref{table:comp}). The value of $X_C$ for a given data set of UHECR arrival directions, the average  value of $X_C$ over realizations of the various models (isotropic, I, un-biased, UB, biased, B), and the distributions of $X_C$ values expected in realizations of the various models in the limit of infinite UHECR source density, can all be straightforwardly calculated using the numerical representations of the UHECR maps at \\verb\"http://www.weizmann.ac.il/\"$\\sim$\\verb\"waxman/criso\". In order to avoid sensitivity to possible distortions of the UHECR intensity map by magnetic fields, we have used $6^\\circ\\times6^\\circ$ bins in calculating $X_C$ and excluded the Galactic plane region, $|b|<12^\\circ$. As can be seen by comparing results in tables~\\ref{table:comp} and~\\ref{table:opt}, the exclusion of the Galactic plane region does not affect the results significantly. Table~\\ref{table:opt} demonstrates that the anisotropy signal is stronger at lower energy: Although the contrast of the fluctuations in the UHECR intensity is higher at high energy (see fig.~\\ref{fig:CRmap1}), the signal becomes weaker at higher energies since the number of observed UHECR drops rapidly with energy. At $E\\sim20$~EeV a significant contamination from Galactic sources may be present and the deflection of cosmic-ray particles may become significant. In order to avoid distortions of the UHECR intensity map by these effects, it is advisable to choose an energy threshold of $E>40$~EeV.  We have shown that the distribution of $X_C$ is not sensitive to the assumed redshift evolution of the source density distribution, to variations in the intrinsic spectrum of protons produced by the sources, and to statistical errors in the experimental determination of event energies (see fig.~\\ref{fig:sens} and table~\\ref{table:opt}). The distribution of $X_C$ does depend on the assumed local ($z=0$) source number density, $\\bar{s}_0$: The average value of $X_C$ obtained in realizations of the various models (I, UB, B) is independent of $\\bar{s}_0$, but the width of the distribution is wider, and hence discrimination between models is more diffcult, for lower $\\bar{s}_0$ (see fig.~\\ref{fig:sens}c and table~\\ref{table:opt}). Thus, in order to determine the probability of obtaining a certain value of $X_C$ in a given model (I, UB, B), one must make an assumption regarding the value of $\\bar{s}_0$. As mentioned in \\S~\\ref{sec:Method}, and discussed in detail in \\cite{Waxman:1996hp}, $\\bar{s}(z=0)$ may be constrained by the number of \"repeaters\", the number of sources producing multiple events. The last column of table~\\ref{table:opt} demonstrates that the presence or absence of repeaters may provide additional constraints on the source density (beyond $\\bar{s}_0\\gtrsim10^{-5}{\\rm Mpc}^{-3}$) once the number of events observed above 40~EeV exceeds several hundreds.  Figure~\\ref{fig:sens}d demonstrates that the distribution of $X_{C}$ values, which is expected to be measured from the distribution of UHECR events with energies inferred {\\it by the experiment} to exceed a certain threshold, is sensitive to uncertainties in the experimental calibration of the absolute energy scale: Since the anisotropy is stronger at higher energy, a systematic underestimate (overestimate) of event energies leads to a shift of the $X_C$ distribution towards larger (smaller) values. Since the anisotropy is also stronger (at a given energy) for stronger bias, this implies that an underestimate (overestimate) of the absolute energy scale will lead to an overestimate (underestimate) of the bias of the UHECR source distribution. In order to distinguish between different bias models, the systematic uncertainty in the absolute energy calibration of the experiments should be reduced to well below the current $\\simeq25\\%$.  We have shown, using the $X_C$ statistic, that the recently published $>5.7\\times10^{19}$~eV Auger data are consistent with a source distribution that traces LSS, with some preference to an UHECR source distribution that is biased with respect to the galaxy distribution, and inconsistent with isotropy at $\\simeq98\\%$ CL (see table~\\ref{table:Auger}). We have noted, however, that the optimization of the energy threshold made by the Auger collaboration analysis raises the concern that the significance with which isotropy is ruled out may be overestimated (see the end of \\S~\\ref{sec:Auger}). In order to confirm our detection of a correlation with LSS, one may search for the anisotropy signature in the Auger data using a lower energy threshold, $4\\times10^{19}$~eV, or repeat the analysis with an energy threshold of $5.7\\times10^{19}$~eV using a larger data set (which will be accumulated on a 1~yr time scale).  According to table~\\ref{table:opt}, the detection of $\\approx300$ events above $4\\times10^{19}$~eV is required in order to enable one to identify the predicted anisotropy signature, and to discriminate between different bias models, with statistical significance exceeding 99\\% CL. The experimental exposure required to accumulate this number of events is uncertain, due to the uncertainty in the absolute energy calibration of the UHECR experiments, which implies an uncertainty in the absolute normalization of the UHECR flux (\\cite{Bahcall:2002wi}, see also fig.~\\ref{fig:spectrum}). Adopting, for example, the absolute energy scale of the HiRes experiment, the average number of events detected above 40~EeV is 300 per year for the Auger exposure of $7000$~km$^2$~sr, and about 50 events per year for the planned Telescope Array exposure of $1200$~km$^2$~sr \\cite{TA:2009}. These numbers are reduced by a factor of approximately 2 if the preliminary absolute energy scale of the Auger experiment is chosen (see fig.~\\ref{fig:spectrum}). The numbers given in table~\\ref{table:opt} imply also that the exposure accumulated within a few years of observation is unlikely to increase the significance of the detection to $>99.9\\%$ CL, unless the UHECR source density is comparable (or larger) than that of galaxies."
    },
    "0801/0801.2383_arXiv.txt": {
        "abstract": "In this work we use a sample of 318 radio-quiet quasars (RQQ) to investigate the dependence of the ratio of optical/UV flux to X-ray flux, $\\alpha_{\\rm ox}$, and the X-ray photon index, $\\Gamma_X$, on black hole mass, UV luminosity relative to Eddington, and X-ray luminosity relative to Eddington. Our sample is drawn from the literature, with X-ray data from \\emph{ROSAT} and \\emph{Chandra}, and optical data mostly from the SDSS; 153 of these sources have estimates of $\\Gamma_X$ from \\emph{Chandra}. We estimate $M_{BH}$ using standard estimates derived from the H$\\beta$, Mg II, and C IV broad emission lines. Our sample spans a broad range in black hole mass ($10^6 \\lesssim M_{BH} / M_{\\odot} \\lesssim 10^{10}$), redshift ($0 < z < 4.8$), and luminosity ($10^{43} \\lesssim \\lambda L_{\\lambda} (2500$\\AA$) [{\\rm erg\\ s^{-1}}] \\lesssim 10^{48}$). We find that $\\alpha_{\\rm ox}$ increases with increasing $M_{BH}$ and $L_{UV} / L_{Edd}$, and decreases with increasing $L_X / L_{Edd}$. In addition, we confirm the correlation seen in previous studies between $\\Gamma_X$ and $M_{BH}$ and both $L_{UV} / L_{Edd}$ and $L_X / L_{Edd}$; however, we also find evidence that the dependence of $\\Gamma_X$ of these quantities is not monotonic, changing sign at $M_{BH} \\sim 3 \\times 10^8 M_{\\odot}$. We argue that the $\\alpha_{\\rm ox}$ correlations imply that the fraction of bolometric luminosity emitted by the accretion disk, as compared to the corona, increases with increasing accretion rate relative to Eddington, $\\dot{m}$. In addition, we argue that the $\\Gamma_X$ trends are caused by a dependence of X-ray spectral index on $\\dot{m}$. We discuss our results within the context of accretion models with comptonizing corona, and discuss the implications of the $\\alpha_{\\rm ox}$ correlations for quasar feedback. To date, this is the largest study of the dependence of RQQ X-ray parameters on black hole mass and related quantities, and the first to attempt to correct for the large statistical uncertainty in the broad line mass estimates. ",
        "introduction": "\\label{s-intro}  The extraordinary activity associated with quasars involves accretion onto a supermassive black hole (SMBH), with the UV/optical emission arising from a geometrically thin, optically thick cold accretion disk \\citep{shakura73}, and the X-ray continuum arising from a hot, optically thin corona that Compton upscatters the disk UV photons \\citep[e.g.,][]{haardt91}. In highly accreting objects, like quasars \\citep[$0.01 \\lesssim L_{bol} / L_{Edd} \\lesssim 1$, e.g.,][]{woo02,vest04,mclure04,koll06}, the X-ray plasma geometry is expected to be that of a hot, possibly patchy, ionized `skin' that sandwiches the cold disk \\citep[e.g.,][]{bis77,liang77,nayak00}. However, the evidence for this is not conclusive, and relies on data from X-ray binaries and low-$z$ sources \\citep[e.g., see the dicussion by ][]{czerny03}. Other geometries are possible, including an accretion disk that evaporates into a hot inner flow \\citep[e.g.,][]{shap76,zdz99}, or a combination of a hot inner flow and a corona that sandwiches the disk \\citep[e.g.,][]{pout97,sob04a}. Furthermore, radiation pressure can drive an outflow from the disk into the corona if the two are cospatial, thus altering the physics of the corona \\citep{proga05}. Investigations of how quasar X-ray parameters depend on black hole mass, $M_{BH}$, and accretion rate relative to Eddington, $\\dot{m}$, offer important constraints on models of the disk/corona system.  There have been attempts to link the evolution of SMBHs to analytic and semi-analytic models of structure formation \\citep[e.g.,][]{kauff00,hatz03,bromley04}, where black holes grow by accreting gas funneled towards the center during a galaxy merger until feedback energy from the SMBH expels gas and shuts off the accretion process \\citep[e.g.,][]{silk98,fabian99,wyithe03,begel05}. This `self-regulated' growth of black holes has recently been successfully applied in smoothed particle hydrodynamics simulations \\citep{dimatt05,spring05}. Within this framework, the AGN or quasar phase occurs during the episode of significant accretion that follows the galaxy merger, persisting until feedback from the black hole `blows' the gas away \\citep[e.g.,][]{hopkins06}. Hydrodynamic calculations have shown that line pressure is more efficient than thermal pressure at driving an outflow \\citep{proga07}, and therefore, the efficiency of AGN feedback depends on the fraction of energy emitted through the UV/disk component as compared to the X-ray/corona component. If the fraction of energy emitted in the UV as compared to the X-ray depends on $M_{BH}$ or $\\dot{m}$, then it follows that the efficiency of AGN feedback will also depend on $M_{BH}$ and $\\dot{m}$. This has important consequences for models of SMBH growth, as the SMBH may become more or less efficient at driving an outflow depending on its mass and accretion rate. Studies of the dependence of quasar X-ray/UV emission on black hole mass and accretion rate are therefore important as they allow us to constrain a $M_{BH}$- or $\\dot{m}$-dependent feedback efficiency.  Numerous previous studies have searched for a luminosity and redshift dependence of $\\alpha_{ox} = -0.384 \\log L_X / L_{UV}$, the ratio of X-ray to UV/optical flux \\citep[e.g.,][]{avni82,wilkes94,yuan98,vig03b,strat05,steffen06,kelly07c}, and $\\Gamma_X$, the X-ray spectra slope \\citep[e.g.,][]{reeves00,bech03,dai04,ris05,grupe06}. Most studies have found a correlation between $\\alpha_{\\rm ox}$ and UV luminosity, $L_{UV}$, while the existence of a correlation between $\\alpha_{\\rm ox}$ and $z$ is still a matter of debate \\citep[e.g.,][]{bech03,vig03b,steffen06,just07,kelly07c}. In addition, studies of $\\Gamma_X$ have produced mixed results. Some authors have claimed a correlation between $\\Gamma_X$ and luminosity \\citep[e.g.,][]{bech03,dai04} or redshift \\citep[e.g.,][]{reeves97, vig99, page03}, while, others find no evidence for a correlation between $\\Gamma_X$ and $L_{UV}$ or $z$ \\citep[e.g.,][]{vig05,ris05,kelly07c}.  A correlation between $\\Gamma_X$ and the $FWHM$ of the H$\\beta$ line has also been found \\citep[e.g.,][]{boller96,brandt97}, suggesting a correlation between $\\Gamma_X$ and black hole mass or Eddington ratio \\citep[e.g.,][]{laor97,brandt98}. Recently, it has become possible to obtain estimates of $M_{BH}$ for broad line AGN by calibrating results from reverberation mapping \\citep{peter04,kaspi05} for use on single-epoch spectra \\citep{wand99,vest02,mclure02,vest06,kelly07b}. This has enabled some authors to confirm a correlation between $\\Gamma_X$ and either $M_{BH}$ or $L_{bol} / L_{Edd}$ \\citep[e.g.,][]{lu99,gier04,porquet04,picon05,shemmer06}, where the X-ray continuum hardens with increasing $M_{BH}$ or softens with increasing $L_{bol} / L_{Edd}$. In addition, previous work has also found evidence for quasars becoming more X-ray quiet as $M_{BH}$ or $L_{bol} / L_{Edd}$ increase \\citep{brunner97,wang04}; however, studies involving the dependence of $\\alpha_{\\rm ox}$ on $M_{BH}$ or $L_{bol} / L_{Edd}$ have remained rare compared to studies of $\\Gamma_X$. It is important to note that the correlations inferred in previous work generally employ broad line mass estimates in combination with a constant bolometric correction. Therefore, most of the correlations found in previous work are, strictly speaking, between $\\Gamma_X$ or $\\alpha_{\\rm ox}$ and the estimates $M_{BH} \\propto L_{\\lambda}^{\\gamma} FWHM^2$ and $L_{bol} / L_{Edd} \\propto L_{\\lambda}^{1 - \\gamma} FWHM^{-2}$, where $\\gamma \\sim 0.5$.  In this work, we investigate the dependence of $\\alpha_{\\rm ox}$ and $\\Gamma_X$ on black hole mass, optical/UV luminosity relative to Eddington, and X-ray luminosity relative to Eddington. We combine the main Sloan Digital Sky Survey (SDSS) sample of \\citet{strat05} with the sample of \\citet{kelly07c}, creating a sample of 318 radio-quiet quasars (RQQ) with X-ray data from \\emph{ROSAT} and \\emph{Chandra}, and optical spectra mostly from the SDSS; 153 of these sources have estimates of $\\Gamma_X$ from \\emph{Chandra}. Because the X-ray emission in radio-loud sources can have an additional component from the jet \\citep[e.g.,][]{zam81,wilkes87}, we focus our analysis on the radio-quiet majority. Our sample has a detection fraction of $87\\%$ and spans a broad range in black hole mass ($10^6 \\lesssim M_{BH} / M_{\\odot} \\lesssim 10^{10}$), redshift ($0 < z < 4.8$), and luminosity ($10^{43} \\lesssim \\lambda L_{\\lambda} (2500$\\AA$) [{\\rm erg\\ s^{-1}}] \\lesssim 10^{48}$), enabling us to effectively look for trends regarding $\\alpha_{\\rm ox}$ and $\\Gamma_X$.  The outline of this paper is as follows. In \\S~\\ref{s-sample} we describe the construction of our sample, and in \\S~\\ref{s-specfits} we describe the procedure we used to fit the optical continuum and emission lines. In \\S~\\ref{s-mbhredd_est} we describe how we obtain broad line mass estimates, our bolometric correction, and argue that a constant bolometric correction provides a poor estimate of the bolometric luminosity. In \\S~\\ref{s-alfox_bh} we describe the results from a regression analysis of $\\alpha_{\\rm ox}$ on $M_{BH}, L_{UV} / L_{Edd},$ and $L_X / L_{Edd}$, and in \\S~\\ref{s-gamx_bh} we report evidence for a non-monotonic dependence of $\\Gamma_X$ on either $M_{BH}, L_{UV} / L_{Edd},$ and $L_X / L_{Edd}$. In \\S~\\ref{s-discussion} we discuss our results within the context of AGN disk/corona models, and we discuss the implications for a dependence of quasar feedback efficiency on black hole mass or accretion rate. In \\S~\\ref{s-summary} we summarize our main results.  We adopt a cosmology based on the the WMAP best-fit parameters \\citep[$h=0.71, \\Omega_m=0.27, \\Omega_{\\Lambda}=0.73$,][]{wmap}. For ease of notation, we define $L_{UV} \\equiv \\nu L_{\\nu} (2500$\\AA$), L_X \\equiv \\nu L_{\\nu} (2\\ {\\rm keV}), l_{UV} \\equiv \\log \\nu L_{\\nu} (2500 $\\AA$),$ and $m_{BH} \\equiv \\log M_{BH} / M_{\\odot}$. ",
        "conclusions": "\\label{s-discussion}  Previous work has found evidence for a correlation between $\\alpha_{\\rm ox}$ and both $M_{BH}$ \\citep{brunner97} and $L_{bol} / L_{Edd}$ \\citep{wang04}, and for a correlation between $\\Gamma_X$ and both $M_{BH}$ \\citep[e.g.,][]{porquet04,picon05} and $L_{bol} / L_{Edd}$ \\citep[e.g.,][]{lu99,gier04,shemmer06}, in aggreement with the results found in this work. However, our study differs from previous work in that we study a large sample of RQQs (318 sources with $\\alpha_{\\rm ox}$, 153 with $\\Gamma_X$) over a broad range in black hole mass ($10^6 \\lesssim M_{BH} / M_{\\odot} \\lesssim 10^{10}$) and redshift $0 < z < 4.8$; to date, this is the largest study of the dependence of the X-ray properties of RQQs on $M_{BH}, L_{UV} / L_{Edd},$ and $L_X / L_{Edd}$. In addition, this the first study of its kind to correct for the intrinsic statistical uncertainty in broad line mass estimates when quantifying the \\emph{intrinsic} trends between the X-ray emission and $M_{BH}, L_{UV} / L_{Edd},$ and $L_X / L_{Edd}$.  Currently, there are two main types of geometries being considered for the comptonizing corona. The first of these is that of a `slab'-type geometry, possibly patchy, that sandwiches the disk \\citep[e.g.,][]{bis77,gal79,nayak00,sob04b}, and the second is that of a hot spherical inner advection dominated flow \\citep[e.g.,][]{shap76,zdz99}; hybrids between the two geometries have also been considered \\citep[e.g.,][]{pout97,sob04a}. There is a growing body of evidence that the advection dominated hot inner flow does not exist in objects with Eddington ratios $L_{bol} / L_{Edd} \\gtrsim 0.01$, as inferred from the existence of a relativistically broadened iron line \\citep[e.g.,][]{mineo00,lee02,fabian02}, relativistically broadened reflection of ionized material \\citep{janiuk01}, and by analogy with galactic black holes \\citep[e.g.,][]{esin97,nowak02}. The range in Eddington ratios probed by our study is at most $0.03 \\lesssim L_{bol} / L_{Edd} \\lesssim 2$, with a mean of $L_{bol} / L_{Edd} \\sim 0.25$. Therefore, the RQQs in our study are likely to have disks that extend approximately down to the last marginally stable orbit, and thus should only have the `slab' type geometries.  \\subsection{Dependence of $\\alpha_{\\rm ox}$ on $M_{BH}$}  \\label{s-alfox_mbh}  In this work we have found that RQQs become more X-ray quiet as $M_{BH}$ increases, and confirmed the well-established relationship between $\\alpha_{\\rm ox}$ and $L_{UV}$. Because $L_{UV}$ increases with $M_{BH}$ and the accretion rate relative to Eddington, $\\dot{m}$, the well-known $\\alpha_{\\rm ox}$--$L_{UV}$ correlation is likely driven by the $\\alpha_{\\rm ox}$--$M_{BH}$ and $\\alpha_{\\rm ox}$--$\\dot{m}$ correlations. A correlation between $\\alpha_{\\rm ox}$ and $M_{BH}$ is expected even if the fraction of the bolometric luminosity emitted by the disk is independent of $M_{BH}$, as the effective temperature of the disk depends on $M_{BH}$. As $M_{BH}$ increases, the effective temperature of the disk decreases, thus shifting the peak of the disk emission toward longer wavelengths. Because the flux density at $2500$\\AA\\ lies redward of the peak in the disk SED over most of the range $M_{BH}$ probed by our study, this shift in the disk SED toward longer wavelengths produces an increase in $L_{UV}$ relative to $L_X$.  We can use the standard thin disk solution to assess the evidence that the fraction of energy emitted by the corona depends on $M_{BH}$. We assume a simple model where the spectrum for the disk emission is that expected for an extended thin accretion disk, and the spectrum for the corona emission is a simple power-law with exponential cutoffs at the low and high energy end. According to \\citet{wandel00}, the spectrum from a radially extended thin accretion disk can be approximated as \\begin{equation} f^{D}_{\\nu} \\approx A_{D} \\left(\\frac{\\nu}{\\nu_{co}}\\right)^{-1/3} e^{-\\nu / \\nu_{co}}, \\label{eq-thindisk} \\end{equation} where $A_{D}$ is the normalization and $\\nu_{co}$ is the cut-off frequency. In this work, we choose the normalization to ensure that Equation (\\ref{eq-thindisk}) integrates to unity, and therefore $f^D_{\\nu}$ gives the shape of the disk emission. For a Kerr black hole, \\citet{malkan91} finds that the cut-off frequency is related to $M_{BH}$ as \\begin{equation} h \\nu_{co} = (6 {\\rm eV}) \\dot{m}^{1/4} (M_{BH} / 10^8 M_{\\odot})^{-1/4}. \\label{eq-cutoff} \\end{equation} We assume that the X-ray emission from the corona can be described by a simple power law with an exponential cutoff at the high and low end: \\begin{equation} f^{C}_{\\nu} = A_C \\nu^{-(\\Gamma_X - 1)} e^{-\\nu / \\nu_{high}} e^{-\\nu_{low} / \\nu}. \\label{eq-corona} \\end{equation} Here, $A_C$ is the corona spectrum normalization, $\\nu_{high}$ is the high energy cutoff, and $\\nu_{low}$ is the low energy cutoff. We choose the low energy cutoff to be $\\nu_{low} = 20\\ {\\rm eV}$, and we choose the high energy cutoff to be $\\nu_{high} = 200\\ {\\rm keV}$ \\citep[e.g.,][]{gilli07}. As with Equation (\\ref{eq-thindisk}), we choose the normalization in Equation (\\ref{eq-corona}) to be equal to unity.  Denoting $f_D$ to be the fraction of bolometric luminosity emitted by the disk, our model RQQ spectrum is then \\begin{equation} L_{\\nu} \\approx L_{bol} \\left[f_D f^D_{\\nu} + (1 - f_D) f^C_{\\nu} \\right]. \\label{eq-modelspec} \\end{equation} We computed Equation (\\ref{eq-modelspec}) assuming a value of $\\dot{m} = 0.2$ and $f_D = 0.85$. We chose the value of the $\\dot{m} = 0.2$ because it is representative of the RQQs in our sample, and we chose the value $f_D = 0.85$ because it gives values of $\\alpha_{\\rm ox}$ typical of the RQQs in our sample. We vary $M_{BH}$ but keep $f_D$ and $\\dot{m}$ constant because we are interested in investigating whether there is evidence that assuming independence between $M_{BH}$ and both $f_D$ and $\\dot{m}$ is inconsistent with our $\\alpha_{\\rm ox}$ results.  We compute Equation (\\ref{eq-modelspec}) for two forms of the dependence of $\\Gamma_X$ on $M_{BH}$. For the first model, we assume a constant value of $\\Gamma_X = 2$. For the second model, we assume that $\\Gamma_X$ depends on $M_{BH}$ according to our best fit regression results, where $\\Gamma_X$ depends on $M_{BH}$ according to Equation (\\ref{eq-gamxmbh_hbeta}) for $M_{BH} \\lesssim 3 \\times 10^8 M_{\\odot}$, and $\\Gamma_X$ depends on $M_{BH}$ according to Equation (\\ref{eq-gamxmbh_civ}) for $M_{BH} \\gtrsim 3 \\times 10^8 M_{\\odot}$. We ignore the intrinsic dispersion in $\\Gamma_X$. In Figure \\ref{f-modelspec} we show the spectra computed from Equation (\\ref{eq-modelspec}) for RQQs with $M_{BH} / M_{\\odot} = 10^7, 10^8, 10^9,$ and $10^{10}$. The dependence of the location of the peak in the disk emission on $M_{BH}$ is clearly illustrated.  In Figure \\ref{f-alfox_model} we compare the $\\alpha_{\\rm ox}$--$M_{BH}$ regression results with the dependence of $\\alpha_{\\rm ox}$ on $M_{BH}$ expected from Equation (\\ref{eq-modelspec}) for both $\\Gamma_X$--$M_{BH}$ models. Under the thin disk approximation, a correlation is expected between $\\alpha_{\\rm ox}$ and $M_{BH}$, even if the fraction of bolometric luminosity emitted by the disk is independent of $M_{BH}$. However, our data are inconsistent with the assumption that $f_D$ and $M_{BH}$ are independent, given the thin disk approximation. Under the assumption that $f_D$ and $M_{BH}$ are independent, the $\\alpha_{\\rm ox}$--$M_{BH}$ correlation is too flat, and a increase in the fraction of bolometric luminosity emitted by the disk with increasing $M_{BH}$ is needed to match the steeper observed dependence of $\\alpha_{\\rm ox}$ on $M_{BH}$. Alternatively, if $f_D$ increases with increasing $\\dot{m}$, as we argue in \\S~\\ref{s-alfox_mdot}, then a steeper $\\alpha_{\\rm ox}$--$M_{BH}$ correlation also results if $M_{BH}$ and $\\dot{m}$ are correlated. In this case, if $f_D$ increases with increasing $\\dot{m}$, and if $M_{BH}$ increases with increasing $\\dot{m}$, then $f_D$ will also increase with increasing $M_{BH}$, thus producing a steeper observed $\\alpha_{\\rm ox}$--$M_{BH}$ correlation.  To the extant that Equations (\\ref{eq-thindisk})--(\\ref{eq-modelspec}) accurately approximate the spectral shape of RQQs, our data imply that either the fraction of bolometric luminosity emitted by the disk increases with increasing $M_{BH}$, that $M_{BH}$ and $\\dot{m}$ are correlated, or both. Some theoretical models have suggested that the fraction of bolometric luminosity emitted by the disk should depend on $\\dot{m}$, but be relatively insensitive to $M_{BH}$ \\citep[e.g.,][]{czerny03,liu03}. Therefore, while a significant dependence of $f_D$ on $M_{BH}$ is not predicted by these disk/corona models, these models are still consistent with the interpretation that a $M_{BH}$--$\\dot{m}$ correlation is driving the steeper $\\alpha_{\\rm ox}$--$M_{BH}$ correlation. Unfortunately, without accurate estimates of $\\dot{m}$ we are unable to distinguish between these two possibilities.  A shift in the peak of the disk SED with $M_{BH}$ may also explain the dependence of $\\alpha_{\\rm ox}$ on redshift observed by K07. K07 speculated that the observed hardening of $\\alpha_{\\rm ox}$ with increasing $z$ at a given $L_{UV}$ may be due to a correlation between $\\alpha_{\\rm ox}$ and $M_{BH}$, manifested through a $M_{BH}$--$z$ correlation. At a given $L_{UV}$, an increase in $M_{BH}$ will result in an increase in $L_X$ relative to $L_{UV}$, assuming that $\\dot{m}$ is not strongly correlated with $M_{BH}$. This is because an increase in $M_{BH}$ decreases the temperature of the disk, shifting the peak in the disk SED toward the red, and thus increasing the luminosity at 2500\\AA. However, since K07 investigated the dependence of $\\alpha_{\\rm ox}$ on $z$ at a given $L_{UV}$, the luminosity at 2500\\AA\\ is held constant. Therefore, the overall disk emission must decrease in order to keep the luminosity at 2500\\AA\\ constant despite the increase in $M_{BH}$. As a result, an increase in $M_{BH}$ at a given $L_{UV}$ will result in an increase in the X-ray luminosity relative to the luminosity at 2500\\AA. Because $M_{BH}$ and $z$ are correlated in our flux limited sample (e.g., see Figure \\ref{f-mbh_vs_z}), an increase in $z$ will probe RQQs with higher $M_{BH}$. As a result, RQQs will become more X-ray loud with increasing $z$, at a given 2500\\AA\\ luminosity. Consequently, deeper surveys that probe a greater range of $M_{BH}$ should not see as strong of a correlation between $M_{BH}$ and $z$, thereby reducing the magnitude of a $\\alpha_{\\rm ox}$--$z$ correlation. Indeed, investigations based on samples that span a greater range in luminosity do not find evidence for a correlation between $\\alpha_{\\rm ox}$ and $z$ \\citep[e.g.,][]{steffen06,just07}, qualitatively consistent with our interpretation of a $\\alpha_{\\rm ox}$--$z$ correlation.  \\subsection{Dependence of $\\alpha_{\\rm ox}$ on $\\dot{m}$}  \\label{s-alfox_mdot}  We have found that $\\alpha_{\\rm ox}$ increases with increasing $L_{UV} / L_{Edd}$, and decreases with increasing $L_X / L_{Edd}$. The mere existence of these correlations is not particularly interesting, as we would expect that the ratio of optical/UV luminosity to X-ray luminosity would increase as the fraction of optical/UV luminosity relative to Eddington increases, and vice versa for an increase in $L_X / L_{Edd}$. However, the relative magnitude of these dependencies carries some information regarding the dependence of $\\alpha_{\\rm ox}$ on $\\dot{m}$. A correlation between $\\alpha_{\\rm ox}$ and $L_{UV} / L_{Edd}$ implies that $L_{UV} / L_X$ increases as the quantity $f_{UV} \\dot{m}$ increases, where $f_{UV}$ is the fraction of bolometric luminosity emitted at 2500\\AA. Likewise, an anti-correlation between $\\alpha_{\\rm ox}$ and $L_X / L_{Edd}$ implies that $L_{UV} / L_X$ decreases as the quantity $f_X \\dot{m}$ increases, where $f_X$ is the fraction of bolometric luminosity emitted at 2 keV. If the fraction of the bolometric luminosity emitted by the disk, as compared to the corona, increases with increasing $\\dot{m}$, then we would expect a strong increase in $L_{UV} / L_X$ with the product $f_{UV} \\dot{m}$, resulting from the dual dependency of $L_{UV} / L_X$ on $f_{UV}$ and $\\dot{m}$. Furthermore, because the fraction of bolometric luminosity emitted by the disk should decrease with increasing $f_X$, then, if the fraction of bolometric luminosity emitted by the disk increases with increasing $\\dot{m}$, we would expect a weaker dependence of $L_{UV} / L_X$ on the quantity $f_X \\dot{m}$. This is because an increase in $\\dot{m}$ causes an increase in the disk emission relative to the corona emission, which will then work against the decrease in disk emission relative to the corona that results from an increase in $f_X$. The end result is a weaker dependence of $L_{UV} / L_X$ on the product $f_X \\dot{m}$. Indeed, this is what we observe, where $L_{UV} / L_X \\propto (L_{UV} / L_{Edd})^{2.5}$ and $L_{UV} / L_X \\propto (L_X / L_{Edd})^{-1.5}$. Therefore, we conclude that the disk emission relative to the corona emission increases with increasing $\\dot{m}$. This is in agreement with some models of corona with a slab geometry \\citep[e.g.,][]{czerny97,janiuk00,merloni02,liu03}, where the $\\alpha_{\\rm ox}$--$\\dot{m}$ correlation arises due to a dependency of the size of the corona on $\\dot{m}$.  Our result that $\\alpha_{\\rm ox}$ is correlated with $L_{UV} / L_{Edd}$ and anti-correlated with $L_X / L_{Edd}$ is inconsistent with a constant bolometric correction to both the optical/UV and X-ray luminosities. Instead, an increase in $L_{UV} / L_X$ with increasing $\\dot{m}$ implies that the bolometric correction depends on $\\dot{m}$. Because we conclude that the fraction of bolometric luminosity emitted by the disk increases with increasing $\\dot{m}$, this implies that the bolometric correction to the optical/UV luminosity decreases with increasing $\\dot{m}$, while the bolometric correction to the X-ray luminosity increases with increasing $\\dot{m}$. The direction of this trend is consistent with the results of \\citet{vasud07}, who find that the bolometric correction to the X-ray luminosity increases with increasing $L_{bol} / L_{Edd}$. Similarly, we have found evidence that the fraction of bolometric luminosity emitted by the disk depends on $M_{BH}$, therefore implying that the bolometric correction also depends on $M_{BH}$. Even if the fraction of bolometric luminosity emitted by the disk is independent of $M_{BH}$, the bolometric correction will still depend on $M_{BH}$ because the location of the peak in the disk emission will shift toward longer wavelengths as $M_{BH}$ increases. As $M_{BH}$ varies, the luminosity at 2500\\AA\\ probes a different region of the quasi-blackbody disk emission, thereby producing a dependence of bolometric correction on $M_{BH}$.  \\subsection{Implications for Black Hole Feedback}  A significant amount of recent work suggests that radiative and mechanical feedback energy from AGN plays an important part in galaxy and supermassive black hole coevolution \\citep[e.g.,][]{fabian99,wyithe03,hopkins05}. Within the context of these models, a nuclear inflow of gas, possibly the result of a galaxy merger, feeds the SMBH, thus igniting a quasar. The SMBH grows until feedback energy from the quasar is able to drive out the accreting gas, thus halting the accretion process. Hydrodynamic calculations of accretion flows have shown that the efficiency of the quasar in driving an outflow depends on the fraction of energy emitted through he UV/disk component as compared to the X-ray/corona component \\citep{proga07}. The disk component produces luminosity in the UV, which is responsible for driving an outflow via radiation pressure on lines, whereas the corona component produces luminosity in the X-rays, which is responsible for driving an outflow via thermal expansion. Calculations by \\citet{proga07} have shown that radiation driving produces an outflow that carries more mass and energy than thermal driving. If the efficiency of black hole feedback depends on the quasar SED, any dependence on $M_{BH}$ and $\\dot{m}$ of the fraction of AGN energy emitted in the UV as compared to the X-ray has important consequences for models of black hole growth.  Because we have found evidence that the fraction of bolometric luminosity emitted by the disk increases with increasing $\\dot{m}$ and $M_{BH}$, this implies that black holes becomes more efficient at driving an outflow with increasing $\\dot{m}$ and $M_{BH}$. However, the $\\alpha_{\\rm ox}$--$M_{BH}$ correlation may be due to the combination of both a correlation between $M_{BH}$ and $\\dot{m}$, and a dependence of the location peak in the disk SED on $M_{BH}$. If the fraction of energy emitted by the disk only depends weakly on $M_{BH}$, as some theoretical models have suggested \\citep[e.g.,][]{czerny03,liu03}, the fraction of energy emitted in the UV will still decrease with increasing $M_{BH}$ becuase the peak of the disk emission will shift away from the UV. In this case, at a given $\\dot{m}$ we would expect that black holes will become less efficient at driving an outflow with increasing $M_{BH}$.  \\subsection{Dependence of $\\Gamma_X$ on $M_{BH}$ and $\\dot{m}$}  In this work we have also found evidence that $\\Gamma_X$ and $M_{BH}, L_{UV} / L_{Edd},$ and $L_X / L_{Edd}$ are not statistically independent. Moreover, the dependence of $\\Gamma_X$ on black hole mass or Eddington ratio appears to follow a non-monotonic form, although the $\\Gamma_X$--$M_{BH}$ trend is weak compared to the dependency of $\\Gamma_X$ on Eddington ratio. For the $\\Gamma_X$--$M_{BH}$ relationship, the X-ray continuum hardens with increasing black hole mass until $M_{BH} \\sim 3 \\times 10^8 M_{\\odot}$, after which the X-ray continuum softens with increasing black hole mass. The opposite is true for the $\\Gamma_X$--$L_{UV} / L_{Edd}$ and $\\Gamma_X$--$L_X / L_{Edd}$ trends, and further work is needed to confirm this result. Previous studies have not seen this non-monotonic trend because they have only employed the H$\\beta$ emission line, and therefore their samples have been dominated by low-$z$, low-$M_{BH}$ sources.  \\subsubsection{Selection Effects}  \\label{s-seleff}  It is unlikely that the dependence of $\\Gamma_X$ on $M_{BH}, L_{UV} / L_{Edd},$ or $L_X / L_{Edd}$ is due to redshifting of the observable spectra range. If this were the case, then as $M_{BH}$ increases, so does $z$ due to selection effects, and thus we would observe a decrease in $\\Gamma_X$ as the `soft excess' shifts out of the observed 0.3--7 keV spectral range, while the compton reflection component shifts into the observed spectral range. However, there are lines of evidence that suggest that the $\\Gamma_X$ correlation are not due to redshifting of the observable spectral region, and that at least some of the observed dependency of $\\Gamma_X$ on $M_{BH}, L_{UV} / L_{Edd},$ and $L_X / L_{Edd}$ is real. First, in \\S~\\ref{s-ztest} we tested whether a sample of nine $z > 1$ test sources with $M_{BH} \\lesssim 3 \\times 10^8 M_{\\odot}$ were better described by a regression fit using the $z > 1.5, M_{BH} \\gtrsim 3 \\times 10^8 M_{\\odot}$ sources, or by a regression fit using the $z < 1, M_{BH} \\lesssim 3 \\times 10^8 M_{\\odot}$ sources. We found that the test sources were better fit using the regression of similar $M_{BH}$, and therefore that the difference in the $\\Gamma_X$--$M_{BH}$ correlations primarily depends on $M_{BH}$. Second, similar trends at low redshift between $\\Gamma_X$ and $M_{BH}$ or $L_{bol} / L_{Edd}$ have been seen in other studies that only analyze the hard X-ray spectral slope \\citep[typically 2--12 keV, e.g.,][]{picon05,shemmer06}, and thus these studies are not effected by the soft excess. Third, the compton reflection hump is unlikely to shift into the observabable spectral range until $z \\sim 1$. However, the contribution to the inferred $\\Gamma_X$ from compton reflection at $z \\gtrsim 1$ is likely weak, if not negligible, as our $z \\gtrsim 1$ sources have $M_{BH} \\gtrsim 10^8 M_{\\odot}$ and are highly luminous, and therefore are expected to only have weak reflection components \\citep{mineo00,ball01,bianchi07}.  There are two scenarios in which the non-monotonic behavior of $\\Gamma_X$ with $M_{BH}$ or Eddington ratio may be artificially caused by selection. We will focus on the Eddington ratio dependency, as it is the strongest; however, our argument also applies to $M_{BH}$. First, the intrinsic dependency of $\\Gamma_X$ on Eddington ratio could be linear with increasing intrinsic scatter at high $L_{bol} / L_{Edd}$. Then, an inferred non-monotonic trend would occur if we were to systematically miss quasars with high $L_{bol} / L_{Edd}$ and steep X-ray spectra. Alternatively, there could be no intrinsic dependency of $\\Gamma_X$ on Eddington ratio. In this case, we would infer a non-monotonic trend if we were to systematically miss quasars with steep X-ray spectra at low and high $L_{bol} / L_{Edd}$, and quasars with flat X-ray spectra at moderate $L_{bol} / L_{Edd}$.  We do not consider it likely that the observed non-monotonic dependence of $\\Gamma_X$ on Eddington ratio is due solely to selection effects. K07 describes the sample selection for sources with $\\Gamma_X$. With the exception of some of the $z > 4$ quasars, all sources from K07 were selected by cross-correlating the SDSS DR3 quasars with public \\emph{Chandra} observations. Almost all SDSS sources in K07 had serendipitious \\emph{Chandra} observations, and therefore were selected without regard to their X-ray properties. K07 estimated $\\Gamma_X$ for all sources that were detected in X-rays at the level of $3\\sigma$ or higher. Therefore, the only additional criterion beyond the SDSS selection imposed by K07 is the requirement that the source had to be detected in X-ray, which was fulfilled by $90\\%$ of the quasars; the undetected sources were slightly more likely to be found at lower redshift, probably due to the presence of unidentified BAL quasars. As a result, the K07 sample selection function is essentially equivalent to the SDSS quasar selection function. Because the SDSS selects quasars based on their optical colors, the most likely cause of selection effects is the optical color selection. There is evidence that $\\Gamma_X$ is correlated with the slope of the optical continuum, where the X-ray continuum flattens (hardens) as the optical continuum steepens (softens) \\citep[][K07]{gall05}. The SDSS selection probability is lower for red sources \\citep{rich06}, so we might expect to systematically miss sources with smaller $\\Gamma_X$. However, for the two scenarios described above, this is opposite the trend needed to explain the $\\Gamma_X$--$L_X / L_{Edd}$ relationship, where we need to at least systematically miss sources with larger $\\Gamma_X$. Furthermore, the drop in SDSS selection efficiency with optical spectral slope only occurs at $2 < z < 4$ \\citep{rich06}, thus we would expect a redshift dependence for this selection effect. As we have argued above, and in \\S~\\ref{s-ztest}, the non-monotonic trends for $\\Gamma_X$ cannot be completely explained as the result of different redshift ranges being probed.  \\subsubsection{Implications for Accretion Physics}  \\label{s-accphys}  The dependence of $\\Gamma_X$ on $L_{UV} / L_{Edd}$ and $L_X / L_{Edd}$ is likely due to a dependence of $\\Gamma_X$ on $\\dot{m}$. If these $\\Gamma_X$ correlations were due to a dependence of $\\Gamma_X$ on $f_{UV}$ or $f_X$, then we would expect opposite trends for $L_{UV} / L_{Edd}$ and $L_X / L_{Edd}$, as $f_{UV}$ and $f_X$ should be anti-correlated. The fact that the regression results for the $\\Gamma_X$--$L_{UV} / L_{Edd}$ and $\\Gamma_X$--$L_X / L_{Edd}$ relationships are similar implies that $\\Gamma_X$ depends on $\\dot{m}$, and at most only weakly on $f_{UV}$ or $f_X$.  A non-monotonic dependence of $\\Gamma_X$ on $\\dot{m}$ is predicted from the accreting corona model of \\citet{janiuk00}, as well as a non-monotonic dependence of $\\Gamma_X$ on the viscosity \\citep{bech03}. In addition, $\\Gamma_X$ is expected to steepen with increasing optical depth \\citep[e.g.,][]{haardt91,haardt93,czerny03}. One could then speculate that the dependence of $\\Gamma_X$ on $M_{BH}$ or $\\dot{m}$ is due to a non-monotonic dependence of the corona optical depth on $\\dot{m}$, which may indicate a change in the structure of the disk/corona system at $\\sim 3 \\times 10^8 M_{\\odot}$ or some critical $\\dot{m}$. Recent work also suggests a non-monotonic dependence of the optical/UV spectral slope, $\\alpha_{UV}$, on $\\dot{m}$ \\citep{bonning07,davis07}. From this work, it has been inferred that the optical/UV continuum becomes more red with increasing $\\dot{m}$ until $L_{bol} / L_{Edd} \\approx 0.3$, after which the optical/UV continuum becomes more blue with increasing $\\dot{m}$. Assuming the bolometric corrections described in \\S~\\ref{s-eddrat}, the turnover in the $\\Gamma_X$--$L_{bol} / L_{Edd}$ relationship also occurs at $L_{bol} / L_{Edd} \\approx 0.3$. \\citet{bonning07} suggested that the turnover in spectral slope at $L_{bol} / L_{Edd} \\sim 0.3$ may be due to a change in accretion disk structure, where the inner part of the accretion disk becomes thicker due to increased radiation pressure \\citep{abram88}. \\citet{bonning07} performed a simple approximation to this `slim disk' solution and found that it is able to produce a non-monotonic trend between optical color and Eddington ration. If the inner disk structure changes at high $\\dot{m}$, this change could alter the corona structure, producing the observed trend between $\\Gamma_X$ and Eddington ratio.  Unfortunately, current models for corona geometry make a number of simplifying assumptions, and do not yet predict a specific relationship between $\\alpha_{\\rm ox}, \\Gamma_X, M_{BH},$ and $\\dot{m}$. Ideally, full magneto-hydrodynamic simulations \\citep[e.g.,][]{devill03,turner04,krolik05} that include accretion disk winds \\citep[e.g.,][]{murray95,proga04} should be used to interpret the results found in this work. However, MHD simulations have not advanced to the point where they predict the dependence of $\\alpha_{\\rm ox}$ and $\\Gamma_X$ on quasar fundamental parameters, but hopefully recent progress in analytical descriptions of the magneto-rotational instability \\citep{pessah06,pessah07} will help to overcome some of the computational difficulties and facilitate further advancement.    In this work we have investigated the dependence of $\\alpha_{\\rm ox}$ and $\\Gamma_X$ on black hole mass and Eddington ratio using a sample of 318 radio-quiet quasars with X-ray data from \\emph{ROSAT} \\citep{strat05} and \\emph{Chandra} \\citep{kelly07c}, and optical data mostly from the SDSS; 153 of these sources have estimates of $\\Gamma_X$ from \\emph{Chandra}. Our sample spans a broad range in black hole mass ($10^6 \\lesssim M_{BH} / M_{\\odot} \\lesssim 10^{10}$), redshift ($0 < z < 4.8$), and luminosity ($10^{43} \\lesssim \\lambda L_{\\lambda} (2500$\\AA$) [{\\rm erg\\ s^{-1}}] \\lesssim 10^{48}$). To date, this is the largest study of the dependence of RQQ X-ray parameters on $M_{BH}, L_{UV} / L_{Edd},$ and $L_X / L_{Edd}$. Our main results are summarized as follows: \\begin{itemize} \\item We show that $\\alpha_{\\rm ox}$ is correlated with $L_{UV} / L_{Edd}$ and anti-correlated with $L_X / L_{Edd}$. This result is inconsistent with a constant bolometric correction being applicable to both the optical/UV luminosity and the X-ray luminosity. This result, when taken in combination with recent work by \\citet{vasud07}, implies that constant bolometric corrections can be considerably unreliable and lead to biased results. Instead, we argue that $L_{UV} / L_X$ increases with increasing $\\dot{m}$ and increasing $M_{BH}$, therefore implying that the bolometric correction depends on $\\dot{m}$ and $M_{BH}$. \\item We performed a linear regression of $\\alpha_{\\rm ox}$ on luminosity, black hole mass, $L_{UV} / L_{Edd}$, and $L_X / L_{Edd}$, and found significant evidence that $\\alpha_{\\rm ox}$ depends on all four quantities: $L_{UV} / L_{X} \\propto L_{UV}^{0.31 \\pm 0.03}, L_{UV} / L_{X} \\propto M_{BH}^{0.43 \\pm 0.06}, L_{UV} / L_{X} \\propto (L_{UV} / L_{Edd})^{2.57 \\pm 0.45},$ and $L_{UV} / L_X \\propto L_{X} / L_{Edd}^{-1.48 \\pm 0.14}$. The dependence of $\\alpha_{\\rm ox}$ on $L_{UV}$ may be due to the dual dependence of $\\alpha_{\\rm ox}$ on $M_{BH}$ and $\\dot{m}$. Because we have attempted to correct for the statistical uncertainties in $\\alpha_{\\rm ox}$ and the broad line estimates of $M_{BH}$, these results refer to the \\emph{intrinsic} relationships involving $\\alpha_{\\rm ox}$ and $M_{BH}$, and are not merely the relationships between $\\alpha_{\\rm ox}$ and the broad line mass estimates, $\\hat{M}_{BL} \\propto L^{\\gamma} FWHM^2$. \\item A correlation between $\\alpha_{\\rm ox}$ and $M_{BH}$ is expected from the fact that the peak in the disk emission will shift to longer wavelengths as $M_{BH}$ increases, even if the fraction of the bolometric luminosity emitted by the disk does not change with $M_{BH}$. Using a simple model for RQQ spectra, we argue that the observed $\\alpha_{\\rm ox}$--$M_{BH}$ correlation is steeper than that expected if both $\\dot{m}$ and the fraction of bolometric luminosity produced by the disk are independent of $M_{BH}$. The observed $\\alpha_{\\rm ox}$--$M_{BH}$ relationship therefore implies that either the fraction of bolometric luminosity emitted by the disk increases with increasing $M_{BH}$, that $M_{BH}$ is correlated with $\\dot{m}$, or both. \\item A correlation between $\\alpha_{\\rm ox}$ and $\\dot{m}$ is predicted from several models of `slab'-type corona. We argue that the weaker dependence of $\\alpha_{\\rm ox}$ on $L_X / L_{Edd}$ implies that $L_{UV} / L_X$ increases with increasing $\\dot{m}$. Considering that the efficiency of quasar feedback energy in driving an outflow may depend on the ratio of UV to X-ray luminosity, a correlation between $\\alpha_{\\rm ox}$ and both $M_{BH}$ and $\\dot{m}$ has important consequences for models of black hole growth. In particular, if supermassive black holes become more X-ray quiet at higher $\\dot{m}$, they will become more efficient at driving away their accreting gas, thus halting their growth. \\item Because of a possible nonlinear dependence of $\\Gamma_X$ on $M_{BH}, L_{UV} / L_{Edd}$, or $L_X / L_{Edd}$, we performed seperate regressions for the black hole mass estimates obtained from each emission line. We confirmed the significant dependence of $\\Gamma_X$ on $L_{UV} / L_{Edd}$ and $L_X / L_{Edd}$ seen in previous studies as inferred from the broad line mass estimates based on the H$\\beta$ line; however, we also find evidence that the $\\Gamma_X$ correlations change direction when including the C IV line. In particular, for the H$\\beta$ sample, the X-ray continuum hardens with increasing $M_{BH}$, while for the C IV sample, the X-ray continuum softens with increasing $M_{BH}$. Similar but opposite trends are seen with respect to $L_{UV} / L_{Edd}$ and $L_X / L_{Edd}$, and we conclude that these relationships can be interpreted as resulting from a correlation between $\\Gamma_X$ and $\\dot{m}$. Results obtained from the Mg II line were too uncertain to interpret. We analyzed two test samples to argue that this non-monotonic behavior is not due to the different redshifts probed by the two samples, or to problems with the estimates of $M_{BH}$ derived from the two lines; the different trends may be due to the difference in $M_{BH}$ probed by the two samples. A non-monotonic dependence of $\\Gamma_X$ on $M_{BH}$ and/or $\\dot{m}$ may imply a change in the disk/corona structure, although a non-monotonic dependence of $\\Gamma_X$ on $\\dot{m}$ and the viscosity is predicted by some models of `slab'-type coronal geometries. \\end{itemize}"
    },
    "0801/0801.3450_arXiv.txt": {
        "abstract": "A simple model of quintessential inflation with the modified exponential potential $e^{-\\alpha \\phi} \\left[A+\\left(\\phi-\\phi_0\\right)^2\\right]$ is analyzed in the braneworld context. Considering reheating via instant preheating, it is shown that the evolution of the scalar field $\\phi$ from inflation to the present epoch is consistent with the observational constraints in a wide region of the parameter space. The model exhibits transient acceleration at late times for $0.96 \\lesssim A \\alpha^2 \\lesssim 1.26$ and $271 \\lesssim \\phi_0\\, \\alpha \\lesssim 273$, while permanent acceleration is obtained for $2.3\\times10^{-8} \\lesssim A \\alpha^2 \\lesssim 0.98$ and $255 \\lesssim \\phi_0\\, \\alpha \\lesssim 273$. The steep parameter $\\alpha$ is constrained to be in the range $5.3 \\lesssim \\alpha \\lesssim 10.8$. ",
        "introduction": "The recent measurements of the Wilkinson Microwave Anisotropy Probe (WMAP)~\\cite{Bennett:2003bz,Spergel:2003cb,Hinshaw:2006ia,Page:2006hz,Spergel:2006hy,Hinshaw:2008kr,Dunkley:2008ie,Komatsu:2008hk} on the cosmic microwave background (CMB) made it clear that the current state of the Universe is very close to a critical density and that the primordial density perturbations that seeded large-scale structure in the Universe are nearly scale-invariant and Gaussian, which is consistent with the inflationary paradigm. Inflation is often implemented in models with a single or multiple scalar fields~\\cite{Lyth:1998xn}, which undergo a slow-roll period allowing an early accelerated expansion of the Universe.  Furthermore, the Universe seems to exhibit an interesting symmetry with regard to the accelerated expansion, namely, it underwent inflation at early epochs and is believed to be accelerating at present. The current acceleration of the Universe is supported by observations of high redshift type Ia supernovae~\\cite{Astier:2005qq,Riess:2006fw} and, more indirectly, by observations of the CMB and galaxy clustering~\\cite{Spergel:2006hy,Komatsu:2008hk,Seljak:2006bg}. Within the framework of general relativity, cosmic acceleration should be sourced by an energy-momentum tensor which has a large negative pressure (dark energy)~\\cite{Sahni:1999gb,Padmanabhan:2002ji,Ratra:1987rm,Wetterich:1987fm,Frieman:1995pm,Ferreira:1997au,Zlatev:1998tr,Brax:1999yv,Barreiro:1999zs,Bento:2001yv}. Therefore, in order to comply with the logical consistency and observations, the standard model should be valid somewhere between inflation at early epochs and quintessence at late times. It is then natural to ask whether one can build a model to join these two ends without disturbing the thermal history of the Universe. Attempts have been made to unify both these concepts using models with a single scalar field~\\cite{Peebles:1998qn,Copeland:2000hn,Huey:2001ae,Majumdar:2001mm,Nunes:2002wz,Sami:2004xk,Leith:2007bu,Neupane:2007mu}, i.e., in which a single scalar field plays the role of the inflaton and quintessence - the so-called quintessential inflation.  On the other hand, in recent years there has been increasing interest in the cosmological implications of a certain class of braneworld scenarios where the Friedmann equation is modified at very high energies. In particular, in the Randall-Sundrum type II (RSII) model~\\cite{Randall:1999vf} the square of the Hubble parameter, $H^2$, acquires a term quadratic in the energy density, allowing slow-roll inflation to occur for potentials that would be too steep to support inflation in the standard Friedmann-Robertson-Walker (FRW) cosmology~\\cite{Cline:1999ts,Kaloper:1999sm,Shiromizu:1999wj,Binetruy:1999ut,Flanagan:1999cu,Maartens:1999hf,Maartens:2003tw,Bento:2006sr}. Indeed, in a cosmological scenario in which the metric projected onto the brane is a spatially flat FRW model, the Friedmann equation in four dimensions reads (after setting the 4D cosmological constant to zero and assuming that inflation rapidly makes any dark radiation term negligible)~\\cite{Cline:1999ts} \\begin{equation} H^2 = {1 \\over 3\\, M_4^2}\\, \\rho\\, \\left[1 + {\\rho \\over 2 \\lambda}\\right]~. \\label{eq:Friedmann} \\end{equation} Here $M_4$ is the 4D reduced Planck mass and $\\rho\\equiv \\rho_{\\phi} = \\frac{1}{2}{\\dot\\phi}^2 + V(\\phi)$ in a Universe dominated by a single minimally coupled homogeneous scalar field. The brane tension $\\lambda$ relates the 4D and 5D Planck masses through the relation \\begin{equation} \\lambda = \\frac{3}{32 \\pi^2} \\frac{M_5^6}{M_4^2}~, \\label{eq:tension} \\end{equation} where $M_5$ is the 5D Planck mass.  We notice that Eq.~(\\ref{eq:Friedmann}) reduces to the usual Friedmann equation at sufficiently low energies, \\mbox{$\\rho \\ll \\lambda$}, while at very high energies $H\\propto\\rho$. In this scenario, all matter fields are confined to the brane and, hence, inflation is driven by a 4D scalar field trapped on the brane with the usual equation of motion \\begin{equation} {\\ddot \\phi} + 3H {\\dot \\phi} + V'(\\phi) = 0~, \\label{eq:kg} \\end{equation} where the prime denotes the derivative with respect to the scalar field $\\phi$. From Eqs.~(\\ref{eq:Friedmann}) and (\\ref{eq:kg}) it becomes clear that the presence of the additional term $\\sim \\rho^2/\\lambda$ increases the damping experienced by the scalar field as it rolls down its potential.  It has been shown that, in the RSII braneworld context, quintessential inflation can occur for a sum of exponentials or cosh potentials~\\cite{Majumdar:2001mm,Nunes:2002wz,Sami:2004xk}. In this paper we show that the modified exponential potential (hereafter we adopt natural units, $M_4=1$, unless stated otherwise) \\begin{equation} V(\\phi)=e^{-\\alpha \\phi} \\left[A+\\left(\\phi-\\phi_0\\right)^2\\right] \\label{eq:albrecht} \\end{equation} also leads to a successful quintessential inflation model.  In the context of quintessence, this potential was first analyzed by Albrecht and Skordis (AS)~\\cite{Albrecht:1999rm}. Afterward, it has been extensively studied in the literature~\\cite{Barrow:2000nc,Skordis:2000dz,Barreiro:2003ua,Blais:2004vt,Barnard:2007ta}. Regarding its motivation, it is worth noticing that exponential potentials naturally appear in 4D field theories coming from string/M-theory~\\cite{Gasperini:2001pc}, where the scalar field $\\phi$ is typically identified with the dilaton field. As far as the origin of the polynomial factor is concerned, it can be associated with a nontrivial K\\\"ahler term in an effective 4D supergravity theory~\\cite{Copeland:2000vh}. The scalar $\\phi$ could also be associated with a modulus (radion) field in curled extra dimensions~\\cite{Albrecht:2001xt}, which need not be universally coupled to matter/gauge fields and, therefore, is not subject to quantum corrections~\\cite{Doran:2002bc}.  The tracking properties of the AS potential are similar to the pure exponential, namely, it allows sufficient radiation domination during big bang nucleosynthesis (BBN), followed by matter domination. Nevertheless, due to the presence of the polynomial factor, the field evolves to quintessence dominance near the present epoch. One should notice that, in order to the transition to happen near the present, the parameter $\\phi_0$ must be suitably chosen. In other words, this model does not explain the so-called coincidence problem. However, the model displays an interesting feature: it can lead to both permanent and transient acceleration regimes. Permanent acceleration occurs for $A\\alpha^2<1$, when the field is trapped in the local minimum of the potential. Transient vacuum domination arises in two ways~\\cite{Barrow:2000nc}: when $A\\alpha^2<1$ and the $\\phi$ field arrives at the minimum of the potential with enough kinetic energy to roll over the barrier and resumes descending the potential where $\\phi \\gg \\phi_0$, or for $A\\alpha^2>1$, when the potential loses its local minimum.  In the models mentioned above, inflation takes place when the pure exponential potential dominates the potential. The exit from inflation takes place naturally when the slow-roll conditions are violated because the high-energy brane corrections become unimportant. Moreover, these models belong to the category of nonoscillating models in which the standard reheating mechanism does not work. In this case, one can employ an alternative mechanism of reheating via  gravitational particle production~\\cite{Ford:1986sy,Spokoiny:1993kt,Grishchuk:1990bj}. However, this mechanism is faced with difficulties associated with excessive production of gravity waves. Indeed, the reheating mechanism based upon this process is extremely inefficient. The energy density of the produced radiation is typically one part in $10^{16}$~\\cite{Copeland:2000hn} to the scalar field energy density at the end of inflation. As a result, these models have a prolonged kinetic regime during which the amplitude of primordial gravity waves is enhanced and the nucleosynthesis constraints are violated~\\cite{Sahni:2001qp}. These problems can be circumvented if one invokes an alternative method of reheating, namely, the so-called instant preheating~\\cite{Felder:1998vq,Kofman:1994rk,Kofman:1997yn}. This mechanism is quite efficient and robust, and is well suited to nonoscillating models~\\cite{Felder:1999pv}. The larger reheating temperature in this model results in a smaller amplitude of relic gravity waves which is consistent with the BBN bounds~\\cite{Sami:2004xk}. ",
        "conclusions": "We have analyzed a simple model of quintessential inflation in the RSII braneworld context with a  modified exponential potential. One of the attractive features of the model is that it can lead to transient acceleration at late times. This is particularly welcome in string theoretical formulations in order to avoid the difficulties arising in the $S$-matrix construction at the asymptotic future in a de-Sitter space~\\cite{Hellerman:2001yi,Fischler:2001yj,Witten:2001kn}. Assuming that the Universe was reheated via the instant preheating mechanism, we have shown that the evolution of the scalar field from inflation till the present epoch is consistent with the observations in a wide region of the parameter space. Requiring that the model meets various cosmological constraints at the different stages of the evolution, we were able to constrain tightly its parameters, as summarized in Eqs.~(\\ref{eq:boundtrans1})-(\\ref{eq:boundM5f}).  In view of the very constrained bounds we obtained from the inflationary period, it is useful to consider how we could circumvent them in a simple and natural way. For instance, theoretical predictions for the inflationary observables may be modified by the presence of fields that are heavier than the Hubble rate during inflation. In this case, the coupling of the inflaton field to such heavy fields introduce corrections~\\cite{Bartolo:2007hx} which can be larger than the second-order contributions in the slow-roll parameters. Another way to change predictions for these observables is to consider the more general framework of Gauss-Bonnet gravity. In the presence of the Gauss-Bonnet term, the value of the spectral index is determined by the Gauss-Bonnet coupling parameter and the tension of the brane and is independent of the slope of the potential, thereby bringing the scenario in closer agreement with the most recent observations~\\cite{Lidsey:2003sj,Tsujikawa:2004dm}."
    },
    "0801/0801.4046_arXiv.txt": {
        "abstract": "Observations  of local X-ray absorbers, high-velocity clouds, and distant quasar absorption line systems  suggest    that a significant fraction  of baryons may reside in multi-phase, low-density, extended, $\\sim 100$  kpc,  gaseous halos around normal galaxies.   We present a pair of high-resolution SPH (smoothed particle hydrodynamics) simulations that explore the nature of cool gas infall into galaxies, and the physical conditions necessary to support the type of gaseous halos that seem to be required by observations.  The two simulations are identical other than their initial gas density distributions: one is initialized with a {\\em standard} hot gas halo that traces the cuspy profile of the dark matter, and the other is initialized with a {\\em cored} hot halo with a high central entropy, as might be expected in models with early pre-heating feedback.  Galaxy formation proceeds in dramatically different fashions in these two cases.  While the standard cuspy halo cools rapidly, primarily from the central region, the  cored halo is quasi-stable for $\\sim 4$ Gyr and eventually cools via the fragmentation and infall of clouds from $\\sim 100$ kpc distances.   After 10 Gyr of cooling, the standard halo's X-ray luminosity  is $\\sim 100$ times current limits and the resultant disk galaxy  is twice as massive as the Milky Way.  In contrast, the cored halo has an X-ray luminosity that is in line with observations, an extended cloud population reminiscent of the high-velocity cloud population of the Milky Way, and  a disk galaxy with half the mass and $\\sim 50\\%$ more specific angular momentum than the disk formed in the low-entropy simulation. These results suggest that  the distribution and character of halo gas provides an important testing ground for galaxy formation models and may be used to constrain the physics of galaxy formation. ",
        "introduction": "It is well known that in the absence of feedback the  majority  of  baryons in galaxy-size    dark matter halos  ($M_{\\rm v} \\sim 10^{12}$  M$_\\odot$)  should have  cooled into halo  centers over a     Hubble  time (e.g. White \\& Rees 1978;  Katz et al. 1992; Benson et al. 2003). In contrast, only $\\sim  20 \\%$ of  the  associated baryons  in Milky-Way size halos are observed to be in a cold, collapsed form (Maller \\& Bullock 2004; Mo et al. 2005; Fukugita \\& Peebles 2006; Nicastro et al. 2007). An understanding of the feedback processes that  act  to  solve this   galaxy  ``overcooling'' problem is a major goal of galaxy formation today. It  is not known  if the  unaccounted baryons exist  primarily as plasmas within normal  galaxy halos  (Maller \\& Bullock 2004; Fukugita \\& Peebles 2006; Sommer-Larsen 2006) or if they  have   been largely expelled as a result of energetic blow-out (see, e.g., Oppenheimer \\& Dav{\\'e}   2006, for a discussion).   Observations  of local X-ray absorbers (Williams et al. 2005; Fang et al. 2006), high-velocity clouds (Collins et al. 2005; Thom et al. 2006; Peek et al. 2007), and distant quasar absorption line systems (Tumlinson \\& Fang 2005; Kacprzak et al. 2007; Tinker \\& Chen 2007), suggest    that a significant fraction  of the missing halo baryons may reside in multi-phase, extended, $\\sim 100$  kpc,  gaseous halos of normal galaxies.   However, any hot gas around disk galaxies must be relatively low density in order to evade X-ray emission bounds ($S_x < 10^{-14}$ erg cm$^{-2}$ s$^{-1}$ arcmin$^{-2}$; Rasmussen et al. 2006, Li et al. 2007). These results, together with the fairly high covering factors in cool clouds implied by absorption line studies  ($\\sim 50 \\%$; Kacprzak et al. 2007), are suggestive of a model where normal galaxies are surrounded by extended, low-density hot ($\\sim 10^6$ K) halos that are filled with fragmented, pressure supported cool ($\\sim 10^4$ K) clouds (Maller \\& Bullock 2004).   Independently, models aimed  at explaining the optical properties of galaxies  have relied    increasingly  on the   idea  that  extended, quasi-stable hot gas halos develop around massive galaxies (Kere{\\v s} et~al. 2005; Bower et al. 2006, Croton et al. 2006). It is suggested that these hot halos may be quite susceptible to feedback mechanisms, which could stabilize the systems to cooling  and help explain  the observed  bimodality in galaxy properties (Dekel \\& Birnboim 2006).   Observational probes of the gaseous halos of galaxies provide a potential means of testing these ideas.  Entropy injected from feedback mechanisms will alter the density distribution of halo gas and affect associated cooling rates (and thus X-ray emission) and the distribution of cooling clouds fragmenting within the hot halos. Similarly, early feedback or {\\em pre-heating} before the halo collapses can affect halo gas profiles in a related manner, with positive consequences for galaxy properties at $z=0$ (Mo \\& Mao 2002; Oh \\& Benson; Lu \\& Mo 2007).   \\begin{figure*} \\center \\includegraphics[height=50mm]{kaufmann1.ps} \\includegraphics[height=50mm]{kaufmann2.ps} \\caption{Renderings of the total gas density after 10 Gyr of cooling in our simulations. {\\bf Left:} results for our \"standard\" halo, initialized with a low-entropy cuspy hot halo. {\\bf Right:} results for the \"cored\" halo, initialized with a high-entropy, low-density hot halo of the same mass. In the first case, gas is concentrated around the galaxy.  In the second case, the dense gas is more extended and distinct clouds are visible out to $\\sim100$ kpc. The boxes are $400$ kpc on a side. \\label{pic} } \\end{figure*} ",
        "conclusions": "Testing the effects of different initial conditions can teach us about the types of gaseous halos that are required to match the observed distributions of halo gas around galaxies.   These preliminary simulations  show that a cored initial gas density profile with a high initial entropy (as expected in pre-heating scenarios) but without any other feedback can produce disk masses, cool gas distributions,  and X-ray emission signals that are in better agreement with observations than a more standard cuspy halo case with low central entropy.  These results suggest that  the distribution and character of gaseous galactic halos provide a powerful tool for understanding the physics of galaxy formation."
    },
    "0801/0801.3516_arXiv.txt": {
        "abstract": "At least one massive binary system containing an energetic pulsar, PSR B1259-63/SS2883, has been recently detected in the TeV $\\gamma$-rays by the HESS telescopes. These $\\gamma$-rays are likely produced by particles accelerated in the vicinity of the pulsar and/or at the pulsar wind shock, in comptonization of soft radiation from the massive star. However, the process of $\\gamma$-ray production in such systems can be quite complicated due to the anisotropy of the radiation field, complex structure of the pulsar wind  termination shock and possible absorption of produced $\\gamma$-rays which might initiate leptonic cascades. In this paper we consider in detail all these effects. We calculate the  $\\gamma$-ray light curves and spectra for different geometries of the binary system PSR B1259-63/SS2883 and compare them with the TeV $\\gamma$-ray observations. We conclude that the leptonic IC model, which takes into account the complex structure of the pulsar wind shock due to the aspherical wind of the massive star, can explain the details of the observed $\\gamma$-ray light curve. ",
        "introduction": "PSR B1259-63/SS 2883 is the best known massive binary system containing a classical radio pulsar in a highly eccentric orbit around the Be type star. The pulsar creates strong relativistic wind which prevents accretion from the wind of the massive star. So then, it is expected that high energy processes, which are characteristic for the isolated radio ($\\gamma$-ray) pulsars, can also occur in the inner pulsar magnetosphere of PSR B1259-63. Due to the short rotational period, equal to $47$ ms, and strong surface magnetic field, the pulsar produces a strong relativistic wind which, in contrary to isolated pulsars, interacts with the wind from the massive star. This interaction process is especially important at the periastron passage when the pulsar is immersed in the inhomogeneous massive star wind. As a result of this interaction process, a relativistic shock in the pulsar wind appears. Particles can be accelerated on the shock to $\\sim$ TeV energies. Therefore, it has been suspected, that the binary system PSR 1259-63/SS 2883 might produce observable high energy $\\gamma$-ray radiation. This expectation has been confirmed during the last periastron passage of the pulsar in 2004, when the HESS detectors discovered the TeV $\\gamma$-ray emission from this binary system (Aharonian et al.~2005).  The production of X-ray and $\\gamma$-ray  radiation in PSR B1259-63 system in the model of interaction of the pulsar and stellar winds was considered for the first time  by Grove et al.~(1997), and Tavani et al.~(1997). In those works $e^\\pm$ pairs, escaping from the pulsar magnetosphere, are accelerated additionally at the pulsar wind termination shock and interact with the magnetic field and thermal radiation of the massive star. Moreover, the conditions for propagation and cooling of $e^+e^-$ pairs have been considered as a function of the orbital phase. The produced X-ray light curves were in good agreement with the observed X-ray flux variability. The photon spectra were fitted under the assumption that the massive star equatorial wind is inclined to the orbital plain with the inclination angle $i > 25^o$. According to these models, the Lorentz factors of $e^\\pm$ pairs in the pulsar wind should be of the order of $\\gamma_e \\sim 10^6$, while the wind magnetization parameter, which determines the magnetic field strength at the shock, should be $\\sigma \\sim 10^{-1} - 10^{-2}$  (Tavani et al.~1996).  A similar model for this binary system has been also considered by Kirk, Ball \\& Skjaeraasen~(1999). These authors argued for the first time that also the TeV $\\gamma$-rays should be produced in PSR 1259-63/SS 2883 binary system in the IC scattering of soft radiation from the massive star.  Models with domination of the adiabatic and radiative energy losses were considered. However, they predicted the maximum of the TeV emission around the periastron passage which seems to be inconsistent with the observations. However, very recently Khangulyan et al.~(2007) argued that under special conditions such a general leptonic model can explain the TeV $\\gamma$-ray light curve observed from this system. Also the hadronic model for production of high energy $\\gamma$-rays has been considered by Kawachi et al.~(2004) who postulate acceleration of hadrons at the termination shock in the pulsar wind. In this model the hadrons can pass to the dense equatorial wind of the massive star and produce pions decaying in $\\gamma$-rays. The $\\gamma$-ray light curve expected in this model shows two maxima corresponding to the situation when the pulsar crosses of the equatorial wind of the massive star and the local minimum at the periastron passage. Such a $\\gamma$-ray light curve describes much better the HESS observations. The hadronic model has been also recently discussed by Neronov \\& Chernyakova~(2007).  The above mentioned leptonic models have not taken into account the influence of the effects connected with an aspherical wind of the Be star. Such stars are characterized by dense but slow equatorial winds and low density but fast poloidal winds. The pulsar orbiting around the massive star has to pass trough these different winds. When it is immersed in  the equatorial wind, the shock appears relatively close to the neutron star surface. In contrast, in the polar wind, the shock extends to relatively close to the massive star surface. In this last case, particles accelerated at the shock propagate through a much stronger radiation field and the ICS process should occur more efficiently. Moreover, the $\\gamma$-rays produced close to the stellar surface can also be partially absorbed, initiating in this way the IC $e^{\\pm}$ pair cascade. The importance of the $\\gamma$-ray production in the cascades occurring close to the surface of the massive star has been considered in the past, but only in the case of the isotropic winds (e.g. Sierpowska \\& Bednarek 2005a). In this paper, we consider those leptonic processes in a much more complicated geometry, i.e. by applying the aspherical wind properties. We calculate the $\\gamma$-ray spectra and light curves expected in such a more complicated and realistic model. The results of our calculations are compared with the recent observations of this binary system by the HESS telescopes in the TeV $\\gamma$-ray energies. The preliminary results has been published in Sierpowska \\& Bednarek~(2005b) and more details of the calculations are included in the PhD thesis of Sierpowska-Bartosik (Sierpowska-Bartosik~2006). ",
        "conclusions": "We considered the model for production of high energy gamma rays in the binary system consisting of a Be type massive star and the energetic pulsar applying the details of complicated geometry of the massive star wind. The model predictions are compared with recent observations of the binary system PSR B1259-63/SS2883. In this system, due to the interaction of the pulsar wind and the massive star wind, a shock wave is created. The geometry of the shock is complex and changes with the orbital phase as the pulsar passes through equatorial or polar wind of the massive star. Such structure of the massive star wind is characteristic for a Be type star. In the considered system, the equatorial wind plane is inclined to the orbital one. $\\gamma$-rays are produced by leptons which are accelerated at the shock region and after that scatter thermal photons from the massive companion. Since the soft photon field is anisotropic, we take into account the details of the geometry of the binary system. The change of the separation between the stars with the phase of the binary and the complex structure of the massive star wind have also big impact on the conditions for acceleration of leptons and their interactions with the stellar radiation. The maximum energies of leptons and their energy losses are determined by the magnetic field strength at the shock region which in turn is defined by the magnetization parameter ($\\sigma$) of the pulsar wind. In our calculations a value of $\\sigma = 10^{-2}$ and $10^{-4}$ is considered. The presented model takes into account the complicated geometry of the system and still it has to depend on many parameters which describe the position of the orbital plane and the equatorial plane of the stellar wind with respect to the observer.  In such a complicated geometrically  model, leptons can be accelerated in some cases relatively close to the stellar surface. Therefore, the possible absorption of the $\\gamma$-rays produced in the IC scattering process has to be also taken into account.  The optical depths for leptons on ICS within the pulsar wind region are clearly too low for effective production of $\\gamma$-ray photons. However, leptons, which arrive from the pulsar to the shock region, are trapped by turbulent magnetic field after the shock and can be additionally accelerated. If the shock region is relatively close to the massive star, IC scattering becomes important and high energy $\\gamma$-rays are produced effectively.  We calculate the $\\gamma$-ray spectra produced in the shock region for power law spectra of leptons in two different geometrical models. By applying the sets of parameters, I: $\\Delta\\omega_d=5^o, \\, \\theta_d=20^o, \\, i_d=40^o$ and II: $\\Delta\\omega_d=20^o, \\, \\theta_d=20^o, \\, i_d=30^o$, we reach the general consistency with the TeV $\\gamma$-ray observations. The best fit of the gamma light curve measured by the HESS telescopes (Aharonian, et al.~2005) is obtained by applying the model I with $\\sigma=10^{-4}$. However, this fit has been reached assuming that the periastron position of the pulsar is defined by the angle $\\omega_{\\rm p} = 148^o$ contrary to the present estimates, i.e.  $\\omega_{\\rm p}\\sim 138^o$. Note, however that the estimation of the angle $\\omega_{\\rm p}$ depends on massive star wind parameters such as the inclination of the disc (applied from Wex et al. 1998) and the mass loss rate, which in case of PSR B1259-63 was fixed to $5\\times 10^{-8}$ M$_\\odot$yr$^{-1}$, and which are in fact not very well known.  The calculated $\\gamma$-ray spectra are in agreement with the TeV HESS data provided that the spectral index of the spectrum of leptons is $\\alpha_i = -2.5$. The slope and also the fluxes of the photon spectra do not change with orbital phase, as was reported by the HESS collaboration.  The final light curve for model I ($\\sigma=10^{-4}, \\omega = 148^o$) is characterized by a maximum for the phases when the pulsar crosses the equatorial disc after periastron. This fit differs from the observational data for phases when pulsar approaches the equatorial disc. This can be due to the complex geometry and structure of the shock region when the pulsar wind interacts with the strong equatorial wind. The model assumes that the acceleration of leptons occurs locally at the shock. In reality, this acceleration process of leptons arriving to the specific place at the shock can extend to other parts of the shock. Such a non-local acceleration process is very difficult to consider, but it might be important and influence some details of the calculated $\\gamma$-ray light curves.  Based on the reported calculations, we conclude that the general shapes of the TeV and GeV $\\gamma$-ray light curves close to the periastron should be quite similar. However, farther out from the periastron, the GeV emission should continue with larger flux. This feature can be tested by the future observations in the GeV energy range by the AGILE and GLAST telescopes."
    },
    "0801/0801.1039_arXiv.txt": {
        "abstract": "In this paper we review the current predictions of numerical simulations for the origin and observability of the warm hot intergalactic medium (WHIM), the diffuse gas that contains up to 50 per cent of the baryons at $z\\sim 0$. During structure formation, gravitational accretion shocks emerging from collapsing regions gradually heat the intergalactic medium (IGM) to temperatures in the range $T\\sim 10^5 - 10^7$ K. The WHIM is predicted to radiate most of its energy in the ultraviolet (UV) and X-ray bands and to contribute a significant fraction of the soft X-ray background emission.  While \\ion{O}{vi} and \\ion{C}{iv} absorption systems arising in the cooler fraction of the WHIM with $T\\sim 10^5 - 10^{5.5}$ K are seen in \\fuse\\ and \\hst\\ observations, models agree that current X-ray telescopes such as \\chandra\\ and \\xmm\\ do not have enough sensitivity to detect the hotter WHIM. However, future missions such as \\constellation\\ and \\xeus\\ might be able to detect both emission lines and absorption systems from highly ionised atoms such as \\ion{O}{vii}, \\ion{O}{viii} and \\ion{Fe}{xvii}. ",
        "introduction": "\\label{Introduction}  The cosmic baryon abundance was first inferred by applying the theory of primordial nucleosynthesis to observations of light element abundances (e.g.\\ \\citep{olive2000}) and has been confirmed by observations of the cosmic microwave background radiation (CMB). A recent measurement from the Wilkinson Microwave Anisotropy Probe reveals that $\\Omega_{\\mathrm b} h^2 = 0.0223_{-0.0009}^{+0.0007}$, where the Hubble parameter is $h=0.73_{-0.04}^{+0.03}$ \\citep{spergel2007}. At $z> 2$ most of the baryons are thought to reside in the diffuse, photoionised intergalactic medium (IGM) with $T\\sim 10^4 - 10^5$ K, traced by the \\lya\\ forest\\footnote{Since the diffuse intergalactic medium is highly ionised, the baryon density inferred from \\lya\\ forest observations is inversely proportional to the assumed ionisation rate. Since independent constraints on the UV intensity are only accurate to a factor of few at best, we cannot prove that most of the baryons reside in the forest.} (e.g.\\ \\citealt{rauch1997}; \\citealt{weinberg1997}; \\citealt{schaye2001}). The gas traced by the \\lya\\ forest, together with the mass in galaxies, fully accounts for the cosmic baryon abundance at $z\\sim 2$ and confirms the prediction of nucleosynthesis theories and the CMB measurements.  However, when the mass of stars, interstellar gas and plasma in clusters of galaxies at redshift $z\\sim 0$ is considered, only a small fraction of the mass budget can be accounted for (e.g.\\ \\citealt{persic1992}; \\citealt{fukugita1998}; \\citealt{fukugita2004}). While the gas traced by \\lya\\ forest may account for about a third of the low-$z$ baryons (e.g.\\ \\citealt{danforth2007}), the majority of the baryons remain invisible. Unless the baryon budget at high redshift has been overestimated by two different measurements, a large amount of the low redshift baryons must be ``missing''.  One of the basic predictions of cosmological, gas-dynamical simulations is the distribution of baryons in the universe. How many baryons are locked in stars? How many reside in clusters and filaments? What is the state of the diffuse gas? Simulations of structure formation predict that the missing baryons problem is measure of our technological limitations rather than a real problem. In other words, the baryons are out there, but we cannot see them because they reside in gas that is too dilute to be detected in emission and too hot to be traced by \\lya\\ absorption.  The idea of gravitational heating of the intergalactic gas was first suggested by \\citet{sunyaev1972} and subsequently developed (e.g.\\ \\citealt{Nath2001}; \\citealt{furlanettoloeb2004}; \\citealt{rasera2006}). While the largest structures in the universe form, the IGM is heated by gravitational shocks that efficiently propagate from the collapsing regions to the surrounding medium.  Simulations predict that gas compressed and heated by shocks can reach temperatures of $10^8$ K in rich clusters of galaxies, while filaments and mildly overdense regions are heated to temperatures in the range $10^5$ to $10^7$ K (e.g.\\ \\citealt{Cen1999}; \\citealt{dave2001}; \\citealt{Cen2006}). The latter are commonly known as the Warm-Hot Intergalactic Medium (WHIM), which represents about half the total baryonic mass in the universe at $z\\sim 0$.  At $T\\sim 10^5 - 10^7$ K, the IGM is collisionally ionised and becomes transparent to \\lya\\ radiation. As a consequence, the \\ion{H}{i} \\lya\\ forest does not trace the bulk of the gas mass at low redshift. Although the shock-heated IGM emits radiation in the UV and in the soft X-ray bands, the total energy radiated away by mildly overdense gas is orders of magnitude too small to be detected by current instruments. Similarly, the absorption from such a gas along the line of sight to a bright X-ray source is too weak to be resolved by current spectrographs (\\citealt{richter2008} - Chapter 3, this volume).  In this Paper, we review the predictions of numerical simulations for the properties and the observability of the WHIM.  In Sect.~\\ref{origin} we briefly mention the numerical techniques used to simulate the IGM and we discuss the mechanisms that produce the WHIM, namely heating by gravitational accretion shocks and other, non-gravitational, heating processes. A more detailed description of the numerical techniques themselves can be found in \\citet{dolag2008} - Chapter 12, this volume.  Sect.~\\ref{background} describes the contribution of the WHIM emission to the soft X-ray background. The current predictions for the detectability of the WHIM in emission and absorption are discussed in Sect.~\\ref{emission} and in Sect.~\\ref{absorption} respectively. Sect.~\\ref{sz} briefly reports on the effect of the WHIM on the CMB radiation and Sect.~\\ref{bfield} discusses the magnetisation of the IGM. We draw our conclusions in Sect.~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  The search for the ``missing baryons'' has opened the way to our understanding of the most elusive baryonic component in the universe: the hot diffuse gas that traces the cosmic web.  Given the paucity of observational findings, most of what we know about the WHIM is based on numerical simulations.  These predict that at $z<1$ up to 50 per cent of the IGM has been heated to temperatures in the range $T \\sim 10^5 - 10^7$ K. While most of the heating is provided by gravitational accretion shocks propagating out of regions undergoing gravitational collapse, non-gravitational processes such as galactic winds and AGN feedback may play an important role in high density regions, mostly by preventing cooling and regulating the intensity of the X-ray background.  The observability of the WHIM by future facilities has received a lot of attention and a wealth of predictions have been made based on numerical results.  The detection of the WHIM seems to be extremely challenging, if not impossible, with current instruments and a new generation of X-ray telescopes with higher sensitivity and spectral resolution, such as \\constellation\\ and \\xeus, is needed to study the WHIM in absorption. In addition to high sensitivity and spectral resolution, high spatial resolution on a large field of view, as proposed for the Explorer of Diffuse emission and Gamma-ray burst Explosions (\\edge), would be ideal for mapping the WHIM emission on the sky.  Although most theoretical predictions claim optimistically that we will be able to detect the WHIM in absorption with \\constellation\\ and \\xeus, and maybe in emission with some luck, there is a long way to go before observations can confirm the simulation predictions, or show us a different picture altogether of what and where the WHIM is."
    },
    "0801/0801.2916_arXiv.txt": {
        "abstract": "{This paper presents a method which can be used to calculate models of the global solar corona from observational data.} % {We present an optimization method for computing nonlinear magnetohydrostatic equilibria in spherical geometry with the aim to obtain self-consistent solutions for the coronal magnetic field, the coronal plasma density and plasma pressure using observational data as input.} { Our code for the self-consistent computation of the coronal magnetic fields and the coronal plasma solves the non-force-free magnetohydrostatic equilibria using an optimization method. Previous versions of the code have been used to compute non-linear force-free coronal magnetic fields from photospheric measurements in Cartesian and spherical geometry, and magnetostatic-equilibria in Cartesian geometry. We test our code with the help of a known analytic 3D equilibrium solution of the magnetohydrostatic equations. The detailed comparison between the numerical calculations and the exact equilibrium solutions is made by using magnetic field line plots, plots of density and pressure and some of the usual quantitative numerical comparison measures.} { {We find that the method reconstructs the equilibrium accurately, with residual forces of the order of the discretisation error of the analytic solution. The correlation with the reference solution is better than $99.9\\%$ and the magnetic energy is computed accurately  with an error of $< 0.1 \\%$. } } { We applied the method so far to an analytic test case. We are planning to use this method with real observational data as input as soon as possible. } ",
        "introduction": "In the recent past, numerical methods based on optimization principles have been used for a number of problems associated with the calculation of solar MHD equilibria. \\citet{wheatland:etal00} were the first to suggest the use of an optimization method to calculate nonlinear force-free fields in the corona from photospheric measurements. Since then the optimization method has been extended in various ways, for example by improving certain aspects of the original method for force-free fields \\citep[e.g.][]{wiegelmann04}, by introducing additional features such as plasma pressure into the method \\citep[e.g.][]{wiegelmann:etal06b} or by reformulating the method in other geometries  \\citep[e.g.][]{wiegelmann07}.  Optimization methods have the advantage of being conceptually straightforward and are reasonably easy to implement \\citep[see e.g.]{wheatland:etal00,wiegelmann:etal03,inhester:etal06}. At the moment they also seem to be very competitive in terms of computational efficiency \\citep[e.g][]{schrijver:etal06,metcalf:etal07}. A slight disadvantage compared to, for example, the Grad-Rubin method \\citep[e.g.][]{amari:etal97,wheatland06,inhester:etal06,amari:etal06} is that they still lack the same degree of rigorous mathematical basis existing for other methods.  In the present paper we describe a further extension of the optimization method to calculate magnetohydrostatic (MHS) equilibria (including pressure and gravity) in spherical geometry. This is important for calculating global models of the corona including information going beyond just the structure of the magnetic field. In section \\ref{equations} we describe the basic equations of the optimization method for problem in hand. We then give a brief description of the analytical 3D MHS equilibria that we use to test the numerical code in section \\ref{3DMHS} and present the test results in section \\ref{results}. Our conclusions are presented in section \\ref{conclusions}.   \\section[]{Basic equations} \\label{equations} The magnetohydrostatic (MHS) equations are given by \\begin{eqnarray} (\\nabla \\times {\\bf B}) \\times {\\bf B} -\\mu_0\\nabla p -\\mu_0 \\, \\rho \\, \\nabla \\Psi &=& {\\mathbf{0}}  \\label{1} \\\\ \\nabla \\cdot {\\bf B} &=& 0,  \\label{2} \\end{eqnarray} where ${\\bf B}$ is the magnetic field, $p$ the plasma pressure, $\\rho$ the mass density and $\\Psi=-\\frac{G M_s}{r}$ the gravitational potential with the gravitational constant $G$, the solar mass $M_s$ and the distance from the sun's center $r$. We do not assume an equation of state for the coronal plasma, but leave $p$ and $\\rho$ to be independent quantities. To find a magnetic field $\\mathbf{B}$, plasma pressure $p$ and plasma density $\\rho$ satisfying Eqs. (\\ref{1}) and (\\ref{2}), we follow the spirit of the previous optimization methods \\citep[e.g.][]{wheatland:etal00,wiegelmann:etal03a,wiegelmann04,wiegelmann:etal06b} and define the functional \\begin{eqnarray} L({\\bf B},p,\\rho) & = & \\int \\Bigg[\\frac{\\left|(\\nabla \\times {\\bf B}) \\times {\\bf B} -\\mu_0 \\, \\nabla p -\\mu_0 \\, \\rho \\,\\nabla \\Psi    \\right|^2}{B^2} \\nonumber \\\\ && \\; + \\; |\\nabla \\cdot {\\bf B}|^2\\Bigg] \\; \\, r^2 \\, \\sin \\theta \\, dr \\, d \\theta \\, d \\phi. \\label{defL_Bprho} \\end{eqnarray}  It is obvious that Eqs. (\\ref{1}) and (\\ref{2}) are satisfied if $L=0$. Here ${\\bf B}$ is a vector field, but not necessarily a solenoidal magnetic field during the iteration. The numerical method is based on an iterative scheme to minimize the functional $L$. To simplify the mathematical derivation we define the quantities \\begin{eqnarray} {\\bf \\Omega_a} &=& \\; B^{-2} \\;\\left[(\\nabla \\times {\\bf B}) \\times {\\bf B} -\\mu_0  \\, \\nabla p -\\mu_0 \\,\\rho\\, \\nabla \\Psi \\right] \\\\ {\\bf \\Omega_b} &=& B^{-2} \\;\\left[(\\nabla \\cdot {\\bf B}) \\; {\\bf B} \\right], \\label{defomega} \\end{eqnarray} and rewrite $L$ as \\begin{equation} L=\\int_{V} \\;  \\left[ B^2 \\Omega_a^2+  B^2 \\Omega_b^2 \\; r^2 \\right] \\, \\sin \\theta \\, dr \\, d \\theta \\, d \\phi. \\end{equation} Taking the derivative of $L$  with respect to an iteration parameter $t$, where ${\\bf B}, \\; p, \\; \\rho$ are assumed to depend upon $t$, we obtain \\begin{eqnarray} \\frac{1}{2} \\; \\frac{d L}{d t} &=&-\\int_{V} \\frac{\\partial {\\bf B}}{\\partial t} \\cdot {\\bf F} \\; dV +\\int_{V} \\frac{\\partial p}{\\partial t} \\; \\mu_0 \\nabla \\cdot {\\bf \\Omega_a} \\; dV \\nonumber \\\\ &&-\\int_{V} \\frac{\\partial \\rho}{\\partial t} \\; \\mu_0 {\\bf \\Omega_a} \\cdot \\nabla \\Psi \\; dV -\\int_{S} \\frac{\\partial {\\bf B}}{\\partial t} \\cdot {\\bf G} \\; dS \\nonumber \\\\ && -\\int_{S} \\frac{\\partial p}{\\partial t} \\; \\mu_0 {\\bf \\Omega_a} \\cdot {\\bf dS}, \\end{eqnarray} where \\begin{eqnarray} {\\bf F} & =& \\nabla \\times ({\\bf \\Omega_a} \\times {\\bf B} ) - {\\bf \\Omega_a} \\times (\\nabla \\times {\\bf B}) \\nonumber \\\\ && + \\nabla({\\bf \\Omega_b} \\cdot {\\bf B})-  {\\bf \\Omega_b}(\\nabla \\cdot {\\bf B}) +(\\Omega_a^2 + \\Omega_b^2) \\, {\\bf B} \\end{eqnarray} and \\begin{eqnarray} {\\bf G} & = & {\\bf \\hat n} \\times ({\\bf \\Omega_a} \\times {\\bf B} ) -{\\bf \\hat n} (\\bf \\Omega_b \\cdot \\bf B). \\end{eqnarray} Assuming that \\begin{eqnarray} \\frac{\\partial {\\bf B}}{\\partial t} &=& \\mu \\; {\\bf F} \\label{Beq}\\\\ \\frac{\\partial p}{\\partial t} &=& -\\nu \\; \\mu_0 \\nabla \\cdot {\\bf \\Omega_a} \\label{peq} \\\\ \\frac{\\partial \\rho}{\\partial t} &=& \\xi \\; \\mu_0 \\; {\\bf \\Omega_a} \\cdot \\nabla \\Psi =\\xi \\; \\mu_0 \\frac{G M_s}{r^2} \\; {\\bf \\Omega_a} \\cdot {\\bf e_r} \\label{rhoeq} \\end{eqnarray} with positive definite parameters $\\mu$, $\\nu$ and $\\xi$ and that the boundary integrals vanish, one can easily see that $L$ is monotonically decreasing with $t$ (note that this does not necessarily imply that $L$ tends to zero).  Discretized versions of Eqs. (\\ref{Beq}) to (\\ref{rhoeq}), together with appropriate boundary conditions, form the basis for the numerical scheme. The boundary conditions have to be consistent with the assumption that the boundary integrals vanish, for example, by keeping the magnetic field, pressure and density fixed on the boundaries during the iteration. For testing the method in this paper we shall take these boundary conditions from the known exact solution. For practical applications these would have to come from observational data. We remark that due to the introduction of additional forces the constraints on the consistency of the boundary conditions for the magnetic field are somewhat different from the force-free case. However, the general theory of magnetohydrostatic equilibria requires for example that the pressure has the same value at both foot points of a closed field line under the general conditions assumed in the present paper. This is similar to the Cartesian case discussed by \\citet{wiegelmann:etal06b}, where the pressure equation was forward integrated along field lines using an upwind method from one foot point to the other (in the test case  described later we make use of the property that the pressure is known to be consistent).  In the numerical implementation based on this method one has to choose the product of the time-step $\\Delta t$ with the three numerical parameters $\\mu$, $\\nu$ and $\\xi$. Usually these products have to be small enough to achieve convergence and this is ensured by an adaptive time-step control. In this paper we have chosen the three parameters (multiplied by $\\Delta t$) to have the same values on all grid points. Previous experience with applying a similar method to force-free magnetic fields in spherical geometry \\citep{wiegelmann07} showed that choosing the same values for the entire box can lead to  long computing times in the polar regions due to the distortion of the numerical grid in spherical polar coordinates  towards the poles (note that the poles $\\theta=0$ and $\\theta=\\pi$ are excluded from the computational domain). This could in principle be compensated by allowing for a spatial variation of the iteration parameters. ",
        "conclusions": "\\label{conclusions} \\begin{figure*} \\includegraphics[angle=-90, width=\\textwidth]{8244f4.eps} \\caption{Outlook: How can the tool become applied to data? The basic idea is to compute first a nonlinear force-free field from observed vector magnetograms and then model the plasma along the loops. Our newly developed magneto-hydrostatic code uses the resulting magnetic field and plasma configuration as input for the computation of a self-consistent equilibrium.} \\label{fig4} \\end{figure*} We have extended the optimization method originally proposed for the reconstruction of force-free magnetic fields \\citep{wheatland:etal00} to global magnetohydrostatic equilibria including the pressure force and the gravitational force in spherical geometry. The proposed generalization of the optimization method leads to two additional equations for the pressure and the density that have to be solved simultaneously to the magnetic field equation. Boundary conditions for the magnetic field, the pressure and the density are necessary to complete the problem.  We have implemented a numerical code based on the proposed method and have tested the code using a known three-dimensional magnetohydrostatic equilibrium \\citep{neukirch95}. The numerical calculation is started from a potential field with the same radial magnetic field component as the analytical equilibrium on the boundary. The initial pressure and density distribution are a spherically symmetric stratified atmosphere in hydrostatic balance. Both visual inspection of the results as well as a quantitative analysis using various diagnostic measures indicate that the method works well and converges to the analytic equilibrium. For the presented tests we used a low spatial resolution and got a relatively long computing time (about 200 000 iteration steps) until convergence. In experiments with the force-free version of our spherical code \\citep[see][]{wiegelmann07} we found that the computing time scales with $N^{5.4}$ regarding the number of grid points $N$ in one spatial direction. This is somewhat slower as the theoretical estimate of $N^5$ for a cartesian optimization code obtained by \\citet{wheatland:etal00}. The spherical magnetohydrostatic code is significantly slower than the cartesian force-free code for two reasons. \\begin{enumerate} \\item  The convergence of the numerical grid towards the poles requires a sufficiently small time-step. \\item The plasma $\\beta$ might vary strongly in the entire region and in particular low-$\\beta$-regions require very small time-steps to compute the magnetic field and plasma simultaneously, because small changes in the Lorentz force can result in considerably large changes in the low $\\beta$ plasma. \\end{enumerate} Point 1.) can be addressed by using a more sophisticated numerical grid, e.g. the Yin-Yang grid developed by \\citet{kageyama:etal04}, which has been applied in geophysics \\citep[see e.g.][]{yoshida:etal04}. This overset grid contains two complementary grids which lead to an almost uniformly spaced spherical grid. An additional advantage of the Yin-Yang grid is that it is suitable for massive parallelization. The Yin-Yang grid has been applied for geophysical simulations on the Earth simulator super-computer in Yokohama. To speed up the 2. point one might compute first the magnetic field alone as a nonlinear force-free field (which is a reasonable approximation in low-$\\beta$-regions) and switch on the self-consistent plasma iteration only after for a fine-tuning. A multi-scale approach, as recently implemented to speed up our force-free cartesian optimization code \\citep[see][for details]{metcalf:etal07} is also an option worth trying for the spherical implementation of our force-free and magnetohydrostatic codes. We are confident, that the above mentioned potential for improvements together with a massive parallelization will allow us to apply our newly developed method to real data with a reasonable grid resolution.  In figure \\ref{fig4} we outline a scheme on how the code might be applied to data. {\\tw As boundary conditions on pressure and density are not directly measured, we propose the following approach.} The basic idea is to compute first a force-free magnetic field and then model the plasma along the magnetic loops, e.g., by the use of scaling laws and optionally with the help of a tomography code. Such an approach has been used by \\citet{schrijver:etal04} by using a global potential magnetic field and specifying free scaling law parameters (e.g., heating function) by comparing artificial plasma images (created by line-of-sight integration from the model plasma) with X-ray and EUV observations. We propose to generalize this approach by using a nonlinear force-free magnetic field model and compare the model plasma with observations from two viewpoints as provided by the STEREO-mission. Optional STEREO-images can be used additionally to approximate the coronal density distribution by a tomographic inversion. As a consequence of this step (modelling the plasma along a magnetic loop) the plasma pressure is consistent along the loops. Different values for the pressure on different field-lines will violate the force-free condition for the magnetic field, however, and the configuration is not exactly in a magnetohydrostatic equilibrium. Finite pressure gradients have to be compensated by a Lorentz force. This computation can be done with the help of the program described in this paper. We propose to use the force-free magnetic field configuration and the model plasma as initial state for our newly developed magnetohydrostatic code to compute a self-consistent MHS-equilibrium. For a low $\\beta$ plasma the back-reaction of the plasma onto the magnetic field will be small, for higher values of $\\beta$ (as found e.g., in helmet streamers) the magnetic field might change significantly. As a result of this approach one has reconstructed the 3D coronal magnetic field and plasma configuration self-consistently within the magnetohydrostatic approach and the model is consistent with measured photospheric vector magnetograms and observed coronal images from different viewpoints as well."
    },
    "0801/0801.2257_arXiv.txt": {
        "abstract": "We analyze the mass distribution of cores formed in an isothermal, magnetized, turbulent, and self-gravitating nearly critical molecular cloud model. Cores are identified at two density threshold levels. Our main results are that the presence of self-gravity modifies the slopes of the core mass function (CMF) at the high mass end. At low thresholds, the slope is shallower than the one predicted by pure turbulent fragmentation. The shallowness of the slope is due to the effects of core coalescence and gas accretion. Most importantly, the slope of the CMF at the high mass end steepens when cores are selected at higher density thresholds, or alternatively, if the CMF is fitted with a lognormal function, the width of the log-normal distribution decreases with increasing threshold. This is due to the fact that gravity plays a more important role in denser structures selected at higher density threshold and leads to the conclusion that the role of gravity is essential in generating a CMF that bears more resemblance with the IMF when cores are selected with an increasing density threshold in the observations. ",
        "introduction": "It is now widely accepted that stars form in dense and gravitationally bound molecular cloud (MC) cores. Thus, understanding the properties of the core mass function (CMF) is the cornerstone of any successful theory of the stellar initial  mass function (IMF). The question whether there is a 1-to-1 mapping between the CMF and the IMF remains an open one. The CMF is usually represented by $dN/dM=M^{\\alpha}$, where $\\alpha$ takes different values in different mass ranges (for the IMF, $\\alpha=-2.35$ for $M \\gtrsim 0.5$ M$_{\\odot}$, Salpeter 1955). The determination of the CMF of cores of several MC regions have been obtained using a variety of wavelengths and techniques such as mm spectroscopy of various molecular lines, submm/mm continuum emission from dust, and stellar near infrared absorption by dust. Usually, no prior assumptions are made on the type of structures identified in these studies. They can be a mixture of prestellar cores, unbound objects, and eventually protostellar envelopes. For nearby MCs, where it is possible to identify dense cores (with average number densities $n \\gtrsim 10^{4}$ cm$^{-3}$; hereafter Cores), the CMF is usually described by a two component power law above and below a turnover point (with exponents $\\alpha_{h}$ and $\\alpha_{l}$, respectively). Starting with the 1.2 mm study of the Ophiuchus starless dense cores by Motte et al. (1998), a number of submm/mm dust continuum observations have found $\\alpha_{h}$ and $\\alpha_{l}$ to be in the range [-2.1,-3] and [-1.3,-1.5], respectively (Motte et 2001; Testi \\& Sargent 2001; Johnstone et al. 2006; Beuther \\& Schilke 2004; Stanke et al. 2006; Reid \\& Wilson 2006; Johnstone \\& Bally 2006; Kirk et al. 2006; Enoch et al. 2006). These estimates bracket the values of the exponents of the IMF (Scalo 1998; Kroupa 2002; Chabrier 2003). Similar slopes are found using high density molecular tracers such as H$^{13}$CO$^{+}$ (Onishi et al. 2002; Ikeda et al. 2007) and the near infrared absorption technique (Alves et al. 2007). On the other hand, submm/mm observations of more distant clouds in which it is only possible to identify lower density structures ($n \\sim 10^{2}-10^{3}$ cm$^{-3}$; hereafter Clumps), exhibit values of $\\alpha_{h}$ in the range of [-1.4,-1.8] if the whole CMF is fitted with a single power law (Coppin et al. 2000; Kerton et al. 2001; Tothill et al. 2002; Mookerjea et al. 2004; Massi et al. 2007; Li et al. 2007, Mu\\~{n}oz et al. 2007) or slightly steeper ($\\alpha_{h} \\sim -2$) if it can be fitted by a two component power law (Reid \\& Wilson 2005). Similar slopes are obtained from molecular line observations using low/intermediate density molecular line tracers (Stutzki \\& G\\\"{u}tsen 1990; Lada et al. 1991; Blitz 1993; Brand \\& Wouterloot 1995; Kramer et al. 1998; Heithausen et al. 1998; Kawamura et al. 1998; Wilson et al. 2003). Values of $\\alpha_{h} \\sim [-1.6,-1.7]$ resemble the slopes of the stellar clusters mass function (Elmegreen 2000), which hints to the fact that these objects cannot be the direct progenitors of individual stars and must undergo additional fragmentation.  The discrepancy in $\\alpha_{h}$ between the Cores and Clumps populations does not seem to be merely an effect of the spatial resolution and hence of mass resolution. Additional evidence might be that, in Motte et al. (1998), the inclusion of the larger size cores in the CMF changes the values of $\\alpha_{h}$ from $-2.5$ to $\\sim -1.5$. An explanation of the origin of the discrepancy would be that different physical processes dominate the evolution of these two populations of cores. The close resemblance between the Cores CMF and the IMF suggests that they are more likely to be gravitationally bound (which is confirmed by mm line observations, Andr\\'{e} et al. 2007) and are the direct progenitors of stars or small stellar systems, whereas Clumps are more likely to be dominated by turbulence. Note that the description of the CMF by power law functions is not the unique option. Reid el al. (2006) argued that a lognormal function is a better fit to the observed CMF in many regions and Goodwin et al. (2007) showed that the CMF of the submm cores obtained by Nutter \\& Ward-Thompson (2007) is well represented by a lognormal distribution (also Fig. 2 in Andr\\'{e} et al. 2008).  Several physical processes are believed to play a role in shaping the CMF: gravitational and thermal fragmentation (Larson 1985; Klessen 2001; Li et al. 2003), supersonic turbulence (Padoan \\& Nordlund 2002, PN02), accretion (Larson 1992; Bonnell et al. 2001; Bate \\& Bonnell 2005), core coalescence (Murray \\& Lin 1996; Dib et al. 2007a,2008), feedback (Adams \\& Fatuzzo 1996), and magnetic fields (Shu et al. 2004). A number of numerical studies of MCs that include some of these physical processes have derived a CMF that bears some resemblance with the CMF of Cores and the IMF (Gammie et al. 2003; Li et al. 2004; Ballesteros-Paredes et al. 2006; Padoan et al. 2007) or with the CMF of Clumps (Hennebelle \\& Audit 2007, Heitsch et al. 2008). The origin of the differences between the Cores and Clumps CMFs and the eventual link between them has not been investigated so far. In this study, we identify cores in a high resolution 3D numerical simulation of a magnetized, turbulent, and self-gravitating MC model at different density thresholds. In Dib et al. (2007b), we have shown, using a detailed virial analysis that the inner parts of dense structures, identified at higher density thresholds, are more gravitationally bound than the same objects identified at lower thresholds. Thus, constructing the CMF at different thresholds is a direct test of the effects of gravity on the CMF. A striking example of how the density threshold can modify the inferred masses of cores is the case of the core L1512. Its derived mass from CO observations is $\\sim 6$ M$_{\\odot}$ (Falgarone et al. 2001), whereas it was estimated from submm observations at $\\sim 1.4$ M$_{\\odot}$ (Ward-Thompson et al. 1999). For the whole CMF, Kramer et al. (1998) derived a value of $\\alpha_{h}=-1.72 \\pm 0.09$ for the Orion B-S region from $^{13}$CO(1-0) observations whereas Johnstone et al. (2006) derived a value of $\\sim -2.5$ from submm observations. Another example is the case of the NGC 7538 region, observed by Kramer et al. (1998) to have a CMF slope of $\\alpha_{h}=-1.65 \\pm 0.05$ in the $^{13}$CO(1-0) line and $\\alpha_{h}=-1.79 \\pm 0.12$ in the C$^{18}$O(1-0) line which traces higher density material. ",
        "conclusions": "\\label{conc} In this work, we have constructed the mass function of cores in a 3D numerical simulation of a magnetized, turbulent, and self-gravitating molecular cloud. Cores are identified at the density threshold levels of $10^{4}$ and $2.5 \\times 10^{4}$ cm$^{-3}$. In agreement with observations, the slope of the core mass function is found to be steeper for cores identified at the higher threshold level. The steepening of the slope with increasing threshold is due to the increasing importance of gravity in the virial balance of the cores at higher densities and a sign of the role played by gravity in the gravo-turbulent scenario of star formation. On the other hand, at lower thresholds, the slope is shallower than what is found for the purely turbulent case. This is due to the effects of mass growth that affect the cores. This mass growth is essentially due in the initial phase to the effect of core coalescence which generates more massive structures particularly for cores defined at lower thresholds since they have larger cross sections."
    },
    "0801/0801.0064_arXiv.txt": {
        "abstract": "{We present results from a kinematical study of the gas in the nucleus of a sample of three LINER galaxies, obtained from archival HST/STIS long-slit spectra. We found that, while for the elliptical galaxy NGC 5077, the observed velocity curves are consistent with gas in regular rotation around the galaxy's center, this is not the case for the two remaining objects. By modeling the surface brightness distribution and rotation curve from the emission lines in NGC 5077, we found that the observed kinematics of the circumnuclear gas can be accurately reproduced by adding to the stellar mass component a black hole mass of $M_{\\rm BH} = 6.8_{-2.8}^{+4.3}\\times 10^8 M_{\\odot}$ (uncertainties at a 1$\\sigma$ level); the radius of its sphere of influence ($R_{\\rm sph} \\sim$ 0\\farcs34) is well-resolved at the HST resolution. The BH mass estimate in NGC 5077 is in fairly good agreement with both the $M_{\\rm BH}-M_{\\rm bul}$ (with an upward scatter of $\\sim$ 0.4 dex) and $M_{\\rm BH}-\\sigma$ correlations (with an upward scatter of 0.5 dex in the Tremaine et al. form and essentially no scatter using the Ferrarese et al. form) and provides further support for the presence of a connection between the {\\sl residuals} \\rm from the $M_{\\rm BH}-\\sigma$ correlation and the bulge effective radius. This indicates the presence of a black hole's ``fundamental plane'' in the sense that a combination of at least $\\sigma$ and $R_{\\rm e}$ drives the correlations between $M_{\\rm BH}$ and host bulge properties. ",
        "introduction": "\\label{intro}  It is now clear that the presence of a supermassive black hole (SMBH) is a common, if not universal, feature in the center of galaxies. In fact, since active galactic nuclei (AGNs) are thought to be powered by mass accretion onto an SMBH, the high incidence of low-luminosity AGN activity in nearby galaxies \\citep{heckman80, maoz95, ho97a, ho97b, braatz97, barth98, barth99,nagar02} has led to the conclusion that a significant fraction of galaxies in the Local Universe must host SMBH. This conclusion is now supported by direct measurements of SMBH masses in the centers of nearby galaxies obtained with different techniques \\citep[see][for a review]{ferrarese05}. These measurements indicate that the BH mass $M_{\\rm BH}$ is related to the properties of the host galaxy, such as bulge luminosity $L_{\\rm bul}$ and mass $M_{\\rm bul}$ \\citep{kormendy95, magorrian98, marconi03}, light concentration \\citep{graham01, graham07}, and bulge velocity dispersion $\\sigma_{\\rm bul}$ \\citep{ferrarese00, gebhardt00a, tremaine02}.  The existence of any correlations between $M_{\\rm BH}$ and host bulge properties supports the idea that the growth of SMBHs and the formation of bulges are closely linked \\citep{silk98, haehnelt00}. This has profound implications for the process of galaxy formation and evolution. Moreover, SMBH mass estimates inferred via the above correlations, when more direct methods are not feasible, enter into a variety of important studies spanning from AGNs physics to the coeval formation and evolution of the host galaxy and its nuclear black hole.  These results need, however, to be further investigated by increasing the number of accurate BH mass determinations in nearby galactic nuclei to set these correlations on a stronger statistical basis. In particular, such a study has the potential of establishing the precise role of the various host galaxy's parameters in setting the resulting BH mass. To date, reliable SMBH detections have been obtained for a limited number of galaxies ($\\sim$ 30, \\citealt{ferrarese05}), with the bulk of $M_{\\rm BH}$ estimates in the range of $10^{7}-10^{9}  M_{\\odot}$. To add reliable new points to the $M_{\\rm BH}-$ host galaxy's properties planes is then a fundamental task for future developments of astronomical and physical studies.  One widely applicable and relatively simple method of detecting BHs is based on gas kinematics (e.g. \\citealt{harms94,ferrarese96, macchetto97, barth01}), through studies of emission lines from circumnuclear gas disks, provided that the gas velocity field is not significantly influenced by non gravitational motions. However, the purely gravitational kinematics of the gas can be established {\\sl a posteriori} from successful modeling of the gas velocity field under the sole influence of the stellar and black hole potential.  The Space Telescope Imaging Spectrograph (STIS) onboard $\\it {HST}$ is still the most suitable instrument for such studies as it provides a high angular resolution ($\\sim 0\\farcs1$) in the optical spectral region. In this band brighter emission lines are found with respect to the infrared, the only band accessible at high resolution with ground-based adaptive optics telescopes. The wealth of unpublished data contained in STIS archives represents an extraordinary and still unexplored resource. With the aim of finding galaxy candidates to provide a successful SMBH mass measurement, we performed a systematic search for unpublished data in the $\\it {HST}$ STIS archive.  In this paper we present the results obtained for a sample of three LINER galaxies. The galaxies were observed with STIS on {$\\it HST$} during Cycle 7 under Proposal ID 7354 and were part of a larger sample of eight selected objects. The results on NGC 3998 have already been published \\citep[][hereafter Paper I]{ngc3998}. In a forthcoming paper, we will complete the proposal targets and discuss the relation between BH mass, host galaxy properties, and nuclear activity. Table \\ref{source} lists the galaxies and their main physical properties. Basic data are from the Lyon/Meudon Extragalactic Database (HyperLeda) or NASA/IPAC Extragalactic Database (NED). Distances are calculated with ${\\it H}_0$ = 75 km s$^{-1}$ Mpc$^{-1}$ and are corrected for Local Group infall onto Virgo.  The paper is organized as follows. In Sec. \\ref{obs} we present HST/STIS data and the reduction that lead to the results described in Sec. \\ref{results}. In Sec. \\ref{fitting} we model the observed emission-line rotation curve for NGC 5077 and show that the dynamics of the circumnuclear gas can be accurately reproduced by circular motions in a thin disk when a point-like dark mass is added to the stellar potential.  Our results are discussed in Sec. \\ref{discussion}, and summarized in Sec. \\ref{summary}. ",
        "conclusions": "\\label{summary}  We have presented results from a gas kinematics study in the nucleus of three nearby LINER galaxies: IC 989, NGC 5077, and NGC 6500 using archival HST/STIS spectra. Only in the case of NGC 5077 does the nuclear velocity curves appear to be associated with gas in regular rotation. For IC 989 our results indicate an inner counter-rotating gas system, while for NGC 6500 the complex trend in the velocity curves suggests a nuclear expanding bubble.  We used our modeling code to fit the observed [N II]$\\lambda$6583 surface brightness distribution and velocity curve of NGC 5077. The dynamics of the rotating gas can be accurately reproduced by motions in a thin disk when a compact dark mass of $M_{\\rm BH} = 6.8_{-2.8}^{+4.3}\\times 10^8  M_{\\odot}$ is added to the stellar mass component. Furthermore, the black hole in NGC 5077 has a sphere of influence radius, $R_{\\rm sph} = GM_{\\rm BH}/\\sigma_{\\rm star}^{2}$, of $\\sim$ 62 pc ($\\simeq$ $0\\farcs34$), well-resolved at the HST resolution (2$R_{\\rm sph}/R_{\\rm res} \\simeq$ 6.8).  For what concerns the connections of this BH mass estimate with the properties of the host galaxy, the $M_{\\rm BH}$ value for NGC 5077 is in good agreement (within a factor of 2.3) with the $M_{\\rm BH}-M_{\\rm bul}$ correlation between BH and host bulge mass. The black hole mass predicted by the $M_{\\rm BH}-\\sigma$ correlation is a factor of 3.4 lower than our measure, adopting the relation found by \\citet{tremaine02} and  in excellent agreement using the parameterization by \\citet{ferrarese05}.  This result, in conjunction with the previous results for NGC 3998 (Paper I), strengthens the possibility of a connection between the residuals from the $M_{\\rm BH}-\\sigma$ relation and the bulge effective radius. While NGC 3998, indeed, has one of the lowest values of $R_{\\rm e}$ among galaxies with measured $M_{\\rm BH}$ and shows a negative residual, NGC 5077 has a larger effective radius and shows a small positive residual. We also recently showed that the same result was found for the Seyfert galaxy NGC 5252: a larger effective radius corresponds in this galaxy to a still larger positive residual.  Apparently, only with a combination of at least $\\sigma$ and $R_{\\rm e}$ is it possible to account for the correlations between $M_{\\rm BH}$ and other bulge properties, indicating the presence of a black hole's ``fundamental plane''. Clearly, the number of direct black-hole mass measurements must be further increased, together with precise determinations of $\\sigma$ and $R_{\\rm e}$, to test these conclusions on a stronger statistical basis. In a forthcoming paper we will present new BH mass determinations and discuss the physical implications of our results."
    },
    "0801/0801.0855_arXiv.txt": {
        "abstract": "In this Letter, we propose that a microphysical process takes a vital role in the shocked region in which the prompt emission of GRBs is emitted. The turbulent energy is included in the internal energy transferred by the kinetic energy of the shock. It dissipates through stochastic acceleration for the electrons to supply the early X-ray emission in the phase of shallow decay. We put the constraints on the time evolution of microphysical parameters. The early X-ray fluxes can be obtained by this scenario and these results are consistent with the Swift observation. ",
        "introduction": "In the Swift era \\citep{gehrels04}, the X-ray telescope (XRT) observation \\citep{burrows05} provides the complete light curves of gamma-ray bursts (GRBs) in the 0.2-10 keV band. One of the most interesting discoveries is the so-called shallow decay segment: the flux plateau of $F\\propto t^{-0.5}$ within $10^3-10^4$ second after the trigger \\citep{campana05,depasquale06}. Recent statistic analyses \\citep{nousek06,obrien06,liang07} have revealed that the phase of shallow decay in the X-ray afterglow might be a common feature of the long GRBs. The shallow decay in the early X-ray light curve is still a mystery, although theoretical explanations have been put forward from several aspects (see Zhang 2007 for a comprehensive review). Most of the models are: hydrodynamics of the shock by energy injection \\citep{granot06a,zhang06}, geometry of the jet \\citep{eichler06,toma06}, varying microphysical parameters \\citep{fan06,granot06b,panaitescu06,ioka06}, late prompt emission \\citep{ghisellini07} or up-scattered forward-shock emission \\citep{panaitescu07}.  From the point of microphysics in the hydrodynamic evolution, the nature of coupling between the electrons, protons and magnetic field is complex \\citep{chiang99}. Usually the simple way is to assume the equipartition between electrons, protons and magnetic field \\citep{panaitescu98}. However, in this Letter, we propose that the kinetic energy of the relativistic shocks in the plasma has been converted to the internal energy as three parts: (1) $\\varepsilon_B$, which is the energy of magnetic field; (2) $\\varepsilon_e$ and $\\varepsilon_p$, it means that the electrons and protons/positrons are accelerated by the shocks, normally, this is the process of first-order Fermi acceleration; (3) $\\varepsilon_t$, which presents the turbulent energy, and this energy would sustain a relatively long time (see Section 2). The last part has not been taken into account by the former research. In our novel scenario, the relativistic electrons accelerated by the first-order Fermi acceleration emit the gamma-ray by synchrotron radiation, after the decrease of the prompt emission tail which is shown as the deep decay in the early X-ray light curve, indicating that the internal shocks are abated, the follow-up turbulence and its effects might be dominated. The turbulence could transfer its energy to the electrons via second-order Fermi acceleration, which is also called as the stochastic acceleration. Therefore, in this turbulent region, due to the resonant interaction between electrons and plasma waves, the electrons buried in the magnetic field are re-accelerated by the stochastic acceleration. The emission of these electrons dominates the shallow decay phase, until the turbulent energy dissipates and the external shock sweeps the surround medium thus the deep decay appears again.  In Section 2, we review the Fokker-Planck equation and list the coefficients associated with the turbulent term. In Section 3, the turbulent parameter $\\varepsilon_t$, as same as $\\varepsilon_e$ and $\\varepsilon_B$, is introduced. Due to the turbulence, these microphysical parameters evolved with time are constrained by the process of stochastic acceleration. Finally, we select these relations to reproduce the feature of shallow decay. The discussions are given in Section 4. ",
        "conclusions": "Due to the turbulent energy and its dissipation in the shocked region, after the gamma-ray radiation, the electrons are re-accelerated by the stochastic acceleration and emit the early X-ray fluxes. Therefore, the energy injection in the central engine is not required to reproduce the emission in the shallow decay phase. Our results manifest that the temporal feature in the shallow decay phase represents the process of turbulent dissipation. The microphysical parameters, $\\varepsilon_e$, $\\varepsilon_B$ and $\\varepsilon_t$, are varied with the time in the shocked region. Given the special values of $p=2.0$ and $a=0.5$ in the adiabatic case, our calculation can represent the typical observational flux as $F\\propto t^{-0.5}$. Moreover, it is noted that the same shallow decay phase is also detected in the optical band as well \\citep{mason06}, our simple interpretation is that the synchrotron emission of X-ray band and optical band may be original from the same shocked region with turbulent energy. Another advantage of our model is that the crisis of radiative efficiency \\citep{lloyd04,ioka06,fan06,zhang07} is therefore dispelled without any additional assumption of ejection/ejecta from the central engine. From this point of view, we support the internal shock pattern of standard fireball model.  In this work, we assume $f(\\gamma)\\propto \\gamma^{-p}$ where $p=2.2$ as the typical value. The detailed calculation of the energy distribution $f(\\gamma)$ is not needed, because in our scenario the timescale of acceleration is much smaller than that of turbulent energy dissipation. Since the timescale of acceleration is also smaller than that of shock hydrodynamics, therefore, during the phase of shallow decay, the index of spectrum does not change \\citep{liang07}.  Generally, the varied values of $\\varepsilon_e$ and $\\varepsilon_B$ lead to the different hydrodynamic evolution and emissions in the entire afterglow \\citep{mao01a,mao01b}. In this Letter, the early X-ray fluxes in the shallow decay phase are estimated either in the adiabatic or radiative case through the turbulent process. The average slope of radiative case is deeper than that of adiabatic case. Compared with the XRT observation, the hydrodynamic evolution of adiabatic case could be the realized regime during the shallow decay phase. Although it is hard to obtain the values of $\\varepsilon_e$ and $\\varepsilon_B$ in our scenario, we put the constraints of $\\varepsilon_e\\ll 1$ and $\\varepsilon_B\\ll 1$. More observational samples are particularly requested for the further investigations."
    },
    "0801/0801.3856_arXiv.txt": {
        "abstract": "We present photometry and spectroscopy of the suspected cataclysmic variable (CV) Lanning 386.  We confirm that it is a CV, and observe deep eclipses, from which we determine the orbital period $P_{\\rm orb}$ to be $0.1640517 \\pm 0.0000001$ d (= 3.94 h). Photometric monitoring over two observing seasons shows a very active system with frequent outbursts of variable amplitude, up to $\\sim 2$ mag.  The spectrum in quiescence is typical of dwarf novae, but in its high state the system shows strong \\ion{He}{2} emission and a broad \\ion{C}{4} Wolf-Rayet feature. This is unusual for dwarf novae in outburst and indicates a high excitation. In its high state the system shows some features reminiscent of an SW Sextantis-type CV, but lacks others.  We discuss the classification of this puzzling object. ",
        "introduction": "Cataclysmic variables (CVs) are close binary stars in which a white dwarf primary accretes matter via Roche lobe overflow from a secondary, which usually resembles a lower main-sequence star \\citep{warner}. In most CVs an accretion disk encircles the white dwarf.  Optical emission can come from several locations in the system.  The gas stream from the secondary star strikes the outer rim of the accretion disk producing a rapidly fluctuating bright spot (or hot spot). The accretion disk often dominates the light, and instabilities in the disk can produce large variations in the system brightness.  Strong white-dwarf magnetic fields can truncate or even entirely disrupt the accretion disk, allowing material to fall directly onto the white dwarf.  In many systems the accumulated hydrogen-rich material can undergo a thermonuclear runaway, creating a classical nova outburst. The range of phenomena seen in CVs means that determining the subclass of any example can require extensive photometric and spectroscopic study.  H. H. Lanning (\\citealt{lanning1973} and subsequent papers) conducted a search for ultraviolet-bright (UV) sources in the galactic plane by visually examining plates taken with the Palomar Oschin Schmidt.  To date, six supplements have been published, revealing a total of 724 UV sources.  Six of these are included in the {\\it Catalog and Atlas of Cataclysmic Variables} \\citep{downes2001} as confirmed or suspected CVs; they are listed in Table \\ref{lanningtable}.  \\citet{eracleous2002} presented a spectrum of Lanning 386 which showed apparent CV features -- a blue continuum, He I absorption lines, and broad Balmer emission lines. The extensive photometric and spectroscopic observations we present here establish that Lanning~386 is indeed a deeply-eclipsing CV with $P_{\\rm orb} = 0.1640517 \\pm 0.0000001$ d.  Intriguingly, it shares characteristics of both the dwarf novae -- CVs that undergo occasional eruptions thought to be caused by a disk instability -- and the SW Sextantis stars \\citep{thorstensen1991}, a subclass of the persistently-bright novalike variables. ",
        "conclusions": "Our photometry shows Lanning~386 to be an active, eclipsing CV, with $P_{\\rm orb} = 3.94$ h and an eclipse depth of $\\sim 1.5$ mag. Monitoring over many months shows bright outbursts reaching $V\\approx 15.2$ from a quiescent state with $V\\approx 17.2$.  Lanning 386 does not fit neatly into any CV subclass. The long-term variability resembles that of a dwarf nova, and at minimum light the spectrum is typical of that class. However, when most dwarf novae outburst, their Balmer emission becomes very weak, and broad absorption wings appear (e.g., \\citealt{thorstensen98}); it is also unusual (though not unknown) for \\ion{He}{2} to become prominent (e.g., \\citealt{hessman84}). The high-state spectrum seen here is very much unlike this, resembling instead that of a novalike variable. The \\ion{He}{1} line profiles, and the unusual strength of \\ion{He}{2}, are reminiscent of an SW Sextantis star \\citep{thorstensen1991}. In addition, Lanning 386 is similar to other SW Sex stars in that the emission lines are eclipsed much more shallowly than the continuum, indicating a spatially extended source.  The rotational disturbance around eclipse is also seen in some other SW Sex stars (e.g., \\citealt{rodriguez04}). However, other canonical SW Sex characteristics are missing here. The emission line radial-velocity phase in SW Sex stars lags far behind (by $> 45^{\\circ}$) the phase expected for the white dwarf motion, based on the eclipse ephemeris, but in Lanning 386 the phase lag is only $\\sim 11^\\circ \\pm 2^\\circ$. In addition, while the \\ion{He}{1} lines do appear in absorption for part of the orbit, in other SW Sex stars the absorption appears near phase 0.5, whereas it appears somewhat later here.  Nonetheless, the unusual maximum-light spectrum suggests that we should consider carefully whether the long-term light curve in Fig.~\\ref{detailplot} forces us to classify this star as a dwarf nova.  Might another mechanism explain the variability in a manner more consistent with the maximum-light spectrum?  Many novalike variables, and many SW Sex stars in particular, are also VY Scl stars -- novalike variables that occasionally fade deeply for varying amounts of time, for reasons that are not entirely understood.  Might Lanning 386 be a VY Scl star masquerading as a dwarf nova?   This seems unlikely, for two reasons.  First, the long-term light curves of VY Scl stars tend to look quite different from those of dwarf novae -- VY Scl stars typically remain in a high state for months or years, and only occasionally fade by large amounts, up to 5 mag (see examples in \\citealt{honeycutt04}). The light curve here, by contrast, consists of frequent, moderate-amplitude outbursts from a faint state.  Second, the spectra of VY Scl stars in their faint states look as if accretion has nearly stopped, with the Balmer emission lines typically becoming very narrow and the underlying stars becoming evident \\citep{shafter1983, schneider81}.  This is very different from the essentially normal dwarf nova spectrum seen here.  The DQ Her stars, also called intermediate polars, are magnetic systems \\citep[for a review, see][]{patterson94}; some of these undergo outbursts that at least superficially resemble dwarf nova outbursts. Typically, DQ Her outbursts are infrequent compared to those of most dwarf novae; they also tend to decay rapidly \\citep{hellier97}. For this reason, there is little spectroscopy available for outbursts of these systems.  However, spectra do exist for EX Hya in outburst, which, show strong emission lines with interesting velocity structure, reminiscent of the SW Sex stars \\citep{hellier89,hellier00}.  In this way, EX Hya resembles Lanning 386, but the morphological fit is imperfect -- the brightenings we see in Lanning 386 appear to be more frequent than those seen in most DQ Her stars. Also, we have (as of yet) no direct evidence that Lanning 386 is a magnetic system, because we are unaware of any searches for the coherent pulsations and/or circular polarization that would prove the case.  It has been suggested that the outbursts seen in some magnetic systems are due to mass-transfer bursts rather than the disk instabilities that are thought to cause normal dwarf nova outbursts \\footnote{It is possible that mass-transfer bursts could be triggered by an event in the disk.}.  During a period of increased mass transfer, material overflowing the disk could give rise to enhanced emission -- there is evidence for this in EX Hya \\citep{hellier00}. It may be that the outbursts seen in Lanning 386, whatever their trigger, involve enhanced mass transfer -- which may or may not be exclusive to magnetic systems.  We can summarize the morphological conundrum presented by Lanning 386 as follows:  Its minimum-light spectrum and the light curve suggest it is a dwarf nova with frequent, relatively low-amplitude outbursts, while at maximum light it resembles a novalike variable, with some (but not all) the characteristics of an SW Sex star.  EX Hya, a magnetic system, shows a similar set of symptoms, but the presentation is not identical."
    },
    "0801/0801.1548_arXiv.txt": {
        "abstract": "The residuals about the standard $M_{\\rm bh}$-$\\sigma$ relation correlate with the effective radius, absolute magnitude, and S\\'ersic index of the host bulge.  Although, it is noted here that the elliptical galaxies do not partake in such correlations. Moreover, it is revealed that barred galaxies (with their relatively small, faint, and low stellar concentration bulges) can deviate from the $M_{\\rm bh}$-$\\sigma$ relation by $\\delta \\log M_{\\rm bh} \\approx -0.5$ to $-1.0$ dex (their $\\sigma$ values are too large) and generate much of the aforementioned correlations.  Removal of the seven barred galaxies from the Tremaine et al.\\ set of 31 galaxies gives a ``barless $M_{\\rm bh}$-$\\sigma$'' relation with an intrinsic scatter of 0.17 dex (cf.\\ 0.27 dex for the 31 galaxies) and a total scatter of 0.25 dex (cf.\\ 0.34 dex for the 31 galaxies).  The introduction of a third parameter does not reduce the scatter. Furthermore, removal of the barred galaxies, or all the disk galaxies, from an expanded and updated set of 40 galaxies with direct black hole mass measurements gives a consistent result, such that $\\log(M_{\\rm bh}/M_{\\sun}) = (8.25\\pm0.05) + (3.68\\pm0.25)\\log [\\sigma/200\\, {\\rm km\\, s}^{-1}]$. The (barless) $\\sigma$-$L$ relation for galaxies with black hole mass measurements is found to be consistent with that from the SDSS sample of early-type galaxies. In addition the barless $M_{\\rm bh}$-$\\sigma$ relation, the $M_{\\rm bh}$-$n$ relation, and the $M_{\\rm bh}$-$L$ relation are shown to yield SMBH masses less than 2-4$\\times 10^9 M_{\\sun}$. ",
        "introduction": "Tight correlations between supermassive black hole (SMBH) masses and large scale properties of the host bulges are interesting for two obvious reasons. They enable us to predict SMBH masses in thousands of galaxies where the black hole's sphere-of-influence is highly unresolved, and they provide clues to the physical processes responsible for the co-evolution of black hole and host bulge.  Recent endeavors have advocated relations involving not one but two bulge parameters and therein claims of ``fundamental planes'', akin to the Fundamental Plane for elliptical galaxies (Djorgovski \\& Davis 1987; Dressler et al.\\ 1987). The existence of such SMBH fundamental planes imply that current theories for the $M_{\\rm bh}$-$\\sigma$ relation (Ferrarese \\& Merritt 2000; Gebhardt et al.\\ 2000) or the $M_{\\rm bh}$-$L$ relation (McLure \\& Dunlop 2002; Marconi \\& Hunt 2003; updated in Graham 2007), which do not include a third parameter, are incomplete.  This article investigates the fundamental planes for SMBHs\\footnote{The ``fundamental plane of black hole activity'' involving radio core luminosity and X-ray luminosity (Merloni et al.\\ 2003; Falcke et al.\\ 2004) is not addressed here.} involving the parameters $M_{\\rm bh}$, $R_{\\rm e}$, and $\\langle \\mu \\rangle_{\\rm e}$ (Barway \\& Kembhavi 2007), $M_{\\rm bh}$, $R_{\\rm e}$, and $\\sigma$ (Marconi \\& Hunt 2003; de Francesco et al.\\ 2006; Aller \\& Richstone 2007; Hopkins et al.\\ 2007), and, for the first time, $M_{\\rm bh}$, $\\sigma$, and $n$. In Section~2 it is explained why a previous claim for a small `total' scatter (0.19 dex) about the $M_{\\rm bh}$-$R_{\\rm e}$-$\\langle \\mu \\rangle _{\\rm e}$ plane was the result of a miscalculation. In Section~3 it is revealed that the galaxies which deviate from the $M_{\\rm bh}$-$\\sigma$ relation, giving rise to the $M_{\\rm bh}$-$\\sigma$-$R_{\\rm e}$, $M_{\\rm bh}$-$\\sigma$-$L$,  and $M_{\\rm bh}$-$\\sigma$-$n$ relations with less scatter than the $M_{\\rm bh}$-$\\sigma$ relation, are predominantly barred galaxies.  A ``barless $M_{\\rm bh}$-$\\sigma$'' relation, and an elliptical-only $M_{\\rm bh}$-$\\sigma$ relation, is subsequently constructed in Section~4 and found to heavily nullify the evidence for fundamental planes for SMBHs and their host bulges.  Given the recent discussion in the literature about biases in the $M_{\\rm bh}$-$\\sigma$ and/or $M_{\\rm bh}$-$L$ relation, and also in the local sample of galaxies with direct SMBH mass measurements, these concerns are explored here.  In Sections~5 a $\\sigma$-$L$ relation is constructed and shown to be equal to that obtained using SDSS early-type galaxy data, thereby laying to rest concerns that the local sample of galaxies with direct SMBH mass measurements may be biased with respect to the greater population (e.g.\\ Yu \\& Tremaine 2002; Bernardi et al.\\ 2007). Furthermore, in section~6, the $K$-band $M_{\\rm bh}$-$L$ relation and the barless $M_{\\rm bh}$-$\\sigma$ relation are shown to yield consistent results with neither giving SMBH masses greater than $\\sim4 \\times 10^9 M_{\\sun}$. ",
        "conclusions": ""
    },
    "0801/0801.1062_arXiv.txt": {
        "abstract": "The study of the metal enrichment of the intra--cluster and inter--galactic media (ICM and IGM) represents a direct means to reconstruct the past history of star formation, the role of feedback processes and the gas--dynamical processes which determine the evolution of the cosmic baryons. In this paper we review the approaches that have been followed so far to model the enrichment of the ICM in a cosmological context. While our presentation will be focused on the role played by hydrodynamical simulations, we will also discuss other approaches based on semi--analytical models of galaxy formation, also critically discussing pros and cons of the different methods. We will first review the concept of the model of chemical evolution to be implemented in any chemo--dynamical description. We will emphasise how the predictions of this model critically depend on the choice of the stellar initial mass function, on the stellar life--times and on the stellar yields. We will then overview the comparisons presented so far between X--ray observations of the ICM enrichment and model predictions. We will show how the most recent chemo--dynamical models are able to capture the basic features of the observed metal content of the ICM and its evolution. We will conclude by highlighting the open questions in this study and the direction of improvements for cosmological chemo--dynamical models of the next generation. ",
        "introduction": "\\label{Introduction} Clusters of galaxies are the ideal cosmological signposts to trace the past history of the inter--galactic medium (IGM), thanks to the high density and temperature reached by the cosmic baryons trapped in their gravitational potential wells (\\citealt{2002ARA&A..40..539R}; \\citealt{2005RvMP...77..207V}; \\citealt{diaferio2008} - Chapter 2, this volume). Observations in the X--ray band with the Chandra and XMM--Newton satellites are providing invaluable information on the thermodynamical properties of the intra--cluster medium (ICM) (\\citealt{kaastra2008} - Chapter 9, this volume). These observations highlight that non--gravitational sources of energy, such as energy feedback from supernovae (SNe) and Active Galactic Nuclei (AGN) have played an important role in determining the ICM physical properties.  Spatially resolved X--ray spectroscopy permits to measure the equivalent width of emission lines associated to transitions of heavily ionised elements and, therefore, to trace the pattern of chemical enrichment (e.g., \\citealt{2004cgpc.symp..123M} for a review). In turn, this information is inextricably linked to the history of formation and evolution of the galaxy population (e.g., \\citealt{1997ApJ...488...35R} and references therein), as inferred from observations in the optical/near-IR band. For instance, \\citet{2004A&A...419....7D} have first shown that cool core clusters are characterised by a significant central enhancement of the iron abundance, which closely correlates with the magnitude of the Brightest Cluster Galaxies (BCGs). This demonstrates that a fully coherent description of the evolution of cosmic baryons in the condensed stellar phase and in the diffuse hot phase requires properly accounting for the mechanisms of production and release of both energy and metals.  The study of how these processes take place during the hierarchical build up of cosmic structures has been tackled with different approaches. Semi--analytical models (SAMs) of galaxy formation provide a flexible tool to explore the space of parameters which describe a number of dynamical and astrophysical processes. In their most recent formulation, such models are coupled to dark matter (DM) cosmological simulations, to trace the merging history of the halos where galaxy formation takes place, and include a treatment of metal production from type-Ia and type-II supernovae (SN\\,Ia and SN\\,II, hereafter; \\citealt{2004MNRAS.349.1101D}; \\citealt{2005MNRAS.358.1247N}), so as to properly address the study of the chemical enrichment of the ICM. The main limitation of this method is that it does not include the gas dynamical processes which causes metals, once produced, to be transported in the ICM. As a consequence, they provide a description of the global metallicity of clusters and its evolution, but not of the details of its spatial distribution.  In order to overcome this limitation, \\citet{2006MNRAS.368.1540C} applied an alternative approach, in which non--radiative hydrodynamical simulations of galaxy clusters are used to trace at the same time the formation history of DM halos and the dynamics of the gas. In this approach, metals are produced by SAM galaxies and then suitably assigned to gas particles, thereby providing a chemo--dynamical description of the ICM.  \\citet{2006A&A...452..795D} used hydrodynamical simulations, which include prescriptions for gas cooling, star formation and feedback, similar to those applied in SAMs, to address the specific role played by ram--pressure stripping to distribute metals, while \\citet{2007A&A...466..813K} used the same approach to study the different roles played by galactic winds and by ram--pressure stripping. While these approaches offer advantages with respect to standard SAMs, they still do not provide a fully self--consistent picture, in which chemical enrichment is the outcome of the process of star formation, associated to the cooling of the gas infalling in high density regions. In this sense, a fully self--consistent approach requires that the simulations include the processes of gas cooling, star formation and evolution, along with the corresponding feedback in energy and metals.  A number of authors have presented hydrodynamical simulations for the formation of cosmic structures, which include treatments of the chemical evolution at different levels of complexity. \\citet{1996A&A...315..105R} presented SPH simulations of the Galaxy, forming in an isolated halo, by following iron and oxygen production from SN\\,II and SN\\,Ia, also accounting for the effect of stellar lifetimes. \\citet{2001MNRAS.325...34M} performed a detailed analysis of chemo--dynamical SPH simulations, aimed at studying both numerical stability of the results and the enrichment properties of galactic objects in a cosmological context. \\citet{2002MNRAS.330..821L} discussed a statistical approach to follow metal production in SPH simulations, which have a large number of star particles, showing applications to simulations of a disc--like galaxy and of a galaxy cluster. \\citet{2003MNRAS.346..135K} carried out cosmological chemo--dynamical simulations of elliptical galaxies, with an SPH code, by including the contribution from SN\\,Ia and SN\\,II, also accounting for stellar lifetimes. \\citet{2003MNRAS.339.1117V} performed an extended set of cluster simulations and showed that profiles of the iron abundance are steeper than the observed ones. A similar analysis has been presented by \\citet{2006MNRAS.371..548R}, who also considered the effect of varying the IMF and the feedback efficiency on the enrichment pattern of the ICM. \\citet{2005MNRAS.364..552S} presented an implementation of a model of chemical enrichment in the {\\tt GADGET-2} code, coupled to a self--consistent model for star formation and feedback. In their model, which was applied to study the enrichment of galaxies, they included the contribution from SN\\,Ia and SN\\,II, assuming that all long--lived stars die after a fixed delay time. \\citet{Tornatore07,2004MNRAS.349L..19T} presented results from an implementation of a detailed model of chemical evolution model in the {\\tt GADGET-2} code \\citep{2005MNRAS.364.1105S}, including metallicity--dependent yields and the contribution from intermediate and low mass stars (ILMS hereafter).  The major advantage of this approach is that the metal production is self--consistently predicted from the rate of gas cooling treated by the hydrodynamical simulations, without resorting to any external model. However, at present the physical scales involved by the processes of star formation and SN explosions are far from being resolved in simulations which sample cosmological scales.  For this reason, such simulations also need to resort to external sub--resolution models, which provide an effective description of a number of relevant astrophysical processes.  The aim of this paper is to provide a review of the results obtained so far in the study of the chemical enrichment of the ICM in a cosmological context. Although we will concentrate the discussion on the results obtained from full hydrodynamical simulations, we will also present results based on SAMs. As such, this paper complements the reviews by \\citealt{borgani2008} - Chapter 13, this volume, which reviews the current status in the study of the thermodynamical properties of the ICM with cosmological hydrodynamical simulations, and by \\citealt{schindler2008} - Chapter 17, this volume, which will focus on the study of the role played by different physical processes in determining the ICM enrichment pattern. Also, this paper  will not present a detailed description of the techniques for simulations of galaxy clusters, which is reviewed by \\citealt{10_dolag2008} - Chapter 12, this volume.  In Sect.~\\ref{chemical_evolution} we review the concept of model of chemical evolution and highlight the main quantities which are required to fully specify this model. In Sect.~\\ref{globab} we review the results on the global metal content of the ICM, while Sect.~\\ref{profiles} concentrates on the study of the metallicity profiles and Sect.~\\ref{evol} on the study of the ICM evolution. Sect.~\\ref{galaxies} discusses the properties of the galaxy population. Finally, Sect.~\\ref{summary} provides a critical summary of the results presented, by highlighting the open problems and lines of developments to be followed by simulations of the next generation. ",
        "conclusions": "\\label{summary} The study of the enrichment pattern of the ICM and its evolution provides a unique means to trace the past history of star formation and the gas--dynamical processes, which determine the evolution of the cosmic baryons. In this overview, we discussed the concept of model of chemical evolution, which represents the pillar of any chemo--dynamical model, and presented the results obtained so far in the literature, based on different approaches.  Restricting the discussion to those methods which follow the chemical enrichment during the cosmological assembly of galaxies and clusters of galaxies, they can be summarised in the following categories.\\\\ {\\bf (1)} Semi--analytical models (SAMs) of galaxy formation \\citep{2004MNRAS.349.1101D,2005MNRAS.358.1247N}: the production of metals is traced by the history of galaxy formation within the evolving DM halos and no explicit gas--dynamical description of the ICM is included. These approaches are computationally very cheap and provide a flexible tool to explore the space of parameters determining galaxy formation and chemical evolution. A limitation of these approaches is that, while they provide predictions on the global metal content of the ICM, the absence of any explicit gas dynamical description causes the lack of useful information on the spatial distribution of metals.\\\\ {\\bf (2)} SAMs coupled to hydrodynamical simulations. In this case, the galaxy formation is followed as in the previous approach, but the coupling with a hydrodynamical simulation allows one to trace the fate of the enriched gas and, therefore, to study the spatial distribution of metals. Different authors have applied different implementations of this hybrid approach, by focusing either on the role of the chemical evolution model \\citep{2006MNRAS.368.1540C} or on the effect of specific gas--dynamical processes, such as ram--pressure stripping and galactic winds, in enriching the ICM (e.g.,  \\citealt{2007A&A...466..813K}). The advantage of this approach is clearly that gas--dynamical processes are now included at some level. However, the galaxy formation process is not followed in a self--consistent way from the cooling of the gas during the cosmic evolution.\\\\ {\\bf (3)} Full hydrodynamical simulations, which self--consistently follow gas cooling, star formation and chemical evolution \\citep{2003MNRAS.339.1117V,2003MNRAS.346..135K,2005MNRAS.361..983R,Tornatore07}. Rather than being described by an external recipe, galaxy formation is now the result of the cooling and of the conversion of cold dense gas into stars, as implemented in the simulation code. The major advantage of this approach is that the enrichment process now comes as the result of the star formation predicted by the hydrodynamical simulation, without resorting to any external model.  The bottleneck, in this case, is represented by the computational cost, which becomes prohibitively high if one wants to resolve the whole dynamic range covered by the processes of metal production and distribution, feedback energy release, etc. For this reason, these simulation codes also need to resort to sub--resolution models, which provide an effective description of a number of physical processes. However, there is no doubt that the ever improving supercomputing technology and code efficiency make, in perspective, this approach the way to study the past history of cosmic baryons.  As already emphasised, developing a reliable model of the cosmic history of metal production is a non--trivial task. This is due to the sensitive dependence of the model predictions on both the parameters defining the chemical evolution model (i.e., IMF, stellar life--times, stellar yields, etc.), and on the processes, such as ejecta from SN explosions and AGN, stripping, etc., which at the same time regulate star formation and transport metals outside galaxies. In the light of these difficulties, it is quite reassuring that the results of the most recent hydrodynamical simulations are in reasonable agreement with the most recent observational data from Chandra and XMM--Newton.  In spite of this success, cluster simulations of the present generation still suffer from important shortcomings. The most important of them is probably represented by the difficulty in regulating overcooling at the cluster centre, so as to reproduce at the same time the observed ``cool cores'' and the passively evolving massive ellipticals. Besides improving the description of the relevant astrophysical processes in simulation codes, another important issue concerns understanding the possible observational biases which complicate any direct comparison between data and model predictions \\citep{2007arXiv0707.1573K,2007arXiv0707.2614R}. In this respect, simulations provide a potentially ideal tool to understand these biases. Mock X--ray observations of simulated clusters, which include the effect of instrumental response, can be analysed exactly in the same way as real observational data. The resulting ``observed'' metallicity can then be compared with the true one to calibrate out possible systematics. There is no doubt that observations and simulations of the chemical enrichment of the ICM should go hand in hand in order to fully exploit the wealth of information provided by X--ray telescopes of the present and of the next generation."
    },
    "0801/0801.3251_arXiv.txt": {
        "abstract": "Tachyonic inflationary universe model in the context of a  Chaplygin gas equation of state is studied. General conditions  for this model to be realizable are discussed. By using an effective exponential potential we describe in great details the characteristic of the inflationary universe model. The parameters of the model are restricted by using recent astronomical  observations. ",
        "introduction": "It is well known that inflation is to date the most compelling solution to many long-standing problems of the Big Bang model (horizon, flatness, monopoles, etc.) \\cite{guth,infla}. One of the success of the inflationary universe model is that it provides a causal interpretation of the origin of the observed anisotropy of the cosmic microwave background (CMB) radiation, and also the distribution of large scale structures \\cite{astro}.  In concern to higher dimensional theories, implications of string/M-theory to Friedmann-Robertson-Walker (FRW) cosmological models have recently attracted  great deal of attention, in particular, those related to brane-antibrane configurations such as space-like branes\\cite{sen1}. In recent times a great amount of work has been invested in studying the inflationary model with a tachyon field. The tachyon field associated with unstable D-branes might be responsible for cosmological inflation in the early evolution of the universe, due to tachyon condensation near the top of the effective scalar potential \\cite{sen2}, which could also add some new form of cosmological dark matter at late times \\cite{Sami_taq}. In fact, historically, as was empathized by Gibbons \\cite{gibbons}, if the tachyon condensate starts to roll down the potential with small initial velocity, then a universe dominated by this new form of matter will smoothly evolve from a phase of accelerated expansion (inflation) to an era dominated by a non-relativistic fluid, which could contribute to the dark matter detected in these days.      On the other hand, the generalized Chaplygin gas    has been proposed as an  alternative model for  describing the accelerating of the universe. The generalized Chaplygin gas is described by an exotic equation of state of the form \\cite{Bento} \\begin{equation} p_{ch} = - \\frac{A}{\\rho_{ch}^\\beta},\\label{1} \\end{equation} where $\\rho_{ch}$ and $p_{ch}$ are the energy density and pressure of the generalized Chaplygin gas, respectively. $\\beta$ is a constant that lies in  the range $0 <\\beta\\leq 1$, and $A$ is a positive constant. The original  Chaplygin gas corresponds to the case $\\beta = 1$ \\cite{2}. Inserting this equation of state into the relativistic energy conservation equation leads to an energy density given by \\cite{Bento} \\begin{equation} \\rho_{ch}=\\left[A+\\frac{B}{a^{3(1+\\beta)}}\\right]^{\\frac{1}{1+\\beta}}, \\label{2} \\end{equation} where $a$ is the scale factor  and $B$ is a positive integration constant.  The Chaplygin gas emerges as a effective fluid of a generalized d-brane in a (d+1, 1) space time, where the   action can be written as a generalized Born-Infeld action \\cite{Bento}. These models have been extensively studied in the literature \\cite{other}. The model parameters were constrained using currents cosmological observations, such as, CMB \\cite{CMB} and  supernova of type Ia (SNIa) \\cite{SIa}.     In the  model of  Chaplygin inspired inflation usually the scalar field, which drives inflation, is the standard inflaton field, where the energy density given by Eq.(\\ref{2}), can be extrapolate for  obtaining a successful inflation period with a Chaplygin gas model\\cite{Ic}. Recently, the dynamics of the early universe and the initial conditions for inflation in a model with radiation and a Chaplygin gas was studied  in Ref.\\cite{Monerat:2007ud}.  As far as we know, a Chaplygin inspired inflationary model in which a tachyonic field is considered has not been  studied. The main goal of the present work is to investigate the possible realization of a Chaplygin inflationary universe model, where the energy density is driven by a tachyonic field.  We use  astronomical data for constraining the parameters appearing in this model.    The outline of the paper is a follows. The next section presents a short review of the modified Friedmann equation by using a Chaplygin gas, and  we present the tachyon-Chaplygin inflationary model. Section \\ref{sectpert} deals with the calculations of cosmological perturbations in general term.  In Section \\ref{exemple} we use an exponential potential for obtaining explicit expression for the model. Finally, Sect.\\ref{conclu} summarizes our findings. ",
        "conclusions": "}  In this paper we have studied the tachyon-Chaplygin inflationary model. In the slow-roll approximation we have found a general relation between the tachyonic potential and its derivative. This has led us to a general criterium for inflation to occur (see Eq.(\\ref{cond})). We  have also obtained   explicit expressions for the corresponding scalar spectrum index $n_s$ and its running $\\alpha_s$.    By using  an exponential potential with $\\alpha$ fixed (see Eq.(\\ref{pot})) and from the WMAP three year data,  we  found the values of the parameter $A$ and an upper limit for the tachyon potential $V_*$. In order to bring  some explicit results we have taken  $\\alpha=4\\times 10^{-6} m_p$ and $n_s=0.97$, from which we get the values $A\\simeq 2\\times10^{-23}m_p^8$, $V_*\\simeq 5\\times 10^{-13}m_p^4$ and $\\alpha_s\\simeq -0.02$. The restrictions imposed  by currents  observational data allowed us to establish a small range for the parameter $\\alpha$, which become $10^{-6}<\\alpha/m_p\\lesssim 10^{-5}$. From this range,  and from Eqs.(\\ref{A}) and (\\ref{constrain}), we obtained the ranges $0<A/m_p^8\\lesssim 10^{-23}$ and $8.5\\times10^{-14}\\lesssim V_*/m_p^4<8.5\\times10^{-12}$.  In the context of the curvaton scenario, we gave an estimation of the reheating temperature, when the curvaton decay occurs  after it dominates. However, a more accurate calculation for the reheating temperature $T_{rh}$ in the curvaton scenario, would  be necessary for establishing  some constrains on other parameters appearing  in our model. We hope to return to this point in the near future."
    },
    "0801/0801.4537_arXiv.txt": {
        "abstract": "We present two-dimensional (radial velocity, orbital phase) spectroscopic results for the very low mass-ratio close binary AW~UMa which strongly indicate that the spectroscopic mass ratio ($q_{sp} = 0.10$) does not agree with the photometrically derived one and that the widely adopted contact binary model appears to experience serious inconsistencies and limitations for this object. AW~UMa is compared with V566~Oph ($q_{sp} = 0.26$) which we found to behave according to the contact model. Observed broadening functions of AW~UMa can be interpreted by a very strong limb darkening and/or non-solid body rotation of the dominant primary component; the former assumption is unphysical while the differential rotation is not supported by an apparent stability of localized, dark features on the outer side of the primary. There are indications of the existence of an equatorial belt encompassing the whole system. All deficiencies in the interpretation and the discrepancy between the photometric and spectroscopic mass ratio of AW~UMa can be solved within a new model of AW~UMa where both components are detached and the system is submerged in a stream of hot, optically thick matter which mimics the stellar contact. While the masses and their ratio are correctly given by spectroscopy, the photometric picture is heavily modified by the matter engulfing both stars in the equatorial plane. ",
        "introduction": "\\label{intro}  Most of the light-curve analyses of contact binary stars of the W~UMa type assume -- following the seminal papers of \\citet{Lucy1968a,Lucy1968b} -- that shapes of these stars are well described by a common equipotential surface of the Roche model. However, light curve analyses have a common limitation: a light curve is a result of a mapping of a complex three-dimensional surface into an one-dimensional function of time. The information content of many contact binary light curves is low, particularly when total eclipses do not occur. The correctness of the assumed model is of a crucial importance for a proper interpretation of these binaries and still remains an open issue.  The exact solid-body rotation of the Roche model is a strong assumption of the Lucy model. While it is a convenient simplification of the problem, we have no theoretical basis to believe that a strictly solid-body rotation is really present. The Sun does not rotate rigidly and there are many indications that other solar type stars also rotate differentially. We have rather vague ideas for rotation rates some hundred times than the solar. Simple arguments based on the presence of horizontal temperature gradients could be used to argue that convection cells should break into small ones leading to an enhanced turbulent viscosity and thus a solid-body rotation. But -- on the other hand -- the large difference in the component nuclear energy production requires an effective transport over common envelope which cannot happen without extensive horizontal motions. Also, the degree of the solid-body rotation may depend on the mass of the secondary component; when it is very small, its tidal influence may be too weak to prevent the big primary from rotating differentially.  Within the Roche model, a common equipotential surface has a relatively simple shape and is described by only three free geometrical parameters (the mass ratio, $q$, the degree of contact, $F$ and the inclination, $i$); this parametric description is even simpler than for detached binaries where sizes of components are unrelated. All extant solutions of contact binary light curves still require a proof of the strict applicability of the Roche model described in such simple terms.  In this paper we attempt to lift the degeneracy of the 3D into 1D light-curve mapping by using the Broadening Functions formalism \\citep{Rci2002}. A broadening function (BF) gives the surface brightness distribution of a stellar object in the velocity space. This is true if all surface points radiate the same spectrum, which appears to be a particularly valid assumption for contact binaries (see \\citet{ander1983}). In the case of a solid body rotation, the observed radial velocity is proportional to the projected distance from the axis of rotation. When combined with the uniform surface brightness, a single BF gives a 1D image of system in the velocity space. We go one step further and analyze several BF's by arranging and combining them in the orbital phase into 2D functions (radial velocity, orbital phase), effectively the phase dependent, 2D radial-velocity maps of a binary system. The approach of arranging spectra in the phase domain, when the phase coherence is an issue, has already been used by several researches, particularly to study rapid changes in cataclysmic variables. In the DDO studies of the binary stars, it was used mostly as an auxiliary tool \\citep{ddo10,ddo11,ddo12} for faint binaries where individual BF's were poor and some phase averaging was advantageous. But here, it is used as the main way to assess correctness of the solid-body rotation and to analyze the well observed contact binary AW~UMa.  The main object of our study, AW~UMa, is presented in Section~\\ref{obj}. The observations are described in Section~\\ref{obs} while the determination of radial velocities through a Gaussian and rotational profile fitting and through direct modelling of the BFs is presented in Section~\\ref{velocities}. The resulting spectroscopic elements and the absolute parameters are presented in Section~\\ref{elements}. Sections~\\ref{mass-ratio} and \\ref{deviate} discuss the critical issue of the applicability of the Roche model and of the deviations from it that we have discovered. A new model of AW~UMa is presented and discussed in Section~\\ref{newmodel}. The results are summarized in Section~\\ref{sum}. Throughout the paper V566~Oph is discussed as a comparison object for AW~UMa; this contact binary was previously analyzed in \\citet{ddo11}. ",
        "conclusions": "\\label{sum}  Our analysis of the broadening functions for the prototype of extreme-mass ratio contact binaries, AW~UMa, indicates that the generally accepted Roche model experiences major problems in interpretation of this system. We see strong indications that the system is not a contact binary and that both stars are smaller than their inner Roche lobes. The interpretation is complicated by the presence of what appears to be a luminous stream of the matter (optically thick and thus giving the same stellar spectrum), encompassing the primary and possibly the whole system and forming an equatorial belt around it. The equatorial extent of the belt is similar to the size of the common envelope derived in previous contact binary solutions.  The secondary strongly changes its spectroscopic appearance with the orbital phase and does not appear to be a normal star but rather a small stellar core surrounded by a complex disturbance of the belt. The primary is also showing unexpected deviations from the contact model: The ``triangular'' or ``pointed'' shape of the broadening functions (too narrow spectral lines) requires invoking an unusually strong limb darkening and/or a strongly differential rotation of the primary (the equatorial regions moving faster than polar). While the first explanation is unphysical, the second is inferior in terms of the overall fit to the data; it is contradicted by the presence of photospheric spots on the primary which were stable over one year. Our results and suggestions are radical ones because AW~UMa -- which appeared to be one of the best examples of the contact model -- would not be then a contact binary in the sense envisaged by Lucy and by many researchers who followed his ideas. But we are driven to this desperate move only after exhausting all other alternatives.  A belt in the equatorial plane, encompassing both components is not an entirely new idea. In 1950, in his book \\citet{Struve1950} presented a picture (p.~188, Figure 28) which is very close to what we have in mind for AW~UMa. Obviously, the new model would question the validity of the Lucy model for at least some binaries now considered to be in good contact. This is the price we appear to pay for having access to good spectral data and to improved methods of spectral analysis. Indeed, with only photometric information, and with the Lucy model, the realm of contact binaries has been not only amazingly appealing, but also much simpler.  We stress that irrespectively of assumptions used to describe the details of the whole picture, the radial velocities of the photocentres give the spectroscopic mass ratio of AW~UMa close to $q_{sp} \\simeq 0.10$. This value is significantly higher than the seemingly incontestable photometric mass ratio of $q_{ph} \\simeq 0.07 - 0.08$. By taking global averages of our results, the masses of the components are approximately $M_1 = 1.5 \\pm 0.15 \\,M_\\odot$ and $M_2 = 0.15 \\pm 0.015 \\,M_\\odot$, where the uncertainties are estimated from systematic effects related to choices made in the analysis and are far larger than any of the formal uncertainty estimates for individual methods. Since the spectral data (in spite of several serious problems which we tried to stress rather than hide) directly give us radial velocities of the mass centres, the spectroscopic mass ratio of $q_{sp} \\simeq 0.10$ must be close to the true one. The mass-ratio discrepancy, together with direct deviations from the expected velocity field, may be pointing to us a major deficiency in our assumptions about contact binaries. At this point we cannot state whether these deficiencies are {\\it present\\/} or are merely {\\it best detectable\\/} only at the very low mass ratio of AW~UMa. In either case, the deviations from the Roche/Lucy model appear to be highly significant for interpretation of contact binary systems.  \\medskip  This paper is dedicated to the memory of Bohdan Paczy\\'nski, the formidable teacher and towering intellect, but also an understanding colleague and friend. As a student, he discovered AW~UMa in 1963 and recognized it as a particularly important one for understanding of contact binary stars. This star accompanied him over 44 years of his immensely productive astronomical life to his last contribution which he directly supervised \\citep{BeP2007}.  We thank Stefan Mochnacki for valuable comments and suggestions and to Kosmas Gazeas for his permission to use his DDO observations. The research made use of the SIMBAD database, operated at the CDS, Strasbourg, France and accessible through the Canadian Astronomy Data Centre, which is operated by the Herzberg Institute of Astrophysics, National Research Council of Canada. The stays of TP at DDO and the research of SMR have been supported by a grant to SMR from the Natural Sciences and Engineering Council of Canada. This research has been supported in part by the Slovak Academy of Sciences under a Grant No. 2/7010/7."
    },
    "0801/0801.3910_arXiv.txt": {
        "abstract": "We present a comparison of the properties of a giant radio galaxy and the ambient intergalactic medium, whose properties are inferred from the large-scale distribution in galaxies.  The double lobes of the radio galaxy MSH~05$-$2{\\it 2} are giant---1.8~Mpc projected linear size---and interacting with the environment outside the interstellar medium and coronal halo associated with the host galaxy. The radio lobes appear to be relicts and the double structure is asymmetric. We have examined the large-scale structure in the galaxy distribution surrounding the radio source. The host galaxy of MSH~05$-$2{\\it 2} is associated with a small group that lies close to the boundary of sheet-like and filamentary density enhancements, and adjacent to a void. Assuming that the galaxies trace gas, the asymmetries in the radio morphology in this case study appear related to the anisotropy in the medium. However, the observed overdensities and structure formation models for the heating of the intergalactic medium (IGM) suggest a density-temperature product for the IGM environment that is an order of magnitude below that expected from the properties of the radio source. The discordance suggests that even sources like MSH~05$-$2{\\it 2}, which are observed in the relatively low-density IGM environment associated with the filamentary large-scale structure and have multiple signatures of being relicts, may be overpressured and evolving towards an equilibrium relaxed state with the ambient IGM. In this picture, relaxed relicts outside cluster and group environments would have surface brightness well below that observed in MSH~05$-$2{\\it 2} and the detection limits of current wide-field radio surveys. Alternately, it is speculated that astrophysical feedback originating in galaxy overdensities observed 1--2~Mpc to the N and NE of MSH~05$-$2{\\it 2} might be the mechanism for the heating of the ambient IGM gas. In such a scenario, MSH~05$-$2{\\it 2} would be a relaxed relict in equilibrium with the IGM environment and such feedback might additionally aid in the formation of the observed asymmetries in the radio morphology. ",
        "introduction": "Young radio sources evolve within the host interstellar medium; as they grow larger the radio structures presumably interact with the coronal X-ray halos \\citep{Fo85, Su04} of their host elliptical galaxy. Not only the lobe morphology, but also the lobe separation asymmetries, are believed to be influenced by the environment in which the radio galaxy evolves. For example, associated extended emission line regions have been shown to retard the growth of the lobe on that side \\citep{Mc91}.  Galaxies that host active galactic nuclei (AGNs) occur in a range of environments and the jets from the central engines in these galaxies encounter a variety of conditions as they emerge into the intergalactic medium (IGM).  Membership and proximity to clusters, groups, and the filaments and sheets in the large scale structure of the Universe might all determine the gaseous intergalactic environment and these would, additionally, depend on cosmic epoch. Within clusters of galaxies, the prevailing intracluster winds and galaxy peculiar velocities shape the lobes of double radio sources into head-tail and wide-angle morphologies \\citep{Bl98}.  Outside clusters and groups, giant radio sources whose linear sizes are in the range of several megaparsecs evolve in the more tenuous intergalactic medium, the nature of which is obscure.  Asymmetries caused by jet-IGM interactions would be expected to be manifest particularly in circumstances where the host galaxies are at the edge of galaxy clusters or filamentary structures or more generally at the edge of galaxy overdensities.  \\citet{Sa86} investigated the pronounced asymmetry in the lobe lengths of a small number of relatively nearby giant radio galaxies and found that the shorter lobe was preferentially located on the side that had excess galaxies over the background. The asymmetry was presumably a result of the higher density of gas associated with galaxy clustering seen on that side; this is consistent with the expectation that the large-scale intergalactic gaseous medium would be traced by the distribution of galaxies.  Large scale redshift surveys have revealed the richness in galaxy distribution structures---galaxies are observed to form criss-crossing chains and sheets several megaparsec wide with regions largely devoid of galaxies in between. Most giant radio galaxies do not belong to rich clusters of galaxies; they tend to populate regions with low galaxy density, corresponding to galaxy filamentary structures. Therefore, the medium associated with this component of the large scale structure is of relevance. While galaxy clusters are known to be permeated with hot thermal gas (seen via X-ray emission), little is known regarding gas associated with galaxy filaments.  Simulation studies of galaxy formation \\citep{Ce99,Da01} have suggested the existence of a warm-hot intergalactic medium (WHIM) at low redshifts. The puzzling deficit in the baryon content of the low-redshift Universe was identified with the  WHIM: a gas that is intermediate in temperature between warm photo-ionized gas in voids and hot intracluster gas, and is too low in density to be easily detectable in X-rays. The bulk of this gas is believed to be distributed along galaxy filaments and traced by moderate overdensities; it mostly lies outside of virialized structures like groups or clusters. The gas is believed to be enriched via outflowing galactic superwinds from supernovae in galaxies; as a consequence, highly ionized oxygen and neon are considered important species that would render the WHIM visible via absorption of the continuum emission of background quasars. There are several ongoing searches for this absorption signature and today there is tantalizing evidence for the WHIM via such methods (see, for example, \\citet{Ni05}). However, we recognise that the background AGN population is sparse, and sparser still are powerful AGN capable of rendering robust absorption detections.  Additional and independent probes of this key component of the intergalactic medium are, therefore, highly desirable.  The WHIM gas is believed to have been heated as a result of hydrodynamic shocks during the non-linear evolution of the large-scale filamentary structure in the Universe. However, there are likely to be additional complex astrophysical feedback mechanisms that might very well be a significant contributor to the entropy content of the WHIM.  This is an additional motivation for observations that probe the WHIM gas properties and provide constraints on feedback mechanisms.  In this paper we have for the first time used radio observations of a 1.8~Mpc giant radio galaxy MSH~05$-$2{\\it 2} (referred to as 0503$-$286 by \\cite{Sa86} and \\cite{Su86}) and galaxy redshift surveys to compare the properties of the synchrotron plasma lobes with the properties of the ambient IGM, which is believed to be the WHIM in this particular case. Motivated by the earlier hypothesis that the lobe extent asymmetry in this source might be related to the asymmetry in the sky distribution in galaxy number density, we have measured redshifts of galaxies over a large area around the radio source. Using these new data, along with archival redshift measurements, we have constructed the 3D distribution of galaxies around the radio source and identified galaxy filaments that might be responsible for the asymmetry. Our modeling of the evolution of this radio source leads to inferences for the properties of the gas associated with the large-scale structure. With this case study we illustrate how giant radio galaxies may be useful probes of the IGM gas.  We organise the paper as follows: we begin with an introduction to the giant radio galaxy MSH~05$-$2{\\it 2} in Section~2. In Section~3 we present our new Very Large Array (VLA) and Molonglo Observatory Synthesis Telescope (MOST) images of this radio source. The next section describes the galaxy redshift data: first we discuss the optical field over nearly $45\\degr$ radius and then we present our 2-degree field (2dF) spectroscopy with the Anglo-Australian telescope (AAT). In Section~5 we derive physical properties of the radio source. Section~6 presents an estimation of the properties of the ambient intergalactic medium for comparison with the expectations based on the internal properties of the source; section~7 examines plausible causes for asymmetries in the source; section~8 is a concluding discussion.  Throughout we adopt a flat cosmology with Hubble constant $H_{o}$ = 71~km~s$^{-1}$~Mpc$^{-1}$ and matter density parameter $\\Omega_{m} = 0.26$. The optical host for the radio source is the luminous elliptical galaxy ESO~422-G028 (also cataloged as MCG~$-$05-13-003; with J2000.0 coordinates R.A. 05$^{h}$ 05$^{m}$ $49\\fs22$, decl. $-$28$\\degr$ 35$\\arcmin$ $19\\farcs4$). The redshift of the giant radio galaxy was first measured by \\citet{Sa86} and \\citet{Su86}; more recently the 6dF redshift survey \\citep{Jo04} gave an improved value of z=0.038286. At this redshift $1\\arcmin=45$~kpc and MSH~05$-$2{\\it 2}, with an angular size of $40\\arcmin$, has a projected linear size of 1.8~Mpc. ",
        "conclusions": "We derive a lower limit of $4.1 \\times 10^{-15}$~N~m$^{-2}$ to the synchrotron pressure within the radio lobes; this suggests a density-temperature product of $3 \\times 10^{8}$~K~m$^{-3}$ for the ambient IGM. The galaxy group, nearby galaxy concentrations, and the large scale sheet at the edge of which the radio source is located, all have fractional density contrasts of order 10 or smaller.  The radio source, which has a projected linear size of 1.8~Mpc, is presumably interacting with the IGM associated with a filamentary large-scale structure in the galaxy distribution: the derived density-temperature product for this structure is a measurement of the physical state of the IGM gas in the vicinity of the radio source. Assuming that galaxies trace the unseen IGM gas, the densities and temperatures we expect for the IGM environment---based on the local galaxy overdensity and assuming that the IGM gas is heated by hydrodynamic shocks associated with the formation of the observed large-scale structure---falls at least an order of magnitude short of the expectations based on the radio properties of the source.  There are several plausible explanations for this basic result that the IGM pressure is at least an order of magnitude below that in the radio lobes, we discuss some of them below.  \\subsection{Do the galaxies in the filaments trace the gas?}  The distribution in the unseen collisionless dark matter dominates the gravitational potential of the large scale structure.  Galaxies display luminosity segregation \\citep{Pe02}---the distribution of luminous galaxies are biased---and, therefore, derivations of mass overdensity factors based on the space distribution of luminous galaxies would be overestimates.  Consequently, the gravitational potential and gravitational accelerations would also be overestimates.  We have assumed that the galaxies trace the gas, and that the galaxy overdensity factors derived represent the IGM gas overdensity. Unlike dark matter and galaxies, gas is dissipative and it is possible that the gas distribution is significantly more biased compared to galaxies. However, this might only be expected in later stages of non-linear structure growth when the overdensity factors attain values exceeding about $10^{2}$ and radiative cooling becomes significant. In the mildly overdense stages when the overdensity factor is a few 10's and the structures are yet to collapse, we might expect the gas and galaxies to have similar overdensity factors.  \\subsection{The ambient medium of the radio lobes: is it an intra-group medium?}  The small group of which the host is a member has a relatively small velocity dispersion compared to that typically observed in groups of galaxies (see, for example, the compilation in \\citet{Mu03}) and, therefore, it is expected that any virialized gas associated with the group would have a temperature at the lower end of the observed distribution: 0.4--1.5~keV. Groups with $k_{\\rm B}T < 1$~keV are detected to have a gas extent less than half the virial radius of the galaxies in the group, or out to a radial distance of less than 0.4~Mpc. It may be noted here that the lobes of MSH~05$-$2{\\it 2} extend well beyond the obvious extent of the group, of which the host galaxy is a member.  Although it is unlikely that the gas associated with the small group---of which the host galaxy is the dominant member---extends over the entire radio source, this expectation needs to be confirmed with sensitive X-ray imaging of the environment of the radio source. Based on the observations to date, it is reasonable to assume that the radio lobes interact with the more diffuse IGM associated with the large-scale filament and not thermal intra-group gas.  \\subsection{Are the radio lobes in equilibrium with the IGM?}  Active radio sources inflate cocoons that are likely overpressured and expanding with speeds limited by ram pressure balance. In this model, the synchrotron cocoons are enveloped in gas that has been shock heated and compressed by the mechanical action of the expanding radio lobe in the vicinity of the contact discontinuity. Once the central beams are switched off, the relict radio sources are believed to rapidly diminish in surface brightness---as evidenced by the paucity of relicts---as the cocoons and enveloping thermal gas pressure equilibrates with the IGM.  An assumption made here is that the radio lobes of MSH~05$-$2{\\it 2} are relicts and, therefore, in pressure equilibrium with the ambient thermal gas outside the contact discontinuity.  In the light of the finding that the pressure in the external IGM might be an order of magnitude underpressured, a possible explanation is that the lobes of MSH~05$-$2{\\it 2} are not relaxed---despite the distinctive aspects of the radio properties that suggest a relaxed relict state for the lobes---but are being observed in a rare transition phase as a fading relict continuing ram-pressure balanced expansion.  An observational support for this picture is the enhanced and circumferential magnetic field along the boundaries of the cocoons, which are indicative of radial compression for the synchrotron plasma.  In this interpretation, the lobes of MSH~05$-$2{\\it 2}, with internal pressure exceeding $4.1 \\times 10^{-15}$~N~m$^{-2}$, are over pressured compared to the external thermal gas, with density less than $2.2 \\times 10^{-27}$~kg~m$^{-3}$ and pressure about $2.3 \\times 10^{-16}$~N~m$^{-2}$. Ram pressure expansion with speed $v/c \\approx 0.0045$ over a timescale 0.3~Gyr, which is similar to the lifetime of the source, would expand the lobes by factor two and result in a relaxed relict in static pressure equilibrium with the IGM associated with the large-scale filament in which the source is embedded.  The surface brightness of the relaxed relict at 1.4~GHz would be less than 0.25~mJy~arcmin$^{2}$, much less than the current surface brightness of about 30~mJy~arcmin$^{2}$ and well below surface brightness detection limits of wide-field surveys. This scenario for the evolution in MSH~05$-$2{\\it 2} would be consistent with the rarity of relaxed relict radio galaxies---the `dead' radio galaxies (\\citet{gi88}; see also the discussion in \\citet{Blundell00}).  \\subsection{Additional sources of heating of the IGM environment}  We discuss here the possibility that there is likely a different cause for the heating of the IGM gas in the vicinity of the radio source, which dominates over hydrodynamic shock heating or mechanical heating of the ambient gas by the radio source itself and might permit the radio lobes to be in pressure equilibrium with the IGM.  The galaxy distribution in the vicinity of the radio source (Section~6.1) shows a concentration of galaxies centered 1.8~Mpc NE of the host galaxy that is significantly closer to the northern lobe than the southern lobe. The observational evidence for an association between the asymmetries in the radio source and the large scale distribution of galaxies indicates that a model worth considering is one in which shocks associated with galactic super winds (GSWs) or mechanical heating associated with any AGNs in this galaxy concentration might enhance the density-temperature product in the IGM environment of the radio source and, additionally, also contribute to the dynamical evolution of MSH~05$-$2{\\it 2} and the displacement of the northern lobe off axis and in the SW direction.  Hydrodynamic simulations of the evolution in large-scale structure, incorporating such GSWs \\citep{Ce06} have indicated the power of such winds in translating gas from higher density regions into lower density peripheral regions of filaments as well as driving shocks into regions with low fractional overdensities that would otherwise not have experienced shock heating during non-linear structure formation.  Apart from heating the ambient gas around the radio lobes of MSH~05$-$2{\\it 2}, the GSWs might also transport gas from the galaxies in the concentrations and into the environment of the radio source.  In this picture, the energy arising from GSW or AGN related activity in the galaxies to the NE and propagating in the low density WHIM might aid the movement of the northern radio lobe towards the SW, as discussed in Section~7.  The estimated age of the relict source, $t = 0.3$~Gyr, implies a fairly large displacement speed of order 600~km~s$^{-1}$. \\citet{Ru05} measure a median `maximum' velocity for the winds in their sample of starburst galaxies to be 350~km~s$^{-1}$; however, one galaxy in their sample was measured to have a wind speed of 1100~km~s$^{-1}$. \\citet{He02} suggests that primary wind fluid velocities might be a few thousand km~s$^{-1}$.  Additionally, the superwinds in the environment of MSH~05$-$2{\\it 2} might be a combined wind from several of the galaxies in the nearby concentration to the NE. GSW speeds are, therefore, in the vicinity of the displacement speed inferred for the lobe. It may also be noted that in simulations, GSWs tend to move towards the low density regions (\\citet{Ka07} and references therein). Since MSH~05$-$2{\\it 2} is located between the galaxy concentration to the NE and a void to the SW, winds from the observed galaxy concentration would be expected to displace the northern lobe in the required direction. Therefore, it is plausible that such astrophysical feedback mechanisms, which may be responsible for the high pressure in the ambient IGM, may also contribute to the radio source dynamics and result in the observed displacement of the northern lobe over the source lifetime.  We have looked for evidence of current AGN activity in the galaxies located in the nearby concentration to the NE of MSH~05$-$2{\\it 2}, which are within $\\pm0.003$ of the host redshift.  A galaxy is coincident with a ROSAT X-ray source 1WGA~J0506.8$-$2810, another may be identified with a radio source in the NRAO VLA sky survey (NVSS) as well as a ROSAT object 1WGA~J0506.6$-$2814. About 10\\% of the galaxies appear to be AGNs, which may be responsible for the feedback mechanisms being considered here.  We have examined an archival ROSAT PSPC image of the sky region in the vicinity of the radio source for evidence for extended emission on the scale of the radio source.  This was a 22~ks pointed observation. First, smoothing the ROSAT image does not indicate any evidence for extended X-ray emission associated with the group of which the host galaxy is a member. Second, if the double radio lobes were enveloped by relaxed hot thermal gas on a scale similar to the size of the radio lobes, and with a pressure as indicated by the radio lobes, the expected counts in the 0.2--2.0~keV ROSAT band would be below the sensitivity to such structure. Existing data do not rule out the GSW or AGN feedback models considered in this section.  Adopting an electron density of $n_{e} = 0.26$~m$^{-3}$ and a temperature of 0.5~keV (see Section~6.2), and the expression for temperature relaxation timescale in RA: 05$^{h}$ 03$^{m}$ $51\\fs0$, DEC: $-$28$\\degr$ 45$\\arcmin$ 00$\\arcsec$ \\citet{Yo05}, we infer that the electron-ion equilibration timescale in the ambient IGM gas is about 0.014~Gyr. A higher ambient density would reduce the equilibration time; however, a higher ambient temperature resulting from heating via shocks associated with GSWs or AGNs could potentially increase the equilibration timescale significantly. It is, therefore, possible---although unlikely---that in the case of a shock heated IGM the equilibration timescale exceeds the age associated with the relict phase, as well as the age of the radio source. As a consequence, we cannot rule out a model where the ambient IGM beyond the contact discontinuity of the radio synchrotron plasma is composed of hot ions and relatively cooler electrons, and hence difficult to detect in its continuum radiation."
    },
    "0801/0801.2707_arXiv.txt": {
        "abstract": "The solar magnetic field is key to understanding the physical processes in the solar atmosphere. Nonlinear force-free codes have been shown to be useful in extrapolating the coronal field upward from underlying vector boundary data.  However, we can only measure the magnetic field vector routinely with high accuracy in the photosphere, and unfortunately these data do not fulfill the force-free condition. We must therefore apply some transformations to these data before nonlinear force-free extrapolation codes can be self-consistently applied.  To this end, we have developed a minimization procedure that yields a more chromosphere-like field, using the measured photospheric field vectors as input.  The procedure includes force-free consistency integrals, spatial smoothing, and --- newly included in the version presented here --- an improved match to the field direction as inferred from fibrils as can be observed in, {\\it e.g.}, chromospheric H$\\alpha$ images. We test the procedure using a model active-region field that included buoyancy forces at the photospheric level.  The proposed preprocessing method allows us to approximate the chromospheric vector field to within a few degrees and the free energy in the coronal field to within one percent. ",
        "introduction": "The solar interior, photosphere and atmosphere are coupled by magnetic fields. It is therefore important to gain insights about the magnetic field structure in all layers of the Sun and solar atmosphere. Direct and accurate measurements of the magnetic field vector are typically carried out only on the photosphere. Although measurements in higher layers are available for a few individual cases, {\\it e.g.} in the chromosphere by Solanki {\\it et al.} (2003) and in the corona by \\inlinecite{lin:etal04}, the line-of-sight integrated character of such chromospheric and coronal magnetic field measurements complicates their interpretation \\cite{Kramar:etal06}. Knowledge of the magnetic field in the corona is essential, however, to understand basic physical processes such as the onset of flares, coronal mass ejections and eruptive prominences.  Inferences of the coronal magnetic field can be obtained by extrapolating measurements of the photospheric magnetic field vector ({\\it e.g.} observed by {\\it Hinode}/SOT, SOLIS or the upcoming SDO/HMI instruments) into the corona. Because the magnetic pressure dominates the plasma pressure in active-region coronae, making the plasma $\\beta$ low, [See work by \\inlinecite{gary01} and \\inlinecite{schrijver:etal05}, which discuss the plasma beta over active regions and over the quiet Sun, respectively], these extrapolations neglect non-magnetic forces and assume the coronal magnetic field ${\\bf B}$ to be force-free, such that it obeys:  \\begin{eqnarray} \\nabla \\cdot {\\bf B} &=& 0 \\label{divB}, \\\\ (\\nabla \\times {\\bf B}) \\times {\\bf B} &=& 0 \\label{jxb}. \\end{eqnarray}  Equation (\\ref{jxb}) implies that the electric current density ${\\mu_0 \\bf j}=\\nabla \\times {\\bf B}$ is parallel to the magnetic field ${\\bf B}$. Starting more than a quarter century ago \\cite{sakurai81}, different mathematical methods and numerical implementations have been developed to solve the nonlinear force-free equations (\\ref{divB}) and (\\ref{jxb}) for the solar case.  See, for example, \\cite{sakurai89,amari:etal97,wiegelmann07a} for review papers and \\cite{schrijver:etal06,metcalf:etal07} for evaluations of the performance of corresponding computer programs with model data. The codes use the magnetic field vector (or quantities derived from the magnetic field vector) on the bottom boundary of a computational domain as input. One would like to prescribe the measured photospheric data as the bottom boundary of nonlinear force-free fields (NLFFF) codes, but there is a problem: the observed photospheric magnetic field is usually not force-free. The relatively high plasma $\\beta$ in the photosphere means that non-magnetic forces cannot be neglected there and that such photospheric magnetic field data are not consistent with well known force-free compatibility conditions defined in \\cite{aly89}.  Recently, \\inlinecite{wiegelmann:etal06} developed a scheme that mitigates this problem, in which the inconsistent and noisy photospheric vector magnetograms used as bottom boundary conditions are preprocessed in order to remove net magnetic forces and torques and to smooth out small-scale noise-like magnetic structures.  The resulting magnetic field data are sufficiently force-free and smooth for use with extrapolation codes, but also are found to bear a high resemblance to chromospheric vector magnetic field data.  This leads us to the question whether we can constrain the preprocessing tool further by taking direct chromospheric observations, such as H$\\alpha$ images, into consideration.  We will investigate this topic in the present work. ",
        "conclusions": "\\label{conlusions} Within this work we developed an improved algorithm for the preprocessing of photospheric vector magnetograms for the purpose of getting suitable boundary conditions for nonlinear force-free extrapolations. We extended the preprocessing routine developed by \\inlinecite{wiegelmann:etal06}, which is referred to here as ``classical preprocessing''. The main motivation for this work is related to the fact that active-region coronal magnetic fields are force-free due to the low $\\beta$ coronal plasma, but the magnetic field vector can be measured with high accuracy only on the photosphere, where the plasma $\\beta$ is about unity and non-magnetic forces cannot be ignored. Our original (``classical'') preprocessing removes these non-magnetic forces and makes the field compatible with the force-free assumption leading to more chromospheric-like configurations.  In this study, we have found that by taking direct chromospheric observations into account (such as by using fibrils seen in H$\\alpha$ images), the preprocessing is improved beyond the classical scheme.  This improved scheme includes a term which minimizes the angle between the preprocessed magnetic field and the fibrils. We tested our method with the help of a model active region developed by van Ballegooijen {\\it et al.} (2007), which includes the forced photospheric and force-free chromospheric and coronal layers. This model has been used by Metcalf {\\it et al.} (2007) for an inter-comparison of nonlinear force-free extrapolation codes. The comparison revealed that the model coronal magnetic field was reconstructed very well if chromospheric magnetic fields have been used as input, but in contrast the reconstructed fields compared poorly when unprocessed model-photospheric data were used. Classical preprocessing significantly improves the result, but the H$\\alpha$ preprocessing developed in this paper is even better as the main features of the model corona are reconstructed with high accuracy. Our extended preprocessing tool provides a fair estimate of the chromospheric magnetic field, which is used as boundary condition for computing the nonlinear force-free coronal magnetic field. In particular, the magnetic energy in the force-free domain above the chromosphere agrees with the model corona within $1\\%$, even if only strong-field regions of the model chromosphere, where the fibrils can be identified with highest accuracy, influence the final solution.  From these tests we conclude that our improved preprocessing routine is a useful tool for providing suitable boundary conditions for the computation of coronal magnetic fields from measured photospheric vector magnetograms as provided for example from {\\it Hinode}. The combination of preprocessing and nonlinear force-free field extrapolations seem likely to provide accurate computation of the magnetic field in the corona.  We will still not get the magnetic field structure in the relative thin layer between/in the photosphere and the chromosphere correct, because here non-magnetic forces cannot be neglected due to the finite $\\beta$ plasma. Although this layer is vertically thin ({\\it e.g.}, 2 vertical grid points in the van Ballegooijen {\\it et al.} (2007) model compared to 256 vertical grid points in the corona) it contains a significant part of the total magnetic energy of the entire domain, see Metcalf {\\it et al.} (2007).  Unfortunately, this part of the energy cannot be recovered by force-free extrapolations, because the region is non-force-free. Our improved preprocessing routine includes chromospheric information and therefore provides us with a closer approximation of the chromospheric magnetic field.  This leads to more accurate estimates of the total magnetic energy in the corona.   A further improvement of the preprocessing routine could be done with the help of additional observations, {\\it e.g.} the line-of-sight chromospheric field, as planned for SOLIS. One could include these measurement directly in the $L_3$-term (\\ref{defL_3}) either as the only information or in some weighted combination with the photospheric field measurement.  An investigation of the true 3D-structure of the thin non-force-free layer between photosphere and chromosphere requires further research. First steps towards non-force-free magnetohydrostatic extrapolation codes \\cite{wiegelmann:etal06b} might help to reveal the secrets of this layer. Non force-free magnetic field extrapolations will require additional observational constraints, because the magnetic field, the plasma density and pressure must be computed self-consistently in one model.     \\begin{acks} The work of T. Wiegelmann was supported by DLR-grant 50 OC 0501 and J.K. Thalmann got financial support by DFG-grant WI 3211/1-1.  M. DeRosa, T. Metcalf, and C.  Schrijver were supported by Lockheed Martin Independent Research funds.  We acknowledge stimulating discussions during the fourth NLFFF-consortium meeting in June, 2007 in Paris. \\end{acks}  \\appendix"
    },
    "0801/0801.0644_arXiv.txt": {
        "abstract": "We investigate the correlation of gravitational lensing of the cosmic microwave background (CMB) with several tracers of large-scale structure, including luminous red galaxies (LRGs), quasars, and radio sources.  The lensing field is reconstructed based on the CMB maps from the Wilkinson Microwave Anisotropy Probe (WMAP) satellite; the LRGs and quasars are observed by the Sloan Digital Sky Survey (SDSS); and the radio sources are observed in the NRAO VLA Sky Survey (NVSS). Combining all three large-scale structure samples, we find evidence for a positive cross-correlation at the $2.5\\sigma$ level ($1.8\\sigma$ for the SDSS samples and $2.1\\sigma$ for NVSS); the cross-correlation amplitude is $1.06\\pm 0.42$ times that expected for the WMAP cosmological parameters. Our analysis extends other recent analyses in that we carefully determine bias weighted redshift distribution of the sources, which is needed for a meaningful cosmological interpretation of the detected signal. We investigate contamination of the signal by Galactic emission, extragalactic radio and infrared sources, thermal and kinetic Sunyaev-Zel'dovich effects, and the Rees-Sciama effect, and find all of them to be negligible. ",
        "introduction": "Great progress has been made in cosmology in the past several years, in large part due to measurements of the anisotropies of the cosmic microwave background (CMB).  Most recently, the Wilkinson Microwave Anisotropy Probe (WMAP) experiment has generated high-resolution, all-sky maps of the CMB \\cite{2007ApJS..170..288H,2007ApJS..170..335P}.  While a great deal of attention has been given to the ``primary'' anisotropies in the CMB imprinted at $z\\sim 10^3$, experiments such as WMAP are also sensitive to secondary anisotropies caused by electron scattering or gravitational potentials at lower redshifts.  Among these secondary effects are the thermal and kinetic Sunyaev-Zel'dovich (tSZ/kSZ) effects, the integrated Sachs-Wolfe (ISW) effect, and gravitational lensing.  These secondary anisotropies are important for two reasons: first, they can act as ``foregrounds'' for primary CMB studies if they are not adequately modeled; and second, they provide information about the growth of structure at low redshift. Since the secondary anisotropies are subdominant on the large angular scales observed by WMAP, they are most easily detected by cross-correlation with large scale structure (LSS). Several groups have cross-correlated WMAP data with LSS tracers in order to study the ISW and tSZ secondaries and extragalactic point sources \\cite{1998NewA....3..275B, 2003MNRAS.346..940D, 2003ApJ...597L..89F, 2003astro.ph..7335S, 2004Natur.427...45B, 2004ApJ...608...10N, 2004MNRAS.347L..67M, 2004MNRAS.350L..37F, 2004PhRvD..69h3524A, 2005PhRvD..72d3525P, 2006MNRAS.365..171G, 2006MNRAS.365..891V, 2006MNRAS.372L..23C, 2006PhRvD..74d3524P, 2006PhRvD..74f3520G, 2007MNRAS.376.1211M, 2007MNRAS.377.1085R}.  Lensing of the CMB is a secondary anisotropy that has attracted considerable attention because of its potential to probe the matter distribution at intermediate redshifts $z\\sim O(1)$ \\cite{1999PhRvD..59l3507Z, 2002ApJ...574..566H, 2003ApJ...590..664S, 2003PhRvL..91x1301K}.  Lensing of the CMB has several potential advantages as a cosmological probe relative to using lensing of galaxies: (1) it probes a redshift range inaccessible to galaxy lensing surveys since its ``source screen'' is at $z\\approx 1100$; (2) its redshift is known to high accuracy; and (3) as a Gaussian random field the CMB does not suffer from intrinsic alignments.  Lensing of the CMB does however have its limitations: (1) the low signal-to-noise ratio with current data; (2) one cannot make tomographic measurements of the redshift dependence of the lensing signal without including external data sets \\cite{2002PhRvD..65b3003H}; (3) foregrounds; and (4) instrumental systematics are a problem, especially for the lensing auto-power spectrum.  In the present paper, we aim to measure the weak gravitational lensing effect in the WMAP data.  We do this in cross-correlation with large-scale structure (LSS) in order to improve the signal-to-noise ratio.  We reconstruct the lensing deflection field \\footnote{To be precise, we compute a filtered version of the lensing deflection field (\\ref{eq:v}) to avoid technical problems due to sky cuts.  See Section~\\ref{ss:cmblens} for details.} from WMAP using quadratic estimator methods \\cite{2001ApJ...557L..79H}.  We then measure the cross-power spectrum of this deflection field with luminous red galaxy (LRG) and quasar maps obtained from the Sloan Digital Sky Survey (SDSS) and with the radio source maps from the NRAO VLA Sky Survey (NVSS). The detailed sample definitions, an analysis of the ISW effect, and parameter constraints are presented in a companion paper (``Paper I'' \\cite{paperi}).  There have been two previous searches for the lensing effect on the CMB in cross-correlation.  Hirata {\\em et~al.} \\cite{2004PhRvD..70j3501H} used the the first-year WMAP data release and a smaller sample of LRGs from the SDSS.  They found a nondetection: $1.0\\pm 1.1$ times the expected signal.  More recently, Smith {\\em et~al.} \\cite{2007PhRvD..76d3510S} used the third-year WMAP data and the NVSS radio sources (although with different selection cuts than those used here); they find a $3.4\\sigma$ result ($1.15\\pm 0.34$ times the expected signal, although their expected signal is 15\\%\\ larger due to a different bias, redshift distribution, and fiducial cosmology).  The analysis presented here draws heavily on that of Hirata {\\em et~al.} \\cite{2004PhRvD..70j3501H}, but benefits from: (i) reduced instrument noise due to the longer integration time in the 3-year WMAP data release; (ii) a larger sky area available in the SDSS; (iii) the inclusion of SDSS quasars and NVSS sources in addition to the LRGs; and (iv) improved treatment of tSZ and point source contamination. The current analysis also extends these previous analyses in that we put a lot of effort in accurate determination of the bias weighted redshift distribution, using cross-correlation information with other galaxy samples. Even though this is not required for claiming a detection, it is needed for any cosmological interpretation of the signal and thus for any consistency check of cosmology based on this analysis.  The analysis of this paper will be mostly based on a ``fiducial'' cosmology, which is the 3-year WMAP-only best-fit 6-parameter cosmology \\cite{2007ApJS..170..377S} from the point in the WMAP Markov chain with the highest likelihood.  The basic objective is to measure the strength of the lensing signal, determine its statistical error and sensitivity to systematics, and establish whether it is consistent with expectations for the fiducial cosmology.  We explore the effects of cosmological parameters on the lensing signal by running a Markov chain in Paper I.  The only place where we vary cosmological parameters is in some of the extragalactic foreground tests to show that even extreme variations in the assumed cosmology do not affect the basic conclusion that the foregrounds are negligible.  The parameter sets we will use are as follows.  The ``fiducial'' (WMAP) cosmology is flat with $\\Omega_bh^2=0.0222$, $\\Omega_mh^2=0.1275$, $h=0.727$, $\\sigma_8=0.743$, and $n_s=0.948$. The ``high-$\\Omega_m$'' is an extreme model that maintains the angular diameter distance to the CMB and keeps $\\Omega_bh^2$, $\\Omega_mh^2$, $n_s$, and the primordial power spectrum the same but increases $\\Omega_m$ until there is no cosmological constant ($\\Omega_\\Lambda=0$, $\\Omega_m=1.33$, $h=0.31$).  The ``high-$\\sigma_8$'' model is the WMAP+SDSS vanilla model from Tegmark {\\em et~al.} \\cite{2004PhRvD..69j3501T} and has $\\Omega_bh^2=0.0232$, $\\Omega_mh^2=0.1454$, $h=0.695$, $\\sigma_8=0.917$, and $n_s=0.977$.  The last model would give a large tSZ effect since it depends strongly on the normalization.  Of these models we consider the fiducial and high-$\\sigma_8$ models to be viable, while the high-$\\Omega_m$ model is already strongly ruled out by supernovae, large-scale structure, and measurements of the Hubble constant \\cite{2001ApJ...553...47F, 2004ApJ...607..665R, 2004PhRvD..69j3501T, 2006A&A...447...31A} and is included only to make the point that the foreground analysis is insensitive to cosmology.  The outline of this paper is as follows: the theoretical predictions for lensing of the CMB are reviewed in Section~\\ref{sec:theory}. The WMAP and SDSS data, including the LSS samples used, are described in Section~\\ref{sec:data}.  The analysis methodology is presented in Section~\\ref{sec:analysis}. We present our results in Section~\\ref{sec:results} and systematic error estimates in Section~\\ref{sec:sys}. Extragalactic foregrounds deserve special consideration and occupy Section~\\ref{sec:exf}. We conclude in Section~\\ref{sec:conc}. The appendices cover the description of point source contamination in bispectrum language (Appendix~\\ref{app:cont}); the halo model description of the bispectrum (Appendix~\\ref{app:bs}); the kSZ and Rees-Sciama foregrounds (Appendix~\\ref{app:efg}); cross-correlations of different foreground components (Appendix~\\ref{app:corr}); and the weak lensing likelihood function (Appendix~\\ref{app:lf}). ",
        "conclusions": "\\label{sec:conc}  In this paper, we have measured the cross-correlation between three samples of galaxies (LRGs, quasars, and radio sources) and lensing of the CMB.  We find evidence for such a correlation at the $2.5\\sigma$ level, with an amplitude of $1.06\\pm 0.42$ times the expected signal for the WMAP cosmology. All sources of systematic error and foreground contamination are believed to be negligible compared to the statistical uncertainties.  Our measurement is consistent with the earlier nondetection by Hirata {\\em et~al.} \\cite{2004PhRvD..70j3501H} using LRGs and the more recent $3.4\\sigma$ result by Smith {\\em et~al.} \\cite{2007PhRvD..76d3510S} using radio sources.  The case for lensing of the CMB is strengthened by having analyses by two different groups (Smith {\\em et~al.} and us).  While the results both draw on WMAP and NVSS data, and hence are not independent, it is worth noting the very significant differences between the analysis procedures.  In the lensing reconstruction procedures, Smith {\\em et~al.} used ${\\bf C}^{-1}$ weighting of the WMAP data, whereas we force the reconstruction to use the same weight function $W_l$ at all frequencies and in regions with different noise levels in order to simplify the foreground analyses.  They also use autocorrelations of the same CMB map and subtract off the noise-induced biases, whereas we use cross-correlations which inherently have no noise bias. Our NVSS sample was constructed with a flux cut at 2.5$\\,$mJy whereas Smith {\\em et~al.} use the full catalog (except for masked regions). We also used cross-correlations with other galaxy samples to estimate the NVSS bias and redshift distribution (see Paper I), instead of fitting the auto-power spectrum; such cross-correlations are usually more robust, especially for a sample such as NVSS that shows large amounts of instrumental power.  While even large uncertainties in the redshift distribution do not alter the detection significance (in number of sigmas), without a solution to the redshift distribution problem there is no hope for CMB lensing in cross-correlation to ever become a useful cosmological probe.  We would also like to test for consistency between the two results.  Smith {\\em et~al.} found a cross-correlation of $1.15\\pm 0.34$ times the expected signal.  However, because their redshift distribution model is peaked at higher $z$, and they have a different fiducial cosmology (higher $\\sigma_8$), their predicted signal amplitude is larger than ours; fitting the Smith {\\em et~al.} cross-spectra $C_l^{g\\kappa}$ to our theoretical prediction gives an amplitude of $A=1.32\\pm 0.40$.  This compares well with our result of $A_{\\rm NVSS}=1.11\\pm 0.52$.  [The results are still not quite comparable because of the flux cut difference (2.5 versus 2.0 mJy), although we suspect this would not cause a huge difference because only 18\\%\\ of the NVSS catalog is at $<2.5\\,$mJy, and the bias and redshift distribution of radio sources is usually believed to change very slowly with flux \\cite{1998AJ....115.1693C}.  In particular to explain the entire difference in amplitude, the lensing cross-correlation signal $A$ from the faint $F<2.5\\,$mJy sources would have to be a factor of $\\sim 2$ higher than for the bright sources, which we consider unlikely.] The results were also arrived at independently: no revisions have been made to our central value since we became aware of the Smith {\\em et~al.} result (although our error bars were not finalized at that time because we had not completed analyzing our simulations).  An obvious question in the comparison is why the Smith {\\em et~al.} error bar is smaller than ours by a factor of 1.3.  Some of this is due to the 20\\%\\ higher number density in their Poisson-limited NVSS map due to their lower flux cut, which accounts for a factor of $\\sqrt{1.2}$.  The remaining difference -- a factor of 1.2 in standard deviation -- must be attributable to all other differences in the analysis, such as their use of Q band and auto-correlations between WMAP maps, different handling of point sources, and (going the other direction) Smith {\\em et~al.} inclusion of the point source and Galactic foreground systematics (which we find in this paper to be negligible).  Because of the different choices made in the analysis the final signal amplitudes end up being different and the Smith {\\em et~al.} central value is higher.  This must be a statistical fluctuation if both methods are unbiased, unless the $b\\,dN/dz$ is dramatically larger for the fainter sources. These sorts of fluctuations can matter if one is trying to claim a detection, e.g. our measured amplitude would only give a 2.8$\\sigma$ signal even with their errors versus their 3.4$\\sigma$. However, in this paper we take the view that what is ultimately relevant is the power of the method to disciminate between the cosmological models and the advantage of having a higher detected amplitude can quickly be turned around if there is an interesting class of models that predicts a larger lensing signal than the standard model, and which would therefore be more strongly ruled out if the measured amplitude is lower. It is for this reason that we put such an effort in determining redshift distribution weighted bias $b\\,dN/dz$ for the samples, because without it no cosmological interpretation is possible, since variations in cosmology and in $b\\,dN/dz$ will be degenerate.  Another difference in the two analyses is in the treatment of extragalactic foregrounds.  We split the foregrounds into 1, 2a, 2b, and 3-halo terms for each foreground component, and considered each of the many resulting terms separately.  Smith {\\em et~al.} did several tests, the most important being (i) a comparison of Q, V, and W band signals, and (ii) a point source bispectrum test which in our terminology accounts for the 1+2a halo terms. There are possible worries with these tests, e.g. in their implementation of (ii) the ``template'' bispectrum for the point sources has the same frequency dependence as the CMB.  Real point sources will have a different frequency dependence, however the template should still have some sensitivity to them and this is demonstrated explicitly by the Smith {\\em et~al.} simulations.  The method also does not consider 2b and 3-halo terms, and both tests can miss kSZ or RS contaminants; however as we have argued here these are negligible.  We also included the SDSS LRGs and quasars.  These together (without NVSS) give an amplitude of $0.99\\pm 0.56$ times the expected signal; this is a $1.8\\sigma$ result and is consistent with the fiducial cosmology. Thus in all cases the measured lensing amplitude is consistent with the WMAP cosmology, which can be counted as a success.  However it is not yet competitive with other cosmological probes. We have constructed a likelihood function (see Appendix~\\ref{app:lf}) and included it in the Markov chains of Paper I as a proof of principle, but it does not add much to the CMB+ISW constraints. The real challenge for lensing of the CMB is to move beyond the ``first detection'' stage to providing interesting cosmological constraints. Several improvements must be made in order to bring lensing of the CMB to the leading edge of cosmology: \\newcounter{issues} \\begin{list}{\\arabic{issues}. }{\\usecounter{issues}} \\item The most obvious issue is that CMB maps with higher resolution and lower noise than WMAP are required (one cannot improve on the sky coverage).  This will get better with the upcoming Planck satellite \\footnote{URL: http://planck.esa.int/science-e/www/area/index.cfm?fareaid=17} and the ground-based Atacama Cosmology Telescope \\footnote{URL: http://www.physics.princeton.edu/act/} and South Pole Telescope \\footnote{URL: http://spt.uchicago.edu/}. \\item Extragalactic foregrounds, while only a minor issue here, get worse as one uses the higher multipoles in the CMB. One possibility is the use of multifrequency information: at the higher frequencies probed by the future experiments, infrared point sources and tSZ have a distinctive spectral signature.  Another possibility is to use spatial information: a modified reconstruction procedure could be used to suppress the point source responsivity function $r_{ps}(l)$, which controls the 1- and 2a-halo contributions from point sources.  (It is however impossible to do a reconstruction that is insensitive to all possible forms of the 3-halo term.)  Yet another possibility, which is probably the best when the data become available, is to use polarization \\cite{2002ApJ...574..566H, 2003PhRvD..68h3002H}.  The point sources, tSZ, and kSZ effects are either observed or theoretically expected to be much weaker in polarization than in temperature \\cite{2000ApJ...530..133T, 2000ApJ...529...12H, 2007ApJS..170..335P}, and the RS effect is absent.  Also for randomly oriented point sources the 2b- and 3-halo contamination is exactly zero because two different sources can have no $QQ$ or $UU$ correlations, thus (aside from correlations of position angles of distinct sources) eliminating $r_{ps}(l)$ or using the multi-template bispectrum estimator of Smith {\\em et~al.} \\cite{2007PhRvD..76d3510S} would remove the point sources entirely. \\item The third issue for lensing in cross-correlation is the theoretical uncertainty in the galaxy-convergence cross-spectrum.  In this paper we have used linear biasing, $P_{g\\delta}(k)=b_gP_{\\delta\\delta}^{\\rm lin}(k)$, and thrown out small-scale information ($k>0.1h\\,$Mpc$^{-1}$) where linear biasing is not valid.  This must be improved in the future: especially for low-redshift mass tracers a large fraction of the information is in the nonlinear regime, and if one aims for constraints at the several percent level, simple ideas like linear biasing may not be adequate even at $k=0.1h\\,$Mpc$^{-1}$.  There is also uncertainty in the redshift distribution of the galaxies.  One obvious method is to do lensing of the CMB in autocorrelation (with the four-point function \\cite{2001PhRvD..64h3005H} or using the iterative approaches \\cite{2003PhRvD..68h3002H}), although this will be very demanding on control of instrumental systematics.  If one chooses to go the cross-correlation route, then when sufficiently high signal-to-noise ratio is available it will be desirable to use halo modeling, and to better calibrate the photometric redshifts for the large scale structure tracers with spectroscopic surveys.  The latter will of course also be valuable to any cosmic shear program using the galaxies as lensing sources. \\end{list}  If these challenges can be met, weak lensing of the CMB will make the transition from being a simple consistency check of the cosmological model to a routinely used cosmological probe."
    },
    "0801/0801.2194_arXiv.txt": {
        "abstract": "We report a comprehensive statistical analysis of the observational data of the cosmic evolution of supernova (SN) rate density, to derive constraints on cosmic star formation history and the nature of type Ia supernova (SN Ia) progenitor. We use all available information of magnitude, SN type, and redshift information of both type Ia and core-collapse (CC) SNe in GOODS and SDF, as well as SN Ia rate densities reported in the literature. Furthermore, we also add 157 SN candidates in the past Subaru/Suprime-Cam data that are newly reported here, to increase the statistics. We find that the current data set of SN rate density evolution already gives a meaningful constraint on the evolution of the cosmic star formation rate (SFR) at $z \\lesssim 1$, though strong constraints cannot be derived for the delay time distribution (DTD) of SNe Ia. We derive a constraint of $\\alpha \\sim $ 3--4 [the evolutionary index of SFR density $\\propto (1+z)^\\alpha$ at $z \\lesssim 1$] with an evidence for a significant evolution of mean extinction of CC SNe [$E(B-V) \\sim 0.5$ at $z \\sim 0.5$ compared with $\\sim 0.2$ at $z=0$], which does not change significantly within a reasonable range of various DTD models. This result is nicely consistent with the systematic trend of $\\alpha$ estimates based on galactic SFR indicators in different wavelengths (ultraviolet, H$\\alpha$, and infrared), indicating that there is a strong evolution in mean extinction of star forming regions in galaxies at relatively low redshift range of $z \\lesssim 0.5$.  These results are obtained by a method that is completely independent of galaxy surveys, and especially, there is no detection limit about the host galaxy luminosity in our analysis, giving a strong constraint on the star formation activity in high-$z$ dwarf galaxies or intergalactic space. ",
        "introduction": "In recent years a number of searches for high redshift supernovae (SNe) have been conducted. Although the primary purpose of most of these surveys is measurement of the cosmic expansion, these surveys also allowed measurements of the cosmic supernova rate density and its evolution \\citep{Pain02, Tonry03, Madgwick03, Gal-Yam04, Blanc04, Maoz04, Dahlen04, Cappellaro05, Barris06, Neill06, Poznanski07, Sharon07, Mannucci07, Kuznetsova07}. Studying these data should provide us with important information not only for the cosmic star formation history (CSFH) but also the still unknown progenitor of type Ia supernovae (SNe Ia). The progenitor of SNe Ia is believed to be a binary system including a white dwarf, and the SN Ia rate density evolution is a convolution of the cosmic star formation history and the delay time distribution (DTD) from star formation to SN Ia events.  DTD depends on the progenitor models, and hence to constrain DTD observationally is a useful approach to reveal the SN Ia progenitor (Madau et al. 1998; Yungelson \\& Livio 1998, 2000;  Dahlen \\& Fransson 1999; Gilliland et al. 1999; Gal-Yam \\& Maoz 2004; Strolger et al. 2004, 2005;  Barris et al. 2004; Oda \\& Totani 2005, hereafter OT05; Strigari et al. 2005).  However, it is not an easy task to actually extract useful constraints from the SN rate density evolution data. Previous studies \\citep{Maoz04, Strolger04, Forster06} mainly concentrated on the determination of DTD, using the rate density evolution of SNe Ia.  In such an analysis, sometimes CSFH models are assumed based on the observational estimates from high-$z$ galaxy surveys. However, as argued by \\citet{Forster06}, the constraint on DTD models sensitively depends on the assumed CSFH, and hence it is difficult to derive a robust constraint on DTD.  The primary purpose of this paper is to perform a comprehensive likelihood analysis using all available SN rate density evolution data in the literature, to derive constraints on DTD and/or CSFH.  After the GOODS high-$z$ supernova survey \\citep{Dahlen04, Strolger04}, whose data was used in \\citet{Strolger04} and \\citet{Forster06}, a number of observational estimates of SN rate density evolution have been published (mostly for SNe Ia, but some data also for CC SNe). Our approach is to derive constraints only by using SN rate data, without using information of CSFH from galaxy surveys. We will perform a simultaneous fit to both the SN Ia and CC SN rate density evolution data, surveying parameters of the CSFH model with a variety of DTD models. We will find that, though a strong constraint on DTD models cannot be derived even from all the available data so far, we can set interesting constraints on CSFH and evolution of mean dust extinction of CC SNe, which can be compared with those inferred from galaxy surveys.  Although there are a number of observational estimates on CSFH at a variety of redshifts from galaxy surveys, there is still a large uncertainty in the star formation rate (SFR) density estimated from galaxy observations, because of extinction, initial mass function, or extrapolation of luminosity functions to fainter magnitudes below the detection limits (see Hopkins 2004 and Hopkins \\& Beacom 2006, and reference therein). Therefore it is useful and important to derive constraints on CSFH from SNe independently of galaxy surveys.  In contrast to SFR density estimates by galaxies, detectability of SNe does not depend on the host galaxy brightness, and even intergalactic star formation activity can be probed by hostless SNe. Searches for $z \\sim 1$ SNe are typically performed at wavelength around the $i'$ and $z'$ band roughly corresponding to the rest-frame visual bands, and hence the effect of extinction by dust is expected to be smaller than the CSFH estimates based on the rest-frame UV emission of galaxies. It is not trivial that a unit mass of star formation always produces the same number of SNe, but it could evolve with redshift or physical properties of galaxies.  If a significant difference between CSFH inferred from galaxy surveys and that from SN surveys is found, it might indicate that the relation between star formation and supernova production is not as simple as normally assumed.  In addition to the available SN rate density data in the literature, we also utilize the photometric sample of SN candidates found in the past observations using Subaru/Sprime-Cam. This Subaru Supernova Survey (SSS) sample includes 157 supernova candidates, 61 out of which have clear offsets from the centers of host galaxies and hence they are most likely SNe. This data set is complementary to GOODS, SNLS and the IfA deep survey \\citep{Strolger04,Neill06,Barris06} in terms of the combination of the survey area and depth; the covered area of SSS, $1.4$ deg$^2$, is wider than the GOODS, and the SSS depth, $i' \\sim 26.0$, is deeper than the SNLS and the IfA deep survey. Though no SN type or redshift information is available for the SSS sample, we add this data set to our likelihood analysis to increase the statistics especially for CC SNe. Compared with SNe Ia, there are not many data of the rate density for CC SNe. Combined analysis of the SSS counts including all SNe and other data for SN Ia rate density evolution should give some constraints on the CC SN rate density evolution and hence CSFH.  The following are the plan of this paper. In \\S \\ref{sec:data} and \\S \\ref{sec:obs_res} we describe the SSS data set and analysis procedure of selecting SN candidates. Formulations of the comparison of the theoretical model and the observational data are given in \\S \\ref{sec:comp_method}. Constraints on the CSFH from our comprehensive parameter survey are derived in \\S \\ref{sec:results}. Conclusions are given in \\S \\ref{sec:summary}. Throughout this paper, the standard $\\Lambda$CDM universe is assumed with the following values of the cosmological parameters: $\\Omega_{\\rm M}$ = 0.3, $\\Omega_{\\rm \\Lambda}$ = 0.7, $h_{70} \\equiv H_{0}$/ (70 km s$^{-1}$ Mpc$^{-1}$) = 1. All magnitudes are given in the AB magnitude system. ",
        "conclusions": "\\label{sec:summary}  We performed a comprehensive likelihood analysis of almost all available data of the cosmic SN rate density evolution both for CC and type Ia supernovae, to get information on the CSFH and the DTD of SNe Ia.  We utilized the variability magnitude and redshift distribution of CC and Ia SNe of the GOODS and SDF-SNS, and other estimates of SN Ia rate density at various redshift in the literature.  Furthermore, we added photometrically found supernova candidates in the past imaging data of Subaru/Suprime-Cam (Subaru Supernova Survey, SSS) to increase the statistics. The analysis of the SSS data is newly reported here, including 157 SN candidates down to $i' \\sim 26.0$ in the total survey area of 1.4 deg$^2$.  61 of the 157 SSS candidates are associated with host galaxies with significant offsets from galaxy centers, and hence they are almost certainly supernovae. Though the type and redshift information is not available for SSS, the total SSS SN counts are useful to constrain the poorly known CC SN rate evolution, by a combination with the relatively well determined SN Ia rate evolution.  We have tested a variety of DTD models; some of them are based on the stellar evolution theory, and others are simple analytic functions often used in the literature. It is found that most of DTD models are consistent with the current data set, and hence we cannot set strong constraint on the type Ia SN progenitor.  On the other hand, this rather week dependence on DTD models is an advantage when one tries to constrain CSFH parameters.  It is required that $\\alpha$ (the SFR evolution index from $z=0$ to $\\sim 1$) is 3--4, with a considerable evolution of mean extinction of CC SNe, as $E(B-V) \\sim 0.2$ at local and $E(B-V) \\sim 0.5$ at $z \\sim 0.5$. Since we did not utilize any information from CSFH estimates by galaxy surveys, we can compare our result with those by galaxy surveys. Recent estimates based on UV luminosity are $\\alpha \\lesssim 2.5$, while those based on H$\\alpha$ or mid-infrared luminosity are close to our result, $\\alpha \\sim$ 3--4.  These are nicely consistent with our finding of the significant evolution of extinction for CC SNe, indicating a strong evolution of extinction of star formation activity in the universe even at $z \\sim$ 0--0.5.  The consistency between CSFH based on SFR in galaxies and SN rates is not trivial. Most indicators of galactic SFR trace the production of UV or ionizing photons and hence the formation of massive stars, which is the same as CC SNe. However, an evolution of IMF of massive stars could change the ratio of UV photon production to CC SN rate. Our result implies a roughly constant production efficiency of SNe per unit mass of star formation, and this would give some constraint on IMF evolution or metallicity effects. We have also demonstrated that, based on the counts of SN candidates without host galaxies, the contribution to the cosmic star formation activity from faint galaxies under detection limit or intergalactic star formation should not be significant. This is a clear advantage of CSFH constraint from supernova surveys, which cannot be obtained by CSFH studies based on galactic SFR estimates.  The authors would like to thank L. Greggio for providing the data of delay time distribution and useful discussions. We also appreciate useful comments from A. Gal-Yam, D. Maoz, D. Poznanski, and an anonymous referee.  A part of the SSS data was obtained as a part of the Supernova Cosmology Project (SCP), and we also thank the SCP members for useful discussions.  This work was supported by the 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. T.O. has been supported by JSPS Research Fellowships for Young Scientists. T.T., N.Y., and M.D. are supported by the JSPS - USA bilateral programme."
    },
    "0801/0801.3192_arXiv.txt": {
        "abstract": "{A good correlation has been found in star-forming galaxies, between the soft X-ray and the far infrared or radio luminosities. The soft X-ray emission in star-forming regions is driven by the heating of the diffuse interstellar medium, and by the mechanical energy released by stellar winds and supernova explosions, both directly linked to the strength of the star formation episode. } { We analyze the relation between the soft X-ray and far infrared luminosities as predicted by evolutionary population synthesis models, aiming first  to test the validity of  the soft X-ray luminosity as a star formation rate estimator, using  the already known calibration of the FIR luminosity as a proxy, and second to propose a calibration based on the predictions of evolutionary synthesis models.} { We have computed the soft X-ray and far infrared luminosities expected for a massive starburst as a function of { evolutionary state, the efficiency  of the conversion of mechanical energy into soft X-ray luminosity, the star formation history (instantaneous or extended) and dust abundance,}{} and we have compared these predictions with observational values for 62 star-forming galaxies taken from the literature. } {The observational \\lsxf\\ ratios are consistent with the model predictions under realistic assumptions (young starbursts, and efficiency in the re-processing of mechanical energy of a few percent), confirming the correlation between the diffuse soft X-ray emission and the star formation episode.  } {The soft X-ray emission of the diffuse, extended gas surrounding massive star-forming regions, can be used as a star formation rate {tracer}.  The empirical calibrations presented in the literature are supported by the predictions of evolutionary synthesis models, and by the analysis of a larger number of star-forming galaxies The calibrations are, however, biased towards galaxies dominated by relatively unevolved starbursts. } ",
        "introduction": "The onset of massive star formation episodes in galaxies drives their observational properties in almost any wavelength range. The UV and optical become dominated by the continuum of massive, hot and young stars, as well as by the presence of nebular emission lines. After a few Myr of evolution, red supergiant stars contribute to most of the near infrared emission. The heating of interstellar dust particles by the powerful UV photons induces the thermal re-emission of large amounts of energy in the { mid and far infrared} domain. The injection of ionizing photons into the surrounding gas generates thermal radio emission, which is replaced by non-thermal emission as the ionizing power of the burst declines and the more massive stars begin to explode as supernovae. The direct relation between the strength of the star formation episode and the intensity of the different observable parameters has allowed a number of star formation rate calibrators to be defined, such as UV continuum, emission lines intensity, far infrared or radio luminosities \\citep{Kennicutt98,Rosa02,Bell2003}. These calibrators have proven to be invaluable for statistical studies of the star formation history of the Universe.  Star-forming regions are also the source of  conspicuous X-ray emission, generated by individual stars, by the injection of large amounts of mechanical energy heating the interstellar medium, by supernova remnants, and by binary systems transferring mass to a compact primary \\citep{Cervino02,Persic02}. All of these individual components are in principle  directly linked to the strength of the star formation episode, so that the X-ray luminosity could also be used as an estimator of star formation rates (\\sfr).  \\begin{figure*} \\begin{center} \\includegraphics[width=8 cm,bb=5 39 786 521 dpi,clip=true]{8398fi1a.eps} \\hspace{1.0 cm} \\includegraphics[width=8 cm,bb=17 39 786 527 dpi,clip=true]{8398fi1b.eps} \\end{center} \\caption{Evolution of \\lfir\\ (left) and \\lsx\\ (right) for IB (dashed line, right axis) and EB (solid line, left axis) models. IB models predictions are shown normalized to $1$~\\msun\\ of gas transformed into stars, while for EB the luminosities are scaled  to \\sfr\\ $= 1$ \\msun\\ yr$^{-1}$. \\lfir\\ has been plotted (from top to bottom) for \\ebv\\ $= 1.0$, $0.5$ and $0.1$. \\lsx\\ has been computed (from top to bottom) for \\xeff\\ $= 0.1$, $0.05$ and $0.01$.  } \\label{lumin} \\end{figure*}    Several authors have in recent years discussed the feasibility of using the X-ray luminosity as an \\sfr\\ estimator. {\\citet{Fabbiano02} already concluded from the analysis of 234 S0/a-Irr galaxies observed with {\\em Einstein}, that the correlation they found between the X-ray and the FIR luminosities in Sc-Irr galaxies was due to the young stellar populations in these objects.} \\citet{Ranalli03} proposed an empirical calibration of the soft ($0.5$--$2.0$ keV) and hard ($2$--$10$ keV) X-ray luminosities, based on their correlation with the far infrared (FIR) and radio luminosities, and using the known calibrations of these parameters as proxies. \\citet{Grimm03} studied the correlation between the number of high-mass X-ray binaries (HMXB) and the \\sfr, deriving different calibrations of the hard X-ray luminosity for low and high star formation rates.  \\citet{Persic04} obtained a calibration of the hard X-ray luminosity as a \\sfr\\ estimator by assuming that most of the emission in this range is associated with HMXB, and scaling from the number of HMXB to the \\sfr\\ of our Galaxy. \\citet{Gilfanov2004} confirmed the {calibration} of the hard X-ray luminosity associated with HMXB, using slightly different slopes at high and low star formation rates.  In a recent paper, \\citet{Persic07} found indeed that the collective hard X-ray emission of young point sources correlates linearly with the star formation rate derived from the far infrared luminosity. \\citet{Strickland04b} demonstrated that the luminosity of diffuse X-ray emission in star-forming galaxies is directly proportional to the rate of mechanical energy injection from the young, massive stars into the interstellar medium of the host galaxies. A similar result was found by \\citet{Grimes05} from the analysis of  a sample of starburst galaxies of different types  (from dwarf starbursts to ultraluminous infrared galaxies), which concluded that the mechanism producing the diffuse X-ray emission in the different types of starbursts was powered by the mechanical energy injected by stellar winds and supernovae into the surrounding medium.  Recently \\citet{Rosa07} confirmed the reliability of the soft X-ray luminosity as an \\sfr\\ estimator from the analysis of a sample of star-forming galaxies in the {\\em Chandra Deep Field South} at redshifts $ z = 0.01 - 0.67$, using the UV continuum luminosity from {\\em GALEX} as a proxy.   While hard X-ray emission from star-forming regions may be dominated by binary systems, diffuse soft X-ray emission is generated by reprocessed, mechanical energy from stellar winds and supernovae explosions. This mechanical energy is related to the strength of the burst of star-formation, and can be calculated using evolutionary population synthesis models. In this paper, we analyze the correlation between the soft X-ray and FIR luminosities in star-forming regions, both predicted by evolutionary synthesis models. We study the \\lsxf\\ ratio as a function of the {evolutionary state, the efficiency  of the conversion of mechanical energy into soft X-ray luminosity, the star formation history (instantaneous or extended), and the dust abundance,} and compare the computed values with observations taken from the literature. { Our objective is to derive a calibration of \\lsx\\ as a tracer of the star formation rate, based on the predictions of evolutionary synthesis models, and to  test the validity of the empirical calibration proposed by \\citet{Ranalli03}.  }  In Sect.~2 we describe the evolutionary synthesis models that we use in the present study, in Sect.~3 we present the observational data taken from the literature, and in Sect.~4 we discuss the predictions and the comparison with the observational values.  {\\ Throughout this work we have assumed $H_{\\rm0} = 73$ km s$^{-1}$ Mpc$ ^{-1}$ to convert fluxes into luminosities.}  \\begin{figure*} \\begin{center} \\includegraphics[width=8cm,bb=20 39 713 527 dpi,clip=true]{8398fi2a.eps} \\hspace{1.0 cm} \\includegraphics[width=8cm,bb=20 39 713 527 dpi,clip=true]{8398fi2b.eps} \\end{center} \\caption{Evolution of the \\lsxf\\ ratio computed for \\xeff\\ $= 0.01, 0.1$ and \\ebv\\ $= 0.1,  1.0$. Left: predictions for IB models; right: predictions for EB models.  } \\label{ratio} \\end{figure*} ",
        "conclusions": "We have compared the \\lsxf\\ values measured in a sample of  62 star-forming galaxies with the predictions by our evolutionary synthesis models,  aiming to analyze the validity of  semiempirical and theoretical calibrations of \\lsx\\ as a star formation rate estimator. The main results can be summarized as follows:  \\begin{enumerate}  \\item The \\lsxf\\ ratios are strongly dependent on the efficiency in the conversion of the mechanical energy released by the young massive starburst into soft X-ray luminosity, by interaction of the stellar winds and supernova ejecta with the surrounding interstellar medium. From theoretical predictions and observational data, we expect an \\xeff\\ value of few percent in starburst galaxies.  \\item \\lsxf\\ is also dependent on the evolutionary status of the star formation episode. It increases rapidly with time during the first $5$ Myr of evolution of a massive starburst, and shows a slower increase afterwards. After a (nearly) instantaneous burst of star formation, \\lsx\\ decreases slower than \\lfir, as long as there remains a population of massive stars able to collapse as supernovae (up to around 35~Myr).  \\item When star formation proceeds at a nearly constant rate during extended periods of time, the \\lsxf\\ ratio is expected to stabilize after around $40$ Myr, when the number of massive stars that produce supernovae has reached an equilibrium between death and birth of new stars.  { \\item The \\lsxf\\ values measured for the sample of star-forming galaxies are consistent with the predictions by the models under realistic conditions: relatively young and unevolved  star formation  episodes and \\xeff\\ values within $1$--$10$\\%.   \\item A calibration of \\lsx\\ as a star formation rate estimator, based on the predictions of evolutionary synthesis models, has been derived. The calibration proposed by  \\citet{Ranalli03} is consistent with the predictions for relatively unevolved, time-extended bursts of star formation. } \\end{enumerate}"
    },
    "0801/0801.1474_arXiv.txt": {
        "abstract": "{} { We present the characteristics and some early scientific results of the first instrument at the Large Binocular Telescope (LBT), the Large Binocular Camera (LBC). Each LBT telescope unit will be equipped with similar prime focus cameras.  The blue channel is optimized for imaging in the UV-B bands and the red channel for imaging in the VRIz bands.  The corrected field-of-view of each camera is approximately 30 arcminutes in diameter, and the chip area is equivalent to a $23\\times 23$ arcmin$^2$ field. In this paper we also present the commissioning results of the blue channel. } { The scientific and technical performance of the blue channel was assessed by measurement of the astrometric distortion, flat fielding, ghosts, and photometric calibrations. These measurements were then used as input to a data reduction pipeline applied to science commissioning data. } { The measurements completed during commissioning show that the technical performance of the blue channel is in agreement with original expectations. Since the red camera is very similar to the blue one we expect similar performance from the commissioning that will be performed in the following  months in binocular configuration. Using deep UV image, acquired during the commissioning of the blue camera, we derived faint UV galaxy-counts in a $\\sim 500$ sq. arcmin. sky area to U(Vega)$=26.5$. These galaxy counts imply that the blue camera is the most powerful UV imager presently available and in the near future in terms of depth and extent of the field-of-view. We emphasize the potential of the blue camera to increase the robustness of the UGR multicolour selection of Lyman break galaxies at redshift $z\\sim 3$. } {} ",
        "introduction": "The sensitivity of an optical system depends on a combination of the aperture and field-of-view (FoV).  The imaging capabilities of existing or planned facilities are often limited by practical constraints. When the large collecting area of a telescope allows detection of faint sources, the field-of-view is typically less than $7\\times 7$ square arcminutes, and the UV sensitivity is low. Alternatively, wide-field imaging cameras onboard smaller telescopes are optimized to target brighter sources over a larger field-of-view (i.e. MegaCam at CFHT, \\cite{boulade}), and are unable to detect sources of faint magnitudes ($\\sim 28$) in particular in the UV.  For these reasons an imager with a large FoV at an 8m class telescope is of fundamental importance to address the presently still open problems in stellar and extragalactic astronomy. The best example is the prime focus camera at the 8m Subaru telescope, Suprimecam (\\cite{subaru}). This imager is fast and has a FoV of $34\\times 27$ arcmin$^2$. Common science projects that have utilized this imager to date are the search of very high redshift galaxies, the study of the formation and evolution of galaxies, the investigation of the structure of the Universe, and the search for Kuiper Belt objects in the Solar system. The optical corrector cannot, however, simultaneously correct radiation of all wavelengths from UV to I-band. Due to this practical limitation, and to its low sensitivity in the blue band, Suprimecam does not provide imaging in the UV.  At the end of the 1990s, it became clear that the binocular configuration of the Large Binocular Telescope (LBT) (\\cite{hill}), coupled with its mechanical design, provided a unique opportunity to justify a double prime focus camera capable of studying the widest-possible wavelength range from the UV down to the NIR H-band.  The Large Binocular Camera (LBC, \\cite{ragazzoni,pedik,pedik04,ragazzoni06}) is a wide FoV instrument at the prime focus of the twin 8.4 meter Large Binocular Telescope (LBT).  The LBT uses two 8.4-meter diameter honeycomb primary mirrors mounted side-by-side to produce a collecting area of 110 square meters equivalent to an 11.8-meter circular aperture.  A unique feature of the LBT is that the light from the two primary mirrors can be combined optically in the center of the telescope to produce phased array imaging of an extended field. This requires minimal path length compensations, thus making interferometry easier than in completely independent telescopes.  The requirement for an instrument such as LBC has been identified by several high-profile scientific programs that call for an increase in FoV and high-UV/IR sensitivity for deep imaging.  These attributes are essential to programs studying a large FoV, to significant depth, over a wide spectral range, and can only be provided by an imager mounted at the prime focus of an 8m-class telescope.  In Section 2 we provide a description of the two LBC cameras, while in Section 3 we detail the technical performance of LBC-Blue during commissioning observations in 2006. In Section 4 we analyze in detail the case for an UV deep imaging survey in an extragalactic field, and compare results with those obtained using different instruments and telescopes. We present our conclusions in Section 5. ",
        "conclusions": "In this paper we have described the first instrument at the LBT telescope, the prime focus large binocular camera (LBC). The instrument has a binocular configuration with two channels, the blue channel with a good overall efficiency in the UV band and the red channel with good efficiency in the V-z bands.  We have also shown the technical characteristics of the blue channel derived from the commissioning data of LBC-Blue.  \\begin{itemize} \\item The corrected optical field is $\\sim$ 30 arcmin of diameter with $<5$\\% loss of energy within $\\sim 25$ arcmin. The total throughput of the optical corrector is 84\\%. \\item The optical distortion is always $<1.75$\\% even at the edge of the field and it is removed with a specific SW package. \\item The optical quality ensures an energy concentration of 80\\% within a pixel of $\\simeq 0.2254$ arcsec in the overall corrected field.  The active control of the optical quality can provide images as sharp as FWHM=0.5 arcsec even in the UV band. \\item Ghost images produced by the whole optical system are always negligible when glass filters are used (Bessel U,B,V). The total intensity of the primary brightest ghost is $\\sim 2.8$\\% when the wide interference UV filter is used. The primary ghost is centered on the original source position. Ghost images produced by the electronics cross-talk are as small as $3\\times 10^{ -5}$ and can be easily removed during data reduction. \\end{itemize}  We have also described some scientific observations planned for the commissioning to assess the performance of LBC-Blue.  \\begin{itemize} \\item We have obtained very deep UV galaxy counts in a deep pointing with a total exposure time of 3h reaching U(Vega)=26.5, after correction for incompleteness.  The wide magnitude interval ($U(Vega)=19-26.5$) in our galaxy counts allowed a direct evaluation of the shape of the UV counts which shows a break in slope at about $U(Vega)=23.2$ with a change in slope from 0.62 to 0.22 at the faint end. \\item The same LBC area includes a quasar field where extensive study of Lyman break galaxies at redshift $z\\sim 3$ is available. We have reproduced with our UGR filter set the well known multicolour selection of LBGs showing that the robustness of the UV dropout method for the selection of star-forming galaxies at $z\\sim 3$ increases when very deep UV images can be used as obtained by LBC at an 8m class telescope like LBT. \\end{itemize}"
    },
    "0801/0801.1368_arXiv.txt": {
        "abstract": "We have investigated the formation of close-in extrasolar giant planets through a coupling effect of mutual scattering, Kozai mechanism, and tidal circularization, by orbital integrations. Close-in gas giants would have been originally formed at several AU's beyond the ice lines in protoplanetary disks and migrated close to their host stars. Although type II migration due to planet-disk interactions may be a major channel for the migration, we show that this scattering process would also give a non-negligible contribution. We have carried out orbital integrations of three planets with Jupiter-mass, directly including the effect of tidal circularization. We have found that in about 30\\% runs close-in planets are formed, which is much higher than suggested by previous studies. Three-planet orbit crossing usually results in one or two planets ejection. The tidal circularization often occurs during the three-planet orbit crossing, but previous studies have monitored only the final stage after the ejection, significantly underestimating the formation probability. We have found that Kozai mechanism by outer planets is responsible for the formation of close-in planets. During the three-planet orbital crossing, the Kozai excitation is repeated and the eccentricity is often increased secularly to values close enough to unity for tidal circularization to transform the inner planet to a close-in planet. Since a moderate eccentricity can remain for the close-in planet, this mechanism may account for the observed close-in planets with moderate eccentricities and without nearby secondary planets. Since these planets also remain a broad range of orbital inclinations (even retrograde ones), the contribution of this process would be clarified by more observations of Rossiter-McLaughlin effects for transiting planets. ",
        "introduction": "\\label{sec:intro}  More than 250 planets have been detected around both solar and non-solar type stars. Recent development of radial velocity techniques and accumulation of observations have revealed detailed orbital distribution of close-in planets. Figure \\ref{fig:exso}$a$ shows the distribution of semi-major axis and eccentricity for 236 planets \\footnote{http://exoplanet.eu/catalog-RV.php} discovered around solar type stars by the radial velocity techniques. The dotted line shows a pericenter distance $q=0.05$ AU. At $q \\la 0.05$ AU, many close-in planets have small eccentricities, which are accounted for by the circularization due to the tidal dissipation of energy within the planetary envelopes (Rasio \\& Ford 1996). The semi-major axis distribution of discovered extrasolar planets shows double peaks around 0.05 AU and 1 AU (e.g., Marcy et al. 2005; Jones et al. 2005). The location of the outer peak is determined by the observational limits, but the inner peak at around 0.05 AU, in other words shortage of planets between 0.06--0.8 AU, would be real. Since the close-in planets have relatively small eccentricities, the peak around 0.05 AU appears as a pile up at $\\log q/(1{\\rm AU})=-1.3$ also in pericenter distribution (Fig.~\\ref{fig:exso}$b$). The close-in planets are also found in the multi-planetary systems. Up to now, 23 multi-planet systems are reported with known planetary eccentricities (Fig.~\\ref{fig:multi}).  The discovered close-in planets would have been formed at large distances beyond the ice line and migrated to shorter-period orbits (e.g., Ida \\& Lin 2004a, 2004b). A promising mechanism for the orbital migration is ``type II'' migration (e.g., Lin, Bodenheimer, \\& Richardson 1996). This is the migration with disk accretion due to gap opening around the planetary orbit caused by gravitational interaction between the planet and the protoplanetary gas disk. This mechanism can account for not only the pile-up of the close-in planets but also planetary pairs locked in mean-motion resonances. The migrating planets may have stalled as they enter the magnetospheric cavity in their nascent disks or interact tidally with their host stars. However, several close-in ($a \\la 0.1$AU) planets such as HAT-P-2, HD118203b, and HD162020b have moderate eccentricities ($\\ga 0.3$), but are not accompanied by nearby secondary large planets. HD17156b with $a \\simeq 0.15$AU may also belong to this group, since it has very large eccentricity ($\\simeq 0.67$) and $q$ as small as 0.05AU. It may be difficult for the planet-disk interaction alone to excite the eccentricities up to these levels. One possible mechanism for the excitation to the moderate eccentricity for close-in planets in multiple-planet systems is the passage of secular resonances during the epoch of disk depletion (Nagasawa, Tanaka, \\& Ida 2000; Nagasawa, Lin, \\& Ida 2003; Nagasawa \\& Lin 2005). Nagasawa and Lin (2005) showed that the relativity effect and the gravity of (distant) secondary planets excite the eccentricity of close-in planet orbiting around a slowly rotating host star.  Although a main channel to form close-in planets may be type II migration, inward scatterings by other giant planets can be another channel, in particular for the close-in planets with moderate eccentricities. In relatively massive and/or metal-rich disks, multiple gas giants are likely to form (e.g., Kokubo \\& Ida 2002; Ida \\& Lin 2004a, 2004b, 2008). On longer timescales than formation timescales, the multiple gas giant systems can be orbitally unstable (Lin \\& Ida 1997; Marzari \\& Weidenschilling 2002). In many cases, a result of orbital crossing is that one of the planets is ejected and the other planets remain in well separated eccentric orbits (``Jumping Jupiters'' model; e.g., Rasio \\& Ford 1996; Weidenschilling \\& Marzari 1996; Lin \\& Ida 1997). Recent studies of the Jumping Jupiter model show pretty good coincidence with the observational eccentricity distribution (Marzari \\& Weidenschilling 2002; Chatterjee, Ford, \\& Rasio 2007;  Ford \\& Rasio 2007).  With the sufficiently small pericenter distance ($\\la 0.05$ AU), the planetary orbit can be circularized into the close-in orbit by tidal effects from the star (Rasio \\& Ford 1996). In this paper, the process of planet-planet scattering followed by the tidal circularization will be referred to as \"scattering model.\" To send a planet into the inner orbit, the other planets have to lose their orbital energy. Since inward scattering of a lighter planet is more energy saving, the systems which experienced the planet-planet scattering tend to have smaller planets inside. We might see the tendency in Figure \\ref{fig:multi}, but it is also true that radial velocity technique tends to detect heavier planets in outer regions. Note that since outer planets have negligible orbital energy, the energy conservation cannot restrict the semi-major axes of the outer planets and the outer planets may be located to the regions far beyond the radial velocity technique limit (Marzari \\& Weidenschilling 2002).  Observation of Rossiter-McLaughlin effect for 5 transiting close-in planets including HAT-P-2 with eccentricity $e \\simeq 0.52$ show that their orbital planes are almost aligned with the stellar spin axes (Narita et al. 2007; Winn et al. 2005, 2006, 2007; Wolf et al. 2007), which may suggest that the scattering mechanism is not a main channel for formation of close-in planets, because the close scattering may excite orbital inclination as well as eccentricity (note that as shown in \\S 3.4, the orbital inclinations of close-in planets formed by the scattering model are not necessarily high). Since HD17156b with $e \\simeq 0.67$ is also transiting, the observation of Rossiter-McLaughlin effect will be important.  More serious problem of the scattering model may be the probability for pericenter distances to become small enough (in other words, for eccentricities to become close enough to unity (e.g., $e \\ga 0.98$)) to allow tidal circularization. In a system with only two giant planets, the energy conservation law keeps the ratio of semi-major axes close to the initial value (see \\S \\ref{subsec:ppscat}). In such system, the probability for final planets with a large ratio of semi-major axes and very large eccentricity is $\\la 3$\\% (Ford, Havlickova, \\& Rasio 2001). In a system with more giant planets, the situation is slightly improved, especially when planets have non-zero inclinations. The close-in planets are very rare just after the scattering, but when they take into account the longer orbital evolution, the probability of the candidates for close-in planets increases to $\\sim$10 \\% (Weidenschilling \\& Marzari 1996; Marzari \\& Weidenschilling 2002; Chatterjee et al. 2007). The contributed planets are planets in Kozai state (see \\S 3.5).  Here we revisit the scattering model through orbital integration of three giant planets with direct inclusion of the tidal circularization effects and analytical argument of Kozai effect. We find that the probability for formation of close-in planets is remarkably increased to $\\sim$30\\%. Although the previous studies were concerned with orbital states after stabilization by ejection of some planet(s), we find that the orbital circularization occurs through ``Kozai migration'' (Wu 2003; Wu \\& Murray 2003) caused by outer planets, mostly during three-planet orbital crossing, after a tentative separation of the inner planet and the outer planets.  In the next section, we describe orbital instability of a planetary system (\\S \\ref{subsec:ppscat}) and the orbital evolution that becomes important after the system enters stable state: tidal circularization (\\S \\ref{subsec:tide}) and Kozai mechanism (\\S \\ref{subsec:kozai}). Following description of its models and assumptions, we present results of numerical simulations in \\S \\ref{sec:numerical}. We show typical outcome of the scattering (\\S \\ref{subsec:resultscat}), how the planets are circularized into short-period orbits (\\S \\ref{subsec:resulteach}), and its probability (\\S \\ref{subsec:resultcir}). In \\S \\ref{sec:analy} we present analytical arguments. We summarize our results in \\S \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  ``Standard'' model for formation of close-in giant planets is that gas giants may be originally formed at several AU's beyond ice lines and migrate to the vicinity of the star. The most referred migration mechanism is type II migration. Here, we have investigated another channel to move planets to the stellar vicinity, ``scattering'' model (Rasio \\& Ford 1996), which is a coupled process of planet-planet scattering and tidal circularization. If orbital eccentricity is pumped up to values close to unity, because of the proximity of the pericenter the eccentricity and the semimajor axis of the planet are damped by the tidal effect from the star, almost keeping the pericenter distance to form a close-in planet with relatively small eccentricity. We newly found that Kozai mechanism is also one of key processes in the model.  We have carried out orbital integration of three planets of Jupiter-mass, including the tidal damping in the integration. We follow the simulation setup by Marzari and Weidenschilling (2002). We found that in $\\sim$30\\% runs, close-in planets are formed, which is much higher probability than suggested by previous studies (Weidenschilling and Marzari 1997; Marzari and Weidenschilling 2002; Chatterjee et al. 2007), because in many cases, the circularization occurs during three-planet scattering, which was not monitored by previous studies. The three-planet scattering is usually terminated by ejection of one planet and the system enters a stable state. Previous studies were concerned only with the stable state.  Since the tidal damping timescale is usually longer than that of orbital change in the chaotic stage of three-planet interaction, the tidal circularization does not occur during the chaotic stage, even if the pericenter distance happens to be small enough. But, we have found that when one planet is scattered inward, it often becomes separated from other outer planets. As long as the tidal dissipation is negligible, the isolation is only tentative, but its duration is long enough for the tidal circularization if the pericenter distance is small enough ($\\la 0.02$--0.04 AU for Jupiter-mass planets). We also found out that the isolated planet usually enters Kozai state. Even if the eccentricity of the planet is not excited enough by the close scattering that injected the planet to an inner orbit, the eccentricity can secularly increase to values close to unity during Kozai cycle, in particular, when the planet has an inclined orbit. Although the probability that the eccentricity takes values between 0.98 and 1.0 is quite low just after the scattering, that probability is significantly enhanced to 30\\% level during the Kozai cycle. Even if the Kozai cycle has relatively large angular momentum and the eccentricity does not take such high value, the outer planets eventually destroy the quasi-stable state and have the inner planet enter another quasi-stable state, in other words, another Kozai state. The new Kozai state can have relatively small angular momentum and raise the eccentricity high enough for the tidal circularization. The repeated Kozai mechanism enhances the probability for the tidal circularization to occur. Thereby, we have found much higher probability of formation of close-in planets than previously expected. Therefore, the scattering would contribute to formation of close-in planets, although a main channel may still be through type II migration.  Non-negligible number of close-in planets without nearby secondary planets but with eccentricities $\\ga 0.3$ has been discovered. It is not easy for type II migration model to account for the high eccentricities, but the scattering model could account for them. As shown in sections 3.3 and 3.4, the tidal circularization can slow down by increase in $q$ before full circularization, leaving moderate eccentricities. The discovered close-in eccentric planets are generally massive (having masses larger than Jupiter-mass) among the close-in planets (Figure \\ref{fig:exso}a), so their circularization timescales are much shorter than those shown in Figure \\ref{fig:timescale}. Therefore, the results in our Sets T could be consistent with these discovered planets.  Our results also suggest that the close-in planets have a broadly spread inclination distribution, including retrograde rotation. More observations of the Rossiter-McLaughlin effect for the transiting close-in planets, in particular for those with relatively high eccentricity such as HD17156b, will impose constraints on the contribution of the scattering model to close-in planets.  The tidally dissipated energy is $\\sim G M_{\\ast} m/2a$ in total, where $a$ is final semi-major axis of the circularized close-in planet and $m$ is its planet mass. Since this energy is much larger than binding energy of the planet, it could be heat source for inflated close-in planets such as HD209458b, OGLE-TR-10, OGLE-TR-56, HAT-P-1b or WASP-1b, if the tidal circularization occurred relatively recently (orbital crossing followed by the circularization can start after long time after the formation of giant planets).  Note that the tidal circularization model we used may not be sufficiently relevant. We did many assumptions for the usage of equations (\\ref{eq:ang}) and (\\ref{eq:eng}). The model is only valid for fully convective planets in highly eccentric orbits, since it consider only $\\ell=2$ f-mode and use impulse approximation. We apply the model even after the orbits are significantly circularized, which may overestimate effect of tidal dissipation. In the multiple encounters to the star, the energy and angular momentum can either increase or decrease (e.g., Press \\& Teukolsky 1997; Mardling 1995a, 1995b), but we include the tidal effect as it always gives decreasing effect. Because of these simplifications, our simulation may be a limiting case that the tidal force works most effectively, although inclusion of g-modes might strengthen the tidal effect. Our purpose of this paper is not a study of the detailed tidal damping process but to show another available channel to the close-in planets. We also need to take into account the relativity precession and J2 potential of the central star which prevent the Kozai mechanism, as well as long-term tidal disruption process (e.g., Gu, Lin, \\& Bodenheimer 2003). These issues are left for further studies."
    },
    "0801/0801.0617.txt": {
        "abstract": "In 1996, a major radio flux-density outburst occured in the broad-line radio galaxy 3C\\,111. It was followed by a particularly bright plasma ejection associated with a superluminal jet component, which has shaped the parsec-scale structure of 3C\\,111 for almost a decade. Here, we present results from 18 epochs of Very Long Baseline Array (VLBA) observations conducted since 1995 as part of the VLBA 2\\,cm Survey and MOJAVE monitoring programs. This major event allows us to study a variety of processes associated with outbursts of radio-loud AGN in much greater detail than has been possible in other cases: the primary perturbation gives rise to the formation of a leading and a following component, which are interpreted as a forward and a backward-shock. Both components evolve in characteristically different ways and allow us to draw conclusions about the work flow of jet-production events; the expansion, acceleration and recollimation of the ejected jet plasma in an environment with steep pressure and density gradients are revealed; trailing components are formed in the wake of the primary perturbation possibly as a result of coupling to Kelvin-Helmholtz instability pinching modes from the interaction of the jet with the external medium. The interaction of the jet with its ambient medium is further described by the linear-polarization signature of jet components traveling along the jet and passing a region of steep pressure/density gradients. ",
        "introduction": "Direct evidence for the existence of bulk relativistic outflows along the jets in blazars and other radio-loud active galactic nuclei (AGN) comes from Very-Long-Baseline Interferometry (VLBI) observations. The first evidence for apparently superluminal structural changes was found from changes in the fringe visibility curves of 3C\\,279 and 3C\\,273 \\citep{Whi71,Coh71}. Subsequent higher-quality VLBI observations \\citep[see, e.g., compilation by ][and references therein]{Ver94} have established the ``core-jet'' type  milliarcsecond-scale structure of compact extragalactic jets: the core being a bright and unresolved flat-spectrum component at the end of a linear structure, and the jet being composed out of individual steep-spectrum components or ``knots''. The knots frequently move away from the core with apparent velocities exceeding the speed of light. Monitoring observations of large source samples \\citep{Ver94,Jor01,Kel04,Pin07} have provided important statistical tools for probing relativistic beaming and the intrinsic properties of extragalactic radio jets \\citep{Coh07}, their intrinsic brightness temperatures \\citep{Hom06}, or their Lorentz factor distribution \\citep{Kel04} and luminosity function \\citep{Car07}.   The relativistic-jet model \\citep[e.g.,][]{Bla79} has become the de-facto paradigm in multiwavelength research on blazars and other AGN, but VLBI observations have demonstrated that the basic concept of ballistically-moving isolated jet knots is clearly oversimplified: jet curvature \\citep[e.g.,][]{Ver94}, stationary components \\citep[e.g.,][]{Jor01}, and non-radial and accelerated motions \\citep[e.g.,][]{Kel04}, are found to be common features of relativistic jets. Within individual jets, there are often characteristic velocities suggesting the presence of an underlying continuous jet flow, but the ``components'' themselves most likely represent patterns moving  at a different speed than the underlying flow, e.g., as hydrodynamically propagating shocks \\citep{Mar85,Hug85}.  Recent years have brought major improvements in numerical simulations of relativistic jets \\citep[see, e.g.,][for a review]{Gom05}. It is now possible to simulate three-dimensional relativistic jets \\citep[e.g.,][]{Alo03} and to compute the relativistic processes \\citep[e.g.,][]{Gom97} that transfer hydrodynamic results into observed brightness distributions (e.g., relativistic light abberation and light travel time delays). In particular, interactions between strong perturbations or shocks with the underlying jet flow and the jet-ambient medium can be simulated \\citep{Agu01}. With these new techniques, it is now possible to compare the generation, propagation and evolution of emission features in simulated and observed relativistic jets.  The nearby (z$=$0.049)\\footnote{Assuming $H_0=71$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\rm M} = 0.3$, $\\Omega_\\Lambda = 0.7$ ($1{\\rm \\,mas} = 1.0{\\rm \\,pc}$).} broad-line radio galaxy 3C\\,111 (PKS B\\,0415+379) shows a classical FR\\,II morphology on kiloparsec-scales spanning more than 200$^{\\prime\\prime}$ with a highly collimated jet connecting the central core and the northeastern lobe in position angle 63$^\\circ$ while no counterjet is observed towards the southwestern lobe \\citep{Lin84}. This asymmetry is usually explained via relativistic boosting of the jet and de-boosting of the counter-jet. 3C\\,111 exhibits the brightest compact radio core at cm/mm wavelengths of all FR\\,II radio galaxies, a blazar-like spectral energy distribution \\citep{Sgu05}, and it was one of the first (and only) radio galaxies in which superluminal motion was detected \\citep{Goe87,Pre88}. Moreover, the (sub-) parsec scale jet of 3C 111 is intimately related to its high-energy emission: \\citet{Mar06} reports a disk-jet connection, similar to the well-established one in 3C 120 \\citep{Mar02}, in the sense that dips in the X-ray light curve indicate accretion events which are followed by VLBI jet component ejections. Recently, R.C.\\,Hartman \\& M. Kadler (in prep.) showed that the gamma-ray source 3EG\\,J0416+3650 can be decomposed into multiple individual sources inside the EGRET full-band point-spread function, revealing a significant signal from the nominal position of 3C\\,111 in the higher-resolution, high-energy band above 1\\,GeV. This association of 3C\\,111 with 3EG\\,J0416+3650, which had originally been suggested by \\citet{Har99} and \\citet{Sgu05}, makes 3C\\,111 one of the very rare radio galaxies detected at gamma-ray energies and supports the view that this source may be considered a lower-luminosity version of powerful radio-loud quasars.  Here, we report the results from ten years of Very-Long-Baseline Interferometry (VLBI) observations of 3C\\,111 as part of the  VLBA 2\\,cm Survey\\footnote{\\tt http://www.cv.nrao.edu/2cmsurvey/} \\citep{Kel98,Zen02,Kel04,Kov05} and its follow-up program MOJAVE\\footnote{\\tt http://www.physics.purdue.edu/MOJAVE} \\citep{Lis05,Hom06b}. We investigate the parsec-scale source structure during a major flux-density outburst and during its aftermath. We find that this outburst was associated with the formation of an exceptionally bright  feature in the jet of 3C\\,111. A variety of processes (beyond the predictions of simple ballistic motion models) are observed and discussed in view of modern relativistic-jet simulations. In Sect.~\\ref{sect:obs}, our observations and the data reduction are described. A detailed report of the observational results is given in Sect.~\\ref{sect:results}. In Sect.~\\ref{sect:discussion}, we discuss the various processes observed in the jet of 3C\\,111 as a result of the outburst and during the propagation of the new jet feature along the jet. In Sect.~\\ref{sect:conclusions}, we put these results into the context of future simulations and observations with the goal of understanding the production mechanisms of AGN jets. ",
        "conclusions": "\\label{sect:conclusions} In this paper, we have investigated the parsec-scale jet kinematics and the interaction of the jet with its ambient medium in the broad-line radio galaxy 3C\\,111. Our analysis has demonstrated that a variety of processes influence the jet dynamics in this source: a plasma injection into the jet beam associated with a major flux-density outburst leads to the formation of multiple shocks that travel at different speeds downstream and interact with each other and with the ambient medium. The primary perturbation causes the formation of a forward and a backward shock (or rarefaction). The latter fades away so fast that is likely to remain undetected in minor ejections. A separate work by \\citet{Per07} focuses on the nature and characteristics of these initial components. Several parsecs downstream, the jet plasma enters a region of rapidly decreasing external pressure, expands into the jet ambient medium and accelerates. In the following, the plasma gets recollimated and trailing features are formed in the wake of the leading component.  A particularly interesting aspect of the source 3C\\,111 in the light of this and other recent works is that it is one of the very rare non-blazar gamma-ray bright AGN. Besides Centaurus\\,A \\citep{Sre99} and the possible identification of NGC\\,6251 with the EGRET source 3EG\\,J1621+8203 \\citep{Muk02}, 3C\\,111 is the only AGN whose jet-system is inclined at a relatively large angle to the line of sight and that has a reliable EGRET identification: \\citet{Sgu05} reconsidered the possible identification of the EGRET source 3EG\\,J0416+3650 with 3C\\,111, which was first suggested by \\citet{Har99} but considered unlikely because of the poor positional coincidence. Very recently, R.\\,C.~Hartman \\& M.~Kadler (in prep.) found that 3EG\\,J0416+3650 is composed out of at least two distinct components. One of them is the dominant source above 1\\,GeV and is in excellent positional agreement with the location of 3C\\,111. Compared to blazars, the large inclination angle and the relatively small distance of 3C\\,111 allow us to resolve structures along the jet that are as small as parsecs in deprojection and which would be heavily blended with adjacent features in blazar jets. As demonstrated in this paper, VLBA observations of 3C\\,111 probe a variety of physically different regions in a relativistic extragalactic jet such as a compact core, superluminal jet components, recollimation shocks and regions of interaction between the jet and its surrounding medium, which are all possible sites of gamma-ray production. From early 2008 on, the gamma-ray satellite GLAST \\citep{Lot07} is going to monitor the sky. If detected by GLAST, 3C\\,111 may become a key source in the quest for an understanding of the origin of gamma-rays from extragalactic jets. In addition, the combination of GLAST and VLBA data with spectral data at intermediate wavelengths (optical, IR, X-ray) may allow a better determination of jet parameters and relativistic beaming effects than in most blazars because of the higher linear resolution offered by this nearby and only weakly projected jet system.   Our observations of 3C\\,111 are qualitatively in remarkable agreement with numerical relativistic hydrodynamic structural and emission simulations of jets such as the ones presented by \\cite{Agu01} and \\cite{Alo03}. Further progress is being made in the transition from two-dimensional to three-dimensional simulations of relativistic jets and in the development of new methods considering magnetic fields \\citep[e.g.,]{Lei05,Miz07,Roc08}, the equation of state for relativistic gases \\citep{Per07b}, and radiative processes \\citep[e.g.,][and Mimica et al. in preparation]{Mim04,Mim07}. But so far neither observational data nor simulations have reached an adequate level of detail and completeness in order to allow us a quantitative direct comparison of numerical models and observed relativistic jet structure and evolution. In particular, it is not feasible today to fit iteratively the parameters of relativistic magneto-hydro-dynamical (RMHD) jet simulations to match the brightness distribution observed for any individual source. The main reasons for this are a) the immense computational power required to conduct a realistic (i.e., sufficiently detailed) modern 3D jet simulation and b) the  highly non-linear nature of RMHD plasmas and their evolution. Simulation results depend critically on the starting conditions like the exact velocity, composition, and transversal structure of the flow, the structure and strength of the magnetic field and the jet environment. Future development of computational power will allow us to use larger resolutions to decrease the numerical viscosities, and to implement nonlinear and microphysics processes into simulations. VLBA observations are capable of putting hard quantitative constraints on the input parameters for RMHD jet simulations if they are densely sampled over several years. Polarimetric observations at multiple radio frequencies may allow the effects of jet-intrinsic magnetic-field variations and external Faraday-screen inhomogenities or temporal variations to be disentangled. Such data at 15\\,GHz are on the way, e.g., as part of the next phase of the MOJAVE program, in which rapidly evolving sources like 3C\\,111 are being observed every two months."
    },
    "0801/0801.1018_arXiv.txt": {
        "abstract": "{}{We study the effects of rotation on the outer convective zones of massive stars.}{We examine the effects of rotation on the thermal gradient and on the Solberg--Hoiland term by analytical developments and by numerical models.}{Writing the criterion for convection in rotating envelopes, we show that the effects of rotation on the thermal gradient are much larger and of  opposite sign to   the effect of the Solberg--Hoiland criterion. On the whole, rotation favors convection in stellar  envelopes at the  equator  and to a smaller extent at the poles. In a rotating 20 M$_{\\odot}$ star at 94\\% of the critical angular velocity, there are two convective envelopes, with the bigger one having a thickness of 13.2\\% of the equatorial radius. In the non-rotating model, the corresponding convective zone has a thickness of only 4.6\\% of the radius. The occurrence of outer convection in massive stars has many consequences.}{}    \\keywords {stars: evolution - convection - rotation } \\titlerunning{Convection in OB Stars} ",
        "introduction": "It  is generally  considered that the Cowling model applies to massive OB stars: i.e., a convective core surrounded by a large radiative envelope. However, long since stellar models  have shown that massive stars have an outer  convective envelope  encompassing several  percent of the stellar radius \\citep{Maeder80}. Also, \\citet{Langer97} has shown that an  Eddington factor $\\Gamma  ={\\kappa L}/{(4 \\pi c GM)}$ tending toward 1.0 implies convection.  Our aim  is to show that fast rotation amplifies the size of the convective envelope in OB stars as well as to develop anisotropic convective envelopes.  Various limits can be considered about the effects of rotation and high luminosity on the stellar stability \\citep{Langer97,MMVI}: the $\\Gamma$--limit, which is the Eddington limit for  $\\Gamma \\rightarrow 1$; the $\\Omega$--limit, which is reached by stars at rotational break--up with a small or negligible effect of the Eddington factor $\\Gamma$; the $\\Omega \\Gamma$-- limit, which applies to stars where both luminosity and rotation play significant roles. We   show that not only the stars at the $\\Gamma$--limit, but also the stars at  the $\\Omega \\Gamma$--limit and at the  $\\Omega$--limit, have amplified external convective zones.   The occurrence of outer  convective envelopes in OB stars and their  anisotropic structure lead to many astrophysical consequences: \\begin{itemize} \\item Convection generates acoustic modes that may allow  asteroseismic observations of OB stars. \\item Convective motions  may play a role in driving mass loss by stellar winds. \\item For stars close to the critical rotation, convective motions lower the effective break--up velocities. \\item An outer convective envelope  may make a dynamo and  contribute to some chromospheric activity generating an  X--ray emission from OB--stars. \\item A convective envelope transports   chemical elements and angular momentum. \\item The occurrence of outer convection may modify the von Zeipel theorem \\citep{vZ24}. \\end{itemize} A closer investigation is justified. We  start by an analytical approach  (Sect. \\ref{analytical}), and finish by  two--dimensional models of 20 M$_{\\odot}$  rotating envelopes (Sect. \\ref{numerical}). ",
        "conclusions": "In stellar envelope of rotating stars, the effects of rotation on the thermal gradient arestronger and with the opposite sign with respect to the Solberg--Hoiland criterion, so that rotation favors convection instead of inhibiting it. The increase of the convective zone occurs mainly at the equator  and  also a bit at the poles. In a fast--rotating 20 M$_{\\odot}$ Pop I star, there are  two  equatorial zones covering a total of 16\\% of the stellar radius at the equator.  There are several consequences of thees results to be examined in future. The outer convective motions may   lower the escape velocity as well as the critical rotation velocity. The matter accelerated in the winds continuously goes through the convective zone in a dynamical process, suggesting  that convection plays a role in accelerating the stellar winds and in producting the clumps in the winds. The convective pistons generate acoustic waves of  periods of several hours to a few days. The density is very low, and it is thus likely that convection injects oscillations into the wind rather than into the interior."
    },
    "0801/0801.4135_arXiv.txt": {
        "abstract": "The Bekenstein-Milgrom gravity theory with a modified Poisson equation is tested here for the existence of triaxial equilibrium solutions. Using the non-negative least square method, we show that self-consistent triaxial galaxies exist for baryonic models with a mild density cusp $\\rho \\sim {\\Sigma \\over r}$. Self-consistency is achieved for a wide range of central concentrations, $\\Sigma \\sim 10-1000\\mathrm{M_{\\odot}pc^{-2}}$, representing low-to-high surface brightness galaxies. Our results demonstrate for the first time that the orbit superposition technique is fruitful for constructing galaxy models beyond Newtonian gravity, and triaxial cuspy galaxies might exist without the help of Cold dark Matter. ",
        "introduction": "Constructing models of galaxies in triaxial equilibrium is a classical challenge in dynamics, either in Newtonian or Modified gravity. However, due to numerical constructions, few such models have been developed in Newtonian since the pioneering work of Schwarzchild(1979) on triaxial elliptical galaxies, and Zhao (1996) on the fast rotating triaxial bar of the Milky Way. Merritt \\& Fridman (1996) (hereafter MF96) considered a model with a central density cusp $r^{-1}$ and found that it is possible to construct a model of triaxial equilibrium galaxies self-consistently in Newtonian dynamics.  Models with a steeper density cusp are not in equilibrium. This has generated interests in understanding cusp-triaxiality relation, the triaxiality-velocity anisotropy relation and the cusp-black hole relation. Recently, Capuzzo-Dolcetta et al. (2007) modeled $r^{-1}$ cuspy triaxial galaxies in $\\Lambda$ Cold Dark Matter (CDM) haloes. They took the model of MF96 as the luminous density distribution, and found that a model of cupsy triaxial galaxies with CDM halos is also self-consistent.  A tough problem for $\\Lambda$CDM ($\\Lambda$ plus Cold Dark Matter) is that there is no physics to justify $\\Lambda$ the tiny cosmological constant or dark energy (see White 2007, Sarkar 2007). Although future particle physics experiments might well prove the existence of new species of particles and vacuum energy, their experimental-determined abundance could easily be factors of a few different from the 1:3 ratio as precisely required by fitting observations of Microwave background (Zhao 2006).  This motivates exploring alternative theories like the co-variant version of Modified Newtonian Dynamics (MOND) in the mean-time.  E.g., instead of more traditional interpretation of modified gravity (Bekenstein 2004), MOND was given the interpretation as a co-variant Dark Energy (DE) the ${\\bf V\\Lambda}$ model of Zhao (2007).  This model uses General Relativity plus a non-uniform Dark Energy (DE) fluid (described by a four-vector flow) to give both the effects of DE and DM.  Perturbations of such DE fluid bends space-time as an effective DM without actually invoking DM.  Unlike $\\Lambda$CDM halos, which are positive and nearly round, the effective DM can sometimes be extremely flattened or negative in some regions of galaxies or clusters (Wu et al. 2007, 2008), although the data are not good enough to tell such negative regions yet (Nipoti et al. 2007a).  Overall MOND and $\\Lambda$CDM halos are comparably successful in explaining the flat rotation curves of high surface brightness discs and satellites orbits at large radii (e.g., Angus et al. 2008). However, both theories are faced with severe challenges: $\\Lambda$CDM hinges on uncertain feedback on galactic scales, and even the maximum feedback falls short of explaining the velocity curves of low-surface brightness systems (Gnedin \\& Zhao 2002); Bekenstein's (2004) co-variant MOND seems to rely on non-relativistic neutrinos on large scale to match the weak lensing data (Angus et al. 2007) and the Cosmic Background Radiation (Skordis et al. 2006). \\footnote{although these problems appear resolvable in the ${\\bf V\\Lambda}$ model (Zhao 2007).}  The tight correlation observed between the mass profiles of baryonic matter and dark matter at all radii in spiral galaxies (e.g., McGaugh et al. 2007; Famaey et al. 2007) is still the best case for the MOND paradigm of Milgrom (1983), which postulates that for accelerations below $a_0 = 1.2 \\times 10^{-10} {\\rm m} \\, {\\rm s}^{-2}$ the effective gravitational attraction approaches $(g_N a_0)^{1/2}$, where $g_N$ is the usual Newtonian gravitational field. Indeed, without resorting to galactic dark matter, this simple prescription is amazingly successful at reproducing galactic rotation curves over five decades in mass ranging from tiny dwarfs (e.g., Gentile et al. 2007) to early-type disk galaxies (e.g., Sanders \\& Noordermeer 2007) to massive ellipticals (Milgrom \\& Sanders 2003). Some outliers exist in models of gravitational lensing (Zhao et al. 2006, Chen \\& Zhao 2006, Shan et al. 2008). Lensing and structure formation are well-defined calculations in co-variant MOND (Halle \\& Zhao 2007). Galaxy formation simulations in classical MOND have been carried out too, and it is possible to form bars (Tiret \\& Combes 2007), and elliptical galaxies by mergers (Nipoti et al. 2007b), and spherical bulges by instability (Zhao, Xu \\& Dobbs 2008).  Very little is known about the triaxial stationary equilibrium in non-Newtonian dynamics, although one speculates that such equilibrium might exist. In this paper, we test the existence of triaxial galaxies in the Bekenstein-Milgrom MOND theory, a well-defined classical theory with relativisitic and cosmological extentions. We apply the same density model given by MF96 with a fixed mass, and axis ratios. Our aim is to examine whether this model is self-consistent in MOND.  The MOND theory is scale-dependent.  When the gravitational acceleration is below $a_0$, the scaling deviates from the $r^{-2}$ Newtonian law.  Here we consider models with different ratio of $a_0$ to the simulation unit G $\\times$ (Total Mass) / (Scale Radius)$^{2}$. We start with a nearly Newtonian system (high surface brightness) and finish with a deep-MOND system (low surface brightness).  The rest of the paper is organized as follows. In \\S2, we present the density distribution of the galaxy and the corresponding potential and forces. The orbits are calculated in \\S3. The construction of self-consistent models is discussed in \\S4, and in \\S5, we give our conclusions and disscussions. ",
        "conclusions": "In practice, a parameter $\\delta$ is adopted to describe the departure form self-consistency, which is defined as in MF96: \\begin{equation} \\delta=\\frac{\\sqrt{\\chi^2}}{\\bar{M}}, \\end{equation} where $\\bar{M}$ is the average mass in each cell. Figure \\ref{delta} shows the departure from self-consistency as a function of the number of orbits. Note that there is a remarkably strong dependence on the number of orbits. The departure parameter $\\delta$ drops quickly when the number of orbits exceeds 4000. $\\delta$ is $~10^{-15}$ when the 6840 orbits are adopted in the optimization routine, which shows that all three triaxial models are self-consistent in MOND. Amazingly, we also find that the departure parameter $\\delta$ is insensitive to $a_0$. In other words, if one of the triaxial galaxies is self-consistent in MOND, then the same model with a different value of $a_0$ will also be self-consistent.  Figure \\ref{ps_den} shows the contribution of mass for various orbital families in the self-consistent model. Chaotic orbits contribute a larger fraction of mass with the increase of energy. This shows a certain degree of consistency with that of Capuzzo-Dolcetta et al. (2007). However, discrepancies between our results and that of Capuzzo-Dolcetta et al. (2007) are apparent. There is a clear variational trend of the cumulative energy distribution of the various orbital families in Newton's case, however, no apparent variational trend for different orbits has been found in MOND.  In Figure \\ref{sigma}, we show the radial profile of the rescaled radial velocity dispersion $\\sigma_r/V_{\\rm cir,\\infty}$(left panel) and the radial profile of the anisotropy parameter $\\beta=1-0.5(<V^2-V_r^2>)/<V_r^2>$ (right panel), where $V_{\\rm cir,\\infty}$ is the circular velocity at infinity, and $V_r$ is radial velocity. The parameter $V^2$ is defined as $V^2=v_x^2+v_y^2+v_z^2$, where $v_x$, $v_y$ and $v_z$ are the three components of velocity. It can be seen that the profiles of the radial velocity dispersion are nearly flat in MOND, except in the central region. Recently, Capuzzo-Dolcetta et al. (2007) studied the same model as that used here, but in the Newtonian gravity. It can be seen that the profile of the velocity dispersion in our model shows a certain degree of consistency with that of Capuzzo-Dolcetta et al. (2007), but the discrepancies between our results and that of Capuzzo-Dolcetta et al. (2007) are apparent. There is central peak in our velocity dispersion profile. The global average of the velocity  $V_{\\rm rms}$ is 170.49, 78.84 and 42.70 km/s for $a_0$=0.083, 0.833, and 8.333, respectively (see the last column of Table~\\ref{parameter}).  The anisotropy parameter $\\beta\\geq0$ suggests that the stars have high radial velocity and the ratio of $<V_r^2>/<V^2-V_r^2>$ reaches a maximum value at $2-3$ kpc. Our anisotropy profiles are similar to that of Nipoti et al. (2007b) form dissapationless collapse in MOND (their figur 2) and in Newtonian gravity van Albada (1982) and Wang et al. (2008a). A gently rising beta is widely used (e.g., Milgrom \\& Sanders 2003). However, the peak in the anisotropy profile in our model is not seen in any simulations and observations to our best knowledge. This is might be indication of problem of MOND. Our models have a sudden change of the amount of box orbit near shells 10-12 (radius~2-3kpc), which could be the reason.  In order to check the stability of our models, we used the Antonov's third law (Binney \\& Tremaine 1987). If the stellar system is stable, then its density $\\rho$ and potential $\\Phi$ satisfy everywhere the inequality ${\\rm d}^3\\rho_b/{\\rm d}\\Phi^3 < 0$. Figure~\\ref{drdp} shows that ${\\rm d}^3\\rho_b(0,y,0)/{\\rm d}|\\Phi(0,y,0)|^3$ is positive everywhere along the intermediate axis for our three models, since the values of potential are negative, our models are consistent with the law, which means the models appear globally stable. We will also be checking the stability of the models in collaboration with groups which have MOND N-body code. Effects due to pattern rotation and external field in MOND remain to be studied (Wang et al. 2008b).   We have used the Schwarzschild approach to examine whether the triaxial galaxies are self-consistent in MOND. Using the Bologna Poisson solver to determine the potential and the accelerations at some discrete points, we used three-dimensional interpolation then obtain the potential and accelerations at arbitrary point. The orbits conserve energy to a very high level of accuracy, which shows that our method is feasible in calculating orbits using a grid-based MONDian potential. The departure parameter $\\delta$ shows that the triaxial galaxy models adopted in this paper in MOND are self-consistent. Our results show that it is theoretically allowed to have triaxial galaxies in rigorous equilibrium in a universe with either non-Newtonian gravity or certain models of Dark Energy (Zhao 2007)."
    },
    "0801/0801.3653_arXiv.txt": {
        "abstract": "\\noindent Many models of baryogenesis rely on anomalous particle physics processes to give baryon number violation. By numerically evolving the electroweak equations on a lattice, we show that baryogenesis in these models creates helical cosmic magnetic fields. After a transitory period, electroweak dynamics is found to conserve the Chern-Simons number and the total electromagnetic helicity. We argue that baryogenesis could lead to magnetic fields of nano-Gauss strength today on astrophysical length scales. In addition to being astrophysically relevant, such helical magnetic fields can provide an independent probe of baryogenesis and CP violation in particle physics. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.4559_arXiv.txt": {
        "abstract": "We present a survey of bright optical dropout sources in two deep, multiwavelength surveys comprising eleven widely-separated fields, aimed at constraining the galaxy luminosity function at $z\\approx7$ for sources at 5-10L*($z=6$).  Our combined survey area is 225\\,arcmin$^2$ to a depth of $J_{AB}=24.2$ (3\\,$\\sigma$) and 135\\,arcmin$^2$ to $J=25.3$ (4\\,$\\sigma$).  We find that infrared data longwards of 2$\\mu$m is essential for classifying optical dropout sources, and in particular for identifying cool Galactic star contaminants. Our limits on the number density of high redshift sources are consistent with current estimates of the Lyman break galaxy luminosity function at $z=6$. ",
        "introduction": "\\label{sec:intro}   Direct observations of galaxies at very high redshift can cast light on the details of structure formation and early galaxy evolution. In recent years, the most commonly observed tracer of star formation at high redshift has been the Lyman Break galaxy (LBG) population \\citep{1999ApJ...519....1S}.  These sources are selected photometrically based on a spectral feature interpreted as being the $\\lambda_\\mathrm{rest}=$1215.67\\AA\\ Lyman break. Neutral hydrogen along the line of sight leads to a forest of redshifted Lyman-$\\alpha$ absorption lines, blue continuum flux is suppressed, and a galaxy will exhibit red colours in a pair of filters bracketing the break. Depending on the relative depth of the imaging in different bands the galaxy will be seen to fall in magnitude between the red and bluewards bands, or even to drop below the detection limit.  As a result, Lyman break galaxies are often termed `drops' or `dropouts' and their approximate redshift is determined by the band bluewards of the break. Hence `U'- and `G'-drops lie at $z\\approx3$ and $z\\approx4$ \\citep{1999ApJ...519....1S}, `V'- and `R'-drops at $z\\approx5$ \\citep{2003ApJ...593..630L,2004ApJ...600L.103G} and `I'-drops at $z\\approx6$ \\citep{2003MNRAS.342..439S,2006ApJ...653...53B}.   The use of LBG samples to study star formation at ever-increasing redshifts has recently encountered a practical barrier.  At $z\\approx7$ the Lyman break is redshifted to 1\\,$\\mu$m, a spectral region with poor sensitivity in both optical CCD detectors and near-infrared instruments. Photometry in the relatively broad photometric filters used longwards of 1\\,$\\mu$m effectively smears out the signatures of sharp-sided spectral features such as the Lyman break, making them difficult to distinguish from more gradual spectral features such as a Balmer break or molecular absorption in low redshift sources. The small field of view, high background, low sensitivity and low multiplexing of current near-infrared spectrographs further complicate the follow-up of Lyman break galaxy candidates.  Sensitive, wide-format near-infrared imagers and multi-object spectrographs on the largest telescopes offer potential for $z=7$ Lyman Break galaxy searches, identifying bright near-infrared sources that drop in the optical bands (as discussed below in section \\ref{sec:col_sel}). Such surveys require knowledge of the $z=7$ luminosity function in order to estimate the required survey depth in a given area. High contamination rates in optical-dropout samples also complicate such surveys. While the most promising candidates will always require spectroscopic follow-up, it is clearly necessary to develop strategies for minimising the number of contaminating lower redshift sources from photometry alone.  An important tool for fulfilling these requirements is the {\\em Spitzer Space Telescope}, which images at infrared wavelengths.  However the 85cm aperture of {\\em Spitzer} and 1$''.$2$\\times$1$''.$2 arcsecond pixels of its IRAC instrument lead to blending and confusion issues at faint limits ($m_{AB}>24$).  Existing deep field surveys incorporate sensitive imaging at wavelengths from 4000\\AA\\ to 8$\\mu$m (the four bands of the IRAC instrument). As a result, they can provide a first constraint on the surface density of bright galaxies at very high redshifts, and a test of the ability of {\\em Spitzer} to distinguish genuine $z>7$ sources from lower redshift contaminants.  In this paper we present an analysis of two multiwavelength data-sets, identifying and classifying optical-dropout sources in eleven widely separated fields, averaging across cosmic variance. In section \\ref{sec:col_sel} we discuss the colour selection criteria which isolate high redshift galaxies. In sections \\ref{sec:ergs} and \\ref{sec:goods} we apply such criteria to the ESO Remote Galaxy Survey and the Great Observatories Origins Deep Survey respectively and consider the resulting candidate sources. Finally, in section \\ref{sec:discussion} we discuss the implications of this analysis for both the $z=7$ Lyman Break Galaxy luminosity function and future wide-area near-infrared survey.  All magnitudes in this paper (optical and infrared) are quoted in the AB system \\citep{1983ApJ...266..713O}.  We use a flat Universe with $\\Omega_{\\Lambda}=0.7$, $\\Omega_{M}=0.3$ and $H_{0}=70 h_{70} {\\rm km\\,s}^{-1}\\,{\\rm Mpc}^{-1}$. ",
        "conclusions": "\\label{sec:conc}  The key points presented in this paper can be summarised as follows:  i) We have identified an analysed a sample of $z$-band dropout sources in deep, multi-wavelength imaging, considering their suitability as potential $z>7$ galaxy candidates.  ii) In an area of 360\\,arcmin$^2$ (the ERGS+GOODS surveys), we determine a strict upper limit of 1.4 potential $z=7$ galaxies and a likely limit of zero candidates to $J=24.2$.  iii) In an area of 135\\,arcmin$^2$ (the GOODS survey), we determine a strict upper limit of 2.8 potential $z=7$ galaxies and a likely limit of no good candidates to $J=25.3$.  iv) These results are consistent with current estimates of the luminosity function of Lyman break galaxies at $z=6$, for galaxies with $L>5L*$.  v) The use of deep data longward of the $K$-band is extremely valuable if the number of high redshift candidates from deep surveys such as UKIDSS UDS from unmanageable to reasonable numbers for spectroscopic follow-up.    \\subsection*{Acknowledgements}  ERS gratefully acknowledges support from the UK Science and Technology Facilities Council (STFC).  Based in part on observations made with the NASA/ESA Hubble Space Telescope, obtained from the Data Archive at the Space Telescope Science Institute. Also based in part on observations made with the Spitzer Space Telescope, which is operated by the Jet Propulsion Laboratory, CalTech. We thank the GOODS team for making these high quality datasets public. Research presented here has benefitted from the M, L, and T dwarf compendium housed at DwarfArchives.org and maintained by Chris Gelino, Davy Kirkpatrick, and Adam Burgasser. Results from the EDisCS/ERGS fields are based on observations made with ESO telescopes under programmes 166.A-0162 and 175.A-0706."
    },
    "0801/0801.3186_arXiv.txt": {
        "abstract": "The unusual mushroom-shaped HI cloud, GW 123.4--1.5, is hundreds of parsecs in size but does not show any correlations to HI shells or chimney structures. To investigate the origin and velocity structure of GW 123.4--1.5, we perform three-dimensional hydrodynamical simulations of the collision of a high-velocity cloud with the Galactic disk. We also perform a parameter study of the density, radius, and incident angle of the impact cloud. The numerical experiments indicate that we reproduce the mushroom-shaped structure which resembles GW 123.4--1.5 in shape, size, position-velocity across the cap of the mushroom, and the density ratio between the mushroom and surrounding gas. GW 123.4--1.5 is expected to be formed by the almost head-on collision of a HVC with velocity $\\sim 100 \\kms$ and mass $\\sim 10^5 \\Msun$ about $5 \\times 10^7 \\yr$ ago. A mushroom-shaped structure like GW 123.4--1.5 must be infrequent on the Galactic plane, because the head-on collision which explains the mushroom structure seems rare for observed HVCs. HVC-disk collision explains not only the origin of the mushroom-shaped structure but also the formation of a variety of structures like shells, loops, and vertical structures in our Galaxy. ",
        "introduction": "GW 123.4--1.5 revealed by the Canadian Galactic Plane Survey \\citep{Tay03} is a mushroom-shaped HI structure composed of a stem and a cap and is dissimilar to any common shells or chimneys \\citep{eng2000}. Using the kinetic distance of the Galactic rotation model, \\citet{eng2000} derived properties of the mushroom cloud such as projected length, mass, velocity, mean excess column-density, and kinetic energy. The mushroom cloud extends a few hundred parsecs ($\\sim 350 \\pc$) at the assumed distance of $3.8 \\pm 1.2 \\kpc$ and has an HI mass of $\\sim 10^5 \\Msun$,  a mean excess column-density of $N_{\\rm H} \\simeq 6 \\times 10^{20} {~{\\rm cm^{-2}}}$, and a kinetic energy of $\\sim 2 \\times 10^{50}$ ergs. As a model for the origin of the mushroom-shaped cloud GW 123.4--1.5, the gas buoyancy model \\citep[]{eng2000,dm2001} and the cloud collision model \\citep[hereafter KB2004]{kb2004} were proposed and numerical simulations were executed, respectively. \\citet{eng2000} carried out preliminary two-dimensional hydrodynamical simulations and showed that the origin of the mushroom-shaped GW 123.4--1.5 is the rise of buoyant gas from the explosion of a supernova above the Galactic midplane. Using three-dimensional hydrodynamical simulations of large-scale modeling of the interstellar gas in the galactic disk \\citep{da2000}, \\citet{dm2001} showed that the mushroom-shaped cloud can result from the buoyant rise of bubbles out of pools of hot gas created from a supernova or supernovae. However there is some difficulty in interpreting the origin of GW 123.4--1.5 by the buoyancy of a hot gas created by a single or multiple supernova explosions, because there are no observations of soft X-ray emission and IRAS sources near the base of GW 123.4--1.5. On the other hand, the collision of a high-velocity cloud (HVC) with the Galaxy, which was studied by \\citet{Ten86} for the first time, is not restricted to the regions of active star formation \\citep{Ten87}. KB2004 performed two-dimensional hydrodynamical simulations of the collision of a cloud with the Galactic disk and proposed that GW 123.4--1.5 is created by the impact of an intermediate-velocity cloud (IVC)\\footnotemark \\footnotetext{Historically, the clouds with $|v_{lsr}| \\geq 70 \\kms$ are classified into HVCs and those with $ < 70 \\kms$ are called intermediate-velocity clouds (IVCs).} with the Galactic disk. Also they showed that the velocity structure across the cap of the mushroom-shaped structure in their simulation is consistent with that of GW 123.4--1.5. In the buoyant models, the velocity characteristic of the cap is not obvious. Unlike the gas buoyancy model, cloud-disk collision can easily explain the formation of GW 123.4--1.5 without inference for X-ray emissions.  In this study we extend the two-dimensional hydrodynamical simulation of KB2004 to three dimensions, in order to investigate the origin, dynamics, and velocity structure of a mushroom-shaped structure. Through numerical simulations and a parameter study, we can reproduce the size, shape, and position-velocity of the mushroom-shaped structure which resemble those of GW 123.4--1.5. Consequently, we can infer (1) how GW 123.4--1.5 was formed in the Galactic disk and (2) why a mushroom-shaped structure is so rare on the Galactic plane. In the next section we describe our models and numerical method. Simulation results are presented in \\S III, followed by a summary and discussion in \\S IV. ",
        "conclusions": "We perform three-dimensional hydrodynamical simulations for the impact of a HVC with the Galactic disk, in order to explore the formation of the mushroom-shaped HI cloud GW 123.4--1.5. The main results can be summarized as follows: \\begin{itemize} \\item A mushroom-shaped structure can be formed not only by a high-density IVC collision with the Galactic disk (KB2004) but also by a low-density HVC collision. \\item The resultant structures formed by cloud-disk collision are greatly affected by the density and incident angle of a colliding cloud. \\item GW 123.4--1.5 is expected to be formed by the almost head-on collision of a HVC with velocity $\\sim 100 \\kms$ and mass $\\sim 10^5 \\Msun$ about $5 \\times 10^7 \\yr$ ago. \\end{itemize}  Recently, \\citet{as2005} have tried to find a mushroom-shaped HI structure like GW 123.4--1.5 in the Canadian Galactic Plane Survey. However, they did not find other mushroom shape structures. Our simulations could explain why mushroom-shaped structures are so infrequent in our Galaxy. A nearly head-on collision is needed to form a mushroom-shaped structure. In a rotating disk, the head-on collision of a cloud with the Galactic plane would be rare in a general way.  HVC-disk collision can not only explain the origin of a mushroom-shaped structure but also the formation of a variety of structures like shells, loops and vertical structures in our Galaxy. For example, in model O30 (see the middle panel of Fig.\\,6), we can find an interesting structure, a loop which resembles that found by NANTEN in the central molecular zone within $\\sim 1 \\kpc$ from the Galactic center \\citep{Fuk2006}. The two loops of CO emission obtained with NANTEN have strong velocity gradients in a longitude-velocity diagram: $\\sim 80 \\kms$ per $250\\pc$ and $\\sim 60 \\kms$ per $150\\pc$ in each loop \\citep[see Fig.\\,2 of][]{Fuk2006}. Although the projected velocity distribution of our simulation ($\\sim 20 \\kms$) is somewhat smaller than that of observation ($\\sim 60 - 80 \\kms$), the size of the loop, $\\sim 300 (H_0/140\\,{\\pc}) \\pc$, is similar to that of the molecular loop. In order to explain the formation of the molecular loop, \\citet{Fuk2006} proposed a magnetic flotation model (due to magnetic buoyancy caused by the \\citet{Park66} instability). From our preliminary result of model O30, we also propose that HVC-disk collision is another possible model of formation of the molecular loop structure at the Galactic center. This issue will be studied in detail in an upcoming paper.  The present study as well as the previous study (KB2004) are restricted to hydrodynamical simulations. When the magnetic field parallel to the Galactic disk exists \\citep[]{santi2000,franco2002,nozawa2005} , we also expect the magnetic buoyant effect (Parker instability). Both the impact of the cloud and the Parker instability \\citep[]{santi1999,santi04,santi07k} may create complicate structures in the Galactic disk such as worms, loops, shells, and chimneys. Moreover the interaction between the infalling HVCs and the magnetized disk must be an elementary process in the interstellar dynamics. Using high-resolution two-dimensional MHD simulation, \\citet{santi07} pointed out a possibility that the rain of compact HVCs onto the disk can maintain transonic turbulent motion in the warm phase of HI. In this sense, we believe that high-resolution three-dimensional MHD simulation will broaden our understanding of the structures formed by HVC collisions and bring us more interesting phenomena to be important in the interstellar dynamics. In the near future, we will perform three-dimensional magnetohydrodynamical simulations."
    },
    "0801/0801.2978_arXiv.txt": {
        "abstract": "HD 179821 is an evolved star of unknown progenitor mass range (either post-Asymptotic Giant Branch or post-Red Supergiant) exhibiting a double peaked spectral energy distribution (SED) with a sharp rise from $\\sim8-20$ $\\mu$m.  Such features have been associated with ejected dust shells or inwardly truncated circumstellar discs.  In order to compare SEDs from both systems, we employ a spherically symmetric radiative transfer code and compare it to a radiative, inwardly truncated disc code.  As a case study, we model the broad-band SED of HD 179821 using both codes.  Shortward of 40 $\\mu$m, we find that both models produce equivalent fits to the data.  However, longward of 40 $\\mu$m, the radial density distribution and corresponding broad range of disc temperatures produce excess emission above our spherically symmetric solutions and the observations.  For HD 179821, our best fit consists of a $T_{eff}=7000$ K central source characterized by $\\tau_V\\sim1.95$ and surrounded by a radiatively driven, spherically symmetric dust shell.  The extinction of the central source reddens the broad-band colours so that they resemble a $T_{eff}=5750$ K photosphere.  We believe that HD 179821 contains a hotter central star than previously thought.  Our results provide an initial step towards a technique to distinguish geometric differences from spectral modeling. ",
        "introduction": "HD 179821 (IRAS 19114+0002; V1427 Aql; SAO 124414; AFGL 2423) is an evolved star surrounded by gas and dust ejected during a phase of mass-loss.  The luminosity of this object is undetermined as the Hipparcos parallax measurement ([0.18 $\\pm$ 1.12 mas]; \\citealt{1997A&A...323L..49P}) allows for any distance greater than 1 kpc.  HD 179821 is either a close post-Asymptotic Giant Branch (post-AGB) star ($D=1$ kpc, $M_i\\sim3-4$ $M_\\odot$) or a distant, massive ($D=5-6$ kpc, $M_i\\sim20-30$ $M_\\odot$) post-Red Supergiant (post-RSG).  The post-AGB phase of stellar evolution is a short ($\\sim10^3$ yrs) period in which initially intermediate-mass main sequence stars ($<8$ $M_\\odot$) transition from the Asymptotic Giant Branch to the planetary nebula (PN) phase.  In contrast, the progenitors of Red Supergiants are massive main-sequence stars ($>8$ $M_\\odot$) evolving toward the Wolf-Rayet stage and may end their lives as supernovae.  For a short review on the post-RSG status of HD 179821, see \\cite{2008arXiv0801.2315O}.  In both phases, mass is ejected into the circumstellar environment.  Distinguishing between these classes of objects can be difficult.  In this paper, we model the broad-band SED of HD 179821 using two distinct radiative transfer codes.  The first code, ``DUSTY\", computes radiative transfer through a spherically symmetric shell with canonical density profile $\\rho(r)\\propto r^{-2}$ (\\citealt{DUSTY}).  The second code computes radiative transfer through an axisymmetric, flared disc with canonical density profile $\\rho(r)\\propto r^{-3/2}$ (\\citealt{2001ApJ...553..321D,2005ApJ...621..461D}).  While we believe the mid-IR and radio imaging of HD 179821 indicates a roughly spherical nebula, the point of this paper is to determine whether spectral modeling, through the use of distinct radiative transfer codes, can constrain geometric features of evolved star nebulae.  Thus, HD 179821 serves as a test case for this investigation.  We model HD 179821 for both nebular geometries and investigate the degeneracies between the SED model fits.  In addition, we present and compare the ISO spectrum of HD 179821 to the post-AGB object HD 161796 and post-RSG object IRC +10420.  All three objects display a strong mid-IR excess and exhibit steep increases in their SED's beyond 8 microns.  In Section 2, we discuss the enigmatic nature of HD 179821 in both the post-RSG and post-AGB scenarios.  In Section 3, we describe our synthetic photospheres and extinction corrections.  In Section 4, we present our SED model fits from the radiative transfer disc code.  In Section 5, we present our model fits from DUSTY and identify wavelength regions in which degeneracies occur between the two codes.  For intermediate mass stars below 8 $M_\\odot$, the late stages of stellar evolution are characterized by transitions from  quasi-spherical mass-losing AGB stars to complex aspherical post-AGB objects.  Many post-AGB objects feature toroidal density enhancements, bipolar jets and an array of other non-spherically symmetric features, the origin of which is unknown (\\citealt{2002ARA&A..40..439B}).  In addition, most post-AGB objects display a large momentum excess above what would be supplied by radiation pressure (\\citealt{2001A&A...377..868B}).  While the origin of the additional source of momentum is unclear, a binary companion is an attractive candidate as energy and angular momentum can transfer from the secondary to the primary (\\citealt{JN2006}).  This hypothesis is supported by recent observational (\\citealt{2004ApJ...602L..93D,2007arXiv0709.1508D}) and theoretical (\\citealt{2006ApJ...650..916M,2006ApJ...645L..57S}) efforts suggesting that a significant fraction of PNe may be descendants of interacting binaries.  Binary companions can influence mass loss in many ways.  A common envelope phase, even for low-mass companions, can lead to equatorial outflows, poloidal outflows and the formation of discs (\\citealt{1999ApJ...524..952R}, \\citealt{JN2006}, \\citealt{JN2007}, \\citealt{2007arXiv0707.3792N}).  For wider binaries, low-mass companions can induce spiral waves and convert amorphous dust to crystalline dust via annealing (\\citealt{2007arXiv0709.2292E}).  For massive stars, the end stages of stellar evolution are poorly understood.  The post-red supergiant (post-RSG) phase is among the most luminous and uncertain epochs of post-main sequence stellar evolution.  Like their intermediate mass conterparts, post-RSG's are thought to be ejecting a substantial portion of their mass (\\citealt{2006A&A...456..549D}).  However, a massive circumstellar envelope has only been detected (scattered light, IR emission, molecular line emission) around the post-RSG IRC +10420, thus making a comparison with HD 179821 difficult (\\citealt{1995ApJ...452..833K}).  Mass-loss may be consistent with a radiation driven spherically symmetric outflow (\\citealt{2007A&A...465..457C}).  However, if HD 179821 does indeed display a momentum excess in its ejected nebula (\\citealt{2001A&A...377..868B}), then it is not unreasonable for a binary companion to have influenced the outflow and produced asymmetries.  Interestingly, the WFPC2 images of \\citealt{2000ApJ...528..861U} show collimated bipolar structures emerging from the dust shell.  These may be remnant jets propagating through the envelope.  There is also evidence for slight clumpy regions in both OH maser and mid-IR emission (\\citealt{2001MNRAS.328..301G,1995ApJ...452..833K}).  A density asymmetry may also be present as the $^{13}$CO line profiles are asymmetric (\\citealt{1992A&A...257..701B}).  The spectral energy distribution of HD 179821 exhibits a double-peaked shape, indicative of a stellar photospheric component and ejected dusty component.  The spectral energy distribution (SED) is consistent with photospheric emission to $\\sim8$ $\\mu$m at which point there is a steep increase until the second peak at $\\sim$ 25 $\\mu$m.   The overall SED of HD 179821 is thus remarkably similar to that of the transitional young stellar object CoKu Tau/4 (see Fig. 2 of \\citealt{2005ApJ...621..461D}), which has been well modeled as a dusty disc with an inner hole.  The sharp, interior wall of the disc is illuminated by the central source and produces the excess infrared emission.  HD 179821 also exhibits an evacuated interior region between the central object and hence, the system might be modeled similarly to CoKu Tau/4.  However, the dust envelope of HD 179821 could also be a detached spherical shell and past modeling efforts have treated it as such (\\citealt{2002Ap&SS.281..751S,2007A&A...462..637B,2007A&A...465..457C}). ",
        "conclusions": "We have modeled the broad-band spectral energy distribution of HD 179821 using two distinct radiative transfer codes: one corresponding to a spherical shell geometry with $\\rho\\left(r\\right)\\propto r^{-2}$, and one to a disc geometry with $\\rho\\left(r\\right)\\propto r^{{-3}/{2}}$.  Under these assumptions, both codes provided equally good fits for similar dust compositions and size distributions shortward of $40$ $\\mu$m.  However, longward of $40$ $\\mu$m, only the spherical shell model reproduces the observations.  The radial dust density profile and corresponding range of temperatures present in the disc provide excess emission above the spherical shell solution and cannot fit the data.  A wavelength of $\\sim$ 40 $\\mu$m marks the boundary for discerning the model fits.  It should be noted that the density profiles need not uniquely indicate a particular geometry.  If the densities deviate from the anticipated theoretical values of $\\rho\\left(r\\right)\\propto r^{-2}$ (shell), $\\rho\\left(r\\right)\\propto r^{-3/2}$ (disc), then the infrared excess is likely coupled to the density distribution rather than the geometry.  The effect of changing the density profile in both codes should be investigated before comparative spectral modeling is established as a reliable method for distinguishing geometric differences.   Based on our detailed spectral modeling, we conclude that it is likely that the nebula around HD 179821 is a radiatively driven shell.  The dust is mainly composed of small ($a_{max}=0.25$ $\\mu$m) amorphous, glassy silicates.  Interior to the shell is an evacuated region in addition to the central radiation source.  The central star is most likely a $T=7000$ K star obscured such that it appears as a $T=5750$ K star (optical depth of $\\tau_V\\sim1.95$ through the spherical shell).  This may aid in explaining previous conflicting spectral classifications.  The dimensions of our shell agree well with previous observations and provide a good fit to the ISO spectrum.  We compared the ISO spectrum to that of HD 161796, a confirmed post-AGB object (e.g. Fig. \\ref{geometry}).  In particular, the ISO spectra and broad-band SED look remarkably similar.  Because our results scale with distance, our spectral modeling does not provide insight into whether HD 179821 is a post-AGB or post-RSG.  However, the similarity of the ISO spectra suggests that HD 161796 and HD 179821 experienced a very similar mass-loss history."
    },
    "0801/0801.0699_arXiv.txt": {
        "abstract": "An asymptotic treatment of thin accretion disks, introduced by Klu\\'zniak \\& Kita (2000) for a steady-state disk flow, is extended to a time-dependent problem. Transient growth of axisymmetric disturbances is analytically shown to occur on the global disk scale. The implications of this result on the theory of hydrodynamical thin accretion disks, as well as future prospects, are discussed. ",
        "introduction": "\\label{1sec} Thin accretion disks (AD) form whenever a sufficiently cool gas, endowed with a significant amount of angular momentum, is gravitationally attracted towards a relatively compact object. This situation is quite common in astrophysics and therefore the observational and theoretical study of accretion disks has been quite intensive. The necessary condition for accretion to take place is that angular momentum be extracted from the fluid swirling around the central object. To achieve accretion rates that are consistent with observations, the physical mechanism for such angular momentum transport must be more efficient (by many orders of magnitude) than just the one resulting from torques caused by microscopic viscosity.  Already at the outset, when AD were theoretically proposed (Prendergast \\& Burbidge 1968, Pringle \\& Rees 1972), the enhanced effective viscosity was postulated to result from turbulence and was parametrized, using mixing length theory or a similar scheme, as detailed theoretical understanding of turbulence and the transition to it was lacking then (a situation that is still with us nowadays). The parametrization of the effective viscosity in disks in terms of a single parameter - $\\alpha$, introduced by Shakura \\& Sunyaev (1973), proved itself to be the most fruitful, giving rise to successful interpretations of many observational results (see Lin \\& Papaloizou 1996, Frank, King \\& Raine 2002, for reviews). In some cases (e.g. dwarf nova models), however, more complex (and cumbersome) prescriptions for the viscosity had to be employed. However, until the early 1990's the question what is the physical origin of turbulence (or, more precisely, the anomalous angular momentum transport) in AD has essentially remained unanswered. No definite linear instability has been identified for thin disks in which the swirling flow consist of Keplerian shear. The magneto-rotational instability (MRI), originally found (Velikhov 1959, Chandrasekhar 1960) for magnetic Taylor-Couette, (i.e., cylindrical) flows, has been shown by Balbus \\& Hawley (1991) to also operate in cylindrical AD, when the fluid is electrically conducting and for not too large initial magnetic fields.  It has thus become the operating paradigm in the astrophysical community that purely hydrodynamical turbulence in Keplerian disks is altogether ruled out and since the problem of angular momentum transport in these objects relies on MRI driven magneto-hydrodynamical (MHD) turbulence, extensive numerical calculations are necessarily the main tool of research on this problem (see Balbus \\& Hawley 1998 and Balbus 2003, for reviews and references). Nonetheless, efforts to find a purely hydrodynamical transition to turbulent activity in thin AD have continued until the present time, more than a decade after Balbus , Hawley \\& Stone (1996) appeared to settle the matter. These have largely been motivated by the fact that a purely hydrodynamical AD flow has a `microscopic' Reynolds number ($\\rm Re$) of the order of $10^{14}$ or so, a quite amazing setting for a laminar flow.  Among the ideas that have been put forward in this context some are based on the application to thin disk flows of the viewpoint, familiar to the fluid-dynamical community, that {\\em transient dynamics induced by perturbations}, i.e. strong transient growth (TG) in linearly stable shear flows may play an important r\\^ole in nonlinearly shaping the final dynamical state. TG is possible because the relevant linear operator governing the behavior of infinitesimal perturbations in these flows is non-normal and thus the usual modal approach essentially fails (see, e.g., Grossman 2000, Schmid \\& Hennigson 2001, and Criminale, Jackson \\& Joslin 2003)  TG of disturbances has been discussed in the astrophysical literature within two quite different (but related) contexts: \\begin{description} \\item[{$\\bullet$}] Local disturbances experiencing large enough TG that can possibly trigger a subcritical nonlinear transition into turbulence in a linearly stable shear flow, i.e. via the so-called bypass transition (e.g., Chagelishvili {\\em et} al. 2003, Tevzadze {\\em et} al. 2003, Yecko 2004, Afshordi {\\em et} al. 2005, Mukhopadhyay {\\em et} al. 2005) \\item[{$\\bullet$}] Perturbations experiencing significant TG that can be excited either by some external agent, perhaps as secondary flows on a pre-existing 3D turbulence itself, however weak, may give rise to intense global dynamical activity (e.g. Ioannaou \\& Kakouris 2001, and see below). \\end{description} Other scenarios for inducing hydrodynamical activity in AD (some of them directly or indirectly related to TG as well) have also been proposed. Among them are those invoking baroclinic instabilities (e.g., Klahr \\& Bodenheimer 2003), strato-rotational instabilities (e.g., Dubrulle {\\em et} al. 2005, Umurhan 2006), formation and long sustenance of vortices and/or waves (e.g., Godon \\& Livio 1999, Bracco {\\em et} al. 1999, Li {\\em et} al. 2001, Umurhan \\& Regev 2004, , Barranco \\& Marcus 2005, Petersen, Stewart \\& Julien 2007, Bodo {\\em et} al. 2007, Lithwick 2007). Even though no definite and undisputed hydrodynamical mechanism, that can give rise to sufficient angular momentum transport in AD, has been identified so far, significant efforts along these lines are continuously being made by several groups of researchers.  It appears that continued study of purely hydrodynamical processes in disks remains viable and worthwhile, in particular in view of the difficulties with MRI driven angular momentum transport in AD, which have recently been pointed out. Questions of insufficent numerical resolution in MHD disk simulations have been convincingly raised (Pesah, Chen \\& Psaltis 2007, Fromang \\& Papaloizou 2007). The decrease of transport with decreasing magnetic Prandtl number (${\\rm Pm}$) for various setups and boundary conditions and, in particular, the vanishing of this transport for ${\\rm Pm}\\ll 1$ (which is the case in many instances of astrophysical AD) has been demonstrated (Umurhan, Menou \\& Regev 2007, Lesur \\& Longaretti 2007, Umurhan, Regev \\& Menou 2007, Fromang {\\em et} al. 2007). Finally, serious doubts as to the viability of the local shearing box approximation (Goldreich \\& Lynden-Bell 1965) to the numerical study of accretion disk MHD turbulence (Coppi \\& Keyes 2003, King, Pringle \\& Livio 2007, Shu {\\em et} al. 2007, Regev \\& Umurhan 2007) have been raised.  The purpose of this contribution is to draw attention to the possibility that a thin AD, in which a sub-critical transition to hydrodynamical turbulence occurs (for a sufficiently high ${\\rm Re}$) large transient amplification of global disturbances may induce recurrent or even persistent secondary flow activity. Lesur \\& Longaretti (2005) have explicitly demonstrated, using high resolution 3D numerical simulations, that such sub-critical transition to turbulence does appear in a model flow having Keplerian shear characteristics, but at a very high ${\\rm Re}$ (see also Rincon {\\em et} al. 2006). However, it appears that the efficiency of turbulent transport in these flows is insufficient for astrophysical purposes (i.e. AD). A similar conclusion can perhaps be drawn from the recent experimental study of Ji {\\em et} al. (2006). The very small value of the effective $\\alpha$ (measuring the angular momentum transport) in these flows is, as we shall show, the key necessary ingredient for the excitation of vigorous global (transient) secondary flows, atop the weak turbulent state. It is quite conceivable that the large (by orders of magnitude in the disturbance energy) TG may (nonlinearly) give rise to persistent dynamical activity, or at least be recurrently re-excited by external perturbations (see Ionnaou \\& Kakouris 2001, who advocated the latter possibility). It however remains to be shown, of course, that in the ultimate state angular momentum transport is appropriate for AD, i.e. the effective $\\alpha$ acquires a high enough value.  The primary tool utilized here in order to facilitate an analytical treatment is the asymptotic expansion, where the dependent variables and governing equations are expanded in powers of a small quantity (here the measure of the disk's `thickness', $\\epsilon$ - see below). Exposing the resulting mathematical system (an initial value problem) to a set of initial conditions, in which the extreme geometry of the disk structure is taken into account (i.e. $\\epsilon \\ll 1$), we  aimed at obtaining an analytically treatable problem. To achieve this goal we also assumed a polytropic relation between the pressure and density. The advantage of such an approach is obvious - the treatment can be essentially analytical and the responsible physical effects leading to any interesting dynamics may be transparently traced. The study reported on here is limited to axi-symmetric disturbances and the question of its ultimate development still remains open. We are now in the process of generalizing it to fully 3-D disturbances and it is clear that the present analytic work should be ultimately complemented with detailed and uncompromising 3-D numerical calculations with a proper treatment of energy generation and transfer. ",
        "conclusions": "\\label{4sec} We conclude with some remarks on possible improvements to the asymptotic analysis of the sort done here and prospects for the future, e.g., extensions to non-axisymmetric perturbations. Asymptotic expansions, when viable, are often very robust and provide a good approximation to the solution when truncation to only few leading terms is done. Obviously, when a term in the series becomes {\\em very} large it may `break its order', that is, become larger than a previous term and as such make the expansion invalid in this region. In our expansions successive terms ratios are of $\\order{\\epsilon^2}$, and thus the procedure's validity should not severely be limited even up to a growth factor of $\\sim 1000$ or so (in the velocity or density perturbations). This matter is further discussed in UNRS, here we just remark that the validity with respect to time for weakly viscous solutions are somewhat influenced by the thinness of the disk: smaller values of $\\epsilon$ mean that the solutions are valid for longer times after the initial disturbance. More importantly, however, for a given value of $\\epsilon$ one must not be too zealous or overreaching by attempting to infer the quantitative behavior of the disk for arbitrarily small values of $\\alpha$ - which, as one will recall, is formally assumed here to be an $\\order1$ quantity.  Despite these caveats, the procedure when carried to higher order introduces corrections which are technically non-linear.  Careful consideration must be undertaken in order to handle the response at these higher orders.  This may entail treating the disturbance amplitudes for the lower order solutions (like $A(r)$ in $u_{1}'$) as {\\em weakly non-linear} governed by a second `slow' time (e.g. the amplitude is instead written as $A = A(r,\\tau)$ where $\\tau = \\epsilon^2 t$) in a manner analogous to the treatment of non-linear thick polytropes (e.g, Balmforth \\& Spiegel, 1996).  The approach used here may be generalized in a number of directions. Allowing for non-axisymmetric perturbations, including the disk inner and outer boundary in some kind of boundary layer analysis and relaxing the polytropic assumption seem to be the most obvious generalizations.  We have found the presence of prominent TG in the simplest cases. It is difficult to imagine that it will be suppressed in the more general conditions although the effect of radiative energy losses on TG must be carefully examined.  The question concerning the ultimate fate of the transiently grown perturbations and their ability to induce a state of sustained complex dynamical activity in the disk remains open. In this context it is worthwhile to notice that since the TG decay times are of the order of hundreds of rotation periods, it is conceivable that AD, which are usually not isolated systems, may experience recurrent external perturbations on such time scales and in this way the dynamical activity may be sustained. Extensive numerical calculations of AD are however needed to decide if TG may lead, through non-linear processes, to sustained turbulence, or at least a dynamical state in which adequate angular momentum transport can be sustained. Such high-resolution global 3-D calculations are, however, still above the ability of the present computer power and it may be thus advantageous to also exploit sophisticated non-linear asymptotic methods to complement and guide them."
    },
    "0801/0801.0520_arXiv.txt": {
        "abstract": "In this paper, we present a correlation  between the  spectral index distribution (SED) and the dimensionless accretion rate defined as $\\dot{m}={L_{bol}/L_{Edd}}$ for AGN. This quantity  is used as a substitute of the physical accretion rate. We select 193 AGN with both broad H$\\alpha$ and broad H$\\beta$, and with absorption lines near MgI$\\lambda5175\\AA$ from SDSS DR4. We determine the spectral index and dimensionless accretion rate after correcting for both host galaxy contribution and internal reddening effects. A correlation is found between the optical spectral index and the dimensionless accretion rate for AGN, including low luminosity AGN ($L_{H\\alpha}<10^{41}{\\rm erg\\cdot s^{-1}}$ sometimes called \"dwarf AGN\" (Ho et al. 1997)). The existence of this correlation provides an independent method to estimate the central BH  masses for all types of AGN. We also find that there is a different correlation between the spectral index and the BH masses for normal AGN and low luminosity AGN, which is perhaps due to the different accretion modes in these two types of nuclei. This in turn may lead to the different correlations between BH masses and optical continuum luminosity reported previously (Zhang et al. 2007a), which invalidates the application of the empirical relationship found by Kaspi et al. (2000, 2005) to low luminosity AGN in order to determine their BLR sizes. ",
        "introduction": "Because of the difficulty to calculate the physical accretion rate to the BH in an AGN, a dimensionless accretion rate can be defined and estimated based on the bolometric luminosity and BH mass, $\\dot{m} = \\frac{L_{bol}}{L_{Edd}}$. This parameter can be used as the substitute for the actual accretion rate. Moreover, the dimensionless accretion rate is an important parameter in the scheme of the so called unified model for AGN (Antonucci 1993, Quintilio \\& Viegas 1997, Urry \\& Padovani 1995), and also the 4D Eigenvector 1 Scheme (e.g., Dultzin-Hacyan et al. (2007). The difference in  accretion rate leads to the principal difference between high luminosity QSOs and low luminosity Seyfert galaxies. It seems to be also one of the main physical parameters underlying the 4D Eigenvector 1 Scheme as explained recently in Dultzin-Hacyan et al. (2007). In order to obtain the dimensionless accretion rate $\\dot{m}$, two other parameters must be calculated first: BH mass $M_{BH}$ and the bolometric luminosity $L_{bol}$.  The common method to estimate the bolometric luminosity $L_{bol}$ of AGN is based on the continuum luminosity from the nuclei: $L_{bol}\\sim 9\\times L_{5100\\AA}$ given by Kaspi et al. (2000) and confirmed by Shang et al. (2005) for QSOs. Recently, the relation was used by Bonning et al. (2007) to study the correlation between the accretion disk temperatures and the continuum colors in QSOs. However, it should be stressed  that the relation does not hold for ALL kinds of AGN, in particular, for low luminosity AGN (Ho et al. 1997a, 1997b, Ho 1999), because of the different Spectral Energy Distribution (SED) (the lack of the Big Blue Bump).  Several methods are used to estimate BH masses of AGN. The most reliable method is based on the stellar velocity dispersion of the bulge of the host galaxy first presented by Ferrarese \\& Merritt (2000) and Gebhardt et al. (2000), then confirmed by Tremaine et al. (2002) and Merritt \\& Ferrarese (2001) etc. \\begin{equation} M_{BH} = 10^{8.13\\pm0.06}(\\frac{\\sigma}{200{\\rm km\\cdot s^{-1}}})^{4.02\\pm0.32} {\\rm M_{\\odot}} \\end{equation} which indicates a strong correlation between BH masses and bulge masses (H\\\"{a}ing \\& Rix 2004, Marconi \\& Hunt 2003, McLure \\& Dunlop 2002, Laor 2001, Kormendy 2001, Wandel 1999) etc. However we should note that the relation of $M_{BH} - \\sigma$ is obtained through the results of nearby inactive galaxies. Whether the relation can be applied to far away active galaxies is an interesting question. So far, there are a few dynamical mass estimates of central black holes of broad line AGN, and the BH masses are consistent with the BH masses estimated from the relation of $M_{BH} - \\sigma$, although the uncertainties are still large. In addition, we should say that the objects in our sample described in the following section are not high luminosity and high redshift QSOs, thus the correlation between central black hole and bulge of host galaxy can be reasonably considered to hold.  The other methods are based on the assumption of virialization of the Broad Line Emitting Regions (BLRs), or at least part of them (Peterson et al. 2004, Onken et al. 2004, Sulentic et al. 2006, Dultzin-Hacyan et al. 2007). In order to calculate the parameter of dimensionless accretion rate, the  BH mass is necessary. However, for high luminosity and high redshift AGN, it is difficult to measure the stellar velocity dispersions. The assumption of virialization is applied for QSOs. In order to estimate BH masses of QSOs based on the assumption of virialization, the most convenient way is to use the equation: \\begin{equation} \\begin{split} M_{BH} &= f\\times\\frac{R_{BLRs}\\times\\sigma_{b}^2}{G} \\\\ &= 2.15\\times10^8(\\frac{\\sigma_b}{3000{\\rm km\\cdot s^{-1}}})^2(\\frac{L_{5100\\AA}}{10^{44}{\\rm erg\\cdot s^{-1}}})^{0.69} {\\rm M_{\\odot}} \\end{split} \\end{equation} There are, however, some caveats with this method. First, the question of whether the relation $R_{BLRs}\\sim L_{5100\\AA}^{0.69}$ found by Kaspi et al. (2000, 2005) can be applied for all AGN, in particular high redshift ones. In an attempt to answer this question, we have found that the relation is not valid for some special kinds of AGN, such as the low luminosity AGN (Zhang, Dultzin-hacyan \\& Wang 2007a, Wang \\& Zhang 2003) and the AGN with double-peaked low ionization emission lines (Zhang, Dultzin-Hacyan \\& Wang 2007b).  Second, the estimation of the BH masses of high redshift AGN by means of Equation (2) will lead to BH masses larger than $10^{10}{\\rm M{\\odot}}$ (meaning $\\sigma>600{\\rm km\\cdot s^{-1}}$), which leads to  unreasonable masses of the bulge larger than $10^{13}{\\rm M_{\\odot}}$ (Netzer 2003, Sulentic et al. 2006). For this reason, finding another parameter which can be observationally determined, related to the dimensionless accretion rate is an important task and is the main objective of this paper. The accretion rate is determined by two properties: the continuum luminosity $L_{5100\\AA}$ and the BH mass $M_{BH}$. The continuum luminosity can be calculated from the observed spectra as discussed in the the next section. Thus, in order to obtain a reliable result, we select Equation (1) to estimate the central BH masses of AGN rather than Equation (2).  The accretion disk model has been widely accepted as the standard model for AGN. In the NLTE (Non Local Thermodynamic Equilibrium) accretion disk mode, the generated SED (Spectral Energy Distribution) is based on three main parameters: BH masses $M_{BH}$, accretion rate $\\dot{M}$ and the viscosity parameter $\\alpha$. An expected result is that there should be a correlation between the spectral index and the accretion rate $\\dot{m}$. In this paper, we answer the question whether the observed spectral index can be used to trace the  dimensionless accretion rate. In section II, we present the data sample. Section III gives the results. Finally  the discussion and conclusions are given in Section IV. In this paper, the cosmological parameters $H_{0}=70{\\rm km\\cdot s}^{-1}{\\rm Mpc}^{-1}$, $\\Omega_{\\Lambda}=0.7$ and $\\Omega_{m}=0.3$ have been adopted. ",
        "conclusions": "There are 38 low luminosity AGN with $L_{H\\alpha}<10^{41}{\\rm erg\\cdot s^{-1}}$ (Ho et al. 1997a, 1997b and Ho 1999), which are shown in solid circles in Figure 2. From the figure, we can see that there is no difference in the correlation between spectral index and $\\frac{9\\times L_{5100\\AA}}{L_{Edd}}$ for normal AGN and low luminosity AGN. If the bolometric luminosity of low luminosity AGN was different from $9\\times L_{5100\\AA}$, all the low luminosity AGN would deviate from the correlation for normal AGN, due (probably) to different accretion modes. Even if there is a different accretion mode for low luminosity AGN (as suggested by the lack of the big blue bump in the spectra of low luminosity AGN as shown in Ho (1999)), the bolometric luminosity of low luminosity AGN can also be calculated using $L_{bol}\\sim k\\times L_{5100\\AA}$. Otherwise, we could not find the same correlation between spectral index and $\\frac{9\\times L_{5100\\AA}}{L_{Edd}}$ for low luminosity AGN and normal AGN.  According to the accretion disk model, the output SED is the result of the convolution of other parameters as well, such as the central BH mass, the viscosity in the disk, and the inclination angle. However there is no correlation between the spectral index and  central BH masses, which is shown in Figure 5. The Spearman Rank Correlation Coefficient is less than 0.1 with $P_{null}>60\\%$ for all objects in our sample. An interesting result is that there is actually a negative trend (anticorrelation)  between BH masses and the spectral indexes for the 38 low luminosity AGN. The coefficient is about -0.54 with $P_{null}\\sim4.97\\times10^{-4}$, -0.52 with $P_{null}\\sim8.62\\times10^{-4}$ and -0.55 with $P_{null}\\sim3.98\\times10^{-4}$ for $\\frac{F_{5100\\AA}}{F_{6800\\AA}}$, $\\frac{F_{4400\\AA}}{F_{5100\\AA}}$ and $\\frac{F_{4400\\AA}}{F_{6800\\AA}}$ respectively.  Because of the positive correlation between the spectral index and the accretion rate, a negative correlation between the spectral index and  central BH masses could be expected for all AGN. However our results indicate that this expectation is only valid for low luminosity AGN. The reason is probably related to the correlation between BH masses and  continuum luminosity. For Normal AGN, there is strong correlation between the  BH masses and the continuum luminosity (Peterson et al. 2004). However for low luminosity AGN, this correlation is much weaker (Zhang, Dultzin-Hacyan \\& Wang 2007a). The correlation between the central BH masses and the internal continuum luminosity is shown in Figure 6. The coefficient is about 0.47 with $P_{null}\\sim9.05\\times10^{-10}$, however, the coefficient is only 0.12 with $P_{null}\\sim49\\%$ for the 38 low luminosity AGN.  The same result for low luminosity AGN can be found in Panessa et al. (2006). In their paper, they selected all the low luminosity Seyfert galaxies from Ho, Filippenko \\& Sargent (1997a, 1997b), and found that there is NO correlation between the X-ray or optical emission line luminosities (especially [OIII]$\\lambda5007\\AA$ line) and BH masses. The result also confirms that there is a different accretion mode for normal and low luminosity AGN.  To estimate the effects of the inclination angle of the accretion disk is difficult. However under the assumption that the narrow line emission region is isotropic, we can check the correlation between the continuum luminosity and the luminosity of narrow emission lines. If the objects have very different inclination angles of the accretion disk, a loose correlation between the continuum luminosity and the luminosity of narrow emission line should be expected. Here we show the correlation between $L_{5100\\AA}$ and the luminosity of narrow H$\\alpha$ in Figure 7. Although it is more common to use the [OIII]$\\lambda5007\\AA$ as an isotropic estimator of AGN luminosity, we prefer to use the narrow component of H$\\alpha$ because of the following reason. [OIII] emission line frequently cannot be fitted by one single gaussian function, because it has extended wings (Greene \\& Ho 2005a). The two components are not emitted from the same region. The extended component is probably emitted from the far-side of the BLRs. Thus when we fit the [OIII] line, two gaussian functions are applied as described in Section II. We thus select the narrow component of H$\\alpha$ rather than [OIII] to test the effects of inclination angle. A strong correlation can be confirmed. The spearman Rank Correlation Coefficient is about 0.89 with $P_{null}\\sim0$ for normal AGN, and about 0.67 with $P_{null}\\sim4.04\\times10^{-26}$ for the 38 low luminosity AGN. The best fit to the correlation (after considering the error in the determination of the luminosity of narrow H$\\alpha$) is given by: \\begin{equation} \\log(L_{H\\alpha_N}) = (1.373\\pm0.032) + (0.915\\pm0.003)\\times\\log{L_{5100\\AA}} {\\rm erg\\cdot s^{-1}} \\end{equation} This result indicates that the effects of the inclination angle can be neglected for the correlation between the spectral index and the dimensionless accretion rate.  The correlation between spectral index and dimensionless accretion rate found in this research provides another independent method to estimate the central BH masses of AGN. The spectral index and continuum luminosity can be directly determined from the observed spectrum in the optical band, and subsequently the Eddington Luminosity, i.e. BH masses, can be determined by means of the correlation we found. This method has the advantage of being  independent of the different correlations between the size of the BLRs and the continuum luminosity in Equation (2). Also, this method can be applied when it is not possible to measure the stellar velocity dispersion of the bulge. In future work, we will estimate the BH masses of QSOs with higher redshift using this method to solve the problem of why  virial BH masses of QSOs estimated by Equation (2) lead to BH masses larger than $10^{10} {M_{\\odot}}$, while observational results (fortunately) seem to contradict this result (Dultzin-Hacyan et al. 2007).  Finally, a simple summary is as follows. We first select 193 AGN with both broad H$\\alpha$ and broad H$\\beta$, and with apparent absorption MgI$\\lambda5175\\AA$ from SDSS DR4. Then after the determination of the spectral index (after the correction of internal reddening effects through the Balmer decrements for broad Balmer emission lines, and after the subtraction of stellar component) and dimensionless accretion rate ($\\frac{9\\times L_{5100\\AA}}{L_{Edd}}$), we find a strong correlation between these parameters for AGN, which provides another independent and method to estimate the central BH masses of AGN."
    },
    "0801/0801.2713_arXiv.txt": {
        "abstract": "{New candidate variable stars have been identified in the Small Magellanic Cloud cluster NGC121, by applying both the image subtraction technique (ISIS, Alard 2000) and the Welch \\& Stetson (1993) detection method to HST WFPC2 archive and ACS proprietary images of the cluster. The new candidate variable stars are located from the cluster's Main Sequence up to Red Giant Branch. Twenty-seven of them fall on the cluster Horizontal Branch and are very likely RR Lyrae stars. They include the few RR Lyrae stars already discussed by Walker \\& Mack (1988). We also detected 20 Dwarf Cepheid candidates in the central region of NGC121. Our results confirm the ``true\" globular cluster nature of NGC121, a cluster that is at the young end of the Galactic globulars' age range. ",
        "introduction": "NGC121 is a key cluster to understand the Star Formation History (SFH) occurred in the SMC through the study of the similarities and differences with respect to the Milky Way clusters. In fact, with an estimated age of about 10 Gyrs, it locates at the transition between open and globular clusters in the MW. Moreover, as emphasized by Stryker et al. (1987), NGC121 very likely marks the boundary between ``old'' Population II (containing RR Lyrae but not carbon stars) and intermediate-age populations (containing carbon stars but not RR Lyrae stars). Up to now, only four RR Lyrae stars with light curves derived from ground-based observations, were known in NGC121 (Walker \\& Mack 1988), and references therein). As part of a coordinated HST and ground-based effort, we have observed the SMC with the ACS (Cycle13, GO prog.10396, PI Gallagher), pointing at 7 star clusters of different age and metallicity (including NGC121) and 7 fields in various galactic locations (center, periphery, wing and bridge towards the LMC), to derive their key evolution parameters and SF histories (see e.g. Glatt et al. 2007, submitted). The ACS observations of NGC121 were taken in time-series fashion in order to use them, along with the existing WFPC2 archive data of the cluster, to detect new variable stars in the cluster. ",
        "conclusions": "\\begin{figure} \\includegraphics[width=6cm]{fig2_fiorentino_poster.ps} \\caption[]{Light curves of fundamental-mode RR Lyrae stars. {\\it Upper panel}: new RR Lyrae star with $P \\sim$ 0.63 days. {\\it Lower panel}: the fundamental-mode RR Lyrae star V37 discovered by Walker \\& Mack (1988).} \\label{map} \\end{figure}  The NGC121 candidate variable stars were identified by applying both the image subtraction technique (ISIS, Alard 2000) and the Welch \\& Stetson (1993) detection method. We identified about 50 candidate variables in NGC121. Period search was performed using GRaTiS (Graphical Analyzer of Time Series) a private software developed at the Bologna Observatory (see e.g. Clementini et al. 2000). Twenty-seven of the candidate variables are located on the cluster Horizontal Branch, thus are very likely RR Lyrae stars. They include the few RR Lyrae stars already discussed by Walker \\& Mack (1988). We also detected 20 Dwarf Cepheid candidates in the central region of NGC121, and recovered 34 Blue Straggler star candidates found by Shara et al. (1998). DCs are late-A and early-F type stars, that populate the instability strip near or slightly above the zero-age main sequence and have magnitude from 0.2 to about 3 mag fainter than the RR Lyrae stars. In the CMD they fall where the instability strip crosses the region of the Blue Straggler stars. These variables have typical periods in the range from 1 to 6 hours, thus our data sampling allow to sample their light curves better than for the RR Lyrae stars.  \\begin{figure} \\includegraphics[width=6cm]{fig3_fiorentino_poster.ps} \\caption[]{Light curves of two Dwarf Cepheids. {\\it Upper panel}: SX Phoenicis star with $P$=0.05 days. {\\it Lower panel}: $\\delta$ Scuti star with $P$=0.19 days. } \\end{figure}  DCs are divided into: 1) $\\delta$ Scuti stars, which are metal rich, young, Population I stars with typical periods longer than 0.1 day usually observed in open clusters and in the field of the MW; and 2) SX Phoenicis stars, which are metal poor, Population II variables, generally observed in GCs, with typical periods in the range from 0.03 to 0.09 days (Poretti et al. 2006). NGC121 seems to host both types of DCs. Fig. 1 shows the CMD of NGC121 for stars on the PC of the WFPC2, obtained using DAOPHOT/ALLFRAME for the data reduction. The variable stars are marked by different symbols (see caption). Examples of light curves for RR Lyrae stars and DCs are shown in Figs. 2-3.\\\\  The conspicuous number of RR Lyrae stars detected in NGC121 confirms the ``true\" globular cluster nature of NGC121, a cluster that is at the young end of the Galactic globulars' age range. The apparent presence of both $\\delta$ Scuti and SX Phoenicis stars in NGC121 is an intriguing feature, that needs to be confirmed by a more detailed analysis of the light curves and a careful study of the contaminations by the SMC field stars."
    },
    "0801/0801.3303_arXiv.txt": {
        "abstract": "A theoretical model framework of spherical symmetry is presented for a composite astrophysical system of two polytropic fluids coupled together by gravity to explore large-scale shocks and flow dynamics in clusters of galaxies or in globular clusters. The existence of such large-scale shocks in clusters of galaxies as inferred by high-resolution X-ray and radio imaging observations implies large-scale systematic flows that are beyond usual static models for clusters of galaxies. Here, we explore self-similar two-fluid flow solutions with shocks for a hot polytropic gas flow in a cluster of galaxies in the presence of a massive dark matter (DM) flow after the initiation of a gravitational core collapse or a central AGN activity or a large-scale merging process. In particular, the possibility of DM shocks or sharp jumps of mass density and of velocity dispersion in dark matter halo is discussed and such DM shocks might be detectable through gravitational lensing effects. To examine various plausible scenarios for clusters of galaxies, we describe three possible classes of shock flows within our model framework for different types of temperature, density and flow speed profiles. Depending upon sensible model parameters and shock locations, the hot ICM and DM halo may have various combinations of asymptotic behaviours of outflow, breeze, inflow, contraction or static envelopes at large radii at a given time. We refer to asymptotic outflows of hot ICM at large radii as the {\\it galaxy cluster wind}. As a result of such galaxy cluster winds and simultaneous contractions of DM halo during the course of galaxy cluster evolution, there would be less hot ICM within clusters of galaxies as compared to the average baryon fraction in the Universe. Physically, it is then expected that such `missing baryons' with lower temperatures reside in the periphery of galaxy clusters on much larger scales. Based on our model analysis, we also predict a limiting (the steepest) radial scaling form for mass density profiles of $r^{-3}$ within clusters of galaxies. ",
        "introduction": "Extensive X-ray observations have revealed that almost completely ionized hot gas medium permeates within clusters of galaxies with a typical temperature range of $\\sim 10^7-10^8$ K and a range of typical electron number density $\\sim 10^{-2}-10^{-4}\\hbox{ cm}^{-3}$ (e.g., Cavaliere \\& Fusco-Femiano 1978; Sarazin 1988; Fabian 1994). Clusters of galaxies are largely gravitationally bound systems on spatial scales of several Mpcs; together with the strong evidence of high velocity dispersions of galaxies (e.g., $\\sim 700-1000\\hbox{ km s}^{-1}$), hot X-ray emitting thermal electron gas, and gravitational lensing effects, we have realized the presence of massive dark matter halo within clusters of galaxies. Physical properties of galaxy clusters with static models of spherical symmetry have been extensively studied in the past (e.g., Lea 1975; Cavaliere \\& Fusco-Femiano 1976, 1978; Sarazin \\& Bahcall 1977; Sarazin 1988; Fabian 1994; Carilli \\& Taylor 2002; Voit 2005).  In the past several years, high-resolution X-ray imaging observations have revealed density, temperature and pressure jumps in the hot intracluster medium (ICM) within clusters of galaxies (e.g., Fabian et al. 2003; Nulsen et al. 2005a; Nulsen et al. 2005b; McNamara et al. 2005), indicating that these structures are likely large-scale shocks rather than cold fronts (e.g., Sanders \\& Fabian 2006). Active galactic nuclei (AGNs) (e.g., Nulsen et al. 2005a; Nulsen et al. 2005b; McNamara et al. 2005) and merging galaxies (e.g., Markevitch \\& Vikhlinin 2001; Markevitch et al. 2002; Gabici \\& Blasi 2003; Markevitch et al. 2005) are proposed to be the driving force and energy source of these large-scale shocks. In addition, large-scale sound waves in clusters of galaxies have been proposed by Sanders \\& Fabian (2007) to explain the observed quasi-concentric ripples in surface brightness of X-ray emissions. Once shocks are identified in clusters of galaxies, there must be large-scale flows involved. In other words, these clusters of galaxies cannot be really static on large scales at least in the spatial region where shocks are presumably identified. Relative to the cluster centre, radial distances of shocks observed vary from tens of kpcs (e.g., McNamara et al. 2005; Nulsen et al. 2005b) to several Mpcs (e.g., A3667 in Rotteringer et al. 1997; A3376 in Bagchi et al. 2006; galaxy clusters A786, A2255, A2256 in Ensslin et al. 1998); at yet smaller radii around the centre, there could be emerging shocks that may not be easily identified. We take the point of view that these shocks are moving in hot ICM and we simply catch them at different epochs of evolution. These shocks may all be born somewhere around the central region and travel outwards to the locations we observe at the present epoch. For those clusters of galaxies without shock signatures, one possibility is that shocks have occurred in the distant past and have disappeared after their energies were dissipated during the process of propagation. So large-scale shocks and flows may well be common phenomena in clusters of galaxies.  On much smaller scales compared to those of clusters of galaxies, X-rays have been also observed in globular clusters (e.g., Verbunt et al. 1984) and these emissions are interpreted by some as associated with flowing gas towards a black hole residing in the centre of globular clusters (e.g., Silk \\& Arons 1975). Moreover, outflows of gas materials from globular clusters have also been discussed in the literature (e.g., VandenBerg 1978). Fully relaxed globular clusters can be well treated as spherically symmetric (e.g., Harris \\& Racine 1979). One may view the collection of stars as one `fluid' and the tenuous gas as another fluid; these two fluids are coupled together by gravity on large scales. Here, we note globular clusters in passing and will mainly focus on large-scale self-similar dynamics for clusters of galaxies.  With these two classes of astrophysical systems in mind, we develop a theoretical model framework of spherical symmetry to study dynamic behaviours of hot ICM and dark matter halo using the two-fluid approximation (e.g., Lou 2005), where the two polytropic fluids are coupled together by gravity. We should note that here the notion of a polytropic fluid is fairly general in the sense of specific entropy conservation along streamlines (Lou \\& Cao 2007).  In the context of large-scale structure formation, extensive numerical works have been carried out to simulate the formation of galaxy clusters in the expanding universe, providing information for the hot ICM and dark matter halo (see, e.g., Bertschinger 1998 for a review on numerical simulations of structure formation in the universe). Evrard (1990) and Thomas \\& Couchman (1992) simulated properties (such as the number density and temperature profiles) of a hot gas in the presence of a dark matter halo. In Katz \\& White (1993) and Frenk et al. (1996), the radial cooling process is simulated. In particular, Evrard et al. (1994) simulated the formation of galaxies with two gravitationally coupled fluids representing dark matter halo and baryon matter, which is similar in essence to the approximation adopted in our semi-analytical model for clusters of galaxies but on much larger scales.  Various self-similar solutions describing hydrodynamic processes of a single self-gravitational isothermal or polytropic gas under spherical symmetry have been investigated previously in contexts of star formation (e.g., Larson 1969a; Larson 1969b; Shu 1977; Hunter 1977; Shu et al. 1987; Lou \\& Shen 2004). Very recently, asymptotic behaviours of novel quasi-static solutions in a single polytropic gas sphere with self-gravity have been reported by Lou \\& Wang (2006, 2007) and was utilized to model rebound (MHD) shocks in supernovae (Wang \\& Lou 2007). For an astrophysical system of two fluids coupled by gravity, we can systematically extend these self-similar solutions, especially the new quasi-static solution, which may be used to describe behaviours of hot ICM and dark matter halo in clusters of galaxies. Except for the gravitational effect in the Newtonian sense, nothing else is known about dark matter particles at present. Using the coupled two-fluid model, we might be able to learn physical properties of dark matter halo through detectable diagnostics of hot ICM and of gravitational lensing effects.  For clusters of galaxies, there is an outstanding problem of `missing baryons'. Extensive X-ray observations have indicated that the baryon mass fraction in clusters of galaxies is typically less than the prediction of primordial nucleosynthesis (e.g., Ettori \\& Fabian 1999; Ettori 2003; He et al. 2005; McGaugh 2007). This discrepancy becomes more difficult to reconcile in the cores of galaxy clusters (e.g., Sand et al. 2003). The best fit of cosmological parameters with tiny temperature fluctuations of the cosmic microwave background (CMB) radiation and large-scale structure clustering shows that relative to the critical mass density $\\rho_c$ in the universe, the mass density of baryon matter is $\\Omega_b=0.0224\\pm 0.0009 h^{-2}_{100}$ and the total matter density is $\\Omega_m=0.135^{+0.008}_{-0.009}h^{-2}_{100}$, where parameter $h_{100}$ is related to the Hubble constant $H_0$ by $H_0=100h_{100}$ km s$^{-1}$ Mpc$^{-1}$. Therefore, the mean cosmic baryon mass fraction is $f_b\\equiv\\Omega_b/\\Omega_m=0.166^{+0.012}_{-0.013}$ (e.g., He et al. 2005 and references therein). While there are different methods in determining the $f_b$ value, the cosmic baryon fraction $f_b$ is around 0.17 (e.g., McGaugh 2007). However in clusters of galaxies, the average gas (baryon) fraction inferred by two methods are about $0.107^{+0.028}_{-0.019}$ and $0.111^{+0.069}_{-0.063}$ (e.g., Ettori 2003). Others estimated that the baryon fraction $f_b$ observed in clusters of galaxies can be lower than the cosmic baryon fraction by about $10\\% - 20\\%$ at $z=0$ (e.g., He et al. 2005). Some even claimed that the value of $f_b$ can be lowered by as much as 30\\% (e.g., Ettori 2003). In conclusion, the baryon fraction in most clusters of galaxies are systematically lower than the average cosmic value $f_b$ except those highest estimates for gas (baryon) mass fraction in some clusters of galaxies (e.g., A426, A2142, RXJ1350; see Ettori 2003). To resolve this important issue, the notion of Warm-Hot Intergalactic Medium (WHIM) has been introduced (e.g., Cen \\& Ostriker 1999, 2006; Ettori 2003). In their opinion, the WHIM may actually exist within clusters of galaxies to account for the mass of `missing baryons', yet the WHIM cannot be detected at present because it does not emit X-rays. These results show that a significant fraction ($\\sim 40\\%-50\\%$) of the baryon component might be found in the form of WHIM in the temperature range of $T\\sim 10^{5-7}$ K (e.g., Cen \\& Ostriker 2006).  As will be discussed in more details, this problem of `missing baryons' in our model scenario is a natural consequence of galaxy cluster winds, be it sustained or sporadic or be it stationary or dynamic during the evolution of galaxy clusters. These so-called `missing baryons' are blown away in the form of hot ICM and cool down gradually with time; with relatively low temperatures, they should mostly reside in the periphery of galaxy clusters and spread out in space on much larger scales. Meanwhile, the dark matter halo may contract within clusters of galaxies in our model. Therefore the mass fraction of baryons $f_b$ (i.e., the mass ratio of total baryons to the total gravitational mass inferred) would be lower than the initial value when a cluster of galaxies was born and started to evolve. The age of galaxy clusters is estimated to fall in the range of $\\sim 10^9\\ -\\ 10^{10}$ yr (e.g., Fabian 1994). As galaxy cluster winds may have existed since galaxy clusters were born, the timescale of galaxy cluster winds would be comparable to or somewhat less than this estimate. In Section 3, we show a few specific examples of numerical shock flow solutions in our model and estimate the loss of baryons within a timescale of $\\sim 10^9$ yr.  As different behaviours of temperature profiles have been inferred from X-ray observations of galaxy clusters (Markevitch 1996 and Markevitch et al. 2005 for decreasing temperatures with increasing radius; Peres et al. 1998 and Sanders \\& Fabian 2006 for nearly constant temperatures in several galaxy clusters; McNamara et al. 2005 and Blanton et al. 2001 for increasing temperatures with increasing radius) and electron number densities are observed to fit a power law fairly well (e.g., Peres et al. 1998; Nulsen et al. 2005b), we shall take the specific entropy conservation along streamlines as the equation of state and see how well this may account for the various observed profiles of thermodynamic variables. By properly choosing model parameters in various regimes, we can describe properties of galaxy clusters to a considerable extent.  This paper is structured as follows. The background and motivation of our model development is introduced in Section 1. Section 2 presents in order the basic formulation for the two-fluid model of spherical symmetry involving two polytropic fluids, self-similar transformation, asymptotic solutions at small and large $x$, singular surfaces and sonic critical curves, and shock conditions. In Section 3, we show numerical examples of quasi-static solutions for three different situations. The major results are summarized in Section 4. Finally we discuss our model results and numerical solutions in Section 5. Certain mathematical details are contained in Appendices A through G for the convenience of reference. ",
        "conclusions": "In this paper, we formulate a two-fluid model framework of spherical symmetry to explore dynamic behaviours of hot ICM and dark matter during the evolution of galaxy clusters. In this scenario, the hot ICM and dark matter halo are approximated as two polytropic `fluids' and are coupled by gravity. For both `fluids' in general situations, specific entropies are conserved along streamlines separately and are related to the enclosed masses. Quasi-static solutions for both `polytropic fluids' can be obtained and are adopted for sufficiently small radii around the central core region at a given time. In order to construct dimensionless quasi-static solutions in the regime of small $x$, we need to specify nine dimensionless parameters $\\{n,\\ \\gamma_1,\\ \\gamma_2,\\ \\kappa,\\ L_1,\\ L_2,\\ x_{ini},\\ x_{d,1},\\ x_{d,2}\\}$ where $x_{ini}$ needs to be properly chosen. Two additional dimensional parameters $K_2$ and $t$ need to be specified for dimensional solutions for a physical description. The ICM temperature is taken to be on the order of $\\sim 10^7-10^8$ K in typical clusters of galaxies and the electron number density is taken to be $\\sim 10^{-2}-10^{-4}$ cm$^{-3}$ in the typical radial range of kpcs to Mpcs. It is possible to construct different types of flow solutions with shocks in both fluids and with various asymptotic scaling features in flow speeds, mass densities, enclosed mass and temperatures. In particular, we can construct dynamic solutions for galaxy cluster winds and discuss the important problem of `missing baryons' during the evolution of galaxy clusters. There are several physical hypotheses in our two-fluid model to simplify the mathematical analysis. First, we take the dark matter halo as a kind of `fluid' for simplicity and explore this alternative theoretical possibility. Secondly, we assume the two-fluid system of galaxy clusters to be grossly spherically symmetric with a common centre and hope to catch major dynamic flow features on large scales. Thirdly, both `fluids' are assumed to be `polytropic' in the most general sense of entropy conservation along streamlines. In other words, the specific entropy distribution is not necessarily constant in space and time but are allowed to vary in $r$ and $t$ in general. Finally, dark matter interacts with the hot ICM only through gravity.  Global semi-complete solutions can be constructed to pass through the two singular surfaces via shocks in both fluids. As large-scale ICM shocks have been identified observationally, there must be large-scale flows of hot ICM. As the dark matter halo and the hot ICM are coupled by gravity, there may be dark matter flows and the possibility of shocks in dark matter halo. Such dark matter shocks are characterized by drastic density jumps and sharp rises of velocity dispersions and may be detected by utilizing gravitational lensing effects. In our model framework, outflows of hot ICM in galaxy clusters actually form galaxy cluster winds, which is a systematic mechanism of reducing the baryon fraction $f_b$. In our model, the self-similar shocks travel outwards in both hot ICM and dark matter halo, respectively. The travel speeds of these shocks actually change with time for $n\\neq 1$ and their radial positions at present can vary from tens of kpcs to a few Mpc. Due to galaxy cluster winds, the baryon fraction $f_b$ in galaxy clusters can be $\\sim 15\\%-40\\%$ lower than the average cosmic baryon fraction in the Universe, which may account for the problem of `missing baryons' in clusters of galaxies. Physically, these `missing baryons' should reside in the periphery of galaxy clusters in the form of warm gas as results of unavoidable radiative cooling. Since the lower baryon fraction $f_b$ is a generic phenomenon in clusters of galaxies, we therefore suggest that galaxy cluster winds would be common and frequent during the evolution of galaxy clusters.  The main features of our self-similar polytropic solution for two fluids coupled by gravity are summarized below. The radial profile of mass density at a given time is $\\rho_i\\propto r^{-2/n}$ in both fluids for either $r\\rightarrow 0^{+}$ or $r\\rightarrow+\\infty$ with $2/3<n<2$ in general. That shows that the asymptotic mass density profiles of hot ICM and dark matter take the same form of power laws with the same index $-2/n$. Meanwhile, the radial flow velocities of both fluids approach zero as $x\\rightarrow 0^{+}$ in the form of quasi-static asymptotic solution (\\ref{quasiv}) and (\\ref{quasia}). At large $x$, the asymptotic flow solution is (\\ref{largea1}), (\\ref{largea2}), (\\ref{largev1}), (\\ref{largev2}) for $2/3<n<2$. When $2/3<n<1$ and $x\\rightarrow+\\infty$, the asymptotic radial flow velocities of both fluids approach zero. When $n=1$ and $x\\rightarrow+\\infty$, the radial flow velocities of both fluids approach constant values, which may be positive, zero or negative for various combinations of two fluids. When $n>1$ and $x\\rightarrow+\\infty$, the radial flow velocities of both fluids become divergent for $H_1\\neq 0$ and $H_2\\neq 0$ in asymptotic solution (\\ref{largev1}) and (\\ref{largev2}). For $H_1=H_2=0$, the radial flow velocities of both fluids remain finite at large $x$ for $n>1$. The radial profile of temperature in the hot ICM is $T_2\\propto r^{2-2/n}$ for either $r\\rightarrow 0^{+}$ or $r\\rightarrow+\\infty$. By mass conservation, the enclosed mass is continuous across shocks."
    },
    "0801/0801.4465_arXiv.txt": {
        "abstract": "Published data for large amplitude asymptotic giant branch variables in the Large Magellanic Cloud are re-analysed to establish the constants for an infrared ($K$) period-luminosity relation of the form: $M_K=\\rho[\\log P-2.38] + \\delta$. A slope of $\\rho=-3.51\\pm0.20$ and a zero point of $\\delta=-7.15\\pm0.06$ are found for oxygen-rich Miras (if a distance modulus of $18.39\\pm0.05$ is used for the LMC). Assuming this slope is applicable to Galactic Miras we discuss the zero-point for these stars using the revised {\\it Hipparcos} parallaxes together with published VLBI parallaxes for OH Masers and Miras in Globular Clusters. These result in a mean zero-point of $\\delta=-7.25\\pm0.07$ for O-rich Galactic Miras.  The zero-point for Miras in the Galactic Bulge is not significantly different from this value.  Carbon-rich stars are also discussed and provide results that are consistent with the above numbers, but with higher uncertainties. Within the uncertainties there is no evidence for a significant difference between the period-luminosity relation zero-points for systems with different metallicity. ",
        "introduction": "Large amplitude Asymptotic Giant Branch (AGB) variables (Miras) are important distance indicators for old and intermediate age populations. They are luminous, both bolometrically and in the near-infrared, and easily identified by their late spectral types (Me, Ce, Se or very rarely Ke), large amplitudes ($\\Delta V > 2.5$ mag, $\\Delta K > 0.4$ mag) and long periods ($100 \\lsim P \\lsim 1000$). The increasing use of adaptive optics on large telescopes at near-infrared wavelengths to study stellar populations (e.g. Da Costa 2004) at large distances will require confidence in the calibration of the AGB variables as distance indicators. Alternatively, if the distance is known the luminosities of the brightest AGB stars provide insight into the intermediate age populations, e.g. Menzies et al. (2008).  Wood et al. (1999) demonstrated that, within the LMC, AGB variables fall on a series of, approximately parallel, period-luminosity (PL) relations at $K$, (see also Cioni et al. 2001, Ita et al. 2004, Fraser et al. 2005, Soszy\\`nski et al. 2007). The large amplitude variables, i.e. the Miras, however, lie only on a singe PL($K$) relation. The existence of a Mira PL($K$) relation has been known for some while (Feast et al. 1989; Hughes \\& Wood 1990) and is now generally thought to represent the relation for fundamental pulsation. While some of the low amplitude variables also pulsate in the fundamental, others show various overtones. Low amplitude variables are of limited use for distance scale studies as it is not simple to establish upon which of the various PL relations they lie.  Complications do arise at periods in excess of about 400 d where some Miras in the LMC have higher luminosities, possibly as a consequence of hot bottom burning (e.g. Whitelock et al. 2003). Feast (2004) provides a recent discussion of AGB stars as distance indicators and of the zero-point of the PL relation.  Following a brief re-examination of the Mira PL($K$) relation in the LMC, we take advantage of the new analysis of the {\\it Hipparcos} data (van Leeuwen 2007a and 2007b, see also van Leeuwen 2005 and van Leeuwen \\& Fantino 2005) to re-examine the distance scale for large amplitude AGB variables within the Galaxy. Analysis of the original {\\it Hipparcos} data was discussed by van Leeuwen et al. (1997), Whitelock \\& Feast (2000 - henceforth Paper~I) and Knapp et al. (2003).  In the last part of this paper we put together all the available information on Mira parallaxes to establish the best value for the zero-point of the Mira PL($K$) relation for the Galaxy and compare this with values for elsewhere.  \\begin{table} \\centering \\caption{LMC variables used to establish the slope for the Mira PL($K$) relation}\\label{lmcdata} \\begin{center} \\begin{tabular}{lccc} \\hline \\multicolumn{1}{c}{name}& P & $K$ & note\\\\ & \\multicolumn{1}{c}{(d)} &  \\multicolumn{1}{c}{(mag)} \\\\ \\hline \\multicolumn{4}{c}{O-rich stars}\\\\ 0517--6551     & 116 & 12.25  & \\\\ C38            & 130 & 12.12  & \\\\ 0512--6559     & 141 & 12.13  & \\\\ W132           & 156 & 11.67  & \\\\ 0526--6754     & 157 & 11.79  & \\\\ W151           & 174 & 11.74  & \\\\ W148           & 183 & 11.82  & \\\\ W158           & 194 & 11.77  & \\\\ 0528--6531     & 195 & 11.48  & \\\\ C11            & 202 & 11.51  & \\\\ GR13           & 202 & 11.59  & \\\\ SHV05220--7012 & 205 & 11.73  & 1\\\\ 0507--6639     & 208 & 11.57  & \\\\ C20            & 210 & 11.54  & \\\\ W77            & 213 & 11.25  & S\\\\ R120           & 217 & 11.38  & \\\\ W94            & 220 & 11.28  & \\\\ W74            & 231 & 11.49  & \\\\ WBP74          & 233 & 11.50  & 1\\\\ W1             & 235 & 11.48  & \\\\ W140           & 243 & 11.19  & \\\\ 0533--6807     & 247 & 11.38  & \\\\ R141           & 258 & 10.99  & \\\\ R110           & 261 & 11.29  & \\\\ W48            & 279 & 10.99  & \\\\ 0537--6607     & 284 & 11.02  & \\\\ 0505--6657     & 307 & 10.67  & \\\\ 0524--6543     & 315 & 10.71  & \\\\ W126           & 318 & 10.89  & K\\\\ SHV05305--7022 & 362 & 10.57  & \\\\ R105           & 413 & 10.33  & \\\\ \\multicolumn{4}{c}{C-rich stars}\\\\ 0530--6437     & 157 & 12.08  & \\\\ 0515--6617     & 226 & 11.16  & \\\\ 0528--6520     & 229 & 11.08  & \\\\ 0520--6528     & 233 & 11.28  & \\\\ 0519--6454     & 242 & 11.09  & \\\\ W220           & 281 & 10.83  & \\\\ 0529--6759     & 283 & 10.91  & \\\\ 0515--6451     & 284 & 10.81  & \\\\ SHV05027--6924 & 298 & 10.82  & \\\\ 0514--6605     & 308 & 10.64  & \\\\ 0534--6531     & 308 & 10.98  & \\\\ 0529--6739     & 319 & 10.60  & \\\\ 0502--6711     & 322 & 10.53  & \\\\ C7             & 327 & 10.69  & \\\\ 0541--6631     & 342 & 10.50  & \\\\ R153           & 347 & 10.52  & \\\\ WBP14          & 351 & 10.62  & \\\\ W103           & 363 & 10.78  & \\\\ 0515--6438     & 365 & 10.90  & \\\\ 0537--6740     & 367 & 10.47  & \\\\ SHV05003--6817 & 369 & 10.58  & \\\\ SHV05260--7011 & 373 & 10.54  & \\\\ \\hline \\multicolumn{4}{l}{notes 1. Period taken from MACHO;}\\\\ \\multicolumn{4}{l}{S and K are spectral types}\\\\ \\hline \\end{tabular} \\end{center} \\end{table}  \\begin{figure*} \\includegraphics[width=16cm]{lmcpl.ps} \\caption{The PL($K$) relation for Mira variables in the LMC. Solid and open symbols are O- and C-rich respectively. The dashed line is the fit to the O-rich stars from Table~\\ref{lmcpl} assuming a distance modulus for the LMC of 18.39 mag, while the dotted line is the relation for C-stars.} \\label{fig_lmcpl} \\end{figure*} ",
        "conclusions": "In Paper~I (section 2.2) we noted a caveat with regard to temporal changes in the light distribution across the stellar disk as problematic for the interpretation of the parallaxes. This is particularly so because the angular diameters of the Miras are two or three times larger than their parallaxes. More recent work (e.g. Ragland et al. 2006; Woodruff et al. 2008) confirms that the diameters are large, variable and non-uniform. These same references further highlight the disagreement between observation and theory, emphasizing our very limited understanding of the atmospheres of these stars or their variations. This is clearly a real problem, but the agreement on the distance scale achieved by the different methods summarized above may indicate that the net effect is small.  As noted, the various results discussed above are in good agreement with each other. UX Cyg, is brighter than we would expect from the PL($K$) relation and this may be because of hot bottom burning. We take a simple mean of solutions 10, 25, 26 from Table~\\ref{zp}, to set a mean value of $\\delta = -7.25 \\pm 0.07$ for the O-rich Mira variables in the Galaxy. This is in reasonable agreement with the LMC value of $\\delta = -7.15 \\pm 0.06$.  Solution 24, for the C-rich Miras, of $\\delta = -7.18 \\pm 0.37$ is in agreement with the above result for O-rich Miras and with the LMC value for C-stars of $\\delta = -7.24 \\pm 0.07$ This supports an identical PL($K$) relation for O- and C-rich variables.  The O-rich result represents the best value for large amplitude variables, but it should be used with caution on stars with periods in excess of 400 days. The $\\rm H_2O$ parallaxes offer the most likely significant improvement in this result in the near future."
    },
    "0801/0801.1847_arXiv.txt": {
        "abstract": "Standard calculations suggest that the entropy of our universe is dominated by black holes, whose entropy is of order their area in Planck units, although they comprise only a tiny fraction of its total energy. Statistical entropy is the logarithm of the number of microstates consistent with the observed macroscopic properties of a system, hence a measure of uncertainty about its precise state. Therefore, assuming unitarity in black hole evaporation, the standard results suggest that the largest uncertainty in the future quantum state of the universe is due to the Hawking radiation from evaporating black holes. However, the entropy of the matter precursors to astrophysical black holes is enormously less than that given by area entropy. If unitarity relates the future radiation states to the black hole precursor states, then the standard results are highly misleading, at least for an observer that can differentiate the individual states of the Hawking radiation. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.1076_arXiv.txt": {
        "abstract": "We report on next phase of our study of rotating accretion flows onto black holes. We consider hydrodynamical (HD) accretion flows with a spherically symmetric density distribution at the outer boundary but with spherical symmetry broken by the introduction of a small, latitude-dependent angular momentum. We study accretion flows by means of numerical two-dimensional, axisymmetric, HD simulations for variety of the adiabatic index, $\\gamma$ and the  gas temperature at infinity, $c_\\infty$. Our work is an extension of work done by Proga \\& Begelman who consider models for only $\\gamma=5/3$. Our main result is that the flow properties such as the topology of the sonic surface and time behavior strongly depend on $\\gamma$ but little on $c_\\infty$. In particular, for $1 < \\gamma < 5/3$, the mass accretion rate shows large amplitude, slow time-variability which is a result of mixing between slow and fast rotating gas. This temporal behavior differs significantly from that in models with $\\gamma\\simless 5/3$ where the accretion rate is relatively constant and from that in models with $\\gamma\\simgreat 1$ where the accretion exhibits small amplitude quasi-periodic oscillations. The key parameter responsible for the differences is the sound speed of the accretion flow which in turn determines whether the flow is dominated by gas pressure, radiation pressure or  rotation. Despite these differences the time-averaged mass accretion rate in units of the corresponding Bondi rate is a weak function of $\\gamma$ and $c_\\infty$. ",
        "introduction": "Many types of astrophysical objects are powered by black hole (BH) accretion. Radiation energy produced by accretion can be very high and can explain such dramatic phenomena as quasars, powerful radio galaxies, X-ray binaries, and gamma ray bursts (GRBs). However, BH accretion does not always result in high radiative output. This is true for both stellar BH and super massive black holes (SMBH). In particular, objects with SMBH appear to spend statistically most of their time in an inactive phase. Inactive SMBHs are not something that one would expect, because these black holes are embedded in the relatively dense environments of galactic nuclei. Therefore it is natural to suppose that the gravity due to an SMBH will draw in matter at high rates, leading to a high system luminosity. Monitoring of X-ray binaries with stellar BHs reveals that these objects often exhibit large time variability in the total energy output in the spectral energy distribution. Then one of the main goals of any theory of BH accretion is two explain why accretion proceeds through very different modes.   Generally, the radiative output from accretion depends on the mass accretion rate $\\MDOT_a$ and an efficiency factor, $\\eta$. In previous papers of this series (Proga \\& Begelman 2003a,b, hereafter PB03a and PB03b, respectively), we studied how physical conditions at large distances from SMBH affect $\\MDOT_a$ in the so-called radiatively inefficient accretion flows (RIAF). RIAF with very low $\\eta$ and also low $\\MDOT_a$ have been proposed to explain very low radiative luminosities in systems as such Sgr~A* (e.g., Ichimaru 1977; Rees et al. 1982; Narayan \\& Yi 1994, 1995; Abramowicz et al. 1995; Blandford \\& Begelman 1999; Sharma et al. 2007 and reference therein). In PB03a, we addressed the issue of how $\\MDOT_a$ depends on the distribution of specific angular momentum, $l$ at large radii assuming that the adiabatic index, $\\gamma=5/3$.  For high $l$, corresponding to the circularization radius larger than the last stable orbit, gas cannot directly accrete onto a BH, unless some physical mechanism like e.g., viscosity or magnetic fields, or both, transports angular momentum towards.  For inviscid accretion flow the matter with too high $l$, either flows outward due to gas pressure and centrifugal forces or accumulates close the central accretor (e.g. Hawley et al. 1984a,b,; Clarke et al. 1985; Chen et al 1997; PB03a; Janiuk $\\&$ Proga 2007). The simulations presented by PB03a illustrate a general flow pattern with an inflow in the polar funnel, and an equatorial outflow (Fig.~1 there).  Such flow pattern is induced by the angular momentum distribution simply because rotating gas tend to converge toward the equator. If $l$ is high, the gas forms a subsonic very dynamic torus with gas flowing out. The mass accretion rate can be significantly smaller compared to corresponding the Bondi rate.  In this work, we address the problem how the properties of the accretion flow presented in PB03a, change when different micro-physical properties of the flow are assumed.  In particular, we explore effects of changing $\\gamma$ in the polytropic equation of state.  In reality, accretion flows with different $\\gamma$'s, may correspond to different types of objects or different phases of activity.  For example, in the weakly active galaxies like e.g. Sgr A*,  $\\gamma$ is usually considered to be around 5/3, because these flows are believed to be radiatively inefficient and gas pressure dominated. On the other hand, the large energetic output observed in GRBs indicates that an accretion flow must be radiation pressure dominated (e.g., Meszaros 2006)  and a relativistic equation of state is required with $\\gamma$=4/3.  Another example are protogalactic disks, which are sometimes considered as being formed by isothermal accretion flow ($\\gamma\\approx 1$; e.g., Mo et al. 1998).  Also, the so-called `high' and `low' accretion states in the X-ray binary systems, may reflect physical proprieties changing in time (for a review see Done et al. 2007).  Our work is a straightforward extension of PB03a's work, who showed results for $\\gamma$=5/3.  We extend their models by considering the flows, for  $\\gamma$ ranging between 1 and 5/3.  We also perform additional calculations with sound speed at infinity $c_{s,\\infty}$ much smaller than that considered by PB03a allowing us to model flows with the Bondi radius as large as that estimated in real systems, for example in Sgr~A*. Taking advantage of faster computers, we perform simulations with higher resolution in the $\\theta$ direction, and with much larger computational domain in comparison to PB03a. We keep other model parameters such as angular momentum distribution as in PB03a.   For $\\gamma=5/3$, our new simulations are consistent with those presented by PB03a. However, we find that $\\MDOT_a$, the flow structure, defined by the gas density and angular momentum distribution, and  the sonic surface topology depend on $\\gamma$. In particular,for $\\gamma=5/3$, $\\MDOT_a$ is nearly constant whereas for $\\gamma=4/3$, $\\MDOT_a$ shows stochastic, large-amplitude time-variability. For $\\gamma=1.01$, $\\MDOT_a$ shows small-amplitude periodic changes. In Sec.2, we describe a general set-up of our simulations. In Sec.3, we present our results. In the last section, we discuss the results, and relate them to results found in previous studies. ",
        "conclusions": "We report on next phase of our study of rotating accretion flows onto black holes. We consider hydrodynamical (HD) accretion flows with a spherically symmetric density distribution at the outer boundary but with spherical symmetry broken by the introduction of a small, latitude-dependent angular momentum. We study accretion flows by means of numerical two-dimensional, axisymmetric, HD simulations for variety of the adiabatic index, $\\gamma$ and the  gas temperature at infinity, $c_\\infty$. Our work is an extension of work done by PB03a who consider models for only $\\gamma=5/3$. We also reran some models from PB03a with higher resolution. The latter are fully consistent with lower resolution PB03a's runs.  Our main result is that the flow properties such as the topology of the sonic surface and time behavior strongly depend on $\\gamma$ but little on $c_\\infty$. In particular, for $1 < \\gamma < 5/3$, the mass accretion rate shows large amplitude, slow time-variability which is a result of mixing between slow and fast rotating gas. This temporal behavior differs significantly from that in models with $\\gamma\\approx 5/3$ where the accretion rate is relatively constant and from that in models with $\\gamma\\approx 1$ where the accretion exhibits small amplitude quasi-periodic oscillations. The key parameter responsible for the differences is the sound speed of the accretion flow which in turn determines the strength of shocks and whether the flow is dominated but gas pressure ($\\gamma\\simless 5/3$), radiation pressure ($1 < \\gamma < 5/3$) ,  or  rotation ($\\gamma\\simgreat 1$). Despite these differences, the time-averaged mass accretion rate in units of the corresponding Bondi rate is a weak function of $\\gamma$ and $c_\\infty$.  We realize that our simulations do not capture several physical processes that may be important in real systems (e.g., magnetic fields, radiative cooling and heating, energy dissipation). However, our simulations are done within a general framework so that some of these processes can be straightforwardly added (e.g., magnetic fields as in PB03b, or radiative processes as in Proga 2007). Therefore, this work could serve as a good reference point to analyze and interpret more complete and complex simulations (e.g., we already reran some of the models from this paper including magnetic fields, Moscibrodzka and Proga, in preparation).    ACKNOWLEDGMENTS: We thank Bozena Czerny and Marek Abramowicz for very useful comments. We acknowledge support provided by the Chandra award TM7-8008X issued by the Chandra X-Ray Observatory Center, which is operated by the Smithsonian Astrophysical Observatory for and on behalf of NASA under contract NAS8-39073. M.M also acknowledges supported in part by grant 1P03D~008~29 of the Polish State Committee for Scientific Research (KBN). While D.P. acknowledges support from NASA under ATP grant NNG05GB68G.    \\newpage"
    },
    "0801/0801.0593_arXiv.txt": {
        "abstract": "Simple one--zone homogeneous synchrotron self--Compton models have severe difficulties is explaining the TeV emission observed in the radiogalaxy M87.  Also the site the TeV emission region is uncertain: it could be the unresolved jet close to the nucleus, analogously to what proposed for blazars, or an active knot, called HST--1, tens of parsec away. We explore the possibility that the TeV emission of M87 is produced in the misaligned subpc scale jet. We base our modelling on a structured jet, with a fast spine surrounded by a slower layer. In this context the main site responsible for the emission of the TeV radiation is the layer, while the (debeamed) spine accounts for the emission from the radio to the GeV band: therefore we expect a more complex correlation with the TeV component than that expected in one--zone scenarios, in which both components are produced by the same region.  Observed from small angles, the spine would dominate the emission, with an overall Spectral Energy Distribution close to those of BL Lac objects with a synchrotron peak located at low energy (LBLs). ",
        "introduction": "An increasing number of extragalactic objects is detected at energies greater than 100 GeV. The list of published sources comprises 20 objects\\footnote{see {\\tt http://www.mppmu.mpg.de/$\\sim$rwagner/sources}}: as expected (e.g., Costamante \\& Ghisellini 2002), the bulk of them (17) belongs to the class of Highly Peaked BL Lac objects. In fact, the position of the synchrotron peak, usually located in (or close to) the X--ray band assures the existence, in these sources, of electrons with extremely large Lorentz factors ($\\gamma=10^5$--$10^6$), a key ingredient for the emission in the TeV band via inverse Compton (IC) scattering. The remaining three sources are BL Lac itself, which belongs to the LBL class, (Albert et al. 2007), 3C279, a FSRQs, (Teshima et al. 2007) and the nearby (16 Mpc, Tonry 1991) radiogalaxy M87 (Aharonian et al. 2003, 2006).  Already suggested as a possible source of high-energy radiation (e.g. Bai \\& Lee 2001) based on its similarity with BL Lac objects (Tsvetanov et al. 1998), M87 was discovered as a TeV source by the HEGRA array of Cherenkov telescopes (Aharonian et al. 2003) and it has been recently confirmed by H.E.S.S. (Aharonian et al. 2006). Due to the limited spatial resolution it is not possible to identify the emission region of this radiation. However, the sensitivity of H.E.S.S. allowed the detection of variability on short timescales ($\\sim$ days), suggesting a compact emission region.  Three main sites have been proposed as emission regions for the TeV radiation: the resolved jet (in particular the so--called knot HST--1); the unresolved base of the jet (in analogy with blazars); the vicinity of the supermassive ($M=3\\times 10^9$ $M_{\\odot}$, Marconi et al. 1997) black hole (BH).  HST--1 is a quite interesting jet feature located at 60 pc (projected) from the core of of M87. Superluminal motions with apparent speeds up to $\\sim 6 c$ were observed with {\\it HST} (Biretta et al. 1999). Continuous multifrequency coverage has shown spectacular activity from this knot in the past years (Perlman et al. 2003; Harris et al. 2006, 2007; Cheung et al. 2007). In particular, the X--ray flux increases by a factor of $\\sim 30$ from 2000 to 2005 and since then it is steadily decreasing. These variations are accompanied by similar changes in the optical and in the radio bands and by the appearance of new superluminal components. This extreme phenomenology is described by Stawarz et al. (2006) on the assumption that HST-1 marks the recollimation shock of the jet. As such, HST--1 is thought to be a rather efficient particle accelerator and thus a possible source of intense TeV radiation. However, the small variability timescales observed by H.E.S.S. put strong limits to the size of the source ($R < c \\Delta t \\, \\delta \\simeq 5\\times 10^{15} \\delta$ cm, where $\\delta=[\\Gamma(1-\\beta \\cos\\theta)]^{-1}$ is the relativistic Doppler factor, $\\theta$ is the viewing angle and $\\Delta t\\sim 2$ days), which seems difficult to accomplish in this scenario (e.g. Neronov \\& Aharonian 2007; Levinson 2007).  The rapid variability could be easily reproduced if the TeV emission is produced close to the base of the jet, in the region associated to the emission from blazars.  However, standard one--zone leptonic models face difficulties in reproducing the observed spectral energy distribution (SED) of the core including the TeV data.  A possible solution advocates that the emission comes from a decelerating jet (Georganopoulos et al. 2005) .  Alternatively, the emission can be produced by high--energy protons (Reimer et al. 2004).  Going further in towards the BH, TeV photons could be emitted through IC by relativistic pairs produced by electromagnetic cascades in the BH magnetosphere (Neronov \\& Aharonian 2007).  The determination of the site producing the high-energy emission is rather important. Indeed, if the emission site will be eventually identified as knot HST-1, this would have a broad impact on the current view of relativistic jets (see the discussion in Cheung et al. 2007). It is thus extremely important to investigate whether the standard view (or minimal variations of it) for the inner jet is able reproduce the observed phenomenology. If not, it would be mandatory to explore the other alternatives.  As discussed in Ghisellini, Tavecchio \\& Chiaberge (2005, hereafter Paper I), there is compelling evidence supporting the view that jets of BL Lac objects are structured at blazar scales (at distances $\\sim 10^{17}$ cm from the BH), with a fast core ({\\it spine}) surrounded by a slower sheet ({\\it layer}).  In Paper I we showed that the radiative coupling between the spine and the layer can naturally account for the deceleration inferred for the jet between the blazar and the VLBI scale.  Furthermore, one component sees the radiation of the other boosted, since the relative velocity can be relativistic. This enhances the IC emission of both the spine and the layer.  A direct consequence of such a structure is that while at small viewing angle (as in the case of blazars) the emission is dominated by the boosted spine emission, at large observing angles ($\\theta>45^\\circ$, typical for radio--galaxies) the emission from the spine would be suppressed, while the layer, characterised by a broader beaming cone, could substantially contribute to (and sometimes dominate) the overall emission. At intermediate angles both components can significantly contribute.  Following this line, in this letter we explore an alternative interpretation for the TeV emission of M87, assuming that the TeV radiation is produced in the layer of the misaligned jet. We use $H_{\\rm 0}\\rm =70\\; km\\; s^{-1}\\; Mpc^{-1} $, $\\Omega_{\\Lambda}=0.7$, $\\Omega_{\\rm M} = 0.3$. ",
        "conclusions": "\\label{sec4}  We propose that the TeV emission detected from M87 originates in the structured jet at blazar scale. The most likely possibility is that the emission from the spine reproduces the low energy component, from radio to X--rays, while the layer contributes to the bulk of the TeV radiation. With this choice, the beamed counterpart of M87 observed at small angle would have a SED closely resembling those of Low--Peaked BL Lac objects. Correspondingly the spine is characterised by physical parameters close to those usually inferred for those sources ($\\gamma _b\\sim 10^3$, $B\\sim 1$ G). The layer, instead, would be characterised by a low magnetic field ($B\\sim 0.1$ G) and large peak Lorentz factors ($\\gamma _b\\sim 10^6$). It is possible that particles in the layer are energized by turbulent acceleration (e.g. Stawarz \\& Ostrowski 2002). In this case the expected distribution would have a pile-up at the energy where acceleration and loss processes are in equilibrium (e.g. Katarzynski et al. 2006). Suggestively, in our conditions (losses dominated by IC) the expected equilibrium Lorentz factor would be close to $\\gamma \\sim 10^6$ for $B\\sim 0.1$ G.  Note that our conclusions are rather different from that of Georganopoulos et al. (2005) who discussed a similar model relying on the possibility that the jet experiences a strong deceleration at the blazar scale. In their model the high-energy emission is produced in the innermost and faster portion of the jet, while the slow external portions mainly contributes at low frequency. Under these conditions the beamed counterpart of M87 would be a TeV BL Lac. However, due to the different beaming patterns or the synchrotron and IC components (analogously to our case), the resulting spectrum would be characterised by a strong dominance of the TeV component with respect to the X--ray one, contrary to what is usually observed in the known TeV BL Lacs.  In our model the optical and the X--rays are produced by a region different from that responsible for the TeV emission. Therefore a strict correlation between low energy and TeV emission is not directly required (even if it is possible). Analogously, the MeV-GeV emission should originate in the spine. Therefore we expect that the emission possibly detected by {\\it GLAST} will not exactly follow the TeV component.  A potential weak point of our model concerns the slope of the TeV spectrum. The slope is mainly dictated by the absorption of TeV photons by the dense optical radiation field and, in this sense, it is rather robust and does not strongly depend on the underlying slope of the intrinsic TeV spectrum. In particular, hard spectra such as that recorded in 2005 are quite difficult to achieve in this scheme. A possibility to avoid important absorption of gamma-rays would be to enlarge the emission region, hence diluting the target photon density and decreasing the absorption optical depth. However the increase of the source size is limited by the observed short variability timescales. Further observations with increased sensitivity could provide stringent constraints to our interpretation."
    },
    "0801/0801.2769_arXiv.txt": {
        "abstract": "We measure the dependence of the AGN fraction on local environment at $z\\sim1$, using spectroscopic data taken from the DEEP2 Galaxy Redshift Survey, and Chandra X-ray data from the All-Wavelength Extended Groth Strip International Survey (AEGIS). To provide a clean sample of AGN we restrict our analysis to the red sequence population; this also reduces additional colour--environment correlations.  We find evidence that high redshift LINERs in DEEP2 tend to favour higher density environments relative to the red population from which they are drawn.  In contrast, Seyferts and X-ray selected AGN at $z\\sim1$ show little (or no) environmental dependencies within the same underlying population. We compare these results with a sample of local AGN drawn from the SDSS. Contrary to the high redshift behaviour, we find that both LINERs and Seyferts in the SDSS show a slowly declining red sequence AGN fraction towards high density environments. Interestingly, at $z\\sim1$ red sequence Seyferts and LINERs are approximately equally abundant.  By $z\\sim0$, however, the red Seyfert population has declined relative to the LINER population by over a factor of $7$. We speculate on possible interpretations of our results. ",
        "introduction": "\\label{sec:intro}  Both active galactic nuclei (AGN) and local environment play key roles in shaping galaxy evolution. It is now understood that AGN are those nuclei in galaxies that emit radiation powered by accretion onto a supermassive black hole. Although this realisation has proved useful for explaining many observed characteristics of these active objects, there are still many unsolved problems, especially related to the physics of the accretion process itself.  In the recent years much effort has been invested in studying the global properties of AGN as a unique population in the context of galaxy formation. In this work, we focus on a fundamental question: the dependence of the fraction of galaxies that have AGN on the density of the local environment at $z\\sim1$, and the evolution of this dependence to $z\\sim0$.  At low redshift, many authors have investigated various correlations between \\textit{galaxy properties} and environment. It is now well established that there exists a relationship between morphology and density (\\citealt{Oemler1974} and \\citealt{Dressler1980}), in that star-forming disk-dominated galaxies tend to inhabit less dense regions of the Universe than ``quiescent'' or inactive elliptical galaxies.  Moreover, additional (and related) dependencies with environment have been found, such as with stellar mass, luminosity, colour, recent and past star formation, star formation quenching, surface brightness, and concentration (to name but a few) \\citep[e.g.][]{Kauffmann2004, Balogh2004, Hogg2004, Blanton2005, Bundy2006}.  In this scenario of entangled correlations it is useful to investigate the dependence of AGN properties on the local environment, especially since AGN are believed to play an important part in shaping galaxy evolution. This has sometimes been a rather controversial issue. In the local Universe, \\cite{Miller2003} found no dependence on environment of the fraction of spectroscopically selected AGN, using the SDSS early data release. This result is in good agreement with \\cite{Sorrentino2006} who used the much larger SDSS DR4. However, many other authors claim the existence of a strong link between nuclear activity and environment, at least for specific AGN types. \\cite{Kauffmann2004} found that intermediate luminosity optically selected AGN (Seyfert IIs) favoured underdense environments, while low-luminosity optically selected AGN (Low-Ionization Nuclear Emission-line Regions; hereafter, LINERs) showed no density dependence, within the SDSS DR1. Similarly, lower-luminosity AGN were found to have a higher clustering amplitude than high-luminosity AGN by \\cite{Wake2004} and \\cite{Constantin2006a}.  Radio-loud AGN have been noted to reside preferentially in mid-to-high density regions and tend to avoid underdense environments \\citep{Zirbel1997, Best2004}.  At high redshift the study of both galaxies and AGN, and their relation to the environment, has been restricted by the lack of adequate data. Only in recent years, with the emergence of quality large-scale probes of the high redshift galaxy population, such as the DEEP2 Galaxy Redshift Survey \\citep{Davis2003} or the VIMOS-VLT Deep Survey \\cite[VVDS,][]{LeFevre2003}, have we reached the stage where we can begin to measure the statistics of galaxy evolution in some detail.  Using DEEP2, \\cite{Cooper2006} found that the many of the low redshift galaxy correlations with environment are already in place at $z\\!\\sim\\!1$. However important differences exist. The colour-density relation, for instance, tends to weaken towards higher redshifts \\citep{Cooper2007a,Cucciati2006}.  Also, bright blue galaxies are found, on average, in much denser regions than at low redshift. Such a population inverts the local star formation-density relation in overdense environments \\citep{Cooper2007b,Elbaz2007}. This inversion may be an early phase in a galaxy's transition onto the red sequence through the process of star formation quenching. The truncation of star formation in massive galaxies is believed to be tightly connected with nuclear activity \\citep[see e.g.][for more information]{Croton2006, Bower2006}. Further investigation reveals that post-starburst (aka. K+A or E+A) galaxies \\citep[e.g.][]{DresslerGunn1983} are galaxies ``caught in the act'' of quenching and are in transit to the red sequence. These predominantly ``green valley'' objects reside in similar environments to regular star forming galaxies (\\citealt{Hogg2006, Nolan2007}; Yan et al. 2007 in prep.) supporting the picture that star formation precedes AGN-triggered quenching, which precedes retirement onto the red sequence.  \\cite{Georgakakis2007} were one of the first to study the environments of X-ray selected AGN at $z\\sim1$ using a sample of 58 sources drawn from the All-Wavelength Extended Groth Strip International Survey (AEGIS, \\citealt{Davis2007}). The authors found that these galaxies avoided underdense regions with a high level of confidence. \\cite{Nandra2007} show that the same AGN reside in host galaxies that populate from the top of the blue cloud to the red sequence in colour-magnitude space.  They speculate that such AGN may be the mechanism through which a galaxy stays red.  Similar ideas have become a popular feature of many galaxy formation models that implement lower luminosity (i.e. non-quasar) AGN to suppress the supply of cooling gas to a galaxy, hence quenching star formation through a process of ``starvation'' \\citep[e.g.][]{Croton2006, Bower2006}.  In this work we study the environmental dependence of nuclear activity in red sequence galaxies within a carefully chosen sample of both X-ray and optically selected AGN, drawn from the AEGIS Chandra catalogue and the DEEP2 Galaxy Redshift Survey, respectively. Our paper is organised as follows.  In Section~\\ref{sec:survey} we describe our AGN selection. In Section~\\ref{sec:results} we present our main result: the AGN fraction of red sequence galaxies at $z\\sim1$ as a function of environment for three types of AGN (LINERs, Seyferts and X-ray selected). We undertake a comparison between our high-z results and those derived from a low-z sample drawn from the SDSS in Section~\\ref{sec:sdss}. Finally, in Sections~\\ref{sec:discussion} and \\ref{sec:summary} we provide a discussion and brief summary. Throughout, unless otherwise stated, we assume a standard $\\Lambda$CDM concordance cosmology, with $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, $w=-1$, and $h=1$. In addition, we use AB magnitudes unless otherwise stated.   \\begin{figure*} \\begin{center} \\begin{tabular}{c} \\includegraphics[scale=0.7]{./figures/FIG_paper1_2.ps} \\end{tabular} \\end{center} \\caption{Two panels that show our AGN selection of Seyferts and LINERS within the DEEP2.  The left panel plots [OII] EW versus ${\\rm H{\\beta}}$ EW for objects with accurate redshifts ($Q\\ge3$), $\\delta_{3}$ environment measures, and covered [OII], [OIII] and ${\\rm H{\\beta}}$ (grey points).  LINERs (black points) are selected using the empirical demarcation of Equation~\\ref{eqn:liners_selection} along with the colour cut defined by Equation~\\ref{eqn:color_deep2}.  The right panel shows the line ratio ${\\rm [OIII]/ H{\\beta}}$ plotted against $(U-B)$ rest-frame colour for the same DEEP2 sample (grey points).  Seyferts (black points) are selected to have ${\\rm [OIII]/H{\\beta}\\ge 3}$ and rest-frame colour $(U\\!-\\!B)>0.8$, as denoted by the horizontal and vertical lines.  See Section~\\ref{sec:optical} for further details.} \\label{fig:OSS_selection} \\end{figure*} ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Previous measures of AGN and environment}  It is difficult to make direct comparisons of our results with previously published works.  This is because past environment studies have tended to focus on the AGN fraction of \\emph{all} colours of host galaxies, and also to mix both LINER and Seyfert classes into a combined AGN population.  Our selection is restricted to the red sequence only (and also green valley), which allows us to compare high and low redshift populations and also study the AGN--environment connection without the colour--environment correlation.  Locally, a number of SDSS measures of AGN and environment have been made.  Using the SDSS early data release, \\cite{Miller2003} found no dependence on environment for the spectroscopically selected AGN fraction in a sample of 4921 objects. Specifically, the authors report no statistically significant decrease in the AGN fraction in the densest regions, although their densest points visually suggest such a trend. This result is broadly consistent with the results of both LINERs and Seyferts in Figure~\\ref{fig:SDSS_result}, even though we only consider red sequence objects.  \\cite{Kauffmann2003} also found little environment dependence of the overall fraction of detected AGN in a sample drawn from the SDSS DR1. However, they do report different behaviour when the sample is broken into strong AGN ($\\log L[{\\rm OIII}]>7$, ``Seyferts'') and weak AGN ($\\log L[{\\rm OIII}]<7$, ``LINERs''). For Seyfert they find a significant preference for low-density environments, especially when hosted by more massive galaxies. This is consistent with our SDSS findings in Figure~\\ref{fig:SDSS_result} and different to what we find at $z\\sim1$ in the DEEP2 fields. For LINERs, \\citeauthor{Kauffmann2003} measure little environment dependence, whereas we find a significant decline in the SDSS LINER fraction in our overdense bins.  The explanation for this difference may come from our removal of possible contaminating star forming galaxies by restricting our analysis to the red sequence. Also, we impose a higher line detection threshold on the SDSS data to provide a fair comparison with DEEP2 (see Section~\\ref{sec:sdss}). Finally, \\cite{Kauffmann2003} required that all lines for the BPT diagnosis were detected and this implies biasing the LINERs sample towards the strongest objects.  Between redshifts $z=0.4$ and $z=1.35$, \\cite{Cooper2007a} show that \\emph{red galaxies} within the DEEP2 survey favour overdense environments, although the blue fraction in clusters does become larger as one moves to higher redshift \\citep[see also][]{Gerke2007, Coil2007b}.  At all redshifts there exists a non-negligible red fraction in underdense environments, which evolves only weakly if at all.  \\cite{Nandra2007} show that the host DEEP2 galaxies of X-ray selected AGN within the EGS field ($\\sim 1/6$ of the DEEP2 survey volume) occupy a unique region of colour-magnitude space.  These objects typically live at the top of the blue cloud, within the green valley, or on the red sequence.  \\cite{Georgakakis2007} measure the mean environment of this population and confirm that, on average, they live in density regions above that of the mean of the survey.  They find this to be true for all host galaxy magnitudes studied ($M_B\\simlt-21$) and colours ($U-B\\simgt0.8$) (note the DEEP2 red sequence begins at $U-B\\sim1$).  However, given limited sample sizes, they were not able to establish whether the environment distribution of the X-ray AGN differed from that of the red population, rather than the DEEP2 population as a whole.   \\subsection{Understanding the sequence of events}  From our results alone a comprehensive understanding of the different environment trends within the AGN population from high to low redshift is not possible.  However, some speculation and interpretation can be made by drawing on our broader knowledge of these active objects from the literature.  One possible scenario posits that LINERs and Seyferts occur in different types of galaxies. In this picture, LINERs are often associated with young red sequence galaxies \\citep{Graves2007} and are especially common among post-starburst (K+A) galaxies \\citep{Yan2006}. These galaxies would already be into the quenched phase of their evolution but still relatively young.  Merger triggered starbursts and subsequent quasar winds are a possible mechanisms to produce rapid star formation shut down in such objects \\citep{Hopkins2006}. The gas rich merging events required in this scenario are common in overdense environments at $z\\sim1$ as clusters and massive groups assemble.  By $z\\sim0$, however, the activity in these environments has mostly ended.  Hence, if this picture is correct, one may expect an over-abundance of red sequence LINERs in dense environments at high redshift (since both star formation and rapid quenching is common) that is not seen locally.  This may be consistent with the trends found in the left panel of Figure~\\ref{fig:SDSS_result}.  Seyfert galaxies, on the other hand, could be objects in transition from the blue cloud to red sequence \\citep{Groves2006}, whose AGN are thought to be initiated by internal processes (and not mergers), inferred from their often found spiral structure (e.g. M77) (mergers act to destroy such structure).  From this, one may expect our red sequence Seyfert population to represent the tail of the colour distribution of transitioning objects whose dependence on environment is determined by secular mechanisms and who would evolve accordingly. At high redshift, disk galaxies are commonly found in all environments, including the most dense.  In contrast, overdense regions in the local Universe are dominated by passive ellipticals and show an absence of spirals.  This would be broadly consistent with our findings in the right panel of Figure~\\ref{fig:SDSS_result}, where the most significant evolution in the red Seyfert fraction arises from a depletion in overdense regions relative to other environments, from high redshift to low.  Alternatively, some authors claim that LINERs and Seyferts form a continuous sequence, with the Eddington rate the primary distinguishing factor \\citep{Kewley2006}. In this scenario, Seyferts are young objects with actively accreting black holes. As the star formation begins to decay so does the accretion rate, and the galaxy enters a transition phase. Eventually, a LINER-like object emerges, with an old stellar population and very low supermassive BH accretion rate. This picture is supported by recent studies in voids from \\cite{Constantin2007}. At high redshift, our results show that red Seyferts and LINERs are approximately equally abundant. By $z\\sim0$ however, the Seyfert population has declined relative to the LINER population by over a factor of $7$. This may be interpreted as the natural transformation of Seyferts into LINERs with time, within a galaxy population which is smoothly reddening from $z\\sim1$ to $z\\sim0$.  Moreover, the fact that high-z LINERs reside preferentially in high-density environments may imply that this Seyfert-LINER transition is more efficient in dense regions of the Universe.      In this paper we measure the dependence of the AGN fraction of red galaxies on environment in the $z\\sim1$ DEEP2 Galaxy Redshift Survey and local $z\\sim0.1$ SDSS.  We restrict our analysis to the red sequence to maintain a clean and consistent selection of AGN at high and low redshift, and this also reduces the additional effects of environment associated with galaxy colour.  Our results can be summarised as follows: \\begin{itemize} \\item[(i)] High redshift LINERs at $z\\sim1$ in DEEP2 appear to favour higher density environments relative to the red sequence from which they are drawn.  In contrast, Seyferts and X-ray selected AGN at $z\\sim1$ show much weaker (or no) environmental dependencies within the same underlying population.  Extending our analysis to include green valley objects has little effect on the results. \\item[(ii)] Low redshift LINER and Seyfert AGN in the SDSS both show a slowly declining red sequence AGN fraction towards high density environments.  This is in contrast to the high redshift result. \\item[(iii)] At $z\\sim1$, Seyferts and LINERs are approximately equally abundant.  By $z\\sim0$ however, the Seyfert population has declined relative to the LINER population by over a factor of $7$. \\end{itemize}  It is important to remember that such measures are difficult to make with current data, and hence we remain limited by statistics to the extent to which we can physically interpret our results.  Regardless, a robust outcome of our analysis is the differences between LINER and Seyfert AGN populations in high density regions, and between high and low redshift in all environments.  Our results indicate that a greater understanding of both AGN and galaxy evolution may be possible if future analyses simultaneously focus on the detailed subdivision of different AGN classes, host galaxy properties, and their environment."
    },
    "0801/0801.2902_arXiv.txt": {
        "abstract": "The coronal magnetic field is an important quantity because the magnetic field dominates the structure of the solar corona. Unfortunately direct measurements of coronal magnetic fields are usually not available. The photospheric magnetic field is measured routinely with vector magnetographs. These photospheric measurements are extrapolated into the solar corona. The extrapolated coronal magnetic field depends on assumptions regarding the coronal plasma, e.g. force-freeness. Force-free means that all non-magnetic forces like pressure gradients and gravity are neglected. This approach is well justified in the solar corona due to the low plasma beta. One has to take care, however, about ambiguities, noise and non-magnetic forces in the photosphere, where the magnetic field vector is measured. Here we review different numerical methods for a nonlinear force-free coronal magnetic field extrapolation: Grad-Rubin codes, upward integration method, MHD-relaxation, optimization and the boundary element approach. We briefly discuss the main features of the different methods and concentrate mainly on recently developed new codes. ",
        "introduction": "$ $\\\\ Information regarding the coronal magnetic field are important for space weather application like the onset of flares and coronal mass ejections (CMEs). Unfortunately we usually cannot measure the coronal magnetic field directly, although recently some progress has been made \\citep[see e.g.,][]{judge98,solanki:etal03,lin:etal04}. Due to the optically thin coronal plasma direct measurements of the coronal magnetic field have a line-of-sight integrated character and to derive the accurate 3D structure of the coronal magnetic field a vector tomographic inversion is required. Corresponding feasibility studies based on coronal Zeeman and Hanle effect measurements have recently been done by \\cite{kramar:etal06} and \\cite{kramar:etal07}. These direct measurements are only available for a few individual cases and usually one has to extrapolate the coronal magnetic field from photospheric magnetic measurements. To do so, one has to make assumptions regarding the coronal plasma. It is helpful that the low solar corona is strongly dominated by the coronal magnetic field and the magnetic pressure is orders of magnitudes higher than the plasma pressure. The quotient of plasma pressure $p$ and magnetic pressure, $B^2/(2 \\mu_0)$    is small compared to unity $(\\beta=2 \\mu_0 p/B^2 \\ll 1)$. In lowest order non-magnetic forces like pressure gradient and gravity can be neglected which leads to the force-free assumption. Force free fields are characterized by the equations:  \\begin{eqnarray} {\\bf j}\\times{\\bf B}  & = & {\\bf 0}        \\label{forcebal}\\\\ \\nabla \\times {\\bf B }& = & \\mu_0 {\\bf j}  \\label{ampere}, \\\\ \\nabla\\cdot{\\bf B}    & = &         0      \\label{solenoidal}, \\end{eqnarray} where ${\\bf B}$ is the magnetic field, ${\\bf j}$ the electric current density and $\\mu_0$ the permeability of vacuum. Equation (\\ref{forcebal}) implies that for force-free fields the current density and the magnetic field are parallel, i.e.  \\begin{equation} \\mu_0 {\\bf j} = \\alpha {\\bf B}, \\label{jparb} \\end{equation}  or by replacing ${\\bf j}$ with Eq. (\\ref{ampere})  \\begin{equation} \\nabla \\times {\\bf B }  =  \\alpha {\\bf  B} \\label{amperealpha}, \\end{equation}  where $\\alpha$ is called the force-free function. To get some insights in the structure of the space dependent function $\\alpha$, we take the divergence of Eq. (\\ref{jparb}) and make use of Eqs. (\\ref{ampere}) and (\\ref{solenoidal}):  \\begin{equation} {\\bf B} \\cdot \\nabla \\alpha  = 0, \\label{alphaeq} \\end{equation}  which tells us that the force-free function $\\alpha$ is constant on every field line, but will usually change from one field line to another. This generic case is called nonlinear force-free approach.  Popular simplifications are $\\alpha=0$ (current free potential fields, see e.g., \\cite{schmidt64,semel67,schatten69,sakurai82}) and $\\alpha={\\rm constant}$ (linear force-free approach, see e.g., \\cite{nakagawa:etal72,chiu:etal77,seehafer78,alissandrakis81,seehafer82,semel88}). These simplified models have been in particular popular due to their relative mathematical simplicity and because only line-of-sight photospheric magnetic field measurements are required. Linear force-free fields still contain one free global parameter $\\alpha$, which can be derived by comparing coronal images with projections of magnetic field lines (e.g., \\cite{carcedo:etal03}). It is also possible to derive an averaged value of $\\alpha$ from transverse photospheric magnetic field measurements (e.g. \\cite{pevtsov:etal94,wheatland99,leka:etal99}). Despite the popularity and frequent use of these simplified models in the past, there are several limitations in these models (see below) which  ask for considering the more sophisticated nonlinear force-free approach.  Our aim is to review recent developments of the extrapolation of nonlinear force-free fields (NLFFF). For earlier reviews on force-free fields we refer to \\citep{sakurai89,aly89,amari:etal97,mcclymont:etal97} and chapter 5 of \\cite{aschwanden:book}. Here we will concentrate mainly on new developments which took place after these earlier reviews. Our main emphasis is to study methods which extrapolate the coronal magnetic field from photospheric vector magnetograms. Several vector magnetographs are currently operating or planed for the nearest future, e.g., ground based: the solar flare telescope/NAOJ \\citep{sakurai:etal95}, the imaging vector magnetograph/MEES Observatory \\citep{mickey:etal96}, Big Bear Solar Observatory, Infrared Polarimeter VTT, SOLIS/NSO  \\citep{henney06} and space born: Hinode/SOT \\citep{shimizu04}, SDO/HMI \\citep{borrero:etal06}. Measurements from these vector magnetograms will provide us eventually with the magnetic field vector on the photosphere, say $B_{z0}$ for the normal and $B_{x0}$ and $B_{y0}$ for the transverse field. Deriving these quantities from the measurements is an involved physical process, which includes measurements based on the Zeeman and Hanle effect, the inversion of Stokes profiles (e.g., \\cite{labonte:etal99}) and removing the $180^\\circ$ ambiguity (e.g., \\cite{metcalf94,metcalf:etal06}) of the horizontal magnetic field component. Special care has to be taken for vector magnetograph measurements which are not close to the solar disk, when the line-of-sight and normal magnetic field component are far apart (e.g. \\cite{gary:etal90}). For the purpose of this paper we do not address the observational methods and recent developments and problems related to deriving the photospheric magnetic field vector. We rather will concentrate on how to use the photospheric $B_{x0}, \\, B_{y0}$ and $B_{z0}$ to derive the coronal magnetic field.  The transverse photospheric magnetic field $(B_{x0}, \\, B_{y0})$ can be used to approximate the normal electric current distribution by  \\begin{equation} \\mu_0 j_{z0}=\\frac{\\partial B_{y0}}{\\partial x}-\\frac{\\partial B_{x0}}{\\partial y} \\label{j_photo} \\end{equation}  and from this one gets the distribution of $\\alpha$ on the photosphere by  \\begin{equation} \\alpha(x,y)=\\mu_0 \\frac{j_{z0}}{B_{z0}} \\label{alpha0direct} \\end{equation}  By using Eq. (\\ref{alpha0direct}) one has to keep in mind that rather large uncertainties in the transverse field component and  numerical derivations used in (\\ref{j_photo}) can cumulate  in significant errors for the current density. The problem becomes even more severe by using (\\ref{alpha0direct}) to compute $\\alpha$ in regions with a low normal magnetic field strength $B_{z0}$. Special care has to be taken at photospheric polarity inversion lines, i.e. lines along which $B_{z0}=0$ \\citep[see e.g.,][]{cuperman:etal91}. The nonlinear force-free coronal magnetic field extrapolation is a boundary value problem. As we will see later some of the NLFFF-codes make use of (\\ref{alpha0direct}) to specify the boundary conditions while other methods use the photospheric magnetic field vector more directly to extrapolate the field into the corona.   Pure mathematical investigations of the nonlinear force-free equations \\citep[see e.g.][]{aly84,boulmezaoud:etal00,aly05} and modelling approaches not based on vector magnetograms are important and occasionally mentioned in this paper. A detailed review of these topics is well outside the scope of this paper, however. Some of the model-approaches not based on vector magnetograms are occasionally used to test the nonlinear force-free extrapolation codes described here. \\subsection{Why do we need nonlinear force-free fields?} \\begin{itemize} \\item A comparison of global potential magnetic field models with TRACE-images by \\cite{schrijver:etal05a} revealed that significant nonpotentially occurs regulary in active regions, in particular when new flux has emerged in or close to the regions. \\item Usually $\\alpha$ changes in space, even inside one active region. This can be seen, if we try to fit for the optimal linear force-free parameter $\\alpha$ by comparing field lines with coronal plasma structures. An example can be seen in \\cite{wiegelmann:etal02} where stereoscopic reconstructed loops by \\cite{aschwanden:etal99} have been compared with a linear force-free field model. The optimal value of $\\alpha$ changes even sign within the investigated active regions, which is a contradiction to the $\\alpha={\\rm constant}$ linear force-free approach (see Fig. \\ref{figure1}). \\item Photospheric $\\alpha$ distributions derived from vector magnetic field measurements by Eq. (\\ref{alpha0direct}) show as well that $\\alpha$ usually changes within an active region \\citep[see, e.g.][]{regnier:etal02}. \\item Potential and linear force-free fields are too simple to estimate the free magnetic energy and magnetic topology accurately. The magnetic energy of linear force-free fields is unbounded in a halfspace \\citep{seehafer78} which makes this approach unsuitable for energy approximations of the coronal magnetic field. Potential fields have a minimum energy for an observed line-of-sight photospheric magnetic field. An estimate of the excess of energy a configuration has above that of a potential field is an important quantity which might help to understand the onset of flares and coronal mass ejections. \\item A direct comparison of measured fields in a newly developed active region by \\cite{solanki:etal03} with extrapolations from the photosphere with a potential, linear and nonlinear force-free model by \\cite{wiegelmann:etal05} showed that nonlinear fields are more accurate than simpler models. Fig. \\ref{figure2} shows some selected magnetic field lines for the original measured field and extrapolations from the photosphere with the help of a potential, linear and nonlinear force-free model. \\end{itemize} These points tell us that nonlinear force-free modelling is required for an accurate reconstruction of the coronal magnetic field. Simpler models have been used frequently in the past. Global potential fields provide some information of the coronal magnetic field structure already, e.g. the location of coronal holes. The generic case of force-free coronal magnetic field models are nonlinear force-free fields, however. Under generic we understand that $\\alpha$ can (and usually will) change in space, but this approach also includes the special cases $\\alpha={\\rm constant}$ and $ \\alpha=0$. Some active regions just happen to be more potential (or linear force-free) and if this is the case they can be described with simpler models. Linear force-free models might provide a rough estimate of the true 3D magnetic field structure if the nonlinearity is weak. The use of simpler models was often justified due to limited observational data, in particular if only the line-of-sight photospheric magnetic field has been measured.  While the assumption of nonlinear force-free fields is well accepted for the coronal magnetic fields in active regions, this is not true for the photosphere. The photospheric plasma is a finite $\\beta$ plasma and nonmagnetic forces like pressure gradient and gravity cannot be neglected here. As a result electric currents have a component perpendicular to the magnetic field, which contradicts the force-free assumption. We will discuss later, how these difficulties can be overcome. ",
        "conclusions": "Within the last few years the scientific community showed a growing interest into coronal magnetic fields. \\footnote{Publications containing the phrase 'coronal magnetic fields' in title or abstract have been cited less than about $50$ times per year until the early 1990th and this number increased to about $150$ citations per year in 2004. A peak year was 2006 (last year) with more than $300$ citations. (Source: ISI Web of Knowledge, march 2007)} The development of new ground based and space born vector magnetographs provide us measurements of the magnetic field vector on the suns photosphere. Accompanied from these hardware development, software has been developed to extrapolate the photospheric measurements into the corona. Special attention has recently been given to nonlinear force-free codes. Five different numerical approaches (Grad-Rubin, upward integration, MHD-relaxation, optimization, boundary elements) have been developed for this aim. It is remarkable that new codes or major updates of existing codes have been published for all five methods within the last two years, mainly in the last year (2006). A workshop series (NLFFF-consortium) since 2004 on nonlinear force-free fields has recently released synergy effects, by bringing modelers of the different numerical implementations together to compare, evaluate and improve the programs. Several of the most recent new codes and utility programs (e.g. preprocessing) have at least been partly inspired by these workshops. The new implementations have been tested with the smooth semi-analytic Low-Lou-equilibrium  and showed reasonable agreement with this reference field. While all methods aim for a reconstruction of the coronal magnetic field from the photospheric magnetic field vector, the way how these measurements are used to prescribe the boundaries of the codes is different. \\begin{itemize} \\item MHD-relaxation and optimization use  $B_{x0}, B_{y0}, B_{z0}$ on the bottom boundary. This over-determines the boundary value problem. Both methods are closely related and compute the magnetic field in a computational box with  \\begin{equation} \\frac{\\partial {\\bf B}}{\\partial t}= \\mu {\\bf F}, \\end{equation}  where the structure of {\\bf F} is somewhat different (the optimization approach has more terms) for both methods. Usually a potential field is used as initial state for both approaches, also the use of a linear force-free initial state is possible. Recently a multiscale version of optimization has been installed, which uses a low resolution NLFFF-field as input for higher resolution computations. Specifying the entire magnetic field vector on the bottom boundary is an over-imposed problem and a unique NLFFF-field (or a solution at all) requires that the boundary data fulfill certain consistency criteria. A recently developed preprocessing-routine helps to find suitable consistent boundary data from inconsistent photospheric measurements. Earlier and current comparisons showed a somewhat higher accuracy for the optimization approach. A practical advantage of the MHD-approach is that in principle any available time-dependent MHD-code can be adjusted to compute the NLFFF-field. \\item The Grad-Rubin approach uses  $B_{z0}$ and the distribution of $\\alpha$ computed with Eq. (\\ref{alpha0direct}) for one polarity, which corresponds to well posed mathematical problem. A practical problem is that the computation of $\\alpha$ requires numerical differences of the noisy and forced transverse photospheric field  $B_{x0}, B_{y0}$ with (\\ref{j_photo}) leading to inaccuracies in the normal electric current distribution and in $\\alpha$. For smooth semi-analytic test cases this is certainly not a problem, but real data require special attention (smoothing, preprocessing, limiting $\\alpha \\not=0$ to regions where $B_{z0}$ is above a certain limit) to derive a meaningful distribution of $\\alpha$. While the method requires only $\\alpha$ for one polarity, the computation from photospheric data provide $\\alpha$ for both polarities. We are not aware of any tests on how well NLFFF-solutions computed from $\\alpha$ prescribed on the positive and negative polarity coincide. It is also unclear how well the computed transverse field components on the bottom boundary agree with the measured values \\footnote{In principle $B_{x0}, B_{y0}$ may have an additional field $(B_{x0}, B_{y0})+(\\partial_x \\partial_y) \\phi$ without making a difference for $\\alpha$ and hence for the Grad-Rubin result.} of $B_{x0}, B_{y0}$. More tests on this topics are necessary, including the recently installed possibility to prescribe $\\alpha$ for both polarities and adjust the boundary by a weighed average of $\\alpha$ on both polarities to fulfill Eq. (\\ref{alphaeq}). As initial state the Grad-Rubin method uses a potential field, which is also true for MHD-relaxation and optimization. \\item The upward integration and the boundary element method prescribe both all components of the bottom boundary magnetic field vector and the $\\alpha$ distribution computed with Eq. (\\ref{alpha0direct}). This approach over imposes the boundary and $B_{x0}, B_{y0}, B_{z0}$ and $\\alpha$ have to be consistent which each other and the force-free assumption. This is certainly not a problem at all for smooth semi-analytic test equilibria and strategies to derive consistent boundary data from measured data have been developed recently. Different from the three approaches discussed above, upward integration and boundary element methods do not require to compute first an initial potential field in the computational domain. It is well known that the upward integration method is based on an ill-posed problem and the method has not been considered for several years, but a recent implementation with smooth analytic functions might help to regularize this method. First tests showed a reasonable results for computations with the smooth semi-analytic Low-Lou solution.  The boundary element method has the problem to be very slow and an earlier implementation of this method could not reach a converged state for a $64^3$ boxed used in the I. NLFFF-consortium paper due to this problem. A new 'direct boundary method' has been developed, which seems to be faster than the original 'boundary element method', but still slower compared with the four other NLFFF-approaches if the task is to compute a 3D magnetic field in an entire 3D-domain. Different from all other described methods the boundary element approach allows to compute the nonlinear force-free field vector at any arbitrary point above the boundary and it is not necessary to compute the entire 3D-field above the photosphere. This might be a very useful feature if one is interested in computing the magnetic field only along a single loop and not interested in an entire active region.  The new implementations of upward integration and boundary element method show both reasonable results for first tests with the smooth semi-analytic Low and Lou equilibrium. Further tests with more sophisticated equilibria, e.g. a solar-like test case as used in the II. NLFFF-consortium paper would be useful to come to more sound conclusions regarding the feasibility of these methods. \\end{itemize} Most of the efforts done in nonlinear force-free modelling until now concentrated mainly on developing these models and testing their accuracy and speed with the help of well known test configuration. Not too many applications of nonlinear force-free models to real data are currently available, from which we learned new physics. One reason was the insufficient access to high accuracy photospheric vector magnetograms and a second one were limitations of the models. Force-free field extrapolation is a mere tool, if properly employed on vector magnetograms, it can help to understand physical, magnetic field dominated processes in the corona. Both the computational methods as well as the accuracy of required measurements (e.g. with Hinode, SDO) are rapidly improving. Within the NLFFF-consortium we just started (since april 2007) to apply the different codes to compute nonlinear force-free coronal magnetic fields from Hinode vectormagnetograms. This project might provide us already some new insights about coronal physics.  To conclude, we can say that the capability of Cartesian nonlinear force-free extrapolation codes has rapidly increased in recent years. Only three years ago most codes run usually on grids of about $64^3$ pixel. Recently developed or updated codes (Grad-Rubin by Wheatland, MHD-relaxation by Valori, optimization by Wiegelmann, optimization by McTiernan) have been applied to grids of about $300^3$ pixel. Although this increase of traceable grid sizes is certainly encouraging, the resolution of current and near future vector magnetographs (which of course measure only data in 2D!) is significantly higher. We should keep in mind, however, that the currently implemented NLFFF-codes have been only moderately parallelized using only a few processors. The CSWM-conference, where this paper has been presented, took place at the 'Earth simulator' in Yokohama, which contains several thousands of processors used for Earth-science computer simulations. An installation of NLFFF-codes on such massive parallel computers (which has been briefly addressed on NLFFF-consortium meetings) combined with adaptive mesh refinements might enable drastically improved grid sizes. One should not underestimate the time and effort necessary to program and install such massive parallelized versions of existing codes. As full disk vectormagnetograms will become available soon (SOLIS, SDO/HMI) it is also an important task to take a spherical geometry into account. First steps in this direction have been carried out with the optimization and boundary element methods. Spherical NLFFF-geometries are currently still in it's infancy and have been tested until now only with smooth semi-analytic Low and Lou equilibria and require further developments.  Attention has also recently been drawn to the problem that the coronal magnetic field is force-free, but the photospheric one is not. Tests with extrapolations from solar-like artificial photospheric and chromospheric measurements within the II. NLFFF-consortium paper revealed that extrapolations from the (force-free) chromospheric field provide significantly better results as extrapolations using directly the (forced) photospheric field. Applying a pre-processing program on the photospheric data, which effectively removes the non-magnetic forces, leads to significantly better results, but they are not as good as by using the chromospheric magnetic field vector as boundary condition. An area of current research is the possibility to use chromospheric images to improve the preprocessing of photospheric magnetic field measurements. Improvements in measuring the chromospheric magnetic field directly \\citep[e.g.][]{lagg:etal04} might further improve to find suitable boundary conditions for NLFFF-extrapolations. Force-free extrapolations are not suitable, however, to understand the details of physical processes on how the magnetic field evolves from the forced photosphere into the chromosphere, because non-magnetic forces are important in the photosphere. For a better understanding of these phenomena more sophisticated models which take pressure gradients and gravity (and maybe also plasma flow) into account are required. Some first steps have been done with a generalization of the optimization method by \\cite{wiegelmann:etal06b}, but such approaches are still in their infancy and have been tested so far only with smooth MHD-equilibria. It is also not entirely clear how well necessary information regarding the plasma (density, pressure, temperature, flow) can be derived from measurements. Non-magnetic forces become important also in quiet sun regions (\\cite{schrijver:etal05}) and in the higher layers of the corona, where the plasma $\\beta$ is of the order of unity. Coronagraph measurements, preferably from two viewpoints as provided by the STEREO-mission, combined with a tomographic inversion might help here to get insights in the required 3D structure of the plasma density. One should also pay attention to the combination of extrapolation methods, as described here, with measurements of the Hanle and Zeeman effects in coronal lines which allows the reconstruction of the coronal magnetic field as proposed in feasibility studies of vector tomography by \\cite{kramar:etal06,kramar:etal07}. Other measurements of coronal features, e.g., coronal plasma images from two STEREO-viewpoints, can be used for observational tests of coronal magnetic field models. Using two  viewpoints provide a much more restrictive test of models as images from only one view direction. While a nonlinear force-free coronal magnetic field model helps us to derive the topology, magnetic field and electric current strength in coronal loops, they do not provide plasma parameters. One way to get insights regarding the coronal plasma is the use of scaling laws to model the plasma along the reconstructed 3D field lines and compare correspondent artificial plasma images with real coronal images. \\cite{schrijver:etal04} applied such an approach to global potential coronal magnetic fields and compared simulated and real coronal images from one viewpoint. A generalization of such methods towards the use of more sophisticated magnetic field models and coronal images from two STEREO-viewpoints will probably provide many insights regarding the structure and physics of the coronal plasma. An important challenge is for example the coronal heating problem. The dominating coronal magnetic field is assumed to play an important role here, because magnetic field configuration containing free energy can under certain circumstances reconnect (\\cite{priest96,priest99}) and supply energy for coronal heating. \\cite{priest:etal05} pointed out that magnetic reconnection at separators and separatrices plays an important role for coronal heating. Nonlinear force-free models can help here to identify the magnetic field topology, magnetic null points, separatrices and localized strong current concentration. While magnetic reconnection \\citep[see e.g.][]{priest:etal99} is a dynamical phenomenon, the static magnetic field models discussed here can help to identify the locations favourable for reconnection. Time sequences of nonlinear force-free models computed from corresponding vector magnetograms will also tell wether the topology of the coronal magnetic field has changed due to reconnection, even if the physics of reconnection is not described by force-free models. Sophisticated 3D coronal magnetic field models and plasma images from two viewpoints might help to constrain the coronal heating function further, which has been done so far with plasma images from one viewpoint \\citep[][by using data from Yokoh, Soho and Trace]{aschwanden01,aschwanden01a}."
    },
    "0801/0801.0907_arXiv.txt": {
        "abstract": "We study dust accumulation by photophoresis in optically thin gas disks. Using formulae of the photophoretic force that are applicable for the free molecular regime and for the slip-flow regime, we calculate dust accumulation distances as a function of the particle size. It is found that photophoresis pushes particles (smaller than 10 cm) outward. For a Sun-like star, these particles are transported to $0.1-100$ AU, depending on the particle size, and forms an inner disk. Radiation pressure pushes out small particles ($\\la 1$ mm) further and forms an extended outer disk. Consequently, an inner hole opens inside $\\sim 0.1$ AU. The radius of the inner hole is determined by the condition that the mean free path of the gas molecules equals the maximum size of the particles that photophoresis effectively works on ($100 \\ \\micron - 10$ cm, depending on the dust property). The dust disk structure formed by photophoresis can be distinguished from the structure of gas-free dust disk models, because the particle sizes of the outer disks are larger, and the inner hole radius depends on the gas density. ",
        "introduction": "Protoplanetary disks are composed of gas and dust. At initial stages of their evolution, the gas of protoplanetary disks is as massive as $10^{-3} - 10^{-1} M_{\\sun}$ (Greaves 2004), and small dust particles are mixed with the gas, making the disks optically thick at optical wavelengths. As the dust particles grow, the number of small particles reduces, and at a certain stage, the disks become optically thin (Tanaka et al. 2005; Dullemond \\& Dominik 2005). The amount of gas also decreases, as the gas dissipates (e.g., Hartmann et al. 1998; Clarke et al. 2001; Takeuchi et al. 2005; Alexander et al. 2006a,b). At late stages of the disk evolution, the disks become gas-free dust disks, which are observed as Vega-type disks. During the transition from protoplanetary disks to Vega-type disks, there may be a stage where the disks are optically thin, but their gas component still remains. An example is HD 141569A, which is a 5 Myr Herbig Be star (Weinberger et al. 2000) and has an optically thin dust disk (Augereau et al. 1999; Weinberger et al. 1999; Fisher et al. 2000; Mouillet et al. 2001; Marsh et al. 2002; Clampin et al. 2003). The gas component of this system has been observed (Zuckerman 1995; Dent et al. 2005; Goto et al. 2006) and its amount is estimated to be $\\la 60 \\ M_{\\earth}$ (Ardila et al. 2005).  Dynamics of dust particles in optically thin gas disks is of interest in order to investigate the structure of transitional disks. Krauss \\& Wurm (2005) considered the motion of dust particles in a gas disk that is optically thin at optical wavelengths. A dust particle receives stellar radiation directly, and the radiation pressure pushes the particle outward. In addition to radiation pressure, interaction between the particle and the surrounding gas molecules induces photophoresis, which also pushes the particle outward (see also Wurm \\& Krauss 2006; Krauss et al. 2007). When these outward accelerations act with the gas drag on the particle, the particle drifts outward (Takeuchi \\& Artymowicz 2001). In the outer part of the disk ($\\ga 10$ AU), where the mean free path of the gas molecules is larger than the particle size (of $\\la 1$ m), the photophoretic force is proportional to the gas density. When a particle moves outward to the point where the gas density is no longer high enough to induce a strong photophoretic force, the particle's outward motion stops. Consequently, the dust particles pile-up at a certain distance from the star that is determined by the density profile of the gas disk. In a gas disk with mass $\\sim 0.01 \\ M_{\\sun}$, particles of $100 \\ \\micron$ to 10 cm pile-up at several tens of AU. This spontaneous ring formation is a characteristic feature that is caused by photophoresis in a gaseous dust disk.  Krauss et al. (2007) have demonstrated that photophoretic dust motion may be the key process for the transition from optically thick protoplanetary disks to optically thin circumstellar disks with ring-like dust distributions via the stage of transitional disks with an inner hole and a more or less sharp transition to the outer disk that is continuously pushed outward. This outward dust migration can also explain the presence of material formed close to the sun like chondrules and CAIs in asteroids of the main asteroid belt (Wurm \\& Krauss 2006) or high temperature crystalline silicates in comets from the Kuiper belt (Petit et al. 2006; Brownlee et al. 2006; Mousis et al. 2007).  In this paper, we seek to investigate other characteristic structures formed by photophoresis, especially in the inner part of the disk, but at a stage when most of the dust has already been built into larger bodies. Krauss \\& Wurm (2005) focus on the dust dynamics in the outer part of the disk ($\\ga 10$ \\ AU) and use a formula for the photophoretic force that is relevant for low gas densities. In order to investigate the structure of the whole disk, we use a photophoresis formula that can be applied to the whole range of gas densities. Details of the formula for the photophoretic force adopted here are described in \\S\\ref{sec:drift}. Using this formula, we find that photophoresis does not induce dust pile-up at the innermost region of the disk ($\\la 0.1$ AU), leading to the formation of an inner hole (see \\S\\ref{sec:accrad}). At the inner part of the gas disk, the gas density may be too high that the gas disk itself is no longer optically thin to the stellar radiation, and thus any radiation effects on the dust such as radiation pressure and photophoresis do not work at all. In order to investigate whether photophoresis still works at the inner disk, in \\S\\ref{sec:opacity}, we calculate the optical depth of the disk due to Rayleigh scattering of the hydrogen molecules. In \\S\\ref{sec:discussion1}, the thermal relaxation time of dust particles is estimated and is compared to the rotation periods induced by gas turbulence and by photophoresis. In \\S\\ref{sec:discussion2}, we discuss what are the characteristic features that photophoresis forms in dust disks, and what are the observable differences from the gas-free dust disk structures. In Appendix A, a simple calculation deriving the photophoretic force is described. ",
        "conclusions": "\\subsection{Thermal Relaxation and Rotation Times of Dust Particles} \\label{sec:discussion1}  As mentioned in \\S\\ref{sec:fdust}, if particles rotate rapidly, photophoresis does not work. Even if particle rotation is not considerably rapid compared to the thermal relaxation time of the particle, photophoretic Yarkovsky effect may prevent the particle accumulation. Such cases where photophoresis does not work are represented by the model of $J_1=0$, where the incident starlight does not cause any temperature gradient in the particles. If photophoresis is suppressed by any reason, particles can reside in the region where ${\\rm Kn} < 1$ (Fig. \\ref{fig:req}), and such cases can be observationally distinguished from the disks in which photophoresis clears the dust in the innermost region. Hence, whether particles in gas disks rotate rapidly or not can be observationally investigated by checking whether the dust in the region of ${\\rm Kn} < 1$ is cleared or not.  In this subsection, we estimate the thermal relaxation time of a dust particle and then compare it to the rotation period induced by gas turbulence or by photophoresis itself. We use a simple model of a cylindrical dust particle described in Appendix A. Consider a cylindrical dust particle with radius $a$ and height $2a$ (see Figure \\ref{fig:cylndust}). The stellar radiation flux $I$ irradiates the front surface. We suppose that, at the beginning, the temperature inside the particle is homogeneous. This initial (or average) temperature $T_d$ is given by equation (\\ref{eq:td}), assuming that the incident flux $I$ balances with thermal radiation from the front and back surfaces. (Radiation from the side surface is ignored for simplicity.) The temperature of the front surface increases to the equilibrium value $T_f$ in a thermal relaxation time $\\tau_{\\rm th}$. During $\\tau_{\\rm th}$, the incident energy at the front surface conducts for a length $\\delta a$ and forms a skin layer of temperature gradient. During $\\tau_{\\rm th}$, the temperature at the front surface is approximated as the initial value $T_d$, and the energy flux inside the skin layer is estimated as $Q_{\\rm con} \\approx I-\\sigma_{\\rm SB} T_d^4 = I/2$. (We set $\\varepsilon=1$.) When the skin layer depth has grown to $\\delta a$ and the temperature of the front surface converges to $T_f$, the energy flux is $Q_{\\rm con} \\approx k_d (T_f - T_d) / \\delta a$. Equating the above two expressions for $Q_{\\rm con}$, the thickness of the temperature gradient layer is \\begin{equation} \\delta a = \\frac{2 k_d T_d \\Delta}{I} \\ , \\end{equation} where $\\Delta = (T_f - T_d)/T_d$. Note that the maximum value of $\\delta a$ is $2a$. The thermal relaxation time is \\begin{equation} \\tau_{\\rm th} = \\frac{c_d \\rho_d \\delta a^2}{k_d} \\ , \\end{equation} where $c_d$ is the specific heat capacity of the particle.  \\begin{figure} \\epsscale{1.2} \\plotone{f7.eps} \\caption{ Thermal relaxation times of cylindrical particles of various sizes are plotted by the solid lines against the distance from the star. The dashed line shows the turnover time, $\\tau_{\\rm ed}$, of the smallest eddies for $\\alpha_{\\rm tur}=10^{-2}$. \\label{fig:tauth} } \\end{figure}  Figure \\ref{fig:tauth} shows the thermal relaxation time $\\tau_{\\rm th}$ for $10 \\ \\micron - 10$ cm particles in the disk of model A. We set the particle bulk density $\\rho_d=1 \\ {\\rm g \\ cm^{-3}}$, the specific heat capacity $c_d=10^7 \\ {\\rm erg \\ g^{-1} \\ K^{-1}}$, and the thermal conductivity $k_d=10^2 \\ {\\rm erg \\ s^{-1} \\ cm^{-1} \\ K^{-1}}$. The temperature difference $\\Delta$ is calculated by equation (\\ref{eq:tempdif}). At small distances from the star, radiative cooling $4 \\varepsilon \\sigma_{\\rm SB} T_d^4$ dominates in determining $\\Delta$ in equation (\\ref{eq:tempdif}), and $\\Delta \\approx 1/4$. Thus, $\\tau_{\\rm th}$ is independent of the particle size $a$, and is proportional to $T_d^{-6}$. In our model $T_d \\propto r^{-1/2}$, and thus $\\tau_{\\rm th} \\propto r^3$. At large distances, on the other hand, internal thermal conduction $k_d T_d /a$ determines $\\Delta$. In this case, the skin depth is $\\delta a=a$. Hence, $\\tau_{\\rm th}$ is proportional to $a^2$ and independent of $r$.  If the rotation period of the particle is smaller than $\\tau_{\\rm th}$, photophoresis is considerably suppressed. Particle rotation can be excited by several mechanisms such as Brownian motion, gas turbulent motion, collisions with other dust particles, and the photophoretic force itself can induce rotation if it has a offset from the mass center. Detailed calculation of the rotation period for each mechanism is beyond the scope of this paper. Here, we make a rough estimate of the rotation speed induced by turbulence and by the photophoretic force.  Suppose that the gas disk has isotropic turbulence with the Kolmogorov energy spectrum. We assume that the largest eddies have size $L=\\alpha_{\\rm tur}^{1/2} h_g$ and velocity $V=\\alpha_{\\rm tur}^{1/2} c_s$, where $\\alpha_{\\rm tur}$ is the ``$\\alpha$-viscosity'' parameter (Cuzzi et al. 2001). The energy of turbulent motion cascades down to smaller eddies, and finally dissipates by the molecular viscosity. From dimensional analysis, the energy dissipation rate per unit mass is $\\dot{e} \\sim V^3 / L \\sim \\alpha_{\\rm tur} c_s^2 \\Omega_{\\rm K}$. The eddy turnover time is faster for smaller eddies and that for the smallest eddies is \\begin{equation} \\tau_{\\rm ed} \\sim \\left( \\frac{\\eta_{\\rm vis}}{\\rho_g \\dot{e}} \\right)^{1/2} \\sim \\alpha_{\\rm tur}^{-1/2} \\left( \\frac{l}{h_g} \\right)^{1/2} \\Omega_{\\rm K}^{-1} \\ , \\end{equation} where $\\eta_{\\rm vis}=l v_T \\rho_g /2$ is the molecular viscosity (Weidenschilling 1984). The smallest eddies can induce a particle rotation period as short as $\\tau_{\\rm ed}$, if the particle is well coupled to the turbulent motion of the smallest eddies. However, if the particle does not strongly couple to the gas, the rotation period is longer than $\\tau_{\\rm ed}$. In Figure \\ref{fig:tauth}, the turnover time, $\\tau_{\\rm ed}$, of the smallest eddies is plotted by the dashed line for $\\alpha_{\\rm tur} = 10^{-2}$. The particle rotation period induced by gas turbulence is expected to be larger than the dashed line, and therefore is much longer than the thermal relaxation time.  \\begin{figure} \\epsscale{1.2} \\plotone{f8a.eps} \\plotone{f8b.eps} \\caption{ Times taken for a particle to rotate 180 degree by photophoresis are plotted by the dashed lines. The photophoretic force is exerted on a point at a distance $b=a$ or $b=0.01a$ from the cylinder axis. The solid line shows the thermal relaxation time. ($a$) For $100 \\ \\micron$ particles. ($b$) For 1 cm particles. \\label{fig:rot} } \\end{figure}  We next consider rotation induced by the torque exerted by photophoresis. Consider again a cylindrical dust particle and suppose that the photophoretic force is exerted on a point at a distance $b$ from the cylinder axis. The torque is $K=b F_{\\rm ph}$. The value of $b$ is unknown, and we treat $b$ as a free parameter. The principal moment of inertia of the cylinder (for the axis perpendicular to the cylinder axis) is $I_{xx}=7 \\pi a^5 \\rho_d / 6$. The time needed for a initially stationary particle to rotate 180 degree is \\begin{equation} \\tau_{\\rm rot} = \\left( \\frac{2 \\pi I_{xx}}{K} \\right)^{1/2} = \\left( \\frac{7 \\pi^2 \\rho_d a^5}{3 b F_{\\rm ph}} \\right)^{1/2} \\ . \\end{equation} In Figure \\ref{fig:rot}, the rotation time $\\tau_{\\rm rot}$ is plotted for assumed off-centers $b=a$ and $b=0.01a$. If the photophoretic force has a large off-center ($b=a$), the rotation time can be shorter than the thermal relaxation time at $0.1 - 2$ AU for $100 \\ \\micron$ particles and at $4-40$ AU for 1 cm particles. In such regions, photophoresis is probably significantly suppressed. In order for photophoresis to work effectively in the whole disk, the off-center of the photophoretic force has to be as small as $b=0.01a$.  \\subsection{Characteristic Structure Due to Photophoresis} \\label{sec:discussion2}  The dust disk structure that photophoresis makes has three zones: the outer disk, the inner disk, and the inner hole. The outer disk is composed of small ($a \\la 1$ mm or $\\beta \\ga 0.01$) particles, and their dynamics is controlled mainly by radiation pressure. The inner disk is composed of large ($a \\ga 100 \\ \\micron$ or $\\beta \\la 0.01$) particles accumulating there due to photophoresis. The boundary between the outer and inner regions is at $10-100$ AU, depending on the gas density profile. The inner hole opens inside $\\sim 0.1$ AU, where the Knudsen number ${\\rm Kn}$ is smaller than unity for particles that photophoresis effectively works on ($a=100 \\ \\micron - 10$ cm).  The structure formed by photophoresis should be compared to the structure of gas-free disks. Modeling of gas-free dust disks also shows zonal structure that consists of the outer extended disk, the inner disk, and the inner hole (Th\\'ebault \\& Augereau 2005; Wyatt 2006; Strubbe \\& Chiang 2006; Krivov et al. 2006). In their models, a planetesimal belt is assumed and dust particles are continuously produced by planetesimal collisions. The inner disk is composed of large particles whose orbits are hardly affected by radiation pressure and are nearly circular. Thus, the location and width of the inner disk are basically similar to those of the planetesimal belt. Outside the planetesimal belt, the outer disk extends and it is composed of small particles whose $\\beta$-value is large. Their orbits are strongly influenced by radiation pressure and are excited to highly eccentric orbits. In the model of Th\\'ebault \\& Augereau (2005), the particles of the outer disk have $\\beta > 0.05$, which is relatively larger than $\\beta$ of our models. In our models, particles' $\\beta$ of the outer disk can be as small as $0.01$. Therefore, infrared to millimeter radio observations of outer disks and determination of the particle size will provide key information for determining which model is plausible.  Our models of photophoresis also predict that an inner hole opens in the dust disk, but the gas still fills the hole. This is a significant feature, but it is difficult to test this with the present observational techniques. We have to observe the gas, but the amount of the gas of optically thin disks is probably quite small, and most of Vega-type stars do not have gas in an amount exceeding the current detection limit. Further, the radius of the inner hole is of the order of $0.1$ AU, and the temperature there is close to the sublimation temperature of the dust. (The temperature exceeds 1500 K inside $0.03$ AU of our model disks.) Therefore, we need a careful observation that can distinguish between the inner holes made by photophoresis and by dust sublimation.  \\subsection{Summary}  Dust accumulation by photophoresis is studied. We use formulae of the photophoretic force that are applicable for the free molecular regime and for the slip-flow regime. The main results are as follows.  1. Particle accumulation occurs at a point where the outward acceleration on the gas by the pressure gradient equals to the outward acceleration on the particle by radiation pressure and photophoresis.  2. Photophoresis makes an inner disk composed of relatively large particles ($a= 100 \\ \\micron - 10$ cm). The inner disk extends between $0.1$ AU and $10-100$ AU, surrounded by the outer disk composed of small particles ($a \\la 1$ mm).  3. An inner hole opens inside $\\sim 0.1$ AU. The inner hole radius is determined by the condition ${\\rm Kn}=1$ for the maximum size particles that photophoresis effectively works on ($a=100 \\ \\micron - 10$ cm).  Photophoresis works effectively only when the disk is optically thin. Most of small ($\\la 10 \\ \\micron$) dust grains must be removed from the disk such that their column density to the star becomes smaller than $10^{-3} \\ {\\rm g \\ cm^{-2}}$. For example, at 1 AU, the dust density  must be smaller than $10^{-16} \\ {\\rm g \\ cm^{-3}}$, i.e., $10^{-5}$ times smaller than the value of the minimum mass solar nebula model (Hayashi et al. 1985). The gas disk also must be optically thin. Figure \\ref{fig:gasden} shows that the gas surface density must be smaller than $10^2 - 10^3 \\ {\\rm g \\ cm^{-2}}$. At $0.1$ AU, where the typical ray from the star enters the disk (for the dust particles at $r \\gg 0.1$ AU), this value is $\\sim 10^{-2}$ times smaller than the value of the minimum mass solar nebula model (Hayashi et al. 1985). Even in such a tenuous gas disk, the photophoretic force is strong enough to change the dust disk structure inside a few AU, as shown by model A (Figure \\ref{fig:req}$a$). If the gas density is more tenuous than model A, the region where photophoresis has a substantial effect shrinks toward the central star. In a gas disk with $0.1$ times smaller gas density than model A, photophoresis on millimeter sized particles works effectively only inside 1 AU, and in a disk with $10^{-2}$ times smaller gas density, the effective region shrinks to $0.3$ AU."
    },
    "0801/0801.2133_arXiv.txt": {
        "abstract": "I calculate the action of a satellite, infalling through dynamical friction, on a coplanar gaseous disk of finite radial extent.  The disk tides, raised by the infalling satellite, couple the satellite and disk.  Dynamical friction acting on the satellite then shrinks the radius of the coupled satellite-disk system.  Thus, the gas is ``shepherded'' to smaller radii.  In addition, gas shepherding produces a large surface density enhancement at the disk edge.  If the disk edge then becomes gravitationally unstable and fragments, it may give rise to enhanced star formation. On the other hand, if the satellite is sufficiently massive and dense, the gas may be transported from $\\sim 100$ pc to inside of a 10 to 10s of parsecs before completely fragmenting into stars.  I argue that gas shepherding may drive the fueling of active galaxies and central starbursts and I compare this scenario to competing scenarios.  I argue that sufficiently large and dense super star clusters (acting as the shepherding satellites) can shepherd a gas disk down to ten to tens of parsecs.  Inside of ten to tens of parsecs, another mechanism may operate, i.e., cloud-cloud collisions or a marginally (gravitationally) stable disk, that drives the gas $\\lesssim 1$ pc, where it can be viscously accreted, feeding a central engine. ",
        "introduction": "\\label{sec:intro}  The tidal interaction between small satellites and the gas or particle disks in which they are embedded is a subject of wide study in the planetary community.  These satellites are tidally coupled via disk tides, i.e., satellites excite spiral density waves at the Lindblad resonances in the disk (Goldreich \\& Tremaine 1978, 1980; Artymowicz 1993).  As a result of these tidal interactions, gaps can be opened up in the embedded disks as in the case of the shepherding moons of Saturn's rings (Goldreich \\& Tremaine 1978) or in type-II migration (Ward 1997).  Alternatively, if a gap does not open up, satellite may migrate inward rapidly due to tidal torque imbalances, i.e., type-I migration (Ward 1997).  The wide applicability of satellite-disk interactions in the planetary community raises an interesting question of whether the same physics may be applicable at larger scales, i.e., galactic scales.  In this paper, I will address this question by considering a very simple model, which illustrate the modification of the physics of satellite-disk interaction when applied on a galactic scale.  Most notably, the critical difference is that in addition to tidal torques, the satellite will experience dynamical friction.  The inclusion of dynamical friction produces a non-trivial effect.  Namely, dynamical friction on the satellite provides a sink of angular momentum in the system.  As a result, the coupled satellite-disk system will continually lose angular momentum and sink toward the center.  To study this physics, I consider a very simple model.  In my model, a satellite starts out in a circular orbit at a large radius and sinks toward the central mass concentration because of dynamical friction on the background stars.  Along its infall, it encounters a coplanar gaseous disk or ring, which initially has a finite radial extent, $r_{\\rm d,0}$.  Tides begin to couple the satellite with the disk. Because the satellite-disk system continues to suffer an ongoing loss of angular momentum from dynamical friction, it will shrink in radius. As a result gas can be transported on the dynamical friction timescale to smaller radii. In addition, I find that this satellite-disk interaction also builds a substantial surface density enhancement at the disk edge. This may lead to enhanced star formation at the disk edge.  This process which I call {\\it gas shepherding} may be a generic feature of gaseous disks around galaxies, if a sufficiently massive and dense satellite is available.  The physics of satellite-disk interactions in galaxies or gas shepherding is not just an interesting exercise in mathematical physics, but may be important in the fueling of active galaxies and central starbursts.  First, the action of dynamical friction on the satellite-disk system will shrink radius of the disk, thereby transporting gas to smaller radii.  This shepherding of gas will continue until the gas is forced into the center or the shepherding satellite is destroyed.  I will discuss this scenario for nuclear fueling and how it compares to competing scenarios in this paper.   I have structured this paper as follows.  In \\S\\ref{sec:basic picture}, I discuss the basic picture of gas shepherding and give a few order of magnitude estimates.  I estimate the timescale for gas shepherding and the expected surface density enhancement.  I present the physics of gas shepherding and solve numerical models in \\S\\ref{sec:shepherding}.  I also present an approximate analytic solution to this problem. I then discuss the application of gas shepherding to the feeding of active galactic nuclei and central starbursts in \\S\\ref{sec:applications}. Finally I present my conclusions in \\S\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  I have calculated the action of an infalling satellite on a coplanar gaseous disk.  The satellite falls inward because of dynamical friction and couples to the gaseous disk via disk tides.  Due to the unrelenting loss of angular momentum, the satellite-disk system shrinks in radius.  Thus, gas is transported from larger radii to smaller radii.  I calculate the structure of the evolving disk as a function of radius, both numerically and analytically and show a significant density enhancement at the disk edge.  This density enhancement is due to the transport of gas initially at larger radii to a ring that is $r_{\\rm H,d}$ thick radially.  The fate of this shepherded gas is not known.  I presented two possible outcomes: 1. the gas will fragment and form stars or 2. the gas is shepherded into the nuclear region.  In the latter case, I suggest that the gas shepherding scenario, which I have presented may be one means by which the ``100 parsec'' problem may be resolved.  For this to occur, gas must be shepherded sufficiently rapidly such that the gas is not completely consumed in star formation, which requires a sufficiently massive satellite ($\\sim 10^6\\,{\\rm M}_{\\odot}$).  I argue that such massive satellites exist on the form of SSCs.  In the model presented, the physics of star formation is a disk environment has not been fully addressed.  Observationally, galactic star-forming disks appear to be marginally unstable to fragmentation, i.e., $Q \\sim 1$ (Thompson et al 2005) and their star formation rate follows the Kennicutt law (Kennicutt 1998).  These two points suggests that some sort of feedback is responsible for maintaining this careful balance between fragmentation and heating.  How this feedback operates is not well understood.  As discussed previously, various models have invoked radiation pressure (Thompson et al. 2005), supernova feedback (Wada \\& Norman 2002), and cosmic ray pressure (Socrates et al. 2008) from star formation.  Alternatively, the turbulent pressure support may result from gravitational instability, i.e., gravitoturbulence (Gammie 2001).  It is unclear that in the presence of gas shepherding if this equilibrium can be maintained at the edge of the disk where the surface density is enhanced.  If it does not, this edge is likely to fragment completely into stars on a few dynamical times.  It is also unclear if the Kennicutt law would continue to hold at the disk edge.  Additional insights from future research on self-gravitating star-forming galactic disks will be needed to help formulate a more self-consistent calculation of the gas shepherding scenario, which will be needed to fully address the fate of the shepherded gas.  Assuming that the gas is not fully consumed in star formation, it would be shepherd down to 10 to 10s of parsecs (see \\S \\ref{sec:final 10 pc}). At that radii cloud-cloud collisions (Krolik \\& Begelman 1988; SBF90) or a starburst supported disk (Thompson et al.  2005) would allow the material to be viscously accreted.   The initial satellite inclination and/or eccentricity is likely to significantly affect gas shepherding.  In this paper, I have studied a special case where the orbit of the satellite is in the plane of the disk and approximately circular for simplicity.  An eccentric satellite would have its eccentricity damped via the excitation of spiral density waves.  If the eccentricity is below $r_{\\rm H, d}/r_{\\rm d}$, i.e., the satellite stays within one ``Hill radius'' of the disk of its outer edge, then the shepherding calculation presented should be a reasonable estimate of the shepherding timescales and velocities.  Compared to a coplanar satellite, an inclined satellite would couple less well to the gaseous disk.  For a sufficiently large inclination, the gas disk and satellite would be effectively decoupled.  However, its inclination could be sufficiently rapidly damped via induced bending waves in the gas disk in addition to spiral density waves (see for instance Ostriker 1994; Terquem 1998).  Once its inclination is damped below the critical value, where the disk and satellite become well coupled, the shepherding scenario presented in this paper would likely follow.  I would expect this critical inclination to be of order $\\sin i = r_{\\rm H,d}/r_{\\rm d}$, i.e., the satellite stays within of order a ``Hill radius'' of the disk above and below it.  Extending the present calculation to satellites of modest inclination would be fruitful and would allow a study of using the nucleated cores from minor mergers as shepherding satellites (see \\S\\ref{sec:shepherd destruction}). As discussed in \\S\\ref{sec:applications}, minor mergers may drive nuclear activity.  The initial inclination of minor mergers is arbitary, but their inclination is damped as they interact with the {\\it large scale} stellar and gaseous disk.  As a result, they may approach the shepherding scenario presented in this paper. Such a study is a subject for future work.  In addition, dynamical friction would also act on the satellite-induced spiral density waves.  Dynamical friction of the extended mass perturbation, i.e., the satellite-induced spiral density waves, may enhance the transfer of angular momentum between the gas disk and background stars (Tremaine \\& Weinberg 1984).  A study of these effects, while fruitful, is beyond the scope of this work.  The physics of gas shepherding is also important in other areas as well.  For instance, it may shape radial distribution of star clusters in nuclear gas rings.  In some of these nuclear ring systems, the location of the star clusters is exterior to the gas ring from which they are presumably formed.  For instance, NGC 4314 clearly show the star clusters exterior to the gas (Benedict et al. 2002).  Martini et al. (2003) found that the eight galaxies with strong bars all have star formation exterior to the dust ring in their sample of 123 galaxies.  How gas shepherding drives this morphology is the subject of a forthcoming work. (Van der Ven \\& Chang, in preparation)."
    },
    "0801/0801.3724.txt": {
        "abstract": "% % Context: ------------------- { Rings or annulus-like features have been observed in most imaged debris discs. Outside the main ring, while some systems (e.g., $\\beta$ Pictoris and AU Mic) exhibit smooth surface brightness profiles (SB) that fall off roughly as $\\sim r^{-3.5}$, others (e.g. HR 4796A and HD 139664) display large drops in luminosity at the ring's outer edge and steeper radial luminosity profiles. } % % Aims: ------------------------ { We seek to understand this diversity of outer edge profiles under the ``natural'' collisional evolution of the system, without invoking external agents such as planets or gas. } % % Methods: { We use a multi-annulus statistical code to follow the evolution of a collisional population, ranging in size from dust grains to planetesimals and initially confined within a belt (the \"birth ring\"). The crucial effect of radiation pressure on the dynamics and spatial distribution of the smallest grains is taken into account. We explore the dependence of the resulting disc surface brightness profile on various parameters. } % % Results: { The disc typically evolves toward a ``standard'' steady state, where the radial surface brightness profile smoothly decreases with radius as $r^{-3.5}$ outside the birth ring. This confirms and extends the semi-analytical study of Strubbe \\& Chiang (2006) and provides a firm basis for interpreting observed discs. Deviations from this typical profile, in the form of a sharp outer edge and a steeper fall-off, occur for two \"extreme\" cases: 1) When the birth ring is so massive that it becomes radially optically thick for the smallest grains. However, the required disc mass is probably too high here to be realistic. 2) When the dynamical excitation of the dust-producing planetesimals is so low ($<e>$ and $<i> \\leq 0.01$) that the smallest grains, which otherwise dominate the optical depth of the system, are preferentially depleted. This low-excitation case, although possibly not generic, cannot be ruled out by observations for most systems, . } %",
        "introduction": "\\subsection{the ubiquity of ring-like features}  Dusty debris discs have been detected by their infrared excess around $\\sim 15\\%$ of nearby main sequence stars \\citep[e.g.][]{back93}. More than a dozen of these discs have also been imaged, mainly in scattered light, since the initial observation of the $\\beta$ Pictoris system by \\citet{smi84}. One unexpected result from these images is that almost no system displays a smooth extended radial profile: the usual morphology is the presence of rings (or annuli) where the bulk of the dust population is located. This ring morphology is in fact so common that \\citet{stru06} pointed out that the debris disc phenomenon could more appropriately be renamed debris ``ring'' phenomenon.  Even for the archetypal debris ``disc'' $\\beta$ Pictoris, which has been imaged from 5 to a few thousand AU, the bulk of the dust is probably concentrated in a rather narrow region between 80 and 120 AU \\citep[e.g.][]{aug01}. One of the few systems actually resembling an extended ``smooth'' disc might be Vega, for which Spitzer mid--infrared images show a rather smooth radial luminosity profile \\citep{su05}. At the other end of disc morphologies, among the most striking ring features are the ones around HR4796 \\citep[e.g][]{Jay98,koer98,schnei99}, HD139664 \\citep{kal06}, and Fomalhaut \\citep{kal05}.   \\begin{table*} \\begin{minipage}{\\textwidth} \\caption[]{Geometry and surface brightness profile for a selection of debris disc systems resolved in scattered light images. } \\label{init} \\renewcommand{\\footnoterule}{} \\begin{tabular}{lcccc} System \\footnote{For many of these systems, many additional features, i.e., warps, clumps, etc., have been observed (it is especially true for \\bp) but we focus here on the main issue of average radial profiles}& Orientation & Detected Radial Extent\\footnote{The radial extents and surface brightness profiles are given for regions beyond the likely ``birth ring'' -- the region where scattering luminosity peaks and where most parent bodies are believed to reside (see Sec.\\,\\ref{sec:approach})} & Surface Brightness $SB \\propto r^{\\alpha}$ & Reference\\\\ \\hline HD 53143 & $\\sim$ face-on & $55-110$AU & $\\alpha \\sim -3$ & Kalas et al. (2006)\\\\ $\\beta$ Pic & edge-on & $127-193$AU & $\\alpha \\simeq -3$ & Golimovski et al.(2006) \\\\ & & $193-258$AU& $\\alpha \\simeq -4$ & \\\\ HD 32297 & edge-on & $80-170$AU & $\\alpha \\sim -3.5$ & \\citet{schnei99}\\\\ & & $170-350$AU & $\\alpha \\sim -3$ (averaged over the 2 ansae)& \\\\ AU Mic & edge-on & $32-210$AU & $\\alpha \\sim -3.8$ (averaged over the 2 ansae) & Fitzgerald et al (2007)  \\\\ HD139664 & edge-on & $83-109$AU& $\\alpha \\sim -4.5$ & Kalas et. al. (2006)\\\\ HD 107146 & $\\sim$ face-on & $130-185$AU & $\\alpha \\sim -4.8 \\pm 0.3$  & Ardila et al. (2004)\\\\ HD 181327 & face on & $100-200$ AU & $\\alpha \\sim -5$ & \\citet{schnei06}\\\\ Fomalhaut &  $66^{o}$ inclination  &  $140-160$ AU & $\\alpha \\sim -6.1$ & Kalas et al. (2006)\\\\ HR 4796A & $73^{o}$ inclination & $70-120$AU & $\\alpha \\sim -7.5$ & \\citet{wahhaj05} \\\\ \\hline \\end{tabular} \\label{table:observed} \\end{minipage} \\end{table*}  The characteristics that most differentiates one ring-like system from the other is the sharpness of the luminosity drop at the inner and outer edges of the rings. We shall in this paper focus on the outer edge issue, for which systems can be basically divided into two categories \\citep[see also][]{kal06}: \\begin{itemize} \\item ``Sharp edge'' rings, displaying abrupt surface brightness drops, sometimes as steep as $r^{-8}$ beyond the ring. The most representative members of this group are HR4796A and Fomalhaut. \\item ``Smooth edge'' rings, with no sharp outer edge and a surface brightness drop beyond the ring in $r^{-3}$ or $r^{-4}$. The most famous examples are here $\\beta$ Pic and AU Mic \\end{itemize} See Tab.\\ref{table:observed} for a list of outer-edge profiles for several resolved debris discs.   The presence of ring features is commonly attributed to some sculpting mechanisms including, in particular, the presence of a massive planet (see, for instance, \\citet{quillen07} for Fomalhaut, \\citet{frey07} for $\\beta$ Pictoris, \\citet{wyatt99} for HR4796A, or the more general study of \\citet{moro05}). The gravitational effect of such a planet can truncate or create gaps in the disc, either directly on the dust particles or indirectly on the planetesimals that produce these particles.   \\subsection{outer edges}  Nevertheless, while planets may be a natural and easy explanation for sculpting the inner edges, outer edges are a different problem. They are difficult to explain  in the light of one well established fact about debris discs, i.e. that the observed dust is not primordial but steadily produced by collisions, through a collisional cascade starting at much larger parent bodies, maybe in the planetesimal size range \\citep[e.g.][]{lag00}. In this respect, even if there is a sharp outer edge for the parent body population, collisions would constantly produce very small grains that would be launched by radiation pressure onto eccentric or even unbound orbits, thus populating the region beyond the outer edge and erasing the appearance of a narrow ring over timescales which might be shorter than those for gravitational sculpting by a planet. This issue is a critical one, since these small grains dominate the total geometric cross section and thus the flux in scattered light \\citep[see the discussion in][]{theb07}.  One possible confining mechanism for the sharp outer edge is gas. For discs transiting between gas-dominated phase (proto-planetary discs) to dust-dominated phase (debris discs), narrow dust rings may arise \\citep{klahr05}. \\citet{beslawu} further demonstrate that there exists an instability with which the residual gas collects grains of various sizes (even those subject to radiation pressure) into a narrow belt. However, this mechanism requires at least comparable amount of gas and dust. This may be difficult to justify for most debris discs, which are evolved systems where the amount of gas is probably too low to prevent the smallest grains to be launched onto very eccentric orbits smoothing out any sharp outer edge.  ``Razor sharp'' outer edges are thus very difficult to explain in the presence of this unavoidable outward launching of small grains. However, although no perfect abrupt outer edge has indeed been observed \\footnote{HD141569A could be a possible exception, but this system is  likely not a true debris disc, being significantly younger and gas rich \\citep{jonk06}, but a member of the loosely defined \"transition object\" category.} (unlike inner edges, which are in some cases, like Fomalhaut, almost razor sharp), a great variety of outer edge profiles does exist, from relatively smooth to very steep (see Tab.\\,\\ref{table:observed}).  In this study, we address the issue of how these different profiles can be physically achieved: can this diversity be explained by the sole ``natural'' evolution of a collisional active disc steadily producing small, radiation-pressure affected grains, or is (are) additional mechanism(s) needed? We consider initial conditions which are {\\it a priori} the most favourable for creating sharp edges, by assuming a population of large parent bodies confined within an annulus with an abrupt cutoff at its outer edge. How this initial confinement may have come about is itself an interesting question but is not the focus of the current paper (see however the discussion in Sec.\\,\\ref{sec:real}). The outcome we consider as a reference for our investigation is the radial surface brightness (SB) profile in scattered light, since this is an observable which is reasonably well constrained for most imaged debris discs (either directly observed or obtained by de-projection). We consider the nominal case of a disc seen edge-on, but results can easily be extrapolated to face-on systems, since SB profiles for both orientations tend towards the same radial dependence far from the birth ring (given the same radial dust distribution).  In several previous studies, all addressing the specific $\\beta$ Pic case \\citep{lec96, aug01, theb05}, it has been argued that the ``natural'' SB profile outside the collisionaly active parent bodies belt, or ``birth ring'', falls off as $\\sim r^{-5}$. This is based on the assumption that all particles produced in the birth ring have a size distribution which scales as $dN/ds \\propto s^{-3.5}$, \\citep[as expected for an idealized infinite collisional cascade at equilibrium, see][]{dohn69}, down to the radiation blow--out limit $s=s_{0.5}$ (where the ratio of radiation pressure to gravity $\\beta = 0.5$). The smallest radiation pressure-affected grains, which dominate the light receiving area, are then diluted along their eccentric orbit. This geometrical spread results in $SB \\propto r^{-5}$. However, \\citet{stru06} argued that, since high-$\\beta$ grains spend a long time in the collisionaly inactive region beyond the birth ring, the $dN/ds \\propto s^{-3.5}$ collisional equilibrium law should only apply to the small fraction $f$ of these grains which are present in the collisionaly active birth ring. This results in a $1/f$ excess of the disc-integrated number of small grains, which in turn results in a flatter $SB$ profile in $r^{-3.5}$. They applied their theory to the AU Mic disc \\footnote{AU Mic is an M star with weak radiation but where stellar wind is believed to act on small grains in an equivalent way as radiation pressure does around more massive A stars \\citep[e.g][]{aug06}.} and reproduced the observed SB profile, spectral energy distribution and disc colour. \\citet{stru06} further argued that the observed SB profile depends only weakly on the radial and size distributions of grains within the birth ring. The discs which exhibit a fall-off sharper than $r^{-3.5}$ are thus puzzling in the face of this theory.  The innovative model of \\citet{stru06} is build on analytical derivations and Monte-Carlo modeling which did not actually treat the collisional evolution of the system and relies on several simplifying assumptions. The main one is that the size distribution is fixed and is assumed to follow the idealized $dN/ds \\propto s^{-3.5}$ scaling (corrected by the fraction $1/f$), whereas several studies have shown that this law cannot hold in real systems \\citep[see][and references therein]{theb07} \\footnote{The need for realistic size distributions departing from fixed power laws has been very recently emphasized by \\citet{fitz07} in their analysis of the AUMic data}. Another issue is that when evaluating collisional life-times, only the vertical velocity of the grains was taken into account, thus neglecting their radial movement which can be appreciable, if not dominant for the smallest grains. Finally, the specific dynamics of the small radiation-pressure-affected grains, in particular the fact that they suffer much more frequent collisions and at much higher velocities, is not taken into account. ",
        "conclusions": "We numerically investigate if the diverse outer surface brightness profiles observed in debris discs can be explained by the ``natural'' collisional evolution of belts made of solid particles ranging from planetesimal to micron-sized dust grains. We consider an initial ring of parent bodies with a razor-sharp outer edge and quantitatively examine to what extent the steady collisional production of small, radiation-pressure-affected grains modifies the initially perfect ring-like structure.  Concurring with the pioneering results of \\citet{stru06}, our numerical explorations have shown that for most ``reasonable'' parameters (mass, ring width, dynamical excitation) of a collisionally evolving debris ring, the surface brightness outside the ring naturally tends towards a standard profile, with no sharp drop at the outer edge and a mid-plane surface brightness profile $\\propto r^{-3.5}$. We confirm this result by simple but robust analytical considerations. This nominal result is in good agreement with the luminosity profiles observed in the outer regions of several debris discs, including \\bp\\ and AUMic.  Some observed systems however exhibit steeper radial luminosity profiles. Our numerical exploration shows that sharper outer edges and profiles steeper than that of the nominal case can only be obtained for two ``extreme'' cases: \\begin{itemize} \\item 1) For a system with very high dust content (typically $M_{dust}\\geq10M_{\\oplus}$), the radial optical depth is raised to near unity for small grains. Most of these high-$\\beta$ grains cannot travel out of the birth ring without suffering a collision. This allows the outer edge to sharpen. However, systems with such high radial optical depth are expected to wear down with time because of strong collisional erosion. Moreover, the existence of such very dusty systems is not backed by observations. This case does not appear as a realistic option. \\item 2) A longer survival of the sharp outer edge is achieved for systems with normal dust content that are \\emph{dynamically cold}, with typically $\\langle e \\rangle =2\\langle i \\rangle \\leq 0.01$. In this case, small grains are destroyed much more efficiently than they are created, leading to a depletion of this population. The system's optical depth and luminosity are then dominated by large grains which do not leave the main birth ring, leading to a sharp outer edge. Even if this case might not correspond to the most generic debris disc configuration, it cannot be ruled out by observations and is thus a possible explanation to some of the observed systems. \\end{itemize}  To numerically investigate the applicability of the dynamically cold case to real sharp-edge systems, we consider the specific case of HR4796A. We find a reasonably good fit of this system's outer region luminosity profile with a dynamical excitation $\\langle e \\rangle \\sim 0.0035$. There is thus the possibility that such a sharp outer edge could be explained by the natural collisional evolution of a confined disc of large parent bodies. %(the confinement of the parent bodies %being possibly due to some past punctual or transitory event or %to the natural outer limit of the protoplanetary nebulae)."
    },
    "0801/0801.3952_arXiv.txt": {
        "abstract": "In 1981, a faint radio source (G$'$) was detected near the center of the lensing galaxy of the famous ``twin quasar'' Q0957+561. It is still unknown whether this central radio source is a third quasar image or an active nucleus of the lensing galaxy, or a combination of both. In an attempt to resolve this ambiguity, we observed Q0957+561 at radio wavelengths of 13~cm, 18~cm, and 21~cm, using the Very Long Baseline Array in combination with the phased Very Large Array and the Green Bank Telescope.  We measured the spectrum of G$'$ for the first time and found it to be significantly different from the spectra of the two bright quasar images. This finding suggests that the central component is primarily or entirely emission from the foreground lens galaxy, but the spectrum is also consistent with the hypothesis of a central quasar image suffering free-free absorption.  In addition, we confirm the previously-reported VLBI position of G$'$ just north of the optical center of the lens galaxy.  The position slightly favors the hypothesis that G$'$ originates in the lens, but is not conclusive. We discuss the prospects for further clarification of this issue. ",
        "introduction": "When a galaxy lies close to the line of sight to a quasar, multiple images of the quasar may be produced.  Almost all known lensed quasars have an even number of images---either two or four---despite a mathematical proof that nonsingular lens models always produce an odd number of images \\citep{dyer80a,burke81a}.  The missing image is the ``central image,'' corresponding to the central maximum in light-travel time.  Central images are hard to detect because they are highly demagnified by the dense core of the lensing galaxy.  They are worth finding, however, because the properties of central images are unique probes of the inner 10-100 parsecs of galaxies that are too distant to resolve with ordinary observations.  Even non-detections of central images can be useful constraints on the central structure of galaxies.  An upper limit on the flux density of a central image corresponds to a lower bound on the central surface density of the lens galaxy.  In the context of particular galaxy models, the absence of a detectable central image can be used to set the maximum size of any constant-density core \\citep{narasimha86a,wallington93a} or the steepness of any central density cusp \\citep{rusin01a,evans02a,keeton03b}.  Supermassive black holes also affect central images: they can prevent the formation of a central image, or produce an additional image \\citep{mao01a,bowman04a,rusin05a}. Microlensing by stars in the lens galaxy can demagnify a central image even further \\citep{bernstein99a,dobler07a}.  Besides the faintness of the central image, there are two other reasons why it has proven difficult to identify central images. One reason is the close proximity between the expected position of the central image and the center of the lensing galaxy. The expected separation is less than $\\sim$10 milliarcseconds for an isolated lensing galaxy, although it can be larger when the galaxy is part of a cluster that provides additional magnification \\citep{inada05a}. Resolving such a small separation would often require space-based imaging at optical or near-infrared wavelengths, or Very Long Baseline Interferometry (VLBI) at radio wavelengths.  A second and related reason is that even when a central point source is detected, it is difficult to decide conclusively whether it is a central image of a background quasar, or an active galactic nucleus (AGN) in the lens, or a combination of these two possibilities. This decision is difficult not only because the point source is faint and challenging to characterize, but also because propagation effects may alter the properties of the central image and cause it to appear different from the other lensed images, thereby interfering with the usual tests for lensed image identification. Such propagation effects include extinction by dust at optical wavelengths, and scintillation and free-free absorption at radio wavelengths; these effects are expected to be more important for lines of sight passing very near the center of the foreground galaxy, where on average the density of dust and plasma should be greatest.  A central image was detected for the bright optical lens APM~08279+5255 \\citep{ibata99a,egami00a,munoz01a}, but it seems probable that this image is the result of a highly elongated mass distribution rather than the high central density of the lensing galaxy \\citep{lewis02a}.  More recently, \\cite{winn04a} presented evidence for a central image in PMN~J1632--0033, thereby setting an upper limit on the mass of any central black hole and a lower limit on the central surface density of the lens galaxy. A central image has also been detected at optical wavelengths for SDSS~J1004+4112 \\citep{inada05a}.  In addition, there are cases in which central radio components have been detected but are most likely to be the active nucleus of the lens galaxy \\citep{chen93a,fassnacht99b}, and cases where recent work has placed stringent upper limits on the flux density of any central component \\citep{boyce06a,zhang07a}.  This paper is concerned with the very first lensed quasar to have been discovered, Q0957+561 \\citep{walsh79a}. More than 20 years ago, a central radio component was detected in this system \\citep{gorenstein83a}, but there is still a lingering uncertainty about its interpretation.  The system consists of a quasar at redshift 1.41 that is gravitationally lensed into 2 images (A and B) by an elliptical galaxy G1 and a surrounding set of galaxies mostly at redshift 0.36 \\citep{stockton80a,young81a}. Quasar images A and B appear as point sources in optical images and as core-jet sources in radio images. Observations with the Very Large Array (VLA) reveal a radio source, dubbed G, close to the center of the optically-identified lens galaxy G1 \\citep{greenfield80a,roberts85a}. Radio observations with VLBI have revealed an unresolved point source near the center of G1 \\citep{gorenstein83a}. The VLBI source is known by the different name G$'$ because its relationship to G is unclear.  The VLA and VLBI observations are sensitive to very different spatial scales of radio emission; typically, VLBI observations have an angular resolution 10-100 times finer than VLA observations and are blind to most structures that are resolved by the VLA.  Modelers of this lens have typically assumed that G$'$ marks the center of the lensing galaxy, and that any central quasar image is undetectably faint \\citep[see, e.g.,][]{barkana99a,bernstein99a,chae99a}. If instead G$'$ is the central image, then the lens galaxy has a shallower central gravitational potential than has been assumed by those modelers. Traditionally this identification has not been a major concern, because the lens modeling literature has been preoccupied with the connection between the Hubble constant and measurements of the A--B time delay, which is not greatly affected by the innermost $\\sim$100~pc of the galaxy mass distribution. However, as noted above, the interpretation of G$'$ does have an impact on our understanding of the mass distribution of the central regions of elliptical galaxies.   We undertook new radio observations of Q0957+561 in an attempt to clarify the nature of G$'$.  Our strategy was to compare the radio spectral-energy distributions of A, B, and G$'$ over as large a range of radio wavelengths as feasible.  For all practical purposes, gravitational lensing is wavelength-independent. Thus, the central image would have the same spectrum as images A and B, were it not for the confounding factors of intrinsic variability (in conjunction with different time delays for the different images), the magnification gradient along the jet images (in conjunction with different spectral energy distributions for the core and jets), and propagation effects specific to the line of sight of each image (such as scintillation and free-free absorption), which must be taken into account. An AGN in the lens would not have the same spectrum as the quasar images except by coincidence.  Although the Q0957+561 system has been studied extensively at radio wavelengths, flux-density measurements of G$'$ have not been well documented.  Previously, the only flux-density measurement with a reported uncertainty was by \\cite{gorenstein83a} at 13~cm. The VLA source G has been documented more completely, including measurements at wavelengths of 2~cm, 3.6~cm, and 6~cm by \\cite{harvanek97a}.   In the following section, we describe our VLBI observations of this system and G$'$ in particular. We describe the data reduction procedures in \\S\\ref{reduction}, and we present the empirical results in \\S\\ref{results}. In \\S\\ref{discussion} we discuss the spectrum, and in \\S\\ref{future} we review the prospects for further progress on this issue. ",
        "conclusions": ""
    },
    "0801/0801.4936.txt": {
        "abstract": "Following the recent abundance measurements of Mg, Al, Ca, Fe, and Ni in the black hole X-ray binary \\mbox{XTE J1118+480} using medium-resolution Keck~II/ESI spectra of the secondary star (Gonz\\'alez Hern\\'andez et al. 2006), we perform a detailed abundance analysis including the abundances of Si and Ti. These element abundances, higher than solar, indicate that the black hole in this system formed in a supernova event, whose nucleosynthetic products could pollute the atmosphere of the secondary star, providing clues on the possible formation region of the system, either Galactic halo, thick disk, or thin disk. We explore a grid of explosion models with different He core masses, metallicities, and geometries. Metal-poor models associated with a formation scenario in the Galactic halo provide unacceptable fits to the observed abundances, allowing us to reject a halo origin for this X-ray binary. The thick-disk scenario produces better fits, although they require substantial fallback and very efficient mixing processes between the inner layers of the explosion and the ejecta, making quite unlikely an origin in the thick disk. The best agreement between the model predictions and the observed abundances is obtained for metal-rich progenitor models. In particular, non-spherically symmetric models are able to explain, without strong assumptions of extensive fallback and mixing, the observed abundances. Moreover, asymmetric mass ejection in a supernova explosion could account for the required impulse necessary to launch the system from its formation region in the Galactic thin disk to its current halo orbit. ",
        "introduction": "The low-mass X-ray binary \\mbox{XTE J1118+480} is the first identified black hole moving in Galactic halo regions (Wagner et al. 2001; Mirabel et al. 2001). Since it was discovered during a faint outburst on UT 2000 March 29 (Remillard et al.\\ 2000), it has been intensively studied in both the X-ray and optical spectral regions. During the decay of the outburst, McClintock et al. (2001) and Wagner et al. (2001) determined the radial-velocity curve of the companion star, yielding a mass function $f(M) \\approx 6$~\\Msuno. The companion star was classified as a late-type main-sequence star with a mass of 0.1--0.5~{\\Msuno} (Wagner et al.\\ 2001).  By modelling the light curve, McClintock et al. (2001) derived a lower limit to the orbital inclination, $i \\ga 55^\\circ$, and consequently an upper limit to the black hole mass of $M_{\\rm BH} \\la 10$~\\Msuno. Additional evidence for a high inclination ($i \\ga 60^\\circ$) comes from measurements of tidal distortion (Frontera et al. 2001), whereas the lack of dips or eclipses for a Roche-lobe filling secondary yields upper limits of $i \\ga 80^\\circ$ and $M_{\\rm BH} \\ga 7.1$~\\Msuno. Later, Gelino et al. (2006) derived an orbital inclination of $68^\\circ\\;\\pm\\;2^\\circ$, by modeling the optical and infrared ellipsoidal light curves of the system in quiescence. This value of the inclination allowed them to better constrain the black hole mass at $M_{\\rm BH} = 8.53 \\pm 0.60$~\\Msuno.  The system is placed in the Galactic halo, with an extraordinarily high Galactic latitude ($b \\approx 62.3^{\\circ}$), and a height of $\\sim$1.6 kpc above the Galactic plane, according to its distance of $1.85\\pm0.36$ kpc (Wagner et al.\\ 2001). This appears surprising since all other black hole binaries are located in the Galactic disk. An accurate measurement of its proper motion coupled with its distance provides space-velocity components ($U$, $V$) which seem consistent with those of some old halo globular clusters (Mirabel et al.\\ 2001). This opened the possibility that the system originated in the Galactic halo, and therefore, that the black hole could be either the remnant of a supernova (SN) in the very early Galaxy or the result of direct collapse of an ancient massive star. However, the galactocentric orbit crossed the Galactic plane many times in the past, and the system could have formed in the Galactic disk and been launched into its present orbit as a consequence of the ``kick'' imparted during the SN explosion of a massive star (Gualandris et al.\\ 2005). Recent observations with the 10-m Keck~II telescope revealed that the secondary star has a supersolar surface metallicity ($[{\\rm Fe/H}]=0.2\\pm0.2$, Gonz\\'alez H\\'ernandez et al. 2006), confirming the origin of the black hole in a SN event. Thus, if the system originated in the Galactic halo, the element abundances of the secondary star must have been enriched by a factor of 5--25 depending on whether its initial metallicity resembled a thick-disk star or a halo star.  Element abundances of secondary stars of X-ray binaries have been studied for the systems Nova Scorpii 1994 (Israelian et al. 1999; Gonz\\'alez Hern\\'andez et al. 2007), A0620--00 (Gonz\\'alez Hern\\'andez et al. 2004), Centaurus X-4 (Gonz\\'alez Hern\\'andez et al. 2005), \\mbox{XTE J1118+480} (Gonz\\'alez Hern\\'andez et al. 2006, hereafter Paper~I), and V4641 Sagittarii (Orosz et al. 2001; Sadakane et al. 2006). All of these X-ray binaries show metallicities close to solar independent of their location with respect to the Galactic plane, and possible scenarios of pollution from a SN or hypernova have been discussed. In this paper, we compare in detail different scenarios of the possible enrichment of the secondary star from SN yields, providing conclusions on the formation region (Galactic halo, thick disk, or thin disk) of this halo black hole X-ray binary.  \\begin{figure}[ht!] %\\epsscale{.80} %\\plotone{f1.eps} \\centering \\includegraphics[width=6.2cm,angle=90]{f1.eps} \\caption{\\footnotesize{Radial velocities of \\mbox{XTE J1118+480} folded on the orbital solution of the data, together with the best-fitting sinusoid. Individual velocity uncertainties are $\\le 7$ ${\\rm km}\\ {\\rm s}^{-1}$ and are not plotted because they are always smaller than the symbol size. The bottom panel shows the residuals of the fit.}} \\label{fig1} \\end{figure} ",
        "conclusions": "We have presented Keck~II/ESI medium-resolution spectroscopy of the black hole binary \\mbox{XTE J1118+480}. The individual spectra of the system allowed us to derive an orbital period of $P = 0.16995 \\pm 0.00012$~d and a radial velocity semiamplitude of the secondary star of $K_2 = 708.8 \\pm 1.4$ \\kmso. The implied updated mass function is $f(M) = 6.27 \\pm 0.04$~\\Msun, consistent with (but more precise than) previous values reported in the literature. Inspection of the high-quality averaged spectrum of the secondary star provides a rotational velocity of $v~\\sin~i = 100^{+3}_{-11}$ \\kmso, and hence a binary mass ratio $q = 0.027 \\pm 0.009$. The derived radial velocity, $\\gamma = 2.7 \\pm 1.1$ \\kmso, of the center of mass of the system agrees, at the 3$\\sigma$ level, with the results of previous studies.  We have performed a detailed chemical analysis of the secondary star. We applied a technique that provides a determination of the stellar parameters, taking into account any possible veiling from the accretion disk. We find $T_{\\mathrm{eff}} = 4700 \\pm 100$~K, $\\log [g/{\\rm cm~s}^2] = 4.6 \\pm 0.3$, $\\mathrm{[Fe/H]} = 0.18 \\pm 0.17$, and a disk veiling (defined as $F_{\\rm  disk}/F_{\\rm total}$) of $\\sim$40\\% at 5000~{\\AA}, decreasing toward longer wavelengths.  We have provided further details on the abundances of Mg, Al, Ca, Fe, Ni, and Li already reported by Gonz\\'alez Hern\\'andez et al. (2006), and we determined new element abundances of Si and Ti. The chemical abundances are typically higher than solar, and in some cases they are slightly enhanced (e.g., Mg, Al, and Si) in comparison with the abundances of these elements in stars of the solar neighborhood having similar iron content.  The present location and kinematics of this binary system had suggested that it could have originated in the Galactic halo. However, the chemical abundances strongly indicate that the black hole formed as a consequence of a supernova/hypernova explosion that occurred within the binary system. This explosive event must have either provided a kick to the system if it was formed in the thin disk or enriched significantly the atmosphere of the secondary star if the system formed in the thick disk or halo.  We have explored a variety of supernova/hypernova explosion models for different metallicities, He core masses, and geometries. We compared the expected abundances in the secondary star after contamination from nucleosynthetic products from different initial metallicities of the secondary star ($-2.2 < \\mathrm{[Fe/H]} < 0.2$), to investigate the formation region in the Galactic halo, thick disk, or thin disk. Metal-poor explosion models ($Z = 0$ and $Z = 0.001$) were not able to fit the observed abundances since they produce inappropriate ratios between $\\alpha$-elements and iron-peak elements, and they are extremely sensitive to the adopted mass-cut. This comparison probably rules out an origin in the Galactic halo for this black hole binary.  For the thick-disk scenario, we carefully inspected the model predictions, and although they provide better fits to the observed abundances, they require substantial fallback (up to 5.5~\\Msun) and very efficient mixing processes between the inner layer of the explosion and the ejecta. We thus conclude that this scenario is unlikely.  Metal-rich spherically symmetric models for the thin-disk scenario were able to fairly reproduce the observed abundances, although they do not easily provide the energy required to launch the system from the Galactic plane to its current halo orbit.  Finally, non-spherically symmetric models produce excellent agreement with the observed element abundances in the secondary star without invoking extensive fallback and mixing. In addition, asymmetric mass ejection would naturally provide the kick to expel this binary system from its birth place in the Galactic thin disk, which seems to be the most plausible explanation for the origin of this halo black hole X-ray binary."
    },
    "0801/0801.2435_arXiv.txt": {
        "abstract": "We present a calculation of the sedimentation of grains in a giant gaseous protoplanet such as that resulting from a disk instability of the type envisioned by Boss (1998).  Boss (1998) has suggested that such protoplanets would form cores through the settling of small grains.  We have tested this suggestion by following the sedimentation of small silicate grains as the protoplanet contracts and evolves.  We find that during the course of the initial contraction of the protoplanet, which lasts some $4\\times 10^5$ years, even very small ($>1\\mu $m) silicate grains can sediment to create a core both for convective and non-convective envelopes, although the sedimentation time is substantially longer if the envelope is convective, and grains are allowed to be carried back up into the envelope by convection. Grains composed of organic material will mostly be evaporated before they get to the core region, while water ice grains will be completely evaporated.  These results suggest that if giant planets are formed via the gravitational instability mechanism, a small heavy element core can be formed due to sedimentation of grains, but it will be composed almost entirely of refractory material.  Including planetesimal capture, we find core masses between 1 and 10 M$_{\\oplus}$, and a total high-Z enhancement of $\\sim$ 40 M$_{\\oplus}$.  The refractories in the envelope will be mostly water vapor and organic residuals. ",
        "introduction": "The mechanism of giant plant formation is still a matter of debate. The two main candidates are core accretion (see, e.g. Pollack et al. 1996) and disk instability (see, e.g. Boss 1997). In the core accretion scenario, kilometer-sized planetesimals collide and accrete to form a core.  As this core grows, it begins to attract the surrounding gas.  By the time the core has reached a mass of some 10 $M_{\\oplus}$, the accretion rate of the gas becomes very high and a Jupiter-mass object is formed (Pollack et al. 1996, Hubickyi et al. 2005). \\newpage The alternative scenario invokes a local gravitational instability in the gas disk surrounding a young star.  Such an instability can lead to the creation of gas clumps which evolve to become giant plants. Such a mechanism of planet formation would seem to imply that the resultant planet has a solar ratio of elements. Observations of Jupiter indicate that the planet's envelope is enriched in heavy elements by a factor of $\\sim 3$ over the solar ratio to hydrogen (Young 2003).  In addition, recent models of Jupiter's interior that fit the gravitational moments (Saumon and Guillot, 2004) indicate that Jupiter's core is between 0--6 $M_{\\oplus}$, smaller than earlier estimations. However, the overall enhancement of heavy elements in the planet, according to these models, is of the order of 20 -- 40 $M_{\\oplus}$, or 3 -- 6 times the solar value.  \\noindent The core instability model requires a core of the order of $\\sim 10 M_{\\oplus}$ to reach the rapid gas accretion stage, while the disk instability can, in principle, form a giant planet with no core at all.  Indeed a simple interpretation of this scenario would seem to require this.  Therefore, if Jupiter does turn out to have a substantial core, the disk instability model must provide a mechanism for forming one.  Boss (1997, 1998) has suggested that a solid core can be formed by sedimentation of dust grains to the center of the protoplanet before the protoplanet contracts to planetary densities and temperatures. In this paper we investigate this possibility by applying the microphysics of grain coagulation and sedimentation to silicate grains in a Jupiter-mass protoplanet with initial conditions similar to those in a newly formed clump. ",
        "conclusions": "We present a calculation of grain sedimentation inside a protoplanet formed via the disk  instability scenario. During the first $\\sim 10^5$ yr the clump is cold enough to let silicate grains settle to the center and create a heavy element core. This is true for all grain sizes we considered ($a_0\\geq 10^{-4}$ cm). Grains smaller than this will have high enough number densities so that they will grow quickly into the size range we considered. Core formation proceeds whether the envelope is convective or not, although the convective envelope delays core formation substantially for smaller grains.  In all cases a core can be formed before the temperatures near the center get high enough to evaporate the silicate grains. If the silicate material that reaches the central region and is incorporated into a core is immune to further dissolution, core masses of some 4 M$_{\\oplus}$ can be obtained.  This includes material added to the planet later via planetesimal capture.  If the planetesimals are of the order of 100 km radius or larger, they will not be captured until the protoplanet contracts sufficiently. This delay means that the envelope temperatures are higher when this additional mass settles, and not all the grains can reach the core before they are vaporized.  For such large planetesimals, we find that the core mass will be only $\\sim$ 1.7 M$_{\\oplus}$.  Grains consisting of ice or more volatile material cannot sediment quickly enough to avoid being vaporized before they reach the core region.  Grains consisting of CHON depend on the volatility chosen.  For hexacosane the grains evaporate before reaching the core at all the sizes we considered.  For CHON that has the volatility suggested by Obrec (2004), only very large grains ($\\gtrsim 10 $ cm) can reach the core, and these add at most a few tenths of an Earth mass to the total core mass.  The remainder of this high-Z material would remain in the envelope, giving a high-Z mass in the envelope of some 30 M$_{\\oplus}$.  Our simulations thus produce both core and envelope high-Z masses that agree very well with the values determined by fitting the gravitational moments of Jupiter (Guillot 1999, Saumon and Guillot 2004).  In this work, two very important effects have been neglected.  The first is radiative heating by the surrounding medium.  As Kovetz et al. (1988) have shown, such radiation tends to be deposited at a point where the optical depth is approximately equal to the ratio of the incoming flux to the outgoing flux.  In our case, the isolated planet initially has a photospheric temperature of $\\sim 25$~K.  If we take the protoplanet to be at 5 AU, the intervening gas and dust will have a high optical depth in the visible [e.g. Chiang and Goldreich (1997)], and the radiation falling onto the protoplanet will be mostly at the ambient temperature of the surrounding gas.  This is model-dependent, but is of the order of 100 K (Lecar et al. 2006) so that the ratio of incoming to outgoing flux can be quite large.  This will cause higher temperatures in the interior, and will make it more difficult for grains to survive long enough to sediment to the core.  On the other hand, in the early stages the planet will stay extended longer, and this will allow more time for the grains to settle.  A second effect is the contraction of the gas near the center by the increased pressure due to the core itself.  This too will heat the gas near the core, and will inhibit the grains from settling.  This second effect seems more serious, since, in this scenario, the formation of the core itself inhibits further core growth.  A proper treatment of this effect requires an equation of state for the core material, and work on this is in progress."
    },
    "0801/0801.2573_arXiv.txt": {
        "abstract": "{ Variability is a common characteristic of magnetically active stars.  Flaring variability is usually interpreted as the observable consequence of transient magnetic reconnection processes happening in the stellar outer atmosphere.  Stellar flares have been observed now across 11 decades in wavelength/frequency/energy; such a large span implies that a range of physical processes takes place during such events. Despite the fact that stellar radio flares have long been recognized and studied, key unanswered questions remain.  I will highlight what, in my opinion, are some of these questions. I will also describe recent results on stellar flare emissions at radio wavelengths, discussing the nature of coherent and incoherent emissions and the prospects of wide-field radio imaging telescopes for studying such events. }  \\FullConference{Bursts, Pulses and Flickering:Wide-field monitoring of the dynamic radio sky\\ June 12-15 2007\\\\ Kerastari, Tripolis, Greece}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.0740_arXiv.txt": {
        "abstract": "\\noindent  We investigate a scenario where the recently discovered non-thermal hard X-ray emission from the Ophiuchus cluster originates from inverse Compton scattering of energetic electrons and positrons produced in weakly  interacting dark matter pair annihilations. We show that this scenario can account for both the X-ray and the radio emission, provided the average magnetic field is of the order of \\mbox{0.1 $\\mu$G}. We demonstrate that GLAST will conclusively test the dark matter annihilation hypothesis. Depending on the particle dark matter model, GLAST might even detect the monochromatic line produced by dark matter pair annihilation into two photons. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.3269_arXiv.txt": {
        "abstract": "We study the effects of substructure in the Galactic halo on direct detection of dark matter, on searches for energetic neutrinos from WIMP annihilation in the Sun and Earth, and on the enhancement in the WIMP annihilation rate in the halo. Our central result is a probability distribution function (PDF) $P(\\rho)$ for the local dark-matter density.  This distribution must be taken into account when using null dark-matter searches to constrain the properties of dark-matter candidates.  We take two approaches to calculating the PDF. The first is an analytic model that capitalizes on the scale-invariant nature of the structure--formation hierarchy in order to address early stages in the hierarchy (very small scales; high densities). Our second approach uses simulation-inspired results to describe the PDF that arises from lower-density larger-scale substructures which formed in more recent stages in the merger hierarchy.  The distributions are skew positive, and they peak at densities lower than the mean density.  The local dark-matter density may be as small as 1/10th the canonical value of $\\simeq0.4\\,\\GeV~\\cm^{-3}$, but it is probably no less than $0.2\\,\\GeV~\\cm^{-3}$. ",
        "introduction": "The nature of dark matter remains a mystery.  Among the plethora of dark-matter candidates, there are two classes that are sufficiently promising to motivate major experimental searches. The first is a weakly interacting massive particles (WIMP) \\cite{Jungman:1995df,BHS05}, which may arise in supersymmetric extensions of the standard model or in theories that include universal extra dimensions (UED) \\cite{Hooper:2007qk}.  The other leading candidate is the axion \\cite{axionreviews}, hypothesized to solve the strong-CP problem.  WIMPs may be detected indirectly through observation of gamma rays or cosmic-ray positrons, antiprotons, and/or antideuterons produced when WIMPs annihilate in the Galactic halo or in the halos of some extragalactic systems.  WIMPs might also be detected via observation of energetic neutrinos produced by annihilation of WIMPs that have accumulated in the Sun and/or Earth. There is also an effort to detect dark-matter particles in low-background experiments via detection of the $O(100\\,{\\rm keV})$ energy they impart to nuclei from which they elastically scatter.  Likewise, dark-matter axions are being sought directly by conversion to photons in resonant-cavity experiments \\cite{axionexperiments}.  If dark matter is composed of WIMPs, we may also get some clue to its nature from forthcoming LHC accelerator experiments \\cite{lhc}.  Predictions for the event rates for any of these detection schemes depend on the distribution of dark matter in the Galactic halo.  The flux of gamma and cosmic rays from WIMP annihilation depends on an integral of the square of the dark matter density over volume, while the rates for energetic neutrinos and direct detection depend only on the local dark-matter density.  In early calculations, it was assumed simply that dark matter was smoothly distributed in the Galactic halo, with either an isothermal or NFW \\cite{NFW96} radial density profile, and with a local dark-matter density fixed by Milky Way dynamics to be $\\rhosun\\simeq0.4~\\GeV~\\cm^{-3}$.  It has become clear in recent years, from theory and N-body simulations, that the distribution of dark matter in the Galactic halo is not likely to be perfectly smooth, and that some of the dark matter in the Galactic halo will be distributed into subhalos with a variety of sizes \\cite{substructure}. This substructure hierarchy may extend to very small scales \\cite{Boehm,Green:2003un,Diemand:2005vz,Diemand:2007qr,Ando:2005xg}. For WIMPs in supersymmetric and UED models, the cutoff scale of the power spectrum is in the range $[10^{-6} - 10^{2}\\, ]M_\\oplus$ \\cite{Chen:2001jz,Profumo:2006bv}. For axions, the growth of perturbations is suppressed on scales smaller than $(m_a H)^{-1/2}$ \\cite{Hu:2000ke,Boyle:2001du}, the geometric mean of the inverse of the axion mass and the Hubble constant; for example, for an axion of mass $m_a\\sim10^{-5}$~eV, the cutoff scale corresponds to about $10^{-11}\\,M_\\oplus$.  The implications of such substructure for searches for cosmic rays from WIMP annihilation in the halo have been explored extensively in the literature \\cite{subgamma}. Simply stated, if a dark matter halo contains substructure, then the volume integral of the density squared is increased by some boost factor $B$.  For example, Ref.~\\cite{Giocoli:2007gf} recently claimed $B\\sim10^3$ enhancements in the direction of the Galactic anticenter, although Ref.~\\cite{Berezinsky:2005py} finds an enhancement $B\\simeq2-5$. An analytic approach presented in Ref.~\\cite{Strigari:2006rd} also finds smaller values of $B$, of order $B \\sim 10$.  The implications of substructure for {\\it direct} detection of dark matter have, however, been comparatively neglected.  This is a serious omission, as substructure implies fluctuations in the local dark-matter density, and these fluctuations imply an uncertainty in the predicted rates for direct-detection experiments and for energetic-neutrino searches.  A simple but illustrative example places all of the dark matter in subhalos of density $B$ times the mean density.  In this case, the total annihilation rate is enhanced by a boost factor $B$, but the probability that the Solar System lives in such a subhalo is only $B^{-1}$.  If $B\\gg 1$, this probability is small, implying dire prospects for direct-detection, an alarming conclusion that warrants further investigation.  Few past studies focused on the implications of substructure for the local dark-matter density. Some focused on implications derived from numerical simulations \\cite{Moore:2001vq,Helmi:2002ss}, thus limited by resolution effects. Others looked at the contribution of tidal streams to the local density \\cite{Stiff:2001dq,Freese:2003tt}.  The goal of this paper is to begin addressing questions such as: What are the possible values of the local density?  How small can they be?  What are the most probable values? To answer these questions, we calculate a probability distribution function (PDF) $P(\\rho)$ for the local dark-matter density. This PDF accounts for fluctuations in the local density due to substructure. The mean of this distribution is $\\simeq0.4~\\GeV~\\cm^{-3}$, the canonical value determined from dynamics, but we find that the local dark-matter density may be as small as 1/10th, but probably no less than half this canonical value. The PDF we calculate will have to be convolved with the results of null searches to infer constraints to particle-dark-matter parameters. In addition, the PDF can be used to assess what a substructure enhancement in the halo WIMP annihilation rate implies for the local dark matter density, and vice-versa.  Our first approach, in Section II, uses an {\\it ansatz} about the survival fraction for subhalos that form early at high densities.  General arguments will be used to bracket the plausible range of PDFs, and also the plausible lower limit to the local dark-matter density.  This analytic approach capitalizes on the scale-invariant nature of the hierarchical formation of galaxies to extrapolate simulation results to the highest-density, smallest-scale substructures that form earliest and that are beyond the reach of simulations.  We then perform a second calculation that assumes substructures with NFW profiles are distributed in the Milky Way halo.  This second approach describes the local PDF due to larger-scale substructures which have formed from more recent stages in the merger hierarchy, and it thus complements the first approach.  In Section III, we improve upon this latter calculation by using ingredients on subhalo mass functions and concentration parameters taken from numerical simulations.  We compare our work to past literature in Section IV, and we state our conclusions in Section V. ",
        "conclusions": "It is evident from simulations and analytic arguments that some fraction of the local Milky Way dark matter may be in subhalos. The implications of substructure for indirect detection of WIMPs have been studied broadly, with the conclusion that there may be large enhancements in annihilation rates over the rates predicted assuming a smooth halo.  The implications for {\\it direct} searches have, however, been largely overlooked.  This is a possibly serious omission, as one consequence of substructure is that the local density will be smaller than the smoothly-distributed local density usually assumed.  We have taken a few first steps to understand the implications of substructure for the local density.  Our central goal is a calculation of the PDF $P(\\rho)$ for the local dark-matter density $\\rho$.  This $P(\\rho)$ will need to be taken into account when interpreting the implications of null dark-matter searches for constraints to the particle-dark-matter parameter space (e.g., couplings and/or elastic-scattering cross sections).  We considered two simple scenarios for substructure:  In the first, early generations of very dense subhalos survive with some probability.  The advantage of this approach is that if subhalo survival fractions can be measured in simulations for recently merged subhalos, the results might be extrapolated to the much earlier generations (much smaller and denser subhalos) that may be below the resolution of simulations.  In the second approach, the halo is assumed to consist of recently formed subhalos, each with an NFW profile.  This approach provides a simple, albeit approximate, way to understand the effects of larger-scale substructure, from more recent stages in the merger hierarchy, on the PDF. We then pursued this approach further using subhalo mass functions and concentration-parameter distributions taken from simulations.  This approach does not take into account the possible contribution of subhalos within subhalos.  In principle, a complete solution for the PDF, including substructures on the largest and smallest scales, can be obtained by convolving our two calculations.  This, however, will be left for future research.  Substructure scenarios that yield larger annihilation enhancements generally imply a smaller local dark-matter density.  Very large annihilation enhancements require that the very densest substructures, which generally form earlier, must survive through all later generations of structure formation.  If earlier substructures are less likely to survive than more recent substructures, then a very large annihilation enhancement is unlikely.  So, how small can the local density be?  The smallest local density in the models we surveyed was one tenth the canonical value of $0.4\\, \\GeV\\,\\cm^{-3}$, usually assumed for a smoothly distributed halo.  This small value was obtained from our simulation-inspired result using what we believe to be an overly conservative estimate of the smooth fraction, which we obtained by truncating NFW subhalos when their density falls below the mean halo density.  More realistic values for the smooth fraction are probably in the range of 50\\%-80\\%, which would then correspond in our simulation-inspired results to the $\\xi=0.8$ and $\\xi=0.5$ distributions, respectively, shown in Fig.~\\ref{fig:fig1}.  The smallest local density in our analytic model for early substructure was 0.3, obtained using the high value, $f(\\rhosun)=0.2$, and assuming a constant $f(\\rho_1)$.  If, however, $f(\\rhosun)$ is lower, then the local density is increased.  More importantly, it is quite likely that $f(\\rho)$ decreases with $\\rho$, and if so, the local density is increased to $\\sim0.8$ times its canonical value, even for $f(\\rhosun)=0.2$, and even higher for smaller $f(\\rhosun)$.  A combination of our two approaches, to take into account both low- and high-density substructures, will likely show that the local density is no less than half the canonical value.  The volume of the halo probed during a three-year direct-detection experiment is very small, and so the halo density $\\rho$ that we have been discussing can be safely assumed to be the density averaged over the duration of a direct-detection experiment. However, the rate for production of energetic neutrinos from WIMP annihilation in the Sun or Earth lags behind the rate for capture of WIMPs from the halo by an equilibration time $t_{\\rm eq}$ \\cite{Griest:1986yu,Kamionkowski:1991nj} that can vary considerably with the WIMP's mass and elastic-scattering and annihilation cross sections; typical values might be $t_{\\rm eq}^{\\odot}\\sim 5\\times 10^7$ yr and $t_{\\rm eq}^\\oplus\\sim10^{10}$ yr \\cite{Kamionkowski:1994dp}.  An energetic-neutrino search thus probes the halo density averaged over a much larger volume (for the latest bounds on the flux of high-energy neutrinos from annihilations in the Earth and the Sun see \\cite{Desai:2004pq,Ackermann:2005fr, Achterberg:2006jf}).The halo density averaged over this volume will thus have a PDF that will be narrower than that for direct detection, and the minimum density will be closer to the mean halo density. However, even though initially it was assumed that the velocity distribution function of dark matter particles mirrors that in free space (thus justifying the use of a Maxwell-Boltzmann distribution) \\cite{Gould:1991}, it is possible that the velocity distribution function is surpressed in the low-velocity tail due to the effects of solar capture and WIMP diffusion in the solar system due to the presence of other planets \\cite{Gould:1999je,Lundberg:2004dn}. This surpression manifests itself as a reduction in the number density of dark matter particles near the Earth for WIMPs with masses greater than few hundred GeV. Nevertheless, there is a possibility that for low-mass WIMPs, given the longer equilibration time for the Earth relative to that for the Sun, the annihilation signal from the Earth could be boosted relative to that for the Sun, if the Solar System passed through a very dense subhalo at a time $t_{\\rm eq}^\\odot <t<t_{\\rm eq}^\\oplus$ ago.  We leave a more detailed calculation of the density smoothed over timescales relevant for energetic-neutrino searches to future work \\cite{KK}.  The PDF $P(\\rho)$ allows us to evaluate also the boost factor in the annihilation rate.  It must be emphasized, though, that the boost factor $B$ we have considered is the boost only in the {\\it local} annihilation rate (per unit volume).  Strictly, speaking, the PDF may vary from one position in the halo to another. Thus, for example, if we assume that the halo is spherically symmetric (after averaging over substructure fluctuations), then the PDF will be a function $P(\\rho;r)$ of Galactocentric radius $r$, as well as the density $\\rho$.  The mean density of this distribution, as a function of $r$, will be the density of the smooth NFW profile that best fits the rotation curve. Qualitatively, we expect that the high-density tail will be more pronounced (more substructure) at larger $r$ and less pronounced (less substructure) at smaller $r$.  The PDF $P(\\rho;r)$ can then be used to determine the boost factor as a function of position in the halo; again, we expect $B(r)$ to vary with $r$ and to generally increase with $r$.  This $B(r)$ will be essential to compute annihilation fluxes along various lines of sight through the Galaxy.  We encourage future authors to describe the effects of substructure in terms of $P(\\rho;r)$, or at least in terms of $B(r)$, as this may facilitate comparison between the conclusions of different studies."
    },
    "0801/0801.4952_arXiv.txt": {
        "abstract": "The Alpha Magnetic Spectrometer (AMS) to be installed on the International Space Station (ISS) will measure charged cosmic ray spectra of elements up to iron, in the rigidity range from 1\\,GV to 1\\,TV, for at least three years. AMS is a large angular spectrometer composed of different subdetectors, including a proximity focusing Ring Imaging CHerenkov (RICH) detector. This will be equipped with a mixed radiator made of aerogel and sodium fluoride (NaF), a lateral conical mirror and a detection plane made of 680 photomultipliers coupled to light guides. The RICH detector allows measurements of particle's electric charge up to iron, and particle's velocity. Two possible methods for reconstructing the \\CK\\ angle and the electric charge with the RICH will be discussed.  A RICH prototype consisting of a detection matrix with 96 photomultipliers, a segment of a conical mirror and samples of the radiator materials was built and its performance was evaluated using ion beam data. Results from the last test beam performed with ion fragments resulting from the collision of a 158\\,GeV/c/nucleon primary beam of indium ions (CERN SPS) on a lead target are reported. The large amount of collected data allowed to test and characterize different aerogel samples and the NaF radiator. In addition, the reflectivity of the mirror was evaluated. The data analysis confirms the design goals. ",
        "introduction": "\\label{AMS-02} The Alpha Magnetic Spectrometer \\cite{bib:ams02-note} (AMS) is a particle detector to be installed in the International Space Station (ISS) for at least three years. The spectrometer  will be able to measure the rigidity ($R\\equiv pc/ |Z| e$), the charge ($Z$), the velocity ($\\beta$) and the energy ($E$) of cosmic rays from some MeV up to $\\sim$1\\,TeV within a geometrical acceptance of $ \\sim$0.5\\,m$^2$.sr. \\mbox{Figure \\ref{fig:ams}} shows a schematic view of the AMS spectrometer. At both ends of the AMS spectrometer exist the Transition Radiation Detector (TRD) (top) and the Electromagnetic Calorimeter (ECAL) (bottom). Both will provide AMS with capability to discriminate between leptons and hadrons. Additionally the calorimeter will trigger and detect photons. The TRD will be followed by the first of the four Time-of-Flight (TOF) system  scintillator planes. The TOF \\mbox{system \\cite{bib:ams02-tof}} is composed of four roughly circular planes of 12\\,cm wide scintillator paddles, one pair of planes above the magnet, the upper TOF, and one pair bellow, the lower TOF. There will be a total of 34 paddles. The TOF will provide a fast trigger within 200\\,ns, charge and velocity measurements for charged particles, as well as information on their direction of incidence. The TOF operation at regions with very intense magnetic fields forces the use of shielded fine mesh phototubes and the optimization of the light guides geometry, with some of them twisted and bent. Moreover the system guarantees redundancy, with two photomultipliers on each end of the paddles and double redundant electronics. A time resolution of 140\\,ps for protons is \\mbox{expected \\cite{bib:ams02-tof}}.  The tracking system will be surrounded by veto counters and embedded in a magnetic field of about 0.9\\,Tesla produced by a superconducting magnet. It will consist on a Silicon \\mbox{Tracker \\cite{bib:ams02-tracker}}, made of 8 layers of double sided silicon sensors with a total area of $\\sim$6.7\\,m$^2$. There will be a total of $\\sim$2500 silicon sensors arranged on 192 ladders. The position of the charged particles crossing the tracker layers is measured with a precision of $\\sim$10\\,$\\mu$m along the bending plane and $\\sim$30\\,$\\mu$m on the transverse direction. With a bending power (BL$^2$) of around 0.9\\,T.m$^2$, particles rigidity is measured with an accuracy better than 2\\% up to 20\\,GV and the maximal detectable rigidity is around 1-2\\,TV. Electric charge is also measured from energy deposition up to Z$\\sim$26.  The Ring Imaging \\CK\\ Detector (RICH) \\cite{bib:baraoICRC} will be located right after the last TOF plane and before the electromagnetic calorimeter. It will be described in detail in next section.  \\begin{figure}[htb] \\begin{center} \\vspace{-0.45cm} \\epsfig{file=AMS2.DET.3D.May03-newII.ps,width=0.65\\linewidth,clip=,bbllx=0,bblly=0,bburx=612,bbury=550} \\caption{A whole view of the AMS Spectrometer.} \\end{center} \\label{fig:ams} \\end{figure}  The long stay of AMS in space and its large acceptance will allow the accumulation of a large statistic of events increasing in several orders of magnitude the sensitivity of the proposed physical measurements. With an average collection rate of 1000 events per second, a total of $10^9$ protons per year and around $10^4$ antiprotons will be accumulated.  The main goals of the AMS-02 experiment are:  \\begin{itemize}  \\item A precise measurement of the charged cosmic ray spectrum between \\mbox{$\\sim$100\\,MeV} and $\\sim$1\\,TeV, and the detection of photons up to a few hundred GeV;  \\item A search for heavy antinuclei ($Z \\ge 2$), which if discovered would signal the existence of primordial antimatter;  \\item Search for non-baryonic dark matter through the detection of annihilation products appearing as anomalies of the cosmic-ray spectra (e$^+$, $\\bar{p}$, $\\gamma$ and $\\bar{d}$);  \\end{itemize} ",
        "conclusions": "AMS-02 will be equipped with a proximity focusing RICH detector based on a mixed radiator of aerogel and sodium fluoride, enabling velocity measurements with a resolution of about 0.1\\% and extending the charge measurements up to the iron element. Velocity reconstruction is made with a likelihood method. Charge reconstruction is made in an event-by-event basis. Evaluation of both algorithms on real data taken with in-beam tests at CERN, in October 2003 was done. The detector design was validated and a refractive index 1.05 aerogel was chosen for the radiator, fulfilling both the demand for a large light yield and a good velocity resolution. The RICH detetector is being constructed and its assembling to the AMS complete setup is foreseen for 2008."
    },
    "0801/0801.4213_arXiv.txt": {
        "abstract": "We review the impact of global helioseismology on key questions concerning the internal structure and dynamics of the Sun, and consider the exciting challenges the field faces as it enters a fourth decade of science exploitation. We do so with an eye on the past, looking at the perspectives global helioseismology offered in its earlier phases, in particular the mid-to-late 1970s and the 1980s.  We look at how modern, higher-quality, longer datasets coupled with new developments in analysis, have altered, refined, and changed some of those perspectives, and opened others that were not previously available for study. We finish by discussing outstanding challenges and questions for the field. ",
        "introduction": "\\label{sec:intro}  The field of global helioseismology -- the use of accurate and precise observations of the globally coherent modes of oscillation of the Sun to make inference on the internal structure and dynamics of our star -- is about to enter its fourth decade. The observational starting point for global helioseismology was marked in the mid-to-late 1970s by several key papers: First, the observational confirmation by Deubner (1975), and independently by Rhodes, Ulrich and Simon (1977), of the standing-wave nature of the five-minute oscillations observed on the surface of the Sun, which was proposed by Ulrich (1970), and Leibacher and Stein (1971); and then the discovery that the oscillations displayed by the Sun were truly global whole-Sun, core-penetrating, radial-mode pulsations (Claverie \\etal, 1979). Previous observations of pulsating stars had revealed many objects that were oscillating in one, or at most a few, modes. The rich spectrum of oscillations displayed by the Sun was a different matter entirely. Christensen-Dalsgaard and Gough (1976) pointed out the great potential that a multi-modal spectrum could offer: the information content of the observations could potentially be so great as to allow a reconstruction of the internal structure of the star. More than 30 years on, exquisite observational reconstructions of the internal structure and dynamics of the Sun are in everyday use, thanks to helioseismology.  Several thousand modes of oscillation of the Sun have to date been observed, identified, and studied. The oscillations are standing acoustic waves, for which the gradient of pressure ($p$) is the principal restoring force. The modes are excited stochastically, and damped intrinsically, by the turbulence in the outermost layers of the sub-surface convection zone. The stochastic excitation mechanism limits the amplitudes of the $p$ modes to intrinsically weak values. However, it gives rise to an extremely rich spectrum of modes, the most prominent generally being high-order overtones. Detection of individual gravity modes -- for which buoyancy acts as the principal restoring force -- remains an important goal for the field. The $g$ modes are confined in cavities in the radiative interior, and if observed would provide a much more sensitive probe of conditions in the core than the $p$ modes.  The small-amplitude solar oscillations may be described in terms of spherical harmonics $Y^{m}_{\\ell}(\\theta,\\phi)$, where $\\theta$ and $\\phi$ are the co-latitude and longitude respectively: \\begin{equation} Y^{m}_{\\ell}(\\theta,\\phi) = (-1)^m c_{\\ell m} P^{m}_{\\ell}(\\cos \\theta) \\exp({\\rm i} m\\phi). \\label{eq:ylm} \\end{equation} In the above, the $P^m_\\ell$ is a Legendre function, and the $c_{\\ell m}$ is a normalization constant. The $p$ modes probe different interior volumes, with the radial and other low-degree (low-$\\ell$) modes probing as deeply as the core. This differential penetration of the modes allows the internal structure and dynamics to be inferred, as a function of position, to high levels of precision not usually encountered in astrophysics. The Sun has not surprisingly been the exemplar for the development of seismic methods for probing stellar interiors.  Extension of the observations to other Sun-like stars (asteroseismology) has demonstrated that the Sun-like oscillations are a ubiquitous feature of stars with sub-surface convection zones.  In this review our aim is to look at some of the recent advances that global helioseismology has made for studies of various aspects of the internal structure of the Sun. We discuss the observational challenge posed by the detection and identification of individual $g$ modes. We also seek to provide in each section some historical context for discussion of the contemporary results and challenges. We round out the review by looking to the future, in particular the need for continued multi-instrument, multi-network observations of the global modes; and we finish by listing some important questions and challenges for global helioseismology. ",
        "conclusions": ""
    },
    "0801/0801.3928_arXiv.txt": {
        "abstract": "By using  most of the present Cosmic Microwave Background (CMB) and  Large Scale Structure (LSS) measurements and the Big Bang Nucleosynthesis (BBN) constraints on the primordial helium abundance, $Y_p$, we set bounds on the radiation content of the Universe and neutrino properties. We consider lepton asymmetric cosmological models parametrized by the neutrino degeneracy parameter $\\xi_{\\nu}$ and the variation of the relativistic degrees of freedom, $\\Delta  N_{oth}^{eff}$, due to possible other physical processes that occurred between BBN and structure formation epoch.\\\\ We found that present CMB and LSS data constraints the neutrino degeneracy parameter at $\\xi_{\\nu} \\leq 0.722$, implying a lepton asymmetry of the neutrino background $ {\\cal L}_{\\nu} \\leq 0.614$ (2-$\\sigma$). We also found  $\\Delta N^{eff}_{oth}=0.572 ^{+1.972}_{-1.780}$ , the contribution to the effective number of relativistic neutrino species $ N^{eff}=3.058^{+1.971}_{-1.178}$ and a primordial helium abundance $Y_p=0.249^{+0.014}_{-0.016}$ (2-$\\sigma$ errors).\\\\ These results bring an important improvement over the similar ones obtained by using WMAP~1-year and older LSS data or the WMAP~3-year data alone and the standard primordial helium abundance value $Y_p=0.24$, relaxing the stringent BBN constraint on the neutrino degeneracy parameter ($\\xi_{\\nu} \\leq 0.07$).  We forecast that the CMB temperature and polarization maps observed with high angular resolutions and sensitivity by the future {\\sc Planck} Mission will constraint the primordial  primordial helium abundance at $Y_p=0.247 \\pm 0.002$ (2-$\\sigma$ errors) in agreement with the most stringent limits on $Y_p$ given by the BBN  and the neutrino degeneracy parameter at $\\xi_{\\nu} \\leq 0.280$ (2-$\\sigma$), not excluding the possibility of larger lepton asymmetry.\\\\   This work has been done on behalf of {\\sc Planck}-LFI activities. ",
        "introduction": "The radiation budget of the Universe relies on a strong theoretical prejudice: apart from the Cosmic Microwave Background (CMB) photons, the relativistic background would consist of neutrinos and of possible contributions from other relativistic relicts. The main constraints on the radiation energy density come either from the very early Universe, where the radiation was the dominant source of energy, or from the observation of cosmological perturbations which carry the information about the time equality between matter and radiation.  \\vspace{0.3cm} In particular, the primordial light element abundance predictions in the standard theory of the Big Bang Nucleosynthesis (BBN) \\cite{Wagoner,Olive,Burles,Eidelman} depend on the baryon-to-photon ratio, $\\eta_B$, and the radiation energy density at the BBN epoch (energy density of order MeV$^4$), usually parametrized by the effective number of relativistic neutrino species, $N^{eff}$.  Meanwhile, the number of active neutrino flavors have been fixed by  $Z^0$ boson decay width to $N_{\\nu}=2.944 \\pm 0.012$ \\cite{Eidelman} and the combined study of the incomplete neutrino decoupling and the QED corrections indicate that the number of relativistic neutrino species is $N^{eff}_{\\nu}=3.046$ \\cite{Mangano02}.  Any departure of  $N^{eff}$ from this last value would be due to non-standard neutrino features or to the contribution of other relativistic relics. \\\\ The solar and atmospheric neutrino oscillation experiments  \\cite{Fuk98,Amb98} indicate the existence of non-zero neutrino masses in eV range.  There are also indications of neutrino oscillations with  larger mass-squared difference, coming from the short base-line oscillation experiments \\cite{Atha,Miniboone}, that can be explained by adding one or two sterile neutrinos with eV-scale mass to the standard scheme with three active neutrino flavors (see Ref.\\cite{Maltoni} for a recent analysis). Such results have impact on cosmology because sterile neutrinos can contribute to the number of relativistic degrees of freedom at the Big Bang Nucleosynthesis \\cite{Cirelli}. These models are subject to strong bounds on the sum of active neutrino masses from the combination of various cosmological data sets \\cite{Rafelt,Find}, ruling out a thermalized sterile neutrino component with eV mass \\cite{Seljak07,Mel}.\\\\ However, there is the possibility to accommodate the cosmological observations with data from short base-line neutrino oscillation experiments by postulating the existence of a sterile neutrino with the mass of few keV having a phase-space distribution significantly suppressed relative to the thermal distribution.\\\\ For both, non-resonant zero lepton number production and enhanced resonant production with initial cosmological lepton number, keV sterile neutrinos are produced via small mixing angle oscillation conversion of thermal active neutrinos \\cite{Aba02}.\\\\ Sterile neutrino with mass of few keV provides also a valuable Dark Matter (DM) candidate \\cite{DodWidr,Dol,Aba01,Sha06}, alleviating the accumulating contradiction between the $\\Lambda$CDM model predictions on small scales and observations, by smearing out the small scale structure.\\\\ On the other hand, the possible existence of new particles such as axions and gravitons, the time variation of the physical constants and  other non-standard scenarios (see e.g. \\cite{Sarkar96} and references therein) could contribute to the radiation energy density at BBN epoch.  At the same time, more phenomenological extensions to the standard neutrino sector have been studied, the most natural being consideration of the leptonic asymmetry \\cite{Freese83,Ruffini83,Ruffini88}, parametrized by the neutrino degeneracy parameter $\\xi_{\\nu}=\\mu_{\\nu}/T_{\\nu_0}$ [$\\mu_{\\nu}$ is the neutrino chemical potential and $T_{\\nu_0}$  is the present temperature of the neutrino background, $T_{\\nu_0}/T_{\\rm cmb}=(4/11)^{1/3}$]. \\\\ Although  the standard model predicts the leptonic asymmetry  of the same order as the  baryonic asymmetry, $B \\sim 10^{-10}$, there are many particle physics scenario in which a leptonic asymmetry much larger can be generated \\cite{Smith06,Cirelli06}. One of the cosmological implications of a larger leptonic asymmetry is the possibility to generate small baryonic asymmetry of the Universe through the non-perturbative (sphaleron) processes \\cite{Kuzmin85,Falcone01,Buch04}. Therefore, distinguishing between a vanishing and non-vanishing $\\xi_{\\nu}$ at the BBN epoch is a crucial test of the standard assumption that sphaleron effects equilibrate the cosmic lepton and baryon asymmetries.\\\\ The measured neutrino mixing parameters implies that neutrinos reach the chemical equilibrium before BBN \\cite{Dolgov02,Wong02,Beacom02} so that  all neutrino flavors are characterized by the same degeneracy parameter, $\\xi_{\\nu}$, at this epoch.\\\\ The most important impact of the leptonic asymmetry on BBN is the shift of the beta equilibrium between protons and neutrons and the increase of the radiation energy density parametrized by: \\begin{equation} \\Delta N^{ eff}(\\xi_{\\nu})=3\\left[ \\frac{30}{7} \\left(\\frac{\\xi_{\\nu}}{\\pi} \\right)^2 +\\frac{15}{7}\\left( \\frac{\\xi_{\\nu}}{\\pi} \\right)^4 \\right] \\,. \\end{equation} The BBN constraints on N$^{eff}$ have been recently reanalyzed by comparing  the theoretical predictions and experimental data on the primordial abundances of light elements, using the baryon abundance derived from  the WMAP~3-year (WMAP3) CMB temperature and  polarization measurements \\cite{Spergel07,Nolta,Page}: $\\eta_B=6.14 \\times 10^{-10}(1.00\\pm 0.04)$. In particular, the $^4$He abundance, $Y_p$,  is quite sensitive to the value of $N^{eff}$.\\\\ The analysis of Ref. \\cite{Mangano07}, the conservative error analysis of  helium abundance, $Y_P=0.249 \\pm 0.009$ \\cite{Olive04}, yielded to $N^{eff}=3.1^{+1.4}_{-1.2}$ (2-$\\sigma$) in good agreement with the standard value, but still leaving some room for non-standard values, while more stringent error bars of helium abundance , $Y_p=0.2516 \\pm 0.0011$ \\cite{Izotov07}, leaded to $N^{eff}=3.32^{+0.23}_{-0.24}$ (2-$\\sigma$) \\cite{Ichikawa07}.  The stronger constraints on the degeneracy parameter obtained from BBN  \\cite{Serpico05} gives $-0.04 < \\xi< 0.07$ (1-$\\sigma$), adopting  the conservative error analysis of $Y_p$ of Ref. \\cite{Olive04} and $\\xi=0.024 \\pm 0.0092$ (1-$\\sigma$), adopting the more stringent error bars of $Y_p$ of Ref. \\cite{Izotov03}.   \\vspace{0.3cm} The CMB anisotropies and LSS matter density fluctuations power spectra carry the signature of the energy density of the Universe at the time of matter-radiation equality (energy density of order eV$^4$), making possible the measurement of $N^{eff}$ through its effects on the growth of cosmological perturbations. \\\\ More effective number of relativistic neutrino species enhances the integrated Sachs-Wolfe effect on the CMB power spectrum, leading to a higher first acoustic Doppler peak amplitude. Also, the delay of the epoch of matter-radiation equality shifts the LSS matter power spectrum turnover position toward larger angular scales, suppressing the power at small scales. In particular, for the leptonic asymmetric models, the neutrino mass is lighter than in the symmetric case. This leads to changes in neutrino free-streaming length and neutrino Jeans mass due to the increase of the neutrino velocity dispersion \\cite{Lattanzi05,Ichiki07}.  After WMAP3 data release, there are many works aiming to constrain $N^{eff}$ from cosmological observations \\cite{Seljak07,Spergel07,Mangano07,Han06,Cirelli06b,Kawa07}. Their results suggest large values for $N^{eff}$ within 2-$\\sigma$ interval, some of them  not including the standard value $N^{eff}_{\\nu}=3.046$ \\cite{Seljak07,Spergel07,Mangano07}. Recently Ref. \\cite{Han07} argues that the discrepancies are due to the treatment of the  scale-dependent biasing in the galaxy power spectrum inferred from the main galaxy sample of the Sloan Digital Sky Survey data release 2 (SDSS-DR2)  \\cite{Teg04a,Teg04b} and the large fluctuation amplitude reconstructed from  the Lyman$-\\alpha$ forest data \\cite{McDonald05} relative to that inferred from WMAP3. \\\\ Discrepancies between  BBN  and cosmological data results on $N^{eff}$ was interpreted as  2-$\\sigma$ evidence of the fact that further relativistic species are produced by particles decay between BBN and structure formation \\cite{Cirelli06b,Kawa07}. Other theoretical scenarios include the violation of the spin-statistics in the neutrino sector \\cite{Dolgov05}, the possibility of an extra interaction between the dark energy and radiation or dark matter, the existence of a Brans-Dicke field which could mimic the effect of adding extra relativistic energy density between BBN and structure formation epochs \\cite{deFelice}.  \\vspace{0.3cm} The extra energy density can be split in two distinct uncorrelated contributions, first due to net lepton asymmetry of the neutrino background and second due to the extra contributions from other unknown processes: \\begin{equation} \\Delta N^{eff}=\\Delta N^{eff}(\\xi)+ \\Delta N_{oth}^{eff}\\, . \\end{equation} The aim of this paper is to obtain bounds on the neutrino lepton asymmetry and on the extra radiation energy density by using most of the existing CMB and LSS measurements and self-consistent BBN priors on $Y_p$. We also to compute the sensitivity of the future {\\sc Planck} experiment \\cite{Blue} for these parameters testing the restrictions on cosmological models with extra relativistic degrees of freedom expected from high precision CMB temperature and polarization anisotropy measurements.     \\\\  The paper is organized as follows: Section~2 contains a review on the lepton asymmetric cosmological models and the BBN theory, Section~3 is devoted to a summary data analysis method while in Section~4 we discuss our results. We draw our main conclusions in Section~5 ",
        "conclusions": "In this paper, we set bounds on the radiation content of the Universe and neutrino properties by using most of the present Cosmic Microwave Background (CMB) and  Large Scale Structure (LSS) measurements. We also take into account the Big Bang Nucleosynthesis (BBN) constraints on the primordial helium abundance, $Y_p$, which prove to be important in the estimation of cosmological parameters from future {\\sc Planck} data, both in non-degenerate and degenerate BBN models; the importance of the self-consistent BBN prior on $Y_p$ was also emphasized in two other recent analysis \\cite{Hamann07, Ichikawa071}.\\\\ We consider lepton asymmetric cosmological models parametrized by the neutrino degeneracy parameter $\\xi_{\\nu}$ and the variation of the relativistic degrees of freedom, $\\Delta  N_{oth}^{eff}$, due to possible other physical processes that occurred between BBN and structure formation epoch.\\\\ We found that present CMB and LSS data together with BBN prior on the primordial helium abundance ($Y_p$) constraints the neutrino degeneracy parameter at $\\xi_{\\nu} \\leq 0.722$, leading to a lepton asymmetric neutrino background of $ {\\cal L}_{\\nu} \\leq 0.614$ (2-$\\sigma$). We also found  $\\Delta N^{eff}_{oth}=0.572 ^{+1.972}_{-1.780}$ , the contribution to the effective number of re-\\\\lativistic neutrino species $ N^{eff}=3.058^{+1.971}_{-1.178}$ and a primordial helium abundance $Y_p=0.249^{+0.014}_{-0.016}$ (2-$\\sigma$ errors). These values represent important improvements over the similar results obtained by using WMAP~1-year together with older LSS data \\cite{Crotty03} or the WMAP3 data alone \\cite{Lattanzi05} and the standard primordial helium abundance value $Y_p=0.24$, relaxing the stringent BBN constraint on the neutrino degeneracy parameter ($\\xi_{\\nu} \\leq 0.07$).  We observe that, when using WMAP3+All data, an additional number of relativistic relicts brings a substantial degeneracy in the $\\Omega_{m} - \\Omega_{r}$ plane and a weaker constraint on the age of the Universe; the same degeneracy occurs also between other cosmological parameters under the same conditions. We therefore conclude that the present cosmological data do not favor the variation of the relativistic degrees of freedom, $\\Delta N_{oth}^{eff}$, due to other possible physical processes that occurred between BBN and matter-radiation equality epoch.\\\\  We forecast that the CMB temperature and polarization maps observed with high angular resolutions and sensitivity by the future {\\sc Planck} Mission will constraint the primordial  primordial helium abundance at $Y_p=0.247 \\pm 0.002$ (2-$\\sigma$ errors) in agreement with the most stringent limits on $Y_p$ given by the BBN  and the neutrino degeneracy parameter at $\\xi_{\\nu} \\leq 0.280$ (2-$\\sigma$), allowing larger lepton asymmetry models. Also, they will reduce the degeneracy between $|\\xi|$ and $\\Delta N_{oth}^{eff}$ allowing a better distinction between extra radiation energy density coming from an additional number of relativistic relicts and from a lepton asymmetric neutrino background.\\\\    {\\bf Acknowledgements} L.P. and A.V. acknowledge the support by the ESA/PECS Contract \"Scientific exploitation of  Planck-LFI data\"  We also acknowledge the use of the GRID computing system facility at the Institute for Space Sciences Bucharest and would like to thank the staff working there."
    },
    "0801/0801.4814_arXiv.txt": {
        "abstract": "Infrared observations of hot massive stars and their environments provide a detailed picture of mass loss histories, dust formation, and dynamical interactions with the local stellar medium that can be unique to the thermal regime.  We have acquired new infrared spectroscopy and imaging with the sensitive instruments onboard the {\\em{Spitzer}} Space Telescope in guaranteed and open time programs comprised of some of the best known examples of hot stars with circumstellar nebulae, supplementing with unpublished Infrared Space Observatory spectroscopy.  Here we present highlights of our work on the environment around the extreme P Cygni-type star HDE316285, providing some defining characteristics of the star's evolution and interactions with the ISM at unprecented detail in the infrared. ",
        "introduction": "Observations of the circumstellar environments of Lumnious Blue Variables (LBVs) with ESA's Infrared Space Observatory (ISO) and modern ground-based imaging devices have demonstrated how sensitive mid-infrared imaging and spectroscopy can substantially improve our knowledge of the distribution of material, the physical and chemical properties of nebular dust and gas, and therefore on the history of mass loss from the surface of the massive central star during its evolution from the Main Sequence.  Other than the bright and compact environments of most LBV nebulae, none of the lower surface brightness nebulae around OB or Wolf-Rayet stars could be observed with ISO due mainly to instrumental sensitivity limitations. The instruments on the {\\em{Spitzer}} Space Telescope have since offered significant gains in sensitivity and sky coverage; for comparison, the {\\em{Spitzer}} Infrared Spectrometer (IRS; 5.3 -- 38 $\\mu$m) is factor of $\\sim$50-100 more sensitive than the ISO Short Wavelength Spectrograph (SWS; 2.4-45.2 $\\mu$m) at 5 $\\mu$m, depending on resolution modes of the two instruments. We have exploited the {\\em{Spitzer}} IRS,  the Infrared Array Camera (IRAC) with imaging bands at 3.4, 4.6, 5.8, and 8.0 $\\mu$m, and the Multiband Imaging Photometer for {\\em{Spitzer}} (MIPS) with imaging capabilities at 24, 70, and 160 $\\mu$m in Guaranteed and Open Time programs to observe the environments around a number of hot, massive stars.  Most of these stars are surrounded by ring nebulae and are well known for their optical and/or radio properties; NGC2359 (WR), M1-67 (WR), G79.29+00.46 (BIe), and HD148937 (Ofp?e) are examples.  In this paper we summarize our study of the massive B supergiant and candidate LBV HDE316285 from our program.   This highly luminous P Cyg-type star has been suspected of being surrounded by an extended, cold nebula (\\cite[McGregor, Hyland, \\& Hillier 1998]{McGregorHylandHillier_98}); we have obtained conclusive observations of the environment around this star, which harbor telling characteristics of the star's mass loss history, wind-wind and wind-ISM interactions, the molecular content of the gas and thermal history of the dust.  \\begin{figure}[b] \\begin{center} \\includegraphics[width=3.4in]{morrisfig1.eps} \\caption{Combined near-IR (1.4-2.4$\\mu$m), ISO/SWS (2.4-10$\\mu$m), and {\\em{Spitzer}}/IRS (10.0-38$\\mu$m) SEDs of HDE316285 (upper) and G79.29+00.46 (lower). From Morris et al. (2008) submitted.} \\label{fig1} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0801/0801.3433_arXiv.txt": {
        "abstract": "We consider evolutionary models for the population of short-period ($\\per \\lta 10$\\,hr) low-mass black-hole binaries (LMBHB) and compare them with observations of soft X-ray transients (SXT). Evolution of LMBHB is determined by nuclear evolution of the donors and/or orbital angular momentum loss due to magnetic braking by the stellar wind of the donors and gravitational wave radiation. We show that the absence of observed stable luminous LMBHB implies that upon RLOF by the low-mass donor angular momentum losses are substantially reduced with respect to the Verbunt \\& Zwaan ``standard'' prescription for magnetic braking. Under this assumption masses and effective temperatures of the model secondaries of LMBHB are in a satisfactory agreement with the masses and effective temperatures (as inferred from their spectra) of the observed donors in LMBHB. Theoretical mass-transfer rates in SXTs are consistent with the observed ones if one assumes that accretion discs in these systems are truncated (``leaky''). We find that the population of short-period SXT is formed mainly by systems which had unevolved or slightly evolved ($X_c \\gta 0.35$) donors at the Roche-lobe overflow. Longer period ($\\per \\simeq (0.5 - 1)$\\ day) SXT might descend from systems with initial donor mass about 1\\,\\ms\\ and $X_c \\lta 0.35$. It is unnecessary to invoke donors with almost hydrogen-depleted cores to explain the origin of LMBHB. Our models suggest that a very high efficiency of common-envelopes ejection is necessary to form LMBHB, unless currently commonly accepted empirical estimates of mass-loss rates by winds for pre-WR and WR-stars are significantly over-evaluated. ",
        "introduction": "\\label{sec:intro} Out of twenty dynamically-confirmed black-hole candidate X-ray binaries ten have K/M spectral type secondaries and orbital periods $\\aplt$1 day \\citep{2006ARA&A..44...49R}. All these systems are transient in X-ray. Their quiescent X-ray luminosity is $\\simeq (3\\,10^{30} -- 3\\,10^{33})$\\,\\es, while in the outburst luminosity may be by a factor $\\sim 10^5 - 10^8$ higher. These low-mass black-hole binary (LMBHB) X-ray systems belong to the class of ``Soft X-ray transients'' (SXT). Recent reviews of observational data on SXT may be found e.g. in \\citet{2006csxs.book..215C,2006ARA&A..44...49R}. From the evolutionary stand-point SXT are considered as semidetached binaries in which matter is transferred to the compact object via accretion disc. The commonly accepted model for the variability of SXT is based on the thermal-viscous instability of irradiated accretion discs \\citep[see ][and references therein]{lasota01}. The estimated number of SXT in the Galaxy ranges from several hundred \\citep{csl97} to several thousand \\citep{romani1998}. \\begin{figure}[ht!] \\includegraphics[angle=-90,width=0.48\\textwidth]{imfm.ps} \\caption[]{Initial -- final mass relations for massive stars in binaries. The mass of stars at terminal main-sequence (TAMS), the initial mass of helium core ${\\rm M}_{\\rm core, i}$, the pre-supernova mass ${\\rm M}_{\\rm pre-SN}$, and the black-hole mass ${\\rm M}_{\\rm BH}$ are shown.} \\label{fig:imfm} \\end{figure} Below, we will describe the models of a population of short-period LMBHB, discuss their observational properties, and consider possible existence of a population of faint LMBHB with quasi-stable, cold accretion discs \\citep{mnl99}. This review is based mainly on the studies by \\citet{yungelson_bh06} (henceforth, Paper I) and Yungelson \\& Lasota (in prep., Paper II). ",
        "conclusions": "\\label{sec:concl}  1. We have modeled populations of short-period semidetached close binary systems containing black holes and low-mass stars under various assumptions about the AML mechanism.  We have found that the \"standard\" model (based on \\citet{sku72} stellar rotation braking law prescription for AML via magnetic braking) produces unobserved persistent LMBHB, while a model with pure GWR AML produces only transient systems. In the latter model it is possible to reproduce, within uncertainty of observations, the number of LMBHB in the Galaxy, the effective temperatures and masses of the donors in these systems (as inferred from the spectra of the latter). Mass-transfer rates in this model are consistent with the upper limits on \\md, if quiescent accretion discs are truncated and, hence, leaky. All this suggests that the strength of MSW AML should be strongly reduced compared to the ``standard'' model.  In the model without strong MSW it is unnecessary to invoke systems with the donors that are almost at TAMS or already left the latter ($X_c \\lesssim 0.01$) in order to explain observed orbital periods of SXT and their mass-exchange rates. More, population of short-period ($\\per \\aplt 8-9$ hr) is formed by systems that overflow Roche lobes unevolved or slightly evolved ($X_c > 0.35$), while evolution of systems with more evolved donors with masses $\\simeq 1\\,\\ms$ opens possibility to explain the origin of LMBHB with longer \\per.  2. An alternative model suggests that if accretion discs are truncated in quiescence they could be cold and stable. Then truncation of discs at radii close to $R_{\\rm circ}$ can make discs in virtually all LMBHB cold and stable. In this case SXT outbursts perhaps are not related to DIM and are random events that result from \"external factors\", like episodes of enhanced mass transfer or variability in truncation radius.  3. Population synthesis models presented above are obtained under assumption of a high value of the product of binding energy parameter of stellar envelopes and expulsion efficiency of common envelopes -- \\al=2. Model with \\al=0.5 gives similar results (but a reduced total number of systems). Further decrease of \\al\\ results in models that are not consistent with observed SXT. The issue of \\al\\ is still controversial \\citep[see, e.g. ][]{Taudew,Podsiadlowski&al03} in the sense that there are no strict criteria for defining binding energy of stellar envelopes and there is no clear understanding whether sources other than gravitational energy may contribute to unbinding common envelopes. Apart of common envelopes, crucial role in determining the possibility of formation of LMBHB is played by stellar wind mass-loss in all stages of stellar evolution, since it defines parameters of the system at the beginning of common envelope stage and immediately before supernova explosion that gives birth to black hole. If general notions on mass-loss will not be revised, our results suggest that existence of short-period low-mass X-ray binaries demands high values of \\al\\ parameter."
    },
    "0801/0801.3427_arXiv.txt": {
        "abstract": "Several $\\gamma$-ray binaries have been recently detected by the High-Energy Stereoscopy Array (H.E.S.S.) and the Major Atmospheric Imaging Cerenkov (MAGIC) telescope. In at least two cases, their nature is unknown. In this paper we aim to provide the details of a theoretical model of close $\\gamma$-ray binaries containing a young energetic pulsar as compact object, earlier presented in recent Letters. This model includes a detailed account of the system geometry, the angular dependence of processes such as Klein-Nishina inverse Compton and $\\gamma\\gamma$ absorption in the anisotropic radiation field of the massive star, and a Monte Carlo simulation of leptonic cascading. We present and derive the used formulae and give all details about their numerical implementation, particularly, on the computation of cascades.  In this model, emphasis is put in the processes occurring in the pulsar wind zone of the binary, since, as we show, opacities in this region can be already important for close systems. We provide a detailed study on all relevant opacities and geometrical dependencies along the orbit of binaries, exemplifying with the case of LS 5039. This is used to understand the formation of the very high-energy lightcurve and phase dependent spectrum. For the particular case of LS 5039, we uncover an interesting behavior of the magnitude representing the shock position in the direction to  the observer along the orbit, and analyze its impact in the predictions. We show that in the case of LS 5039, the H.E.S.S. phenomenology is matched by the presented model, and explore the reasons why this happens while discussing future ways of testing the model. ",
        "introduction": "Very recently, a few massive binaries have been identified as variable very-high-energy (VHE) $\\gamma$-ray sources. They are PSR B1259-63 (Aharonian et al. 2005a), LS 5039 (Aharonian et al. 2005b, 2006), LS I +61 303 (Albert et al. 2006, 2008a,b), and Cyg X-1 (Albert et al. 2007). The nature of only two of these binaries is considered known: PSR B 1259-63 is formed with a pulsar whereas Cyg X-1 is formed with a black hole compact object. The nature of the two remaining systems is under discussion. The high-energy phenomenology of Cyg X-1 is different from that of the others. It has been detected just once in a flare state for which a duty cycle is yet unknown. The three other sources, instead, present a behavior that is fully correlated with the orbital period. The latter varies from about 4 days in the case of LS 5039 to several years in the case of PSR B1259-63: this span of orbital periodicities introduces its own complications in analyzing the similarities among the three  systems.   LS I +61 303 shares with LS 5039 the quality of being the only two known microquasars/$\\gamma$-ray binaries that are spatially coincident with sources above 100 MeV listed in the Third Energetic Gamma-Ray Experiment (EGRET) catalog (Hartman et al. 1999). These sources both show low X-ray emission and variability, and no signs of emission lines or disk accretion. For LS I +61 303, extended, apparently precessing, radio emitting structures at angular extensions of 0.01-0.05 arcsec have been reported by Massi et al. (2001, 2004); this discovery has earlier supported its microquasar interpretation. But the uncertainty as to what kind of compact object, a black hole or a neutron star, is part of the system (e.g., Casares et al. 2005a), seems settled for many after the results presented by Dhawan et al. (2006). These authors have presented observations from a July 2006 VLBI campaign in which rapid changes are seen in the orientation of what seems to be a cometary tail at periastron. This tail is consistent with it being the result of a pulsar wind. Indeed, no large features or high-velocity flows were noted on any of the observing days, which implies at least its non-permanent nature. The changes within 3 hours were found to be insignificant, so the velocity can not be much over 0.05$c$. Still, discussion is on-going (e.g. see Romero et al. 2007, Zdziarski et al. 2008). New campaigns with similar radio resolution, as well as new observations in the $\\gamma$-ray domain have been obtained since the Dhawan's et al. original results (Albert et al. 2008b). A key aspect  in these high-angular-resolution campaigns is the observed maintenance in time of the morphology of the radio emission of the system: the changing morphology of the radio emission along the orbit would require a highly unstable jet, which details are not expected to be reproduced orbit after orbit as indicated by current results (Albert et al. 2008b). The absence of accretion signatures in X-rays in Chandra and XMM-Newton observations (as reported by Sidoli et al. 2006, Chernyakova et al. 2006, and Paredes et al. 2007) is another relevant aspect of the discussion about the compact object companion. Finally, it is interesting to note that neutrino detection or non-detection with ICECUBE will shed light on the nature of the $\\gamma$-ray emission irrespective of the system composition (e.g.,  Aharonian et al. 2006b, Torres and Halzen 2007).   For LS 5039, a periodicity in the $\\gamma$-ray flux, consistent with the orbital timescale as determined by Casares et al. (2005b), was found with amazing precision (Aharonian et al. 2006). Short timescale variability displayed  on top of this periodic behavior, both in flux and spectrum, was also reported. It was found that the parameters of power-law fits to the $\\gamma$-ray data obtained in 0.1 phase binning already displayed significant variability. Current H.E.S.S. observations of LS 5039  ($\\sim 70$ hours distributed over many orbital cycles, Aharonian et al. 2006) constitute one of the most detailed datasets of high-energy astrophysics. Similarly to LS I +61 303, the discovery of a jet-like radio structure in LS 5039 and the fact of it being the only radio/X-ray source co-localized with a mildly variable (Torres et al. 2001a,b) EGRET detection, prompted a microquasar interpretation (advanced already by Paredes et al. 2000). However, the current mentioned findings at radio and VHE $\\gamma$-rays in the cases of LS I +61 303 (Dhawan et al. 2006, Albert et al. 2006, 2008b) or PSR B1259-63 (Aharonian et al. 2005), gave the perspective that all three systems are different realizations of the same scenario: a pulsar-massive star binary. Dubus (2006a,b) has studied these similarities. He provided simulations of the extended radio emission of LS 5039 showing that the features found in high resolution radio observations could also be interpreted as the result of a pulsar wind. Recently, Rib\\'o et al. (2008) provided VLBA radio observations of LS 5039 with morphological and astrometric information at milliarcsecond scales. They showed that a microquasar scenario cannot easily explain the observed changes in morphology. All these results, together with  the assessment of the low X-ray state (Martocchia et al. 2005) made the pulsar hypothesis tenable, and the possibility of explaining the H.E.S.S. phenomenology in such a case, an interesting working hypothesis.      High energy emission from pulsar binaries has been subject of study for a long time (just to quote a non-exhaustive list of references note the works of Maraschi and Treves 1981; Protheroe and Stanev 1987, Arons and Tavani 1993, 1994;  Moskalenko et al. 1993; Bednarek 1997, Kirk et al. 1999, Ball and Kirk 2000, Romero et al. 2001,  Anchordoqui et al. 2003, and others already cited above). LS 5039 has been recently subject of intense theoretical studies (e.g., Bednarek 2006, 2007; which we comment on in more detail below, Bosch-Ramon et al. 2005; B\\\"ottcher 2007; B\\\"ottcher and Dermer 2005; Dermer and B\\\"ottcher 2006; Dubus 2006a,b; Paredes et al. 2006; Khangulyan et al. 2007; Dubus et al. 2007).    In the penultimate paper mentioned in the list above, Khangulyan et al. (2007), and contrary to the assumption here, authors assumed a jet structure   perpendicular to the orbital plane of the system. The energy spectrum and lightcurves were computed, accounting for the acceleration efficiency, the location of the accelerator along the jet, the speed of the emitting flow, the inclination angle of the system, as well as specific features related to anisotropic inverse Compton (IC) scattering and pair production. Different magnetic fields, affecting Synchrotron emission, and the losses they produced, were also tested given a large model parameter space. Authors found a good agreement between H.E.S.S. data for some of  their models.  In the last of these papers, Dubus et al. (2007) computed the phase dependent lightcurve and spectra expected from inverse Compton interactions from electrons injected close to the compact object, assumed as a likely rotation-powered pulsar. Since the angle at which an observer sees the binary and propagating electrons changes with the orbit (see below), a phase dependence of the spectrum is expected, and anisotropic inverse Compton is needed to compute it. In general, they found that the  lightcurve is a good fit to the observations, except at the phases of maximum attenuation where pair cascade emission plays a role.  Dubus et al. (2007) do not consider cascading in their models, as we do here. Without cascading, zero flux is expected at a broad phase around periastron, which is not found. This lack of cascading in their model also affects the spectra, which are not reproduced well, particularly at the superior conjunction broad phases of the orbit. They mentioned that both, cascading and/or a change in the slope of the power-law injection for the interacting electron distribution  could be needed to explain the spectrum in these phases, what we explore in detail in this work.  In order to compute inverse Compton emission from LS 5039, we use, as in previous works, leptons interacting with the star photon field. Geometry is described there with different levels of detail, what influence the results. In general, cascading processes were not taken into account, and the goodness of fitting the H.E.S.S. data is arguable in most cases, both for the lightcurve and spectrum.  In none of the papers mentioned above, the theoretical predictions for the short timescale spectral variability found by H.E.S.S. in 0.1 phase binning was shown and compared with data.  We discuss these results from our model below.      In recent Letters (Sierpowska-Bartosik and Torres 2007, 2008) under the assumption that LS 5039 is composed by a pulsar rotating around an O6.5V star in the $\\sim 3.9$ day orbit,  we presented the results of a leptonic (for a generic hadronic model see Romero et al. 2003) theoretical modeling for the high-energy phenomenology observed by  H.E.S.S. These works studied the lightcurve, the spectral orbital variability in both broad orbital phases and in shorter (0.1 phase binning) timescales and have found a complete agreement between H.E.S.S. observations and our predictions. We have also analyzed how this model could be tested by Gamma-ray Large Area Space Telescope (GLAST), and how much time would be needed for this satellite in order to rule the model out in case theory significantly departs from reality. But many details of implementation which are not only useful for the case of LS 5039 but for all others close massive $\\gamma$-ray binaries, as well as many interesting results concerning the binary geometry, wind termination,  opacities to different processes along the orbit of the system, and further testing at the highest energy $\\gamma$-ray domain were left without discussion in our previous works. Here, we provide these details,  together with benchmark cases that are useful to understand the  formation of the very high-energy lightcurve and phase dependent spectra.       The rest of this paper is organized as follows. Next Section introduces the model concept and its main properties. It provides a discussion of geometry, wind termination, and opacities along the orbit of the system (we focus on LS 5039). An accompanying Appendix provides mathematical derivations of the formulae used and useful intermediate results that are key for the model, but too cumbersome to include them as part of the main text. It also deals with numerical implementation, and describes in detail the Monte Carlo simulation of the cascading processes. The results follow: Section 3 deals with a mono-energetic interacting particle population, and Section 4, with power-law primary distributions. Comparison with H.E.S.S. results is made in these Sections and details about additional tests are given. Final concluding remarks are provided at the end. ",
        "conclusions": "We presented the details of a theoretical model for the high-energy emission from close $\\gamma$-ray binaries, and applied it to the particular case of LS 5039. The model assumes a pulsar scenario, where either the pulsar or a close-to-the-pulsar shock injects leptons that after being reprocessed by losses to constitute a steady population, are assumed to interact with the target photon field provided by the companion star within the PWZ. The model accounts for the highly variable system geometry with respect to the observer, and radiative processes; essentially, anisotropic Klein-Nishina ICS and $\\gamma\\gamma$ absorption, put together with a Monte Carlo computation of cascading. The formation of lightcurve and spectra in this model was discussed in detail for the case where the interacting leptons are assumed mono-energetic and described by power-laws.  Comparing the interacting models of mono-energetic leptons and power-law distributions we can see similar dependencies for the spectral changes along the orbit and the GeV to TeV lightcurves. Thus, the case of a mono-energetic population is useful to understand some of the aspects regarding the formation of the observational features, although it does not match observational data. For the mono-energetic case, we have shown the spectra produced at characteristic phases (INFC, SUPC, periastron, apastron) for primary energies of 1 TeV and 10 TeV. In the power-law distribution model, the specific features of the photon spectra and the lightcurves produced for an specific primary electron energy  overlap. We can still notice the absorption features in the spectra produced at SUPC. We have discussed effects solely based on the optical depths and the general geometrical dependence along the orbit, as well as on the presence of the shock which terminates the pulsar wind in the direction to the observer at SUPC phases.  A power-law lepton distribution interacting in the PWZ describes very well the phenomenology found in the LS 5039 system at all timescales, both flux and spectrum-wise, even at the shortest timescales measured. This latter result is unexpected: we find that there is nothing a priori in the model that allows one to predict that when broad phase spectra data (INFC and SUPC) are reproduced so will be the data at the individual and much shorter phase-binning, less with such a good agreement. This result point perhaps to some reliability of the model, at least in its essential ingredients: geometry, cascading, interacting electron population.  However, this model certainly has room for improvement. We emphasize here that we do not have an a priori model for the interacting lepton population itself, although we have discussed the research on dissipation processes in the PWZ which may give raise to such distributions if it results from pulsar injection. In any case, the assumption of power-laws is an approximation to a more complex scenario where the real interacting lepton population is the result of a full escape-loss equation. In addition, we are not considering yet the multiwavelength emission at lower energies, since we left out of our description the synchrotron emission of electrons accelerated at the shock and the morphology of the shock along the orbit, what we expect to discuss elsewhere. Other than system scalings that are fixed by multiwavelength observations, the model is based on just a handful of free parameters, and it is subject to tests at high and very high-energy $\\gamma$-ray observations with both GLAST (described in more detail in Sierpowska-Bartosik \\& Torres 2008) and future samples of data at higher energies, where more statistics at finer phase bins can determine better the spectral evolution along the orbit of this interesting system.\\\\     {\\sl We acknowledge  extended use of IEEC-CSIC parallel computers cluster. We acknowledge W. Bednarek for discussions, and the Referee for useful comments. This work was supported by grants AYA 2006-00530 and CSIC-PIE 200750I029.}     \\newpage"
    },
    "0801/0801.4031_arXiv.txt": {
        "abstract": "We revisit the distant future of the Sun and the solar system, based on stellar models computed with a thoroughly tested evolution code. For the solar giant stages, mass-loss by the cool (but not dust-driven) wind is considered in detail. Using the new and well-calibrated mass-loss formula of Schr\\\"oder \\& Cuntz (2005, 2007), we find that the mass lost by the Sun as an RGB giant (0.332 $M_{\\odot}$, 7.59 Gy from now) potentially gives planet Earth a significant orbital expansion, inversely proportional to the remaining solar mass.  According to these solar evolution models, the closest encounter of planet Earth with the solar cool giant photosphere will occur during the tip-RGB phase. During this critical episode, for each time-step of the evolution model, we consider the loss of orbital angular momentum suffered by planet Earth from tidal interaction with the giant Sun, as well as dynamical drag in the lower chromosphere. As a result of this, we find that planet Earth will not be able to escape engulfment, despite the positive effect of solar mass-loss. In order to survive the solar tip-RGB phase, any hypothetical planet would require a present-day minimum orbital radius of about 1.15 AU. The latter result may help to estimate the chances of finding planets around White Dwarfs.  Furthermore, our solar evolution models with detailed mass-loss description predict that the resulting tip-AGB giant will not reach its tip-RGB size. Compared to other solar evolution models, the main reason is the more significant amount of mass lost already in the RGB phase of the Sun. Hence, the tip-AGB luminosity will come short of driving a final, dust-driven superwind, and there will be no regular solar planetary nebula (PN). The tip-AGB is marked by a last thermal pulse and the final mass loss of the giant may produce a circumstellar (CS) shell similar to, but rather smaller than, that of the peculiar PN IC 2149 with an estimated total CS shell mass of just a few hundredths of a solar mass. ",
        "introduction": "Climate change and global warming may have drastic effects on the human race in the near future, over human time-scales of decades or centuries. However, it is also of interest, and of relevance to the far future of all living species, to consider the much longer-term effects of the gradual heating of the Earth by a more luminous Sun as it evolves towards its final stage as a white dwarf star. This topic has been explored on several occasions (e.g. Sackmann, Boothroyd \\&\\ Kraemer 1993, Rybicki \\&\\ Denis 2001, Schr\\\"oder, Smith \\&\\ Apps 2001 (hereafter SSA)), and has been discussed very recently by Laughlin (2007).  Theoretical models of solar evolution tell us that the Sun started on the zero-age main sequence (ZAMS) with a luminosity only about 70\\%\\ of its current value, and it has been a long-standing puzzle that the Earth seems none the less to have maintained a roughly constant temperature over its life-time, in contrast to what an atmosphere-free model of irradiation would predict. Part of the explanation may be that the early atmosphere, rich in CO$_2$ that was subsequently locked up in carbonates, kept the temperature up by a greenhouse effect which decreased in effectiveness at just the right rate to compensate for the increasing solar flux. The r\\^ole of clouds, and their interaction with galactic cosmic rays (CR), may also be important: there is now some evidence (Svensmark 2007; but see Harrison et al. 2007 and Priest et al.  2007) that cosmic rays encourage cloud cover at low altitudes, so that a higher CR flux would lead to a higher albedo and lower surface temperature. The stronger solar wind from the young Sun would have excluded galactic cosmic rays, so cloud cover on the early Earth may have been less than now, allowing the full effect of the solar flux to be felt.   What of the future? Although the Earth's atmosphere may not be able to respond adequately on a short time-scale to the increased greenhouse effect of carbon dioxide and methane released into the atmosphere by human activity, there is still the possibility, represented by James Lovelock's Gaia hypothesis (Lovelock 1979, 1988, 2006), that the biosphere may on a longer time-scale be able to adjust itself to maintain life. Some doubt has been cast on that view by recent calculations (Scaife, private communication, 2007; for details, see e.g. Cox et al. 2004, Betts et al. 2004) which suggest that, on the century timescale, the inclusion of biospheric processes in climate models actually leads to an increase in carbon dioxide emissions, partly through a feedback that starts to dominate as vegetation dies back. In any case, it is clear that the time will come when the increasing solar flux will raise the mean temperature of the Earth to a level that not even biological or other feedback mechanisms can prevent. There will certainly be a point at which life is no longer sustainable, and we shall discuss this further in Section \\ref{habitablezone}.   After that, the fate of the Earth is of interest mainly insofar as it tells us what we might expect to see in systems that we observe now at a more advanced stage of evolution. We expect the Sun to end up as a white dwarf -- do we expect there to be any planets around it, and in particular do we expect any small rocky planets like the Earth?  The question of whether the Earth survives has proved somewhat tricky to determine, with some authors arguing that the Earth survives (e.g. SSA) and others (e.g. Sackmann et al. 1993) claiming that even Venus survives, while general textbooks (e.g. Prialnik 2000, p.10) tend to say that the Earth is engulfed. A simple model (e.g. SSA), ignoring mass loss from the Sun, shows clearly that all the planets out to and including Mars are engulfed, either at the red giant branch (RGB) phase -- Mercury and Venus -- or at the later asymptotic giant branch (AGB) phase -- the Earth and Mars. However, the Sun loses a significant amount of mass during its giant branch evolution, and that has the effect that the planetary orbits expand, and some of them keep ahead of the advancing solar photosphere. The effect is enhanced by the fact (SSA) that when mass loss is included the solar radius at the tip of the AGB is comparable to that at the tip of the RGB, instead of being much larger; Mars certainly survives, and it appears (SSA) that the Earth does also.  The crucial question here is: what is the rate of mass loss in real stars? Ultimately this must be determined from observations, but in practice these must be represented by some empirical formula. Most people use the classical Reimers' formula (Reimers 1975, 1977), but there is considerable uncertainty in the value to be used for his parameter $\\eta$, and different values are needed to reproduce the observations in different parameter regimes. In our own calculations (SSA) we used a modification of the Reimers' formula, which has since been further improved and calibrated rather carefully against observation, so that we believe that it is currently the best available representation of mass loss from stars with non-dusty winds (Schr\\\"oder \\&\\ Cuntz 2005, 2007 -- see Section \\ref{evolution}, where we explore the consequences of this improved mass-loss formulation).  However, although we have considerably reduced the uncertainties in the mass-loss rate, there is another factor that works against the favourable effects of mass loss: tidal interactions. Expansion of the Sun will cause it to slow its rotation, and even simple conservation of angular momentum predicts that by the time the radius has reached some 250 times its present value (cf. Table \\ref{solarmodels}) the rotation period of the Sun will have increased to several thousand years instead of its present value of under a month; effects of magnetic braking will lengthen this period even more. This is so much longer than the orbital period of the Earth, even in its expanded orbit, that the tidal bulge raised on the Sun's surface by the Earth will pull the Earth back in its orbit, causing it to spiral inwards.  This effect was considered by Rybicki \\&\\ Denis (2001), who argued that Venus was probably engulfed, but that the Earth might survive. An earlier paper by Rasio et al. (1996) also considered tidal effects and concluded on the contrary that the Earth would probably be engulfed. However, the Rybicki \\&\\ Denis calculations were based on combining analytic representations of evolution models (of Hurley, Pols \\&\\ Tout 2000) with the original Reimers' mass-loss formula rather than on full solar evolution calculations with a well-calibrated mass-loss formulation. The Rasio et al. paper also employed the original Reimers' formula, and both papers use somewhat different treatments of tidal drag.  We have therefore re-considered this problem in detail, with our own evolutionary calculations and an improved mass-loss description as the basis; full details are given in Sections \\ref{evolution} and \\ref{tiprgb}. ",
        "conclusions": "We have applied an improved and well-tested mass-loss relation to RGB and AGB solar evolution models, using a well-tested evolution code. While the habitable zone in the inner solar system will already move outwards considerably in the next 5 billion years of solar MS evolution, marking the end of life on Earth, the most critical and fatal phase for the inner planetary system is bound to come with the final ascent of the Sun to the tip of the RGB.  Considering in detail the loss of angular momentum by tidal interaction and dynamical drag in the lower chromosphere of the solar giant, we have been able to compare the evolution of the RGB solar radius with that of the orbit of planet Earth. Our computations reveal that planet Earth will be engulfed by the tip-RGB Sun, just half a million years before the Sun will have reached its largest radius of 1.2 AU, and 1.0 (3.8) million years after Venus (and Mercury) have suffered the same fate. While solar mass loss alone would allow the orbital radius of planet Earth to grow sufficiently to avoid this ``doomsday'' scenario, it is mainly tidal interaction of the giant convective envelope with the closely orbiting planet which will lead to a fatal decrease of its orbital size.  The loss of about 1/3 of the solar mass already on the RGB has significant consequences for the solar AGB evolution. The tip-AGB Sun will not qualify for an intense, dust-driven wind and, hence, will not produce a regular PN. Instead, an insubstantial circumstellar shell of just under 1/100\\,$M_{\\odot}$ will result, and perhaps a peculiar PN similar to IC\\,2149."
    },
    "0801/0801.1422_arXiv.txt": {
        "abstract": "We present a spatially resolved 894\\,$\\mu$m map of the circumstellar disk of the Butterfly star in Taurus (IRAS\\,04302+2247), obtained with the Submillimeter Array (SMA). The predicted and observed radial brightness profile agree well in the outer disk region, but differ in the inner region with an outer radius of \\about80-120\\,AU. In particular, we find a local minimum of the radial brightness distribution at the center, which can be explained by an increasing density / optical depth combined with the decreasing vertical extent of the disk towards the center. Our finding indicates that young circumstellar disks can be optically thick at wavelengths as long as 894\\,$\\mu$m. While earlier modeling lead to general conclusions about the global disk structure and, most importantly, evidence for grain growth in the disk (Wolf, Padgett, \\& Stapelfeldt~2003), the presented SMA observations provide more detailed constraints for the disk structure and dust grain properties in the inner, potentially planet-forming region (\\aboutless80-120\\,AU) vs.\\ the outer disk region (\\about120-300\\,AU). ",
        "introduction": "IRAS~04302+2247 is a Class~I protostar in the Taurus-Auriga molecular cloud complex whose equatorial plane is inclined edge-on to the line of sight (inclination = 90$^{\\rm{o}} \\pm 3^{\\rm{o}}$; Wolf et al.~2003, hereafter WPS03). Parallel to the increasing amount of observational constraints, such as ground-based near-infrared images and polarization maps (Lucas \\& Roche~1997, 1998), near-infrared images obtained with the Hubble Space Telescope (Padgett et al.~1999) and spatially resolved 1.3\\,mm and 2.7\\,mm maps (WPS03), several attempts have been undertaken to model the structure and physical conditions in the circumstellar disk and envelope of this object (e.g., Lucas \\& Roche~1998, WPS03, Stark et al.~2006). In particular, spatially resolved images of the circumstellar environment of the Butterfly Star, obtained in the near-infrared and millimeter wavelength range, allowed WPS03 to conclude that the grains in the envelope of this object cannot be distinguished from those of the interstellar medium, while grains have grown via coagulation by up to 2 - 3 orders of magnitude in the much denser circumstellar disk. The separated dust grain evolution is in agreement with the theoretical prediction of a sensitive dependence of grain growth on the location in the circumstellar environment of young (proto)stars: Grain growth is expected to occur on much shorter timescales in the dense region of circumstellar disks than in the thin circumstellar envelope. For the same reason a radial dependence of the dust grain evolution in the disk itself is expected. However, the observational data presented by WPS03 did not allow to constrain the spatial dependence of the dust grain properties in the disk.  Based on radiative transfer simulations, using the disk model by WPS03, we found that further insights into and constraints for the dust grain growth as the first stage of planet formation in the circumstellar disk of the Butterfly Star can be obtained with high-resolution submillimeter observations. As outlined in Sect.~\\ref{sect.ana}, the apparent structure of the disk is predicted to change significantly as the observing wavelength is decreased from millimeter to submillimeter wavelengths, allowing to constrain the radial and vertical disk structure and distribution of the dust grain properties.  In this paper we present and discuss new, spatially resolved observations of the circumstellar disk of the Butterfly Star, obtained with the Submillimeter Array\\footnote{ The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics and is funded by the Smithsonian Institution and the Academia Sinica.} (SMA) at 894\\,$\\mu$m. The observations and data reduction are described in Sect.~\\ref{sect.obs}, followed by a description of the data analysis (Sect.~\\ref{sect.ana}) and conclusions in Sect.~\\ref{sect.sum}. ",
        "conclusions": "\\label{sect.sum}  We obtained the first spatially resolved submillimeter map of circumstellar disk of the prominent Butterfly Star in Taurus. We find a good agreement between the observed and the predicted brightness profile in the outer region of the disk. However, a discrepancy between the predicted and observed radial flux distribution was found in the inner region (inside \\about80-120\\,AU), showing a decrease of the flux towards the center. This discrepancy -- a local minimum at the stellar position -- amounts to only \\about2$\\sigma$. However, the specific location of the minimum at the disk center and the slight symmetry of the corresponding local maxima with respect to the disk center indicate that this profile is an intrinsic feature of the real disk brightness distribution. Based on radiative transfer simulations and additional observations at 1.36\\,mm we exclude a large inner hole (strongly depleted from small dust grains) as a possible explanation of this local minimum in the radial brightness distribution.  These observations may provide the basis for a detailed model of the inner structure of the disk of the Butterfly Star. While earlier modeling lead to conclusion about grain growth in the disk (WPS03), the new SMA observations provide constraints for the disk structure and dust grain properties in the inner, potentially planet-forming region (\\about80-120\\,AU) vs.\\ the outer disk region (\\about120-300\\,AU). It was found that the optical depth in the inner disk region is higher than predicted by WPS03. The column density of dust along the line of sight is therefore higher in the inner disk region than derived from the previous disk model in which a perfect mixing of dust and gas throughout the entire disk is assumed. These observations indicate that there exists a higher dust density and therefore a higher dust-to-gas mass ratio in the inner disk region. This conclusion is based on the assumption that the density profile of the gas phase of the disk can be described by the same approach in the inner and outer disk region (Eq.~\\ref{dendisk}) and that there exists no discontinuity in the gas density profile.  Since the optical depth effect discussed above is constrained by both the radial and vertical disk structure and dust grain properties, such a model will have to take the predicted grain evolution and its dependence on the radial and vertical position in the disk into account. Furthermore, dust settling and the resulting increase of the grain-grain interaction probability will result in a further vertical dependence of the grain size distribution (e.g., Weidenschilling~1997). Beside the grain evolution, mixing processes, such as convection and radial mixing within circumstellar disks have to be considered since these processes lead to a redistribution of processed, evolved dust grains to outer, less dense and colder disk regions (see, e.g., Klahr et al.~1999, Gail~2001).  Our observations illustrate the high potential of submillimeter observations for studying circumstellar disks around young (proto)stars. Although the spatial resolution is significantly lower than that aimed for in the case of ALMA in a few years from now, the tight correlation between the density and temperature structure in the disk is already able to constrain the radial and vertical disk structure and thus the dust grain properties as a function of the distance from the disk midplane, based on the spatially resolved radial brightness profile obtained with the SMA in the submillimeter wavelength range."
    },
    "0801/0801.1108_arXiv.txt": {
        "abstract": "Passive spiral galaxies, despite their spiral morphological appearance, do not have any emission lines indicative of ongoing star formation in their optical spectra. Previous studies have suggested that passive spiral galaxies preferentially exist in infall regions of galaxy clusters, suggesting that the cluster environment is likely to be responsible for creating these galaxies. By carrying out spatially resolved long-slit spectroscopy on four nearby passive spiral galaxies with the Apache Point Observatory 3.5-m telescope, we investigated the stellar populations of passive spiral galaxies separately for their inner and outer regions. In the two unambiguously passive spiral galaxies among the four observed galaxies, H$\\delta$ absorption lines are more prominent in the outer regions of the galaxies, whereas the 4000-{\\AA} breaks (D$_{4000}$) are strongest in the inner regions of the galaxies. A comparison with a simple stellar population model for the two passive spiral galaxies indicates that the outer regions of the galaxies tend to harbour younger populations of stars. The strong H$\\delta$ absorption observed in the outer regions of the sample galaxies is consistent with that of galaxies whose star formation ceased a few Gyrs ago. Because of the large uncertainty in the absorption indices in our samples, further observations are needed in order to place constraints on the mechanisms that quench star formation in passive spiral galaxies. ",
        "introduction": "\\label{sec:section1} Passive spiral galaxies have peculiar spectroscopic characteristics among the galaxy populations having spiral morphologies. They show few or neither of the emission lines in H$\\alpha$ nor [O II] in optical spectra that would be indicative of ongoing star formation (Couch et al. 1998; Dressler et al. 1999; Poggianti et al. 1999; Goto et al. 2003b). The optical g $-$ r colours observed in such galaxies are found to be significantly redder than those of spiral galaxies with emission lines, which confirms the lower star formation rate among passive spiral galaxies (Poggianti et al. 1999).  A similar population of galaxies, known as 'anemic' spiral galaxies, has been found in the Virgo Cluster (van den Bergh 1976). They have smoothed spiral arms showing less prominent star formation activity than in other galaxies of the same Hubble type. Elmegreen et al. (2002) observed anemic spirals in the Virgo Cluster and found lower gas surface densities than for normal spirals for these galaxies. They showed that the gas surface density in these galaxies is below the threshold for star formation (Kennicutt 1989), suggesting that the lack of star formation is caused by the stripping of gas in the environment of the cluster (Elmegreen et al. 2002).  In order to investigate the effect of environment on the formation of such galaxies, Goto et al. (2003b) searched for passive spiral galaxies in all environments, including in dense cluster cores and field regions, using the large samples in the Sloan Digital Sky Survey (SDSS) data (Strauss et al. 2002). They found that passive spiral galaxies preferentially exist in the environment where a local galaxy density is intermediate between that of the cluster cores and the field regions (Goto et al. 2003b). This characteristic environment corresponds to the region where a significant decline in the star formation rate has been observed (Lewis et al. 2002; Gomez et al. 2003; Tanaka et al. 2004). These studies imply that cluster-related phenomena could be the main factors responsible for the formation of passive spiral galaxies.  Recent observations have revealed that the cluster environment may bring about a transformation in the galaxy population from star-forming spirals to passive galaxies, and thus passive spiral galaxies have been suggested to be a population in a transitional phase between these two populations (Poggianti et al. 1999). For instance, the fraction of blue galaxies is found to be larger in distant clusters than in local clusters, a phenomenon known as the Butcher-Oemler effect (Butcher \\& Oemler 1978; Couch et al. 1998; Kodama \\& Bower 2001; Goto et al. 2003d, 2004). The Butcher-Oemler effect implies that the environment of a cluster may affect the star formation activity of its member galaxies. Furthermore, morphological type (Goto et al. 2003c; Treu et al. 2003; Postman et al. 2005) is found to be correlated with various functions of cluster environment, such as local density or cluster centric radius (Goto et al. 2004; Tanaka et al. 2004). Numerical simulations (Bekki et al. 2002) have shown that gas-stripped galaxies may finally become S0 galaxies if no further accretion onto the disc occurs after the stripping. UV observations suggest that the cessation of star formation could take place before this morphological transformation into S0 galaxies (Moran, Ellis \\& Treu 2006). Indeed, Goto et al. (2003b) found that the Petrosian radius of S0s is smaller than that of spirals, which is consistent with the above scenario. Although these studies have provided significant broad implications concerning the origin of passive spiral galaxies, the details of the underlying physical mechanisms, especially the time-scale on which in-falling cluster galaxies terminate their star formation, are still uncertain.  Various mechanisms with different time-scales have been proposed. Because star formation could be sustained by cold gas accreted onto the disc, the time-scale is closely related to the time at which the reservoir of cold gas is removed by means of galaxy-intracluster medium (ICM) interactions (Treu et al. 2003; Diaferio et al. 2001; Kauffmann et al. 1999).  A possible mechanism that occurs on a relatively short time-scale is ram-pressure stripping (Gunn \\& Gott 1972; Fujita 2001; Fujita \\& Nagashima 1999; Fujita \\& Goto 2004; Abadi, Moore \\& Bower 1999; Quilis, Moore \\& Bower 2000). When a galaxy falls into a dense region of the cluster, ram pressure caused by the motion of the galaxy relative to the dense ICM removes the cold interstellar gas in the disc that is the fuel for the star formation (Gunn \\& Gott 1972). Vollmer et al. (2006) modelled the heavily ram-pressure-stripped galaxy NGC 4522 in the Virgo Cluster (Kenney, van Gorkom \\& Vollmer 2004). The ram-pressure-stripping model successfully reproduces the observed H I gas deficiency and the truncated gas disc of the galaxy (Vollmer et al. 2006). The time-scale for the ram-pressure stripping to terminate star formation was estimated to be 0.01-0.1 Gyr from numerical simulations (Abadi et al. 1999; Quilis et al. 2000; Fujita \\& Nagashima 1999).  Some authors, however, argue that ram pressure alone cannot explain the observed decline in star formation in cluster galaxies (Treu et al. 2003; Balogh et al. 2002; Kodama \\& Bower 2001). Treu et al. (2003) reported the mild decline in star formation at the periphery of the cluster Cl 0024$+$16, where the ram pressure may not be effective in stripping cold disc gas.  Other possible mechanisms have been proposed that require relatively longer time-scales to strip the galactic halo gas (Larson, Tinsley \\& Caldwell 1980), and these have been termed 'strangulation' (Fujita 2004; Tanaka et al. 2004) or 'starvation' (Treu et al. 2003; Boselli et al. 2006; Bekki et al. 2002). 'Strangulation' involves the stripping of warm halo gas through the interaction with the ICM in a situation where further supply of halo gases to the disc is disrupted. Because this can occur even in less dense environments, this may explain the observed decline of star formation among galaxies in the outer regions of clusters.  Mechanisms such as mergers cannot be responsible for creating passive spiral galaxies because they may disturb the spiral arms. Similarly, 'harassment', high-speed gravitational encounters between galaxies (Moore et al. 1996), may also lead to changes in morphology and thus cannot be a dominant mechanism for transforming normal spiral galaxies into passive ones.  In this paper, we perform spatially resolved spectroscopy on four of the passive spiral galaxies identified in our previous paper (Goto et al. 2003b) in order to obtain further constraints on the mechanism responsible for halting the star formation. The strength of the H$\\delta$ absorption line and 4000-{\\AA} break are compared with the simple stellar population (SSP) model constructed by Bruzual \\& Charlot (2003) in order to estimate the age of the stellar population. We then attempt to estimate the time-scale for the mechanism to create passive spiral galaxies.  The method of data reduction and analysis are described in Section 2, the results are in Section 3, the discussion based on the comparison with the SSP model is presented in Section 4, and conclusions are given in Section 5. Unless otherwise stated, we adopt the best-fitting WMAP cosmology: (h, $\\Omega_{m}$, $\\Omega_{L}$) = (0.71, 0.27, 0.73) (Bennett et al. 2003).     \\begin{table*} \\centering \\begin{minipage}{140mm} \\label{tab1:1} \\begin{tabular}{@{}lllcccccc@{}} \\hline & & & & & & \\multicolumn{2}{c}{Signal-to-noise ratio\\footnote{Signal-to-noise ratio measured by taking the standard deviation of the continuum waveband defined for spectral indices}}\\\\ Name & Ra(J2000)&Dec(J2000)&$z$&$\\sigma_{V}$ km s$^{-1}$\\footnote{Velocity dispersion $\\sigma_{V}$ and Petrosian radius $r_{\\rm P}$ in $r'$ band are taken from SDSS Data Release 5 \\citep{b3,Adelman}  }&$r_{\\rm P}$ arcsec & ${\\rm AP}_{in}$&${\\rm AP}_{out}$ & $M_r$\\footnote{Absolute magnitude in $r'$ band}\\\\ \\hline SDSS J021534.35$-$090537.0 & 02 15 34.35&$-$09 05 37.0&0.0687 & 112$\\pm$7&15.62 &19.7& 9.8& $-$22.58\\\\ SDSS J024732.02$-$065137.5 & 02 47 32.02&$-$06 51 37.5&0.0705 &131$\\pm$9 &10.31 & 17.2&12.0& $-$21.75\\\\ SDSS J033322.66$-$000907.5&03 33 22.66&$-$00 09 07.5&  0.0838 & - & 6.615 & 11.1 & 7.0 & $-$20.89\\\\ SDSS J074452.51$+$373852.7&07 44 52.51&$+$37 38 52.7&  0.0743 &94$\\pm$9& 8.329 & 9.4 & 4.6 &  $-$20.80\\\\ \\hline \\end{tabular} \\caption{Properties of target galaxies. Aperture radii for the spatially resolved analysis ($r_{\\rm P}$) are also shown.} \\end{minipage} \\end{table*}  \\begin{figure*} \\begin{minipage}{140mm} \\includegraphics[width=60mm,clip]{033.eps} \\includegraphics[width=60mm,clip]{074.eps} \\caption{SDSS $g', r', i'$ composite images \\citep{Fukugita} of target galaxies (left: SDSS J033322.65$-$000907.5, right: SDSS J074452.52$+$373852.7)} \\label{SDSSpictures} \\end{minipage} \\end{figure*}  \\section[]{The Method} \\label{sec:section2} \\subsection{Sample selection} The target galaxies, SDSS J021534.35$-$090537, J024732.02$-$065137.5, J033322.65$-$000907.5 (the left panel of Fig. 1) and J074452.52$+$373852.7 (the right panel of Fig. 1), are a subset of the passive spiral galaxies selected in Goto et al. (2003b) from the volume-limited sample of the SDSS data. Galaxies are selected based on the following criteria: (1) the inverse concentration parameter, which is defined as the ratio of the Petrosian 50 per cent radius to the Petrosian 90 per cent radius, is less than 0.5 (details of the use of this parameter for classifying morphological type are given in Shimasaku et al. 2001) in order to select galaxies having spiral morphology; (2) the absence of [O II] and H$\\alpha$ emission lines (the measured values of equivalent width are less than 1 $\\sigma$ error, Goto et al. 2003b), which are indicative of ongoing star formation activity.    \\subsection{Observation and Data reduction} \\label{subsec:2} The observations were carried out using the Dual Imaging Spectrograph (DIS) installed on the Apache Point Observatory 3.5-m telescope on 2004 October 18. Both the blue and red cameras on the DIS were used in a medium-dispersion mode, with the dispersion covering the wavelength range 3000-9000 {\\AA}. The pixel scales of the spatial axis for the blue and the red cameras were 0.42 and 0.40 arcsec pix$^{-1}$, respectively. To perform spatially resolved spectroscopy, a long slit with a slit width of 1.5 arcsec was used. Each sample galaxy was observed three times with an exposure time of 1000-1500 s. After bias-subtraction and flat-fielding had been applied, the three frames were combined into one frame. Spectra of standard stars, HR 718 and Hilt 600, were taken with an exposure time of 20 and 1 s, respectively, and used for flux calibration. The seeing size measured using the full width at half maximum (FWHM) of the point spread function of Hilt 600, which was monitored during the observation of the target galaxies, was $\\sim$1.7 and $\\sim$1.3 arcsec for the red and the blue camera, respectively.  To perform the spatially resolved analysis, we refer to the Petrosian radius in r' band ($r_{\\rm P}$), which is a measure of the surface brightness profile of galaxies, obtained through the SDSS (Blanton et al. 2001). The values are presented in Table 1. The long-slit data were divided into 11 spatial bins using the iraf routine apall, which outputs 11 spectra for all bins. Then, the three bins around the centre and the remaining eight bins sampling the spectra of two sides of the galaxies were summed, respectively, to increase the signal-to-noise ratio. Hereafter, these summed data for the inner and outer parts of the galaxies are designated as AP$_{in}$ and AP$_{out}$, respectively. AP$_{in}$ samples approximately $r$ $\\leq$ 0.27$r_{\\rm P}$, whereas AP$_{out}$ samples 0.27$r_{\\rm P}$ $<$ r $\\leq$ 1$r_{\\rm P}$, where $r$ represents the distance from the galaxy centre. The diameter of the inner part, $\\sim$3.7 arcsec, is much larger than the seeing size ($<$1.7 arcsec). Wavelength calibration was performed using observations of a HeNeAr lamp. The sensitivity function was obtained using the spectra of the standard stars HR 718 and Hilt 600. Because it was difficult to fit the whole wavelength range with a single sensitivity function, we excluded the data points at the edges of each red and blue spectrum. The resulting spectral coverage, included in the fitting of the sensitivity correction, is 3700-5650 {\\AA} for the blue spectrum and 6300-9500 {\\AA} for the red spectrum.  The wavelength resolution was measured using the FWHM of the emission lines of the HeNeAr lamp spectra. The result was $\\sim 8.5 \\pm 0.5$ {\\AA} FWHM at approximately 5000 {\\AA}.    \\subsection{Measurement of spectral features} \\label{subsec:3} \\begin{figure*} \\begin{center} \\label{fig:ha} \\includegraphics[width=45mm,angle=270]{ha_021.ps} \\includegraphics[width=45mm,angle=270]{ha_024.ps} \\includegraphics[width=45mm,angle=270]{ha_033.ps} \\includegraphics[width=45mm,angle=270]{ha_074.ps} \\caption{The rest-frame spectra around the ${\\rm H}\\alpha$ emission line for the observed galaxies. Solid lines show the spectra for ${\\rm AP}_{out}$ and dotted lines show the spectra for ${\\rm AP}_{in}$.  } \\end{center} \\end{figure*}  Before we measured the spectral features, the redshift for each galaxy was computed from the Ca H line at 3970 {\\AA}, one of the most prominent features in the obtained spectra, by fitting a Gaussian for this line with the iraf splot routine. The wavelengths of the observed spectrum were shifted to the rest-frame based on the derived redshifts. The derived redshifts are shown in Table 1.  The Lick/IDS absorption-line index system (Trager et al. 1998; Worthey 1994) was used to measure the strengths of absorption features in the spectra. The wavelength definitions for the index measurement were taken from Worthey \\& Ottaviani (1997) for the H$\\gamma$ and H$\\delta$ indices, and from Trager et al. (1998) for the other 21 indices. The equivalent width or magnitudes for each index were computed as defined in Trager et al. (1998). We also measured the 4000-{\\AA} break (D$_{4000}$) to estimate the age of the stellar population, as it is a broad feature and can be measured with a greater signal-to-noise ratio than can individual lines. D$_{4000}$ is widely used as a diagnostic of the age and metallicity of stellar populations and can be compared with values in the literature (Bruzual 1983; Gorgas et al. 1999; Kauffmann et al. 2003). D$_{4000}$ is calculated as defined in Balogh et al. (1999), who take a narrower spectral region for the red and blue continua than do Bruzual (1983). The reason for using this narrower definition is that it is less affected by the uncertainty in the sensitivity correction, and, more importantly, by the reddening effect (Balogh et al. 1999).  In order to confirm the absence of star formation activity for the observed candidates of passive spiral galaxies, equivalent widths of H$\\alpha$ and [O II] emission line were measured. Fig. 2 shows the spectral regions around the H$\\alpha$ emission lines. Two of the four target galaxies (SDSS J021534.35$-$090537 and SDSS J024732.02$-$065137.5) are found to show detectable H$\\alpha$ emission (H$\\alpha$ $-$ 1$\\sigma$ error $\\geq$ 10 {\\AA}) at AP$_{out}$, and thus do not meet the criteria for passive spiral galaxies defined in Goto et al. (2003b). The [O II] emission lines, whose equivalent widths are less than 7 {\\AA}, were also detected at AP out, indicating the presence of current star formation activity in the exterior regions of these galaxies. The non-detection of [O II] in the SDSS data was presumably the result of the aperture of the SDSS spectroscopic fiber (diameter of 3 arcsec), which samples only the inner parts of the galaxies, where the observed light may be dominated by the bulge component over the disc component (Abazajian et al. 2005). The results indicate the importance of using the whole light, including that from the outer regions of the galaxies, when identifying and investigating passive spiral galaxies.  We restrict our discussion to galaxies with no prominent star formation activity over the whole galaxy. The remaining two galaxies, SDSS J033322.65$-$000907.5 and SDSS J074452.52$+$373852.7, are hereafter denoted as SDSSJ0333$-$0009 and SDSSJ0744$+$3738. The velocity dispersion (available only for SDSSJ0744$+$3738), Petrosian radius, absolute magnitude in the r' band obtained by the SDSS, and the measured values of redshift and signal-to-noise ratio are shown in Table 1. The g', r', i'-composite SDSS images (Fukugita et al. 1996) of SDSSJ0333$-$0009 and SDSSJ0744$+$3738 are shown in Fig. 1. ",
        "conclusions": "\\label{sec:section4} In the following subsection, we first try to estimate the metallicity and the effects of $\\alpha$-enhancement on the absorption strengths for the observed galaxies using the measured Lick indices. The super-solar $\\alpha$/Fe ratio could affect the strength of the H$\\delta_{\\rm A}$ and the H$\\delta_{\\rm F}$ indices, as discussed in Thomas, Maraston \\& Bender (2003), which leads to an underestimate of the age of stellar populations. Considering the estimated metallicity and the effect of $\\alpha$-enhancement, Section 4.2 discusses the light-averaged ages of the stellar populations using a H$\\delta_{\\rm F}$-D$_{4000}$ plane.  The Lick indices measured with $>$1$\\sigma$ for both the inner and outer regions were Fe4383, Fe4668, H$\\beta$, Fe5270 and H$\\delta_{\\rm F}$ for SDSSJ0333$-$0009. For SDSSJ0744$+$3738, G4300, Fe4531 and Fe4668 were detected with $>$1$\\sigma$. Unfortunately, the H$\\delta_{\\rm F}$ index, which is used as an age indicator, does not have a good enough signal-to-noise ratio for SDSSJ0744$+$3738. Furthermore, neither the inner nor the outer region of this galaxy has an Mgb index, which is used as an indicator of metallicity and $\\alpha$-enhancement, with a signal-to-noise ratio $>$1. We therefore focus our discussion of the metallicity and $\\alpha$-enhancement on SDSSJ0333$-$0009, for which the iron indices and the Mgb indices were measured with a relatively high signal-to-noise ratio.    \\subsection{Metallicity and $\\alpha$-enhancement} \\label{sec:metal} The measured Lick indices (Fe4383, Fe4668, H$\\beta$, Fe5270 and H$\\delta_{\\rm F}$) for SDSSJ0333$-$0009 indicate that the inner region is more metal-rich than the outer region. Fig. 6 shows the strength of these indices for AP$_{in}$ (open triangles) and AP$_{out}$ (filled triangles) as a function of the Mgb index. The SSP models of Thomas et al. (2003) with age = 1.0, 5.0, 10.0, 15.0 Gyr and metallicity Z = 0.67, 0.00(Z$_{\\odot}$), -2.25, assuming [$\\alpha$/Fe] = 0.0, are overlaid. A comparison of the strength of the observed indices with that of the SSP models suggests that the observed indices for both the inner and outer regions are consistent with the model with super-solar metallicity. The iron indices (Fe4383, Fe4668 and Fe5270), especially, suggest that there is a radial gradient in the metallicity; the observed indices for the inner region are marginally consistent with [Z/H] = 0.67 models, whereas for the outer region the observed indices are reproduced by the model with [Z/H] = 0.00 in the Fe4668-Mgb plane and [Z/H] = 0.00-0.67 in the Fe5270-Mgb plane. Therefore, although the index strengths are associated with large uncertainties, they suggest that the inner region is more metal-rich than the outer region. It should be noted that H$\\delta_{\\rm F}$ for both regions and Fe4383 for the outer region were not reproduced by the models for [$\\alpha$/Fe] = 0.0. Rather, they seem to be consistent with models with [$\\alpha$/Fe] = 0.5, as discussed in the next section.  Comparisons of the Fe4383 indices, which are especially sensitive to the $\\alpha$-enhancement (Thomas et al. 2003), with the model suggest that the outer region of this galaxy possibly has a super-solar $\\alpha$/Fe ratio. Fig. 7 shows the strength of the Fe4383 index as a function of the Mgb index. Models with [$\\alpha$/Fe] = 0.0 and 0.5 for ages and metallicities the same as those shown in Fig. 6 are overlaid. Open and filled triangles show the observed index strength for AP$_{in}$ and AP$_{out}$, respectively. This figure indicates that the strength of Fe4383 for AP$_{out}$ is more consistent with the models with [$\\alpha$/Fe] = 0.5 than it is with those with [$\\alpha$/Fe] = 0.0.  In order to elucidate the dependence of the H$\\delta$ absorption strength on [$\\alpha$/Fe], Fig. 8 shows the model predictions for the H$\\delta_{\\rm F}$ index as a function of the Mgb index for [$\\alpha$/Fe] = 0.0 and 0.5. This illustrates that the difference in H$\\delta_{\\rm F}$ absorption owing to the $\\alpha$-enhancement is expected to be $\\sim$ 0.8-1.2 {\\AA} , depending on the age and metallicity of the stellar population.  The inner region (open triangles) is roughly consistent with the model with age $\\sim$ 1.0-1.5 Gyr, metallicity Z $\\sim$ 0.67, and [$\\alpha$/Fe] = 0.5. The outer region (filled triangles), by contrast, shows stronger H$\\delta_{\\rm F}$ absorption, which suggests that the stellar population in the outer region is, on average, younger than that in the inner region. The age and the metallicity implied from the plot for H$\\delta_{\\rm F}$ against Mgb are $<$ 1.0 Gyr and Z $\\sim$ 0.00-0.67, respectively. Although the metallicity may be difficult to determine quantitatively from Fig. 8 because of the quite large uncertainty in the Mgb index, Z $\\sim$ 0.00-0.67 is consistent with the value implied from the strength of the Fe4668 and Fe5270 indices against the Mgb index as shown in Fig. 6.  \\begin{figure*} \\begin{center} \\begin{tabular}{ccc} \\rotatebox{270}{\\resizebox{38mm}{!}{\\includegraphics{mgb_metals_SDSS033322_Fe4383.ps}}} & \\rotatebox{270}{\\resizebox{38mm}{!}{\\includegraphics{mgb_metals_SDSS033322_Fe4668.ps}}} & \\rotatebox{270}{\\resizebox{38mm}{!}{\\includegraphics{mgb_metals_SDSS033322_Hb.ps}}} \\\\ \\rotatebox{270}{\\resizebox{38mm}{!}{\\includegraphics{mgb_metals_SDSS033322_Fe5270.ps}}} & \\rotatebox{270}{\\resizebox{38mm}{!}{\\includegraphics{mgb_metals_SDSS033322_HdF.ps}}} &   \\\\ \\end{tabular} \\caption{he Lick indices measured with $>$1$\\sigma$ as a function of the Mgb index for SDSSJ0333$-$0009. Open and filled triangles show the measured index strength for AP$_{in}$ and AP$_{out}$, respectively. The solid lines show the models with [Z/H] = $-2.25$, $0.00$(Z$_{\\odot}$) and 0.67 (from left to right). The dotted lines show the models with ages = 1, 5, 10 and 15 Gyr (for the Fe4383, Fe4668 and Fe5270 indices, the age increases from left to right; for the H$\\beta$ and H$\\delta_{\\rm F}$ indices, the age increase, from top to bottom).} \\label{mgb_metals} \\end{center} \\end{figure*}   \\begin{figure*} \\begin{tabular}{c@{\\hspace{1cm}}c} \\begin{minipage}{0.45\\hsize} \\begin{center} \\includegraphics[width=55mm,angle=270]{mgb_fe4383.ps} \\end{center} \\caption{The Fe4383 index for AP$_{in}$ and AP$_{out}$ as a function of the Mgb index. The models with [$\\alpha$/H] = 0.0 and 0.5 are overlaid. For both values of [$\\alpha$/H], the solid lines show the models with [Z/H] = $-2.25$, 0.00 and 0.67 (from bottom to top) and the dotted lines show the models with ages = 1, 5, 10 and 15 Gyr (from bottom to top)} \\label{mgb_fe4383} \\end{minipage} & \\begin{minipage}{0.45\\hsize} \\begin{center} \\includegraphics[width=55mm,angle=270]{mgb_HdF.ps} \\end{center} \\caption{The H$\\delta_{\\rm F}$ index for AP$_{in}$ and AP$_{out}$ as a function of the Mgb index. The models with [$\\alpha$/H] = 0.0 and 0.5 are overlaid. For both values of [$\\alpha$/H], the solid lines show the models with [Z/H] = $-2.25$, 0.00 and 0.67 (from left to right) and the dotted lines show the models with ages = 1.0, 5.0, 10.0 and 15.0 Gyr (from bottom to top).} \\label{mgb_hdF} \\end{minipage} \\end{tabular} \\end{figure*}     \\subsection{Comparison with the SSP models} \\label{sec:age} In the following discussions, we use the H$\\delta_{\\rm F}$ index and the D$_{4000}$ feature as age indicators of the stellar population. The H$\\delta_{\\rm F}$-D$_{4000}$ plot is widely used as a diagnostic tool to estimate the age of a stellar population in galaxies (Kauffmann et al. 2003; Poggianti \\& Barbaro 1997; Balogh et al. 1999). The great advantage in using the H$\\delta_{\\rm F}$ and D$_{4000}$ indices as age indicators is that they are less affected by the reddening caused by interstellar dust (Kauffmann et al. 2003). We note that the ages discussed in this section refer to the light-averaged ages of stellar populations, and thus we aimed to estimate the time at which star formation ceased in the samples qualitatively.  In Fig. 9, the measured H$\\delta_{\\rm F}$ index and the values of D$_{4000}$ are compared with predictions from the simple stellar population (SSP) models in the galaxev package of Bruzual \\& Charlot (2003). We used the model with the Salpeter initial mass function with lower and upper cut-offs of stellar masses of 0.1 ${\\rm M}_{\\odot}$ and 100 ${\\rm M}_{\\odot}$, respectively. The model assumes an instantaneous burst of star formation at age = 0 for the various metallicities (Bruzual \\& Charlot 2003).  The spectra of the SSP model used in Fig. 9 were broadened to the instrumental resolution of the observations (8.5 {\\AA} FWHM), which corresponds to a velocity dispersion of $\\sigma$ $\\sim$ 197 km s$^{-1}$. The spectral broadening was applied using the program vel\\_disp included in the galaxev package (Bruzual \\& Charlot 2003). The resulting SSP model tracks for metallicities [Fe/H] = $-1.6464$, $-0.6392$, $+0.0932$, $+0.5595$ are overlaid in Fig. 9. Points with the same ages (1.0, 1.6, 2.5, 4.0, 6.2, 10.0 Gyr) are joined with dotted lines.  Based on the Fe4383 index, the outer regions of SDSSJ0333$-$0009 have [$\\alpha$/Fe] $>$ 0.0, as discussed in the previous subsection. Because a stellar population which has [$\\alpha$/Fe] $>$ 0.0 is reported to be enhanced in H$\\delta$ absorption strength compared to that having solar $\\alpha$/Fe ratios with similar age (Thomas et al. 2004), we have to take this effect into account to estimate the stellar population ages using H$\\delta_{\\rm F}$ indices. According to the stellar population model of Thomas et al. (2003), the enhanced strength in H$\\delta_{F}$ indices is up to $\\sim$ 1.2 {\\AA} for the model with [$\\alpha$/Fe] = 0.5. Therefore, the data points should be shifted downwards from those actually plotted in Fig. 6 by up to $\\sim$ 1.2 {\\AA} following the correction for the $\\alpha$-enhancement. For SDSSJ0744$-$3738, [$\\alpha$/Fe] could not be determined from the observed spectra, because the Mg indices were not measured with a high enough signal-to-noise ratio. Therefore, the additional uncertainty arising from the $\\alpha$-enhancement should be taken into account when considering H$\\delta_{\\rm F}$ versus D$_{4000}$ diagnostics.  Another point of caution in using the H$\\delta_{\\rm F}$-D$_{4000}$ plot is that the stellar absorption of H$\\delta_{\\rm F}$ could be contaminated by nebular emission (Kauffmann et al. 2003; Goto et al. 2003a). The nebular emission of H$\\alpha$ is weak for SDSSJ0333$-$0009 and SDSSJ0744$+$3738 ($<-4$ {\\AA}), and, furthermore, H$\\beta$ was detected in absorption. The effect of emission filling on the H$\\delta_{\\rm F}$ absorption is expected to be negligible ($\\sim$ 0.3 {\\AA} assuming the Case B recombination value for the H$\\alpha$/H$\\delta$ nebulae emission line ratio without dust extinction). In the following subsection, we estimate the light-averaged age of the stellar populations in the inner and outer regions for the observed galaxies.    \\subsubsection{SDSS J033322.65$-$000907.5} The measured H$\\delta_{\\rm F}$ and D$_{4000}$ for AP$_{in}$ and AP$_{out}$ are shown in Fig. 9 by the open and filled triangles, respectively. If we assume a metallicity [Fe/H] $> + 0.0932$ and a correction for the $\\alpha$-enhancement of $\\sim$1.2 {\\AA} as suggested in Section 4.1, the data for APout are broadly consistent with a model with age older than $\\sim$ 1.0 Gyr. By contrast, for the inner region (AP$_{in}$) the observed H$\\delta_{\\rm F}$ index seems stronger than the value that can be reproduced by the model with super-solar metallicity ([Fe/H] = $+0.5595$). However, the observed strength of D$_{4000}$ for AP$_{in}$ suggests an age older than $\\sim$ 1.6 Gyr. Overall, the H$\\delta_{\\rm F}$-D$_{4000}$ plane suggests that the average age of the stellar population is younger in the outer region than in the inner region for this galaxy.    \\subsubsection{SDSS J074452.52$+$373852.7} The open and filled squares in \\ref{fig:4} correspond to the data for AP$_{in}$ and AP$_{out}$, respectively, for SDSSJ0744$+$3738. For this galaxy, the metallicity and [$\\alpha$/Fe] are not available because of the large errors in the Lick indices. Although the determination of the age of the stellar population for SDSSJ0744$+$3738 is quite uncertain, the observed large difference in D$_{4000}$ between the core and the exterior of the galaxy imply that the two regions have distinct stellar populations in terms of their age and/or metallicity. For APin, the large value of D$_{4000}$ is marginally consistent with the model with metallicity [Fe/H] = $+0.5595$ in the H$\\delta_{\\rm F}$-D$_{4000}$ plane. Furthermore, in the outer region (AP$_{out}$), the H$\\delta_{\\rm F}$ index is quite strong (5.5$\\pm$11.1{\\AA}), although we note that the index strength is associated with large errors.  Galaxies with strong H$\\delta$ absorption are classified as 'H$\\delta$-strong' (HDS) galaxies \\citep{b57,b20,b9,2004MNRAS.348..515G,b61}. As strong H$\\delta$ absorption could arise from A-type stars in the main-sequence phase \\citep{b2}, HDS galaxies probably terminated their starburst activity $\\sim$0.1-1.5 Gyr ago \\citep{b32,b9}. Therefore, the strong H$\\delta_{\\rm F}$ absorption in the outer regions of the observed galaxies could be clear evidence that they terminated their star formation a few Gyrs ago. Nevertheless, because of the large uncertainty associated with H$\\delta_{F}$, we need further observations with higher signal-to-noise ratios to detect a clear signature of quenched star formation.    \\begin{figure} \\begin{center} \\includegraphics[width=80mm]{inst_burst.ps} \\caption{Equivalent width of H$\\delta_{\\rm F}$ as a function of D$_{4000}$. The open and filled triangles represent the data for AP$_{in}$ and AP$_{out}$, respectively, for SDSSJ0333$-$0009. The open and filled squares represent the data for APin and APout, respectively, for SDSSJ0744$+$3738. Overlaid lines show predictions of SSP models broadened to the instrumental resolution ($\\sigma$ = 197 km s$^{-1}$) with various metallicities [Fe/H] = $-1.6464$, $-0.6392$, $+0.0932$, $+0.5595$ (Bruzual \\& Charlot 2003). Dotted lines connect points of the same age (from bottom to top, 10.0, 6.2, 4.0, 2.5, 1.6 and 1.0 Gyr, respectively).} \\label{fig:4} \\end{center} \\end{figure}   \\subsubsection{Implications} It is interesting that through the spectral diagnostics of spatially separated regions of passive spiral galaxies we can estimate the history of these galaxies. Even our spectra with relatively low signal-to-noise ratios suggest that passive spiral galaxies were in a star-forming phase for several gigayears before ceasing activity. Moreover, one of the passive spirals, SDSSJ0744$+$3738, probably terminated its star formation approximately 1-2 Gyr ago. If confirmed, this result will provide us with an important constraint in uncovering the underlying (cluster-related) physical mechanism responsible for the creation of passive spiral galaxies; for example, ram-pressure stripping \\citep{b24,b37, b38,b47,b48,b25} stops star formation much quicker than the strangulation scenario \\citep{b52,b46}. To reach a firm conclusion on the subject, however, further observations are needed: it is well known that the age/metallicity estimates are degenerate, for example the spectra of an old ($>$ 2 Gyr) stellar population looks almost identical when the age is doubled and the total metallicity reduced by a factor of 3 \\citep{b62}. It is therefore important to break this degeneracy by obtaining spectra at a higher signal-to-noise ratio to measure indices sensitive only to age \\citep[e.g. Balmer lines;][]{b4} and only to metallicity \\citep[e.g.][]{b19}. Furthermore, we only had enough observing time for two passive spiral galaxies. As larger samples of passive spiral galaxies are now becoming available \\citep[e.g.][]{b1}, it is important to investigate the statistical proportion of passive spiral galaxies that do not suffer from shot-noise to draw firm conclusions on this subject."
    },
    "0801/0801.0167.txt": {
        "abstract": "Particle physics candidates for cosmological dark matter are usually considered as neutral and weakly interacting. However stable charged leptons and quarks can also exist and, hidden in elusive atoms, play the role of dark matter. The necessary condition for such scenario is absence of stable particles with charge -1 and effective mechanism for  suppression of free positively charged heavy species. These conditions are realized in several recently developed scenarios. In scenario based on Walking Technicolor model excess of stable particles with charge -2 and the corresponding dark matter density is naturally related with the value and sign of cosmological baryon asymmetry. The excessive charged particles are bound with primordial helium in techni-O-helium \"atoms\", maintaining specific nuclear-interacting form of dark matter. Some properties of techni-O-helium Universe are discussed. {\\bigbreak\\bgroup\\narrower\\small \\baselineskip=10pt plus .2pt \\parindent 0pt} \\def\\endabstract{\\par\\egroup\\bigbreak}  % --------------------- References ----------------------- % Environnement r\u00e9f\u00e9rences \\newenvironment{fref}{ \\renewcommand\\refname{R\\'ef\\'erences} \\begin{thebibliography}{99} \\setlength\\itemsep {-3pt} \\small}{\\end{thebibliography}} % id en anglais \\newenvironment{eref}{ \\renewcommand\\refname{References} \\begin{thebibliography}{99} \\setlength\\itemsep {-3pt} \\small}{\\end{thebibliography}} % ------------------ Date reception ----------------------- % \\outer\\def\\man#1{\\medskip \\noindent{\\it (Manuscrit re\\c cu le #1)}} %  % Macros possibles \\def\\eqalign#1{\\setlength{\\arraycolsep}{.2em} \\renewcommand{\\arraystretch}{1.5} \\begin{array}{rl} #1\\end{array}} %-------------------- Figures ----------------------------- \\long\\def\\Fig#1{{\\small \\noindent{\\bf Figure #1. }\\par}}  %\\makeatother \u0000\u00faS\u0001\u0088XJ\u0003%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%\u0000\u0000\u0000\u0000\u0000\u0000\u0000\u0000\u0000\u0000\u0000\u0000\u0000\u0000\u0000\u0000\u000f\u0000\u0001\u000e&\u008e\u0010\b\\bibitem{BKSR} Belotsky, K.M., Fargion, % A adapter \\`a l'ordinateur, ici pour un Macintosh %\\usepackage[T1]{fontenc} \\usepackage[applemac]{inputenc} \\usepackage{amssymb,amsmath} \\def\\volnum{Manuscrit} %\\def\\volnum{Volume xx no y, 200z} % pour la c\\'esure %\\french  % Format des AFLB \\usepackage{aflbcours} \\pdebut{1} \\def\\pacs#1{\\LP P.A.C.S.: #1} \\def\\s{{\\,\\rm s}} \\def\\g{{\\,\\rm g}} \\def\\eV{\\,{\\rm eV}} \\def\\keV{\\,{\\rm keV}} \\def\\MeV{\\,{\\rm MeV}} \\def\\GeV{\\,{\\rm GeV}} \\def\\TeV{\\,{\\rm TeV}} \\def\\sv{\\left<\\sigma v\\right>} \\def\\({\\left(} \\def\\){\\right)} \\def\\cm{{\\,\\rm cm}} \\def\\K{{\\,\\rm K}}\\title{Dark matter from stable charged particles?} \\titleshort{Dark matter from stable charged particles \\dots} \\author{Maxim Yu. Khlopov} \\authorshort{Maxim Khlopov} \\address{Centre for Cosmoparticle physics \"Cosmion\", \\\\ Miusskaya Pl. 4, 125047 Moscow, Russia\\\\ Moscow Engineering Physics Institute, 115409 Moscow, Russia, and \\\\ APC laboratory 10, rue Alice Domon et L\u00e9onie Duquet 75205 Paris Cedex 13, France\\\\Maxim.Khlopov@roma1.infn.it}  \\setlength{\\arraycolsep}{0pt}   \\begin{document}  A widely accepted viewpoint is that the approriate candidates for cosmological dark matter are neutral and weakly interacting particles. However it turns out that even stable charged leptons and quarks can exist and, hidden in elusive atoms, play the role of dark matter. The necessary condition for such scenario is absence of stable particles with charge -1 and effective mechanism of suppression for free positively charged heavy species. These conditions are realised in several recently developed scenarios. In scenario based on Walking Technicolor model excess of stable particles with charge -2 (without stable particles with charge -1) is naturally related with a cosmological baryon excess. ",
        "introduction": "The modern theory of Universe, based on General Relativity, has evolved the triumph of Einstein's ideas by putting cosmological term, first introduced by A.Einstein in 1917 \\cite{Lambda}, in the \"standard\" $\\Lambda$CDM model. The corresponding dark energy is the dominant element of the modern Universe, maintaining $70\\%$ of its total density.  General Relativity and Dark Energy maintain the frame for the portrait of elementary particles in the Universe. To survive in the Universe the particles should be stable, as are nuclei and electrons, composing the visible matter. However, one must also explain the modern dark matter density, corresponding to $25\\%$ of total density and exceeding the baryonic matter density by a factor of 5. The widely shared belief is that dark matter is nonbaryonic and consists of new stable particles.  For a particle with the mass $m$ the particle physics time scale is $t \\sim 1/m$ (here and further, if not indicated otherwise, we use the units $\\hbar = c = k = 1$), so in particle world we refer to particles with lifetime $\\tau \\gg 1/m$ as to metastable. To be of cosmological significance metastable particle should survive after the temperature of the Universe $T$ fell down below $T \\sim m$, what means that the particle lifetime should exceed $t \\sim (m_{Pl}/m) \\cdot (1/m)$. Such a long lifetime should find reason in the existence of an (approximate) symmetry. From this viewpoint, cosmology of dark matter is sensitive to the most fundamental properties of microworld, to the conservation laws reflecting strict or nearly strict symmetries of particle theory.   One can formulate the set of conditions under which new particles can be considered as candidates to dark matter (see e.g. \\cite{book,Cosmoarcheology,Bled06,Bled07} for review and reference).  \\begin{itemize} \\item[$\\bullet$] The particles should be stable or have lifetime larger,  than age of the Universe. \\item[$\\bullet$] They should fit the measured density of dark matter. Effects of their decay or annihilation should be compatible with the observed fluxes of electromagnetic background radiation and cosmic rays. \\item[$\\bullet$] More complicated forms of scalar fields, primordial black holes and even evolved primordial large scale structures are also possible, but in the latter case the contribution to the total density is restricted by the condition of the observed homogeneity and isotropy of the Universe\\item[$\\bullet$] The candidates for dark matter should decouple from plasma and radiation at least before the beginning of matter dominated stage. This is necessary to provide formation of large scale structure at the observed level of anisotropy of the cosmic microwave background radiation. \\end{itemize} The easiest way to satisfy these conditions is to involve neutral weakly interacting particles. However it is not the only particle physics solution for the dark matter problem. As we show here, new stable particles can have electric charge, but effectively behave as neutral and sufficiently weakly interacting.  Recently several elementary particle frames for heavy stable charged particles were proposed:  \\begin{itemize} \\item[(a)]  A heavy quark and heavy neutral lepton (neutrino with mass above half the Z-Boson mass) of fourth generation \\cite{N,Legonkov}; which can avoid experimental constraints \\cite{Q,Okun} and form composite dark matter species \\cite{I,lom,KPS06,Khlopov:2006dk}; \\item[(b)] A Glashow's ``sinister'' heavy tera-quark $U$ and tera-electron $E$, which can form a tower of tera-hadronic and tera-atomic bound states with ``tera-helium atoms'' $(UUUEE)$ considered as dominant dark matter \\cite{Glashow,Fargion:2005xz}. \\item[(c)] AC-leptons, predicted in the extension \\cite{5} of standard model, based on the approach of almost-commutative geometry \\cite{bookAC}, can form evanescent AC-atoms, playing the role of dark matter \\cite{5,FKS,Khlopov:2006uv,Khlopov:2006dk}.\\end{itemize}  Finally, it was recently shown in \\cite{KK} that an elegant solution is possible in the framework of walking technicolor models \\cite{Sannino:2004qp,Hong:2004td,Dietrich:2005jn,Dietrich:2005wk,Gudnason:2006ug,Gudnason:2006yj} and can be realized without an {\\it ad hoc} assumption on charged particle excess, made in the approaches (a)-(c).  In all these models, predicting stable charged particles, the particles escape experimental discovery, because they are hidden in elusive atoms, maintaining dark matter of the modern Universe. It offers new solution for the physical nature of the cosmological dark matter.  This approach differs from the idea of dark matter, composed of primordial bound systems of super heavy charged particles and antiparticles, proposed earlier to explain the origin of Ultra High Energy Cosmic Rays (UHECR) \\cite{UHECR}. To survive to the present time and to be simultaneously the source of UHECR super heavy particles should satisfy a set of constraints, which in particular exclude the possibility that they possess gauge charges of the standard model.  The particles, considered here, participate in the standard model interactions and we discuss the problems, related with various scenarios of composite atom-like dark matter, formed by heavy electrically charged stable particles.  \\vskip 30pt \\begin{fref} ",
        "conclusions": "Discussion} To conclude, the existence of heavy stable particles is one of the popular solutions for dark matter problem. If stable particles have electric charge, dark matter candidates can be atom-like states, in which negatively and positively charged particles are bound by Coulomb attraction. In this case there is a serious problem to prevent overproduction of accompanying anomalous forms of atomic matter.  Indeed, recombination of charged species is never complete in the expanding Universe, and significant fraction of free charged particles should remain unbound. Free positively charged species behave as nuclei of anomalous isotopes, giving rise to a danger of their over-production. Moreover, as soon as $^4He$ is formed in Big Bang nucleosynthesis it captures all the free negatively charged heavy particles. If the charge of such particles is -1 (as it is the case for tera-electron in \\cite{Glashow}) positively charged ion $(^4He^{++}E^{-})^+$ puts Coulomb barrier for any successive decrease of abundance of species, over-polluting modern Universe by anomalous isotopes. It excludes the possibility of composite dark matter with $-1$ charged constituents and only $-2$ charged constituents avoid these troubles, being trapped by helium in neutral OLe-helium, O-helium (ANO-helium) or techni-O-helium states.  The existence of $-2$ charged states and the absence of stable $-1$ charged constituents can take place in AC-model, in charge asymmetric model of 4th generation and walking technicolor model with stable doubly charged technibaryons and/or technileptons.   To avoid overproduction of anomalous isotopes, an excess of $-2$ charged particles over their antiparticles should be generated in the Universe. In all the earlier realizations of composite dark matter scenario, this excess was put by hand to saturate the observed dark matter density. In walking technicolor model this abundance of -2 charged techibaryons and/or technileptons is connected naturally to the value and sign of cosmological baryon asymmetry. These doubly charged $A^{--}$ techniparticles bind with $^4He$ in the neutral techni-O-helium states. For reasonable values of the techniparticle mass, the amount of primordial $^4He$, bound in this atom like state is significant and should be taken into account in comparison to observations.   A challenging problem is the nuclear transformations, catalyzed by techni-O-helium. The question about their consistency with observations remains open, since special nuclear physics analysis is needed to reveal what are the actual techni-O-helium effects in SBBN and in terrestrial matter. However, qualitatively one can expect a path for O-helium catalysis of heavy elements, making primordial heavy elements a signature for composite dark matter scenarios.  The destruction of techni-O-helium by cosmic rays in the Galaxy releases free charged techniparticles, which can be accelerated and contribute to the flux of cosmic rays. In this context, the search for techniparticles at accelerators and in cosmic rays acquires the meaning of a crucial test for the existence of the basic components of the composite dark matter. At accelerators, techniparticles would look like stable doubly charged heavy leptons, while in cosmic rays, they represent a heavy $-2$ charge component with anomalously low ratio of electric charge to mass.  The presented arguments enrich the class of possible particles, which can follow from extensions of the Standard Model and be considered as dark matter candidates. One can extend the generally accepted viewpoint that new stable particles should be neutral and weakly interacting. We have seen that they can also be charged and play the role of dark matter because they are hidden in atom-like states, which are not the source of visible light. The constraints on such particles are very strict and open a very narrow window for this new cosmologically interesting degree of freedom in particle theory. However, taking into account the exciting ability of O-helium to catalyze nuclear transformations of chemical elements, it is hard to estimate the expectation value of its discovery.   To conclude, the existence of heavy stable charged particles can offer new solution for dark matter problem. Dark matter candidates can be atom-like states, in which negatively and positively stable charged particles are bound by Coulomb attraction. Primordial excess of these particles over their antiparticles implies the mechanism of its generation and is still an open problem for all the considered models. However, even if such mechanism exists, there is a serious problem of accompanying anomalous forms of atomic matter.  Indeed, recombination of charged species is never complete in the expanding Universe, and significant fraction of free charged particles should remain unbound. Free positively charged species behave as nuclei of anomalous isotopes, giving rise to a danger of their over-production. Moreover, as soon as $^4He$ is formed in Big Bang nucleosynthesis it captures all the free negatively charged heavy particles. If the charge of such particles is -1 (as it is the case for tera-electron in \\cite{Glashow}) positively charged ion $(^4He^{++}E^{-})^+$ puts Coulomb barrier for any successive decrease of abundance of species, over-polluting modern Universe by anomalous isotopes. It excludes the possibility of composite dark matter with $-1$ charged constituents and only $-2$ charged constituents avoid these troubles, being trapped by helium in neutral OLe-helium or O-helium (ANO-helium) states.  The existence of $-2$ charged states and the absence of stable $-1$ charged constituents can take place in AC-model and in charge asymmetric model of 4th generation. In the first case, pregalactic abundance of $C^{++}$ exceeds by ten orders of magnitude the terrestrial upper limit on anomalous helium and the mechanism of suppression of this abundance is inevitably accompanied by observable effects of recombination and implies the existence of $y$ charge, possessed by AC leptons. In the second case, owing to excess of $\\bar U$ anti-quarks primordial abundance of positively charged $U$-baryons is exponentially suppressed and anomalous isotope over-production is avoided. Excessive anti-$U$-quarks should retain dominantly in the form of anutium, which binds with excessive $\\bar N$ and then with $^4He$ in neutral ANO-helium. In the both cases, OLe-helium (or ANO-helium) should exist and its possible role in nuclear transformation is the serious danger (or exciting advantage?) for composite dark matter scenario.  Galactic cosmic rays destroy ANO-helium (as well as OLe-helium), striking off $^4He$. It can lead to appearance of a free [anutium-$\\bar N$] component in cosmic rays, which can be as large as $[\\bar N \\Delta^{--}_{3 \\bar U}]/^4He \\sim 10^{-7}$ and accessible to PAMELA and AMS experiments. The estimation is two orders less in the case of free $A^{--}$ from cosmic ray destruction of OLe-helium \\cite{Khlopov:2006uv}.  In the context of composite dark matter like \\cite{I,lom} or \\cite{FKS,Khlopov:2006uv} accelerator search for new stable quarks and leptons acquires the meaning of critical test for existence of its charged components. Such test will be possible in experiment ATLAS/LHC  \\cite{lom,KPS06}.   %\\Figure{crossec.eps}{Taken from \\cite{lom} cross sections of %production of heavy stable quarks and leptons in experiment ATLAS. %Dash line shows the lower (conservative) level of LHC sensitivity %expected for 1st year of its operation.}"
    },
    "0801/0801.1564_arXiv.txt": {
        "abstract": "Observations show that the central black hole in galaxies has a mass $M$ only $\\sim 10^{-3}$ of the stellar bulge mass. Thus whatever process grows the black hole also promotes star formation with far higher efficiency. We interpret this in terms of the generic tendency of AGN accretion discs to become self--gravitating outside some small radius $R_{\\rm sg} \\sim 0.01 - 0.1$~pc from the black hole. We argue that mergers consist of sequences of such episodes, each limited by self--gravity to a mass $\\dme \\sim 10^{-3}M$, with angular momentum characteristic of the small part of the accretion flow which formed it. In this picture a major merger with $\\dmm \\sim M$ gives rise to a long series of low--mass accretion disc episodes, all chaotically oriented with respect to one another. Thus the angular momentum vector oscillates randomly during the accretion process, on mass scales $\\sim 10^3$ times smaller than the total mass accreted in a major merger event.  We show that for essentially all AGN parameters, the disc produced by any accretion episode of this type has lower angular momentum than the hole, allowing stable co-- and counter--alignment of the discs through the Lense--Thirring effect. A sequence of randomly--oriented accretion episodes as envisaged above then produces accretion discs stably co-- or counter--aligned with the black hole spin with almost equal frequency. Accretion from these discs very rapidly adjusts the hole's spin parameter to average values $\\bar a \\sim 0.1-0.3$ (the precise range depending slightly on the disc vertical viscosity coefficient $\\alpha_2$) from any initial conditions, but with significant fluctuations ($\\Delta a\\sim \\pm 0.2$) about these. We conclude (a) AGN black holes should on average spin moderately, with the mean value $\\bar a$ decreasing slowly as the mass increases; (b) SMBH coalescences leave little long--term effect on $\\bar a$; (c) SNBH coalescence products in general have modest recoil velocities, so that there is little likelihood of their being ejected from the host galaxy; (d) black holes can grow even from stellar masses to $\\sim 5\\times 10^9\\, \\msun$ at high redshift $z\\sim 6$; (e) jets produced in successive accretion episodes can have similar directions, but after several episodes the jet direction deviates significantly. Rare examples of massive holes with larger spin parameters could result from prograde coalescences with SMBH of similar mass, and are most likely to be found in giant ellipticals. We compare these results with observation. ",
        "introduction": "Astronomers generally agree that the nuclei of most galaxies contain supermassive black holes (SMBH), but as yet have little clear idea of how they grow. Cosmological large--scale structure simulations strongly suggest that galaxy mergers are the basic motor of growth, and predict that a black hole typically acquires a mass $\\dmm$ of order its own mass $M$ in such mergers (e.g. di Matteo et al., 2005, Li et al., 2007, and references therein). This is reasonable, as growth rates of this order are needed in order to reach the hole masses observed in the nearby universe. If one assumes further that all the mass gain $\\dmm$ carries the same sense of specific angular momentum, the hole always spins rapidly up to Kerr parameters $a$ close to unity (e.g. Volonteri et al., 2005; Volonteri \\& Rees, 2005). The resulting high radiative efficiency of accretion means that the Eddington limit severely restricts the rate of black hole mass growth, creating difficulties in understanding observations of high masses at early cosmic times.  However it is sigmificant that $\\dmm$ is only a very small part of the total mass involved in a major galaxy merger, in line with the observation that the black hole mass in nearby galaxies is typically only $\\sim 10^{-3}$ of the galaxy bulge mass. Moreover, current cosmological simulations are entirely unable to resolve the hydrodynamics of matter accretion on to the central black hole. The mass inflow time through an accretion disc of 1 pc radius is already $\\ga 10^9$~yr (Shlosman, Begelman \\& Frank, 1990), a factor $\\ga 10$ greater than the timescale on which rapid hole growth occurs. The best cosmological simulations resolve only lengthscales at least 100 times larger. Inevitably the simulations have to resort to sub--resolution recipes in trying to model the accretion process. For example it is usual to assume accretion at the Bondi rate from some large radius, whereas in reality the infalling gas must possess angular momentum which will complicate things.  In this picture the very small mass $\\dmm$ which accretes on to the central black hole must have almost zero angular momentum about it. Moreover, since the hole is growing its mass as rapidly as possible, it must be close to its Eddington limit and thus feeding back energy and momentum into its surroundings. Indeed simple models of the resulting momentum--driven feedback on to the host galaxy correctly reproduce the $M - \\sigma$ relation between black hole mass and velocity dispersion without any free parameter (King, 2003; 2005, Murray et al., 2005).  Along with this, there must be copious star formation from material close to the hole, i.e. at distances $\\sim 0.01 - 0.1$~pc (see below) where the angular momentum neglected in estimating the Bondi accretion rate becomes significant in supporting the captured gas against gravity. This star formation also injects energy and momentum into the gas near the hole. Hence although it is very difficult to predict the gas flow near the black hole it seems more likely to be chaotic than well--ordered, and rather inefficient in growing the hole compared with forming stars, thus keeping the latter a fairly constant fraction $\\sim 10^{-3}$ of the stellar bulge mass.  An obvious reason for the relative inefficiency of hole growth versus star formation is that an accretion disc becomes self--gravitating at radii $R$ where its mass exceeds $M_{\\rm sg} \\sim (H/R)M$, where $H$ is the disc scaleheight and $M$ the central (here black hole) mass (e.g. Pringle, 1981, Frank et al. 2002). For AGN this means that discs must be self--gravitating outside some radius $R_{\\rm sg} \\sim 0.01 - 0.1$~pc (Shlosman et al., 1990; Collin--Souffrin \\& Dumont, 1990; Hur\\'e et al., 1994). Cooling times in these regions are so short that self--gravity is likely to result in star formation (Shlosman \\& Begelman, 1989; Collin \\& Zahn, 1999). As discussed by King \\& Pringle, 2007 (hereafter KP07) we expect that most of the gas initially at radii $R > R_{\\rm sg}$ is either turned into stars, or expelled by those stars which do form, on a rapid (almost dynamical) timescale, while the gas which is initially at radii $R < R_{\\rm sg}$ forms a standard accretion disc, which slowly drains on to the black hole and powers the AGN. KP07 show that for local, low--luminosity AGN, fuelling by {\\it well separated} episodes of this type explains observational features such as the luminosity function for moderate--mass black holes, and the presence and location of a ring of young stars observed about the Galactic Centre.  This paper deals not with such low--luminosity AGN but instead with high--luminosity events causing major black--hole mass growth. In these cases our discussion suggests a picture in which, within an event, the flow on to the hole is episodic, via a rapid succession of accretion discs limited by self--gravity (as in the well--separated episodes in low--luminosity AGN), and thus with masses $\\dme \\sim \\msg$. The number of such episodes is evidently $\\sim \\dmm/\\msg$. Given that the flow is constantly stirred up by energy and momentum input, and that we are concerned only with the very small fraction of the gas with almost zero angular momentum, it is reasonable to expect the disc orientations to be essentially random.  We thus arrive at a picture of AGN accretion close to that suggested by Sanders (1984) (see also Heckman et al., 2004, Greene \\& Ho, 2007, and KP07). This picture reproduces several characteristic features of nearby AGN. In particular the random orientations of the accretion episodes mean that the radio jets show no correlation with the grand--design structure of the host, as observed for low--redshift AGN (Kinney et al., 2000), and inferred at higher redshift by Sajina et al (2007). Moreover the picture implies a luminosity function in broad agreement with that observed for moderate mass central black holes (KP07). Finally this picture provides a plausible explanation for the ring(s) of young stars seen around the black hole in the centre of our own Galaxy (Genzel et al., 2003; Lu et al., 2006).  Since observational properties such as the luminosity function of AGN and quasars are sensitive only to the total mass increase $\\Delta M_{\\rm merger}$, and not to its internal angular momentum budget, our picture predicts the same values for these observables as found in previous studies (e.g. Wyithe \\& Loeb, 2003) and is therefore in broad agreement with expected galaxy merger rates.  Clearly the proper hydrodynamic treatment of the kind of complex event we envisage must await advances in computing power. However we can gain some insight by simple arguments. ",
        "conclusions": "We have argued that accretion on to supermassive black holes in AGN occurs through events with total mass $\\dmm \\sim M$, consisting of long sequences of randomly oriented disc accretion episodes with masses $\\dme \\sim 10^{-3}M$. The mass of these disc episodes is strongly constrained by the generic tendency of such discs to become self--gravitating and form stars outside quite small radii $\\sim 0.01 - 0.1$~pc. The self--gravity constraint also limits the disc angular momentum. This is always less than the maximum that the hole's mass would allow it to have. A hole with significant spin therefore stably co-- and counter--aligns initially prograde and retrograde discs respectively, producing a net spindown until its angular momentum is reduced to a relatively modest value, with fluctuations $\\Delta a \\sim \\pm 0.2$ about the average $\\bar a$ given by (\\ref{av}). We note from (\\ref{av}) that $\\bar a$ itself decreases slowly as the black hole mass decreases.  We can compare these results with those for coalescences of SMBH as a consequence of the mergers of their host galaxies. Hughes and Blandford (2003) show that the average secular evolution of black hole spin under coalescences is \\begin{equation} a = a_0\\biggl({M_0\\over M}\\biggr)^{2.4} \\end{equation} where $a_0$ is the original spin parameter of a hole within initial mass $M_0$. Thus coalescences doubling the mass of a hole on average reduce $a$ by a factor $\\sim 5.3$. This is an even faster decline of spin than given by assuming fixed angular momentum but growing mass, which would produce $a \\propto M^{-2}$, reflecting the fact that retrograde coalescences are more effective in reducing spin than prograde ones are in growing it. Evidently coalescences can reduce $a$ quite quickly to zero. However we have seen that accretion very rapidly pushes $a$ back towards the mean value $\\bar a$ (cf eqns \\ref{equil}, \\ref{av}) from any starting point. Hence on average coalescences have little net effect on the mean value of $a$.  We conclude that repeated accretion episodes and coalescences on average drive the black holes in AGN towards fairly low values of $a$, with statistical fluctuations around them of order $\\Delta a = \\pm 0.2$.  Given this result, we should consider the observational evidence constraining values of $a$ for SMBH. Much of this is derived from the Soltan (1982) argument that the totality of background light in the low--redshift Universe is compatible with black--hole growth with radiative efficiency $\\epsilon \\simeq 0.1$ (e.g. Wang et al., 2006; Treister \\& Urry, 2006; Hopkins, Richard \\& Hernquist, 2007), which formally requires $a \\simeq 0.67$. However $\\epsilon$ is such a slow function of $a$ except near $a=1$ (see Figure \\ref{epsilon}) that this conclusion is vulnerable to the systematic effects inherent in trying to estimate the efficiency from observation. Evidently we have to find $\\epsilon \\simeq 0.1$ from observation to an accuracy of a factor $\\la 2$ in order to exclude {\\it any} value of $a$ other than those very near unity. Since retrograde and prograde accretion appear equally probable, the effective efficiency for the Soltan argument is actually the symmetrized curve \\begin{equation} \\epsilon_{\\pm} = {1\\over 2}[\\epsilon(a)+\\epsilon(-a)] \\label{epspm} \\end{equation} in Figure (\\ref{epsilon}), whose vertical range is now compressed to $0.057 - 0.230$.  In this case only a very accurate estimate of $\\epsilon$ can give real information about $a$ at all.  \\begin{figure} \\centerline{\\epsfxsize9cm \\epsfbox{epsilonbw.eps}} \\caption{Accretion efficiency $\\epsilon$ versus spin parameter $a$ (solid curve). Negative values of $a$ denote retrograde accretion. The dashed curve shows the effective $\\epsilon=\\epsilon_{\\pm}$ given by assuming that retrograde and prograde accretion are equally probable. This is the efficiency relevant to the Soltan argument.} \\label{epsilon} \\end{figure}  Streblyanska et al. (2005) suggest a different method of estimating the typical spin $a$ for SMBH. They find that the composite X--ray spectrum of a sample of AGN shows evidence for a broad iron line. They fit lines produced with assumed emissivity and line centre energy from both non--rotating and rotating black holes, and argue that the outer radius of the emission region for the non--rotating hole is unreasonably small in the case of the Type I AGN in their sample, although the statistical fit is marginally preferred over a rotating hole. The best--fitting rotating hole for this group has an inner emission radius of $3GM/c^2$, which requires $a > 0.78$. Since holes with such values of $a$ would automatically have higher efficiencies and thus presumably stronger iron lines, it is difficult to know what fraction of the SMBH in the sample actually have such fairly large spin parameters (note that, as discussed below, our arguments only refer to the statistical mean $\\bar a$, and individual SMBH can differ significantly from this).  We remark as a general caveat that the search for any effect depending on $a$, or the interpretation of a particular observational effect as being due to $a$, is likely to produce a bias towards $a \\sim 1$. This effect is particularly marked because of the very high efficiency of prograde accretion for $a \\simeq 1$ (see Fig \\ref{epsilon}). We note that Brenneman \\& Reynolds (2006) estimate $a = 0.989^{+0.009}_{-0.002}$ in the specific case of the broad iron line in MCG-06-30-15, remarkably close to the maximum allowed value.  Of course since our treatment is statistical it does not rule out larger spin values in individual cases. The most important way of producing large values of $a$ in individual cases is in coalescences of two SMBH. Hughes \\& Blandford (2003) show that this can produce significant $a$ for prograde coalescences of two black holes with similar mass. This suggests that giant ellipticals offer a promising site for significant values of $a$. Another conceivable way of producing exceptions to our general conclusions is that we have assumed accretion from thin discs, in which cooling is efficient, so that the total disc mass is limited to $\\sim H/R$ times that of the hole. In some situations it may happen that cooling is inefficient and our conclusions do not hold. More work is needed here, although such cases are likely to be rare. Finally one might argue that in some cases an accretion event with mass $\\gg M_{\\rm sg}$ may contrive to retain some memory of an overall angular momentum, despite the fact that it must have a near--radial orbit within the host galaxy in order to accrete at all (cf. Kendall et al., 2003: recall that accretion within a Hubble time requires disc radii $\\la 1$~pc). In this case individual episodes no longer have uncorrelated angular momenta as we assumed. This is effectively what is assumed by Volonteri \\& Rees (2005), and Volonteri et al (2007) for {\\it all} episodes."
    },
    "0801/0801.4177_arXiv.txt": {
        "abstract": "The Newtonian solid-mechanical theory of non-compressional spheroidal and torsional nodeless elastic vibrations in the homogenous crust model of a quaking neutron star is developed and applied to the modal classification of the quasi-periodic oscillations (QPOs) of X-ray luminosity in the aftermath of giant flares in SGR 1806-20 and SGR 1900+14. Particular attention is given to the low-frequency QPOs in the data for SGR 1806-20 whose physical origin has been called into question. Our calculations suggest that unspecified QPOs are due to nodeless dipole torsional and dipole spheroidal elastic shear vibrations. ",
        "introduction": "This study was undertaken as a part of current extensive theoretical investigations of asteroseismology of a neutron star\\cite{P-05,GSA-06,SA-07,Lee-07,Lev-07,Bast-07} (see also references therein) that have been boosted by recent discovery of quasi-periodic oscillations (QPOs) of X-ray luminosity in the decaying flares of two magnetars\\cite{Isr-05,SW-06}, SGR 1806-20 and SGR 1900+14, with concomitant suggestion to interpret this variability as caused by quake induced differentially rotational, torsional, oscillations.  Following this suggestion the focus of  most theoretical works is on computing the frequency spectra of odd-parity torsional mode of shear vibrations and less attention is paid to the even parity spheroidal elastic mode. However, from the viewpoint of modern global seismology\\cite{LW-95,AR-03}, the spheroidal vibrational mode in a solid star and planet has the same physical significance as the toroidal one in the sense that these two fundamental modes owe their existence to one and the same restoring force\\cite{McD-88,Bast-99,Bast-IJ-07}. In this light there is a possibility that, by not considering both these modes on an equal footing, we may miss discovering certain essential novelties which are consequences of solid mechanical laws governing seismic vibrations of superdense matter of neutron stars. Adhering to this attitude and continuing our current investigations\\cite{Bast-IJ-07}, we derive here spectral equations for the frequency of both spheroidal and torsional elastic nodeless vibrations in the solid crust of quaking neutron star and examine what conclusions can be drawn regarding low-frequency QPOs whose physical nature still remain unclear. ",
        "conclusions": "The obtained spectral formulae for the frequency of nodeless elastic vibrations trapped in the finite-depth seismogenic layer may be of some interest in its own right from the viewpoint of general theoretical seismology\\cite{LW-95} in the sense that they can be utilized in the study of seismic vibrations of more wide class of solid celestial objects such as Earth-like planets. One of the remarkable feature of considered model is that the dipole nodeless overtones possess properties of Goldstone soft modes, that is, the dipole overtones emerge if and on only if the elastic vibrations turn out to be locked in the peripheral layer of finite thickness. It is shown that obtained spectral equations are consistent with the existence treatment of low-frequency QPOs in the X-ray luminosity of flares SGR 1900+14 and  SGR 1806-20 as caused by quake-induced torsional nodeless vibrations. What is newly disclosed here is that previously non-identified low-frequency QPOs in data for SGR 1806-20 can be attributed to nodeless dipole torsional and spheroidal vibrations, namely, $\\nu(_0t_1)=18\\,\\mbox{Hz}$ and $\\nu(_0s_1)=26\\,\\mbox{Hz}$."
    },
    "0801/0801.4388_arXiv.txt": {
        "abstract": "Studies of nearby elliptical and S0 galaxies reveal that roughly half of the low mass X-ray binaries (LMXBs), which are luminous tracers of accreting neutron star or black hole systems, are in clusters. There is a surprising tendency of LMXBs to be preferentially associated with metal-rich globular clusters (GCs), with metal-rich GCs hosting three times as many LMXBs as metal-poor ones. There is no convincing evidence of a correlation with GC age so far. In some  galaxies the LMXB formation rate varies with GC color even within the metal-rich peak of the typical bimodal cluster metallicity distribution. This provides some of the strongest evidence to date that there are metallicity variations within the metal-rich GC peak, as is expected in hierarchical galaxy formation scenarios. We also note that apparent correlations between the interaction rates in GCs and LMXB frequency may not be reliable because of the uncertainties in some GC parameters. We argue in fact that there are considerable uncertainties in the integrated properties of even the Milky Way clusters that are often overlooked. ",
        "introduction": "Chandra X-ray images of nearby ellipticals and S0s reveal large numbers of point sources, confirming a long-standing suggestion that the hard X-ray emission in many of these galaxies comes predominantly from X-ray binaries. In the old stellar populations of these galaxies the bright, L$_X$$\\geq$10$^{37}$ erg s$^{-1}$ sources detected in typical Chandra observations are low mass X-ray binaries, binary systems comprising a neutron star or black hole accreting via Roche lobe overflow from a low mass companion.  An interesting characteristic of LMXBs is that they are disproportionately abundant in globular clusters. In the Milky Way (MW) the probability of finding a LMXB among the stars in the globular clusters is a few hundred times larger than for field stars. This has long been attributed to efficient formation of LMXBs in GCs due to dynamical interactions in the core. Early type galaxies are ideal for studies of the LMXB-GC link as they are particularly abundant in clusters. While there are $\\approx$150 globular clusters in the Milky Way, bright elliptical galaxies typically host several thousand GCs. Globular clusters are among the simplest stellar systems known; Thus the stellar population in each of these compact, $\\sim$10$^5$ M$_\\odot$, luminous systems has well defined characteristic properties such as age and metallicity.  The identification of LMXBs with GCs provides a unique opportunity to probe the effects of these parameters on LMXB formation and evolution.  \\begin{figure}[!ht] \\begin{center} \\centerline{\\psfig{figure=kundu_a_fig1.ps,width=4.5in}} \\caption{ Top: The V-I colors of GCs vs. distance from the center of NGC 4472 and GC color distribution. LMXB-GC matches are indicated by large filled circles/bins. The colors of these clusters primarily trace the underlying (bimodal) metallicity distribution. Bottom: V magnitude of globular clusters vs. half light radius and the globular cluster luminosity function. LMXBs are preferentially located in the high mass and metal-rich GCs. There is a weak anti-correlation with GC half-light radius, but no obvious correlation with galactocentric distance. Each of these trends has been confirmed in other galaxies. } \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0801/0801.2084_arXiv.txt": {
        "abstract": "We performed a photometric multicolor survey of the core of the Canis Major over-density at $l \\approx 244^{o}$, $b\\approx -8.0^{o}$, reaching V $\\sim$ 22 and covering $0^{o}.3 \\times 1^{o}.0$. The main aim is to unravel the complex mixture of stellar populations toward this Galactic direction, where in the recent past important signatures of an accretion event have been claimed to be detected. While our previous investigations were based on disjointed pointings aimed at revealing the large scale structure of the third Galactic Quadrant, we now focus on a complete coverage of a smaller field centered on the Canis Major over-density.  A large wave-length baseline, in the $UBVRI$ bands, allows us to build up a suite of colour colour and colour magnitude diagrams, providing a much better diagnostic tool to disentangle the stellar populations of the region. In fact, the simple use of one colour magnitude diagram, widely employed in all the previous studies defending the existence of the Canis Major galaxy, does not allow one to separate the effects of the different parameters (reddening, age, metallicity, and distance) involved in the interpretation of data, forcing to rely on heavy modeling.  In agreement with our previous studies in the same general region of the Milky Way, we recognize a young stellar population compatible with the expected structure and extension of the Local (Orion) and Outer (Norma-Cygnus) spiral arms in the Third Galactic Quadrant. Moreover we interpret the conspicuous intermediate-age metal poor population as belonging to the Galactic thick disk, distorted by the effect of strong disk warping at this latitude, and to the Galactic halo. ",
        "introduction": "In the last years, we have used photometric observations of young open cluster fields to probe the spiral structure in the third Galactic Quadrant (TGQ, $180^{o} \\leq l \\leq 270^{o}$; Carraro et al. 2005a, Moitinho et al. 2006, V\\'azquez et al. 2008), motivated by the very poor knowledge of this portion of the Galaxy's periphery.  Interestingly, important low latitude accretion phenomena have been recently claimed to be ongoing in this part of the Galaxy, such as the Canis Major over-density (CMa, Bellazzini et al. 2004), and the Monoceros Ring (MRi, Newberg et al. 2002). Clearly, a detailed description of the structure and stellar populations of the Galactic disc (thin plus thick) is mandatory to discriminate between Galactic and extragalactic material.  The TGQ is a special region of the Milky Way's outskirts, characterized by significant absorption windows as the Puppis (l $\\sim 243^{o}$) window (Fitzgerald 1968, Moffat et al. 1979, Janes 1991, Moitinho 2001), which allows one to detect very distant star clusters (Baume et al. 2006). Besides, and interestingly, young star clusters are found at low Galactic latitudes, underlining the fact that the young Galactic disk is significantly warped in these directions (May et al. 1997, Momany et al. 2004, Moitinho et al. 2006, Momany et al. 2006, L\\'opez-Corredoira et al. 2007).  In V\\'azquez et al (2008), by combining optical and CO observations, we have provided a fresh and very detailed picture of the spiral structure in the TGQ, showing that this region is characterized by a complicated spiral pattern. The outer (Norma-Cygnus) arm is found to be a grand design spiral feature defined by young stars, whereas the region closer to the Sun (d$_{\\odot}$ less than 9 kpc) is dominated by a conspicuous inter-arm structure, at l $\\sim$ 245$^o$, the Local spiral arm.  In this region, Perseus is apparently defined by gas and dust, and does not appear to be traced by an evident optical young stellar population, similarly to what can be found in other galaxies such as M~74 (V\\'azquez et al. 2008).  The analysis carried out on a substantial fraction of the stellar fields we observed revealed a complicated mixture of young and old populations.  Although centered on catalogued star clusters (Dias et al. 2002), a few colour magnitude diagrams (CMD) do not reveal star clusters but, and more interestingly, show hints of a young, diffuse, and distant stellar populations, which has become recently referred to as {\\it Blue Plume} ( Bellazzini et al. 2004, Dinescu et al. 2005, Mart\\'inez-Delgado et al. 2005,Carraro et al. 2005 ). Since the disk is warped and flared in these directions (Momany et al. 2006), the lines of sight are expected to cross both the thin and thick disk population in front of a particular target, in a way that the analysis of the CMD becomes very challenging (see for instance the analysis of the field toward the star cluster Auner~1, Carraro et al. 2007).  In this paper we present a photometric analysis in the $UBVRI$ filters of the stellar populations in 3 wide field pointings toward the CMa over-density.  Sect.~2 describes the observation and data reduction strategies. In Sect.~3 we discuss various colour combination CMDs, while Sects.~4 and 5 are dedicated to illustrate and analyze the TCD as a function of magnitude. Finally, Sect.~6 summarizes our findings.   \\begin{table*} \\centering \\begin{tabular}{cccccccccccc} \\hline\\hline Designation     & $\\alpha (2000.0)$ & $\\delta(2000.0)$ & l & b & U & B & V & R & I  & Airmass & Seeing\\\\ & & & [deg] & [deg] & secs & secs & secs & secs & secs & & arcsec\\\\ \\hline Field 1&07:22:51&-30:59:20&244.00&-07.50&20,180,1800&10,150,1500&5,60,900&5,60,900&5,60,900&1.00-1.30&0.83-1.02\\\\ Field 2&07:20:46&-31:09:36&244.00&-08.00&20,180,1800&10,150,1500&5,60,900&5,60,900&5,60,900&1.00-1.30&0.90-1.12\\\\ Field 33&07:20:21&-31:15:43&244.00&-08.10&20,180,1800&10,150,1500&5,60,900&5,60,900&5,60,900&1.20-1.80&1.28-2.11\\\\ \\hline\\hline \\end{tabular} \\caption{ List of pointings discussed in this paper. For each pointing, equatorial and galactic coordinates are reported together with the set of filters used, and the range of exposure time, air-mass and typical seeing. The fields are shown in Fig.~1 to 3} \\end{table*}  \\section[]{Observations and Data Reduction}    \\subsection{Observational details} $UBVRI$ images of 3 overlapping fields  (see Figs.~1 to 3) in the Third Quadrant of the Milky Way toward the CMa over-density were obtained at the Cerro Tololo Inter-American Observatory 1.0m telescope, which is operated by the SMARTS\\footnote{http://www.astro.yale.edu/smarts/} consortium. The telescope is equipped with a new 4k$\\times$4k CCD camera having a pixel scale of 0$^{\\prime\\prime}$.289/pixel which allows to cover a field of $20^{\\prime} \\times 20^{\\prime}$ on the sky.  Observations were carried out on the nights of November 28 and December 3, 2005. The two nights were part of a 6 night run.  In the first night we observed the fields $\\#1$ and $\\#2$ (see Table~1) while field $\\#3$ was observed in the last night of the run.  The CCD is read out through 4 amplifiers, each one with slightly different bias levels and gains.  Pre-processing was done using the procedure developed by Philip Massey\\footnote{http://www.lowell.edu/users/massey/obins/y4kcamred.html}. Briefly, the procedure trims and corrects the images for bias, flat-field, and bad pixels, preparing them from photometric extraction. A series of skyflats was employed in all the filters.  \\begin{figure} \\includegraphics[width=\\columnwidth]{MN1806fig4.eps} \\caption{Photometric solution in $UBVRI$ for standard stars. See Table~2 for details. $\\sigma$, on the right, indicates the {\\it rms} of the fit.} \\end{figure}  \\begin{table} \\caption{Calibration coefficients.} \\begin{tabular}{rrr} \\hline \\multicolumn {3}{c}{$u_1 = +3.292 \\pm 0.005$, $u_2 = +0.026 \\pm 0.006$, $u_3 = +0.49$} \\\\ \\multicolumn {3}{c}{$b_1 = +2.187 \\pm 0.004$, $b_2 = -0.164 \\pm 0.005$, $b_3 = +0.25$} \\\\ \\multicolumn {3}{c}{$v_1 = +1.930 \\pm 0.003$, $v_2 = +0.010 \\pm 0.003$, $v_3 = +0.16$}\\\\ \\multicolumn {3}{c}{$r_1 = +1.936 \\pm 0.004$, $r_2 = -0.012 \\pm 0.009$, $r_3 = +0.09$}\\\\ \\multicolumn {3}{c}{$i_1 = +2.786 \\pm 0.004$, $i_2 = +0.015 \\pm 0.004$, $i_3 = +0.08$} \\\\ \\hline \\end{tabular} \\end{table}  \\subsection{Standard Stars} Three Landolt (1992) areas (TPhoenix, Rubin~149, and PG~0231+006) were observed several times each night to tie instrumental magnitudes to the standard system.  All nights, except the last one, were stable and photometric with seeing between 0.8 and 1.2 arcsec. The last night was non-photometric with bad seeing conditions (see Table~1). Photometry from this last night was tied to the other nights through the comparison of stars in common.\\\\ Since the photometric solutions were identical, all the standard star measurements were used together in obtaining a single photometric solution for the entire run.  This resulted in calibration coefficients derived using about 200 standard stars. Photometric solutions have been calculated following Patat \\& Carraro (2001). Fig.~4 shows the run of magnitude differences (standard versus instrumental) for the whole standard set. Notice that the colour baseline is sufficiently broad. On the right, the {\\it rms} of the fit is shown for each colour.  The calibration equations read:   \\begin{center} \\begin{tabular}{lc} $u = U + u_1 + u_2 (U-B) + u_3 X$         & (1) \\\\ $b = B + b_1 + b_2 (B-V) + b_3 X$         & (2) \\\\ $v = V + v_1 + v_2 (B-V) + v_3 X$         & (3) \\\\ $r = R + r_1 + v_2 (V-R) + r_3 X$         & (4) \\\\ $i = I + i_1 + i_2 (V-I) + i_3 X$         & (5), \\\\ \\end{tabular} \\end{center}  \\noindent where $UBVRI$ are standard magnitudes, $ubvri$ are the instrumental magnitudes, $X$ is the airmass, and the derived coefficients are presented in Table~2.  We adopted the extinction coefficients typical of the site (Carraro et al. 2005b).    \\subsection{Photometry extraction} The covered areas are shown in Figs.~1 to 3.  Data have been reduced using IRAF\\footnote{IRAF is distributed by NOAO, which is operated by AURA under cooperative agreement with the NSF.}  packages CCDRED and DAOPHOT. Photometry was done employing the point spread function (PSF) fitting method (Stetson 1987).  Particular care has been put in defining the PSF model. A variable PSF was adopted due to PSF variations across the CCD. In general, up to 40 bright stars have been selected for defining the PSF model.  Aperture corrections were estimated from samples of bright PSF stars (typically 15), and then applied to all the stars. The corrections amounted to 0.250-0.315, 0.280-0.300, 0.200-0.280, 0.190-0.270, and 0.210-0.280 mag for the $UBVRI$ filters, respectively, over the entire run. Photometric completeness was estimated following Baume et al. (2006) and was determined to be higher than 50$\\%$ at V $\\sim$ 20. mag. ",
        "conclusions": "We have presented a photometric analysis in the $UBVRI$ filters of 3 wide field pointings close to the center of the Canis Major over-density. The goal was to study the stellar populations in this region of the Milky Way, where a putative dwarf galaxy in the act of being cannibalized by the Milky Way, is claimed to exist. The analysis presented in this paper followed a different strategy from previous investigations of the CMa over-density.  Instead of studying very large fields in two filters, we have concentrated on a smaller area, but observed in several filters. This approach, frequently employed in the study of star clusters, allowed us to construct several CMDs and the classical (B-V) vs (U-B) TCD, which together constitute a very powerful tool for detecting young stellar populations.  As in our previous studies of stellar fields in the TGQ we found evidence of a diffuse young stellar population, as expected from the presence of the Local and Outer Galactic spiral arms (Carraro et al. 2005, Moitinho et al. 2006, V\\'azquez et al. 2008). Once again, no indication has been found of an ongoing accretion event in this direction of the Galactic disk.  In addition, the estimated ranges of distance, age and metallicity of the older metal poor population are consistent with those of thick disk stars at different distances from the Sun. These findings, together with the results of previous papers by us and other authors (Momany et al 2004, 2007), significantly weaken the proposed scenario of a dwarf galaxy in CMa being cannibalized by the Milky Way.  Instead, all the observational evidence fits our current knowledge of the Galactic disk. The TGQ is indeed a complicated region due to the warp and the existence of the Local Arm. Only the detailed multicolour analysis we have been conducting in the last few years could provide a clear picture of the structure of the outer disk in the TGQ.   \\begin{figure} \\includegraphics[width=\\columnwidth]{MN1806fig12.eps} \\caption{V vs V-I CMD of the CMa over-density. Two isochrones (dashed lines, red when printed in color) have been superimposed to guide the eye, illustrate the position of the TO, and provide a rough estimate of the mean age of the population.} \\end{figure}"
    },
    "0801/0801.0424_arXiv.txt": {
        "abstract": "In the local Universe, high-power radio galaxies live in lower density environments than low-luminosity radio galaxies. If this trend continues to higher redshifts, powerful radio galaxies would serve as efficient probes of moderate redshift groups and poor clusters. Photometric studies of radio galaxies at $0.3 \\lesssim z \\lesssim 0.5$ suggest that the radio luminosity-environment correlation disappears at moderate redshifts, though this could be the result of foreground/background contamination affecting the photometric measures of environment. We have obtained multi-object spectroscopy of in the fields of 14 lower luminosity ($L_{\\rm{1.4~GHz}} < 4\\times10^{24}~\\rm{W~Hz^{-1}}$) and higher luminosity ($L_{\\rm{1.4~GHz}} > 1.2\\times10^{25}~\\rm{W~Hz^{-1}}$) radio galaxies at $z \\approx 0.3$ to spectroscopically investigate the link between the environment and the radio luminosity of radio galaxies at moderate redshifts. Our results support the photometric analyses; there does not appear to be a correlation between the luminosity of a radio galaxy and its environment at moderate redshifts. Hence, radio galaxies are not efficient signposts for group environments at moderate redshifts. ",
        "introduction": "Photometric studies of radio galaxies have shown that Fanaroff-Riley type I \\citep[FR I;][]{fr} (low luminosity and edge-dimmed radio lobes) and FR II (high luminosity and edge-brightened lobes) radio galaxies exist in different environments in the local Universe; FR I galaxies reside in rich clusters while FR II galaxies tend to inhabit poor clusters or rich groups \\citep[e.g.,][]{heckman,prestage,lilly}. Studies of higher redshift radio galaxies indicate that this difference disappears at moderate redshifts ($z \\sim 0.5$), where both FR types are associated with cluster-like environments \\citep[e.g.,][]{hill,bahcall,allington}. In contrast, \\citet{zirbel} suggests that, although the environments of both types of galaxies are more dense at moderate redshifts than low redshifts, the environments of FR II galaxies remain less dense than FR I sources out to $z \\sim 0.5$. Only the low redshift trend of FR I environments being more dense than the environments of FR II galaxies has been confirmed with X-ray observations \\citep{miller99} and optical spectroscopy \\citep{miller}. It is quite difficult to detect and characterize group-like over-densities at moderate redshifts without spectroscopy. Photometric properties are difficult to use because groups tend not to have substantial bright early-type populations \\citep[e.g.,][]{wilman} that form a red sequence, and foreground and background interlopers are common. Therefore, field spectroscopy is necessary to determine whether FR II galaxies may be useful signposts to identify moderate redshift galaxy groups.  These galaxy groups and poor clusters are excellent laboratories for studying galaxy evolution. Over half of all galaxies in the local Universe are members of groups \\citep{huchra,tully}, and these groups have moderate densities and low velocity dispersions that make them the preferred environments for dynamic galaxy evolution \\citep[e.g.,][]{zabludoff,carlberg,aarseth,barnes} compared to clusters or the field \\citep[e.g.,][]{aceves,lin,conselice,patton}. Though cluster cores also promote some galaxy processing and evolution, the high velocity dispersions associated with these cores decrease the effective interaction cross-section between galaxies and limit the evolutionary effects of galaxy-galaxy interactions \\citep[e.g.,][]{bower}. Therefore galaxy groups provide excellent conditions for studying interactions that affect the star formation rate of galaxies and AGN ignition \\citep[e.g.,][]{mihos,iono,bahcallj,best}.  The most significant obstacle to studying galaxy evolution in group environments is identifying non-local groups. The Sloan Digital Sky Survey \\cite[SDSS;][]{york} and the 2 Degree Field Galaxy Redshift Survey \\cite[2dFGRS;][]{colless} have extended local group samples \\citep[e.g.,][]{ramella,zabludoff} to redshifts $z \\sim 0.2$ \\citep[e.g.,][]{weinmann,collister,padilla}. Blind redshift surveys at higher redshifts have successfully identified galaxy groups to $z \\sim 0.7$ \\citep[e.g.,][]{wilman,gerke}, but the substantial amount of telescope time invested in these surveys makes it desirable to find alternative methods of identifying moderate redshift groups. \\citet{mulchaey} use deep X-ray observations to select groups at moderate redshifts, though this method may preferentially select the most massive groups. Gravitational lenses have also been used to identify moderate redshift groups \\citep{momcheva,williams,augera,augerb,fassnacht}, though the scarcity of lenses limits the utility of this method. Radio galaxies have been used extensively to find high redshift clusters \\citep[e.g.,][]{deltorn,blantona,blantonb} and protoclusters \\citep[e.g.,][]{miley,venemans}, but the FR I/FR II environment dichotomy suggests that they may also be useful for identifying less rich galaxy groups \\citep[e.g.,][]{allington}.  In this paper we present spectroscopic observations of the fields of a sample of 14 moderate redshift radio galaxies. We investigate the relationship between the radio luminosity and spectroscopic group properties, and we analyze the effectiveness of using high luminosity radio galaxies to preferentially locate moderate redshift groups. Throughout this paper we assume a radio spectral index of $\\alpha = 0.5$, where $S_{\\nu} \\propto \\nu^{-\\alpha}$, and we use a $\\Lambda$CDM cosmology with $\\Omega_{\\Lambda} = 0.73$ and $\\Omega_M = 0.27$. ",
        "conclusions": "\\citet{geach} have examined the environments of a set of four radio galaxies at redshifts $z \\sim 0.35$ and $z \\approx 0.65$ with spectroscopy and X-ray observations. This sample includes three galaxies that we would classify as low-power and one higher-power radio galaxy, and the sample spans a range of environments. We note, however, that none of the galaxies in our sample or the \\citet{geach} sample are truly isolated; all have at least two spectroscopic companions. This could be due to selection effects. \\citet{geach} investigated radio galaxies that appeared to have a red sequence, while our sample required the galaxy to be in the SDSS spectroscopic sample. The SDSS preferentially observes more luminous galaxies at higher redshifts, and these galaxies tend to exist in over-dense environments. However, photometric surveys of purely radio-selected galaxies have also suggested that nearly all radio-galaxies reside in groups \\citep[e.g.,][]{allington}, and we therefore do not expect this bias to be significant.  Our results spectroscopically support the results from photometric surveys of radio galaxies; low-luminosity and high-luminosity sources tend to exist in the same environments at moderate redshifts. The dichotomy in environments at low redshifts is likely due to the density of the intergalactic medium (IGM); high luminosity galaxies tend to have large extended lobes that do not form if the IGM is too dense \\citep[e.g.,][]{prestage}. The absence of this dichotomy at higher redshift is probably results from cosmological evolution of the density of the IGM \\citep[e.g.,][]{ettori}.  There is a substantial amount of scatter in the properties of the radio galaxy environments, though the typical radio galaxy is in a group with approximately 12 members and a velocity dispersion of $\\sim 500~\\rm{km~s^{-1}}$. The systems 222821.192$+$011412.29 and 013352.730$+$011343.63 are found to have only 3 companions in addition to the radio galaxy. These systems may not belong to bound groups but rather may be embedded in filaments; the large velocity spread for the low-luminosity system 222821.192$+$011412.29 further supports this conclusion. We therefore conclude that at moderate redshifts, radio galaxies reside in all types of environments, from filaments to clusters, without regard to the radio luminosity. \\citet{belsole} reach a similar conclusion from a study of the X-ray environments of higher redshift radio galaxies and quasars. Thus, while radio galaxies are shown to trace large-scale-structure, they are not useful for targeting environments with specific properties (e.g., lower velocity dispersion or less rich environments) at higher redshifts."
    },
    "0801/0801.1034_arXiv.txt": {
        "abstract": "{Using extremely deep (rms $\\sim$3.3$\\mu$Jy/bm) 1.4\\,GHz sub-arcsecond resolution MERLIN$+$VLA radio observations of a 8\\farms5$\\times$8\\farms5 field centred upon the Hubble Deep Field North, in conjunction with {\\it Spitzer} 24\\,$\\mu$m data we present an investigation of the radio-MIR correlation at very low flux densities.  By stacking individual sources within these data we are able to extend the MIR-radio correlation to the extremely faint ($\\sim$microJy and even sub-microJy) radio source population. Tentatively we demonstrate a small deviation from the correlation for the faintest MIR sources. We suggest that this small observed change in the gradient of the correlation is the result of a suppression of the MIR emission in faint star-forming galaxies.  This deviation potentially has significant implications for using either the MIR or non-thermal radio emission as a star-formation tracer at low luminosities. }  \\FullConference{From planets to dark energy: the modern radio universe\\\\ October 1-5 2007\\\\ University of Manchester, Manchester, UK}  \\begin{document} ",
        "introduction": "Radio and infrared emission from galaxies in both the nearby and distant Universe is thought to arise from processes related to star-formation, hence resulting in the correlation between these two observing bands. The infrared emission is produced from dust heated by photons from young stars and the radio emission predominately arises from synchrotron radiation produced by the acceleration of charged particles from supernovae explosions. It has however recently been suggested that at low flux density and luminosities there may be some deviation from the tight well-known radio-IR correlation seen for brighter galaxies \\cite{bell03,boyle07}.  \\cite{bell03} argue that while the IR emission from luminous galaxies will trace the majority of the star-formation in these sources, in low luminosity galaxies the IR emission will be less luminous than expected considering the rate of star-formation within the source (i.e. the IR emission will not fully trace the star-formation). In this scenario the reduced efficiency of IR production relative to the source star-formation rate (SFR) would be the result of inherently lower dust opacities in lower luminosity sources and consequently less efficient reprocessing of UV photons from hot young stars into IR emission. The simple consequence of this is that at lower luminosities the near linear radio-IR correlation L$_{\\rm radio}\\propto$L$_{\\rm IR}^{\\!\\gamma}$, with $\\gamma >1$ (e.g. \\cite{cox88,price92}) will be deviated from. {\\it Of course such an assertion is dependent upon the radio emission providing a reliable tracer of star-formation at low luminosities which may be equally invalid.}  Recently \\cite{boyle07} have presented a statistical analysis of Australia Telescope Compact Array (ATCA) 20\\,cm observations of the 24\\,\\mum\\ sources within  the \\spitzer\\ Wide Field Survey (SWIRE). In this work \\cite{boyle07} have co-added sensitive (rms$\\sim$30\\,$\\mu$Jy) radio data at the locations of several thousand 24\\,\\mum\\ sources. Using this method they have statistically detected the microJy radio counterparts of faint 24\\,\\mum\\ sources. At low flux densities (S$_{\\rm 24\\, \\mu m }=100\\,\\mu$Jy) they confirm the IR-radio correlation but find it to have a lower coefficient (S$_{\\rm 1.4\\,GHz}$\\,=\\,0.039\\,S$_{\\rm 24\\,\\mu m}$) than had previously been reported at higher flux densities. This coefficient is significantly different from results previously derived from detections of individual objects (e.g. \\cite{appleton04}) and is speculated by \\cite{boyle07} to be the result of a change in the slope of the radio-IR correlation at low flux densities.   In this work (which is descirbed in more detail in \\cite{beswick06, beswick08}) we have utilised very deep, high resolution 20\\,cm observations of the Hubble Deep Field North and surrounding area made using MERLIN and the VLA \\cite{muxlow05} in combination with publicly available 24\\,\\mum\\ \\spitzer\\ source catalogues from GOODS to study the MIR-Radio correlation for microjansky radio sources. This study extends the flux density limits of the radio-IR correlation by more than an order of magnitude for individual sources and overlaps the flux density regime studied using statistical stacking methods by other authors.  \\begin{figure*} \\begin{center}$ \\begin{array}{cc} \\includegraphics[width=8cm]{Beswick-fig3a.ps} & \\includegraphics[width=8cm]{Beswick-fig3b.ps} \\end{array}$ \\caption{Mean (left) and median (right) images of the 1.4\\,GHz radio emission for all sources within the six faintest 24\\,\\mum\\ flux density logarithmic bins plotted in Fig.\\,1 in descending flux density order from top-left to bottom-right. The range of 24\\,\\mum\\ flux density over which each image has been stacked is shown at the top of individual panels. Each image is contoured with levels of $-$2, $-$1.414, $-$1,1, 1.414, 2, 2.828, 4, 5.657, 8, 11.31, 16, 22.63 and 32 times 3$\\times$(3.3/$\\sqrt {\\rm N}$)\\,$\\mu$Jy\\,bm$^{-1}$, where N equals the number of 24\\,\\mum\\ source positions averaged in the map. The peak flux density, lowest plotted contour and number of IR sources which have been averaged over (N) is listed at the bottom of each image panel.} \\label{maps} \\end{center} \\end{figure*} ",
        "conclusions": "Using one of the deepest high-resolution 1.4\\,GHz observations made to date, in conjunction with deep 24\\,\\mum\\ \\spitzer\\ source catalogues from GOODS, we have investigated the microJy radio counterparts of faint MIR sources. These observations confirm that the microJy radio source population follow the MIR-radio correlation and extend this correlation by several orders of magnitude to very low flux densities and luminosities, and out to moderate redshifts. This extension of the MIR-radio correlation confirms that the majority of these extremely faint radio and 24\\,\\mum\\ sources are predominantly powered by star-formation with little AGN contamination.  Statistically stacking the radio emission from many tens of faint 24\\,\\mum\\ sources has been used to characterise the size and nature of the radio emission from very faint IR galaxies well below the nominal radio sensitivity of these data. Using these methods the MIR-radio correlation has been further extended and a tentative deviation in this correlation at very low 24\\,\\mum\\ flux densities has been identified."
    },
    "0801/0801.3177_arXiv.txt": {
        "abstract": "{ The HST/ACS Survey  of Galactic globular clusters (GGCs) is a HST Treasury project  aimed at  obtaining high  precision photometry  in a large  sample  of  globular  clusters.   The  homogeneous  photometric catalogs  that  has been  obtained  from  these  data by  Anderson  et al.  (2008)  represents a  golden  mine  for  a lot  of  astrophysical studies.  In  this paper we  used the catalog to  analyse the  properties of MS-MS  binary systems  from  a sample  of  fifty GGCs.   We measured  the fraction  of binaries  (divided  in different  groups), studied  their radial distribution  and constrained the mass  ratio distribution.  We investigated possible  relations between the fraction  of binaries and the main parameters of their host GGCs.  We found a significant anti-correlation between the binary fraction in a cluster and its absolute luminosity (mass). ",
        "introduction": "Knowledge of the binary frequency in Globular Clusters (GCs) is of foundamental importance for a lot of astrophisical studies.  Binaries  play an  important  role  in the  dynamical  evolution of  a clusters. Interactions with hard binaries pump kinetic energy into the cluster core, slowing the core  collaps and, eventually, causing the core to reexpand, if the number of binaries is large enough. In general, binaries are a foundamental ingredient in any dynamical evolution model of a GC.  Exotic stellar  objects, like Blue  Stragglers, cataclismic variables, millisecond  pulsars and  low  mass  X ray  binaries  are believed  to represent   evolutionary   stages   of   close  binary   system.   The determination of  the fraction of  binaries plays a  foundamental step towards the understanding of the evolution of these peculiar objects. Furthermore,   binary  stars  introduce   systematic  errors   in  the determination  of the  main sequence  (MS) fiducial  line and  move it toward red colors with respect to its correct position.  Finally, a correct determination  of the mass and luminosity functions requires a correct measure of the fraction of binaries.  Up to now, three main techniques have been used to measure the fraction ob binaries in GGCs (Hut et al. 1992).  The first  one identifies binaries by measuring  their radial velocity variation (eg. Latham 1996). This  method relies with the detection of each individual binary system but, due to the limits in sensibility of spectroscopy, these studies are possible only for the brightest GGCs stars.  The second tecnique  is based on the search  for photometric variables (eg. Mateo  1996). As well  as the previous  one, it is able  to infer specific properties of each binary system (like the measure of orbital period, mass  ratio, orbital inclination). Unfortunatly,  it is biased towards  binaries with  short  periods and  large orbital  inclination. Moreover these tecniques  have a low discovery efficency  and are very expensive in terms of telescope time because it is necessary to repeat measures in time.  A thirth  approach, that  is based  on the analysis  of the  number of stars  located  on the  red  side  of the main sequence (MS)  ridge line (MSRL)  may represent a more efficient method  to measure the fraction of binaries in a cluster for several reasons: \\begin{itemize} \\item{availability of a large number (thousands) of stars makes it a statistically roboust method;} \\item{it is cheep in terms of observational time: two filters are enough for detecting binaries and repeated measuremets are not needed.} \\item{it is sensitive to binaries with any orbital period and inclination} \\end{itemize} This approach have been used by other groups ( see Sollima et al. 2007 and references within) to study the population of binaries in GGCs. The relative small number of clusters that have been analized is consequence of the intrinsic difficulties of the method: \\begin{itemize} \\item{high photometric quality is required;} \\item{in some cases, the differential reddening spreads the MS and makes it more difficult to isolate the binary sequence;} \\item{an accurate analysis of photometric errors as well as a correct estimate of field contamination are necessary to disantangle real binaries from bad photometry and field stars. } \\end{itemize} In this paper, we analyse the catalogs obtained by Anderson et al. (2008) from HST ACS/WFC data. We exploited both the homegeneity of this dataset, and the high photometric accuracy of the measures to derive the fraction of binaries in the central regions of fifty GGCs. ",
        "conclusions": ""
    },
    "0801/0801.4561_arXiv.txt": {
        "abstract": "Planets embedded in optically thick passive accretion disks are expected to produce perturbations in the density and temperature structure of the disk.  We calculate the magnitudes of these perturbations for a range of planet masses and distances. The model predicts the formation of a shadow at the position of the planet paired with a brightening just beyond the shadow. We improve on previous work on the subject by self-consistently calculating the temperature and density structures under the assumption of hydrostatic equilibrium and taking the full three-dimensional shape of the disk into account rather than assuming a plane-parallel disk. While the excursion in temperatures is less than in previous models, the spatial size of the perturbation is larger. We demonstrate that a self-consistent calculation of the density and temperature structure of the disk has a large effect on the disk model. In addition, the temperature structure in the disk is highly sensitive to the angle of incidence of stellar irradition at the surface, so accurately calculating the shape of the disk surface is crucial for modeling the thermal structure of the disk. ",
        "introduction": "Giant planets forming by core accretion need to have cores of 10-20 M$_{\\earth}$ to be massive enough to accrete a gaseous envelope \\citep{2005Hubickyj_etal}. This predicts that sizeable planet embryos form before circumstellar gas disks dissipate. These disks are typically modeled as passive accretion disks, optically thick and gas-dominated.  Their temperature structure is strongly dependent on heating from stellar irradiation \\citep{CG, calvet, vertstruct, dalessio2, dalessio3}. In particular, the disk temperatures depend strongly on the angle of incidence of the stellar irradiation on the disk surface.  While a substantial amount of work on numerical hydrodynamic simulations of planets embedded in disks has been carried out, calculating the radiative transfer of stellar irradiation in conjunction with these simulations is prohibitively difficult. Simulations by \\citet{bate} illustrate the effect of hydrodynamics on the disk structure around a small embedded planet, but use an isothermal equation of state. The effects of MHD turbulence have been studied by \\citet{2004MNRAS.350..829P} and \\citet{2007Oishi_etal}, but again, assuming a simple and unrealistic equation of state. \\citet{2006KlahrKley} and \\citet{2006PaardekooperMellema,2008PaardekooperMellema} have made an effort to include radiative transfer as well as hydrodynamics in their calculations, but do not include the effects of stellar irradiaton on their models.  The work presented in this paper does not include hydrodynamics, but rather focuses on the effects of stellar irradiation, which is a particularly important source of heating in the photospheres of disks.  Since it is the photospheres of optically thick disks that are observed, the effects of stellar irradiation at the surface of disks are an important consideration in predicting observations of protoplanetary disks.  A limitation of the hydrodynamic simulations described above is that in order to adequately model the densest regions of the gas they are necessarily limited in spatial range above the midplane.  The photosphere and surface of a disk are often several scale heights above the midplane.  The models presented here are complementary to detailed hydrodynamic simulations for this reason.  We are particularly interested in determining if the growing cores of giant planets produce effects that are observable.  If these effects are observed, they would affirm the core accretion model as the paradigm for giant planet formation. While a fully-formed Jupiter-mass planet would produce a larger feature in a disk, it would reveal little about how it formed.  Previous work on small planets embedded in optically thick gas disks indicates that sub-Jovian mass planet cores can perturb the disk enough to alter the temperature structure of the disk in the immediate vicinity of the planet, with consequences for further evolution of the planet \\citep{paper1, paper2, HJCmigration}. In \\citet{paper1,paper2} (henceforth Paper I and Paper II, respectively), the planet is predicted to gravitationally compress the disk in the vertical direction, creating a shadow paired with a bright spot, leading to temperature variations.  However, those models were limited by being plane-parallel, despite the sizes of the simulation boxes used.  In addition, while the density perturbations were calculated under hydrostatic equilibrium, they were not calculated self-consistently with the temperature perturbations.  In this paper, we improve upon the model presented in Papers I and II by iteratively calculating the density and temperature structure of the disk for self-consistency and eliminating the assumption of a locally plane parallel model. The paper is organized as follows: in \\S \\ref{model} we describe the basic model and the improvements we have made, in \\S \\ref{stellarheating} we elaborate in detail on the method we use for modeling radiative transfer, in \\S \\ref{results} we describe our results for a different planet masses and distances, in \\S \\ref{discussion} we discuss our results in comparison to previous models, and in \\S \\ref{conclusion} we present our conclusions. ",
        "conclusions": "This paper presents a model for calculating the density and temperature perturbations imposed on a protoplanetary disk by an embedded protoplanet.  The basic radiative transfer model is adopted from Papers I and II, but a number of improvements have been made on that work such as density-temperature self-consistency and eliminating the assumption of a plane-parallel system.  We have shown that self-consistently calculating the temperature and density significantly increases the effect of planet perturbations by means of postive feedback, where shadowed regions cool and contract and brightened regions heat and expand.  This demonstrates the importance of self-consistency when calculating disk structure with radiative transfer.  While it has already been acknowledged that stellar irradiation heating is important in setting overall disk structure \\citep[e.g.][]{CG,thermstab,vertstruct,2004DullemondDominik}, this work shows that it is important for considering local perturbations on the disk.  Another important result of this paper is that the temperature structure of the disk is extremely sensitive to the angle of incidence of stellar irradiation at the surface.  The precise determination of the surface of the disk is critical to accurately calculating the temperature structure of the disk.  In previous work we had assumed that the surface of constant density was a sufficiently good approximatation for the surface, but in this work we show that that overestimated the temperature perturbation to the disk.  We have omitted some important physics in this calculation.  We do not include heating from accretion onto the planet.  We consider the embedded planet to act simply as a gravitational point mass, and focus only on the vertical component of gravity.  We do not account for non-linearities such as spiral density waves.  All these effects are more adequately addressed using a three-dimensional hydrodynamic simulation of a planet embedded in a disk of gas.  However, since simulations of this sort focus on the bulk flow of gas at the midplane, they typically have insufficient resolution above the midplane to accurately calculate the surface of the disk \\citep[e.g.][]{bate,2004MNRAS.350..829P,2006KlahrKley,2007Oishi_etal}. Calculation of radiative transfer, even without iterating for self-consistency, is very computationally intensive.  To do this iteratively and coupled with three-dimensional hydrodynamics is challenging, but is the next logical step to improving the accuracy of our results."
    },
    "0801/0801.1491_arXiv.txt": {
        "abstract": "In this paper, we present results from the complete set of cosmic microwave background (CMB) radiation temperature anisotropy observations made with the Arcminute Cosmology Bolometer Array Receiver (ACBAR) operating at $150\\,$GHz.  We include new data from the final 2005 observing season, expanding the number of detector-hours by 210\\% and the sky coverage by 490\\% over that used for the previous ACBAR release.  As a result, the band-power uncertainties have been reduced by more than a factor of two on angular scales encompassing the third to fifth acoustic peaks as well as the damping tail of the CMB power spectrum.  The calibration uncertainty has been reduced from 6\\% to 2.1\\% in temperature through a direct comparison of the CMB anisotropy measured by ACBAR with that of the dipole-calibrated WMAP5 experiment.  The measured power spectrum is consistent with a spatially flat, $\\Lambda$CDM cosmological model. We include the effects of weak lensing in the power spectrum model computations and find that this significantly improves the fits of the models to the combined ACBAR+WMAP5 power spectrum. The preferred strength of the lensing is consistent with theoretical expectations. On fine angular scales, there is weak evidence ($1.1\\sigma$) for excess power above the level expected from primary anisotropies. We expect any excess power to be dominated by the combination of emission from dusty protogalaxies and the Sunyaev-Zel'dovich effect (SZE). However, the excess observed by ACBAR is significantly smaller than the excess power at $\\ell >2000$ reported by the CBI experiment operating at $30\\,$GHz. Therefore, while it is unlikely that the CBI excess has a primordial origin; the combined ACBAR and CBI results are consistent with the source of the CBI excess being either the SZE or radio source contamination. ",
        "introduction": "\\label{sec:intro}  Observations of the cosmic microwave background (CMB) are among the most powerful and important tests of cosmological theory. Measurements of the angular power spectrum of CMB temperature anisotropies on angular scales $> 10^\\prime$ - corresponding to multipoles $\\ell \\lesssim 1000$ - \\citep{spergel06} in conjunction with other cosmological probes \\citep{burles01,cole05,tegmark06,riess07} have produced compelling evidence for the $\\Lambda$CDM cosmological model. At higher multipoles, measurements probe the Silk damping tail of the power spectrum and provide an independent check of the cosmological model.  At smaller angular scales, the primary CMB anisotropies originating at redshift z = 1100 are exponentially damped by photon diffusion.  This effect, known as Silk damping, makes secondary anisotropies - those induced along the line of sight at lower redshift - increasingly important at higher $\\ell$.  At 150$\\,$GHz, for example, the Sunyaev-Zel'dovich effect (SZE) is expected to be brighter than the primary CMB anisotropy at $\\ell \\gtrsim 2500$.  The amplitude of the SZE depends sensitively on the amplitude of the matter perturbations, scaling as $\\sigma_8^7$.  Measurements of the CMB power spectrum with sufficient sensitivity on arcminute scales not only extend tests of the $\\Lambda$CDM model's ability to accurately predict the features in the power spectrum of primary CMB anisotropy, but also probe the epoch of cluster formation and provide an independent measure of $\\sigma_8$.  In this paper, we present the complete results of observations of CMB temperature anisotropies at 150 GHz with 5$^\\prime$ resolution from the Arcminute Cosmology Bolometer Array Receiver (ACBAR) experiment at the South Pole station.  Previous measurements of the CMB power spectrum by ACBAR have been presented in \\cite{kuo04} (hereafter K04) and \\cite{kuo07} (hereafter K07).  In addition, the angular power spectrum on these angular scales has been measured at 30 GHz by CBI   \\citep{readhead04}, VSA \\citep{dickinson04}, and BIMA \\citep{dawson06}, and at 100 and 150 GHz by QuAD \\citep{quad07}.     To date, measurements at angular scales $<10^\\prime$  have been consistent with predictions of the primary anisotropy based on measurements at larger angular scales, with one exception.  Both CBI~\\citep{mason03,bond05} and BIMA~\\citep{dawson06} observe excess power for $\\ell > 2000$ at 30 GHz compared to the predictions of the $\\Lambda$CDM  model. This excess can be explained by the SZE if $\\sigma_8 \\approx 1$, but this value is in tension with the best-fit WMAP5 value of  $\\sigma_8 \\approx 0.8$. In K07, we found that while the frequency dependence of the excess is consistent with the SZE, the ACBAR and CBI data could not be used to rule out radio source contamination or systematic errors as the source of the CBI excess. Careful measurements over a broad range of frequencies and angular scales are needed to provide a definitive answer.   Current estimates of the primordial power spectrum are consistent with the predictions of slow-roll inflation for a nearly scale-invariant spectrum which may also include a small running of the spectral index. Sparked by the modest evidence for negative running in the WMAP first-year data, a number of authors have investigated how existing data sets limit the allowed inflationary scenarios \\citep{peiris03,mukherjee03,bridle03,leach03}. Small-scale data extend the range over which the primordial power spectrum is measured and can potentially yield information about the mechanism of inflation.  This is the third and final ACBAR power spectrum release.  The first release in K04 analyzed two fields from the 2001 and 2002 seasons with a conservative field differencing algorithm.  The second ACBAR power spectrum, presented by K07, added two more fields from the 2002 season and implemented an improved, undifferenced Lead-Main-Trail (no-LMT) analysis of the dataset.  The results presented here improve on the previous work in two ways.  First, we include seven additional fields observed in the 2005 Austral winter.  These fields double the total number of detector hours and substantially improve the precision of the band-power estimates.  In particular, the new fields were selected to dramatically expand ACBAR's sky coverage in order to reduce the cosmic variance contribution to the uncertainty and to improve the multipole resolution on angular scales below $\\ell \\lesssim 1800$.  This angular range covers the third to fifth acoustic peaks, making it especially interesting for constraining cosmological models.  Second, we implement a new temperature calibration based on a comparison of CMB fluctuations as measured by ACBAR and the WMAP satellite~\\citep{hinshaw08}. This improved calibration tightens constraints on cosmological models found from the combination of high-$\\ell$ ACBAR band-powers with low-$\\ell$ results from other experiments.    This paper is organized as follows. In \\S~\\ref{sec:instrument} we review the ACBAR instrument and the CMB observation program. The analysis algorithm is explained in \\S~\\ref{sec:analysis}. Section \\S~\\ref{sec:calib} is an overview of the calibration; the details of cross-calibration between WMAP5 and ACBAR are discussed in Appendix A. Information on ACBAR's beams can be found in \\S~\\ref{sec:beam}. Systematic tests and foreground contamination are discussed in \\S~\\ref{sec:sys}. We present the band-power results in \\S~\\ref{sec:results}, including a discussion of the scientific interpretation. The ACBAR band-powers are combined with the results of other experiments to place constraints on the parameters of cosmological models in \\S~\\ref{sec:parameters}. In \\S~\\ref{sec:conclusion}, we summarize the main results of this paper. ",
        "conclusions": ""
    },
    "0801/0801.3341_arXiv.txt": {
        "abstract": "Our \\xmm\\ spectrum of the giant, high-redshift ($z=1.88$) radio galaxy \\qc\\ shows that it contains one of the most powerful, high-redshift, Compton-thick quasars known.  Its spectrum is very hard above 2\\kev.  The steep \\xmm\\ spectrum below that energy is shown to be due to extended emission from the radio bridge using \\chandra\\ data. The nucleus of \\qc\\ has a column density of $3.5^{+1.4}_{-0.4}\\times 10^{24}$\\pcmsq\\ and absorption-corrected X-ray luminosity of $1.7^{+0.9}_{-0.1}\\times 10^{45}$\\ergps\\ in the $2-10$\\kev\\ band.  A lower redshift active galaxy in the same field, SDSS J090808.36$+$394313.6, may also be Compton-thick. ",
        "introduction": "The powerful, classical double \\frii\\ \\citep{FR} radio galaxy \\qc\\ is unusual as it is one of the highest redshift ($z=1.88$) giant radio galaxies known, spanning $111$\\as\\ on the sky \\citep{lawgreen95} corresponding to a projected size of $945$\\kpc\\ in the cosmology assumed in this paper. Radio galaxies have long been associated with obscured, (optically defined) Type 2 quasars, because they are considered to be the same objects as radio-loud quasars seen through an obscuring torus (e.g. \\citealt{barthel89} and \\citealt{antonucci93}).  X-ray observations are particularly efficient at detecting obscured AGN as they provide the observer with several diagnostics, such as the photoelectric cut-off, the iron line equivalent width (EW) and the X-ray to optical or [O III] $\\lambda 5007$ ratio (e.g. \\citealt{bassani99} or, for a review, \\citealt{comastri04}).  Sources with an obscuring column greater than $1.5 \\times 10^{24}$\\pcmsq, at which the Thomson scattering optical depth reaches unity, are known as Compton-thick AGN. It is non-trivial to calculate the column density of such sources as there is little (see \\citealt{wilmanfabian99}) direct emission below 10\\kev\\ meaning that analysis relies on interpreting the reflected spectrum at lower X-ray energies. \\suzaku\\ is starting to provide a window into the X-ray regime above 10\\kev\\ enabling direct detection of nuclear emission in the brightest Compton-thick sources such as NGC\\,5728 \\citep{comastri07}. Such sources are important since the X-ray background gives strong evidence for a significant population of Compton-thick AGN (e.g. \\citealt{comastri04}, \\citealt{worsley05}).  In this paper, we present recent \\xmm\\ observations of \\qc\\ (which is located at RA 09h08m16.9s, Dec +39d43m26s in J2000).  Throughout this paper, all errors are quoted at $1\\sigma$ unless otherwise stated and the assumed cosmology is $\\rm H_{\\rm 0} = 71$\\kmpspmpc, $\\Omega_{0}=1$ and $\\Lambda_{0} = 0.73$.  One arcsecond represents $8.518$\\kpc\\ on the plane of the sky at the redshift ($z=1.883\\pm 0.003$) of \\qc\\ and the Galactic absorption along the line-of-sight is taken to be $1.91 \\times 10^{20}$\\pcmsq\\ \\citep{dickeylockman90}. ",
        "conclusions": "\\label{sec:conclude}  Our \\xmm\\ spectrum of \\qc, which is one of the highest redshift giant radio galaxies known, shows it to contain a Compton-thick quasar. \\qc\\ is one of very few such sources currently known and one of even fewer at high redshift.  It is likely to be a high-Eddington source with a bolometric luminosity of $\\sim 10^{47}$\\ergps, making it one of the most powerful sources known."
    },
    "0801/0801.0204_arXiv.txt": {
        "abstract": "%  We are conducting a deep near-infrared (NIR) imaging survey of young embedded clusters in the extreme outer Galaxy (hereafter EOG), at the Galactic radius ($R_{\\rm g}$) of more than 18 kpc. The EOG is an excellent laboratory to study the nature of the IMF in a low-metallicity environment with a great advantage of the proximity compared to nearby dwarf galaxies, such as LMC \\& SMC. As a first step, we obtained deep NIR images of Digel Cloud 2 clusters at $R_{\\rm g} \\simeq 19$ kpc using the Subaru 8.2-m telescope. The observed $K$-band luminosity function shows that {\\it IMF in the low metallicity environment down to $\\sim 0.1$ $M_\\odot$ is not significantly different from the typical IMFs in the field and in the nearby star clusters} as was suggested in our earlier work \\citep{Yasui2006}. ",
        "introduction": "The initial mass function (IMF) in nearby galaxies has been extensively studied with high-resolution imaging using {\\it HST} \\citep[e.g.,][]{{Sirianni2000},{Gouliermis2005}}. It is critical to reach to the substellar mass regime for the complete census of the IMF, because {\\it a characteristic peak mass, which is important to determine the shape of the IMF, is at around $\\sim$ 0.5 $M_\\odot$} \\citep{AnnualReview}. However it is difficult to reach this limit even for the nearest galaxies with active star formation, such as LMC \\& SMC, because of the lack of spatial resolution. Also, observations of embedded clusters at NIR wavelengths are more ideal for investigating the ``true'' IMF right after the cluster formation without being hampered by dust extinction \\citep{AnnualReview}.  Therefore, we started a program to study the IMF of the young embedded clusters in the EOG defined here as the region at $R_{\\rm g}$ greater than 18 kpc. The heliocentric distance to EOG is as close as 10 kpc, which is five times closer than LMC ($D \\sim 50$ kpc). The EOG has a similar environment as nearby irregular dwarf galaxies (Im) and damped Ly-$\\alpha$ systems (DLAs): low metallicity \\citep[$\\sim - 1$ dex; e.g.,][]{Rudolph2006}, low gas density ($< 0.01$ cm$^{-3}$), and no or little perturbation from the spiral arms. ",
        "conclusions": ""
    },
    "0801/0801.2037_arXiv.txt": {
        "abstract": "{We discuss the equatorial imaging benefits that arise from the addition of the 25-metre dish at Chilbolton to the e-MERLIN array. Its inclusion considerably enhances the capabilities of e-MERLIN on and below the equator. This will become particularly important in the era of ALMA and other upcoming southern hemisphere facilities. We present simulated observations of point sources in the equatorial region of the sky which is the target area for many existing sky surveys. We find that the additional baselines created by the inclusion of the Chilbolton dish favourably adjust the beam shape of e-MERLIN to a more compact and circular shape, with significantly reduced sidelobe structure. Putting aside the benefits of increased collecting area, the modified beam shape has implications for more rapidly reaching a given completeness limit for equatorial surveys.} ",
        "introduction": "The facilities at Chilbolton Observatory in Hampshire, UK, include a fully steerable 25-metre parabolic antenna which is mainly used for meteorological Doppler-polarisation radar. The use of this antenna in the MERLIN array (Thomasson, 1986) has been discussed for many years, and with the provision of a fast fibre link it would become a prime candidate for inclusion in the e-MERLIN\\footnote{\\tt{http://www.merlin.ac.uk/e-merlin/}} array.  Figure~\\ref{fig:merlin} is a version of the diagram presented on the online MERLIN user guide\\footnote{\\tt{http://www.merlin.ac.uk/user\\_guide/OnlineMUG/}} which has been modified in order to show the location of the Chilbolton antenna in addition to the locations of the existing MERLIN stations. Antenna numbers for the existing MERLIN stations have also been added. As can be seen, inclusion of the Chilbolton antenna will boost the number of long and intermediate-length baselines, as well as extending the north-south span of the array, facilitating the sampling of a much greater range of spatial frequencies. This is particularly crucial for `snapshot' mode observations, and large-scale survey programmes where on-source time may necessarily be rather brief.  In this article we demonstrate the improved $uv$-plane sampling of an e-MERLIN+Chilbolton array for both snapshot and full synthesis equatorial observations. We present simulations demonstrating the response of both arrays to an unpolarized point source, and discuss the implications of the inclusion of the Chilbolton antenna when carrying out equatorial surveys.  All observations are simulated using AIPS with standard imaging procedures. All simulations are at L-band, using 100 frequency channels of 5 MHz each to simulate a 1.3 - 1.8 GHz contiguous band, assuming a spectrally flat source. Each antenna has an efficiency of 0.8 and system temperatures in the range 20 - 33~K. The 76-metre Lovell telescope is not included in the simulated array.  \\begin{figure} \\centering \\vspace{215pt} \\special{psfile=merlin2.eps hscale=35 vscale=35 hoffset=0 voffset=0} \\caption{Map showing the current MERLIN array, with the location of the Chilbolton Observatory highlighted in grey. \\label{fig:merlin} } \\end{figure} ",
        "conclusions": "The geographical location of the Chilbolton antenna introduces baselines which complement the existing e-MERLIN array. This is particularly evident in the case of equatorial observations, where the dominant east-west layout of the existing array biases the $uv$ coverage, resulting in strong linear structure in the sidelobes of a synthesised beam which is highly eccentric.  Putting aside the increased sensitivity due to the $\\sim$10\\% increase in collecting area with the additional antenna, the synthesised beam is much more compact and circular for equatorial observations, facilitating more efficient detection of characteristic $\\mu$Jy radio sources, reducing the mapping speed by up to a factor of two (see Table~\\ref{tab:maptimes} for favourably oriented sources.  The addition of the Chilbolton antenna would significantly enhance the power of e-MERLIN. There are many experiments that would benefit significantly from this enhancement, such as high-resolution radio surveys complementing optical and infrared deep-fields such as those to be undertaken by VISTA, and high-resolution radio observations of galactic and extra-galactic radio sources targetted with ALMA. Both e-MERLIN and ALMA naturally complement each other due to their similar sub-arcsecond resolutions, despite their operating frequencies differing by a factor of $\\sim$100."
    },
    "0801/0801.3630_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{intro}  Ordinary fluids (e.g., gases and liquids) may be bounded by rigid walls which allow particle number conservation, avoiding evaporation.   Macroscopical parameters (pressure, density, and temperature) remain uniform within the box, due to its reduced dimensions.   On the other hand, astrophysical fluids (e.g., stars and galaxies) may be conceived as bounded by ``gravitational'' walls which violate particle number conservation by allowing evaporation.   The macroscopical parameters exhibit gradients because they cannot remain uniform within the ``gravitational'' box, due to its large-scale dimensions.  In particular, sufficiently extended celestial objects show at least two distinct components: core-envelope for stars, core-halo for elliptical galaxies, bulge-disk for spiral and lenticular galaxies, baryonic-nonbaryonic for virialized (matter) density perturbations, and matter-dark energy for virialized (matter + dark energy) density perturbations.   On this basis, an investigation on two-component astrophysical fluids appears useful for the comprehension and the interpretation of what is inferred from observations.   To this aim, the choice of the density profiles is of basic importance. The laws of ideal and real gases were deduced from ordinary fluids, characterized by uniform density profiles.   Accordingly, it is expected that astrophysical fluid laws are related to the specified density profiles, and different laws hold for different matter distributions.  Strictly speaking, a density profile should be deduced from the distribution function, or vice versa.   Unfortunately, the determination of the distribution function is much more difficult than in one-component systems, and only a few cases have been studied in detail at present (e.g., Ciotti, 1996, 1999). On the other hand, global properties exhibited by simple density profiles (with somewhere negative distribution function) are expected to maintain a similar trend in dealing with much more complex density profiles (with nonnegative distribution function).  As the current attempt is mainly aimed to explore global properties instead of local properties, density profiles shall be selected according to their intrinsic simplicity, regardless from the physical meaning of the distribution function. Configurations described by simple density profiles could be sufficiently close to their counterparts described by self-consistent density profiles, and related results hold to a first extent. In any case, the self-consistency of density profiles with respect to nonnegativity of the distribution function, can be checked using a specific theorem (Ciotti and Pellegrini, 1992).  To a first extent, the particle number shall be supposed to be conserved, which is equivalent to conceive the boundary of each subsystem as a perfectly reflecting surface in order to avoid evaporation.   In this view, a possible choice of macroscopical parameters is: the fractional virial potential energy, $\\phi$; the fractional truncation radius, $y$; and the fractional mass, $m$; as done in a pioneering paper with respect to uniform density profiles i.e. homogeneous configurations (Caimmi and Secco, 1990, hereafter quoted as CS90).   Accordingly, each subsystem is supposed to be virialized, in the sense that the virial equations are satisfied by averaging over a sufficiently long time, and particles move within a bounded region (e.g., Landau and Lifchitz, 1966, Chap.\\,II, \\S10; Caimmi, 2007a). Then virial equilibrium is a necessary (but not sufficient) condition for dynamical or hydrostatic equilibrium, which, on the other hand, does not imply pressureless configurations, as the stress tensor is related to the kinetic-energy tensor (e.g., Binney and Tremaine, 1987, Chap.\\,4, \\S2).  For sake of simplicity, the applications of the general theory shall be restricted to homeoidally striated density profiles (e.g., Roberts, 1962; Caimmi, 1993; Caimmi and Marmo, 2003, hereafter quoted as CM03).   The larger effects of asphericity are expected to occurr in homogeneous configurations, which have widely been investigated (Brosche et al., 1983; Caimmi et al., 1984; CS90; Caimmi and Secco, 1992).   Focaloidally striated density profiles involve far larger difficulty (e.g., Caimmi, 1995, 2003).  The current investigation is mainly devoted to the following points: (i) expression of an equation of state for two-component systems; (ii) description of global properties deduced from selected density profiles; and (iii) application to elliptical galaxies belonging to a restricted sample, to be represented on the $({\\sf O}y\\phi)$ plane, for fiducial values of model parameters.  The work is organized as follows.   The basic theory of two-component systems with homeoidally striated density profiles, is reviewed and extended (to include both cored and cuspy matter distributions) in Section \\ref{bath}.   The particularization to selected density profiles, involving explicit expressions, is made in Section \\ref{spec}.   The results and related global properties are described and discussed in Section \\ref{resu}. Elliptical galaxies belonging to a restricted sample, are represented on the $({\\sf O}y\\phi)$ plane, for fiducial values of model parameters, in Section \\ref{appl}, where some considerations are drawn.   The concluding remarks are reported in Section \\ref{conc}. ",
        "conclusions": "\\label{conc}  Two-component systems have been conceived as (two-component) macrogases, and the related equation of state has been formulated using the virial theorem for subsystems (Limber, 1959; Brosche et al., 1983; Caimmi et al., 1984; Caimmi and Secco, 1992), under the restriction of (i) homeoidally striated ellipsoids (Roberts, 1962) and (ii) similar and similarly placed boundaries.  Explicit calculations have been performed for a useful reference case and a few cases of astrophysical interest, both in presence and in absence of truncation radius. More specifically, the following cases have been dealt with: IJ= UU, PP, HH, HP, HN, where I and J denote the inner and the outer density profile, respectively, and the other captions relate to the following density profiles: U ($\\rho=$const), P (Plummer, 1911), H (Hernquist, 1990), N (Navarro et al., 1995, 1996, 1997).   Shallower density profiles (UU, PP), have been found to yield an equation of state, $\\phi=\\phi(y,m)$, characterized by the occurrence of two extremum points, one maximum and one minimum, as in an earlier attempt (CS90). Steeper density profiles (HH, HP, HN), have been found to produce a similar equation of state where, in addition, a single horizontal inflexion point occurs in a critical isofractional mass curve, and isofractional mass curves related to lower values, $m=M_j/M_i<m_{\\rm crit}$, show no extremum point. The similarity between isofractional mass curves and van der Waals' isothermal curves, has suggested the possibility that a phase transition could take place in a bell-shaped region of the $({\\sf O}y\\phi)$ plane.  Further investigation has been devoted to HH density profiles for which a numerical algorithm (Ciotti et al., 1996) has been used to represent EGs and their hosting DM haloes along selected isofractional mass curves on the $({\\sf O}y\\phi)$ plane, under the assumption that the related fractional mass has the same value in different systems.    In the light of the model, the evolution of EGs has been found to be non strictly homologous, in the sense that the end of evolutionary tracks on the $({\\sf O}y\\phi)$ plane occur at different points along the related isofractional mass curve, instead of being close to a single point."
    },
    "0801/0801.4226_arXiv.txt": {
        "abstract": "{} {We present an extension to our \\xmm\\ X-ray source catalogue of \\m31\\, containing 39 newly found sources. In order to classify and identify more of the sources we search for X-ray time variability in \\xmm\\ archival data of the \\m31\\ centre field.} {As a source list we used our extended catalogue based on observations covering the time span from June 2000 to July 2004. We then determined the flux or at least an upper limit at the source positions for each observation. Deriving the flux ratios for the different observations and searching for the maximum flux difference we determined variability factors. We also calculated the significance of the flux ratios.} {Using hardness ratios, X-ray variability and cross correlations with catalogues in the X-ray, optical, infrared and radio regimes, we detected three super soft source candidates, one supernova remnant and six supernova remnant candidates, one globular cluster candidate, three X-ray binaries and four X-ray binary candidates. Additionally we identified one foreground star candidate and classified fifteen sources with hard spectra, which may either be X-ray binaries or Crab-like supernova remnants in \\m31\\ or background active galactic nuclei. The remaining five sources stay unidentified or without classification. Based on the time variability results we suggest six sources, which were formerly classified as ``hard\", to be X-ray binary candidates. The classification of one other source (XMMM31~J004236.7+411349) as a supernova remnant, has to be rejected due to the distinct time variability we found. We now classify this source as a foreground star. } {} ",
        "introduction": "An ideal target for a search for time variability of X-ray sources is the bright Local Group spiral galaxy \\m31\\ \\citep[distance 780 kpc,][]{1998AJ....115.1916H,1998ApJ...503L.131S} with its moderate Galactic foreground absorption  \\citep[\\nh = 7\\hcm{20}, ][]{1992ApJS...79...77S}.\\@  The \\ein\\ X-ray observatory found 16 sources in \\m31, which showed variability comparing the individual observations with each other \\citep[][ hereafter TF91]{1979ApJ...234L..45V, 1990ApJ...356..119C,1991ApJ...382...82T}.\\@ \\citet[][ hereafter PFJ93]{1993ApJ...410..615P} compared ROSAT HRI to previous \\ein\\ observations and found several variable sources. The two ROSAT PSPC surveys of \\m31, covered the entire galaxy and were separated by about one year. \\citet{1997A&A...317..328S,2001A&A...373...63S} found, that the intensity of 34 sources varied significantly between the observations.  \\citet{2000ApJ...537L..23G} reported on first observations of the nuclear region of \\m31\\ with \\chandra. They found that the nuclear source has an unusual X-ray spectrum compared to the other point sources in \\m31. Source catalogues, based on \\chandra\\ observations, of the central field of \\m31\\ are provided by \\citet{2002ApJ...577..738K} and \\citet{2002ApJ...578..114K}. Three different \\m31\\ disk fields, spanning a range of stellar populations, were observed by \\chandra\\ to compare their point source luminosity functions to that of the galaxy's bulge \\citep{2003ApJ...585..298K}.\\@ In a synoptic study of the disk ($\\approx 0.9$ square degree) of \\m31, \\citet{2004ApJ...609..735W} measured the mean flux and long-term light curves for 166 objects. At least 25\\% of the sources show significant variability. Bright X-ray binaries (XRBs) in globular clusters and supersoft sources (SSSs) and quasisoft sources (QSSs) were investigated by \\citet[][]{2002ApJ...570..618D,2004ApJ...610..247D} and \\citet{2004ApJ...610..261G}.\\@ The discovery of an X-ray nova was reported by \\citet{2005ApJ...620..723W}.\\@ \\citet{2007A&A...468...49V} used \\chandra\\ data to examine the low mass X-ray binaries (LMXBs) in the bulge of \\m31. Good candidates for LMXBs are the so-called transient sources. Studies of transient sources in \\m31\\ are presented in numerous papers, e.\\,g.~\\citet{2006ApJ...643..356W}, \\citet[][ hereafter TPC06]{2006ApJ...645..277T}, \\citet{2005ApJ...632.1086W}, \\citet[][ hereafter WGM06]{2006ApJ...637..479W}. Using \\xmm\\ and \\chandra\\ data, \\citet{2004ApJ...616..821T} detected 43 X-ray sources coincident with globular cluster candidates from various optical surveys. They studied their spectral properties, time variability and luminosity functions. \\citet{2001A&A...378..800O} used \\xmm\\ Performance Verification observations to study the variability of X-ray sources in the central region of \\m31. They found 116 sources brighter than a limiting luminosity of 6\\ergs{35} and examined the $\\sim60$ brightest sources for periodic and non-periodic variability. At least 15\\% of these sources appear to be variable on a time scale of several months. \\citet{2003A&A...411..553B} used \\xmm\\ to study the X-ray binary RX J0042.6+4115 and suggested it as a Z-source. \\citet{2006ApJ...643..844O} studied the population of SSSs and QSSs with \\xmm. \\citet{2006astro.ph.10809T} provide a study of the bright sources in the central region of \\m31, including spectral properties, variability and source classification. It is based on the same \\xmm\\ observations analysed in this paper. Recently \\citet{2007arXiv0708.0874T} reported the discovery of 217s pulsations in a bright persistent SSS.  \\citet[][ hereafter PFH2005]{2005A&A...434..483P} prepared a catalogue of \\m31\\ point-like X-ray sources analysing all observations available at that time in the \\xmm\\ archive which overlap at least in part with the optical $\\mr{D}_{25}$ extent of the galaxy. In total, they detected 856 sources. The centre part of the galaxy was covered four times with a separation of the observations of about half a year starting in June 2000 (some of the regions at the boundary of the centre area are even covered five times). PFH2005 give only source properties derived from an analysis of the combined centre observations. In follow-up work (i) \\citet[][]{2005A&A...430L..45P} searched for X-ray burst sources in globular cluster (GlC) sources and candidates and identified two X-ray bursters and a few more candidates, and (ii) \\citet[][ hereafter PFF2005]{2005A&A...442..879P} searched for correlations with optical novae. They identified seven SSSs and one symbiotic star from the list of PFH2005 with optical novae and identified one additional \\xmm\\ source with an optical nova. This work was continued in \\citet[][ hereafter PHS2007]{2007A&A...465..375P}.  Similar to the \\object{M~33} work of \\citet{2004A&A...426...11P} PFH2005 used the hardness ratios, i.\\,e.\\ X-ray colours, and correlations with sources in other wavelength regimes to identify and classify the detected sources. \\citet{2006A&A...448.1247M} showed, for a source population study of \\object{M~33}, that X-ray flux variability on different time scales allows us to further distinguish between different source classes. Phenomena such as bursts of X-ray binaries, flares of stars or the periodic variability of pulsars occur on short time scales and can therefore be observed during one single observation. On the other hand there is long term variability. Those time scales can be covered, comparing different observations of the same source. In the field of view of \\m31\\ there are mainly two source classes, which are known to show strong variability (variability factor $> 10$) on time scales of years. These are X-ray binaries and SSS.\\@ Among the active galactic nuclei (AGNs) narrow-line Seyfert 1 galaxies and BL Lac objects show the strongest variabilities. However only a few narrow-line Seyfert 1 galaxies are known in the entire sky, which show flux variability factors larger than 10 on time scales of half a year up to several years. Hence it is very unlikely that one of the strongly variable sources in \\m31\\ would be an AGN.  In this paper we report a search for new X-ray sources in the \\xmm\\ observations to the \\m31\\ centre to extend the source catalogue of PFH2005 and a time variability analysis of all the \\m31\\ centre sources. In Sect.~2 information about the used observations and accomplished analysis is provided. Sect.~3 describes the source catalogue extension. The results of the temporal variability analysis are discussed in Sect.~4. Discussion of the individual source classes, including X-ray identifications, are provided in Sect.~5. We draw our conclusions in Sect.~6. ",
        "conclusions": "In this paper we present an updated source list of the central area of the bright Local Group spiral galaxy \\m31, using the observations from June 2000 to July 2004 available from the \\xmm\\ archive. We extended the source catalogue by PFH2005, based on the merged images of the observations from 2000 to 2002 by searching sources in the observations of 2004 and reexamining the observations used in PFH2005 individually. To classify or identify more of the sources, we examined their long term time variability.  We obtained 39 sources in addition to the 265 reported by PFH2005 in the field. The identification and classification of these sources is based on properties in the X-ray wavelength regime: hardness ratios and temporal variability. In addition, information from cross correlations with \\m31\\ catalogues in the radio, infra-red, optical and X-ray wavelength regimes are used.  We detected three SSS candidates, one SNR and six SNR candidates, one GlC candidate, three XRBs and four XRB candidates. Additionally we identified one foreground star candidate and classified fifteen sources as $<$hard$>$, which may either be XRBs or Crab-like SNRs in \\m31\\ or background AGNs. The remaining five sources remain unidentified and without classification. Two sources were found to be extended. One of them was classified as $<$hard$>$. The other stays without classification.  To examine the time variability we calculated the flux or at least an upper limit at the source positions in each observation. We determined the variability factor and significance parameter for each source, comparing the XID flux ratios of the different observations with each other. The time variability helped us to decide if a source classified as $<$hard$>$ in PFH2005 can be an XRB candidate. In addition we could use time variability to distinguish between foreground star and SNR candidates.  Six sources of PFH2005, which were classified as $<$hard$>$, show distinct time variability. Based on that variability, their hardness ratios and the strong absorption in the centre of \\m31\\ we suggest these sources as XRB candidates. The SNR classification from source 295 was changed to foreground star due to the distinct time variability we found and its identification with a faint stellar object. Other SNR classifications (sources 316, 318) were rejected due to time variability of the sources.  To verify our suggested classifications further investigations, including at other wavelengths will be necessary."
    },
    "0801/0801.1403_arXiv.txt": {
        "abstract": "We explore the importance of magnetic-field-oriented thermal conduction in the interaction of supernova remnant (SNR) shocks with radiative gas clouds and in determining the mass and energy exchange between the clouds and the hot surrounding medium. We perform 2.5D MHD simulations of a shock impacting on an isolated gas cloud, including anisotropic thermal conduction and radiative cooling; we consider the representative case of a Mach 50 shock impacting on a cloud ten-fold denser than the ambient medium. We consider different configurations of the ambient magnetic field and compare MHD models with or without the thermal conduction. The efficiency of the thermal conduction in the presence of magnetic field is, in general, reduced with respect to the unmagnetized case. The reduction factor strongly depends on the initial magnetic field orientation, and it is minimum when the magnetic field is initially aligned with the direction of shock propagation. The thermal conduction contributes to suppress hydrodynamic instabilities, reducing the mass mixing of the cloud and preserving the cloud from complete fragmentation. Depending on the magnetic field orientation, the heat conduction may determine a significant energy exchange between the cloud and the hot surrounding medium which, while remaining always at levels less than those in the unmagnetized case, leads to a progressive heating and evaporation of the cloud. This additional heating may contrast the radiative cooling of some parts of the cloud, preventing the onset of thermal instabilities. ",
        "introduction": "\\label{sec1}  The interaction of the shock waves of supernova remnants (SNRs) with the magnetized and inhomogeneous interstellar medium (ISM) is responsible of the great morphological complexity of SNRs and certainly plays a major role in determining the exchange of mass, momentum, and energy between diffuse hot plasma and dense clouds or clumps. These exchanges may occur through, for example, hydrodynamic ablation and thermal conduction and, among other things, lead to the cloud crushing and to the reduction of the Jeans mass causing star formation.  The propagation of hot SNR shock fronts in the ISM and their interaction with local over-dense gas clouds have been investigated with detailed hydrodynamic and MHD modeling. The most complete review of this problem in the unmagnetized, non-conducting, and non-radiative limits is provided by \\cite{1994ApJ...420..213K}. These studies have shown that the cloud is disrupted by the action of both Kelvin-Helmholtz (KH) and Rayleigh-Taylor (RT) instabilities after several crushing times, with the cloud material expanding and diffusing into the ambient medium. An ambient magnetic field can both act as a confinement mechanism of the plasma and be modified by the interstellar flow and by local field stretching. Also, a strong magnetic field is known to limit hydrodynamic instabilities developing during the shock-cloud interaction by providing an additional tension at the interface between the cloud and the surrounding medium \\citep[e.g][]{1994ApJ...433..757M, 1996ApJ...473..365J}.  The interaction of the shock with a {\\it radiative} cloud has been only recently analyzed in detail \\citep[e.g.][]{2002A&A...395L..13M, 2004ApJ...604...74F}. 2D calculations have shown that the effect of the radiative cooling is to break up the clouds into numerous dense and cold fragments that survive for many dynamical timescales. In the case of the interaction between magnetized shocks and radiative clouds, the magnetic field may enhance the efficiency of the radiative cooling, influencing the size and distribution of condensed cooled fragments \\citep{2005ApJ...619..327F}.  The role played by the thermal conduction during the shock-cloud interaction has been less studied so far. In a previous paper, \\cite{2005A&A...444..505O} (hereafter Paper I) have addressed this point in the unmagnetized limit. In particular, we have investigated the effect of thermal conduction and radiative cooling on the cloud evolution and on the mass and energy exchange between the cloud and the surrounding medium; we have selected and explored two different physical regimes chosen so that either one of the processes is dominant. In the case dominated by the radiative losses, we have found that the shocked cloud fragments into cold, dense, and compact filaments surrounded by a hot corona which is ablated by the thermal conduction. Instead, in the case dominated by thermal conduction, the shocked cloud evaporates in a few dynamical timescales. In both cases, we have found that the thermal conduction is very effective in suppressing the hydrodynamic instabilities that would develop at the cloud boundaries, preserving the cloud from complete destruction. \\cite{orlando2} and \\cite{2006A&A...458..213M} have studied the observable effects of thermal conduction on the evolution of the shocked cloud in the X-ray band.  Here, we extend the previous studies by investigating the effect of the thermal conduction in a magnetized medium, unexplored so far. Of special interest to us is to investigate the role of anisotropic thermal conduction - funneled by locally organized magnetic fields - in the mass and energy exchange between ISM phases. In particular, we aim at addressing the following questions: How and under which physical conditions does the magnetic-field-oriented thermal conduction influence the evolution of the shocked cloud? How do the mass mixing of the cloud material and the energy exchange between the cloud and the surrounding medium depend on the orientation and strength of the magnetic field and on the efficiency of the thermal conduction?  To answer these questions, we take as representative the model case of a shock with Mach number $\\mach = 50$ (corresponding to a post-shock temperature $T\\approx 4.7\\times 10^6$ K for an unperturbed medium with $T=10^4$ K) impacting on an isolated cloud ten-fold denser than the ambient medium. Paper I has shown that, in this case, the thermal conduction dominates the evolution of the shocked cloud in the absence of magnetic field. Around this basic configuration, we perform a set of MHD simulations, with different interstellar magnetic field orientations, and compare models calculated with thermal conduction turned either ``on'' or ``off'' in order to identify its effects on the cloud evolution.  The paper is organized as follows: in Sect. \\ref{sec2} we describe the MHD model and the numerical setup; in Sect. \\ref{sec3} we discuss the results; and finally in Sect. \\ref{sec4} we draw our conclusions. ",
        "conclusions": "  \\begin{enumerate} \\item In general, we found that the effects of thermal conduction on the evolution of the shocked cloud are reduced in the presence of an ambient magnetic field with respect to the unmagnetized cases investigated in Paper I. The efficiency of anisotropic thermal conduction strongly depends on the initial magnetic field orientation and configuration. This efficiency is the largest when the initial {\\mag} is aligned with the direction of propagation of the shock front, and is the smallest when {\\mag} is aligned with the cylindrical cloud, namely when the heat conduction is completely suppressed by the magnetic field.  \\item We found that the hydrodynamic instabilities are suppressed efficiently by the anisotropic thermal conduction when the initial magnetic field is perpendicular to the cylindrical cloud (a configuration referred to as ``external fields''). On the contrary, in the case of {\\mag} parallel to the cylindrical axis of the cloud (i.e. when the field has component only along the $z$ axis - internal field), hydrodynamic instabilities develop at the cloud boundary. We found that, for the parameters of the simulations chosen, the magnetic tension is unable to suppress alone the hydrodynamic instabilities.  \\item As for thermal instabilities, we found that, depending on the magnetic field orientation, the heat flux contributes to the heating of some parts of the cloud, reducing the efficiency of radiative cooling there, and preventing any thermal instability.  \\item The mass loss of the cloud due to mixing with the surrounding medium is mainly driven by hydrodynamic instabilities; in the case of external fields (initial {\\mag} perpendicular to the cylindrical cloud) the anisotropic thermal conduction reduces the mass mixing of the cloud.  In any case, the mass loss rate is larger than that in the corresponding unmagnetized case ($\\dot{m}_{\\rm cl}\\approx 1.5\\times 10^{-7} L_{\\rm pc}~~M_{\\odot}$~yr$^{-1}$, i.e. $\\sim 5$\\% of the cloud mass is in mixed zones at $t = 3.5\\;\\tau_{\\rm cc}$), but can get very high when the thermal conduction is completely suppressed ($\\dot{m}_{\\rm cl}\\approx 4\\times 10^{-6} L_{\\rm pc}~~M_{\\odot}$~yr$^{-1}$, i.e. $\\sim 45$\\% of the cloud mass is in mixed zones at $t = 3.5\\;\\tau_{\\rm cc}$).  \\item The thermal conduction mostly rules the energy exchange between the cloud and surrounding medium. The exchange is favored when the magnetic field configuration is such that the conductive flow is not suppressed (i.e. external field configurations, $B_{\\rm x}$ and $B_{\\rm y}$ cases), but it is never as high as in the absence of magnetic field. In the $B_{\\rm y}$ case, the cloud core is efficiently heated and evaporates in few dynamical timescales.  \\item In general, the initial magnetic field strength has a small influence on the dynamic and thermal evolution of the shocked cloud for the ranges of values explored in this paper (namely $0.26~\\mu\\mbox{G} \\leq |\\mag| \\leq 2.63~\\mu\\mbox{G}$). \\end{enumerate}  It is worth noting that some details of our simulations depend on the choice of the model parameters. For instance, the onset of thermal instabilities or the evaporation of the whole cloud depends on the initial shock Mach number, and on the density and dimensions of the cloud. The cases that we present here (i.e. $\\mach = 50$, $\\chi=10$, and different configurations of \\mag) are representative of a regime in which both the thermal conduction and the radiative cooling play an important role in the evolution of the shocked cloud. Nevertheless, our analysis proves that anisotropic thermal conduction can not be neglected in investigations of the evolution of shocked interstellar clouds.  In our simulations, we consider laminar thermal conduction, although regions of strong turbulence of different strength and extent develop in the system (for instance, at the shear layers along the cloud boundary or at the vortex sheets in the cloud wake). In fact, the turbulence in these regions may have a significant effect on thermal conduction, leading to significant deviations of thermal conductivity from its laminar values (e.g. \\citealt{2001ApJ...562L.129N}; \\citealt{2006ApJ...645L..25L}); in some cases, the turbulence may enhance the heat transfer, exceeding the classical Spitzer value (\\citealt{2006ApJ...645L..25L}). As a result, thermal conduction may be not only anisotropic (in the presence of the magnetic field) but also ``inhomogeneous'' due to the presence of turbulence. However, even modeling accurately the turbulent thermal conductivity, we do not expect significant changes in the results of our $B_{\\rm z}$ case, being the thermal conduction strongly ineffective in the whole spatial domain; in the remaining cases ($B_{\\rm x}$ and $B_{\\rm y}$), our modeled thermal conductivity could be underestimated in regions of strong turbulence, affecting some details of the simulations but not the main conclusion of the paper that, in general, anisotropic thermal conduction can play an important role in the evolution of the shocked cloud.  Note also that the field configurations studied in this work are highly idealized. More realistic fields are expected to have more complex topologies and, often, the field can be tangled and chaotic. In the latter case, the thermal conduction will approach isotropy, whereas the effect of MHD turbulence is expected to partially suppress the heat transfer within a factor $\\sim 5$ below the classical Spitzer estimate\\footnote{As already discussed, the MHD turbulence can even enhance the heat transfer in some cases (see \\citealt{2006ApJ...645L..25L}).} (\\citealt{2001ApJ...562L.129N}; \\citealt{2006ApJ...645L..25L}). The shock-cloud collision in the presence of an organized ambient magnetic field, discussed here, and that in the absence of magnetic field can be considered as extreme cases: the former leading to highly anisotropic thermal conduction, the latter to the classical Spitzer thermal conduction. The case of chaotic magnetic field is expected to fall in between these two.  Our simulations were carried out in 2.5D Cartesian geometry, implying that the modeled clouds are elongated along the $z$ axis. This choice is expected to affect some details of the simulations but not our main conclusions. Adopting a 3D Cartesian geometry and modeling a spherical cloud, the highly symmetric shock transmitted into the cloud converging on the symmetry axis would lead to compression stronger than those found in our 2.5D simulations, enhancing the radiative cooling. Also, 3D simulations would provide an additional degree of freedom for hydrodynamic instabilities, increasing the mass loss rate of the cloud in the cases in which the mass mixing of cloud material is driven by instabilities. Note that, for a spherical cloud, our $B_{\\rm x}$ and $B_{\\rm z}$ cases no longer differ.  Finally we assume, in our simulations, that the cloud and the ambient material have the same composition, implying that microscopic mass mixing due to shear instabilities would be irrelevant. In a more realistic condition, a cold dense cloud may have a different composition from the hot ambient flow and the degree of microscopic mixing may translate into different spectral signatures of the system. In this case, species diffusion could also be important, along with thermal conduction, to determine the degree of microscopic mixing of the materials and, consequently, one would have to ask about the typical values of the Lewis number (i.e. the ratio of thermal diffusivity to mass diffusivity) in the system.  It is worth emphasizing that the quantitative results of our simulations depend on the physical parameters of the model (shock Mach number, density contrast and dimension of the cloud, etc.) as well as on the basic assumptions of the model (geometry of the cloud, geometry of the ambient magnetic field, laminar thermal conduction, composition of the cloud and of the ambient medium, etc.). Nevertheless, our results undoubtedly show that the magnetic-field-oriented thermal conduction can play an important role in the evolution of the shock-cloud interaction (which depends on the magnetic field orientation and configuration) and, in particular, in the mass and energy exchange between the cloud and the hot surrounding medium. We conclude, therefore, that a self-consistent and quantitative description of the interaction between magnetized shock-waves and interstellar gas clouds should include the effects of thermal conduction.  The results presented here are interesting for the study of middle-aged SNR shells expanding into a magnetized ISM and whose morphology is affected by ISM inhomogeneities (for instance, G272.2-3.2, e.g. \\citealt{1996rftu.proc..247E}; Cygnus Loop, e.g. \\citealt{2002AJ....124.2118P}; Vela SNR, e.g. \\citealt{2005A&A...442..513M}). It will be further interesting to extend the present study, by modeling the shock-cloud interaction in 3D with radiative cooling, anisotropic thermal conduction, and magnetic field included and, even, considering detailed comparisons of model results with observations."
    },
    "0801/0801.4010_arXiv.txt": {
        "abstract": "I summarize the highlights of the conference. First I provide a brief history of the beach symposia series our massive star community has been organizing. Then I use most of my allocated space discussing what I believe are the main answered and open questions in the field. Finally I conclude with a perspective of the future of massive star research. ",
        "introduction": "This is the ninth meeting in what has become the traditional beach symposia series of the massive star community. The series had its origins in 1968 at a workshop on Wolf-Rayet (W-R) stars in Boulder (definitely not a beach location). At that time, the cosmic significance of W-R stars was largely unknown but understanding their nature was deemed important enough for follow-up meetings. This then led to IAU Symp. 49 in Argentina in 1971, and to the subsequent tradition of holding massive star symposia at roughly five-year intervals. For those who like statistics, here is the complete tally: IAU Symp. 83 (1978, Canada), IAU Symp. 99 (1981, Mexico), IAU Symp. 116 (1985, Greece), IAU Symp. 143 (1990, Indonesia), IAU Symp. 163 (1994, Italy), IAU Symp. 193 (1998, Mexico), IAU Symp. 212 (2002, Spain), and finally IAU Symp. 250 (2007, USA).  Each meeting had its distinct flavor and theme. For instance, IAU Symp. 163 had the emphasis on binary stars, and IAU Symp. 116 allocated a large fraction of time to massive stars in Local Group galaxies. Overall, a clear trend is apparent: the median distance of the astronomical objects at each meeting has been steadily increasing. Distances were expressed in kpc in the early symposia, then became Mpc, and during this meeting redshifts larger than 3 were commonly mentioned.  Those who attended IAU Symp. 49 would have been astounded had they known how the field would have progressed by the time of IAU Symp. 250. Even as recently as during IAU Symp. 193 (which I attended as well), it would have been preposterous to discuss massive stars in connection with gamma-ray bursts, Lyman-break galaxies, and Population III stars.  I will finish my historical notes with a bit of trivia. Three participants of this symposium have witnessed this rapid development from the very beginning. They already participated in the 1971 meeting: Peter Conti, Lindsay Smith\\footnote{I apologize to Lindsay Smith whom I failed to name in my original talk.}, and Nolan Walborn. Peter Conti distinguished himself with an impressive feat: he attended {\\em all} nine beach symposia. His record would have been tied were it not for Virpi Niemela's untimely death in 2006. ",
        "conclusions": ""
    },
    "0801/0801.3295_arXiv.txt": {
        "abstract": " ",
        "introduction": "The primordial power spectrum or 2-point correlation function  of the temperature anisot\\-ropies in the Cosmic Microwave Background (CMB) is well-measured out to large multipoles. The higher moments of a distribution are, in general, independent of the 2-point function, but the CMB anisotropies are at least approximately Gaussian.   Theoretically, we know that non-Gaussian component to the CMB will always be induced by  non-linear gravitational couplings between modes after they reenter the horizon \\cite{Pyne:1995bs}  while  single field slow-roll inflation yields a primordial  non-Gaussian component roughly 100 times smaller than that induced by gravitational mode-couplings  \\cite{Acquaviva:2002ud,Maldacena:2002vr}. The two terms are additive, so recovering this latter primordial signal is next to impossible, even before contending with cosmic variance.  Multi-field models may generate larger non-Gaussianities \\cite{Rigopoulos:2005us,Seery:2005gb,Vernizzi:2006ve,Battefeld:2006sz,Yokoyama:2007uu}, but this is by no means a generic property of these scenarios. Consequently, the detection of a primordial 3-point function would immediately falsify a very large class of inflationary models. Conversely, non-slow-roll models with higher order derivative terms, such as DBI inflation \\cite{Silverstein:2003hf,Alishahiha:2004eh,Chen:2004gc,Chen:2005ad} and k-inflation \\cite{ArmendarizPicon:1999rj,Garriga:1999vw}, do typically generate large non-Gaussianities  \\cite{Chen:2006nt}. Further references, include more complicated multi-field, non-local or ghost theories, can be found in Refs.~\\cite{Li:2008qc,Bean:2008ga,Lyth:2002my,Chen:2007gd,Gupta:2002kn,Moss:2007cv,Barnaby:2007yb,ArkaniHamed:2003uz,Cheung:2007st}. In this paper we investigate simple models -- in the sense that they are driven by a single, minimally coupled scalar field with a canonical kinetic term -- which generate substantial non-Gaussianities.  Constraining the non-Gaussian signal in a CMB dataset is a highly non-trivial problem.  Firstly, it depends on the choice of estimators. At the moment, only two concrete estimators have been constructed: the $f_{NL}^{local}$ and $f_{NL}^{equil}$ forms \\cite{Komatsu:2001rj,Komatsu:2003iq,Babich:2004gb,Spergel:2006hy,Creminelli:2006gc,Creminelli:2006rz,Yadav:2007ny}, and both are scale-invariant. The recent claim of a detection of a non-zero 3-point function \\cite{Yadav:2007yy} in the WMAP 3-year data \\cite{Spergel:2006hy} relies on the estimator developed in Ref.~\\cite{Yadav:2007ny,Creminelli:2006gc,Komatsu:2003iq} which is of the ``local'' form  \\cite{Salopek:1990jq,Komatsu:2001rj}. However non-Gaussianities that have strong scale dependence are not well-described by $f_{NL}^{local}$ and $f_{NL}^{equil}$, and will require  scale-dependent  estimators, based on theoretically motivated ansatzen % for the primordial non-Gaussianities.  In addition to computing the primordial 3-point term, comparing the predictions of a specific inflationary model to the CMB requires us to evolve this signal through to recombination, and to project it onto the sphere on the sky, which is a convolution of the primordial 3-point term with 3 $l$-valued spherical Bessel functions. In general, this process is computationally expensive but  simplifies dramatically if the 3-point function has special algebraic properties \\cite{Smith:2006ud,Fergusson:2006pr}.  Using Maldacena's elegant formalism \\cite{Maldacena:2002vr}, the primordial 3-point curvature correlation function is computed via a set of integrals (over time) of products of three mode functions (or their derivatives) and slow roll parameters.  Looking at these integrals, we identify two mechanisms which can create substantial non-Gaussianities.  The first class consists of potentials with a localized violation of slow roll. This can take the form of  a small localized feature, including the step models \\cite{Adams:2001vc,Komatsu:2003fd} whose 3-point term was first accurately computed by the present authors in \\cite{Chen:2006xj}. In addition models with a short inflationary phase can have initial transients in their dynamics, which will not be fully erased before the longest modes leave the horizon.  In these cases, the 3-point term for modes which are leaving the horizon during the violation of slow roll can be magnified by three orders of magnitude, without ruining the fit to the 2-point function.  The second class arises when  a small ripple is superimposed on top of an otherwise smooth potential.  This induces a ``resonance'' inside one of these integrals, giving the 3-point function a substantial amplitude {\\em before} the modes cross the horizon. In the former case the 2-point function may be substantially modified; in the second case the modification of the 2-point function is very small, even though the 3-point function is large. The latter mechanism has not been described previously and the seemingly contrived field theoretic potential may in fact arise naturally in brane inflation models \\cite{Bean:2008na}.  With this information in hand, we construct rough analytic approximations to the corresponding 3-point terms.  These expressions have the factorizable form required by  \\cite{Smith:2006ud}, which means that we can efficiently compute constraints on step potentials or similar models from CMB data, although at this point we are only interested in a qualitative match to the numerically computed 3-point term.    % We check our semi-analytic estimates for the 3-point function using an improved version of the code described in \\cite{Chen:2006xj}, which has a  much cleaner scheme for removing numerical divergences in the 3-point integrals. The integrands are products of large, rapidly oscillating terms. Analytically, these are finite, but their {\\em numerical\\/} evaluation is non-trivial, and we show how to transform them into an explicitly finite form before the numerical evaluation is carried out.   From a  practical perspective, this means we can compute the 3-point function for ``triangles'' which contain two very different scales.  The paper is organized as follows. In Section \\ref{sect:MW} we review the computation of  the 3-point function, and in Section  \\ref{sect:NumInt} we describe the numerical algorithm. Sections \\ref{sect:horizon} and \\ref{sect:subhorizon}  discuss the two distinct \\emph{mechanisms} for generating large primordial non-Gaussianities within minimally coupled  single scalar field inflation,  and  we derive the approximate analytic ansatzen for the 3-point function. We then use our numerical methods to compute the exact 3-point function for these models, and show that these approximations yield   fair representations of the underlying non-Gaussianities. In Section \\ref{sect:Conclusions} we summarize and discuss our results and plans for future work.   We follow the notational conventions of \\cite{Chen:2006xj} and  set the reduced Planck mass $M_p$ to unity except when presenting final results or defining parameter values. ",
        "conclusions": "\\label{sect:Conclusions}  We studied two distinctly mechanisms for generating large non-Gaussianities  within single field inflation. We derive the approximate 3-point correlation functions using semi-analytic methods, which are in the computationally useful \\emph{factorizable} form.  In the first mechanism, non-Gaussianity is generated at horizon crossing by either a feature in the potential, or an initial transient in the inflationary dynamics.  In the second case, oscillating slow roll parameters induce a resonance which leads to the generation of a non-Gaussian signal well before horizon crossing. In both cases, the 3-point function depends strongly on the individual wavelengths of the modes in the ``triangle''.   With a single, sharp feature or non-standard initial conditions, the 3-point is only enhanced in modes which are crossing the horizon as the inflaton traverses the ``feature''.   In a resonance model such as (\\ref{Vexample}),  the physical non-Gaussianity is present at all scales.  In the resonance case, we showed that the 3-point function is periodic with a period $\\Delta K$  proportional to, and smaller than, $K=k_1+ k_2+ k_3$.  This non-Gaussianity will peak starting from a fixed scale when projected onto the CMB sky. Assume for simplicity that $K \\sim \\ell$ where $\\ell$ is the CMB multipole, and denote $\\Delta \\ell$ as the oscillation period. At larger scales where the oscillation spanning $\\Delta \\ell <1$, this non-Gaussianity presumably cannot be resolved experimentally. For the numerical example considered here, it would become  visible at $\\ell \\sim \\CO(100)$ where $\\Delta \\ell$ starts to exceed $\\CO(1)$, inducing an effective scale-dependence in this signal.  With a feature in the potential, the resulting transient violation of slow-roll  generically leads to an oscillatory and scale-dependent  3-point function.  We wrote down a heuristic and factorizable scale-dependent expression  for this signal and showed that it captured the qualitative properties seen in the exact numerical evaluations of the corresponding integrals. The 3-point correlation decays as we move away from $K \\sim k_{*}$, where $k_{*}$ corresponds to the scale leaving the horizon as the inflaton traverses the feature.  We also show that while non-standard initial conditions such as the fast-roll model of \\cite{Contaldi:2003zv} can generate large non-Gaussianities relative to our usual expectations for single field inflation, the amplification is not expected to lift them above the ``noise'' of non-linearities produced by post-inflationary gravitational evolution alone.  In addition,  we have presented detailed description of a general numerical method for computing the 3-point correlation functions of primordial perturbations from canonical single scalar field inflationary models with arbitrary potentials. We show that while the integrals themselves are formally convergent, they need to be regularized as the integrands are oscillatory, and show how this can be accomplished analytically, rendering the numerical integrals rapidly convergent.  There is much further work to be done. Our immediate goal \\cite{preparation} is to use our heuristic ansatzen to construct an optimal estimator with which to search for scale-dependent non-Gaussianities in the cosmic microwave background, and estimate the likely bounds that future missions can put on this signal. On the theoretical front,  we plan to investigate the details of non-Gaussianities generated by multi-field models \\cite{Seery:2005gb,Battefeld:2006sz} within the horizon, and with non-standard kinetic terms.  Moreover, while the analytic approximations we present here capture the qualitative form of the 3-point function, they are not intended to provide a precise quantitative match to the numerically computed values, but these approximations can certainly be improved."
    },
    "0801/0801.2770_arXiv.txt": {
        "abstract": "Feedback from star formation is thought to play a key role in the formation and evolution of galaxies, but its implementation in cosmological simulations is currently hampered by a lack of numerical resolution. We present and test a sub-grid recipe to model feedback from massive stars in cosmological smoothed particle hydrodynamics simulations. The energy is distributed in kinetic form among the gas particles surrounding recently formed stars. The impact of the feedback is studied using a suite of high-resolution simulations of isolated disc galaxies embedded in dark halos with total mass $10^{10}$ and $10^{12}~\\Msolh$. We focus in particular on the effect of pressure forces on wind particles within the disc, which we turn off temporarily in some of our runs to mimic a recipe that has been widely used in the literature. We find that this popular recipe gives dramatically different results because (ram) pressure forces on expanding superbubbles determine both the structure of the disc and the development of large-scale outflows. Pressure forces exerted by expanding superbubbles puff up the disc, giving the dwarf galaxy an irregular morphology and creating a galactic fountain in the massive galaxy. Hydrodynamic drag within the disc results in a strong increase of the effective mass loading of the wind for the dwarf galaxy, but quenches much of the outflow in the case of the high-mass galaxy. \\endabstract  \\keywords methods: numerical --- ISM: bubbles --- ISM: jets and outflows --- galaxies: evolution --- galaxies: formation --- galaxies: ISM \\endkeywords ",
        "introduction": "Core-collapse supernovae (SNe) and winds from massive stars feed back energy into the interstellar medium (ISM). The energy released by massive stars can efficiently suppress subsequent star formation by destroying dense, star forming clouds, by generating supersonic turbulence and, if the star formation density is sufficiently high for individual supernova remnants to overlap, by blowing gas out of the disc. In starburst galaxies feedback from star formation may result in the development of a galaxy-wide superwind which may (temporarily) remove a large fraction of the gas, while in galaxies with less intense star formation feedback may lead to the development of a galactic fountain. Because these feedback mechanisms operate on time scales that are very short compared to the age of the universe, they can lead to self-regulation.  Feedback from star formation is thought to be a crucial ingredient for models of the formation and evolution of galaxies. Without it, star formation becomes much more efficient than observed, particularly in low-mass galaxies \\cite[e.g.][]{White&Frenk1991}. Among other things, galactic winds are also thought to be responsible for the enrichment of the intergalactic medium with heavy elements \\cite[e.g.][]{Aguirre2001} and for pre-heating the gas that ends up in groups and clusters of galaxies \\cite[e.g.][]{Ponman1999}.  Modelling SN feedback in simulations of the formation of galaxies is known to be a difficult task, mostly because the injected thermal energy tends to be radiated away well before it has any hydrodynamical effect \\citep[e.g.][]{Katz1996,Balogh2001}. This overcooling problem is probably caused by the fact that current state-of-the-art simulations still lack the resolution to capture the physics of the multiphase ISM \\cite[e.g.][]{Ceverino&Klypin2007}.  A typical SN ejects $\\sim 1~\\Msun$ at $\\sim 10^4~\\kms$, which corresponds to a kinetic energy $\\sim 10^{51}~\\erg$. Since the sound-crossing time is initially much smaller than the radiative cooling time, the remnant starts out as an adiabatic blast wave. Once radiative losses become important, the energy-conserving blast wave gives way to a momentum conserving snow-plough phase, whose deceleration is determined mostly by the amount of mass that is swept up by the wind. The initial phases in the evolution of a superbubble driven by multiple SNe are very similar to the evolution of an individual SN remnant. However, if a superbubble blows out of the disc, its subsequent evolution may be strongly affected by the ram pressure of infalling gas and the gravitational attraction of the galaxy \\cite[e.g.][]{Silich1998,Maclow1999,Dercole1999,Silich2001,Fujita2004,Dubois2008}.   Current simulations lack the resolution to resolve individual SNe or, in the case of cosmological simulations\\footnote{By cosmological simulations we mean simulations that model the evolution of a representative volume of the universe, as opposed to simulations that zoom in on one or a few galaxies.} even superbubbles. The energy released by dying stars in a single resolution element per time step is typically distributed over a mass that exceeds the mass of SN ejecta by many orders of magnitude. The initial expansion velocity and post-shock temperature are therefore underestimated by large factors. Since the radiative cooling time scales as $t_c \\propto T^{1/2}$ above 1~keV, this implies that the cooling time is greatly underestimated. Because the injection radius is also far too large, the initial cooling time tends to be smaller than the bubble sound-crossing time. Thus, the simulation will skip the adiabatic blast-wave phase and the energy will typically be radiated away before a significant fraction of the thermal energy has been converted into kinetic energy.  The fact that poor resolution results in inefficient thermal feedback is generally attributed to a lack of spatial resolution: real SNe explode in hot bubbles of low-density gas, whereas the ISM in cosmological simulations consists of a single dense, ``warm'' phase. Mass resolution is, however, more fundamental. Without sufficient mass resolution, the first SNe will not be able to create a hot, low-density ISM phase in the first place.  Cosmological simulations must resort to sub-grid recipes to solve the overcooling problem.\\footnote{In the absence of efficient feedback mechanisms, the amount of cold gas predicted by cosmological simulations is limited by resolution. Hence, it is possible to roughly match the observed amount of cold gas at a particular redshift by tuning the resolution.} Two types of solutions appear to work: injecting (part of) the SN energy in kinetic rather than thermal form \\cite[e.g.][]{Navarro1993,Mihos&Hernquist1994,Kawata2001,Kay2002,Springel2003,Oppenheimer2006,Dubois2008} and/or suppressing radiative cooling by hand \\cite[e.g.][]{Gerritsen1997,Mori1997,Thacker2000,Kay2002,Sommer-Larsen2003,Brook2004,Stinson2006}. Although the suppression of cooling enables the efficient conversion of thermal energy to kinetic energy, the maximum wind velocity will still be underestimated if the total mass of the neighboring resolution elements exceeds that of a superbubble. Kinetic feedback can alleviate this problem by kicking only a small fraction of the resolution elements near the star particle. Kinetic feedback thus gives us the freedom to distribute a fixed amount of kinetic energy over a varying amount of mass. However, using this freedom to increase the initial wind velocity has the drawback that the imposed winds become more poorly sampled and therefore less isotropic. With increasing resolution, the two types of sub-grid models for the generation of galactic winds are expected to converge.  A third approach, which is often combined with one of the above, is to employ a sub-grid model to describe the multiphase ISM \\cite[e.g.][]{Yepes1997}. Examples include imposing an effective equation of state which specifies the total pressure of the ISM as a function of its mean density \\cite[e.g.][]{Springel2003} and, at least for the case of smoothed particle hydrodynamics simulations, explictly decoupling thermal phases by using different types of particles to represent the hot and cold gas phases \\cite[e.g.][]{Marri&White2003,Scannapieco2006}. While galactic winds can in principle be triggered naturally if the latter method is used, this again requires ad-hoc solutions in the absence of sufficient resolution.  To prevent the overcooling problem, simulations that impose an effective equation of state for the ISM must either make it extremely stiff, resulting in discs that are much thicker and smoother than observed, or employ a sub-grid recipe for galactic winds. However, even if it does not directly generate winds, the use of an effective equation of state for dense gas can be considered a necessary ingredient for simulations that lack the resolution and/or physics to model the multiphase ISM. If the equation of state is not modified by hand, gas will accumulate at unrealistically low temperatures and high densities in the absence of resolved feedback processes. As we discussed in \\cite{Schaye2008}, using a power-law equation of state with a polytropic index of $4/3$ has the advantage of yielding a Jeans mass that is independent of the density which makes it possible to suppress spurious fragmentation due to a lack of resolution.  It is, however, important to note that maintaining an effective equation of state in the presence of radiative losses would, in reality, require energy. The amount of energy that is required depends on unresolved physical processes and can therefore not be reliably determined. Imposing an equation of state for dense gas therefore implies that the energy available for any wind sub-grid model must be less than the energy provided by star formation.  The most widely used recipe for galactic winds in SPH simulations is the kinetic feedback model implemented by \\cite{Springel2003} (hereafter SH03).  Motivated by the desire to impose the net mass loading and velocity of the wind after it has escaped from the disc and by the wish to make the recipe insensitive to numerical resolution, hydrodynamical forces on the wind particles are temporarily switched off in the SH03 model. Thus, winds cannot blow bubbles in the disc, drive turbulence or create channels in gas with densities typical of the ISM. Their sole effect on the disc is the removal of fuel for star formation by the ejection of wind particles.  Another aspect of the SH03 recipe is that wind particles are selected stochastically from all the dense (i.e.\\ star-forming) particles in the simulation and are therefore not constrained to be neighbours of newly-formed stars. While this non-local feedback solves some numerical problems, we will argue that it can lead to undesirable behavior in galaxies that do not contain large numbers of particles (i.e.\\ most galaxies in cosmological simulations), particularly if metal enrichment is included.  We present and test a modified a variation of the SH03 recipe in which the winds are local and not decoupled hydrodynamically. In this paper we will focus on the effects of hydrodynamical drag within the ISM. Using high-resolution SPH simulations of isolated disc galaxies we show that pressure forces exerted by and on wind particles have a dramatic effect on both the structure of the ISM and the development of galactic winds. Hydrodynamical drag on wind particles converts low-mass disc galaxies into irregulars and results in a strong increase of the net wind mass loading, while for high-mass galaxies it leads to the creation of bubbles and the development of a galactic fountain. Pressure forces within the disc reduce the kinetic energy in the wind above the disc by orders of magnitude.  This paper is organized as follows. We present our recipe for galactic winds and its implementation in SPH simulations in section~\\ref{sec:recipe}. After describing our simulations of isolated disc galaxies in section~\\ref{sec:sims}, we present the results on the galaxies' morphology, star formation histories and large-scale winds in sections \\ref{sec:morphology} to \\ref{sec:winds}, respectively. In section \\ref{sec:resol} we present resolution tests which show that the effects of hydrodynamical drag are underestimated if the Jeans scale is unresolved. Finally, we summarize and discuss our conclusions in section~\\ref{sec:disc}.  \\begin{table*} \\begin{center} \\caption{Simulation parameters: total mass, $M$; input mass loading, $\\eta$; input wind velocity, $v_{\\rm w}$; total number of particles, $N_{\\rm tot}$; total number of gas particles in the disc, $N_{\\rm disc}$; mass of baryonic particles, $m_{\\rm b}$; mass of dark matter particles, $m_{\\rm DM}$; gravitational softening of baryonic particles, $\\epsilon_{\\rm b}$; gravitational softening of dark matter particles, $\\epsilon_{\\rm DM}$; wind feedback included, (Wind); wind particles hydrodynamically decoupled, (Decoupled). Values different from the fiducial ones are shown in bold. \\label{tbl:params}} \\begin{tabular}{cccccccccccc}  \\hline Simulation & $M_{\\rm halo}$ & $\\eta$ & $v_{\\rm w}$ & $N_{\\rm tot}$ & $N_{\\rm disc}$ & $m_{\\rm b}$ & $m_{\\rm DM}$ & $\\epsilon_b$ & $\\epsilon_{\\rm DM}$ & Wind & Decoupled\\\\ & $(\\Msolh)$      &        & $(\\kms)$     &              &              & $(\\Msolh)$   & $(\\Msolh)$    & $(\\pch)$       & $(\\pch)$             &     &       \\\\ \\hline \\textit{m10}        & $10^{10}$ & 2 & 600 & 5000$\\,$494 & 235$\\,$294 & $5.1\\times 10^2$ & $2.4\\times 10^3$ & 10 &  17 & Y & N \\\\ \\textit{m10nowind}  & $10^{10}$ & -- & -- & 5000$\\,$494 & 235$\\,$294 & $5.1\\times 10^2$ & $2.4\\times 10^3$ & 10 &  17 & \\textbf{N} & -- \\\\ \\textit{m10$\\eta$1v848} & $10^{10}$ & \\textbf{1} & \\textbf{848} & 5000$\\,$494 & 235$\\,$294 & $5.1\\times 10^2$ & $2.4\\times 10^3$ & 10 &  17 & Y & N \\\\ \\textit{m10$\\eta$4v424} & $10^{10}$ & \\textbf{4} & \\textbf{424} & 5000$\\,$494 & 235$\\,$294 & $5.1\\times 10^2$ & $2.4\\times 10^3$ & 10 &  17 & Y & N \\\\ \\textit{m10dec}     & $10^{10}$ & 2 & 600 & 5000$\\,$494 & 235$\\,$294 & $5.1\\times 10^2$ & $2.4\\times 10^3$ & 10 &  17 & Y & \\textbf{Y} \\\\ \\hline \\textit{m10lr008}   & $10^{10}$ & 2 & 600 &  \\textbf{625$\\,$061} &  \\textbf{29$\\,$411} & $\\mathbf{4.1\\times 10^3}$ & $\\mathbf{1.9\\times 10^4}$ & \\textbf{20} &  \\textbf{34} & Y & N \\\\ \\textit{m10lr064}   & $10^{10}$ & 2 & 600 &   \\textbf{78$\\,$132} &   \\textbf{3$\\,$676} & $\\mathbf{3.3\\times 10^4}$ & $\\mathbf{1.5\\times 10^5}$ & \\textbf{40} &  \\textbf{68} & Y & N \\\\ \\textit{m10lr512}   & $10^{10}$ & 2 & 600 &    \\textbf{9$\\,$766} &    \\textbf{459} & $\\mathbf{2.6\\times 10^5}$ & $\\mathbf{1.2\\times 10^6}$ & \\textbf{80} & \\textbf{136} & Y & N \\\\ \\hline \\textit{m12}        & $10^{12}$ & 2 & 600 & 5000$\\,$494 & 235$\\,$294 & $5.1\\times 10^4$ & $2.4\\times 10^5$ & 10 &  17 & Y & N \\\\ \\textit{m12nowind}  & $10^{12}$ & -- & -- & 5000$\\,$494 & 235$\\,$294 & $5.1\\times 10^4$ & $2.4\\times 10^5$ & 10 &  17 & \\textbf{N} & -- \\\\ \\textit{m12$\\eta$1v848} & $10^{12}$ & \\textbf{1} & \\textbf{848} & 5000$\\,$494 & 235$\\,$294 & $5.1\\times 10^4$ & $2.4\\times 10^5$ & 10 &  17 & Y & N \\\\ \\textit{m12$\\eta$4v424} & $10^{12}$ & \\textbf{4} & \\textbf{424} & 5000$\\,$494 & 235$\\,$294 & $5.1\\times 10^4$ & $2.4\\times 10^5$ & 10 &  17 & Y & N \\\\ \\textit{m12dec}     & $10^{12}$ & 2 & 600 & 5000$\\,$494 & 235$\\,$294 & $5.1\\times 10^4$ & $2.4\\times 10^5$ & 10 &  17 & Y & \\textbf{Y}\\\\ \\hline \\textit{m12lr008}   & $10^{12}$ & 2 & 600 &  \\textbf{625$\\,$061} &  \\textbf{29$\\,$411} & $\\mathbf{4.1\\times 10^5}$ & $\\mathbf{1.9\\times 10^6}$ & \\textbf{20} &  \\textbf{34} & Y & N \\\\ \\textit{m12lr064}   & $10^{12}$ & 2 & 600 &   \\textbf{78$\\,$132} &   \\textbf{3$\\,$676} & $\\mathbf{3.3\\times 10^6}$ & $\\mathbf{1.5\\times 10^7}$ & \\textbf{40} &  \\textbf{68} & Y & N \\\\ \\textit{m12lr512}   & $10^{12}$ & 2 & 600 &    \\textbf{9$\\,$766} &    \\textbf{459} & $\\mathbf{2.6\\times 10^7}$ & $\\mathbf{1.2\\times 10^8}$ & \\textbf{80} & \\textbf{136} & Y & N \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "\\label{sec:disc}  Feedback from star formation is thought to play a key role in determining the observed properties of galaxies. It is therefore essential to include it in simulations of galaxy formation and evolution. Current simulations lack the resolution to resolve the energy conserving phase in the evolution of supernova remnants, causing any thermal energy input to be mostly radiated away before it has any significant hydrodynamical effect. Several types of sub-grid prescriptions for the generation of galactic winds have been proposed with the intent of overcoming numerical limitations. The most widely used methods use either thermal feedback combined with a temporary suppression of radiative cooling or kinetic feedback.  In this work we introduced a sub-grid recipe to model feedback from massive stars in cosmological SPH simulations. The energy is distributed in kinetic form among the neighbours of recently formed stars. We implemented our prescription in the SPH code \\textsc{gadget} and tested it using high-resolution simulations of isolated disc galaxies of total mass $10^{10}$ and $10^{12}~\\Msolh$. Our kinetic feedback scheme strongly reduces the star formation rates and has a dramatic impact on the morphology of the galaxies.  The disc of the dwarf galaxy becomes puffed up and punctuated with low-density bubbles, resembling the HI observations of nearby galaxies. A bipolar outflow develops naturally, with the largest velocities along the minor axis. The end result is a rather diffuse and irregular galaxy. Winds make the disc of the massive galaxy more stable, yielding more diffuse and less fragmented spiral arms. Its outer parts become more extended and also contain a large number of bubbles. While the wind is initially fastest along the minor axis, after a few hundred million years the velocities are higher at large opening angles because of strong infall along the minor axis. The edge-on gas density, temperature and pressure profiles remain, however, highly bi-conical. The infall is highly clumpy, consisting of cold gas clouds that have formed through thermal instabilities in the hot wind.  The mean outflow velocity decreases with time, because gas that is blown out of the disc later must plough through gas that was blown out earlier, some of which has turned around and is falling back. At large radii the velocity increases with radius due to the fact that only gas with some minimum velocity can reach a given radius in the time since the initial starburst. The mean outflow velocities are much lower than the input values ($600~\\kms$ for our fiducial model), except at the largest radii reached by the wind.  While the net mass outflow rate exceeds the input value ($ \\dot{M}_{\\rm w} / \\dot{M}_\\ast=2$ for our fiducial model) by about an order of magnitude for the dwarf galaxy, the reverse is true for the massive galaxy. Apparently, (ram) pressure forces in the disc enable the wind particles to drag large amounts of ISM out of the dwarf galaxy's disc, but are able to confine most wind particles to regions close to the disc in the case of the massive galaxy, whose ISM has both a higher density and pressure.  We varied the input wind velocity from 424 to $848~\\kms$ while keeping the total kinetic energy per unit stellar mass formed constant by adjusting the input mass loading accordingly. For the dwarf galaxy the resulting star formation histories are all nearly identical and in the case of the massive galaxy only the lowest velocity run differed significantly. Thus, provided the input wind velocity is higher than some minimum value, which increases with the mass (and thus ISM pressure) of the galaxy, the results depend on the input kinetic energy but are insensitive to the amount of mass the energy is distributed over.  While the star formation histories are relatively insensitive to numerical resolution, convergence of the predictions for the outflows requires resolving the Jeans scale.  We contrasted our scheme with the sub-grid model of SH03 which has been widely used in the literature. In the SH03 prescription, the wind particles are selected stochastically from all the star-forming (i.e.\\ dense) particles in the simulation and are therefore non-local. These wind particles are subsequently decoupled from the hydrodynamics for 50~Myr (i.e.\\ 31~kpc if traveling at $600~\\kms$) or until their density has fallen below 10 percent of the threshold for star formation.  In this work we have not tested the effects of non-locality because we expect them to be small for high-resolution simulations of individual galaxies such as those presented here. However, we do expect significant differences in cosmological simulations for which most galaxies will contain only a small number of star particles. In such galaxies the formation of star particles and the injection of kinetic energy will become essentially uncorrelated locally. To see this, note that for a galaxy that contains only a few star-forming particles, the typical time difference between the kicking of a wind particle and the creation of a star particle will be of order the gas consumption time scale divided by the input mass loading, or $\\sim 10^9/\\eta~\\yr$ at the threshold for star formation. This time scale greatly exceeds both the lifetime of massive stars and the simulation time step. The disconnection between the creation of a star particle and the injection of kinetic energy in its surroundings may have some undesirable consequences. For example, predictions for the chemical enrichment of the intergalactic medium may be affected because wind particles will contain an amount of heavy elements typical of particles in their host galaxies (which may often be zero) rather than abundances typical of the gas surrounding newly-formed stars.  We carried out simulations in which the wind particles were temporarily decoupled from the hydrodynamics as in the SH03 prescription. The difference between the predictions for the fiducial and decoupled models is dramatic. Decoupled winds have almost no effect on the morphology of the disc. Compared with the fiducial model and HI observations of nearby galaxies, the dwarf galaxy, which has a low enough surface density to be stable, is much smoother while the massive galaxy, which is unstable without the injection of turbulence, is more clumpy but lacks low-density bubbles. While the winds in the fiducial model slightly increase the size of the gas disc, the decoupled winds continuously shrink the disc. While the coupled winds drive a large-scale bipolar outflow from the dwarf galaxy and a clumpy galactic fountain in the massive galaxy, the decoupled winds produce isotropic outflows in both cases.  For the decoupled winds the outflow velocities (at least for $r < r_{\\rm vir}$) are constant in time and nearly independent of the galaxy mass. The mean, projected outflow velocity is about 70~percent of the input value, which greatly exceeds the outflow velocities observed in starbursting dwarf galaxies. The net mass outflow rates are in good agreement with the input values. The kinetic energy in the wind escaping the disc is orders of magnitude higher than in the case of the coupled winds. Compared with the fiducial model, the decoupled winds are less efficient at suppressing the SFR in the dwarf galaxy, but more efficient for the massive galaxy. The wind properties in the decoupled runs are insensitive to numerical resolution, even when the Jeans scale is completely unresolved. The lower the resolution of the runs with coupled winds, the more they resemble the case of decoupled winds.  SH03 had several motivations for decoupling the winds hydrodynamically. They wanted to calibrate the sub-grid model with observations outside the disc, because they had no hope of resolving the structure of the ISM in their cosmological simulations, and they wanted a recipe that was insensitive to numerical resolution. Our tests clearly demonstrate that their method satisfies both these requirements.  One may question, however, whether the observational uncertainties are not far too large to enable a calibration of the wind velocity and mass loading outside the disc. It is currently not even clear how the observed values depend on radius and gas phase. It is also questionable whether it is desirable for the predictions of hydrodynamical simulations to be insensitive to resolution if the simulations do not resolve the Jeans scale.  Even if one is not interested in the internal structure of galaxies, which agrees less well with observations if the winds are decoupled, there are likely to be important differences in other types of predictions. The fact that hydrodynamic drag makes low-mass galaxies much more diffuse may, for example, greatly affect predictions for quasar absorption line observations and may also alleviate the angular momentum problem (i.e.\\ simulated discs are too small, probably due to excessive transfer of angular momentum to the dark matter halos). The fact that neglecting pressure forces on wind particles within the disc increases the kinetic energy of the escaping gas by very large factors, may have a large impact on predictions for the chemical enrichment of the intergalactic medium.  However, we stress that our simulations lack the resolution and the physics necessary to predict the structure of the multiphase ISM and to model the small-scale effects that ultimately lead to the development of galactic winds. It is also important to note that while our artificial set-up of isolated, thin discs is useful for numerical experiments such as those presented here, it may exaggerate the differences between the coupled and decoupled winds. Galaxies in cosmological simulations are surrounded by gaseous halos and in the SH03 prescription wind particles will be re-coupled to the hydrodynamics before they leave the halos.  Summarizing, our results suggest that (ram) pressure forces in the disc and the inner halo have a very strong impact on the structure of the ISM and the properties of galactic winds. Pressure forces exerted by expanding superbubbles puff up disc galaxies, give low-mass starbursting galaxies irregular morphologies and stabilize the discs of massive galaxies. The energy lost in this process strongly reduces the kinetic energy carried by the outflowing gas. Even if the first bubbles open up a channel in the disc through which the gas can be efficiently ejected, it will run into the gas which was blown out earlier and has been decelerated in the process. For massive galaxies the reduction in the kinetic energy results in the development of a galactic fountain. When the resolution is too low to resolve the Jeans scale, the effects of hydrodynamic drag on the galactic winds will be underestimated."
    },
    "0801/0801.2400_arXiv.txt": {
        "abstract": "We analyze the absolute magnitude ($M_r$) and color ($u-r$) of low redshift ($z < 0.06$) galaxies in the Sloan Digital Sky Survey Data Release 6. Galaxies with nearly exponential profiles (Sloan parameter ${\\rm fracDeV} < 0.1$) all fall on the blue sequence of the color -- magnitude diagram; if, in addition, these exponential galaxies have $M_r < -19$, they show a dependence of $u-r$ color on apparent axis ratio $q$ expected for a dusty disk galaxy. By fitting luminosity functions for exponential galaxies with different values of $q$, we find that the dimming is well described by the relation $\\Delta M_r = 1.27 ( \\log q )^2$, rather than the $\\Delta M \\propto \\log q$ law that is frequently assumed. When the absolute magnitudes of bright exponential galaxies are corrected to their ``face-on'' value, $M_r^f = M_r - \\Delta M_r$, the average $u-r$ color is linearly dependent on $M_r^f$ for a given value of $q$. Nearly face-on exponential galaxies ($q > 0.9$) have a shallow dependence of mean $u-r$ color on $M_r^f$ (0.096 magnitudes redder for every magnitude brighter); by comparison, nearly edge-on exponential galaxies ($q < 0.3$) are 0.265 magnitudes redder for every magnitude brighter. When the dimming law $\\Delta M_r \\propto ( \\log q )^2$ is used to create an inclination-corrected sample of bright exponential galaxies, their apparent shapes are confirmed to be consistent with a distribution of mildly non-circular disks, with median short-to-long axis ratio $\\gamma \\approx 0.22$ and median disk ellipticity $\\epsilon \\approx 0.08$. ",
        "introduction": "\\label{sec-intro}  Luminous galaxies ($M_r < -18$ or so) can be coarsely divided into two classes, conventionally labeled ``early-type'' and ``late-type''. Early-type galaxies have redder stellar populations and a scarcity of interstellar gas and dust. The majority of luminous early-type galaxies are elliptical galaxies, characterized by smooth isophotes and concentrated light profiles, well described by a \\citet{dV48} profile: $\\log I \\propto - r^{1/4}$. Highly luminous elliptical galaxies tend to be mildly triaxial ellipsoids (as opposed to perfectly oblate spheroids); their intrinsic short-to-long axis ratio is typically $c / a \\sim 0.7$ \\citep{ry92,vi05}.  Late-type galaxies have blue stellar populations and relatively large amounts of interstellar gas and dust. The majority of luminous late-type galaxies are spiral galaxies, characterized by spiral structure within flattened disks. The disk light profile is generally well described by an exponential profile: $\\log I \\propto -r$. Luminous spiral galaxies are mildly elliptical (as opposed to perfectly circular) when seen face-on; their intrinsic short-to-long axis ratio is color-dependent, but at visible wavelengths is typically $c / a \\sim 0.25$ \\citep{bi81,gr85,la92,fa93,ry04,ry06}.  In a color-magnitude (CM) diagram, if the color index is chosen correctly, the early-type and late-type galaxies manifest themselves as a ``red sequence'' and a ``blue sequence'', respectively. In the Sloan Digital Sky Survey (SDSS), the distribution of $u-r$ colors for low-redshift galaxies is bimodal; \\citet{st01} find an optimal color separator of $u-r = 2.22$, when color alone is used as a discriminator between early-type and late-type galaxies. Using a full CM diagram, the color separator is found to be dependent on $M_r$ \\citep{ba04}, with the color separator ranging from $u-r \\approx 2.3$ for galaxies with $M_r < -21$ to $u-r \\approx 1.8$ for galaxies with $M_r > -18$.  However, a clean separation between early-type and late-type galaxies using color and absolute magnitude information alone is impossible; the red sequence and blue sequence overlap in a CM diagram. This overlap results partly from the fact that the color and apparent magnitude of spiral galaxies are inclination dependent. Since spiral galaxies are disk-shaped and contain dust, an edge-on spiral is both redder and fainter than the same spiral would be if seen face-on. Thus, as noted by \\citet{al02}, a sample chosen solely by the color criterion $u-r \\geq 2.22$ will contain dust-reddened edge-on spirals as well as intrinsically red ellipticals.  The overlap between the red sequence and blue sequence would be reduced if we could perform an inclination correction on the colors and apparent magnitudes of spiral galaxies; that is, if we could convert observed apparent magnitudes into what they would be if the spiral galaxy were face-on. Since pivoting galaxies so that we can see them face-on is an impracticable task, we will take a statistical approach to finding the average dimming ($\\Delta M_r$) and reddening ($\\Delta (u-r)$) of a spiral galaxy as a function of its inclination $i$.  In addition to allowing a cleaner separation between early-type and late-type galaxies in the CM diagram, a statistical correction for dimming and reddening has other practical uses. For instance, a flux-limited survey will undersample edge-on spirals with respect to face-on spirals; a galaxy that is just above the flux limit when it is face on would fall below the limit if it were edge on. With a knowledge of $\\Delta M_r$ as a function of inclination, it is possible to create an inclination-corrected flux-limited sample, by retaining only those galaxies that would still be above the flux limit if they were tilted to be seen edge on. Since the standard technique for finding the distribution of intrinsic axis ratios of galaxies assumes a random distribution of inclinations (see, for instance, \\citep{vi05} and references therein), such an inclination-corrected sample is essential for determining the true distribution of disk flattening. Without the inclination correction, the scarcity of edge-on galaxies would lead to an overestimate of the typical disk thickness.  In section~\\ref{sec-data}, we describe how we select a sample of SDSS galaxies with exponential profiles; this gives us a population of disk-dominated spiral galaxies. The apparent axis ratio $q$ of the 25 mag arcsec$^{-2}$ isophote is chosen as our surrogate for the inclination of a galaxy. In section~\\ref{sec-analysis}, we examine the luminosity function of the SDSS exponential galaxies as a function of $q$. By shifting the luminosity functions until their high-luminosity cutoffs align, we can estimate the dimming $\\Delta M_r$ of the cutoff as a function of $q$. By brightening each galaxy by the $\\Delta M_r$ appropriate to its observed value of $q$, we can find its approximate ``face-on'' absolute magnitude $M_r^f$. We then provide linear fits for the mean $u-r$ color as a function of $M_r^f$ for different ranges of $q$. In section~\\ref{sec-shapes}, we create an inclination-corrected flux-limited sample, which we then use to find the distribution of intrinsic short-to-long axis ratios of the SDSS exponential galaxies, confirming that they are, in fact, disks, as our analysis assumed all along. Finally, in section~\\ref{sec-disc}, we provide a brief discussion of what the form of $\\Delta M_r$ as a function of $q$ and of $\\Delta (u-r)$ as a function of $q$ and $M_r^f$ implies for the properties of spiral galaxies and the dust they contain. ",
        "conclusions": "\\label{sec-disc}  We have selected a population of galaxies from the Sloan Digital Sky Survey which are at low redshift ($z < 0.06$), which are relatively luminous ($M_r \\leq -19$), and which are well described by an exponential surface brightness profile (${\\rm fracDeV} < 0.1$). As we have shown, the properties of these galaxies are consistent with their being a population of slightly elliptical disks containing dust. The median dimensionless disk thickness for these galaxies in the $r$ band is $\\gamma \\approx 0.22$; the median disk ellipticity is $\\epsilon \\approx 0.08$.  By fitting the luminosity function for galaxies with different apparent axis ratio $q$, we found that the apparent dimming $\\Delta M_r$ is not linearly proportional to $\\log q$, but instead is much better fitted by $\\Delta M_r \\propto ( \\log q )^2$. The dependence of dimming on inclination is a valuable clue to the dust properties within disk galaxies. If certain simplifying assumptions are made, the expected attenuation as a function of inclination can be computed for model galaxies. For instance, \\citet{fe99} assumed that dust had either the extinction curve found for Milky Way dust or for Small Magellanic Cloud dust \\citep{go97}. They assumed that disks were perfectly axisymmetric, with a horizontal scale length $r = 4 {\\rm\\,kpc}$ that was the same for both stars and dust. The scale height of the stars was assumed to be $z_\\star = 0.35$, but the scale height of the dust was allowed to vary. \\citet{fe99} found that dimming in the $B$ and $I$ bands were proportional to $\\log q$ only when the dust scale height was greater than that of the stars. However, observation of nearby edge-on disk galaxies \\citep{xi99} indicates that the dust scale height is about half the star scale height.  \\citet{ro08}, using a Monte-Carlo radiative-transfer code to make calculations of internal extinction in dusty galaxies, found that a quadratic dependence of dimming on $\\log q$ provides a good fit for all plausible dust scale heights, scale lengths, and metallicity gradients. Since they were using hydrodynamic galaxy models with spiral structure, they were able to confirm that nonaxisymmetric structures such as spiral arms did not significantly affect the dependence of the total dimming $\\Delta M$ on the apparent axis ratio $q$.  In our sample of exponential galaxies from the Sloan Digital Sky Survey, once the absolute magnitude of a galaxy is corrected for the inclination-dependent dimming, the mean $u-r$ color observed is linearly dependent on the corrected $M_r^f$. For nearly face-on galaxies, with $q > 0.9$, the dependence of $u-r$ on $M_r^f$ is relatively weak. We find $b = -0.096$; that is, less than 0.1 magnitude of reddening in $u-r$ for each magnitude brighter in $M_r^f$. For the edge-on galaxies, with $q < 0.3$, the dependence of $u-r$ on $M_r^f$ is much stronger, with $b = -0.265$. The mean color -- absolute magnitude dependence is a manifestation of the metallicity -- luminosity dependence. High-metallicity galaxies have both redder stellar populations and higher dust contents. For face-on exponential galaxies, the dust effects are minimized, and we see the effect of metallicity on stellar populations. For edge-on exponential galaxies, we see, in addition, the effect of metallicity on the dust content. Exponential galaxies with $M_r^f \\sim -21.5$ have a typical color $\\langle u - r \\rangle \\sim 1.8$ when seen face-on, placing them at the tip of the blue sequence in a color --  magnitude diagram. However, the same bright exponential galaxies when seen edge-on will have $M_r \\sim -21$ and $\\langle u -r \\rangle \\sim 2.5$, a degree of reddening that smuggles them into the red sequence, as usually defined.    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 website is \\url{http://www.sdss.org/}. The SDSS is managed by the Astrophysical Research Consortium (ARC) 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, The Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington.  \\newpage"
    },
    "0801/0801.3226_arXiv.txt": {
        "abstract": "The dynamical interactions of planetary systems may be a clue to their formation histories. Therefore, the distribution of these interactions provides important constraints on models of planet formation. We focus on each system's apsidal motion and proximity to dynamical instability. Although only $\\sim$25 multiple planet systems have been discovered to date, our analyses in these terms have revealed several important features of planetary interactions. 1) Many systems interact such that they are near the boundary between stability and instability. 2) Planets tend to form such that at least one planet's eccentricity periodically drops to near zero. 3) Mean-motion resonant pairs would be unstable if not for the resonance. 4) Scattering of approximately equal mass planets is unlikely to produce the observed distribution of apsidal behavior. 5) Resonant interactions may be identified through calculating a system's proximity to instability, regardless of knowledge of angles such as mean longitude and longitude of periastron (\\eg GJ 317 b and c are probably in a 4:1 resonance). These properties of planetary systems have been identified through calculation of two parameters that describe the interaction. The apsidal interaction can be quantified by determining how close a planet is to an apsidal separatrix (a boundary between qualitatively different types of apsidal oscillations, \\eg libration or circulation of the major axes). This value can be calculated through short numerical integrations. The proximity to instability can be measured by comparing the observed orbital elements to an analytic boundary that describes a type of stability known as Hill stability. We have set up a website dedicated to presenting the most up-to-date information on dynamical interactions: http://www.lpl.arizona.edu/$\\sim$rory/research/xsp/dynamics. ",
        "introduction": "\\label{sec:intro}  One of the most striking differences between known exoplanets and the giant planets of our Solar System involves the observed orbits: Observed exoplanets tend to have large eccentricities $e$, and small semi-major axes $a$, whereas the gas giants of the Solar System have small $e$ and large $a$. Recently, however, it has been shown that the dynamical interactions in many multiple planet systems (including the Solar System) show certain features in common (Barnes \\& Quinn 2004; Barnes \\& Greenberg 2006a [BG06a]; Barnes \\& Greenberg 2006c [BG06c]; Barnes \\& Greenberg 2007b [BG07b]). These shared traits suggest that the character of dynamical interactions (over $10^3$ -- $10^4$ years), rather than the present orbits, may be a more meaningful constraint on the origins of planetary systems (Barnes \\& Greenberg 2007a [BG07a], BG07b). Considerations of shared dynamical properties have even resulted in the first successful prediction of the mass and orbit of an extrasolar planet, HD 74156 d (predicted by Barnes \\& Raymond [2004] and Raymond \\& Barnes [2005]; found by Bean \\etal 2008).  Now with over 200 extra-solar planets known, including $\\sim 25$ multi-planet systems, we can look at this population as a whole, with increasing confidence regarding which common characteristics may be more than statistical flukes. We have identified commonalities among planetary interactions that may be key constraints on the origins of planetary systems.  We have identified parameters that quantify an interaction's proximity to boundaries between qualitatively different types of motion. These two types of boundaries are the ``apsidal separatrix'' and the dynamical stability boundary. The apsidal separatrix is the boundary between different types of apsidal oscillations, \\eg libration and circulation. The stability boundary separates regions in which all planets are bound to the host star from those in which at least one planet is liable to be ejected. These two parameters are fixed quantities that do not vary over time, but they constrain how the systems evolve. Systems tend to lie close to these two boundaries.  In this chapter we review our derivation of $\\epsilon$, which quantifies proximity to an apsidal separatrix, in $\\S$ 2. Then we describe $\\beta$, which parameterizes a system's proximity to dynamical instability, in $\\S$ 3. In $\\S$ 4 we present the observed distributions of these quantities and use them to constrain formation models. Finally in $\\S$ 5 we draw our general conclusions. ",
        "conclusions": "\\label{sec:conclusions} We have described here new approaches for parameterizing the dynamical interactions of extrasolar planet systems. Many systems interact such that they are near the stability limit and the apsidal separatrix.  Most mean-motion resonance interactions would be unstable if not for the resonance, a feature which may be used to identify resonant interactions (\\eg GJ 317 b and c). The distribution of $\\epsilon$ shows that the scattering of approximately equal mass planets is unlikely to produce the observed distribution of apsidal behavior. The distribution of $\\beta$ values shows that many systems formed near the limit of dynamical stability.  These results demonstrate the benefits of the consideration of the dynamical interactions of multiple planet systems, which may constrain models of planet formation. For example, as we discussed above, the proximity of so many systems to an apsidal separatrix ($\\epsilon = 0$) suggests that ``rogue planets'' may have played a major role in scattering planets to higher eccentricities. The proximity of so many systems to the stability limit ($\\beta = 1$) suggests many systems form in densely packed configurations. Moreover, consideration of planetary systems in these terms has revealed that our Solar System shares dynamical traits with the known multiple planet systems. Perhaps constraining planet formation models through dynamical properties, rather than observed orbital elements, will lead to a universal model of planet formation.   About half of planetary systems are multiple (Wright \\etal 2007), predictions of additional companions are being borne out (Barnes \\& Raymond 2004; Raymond \\& Barnes 2005; Bean \\etal 2008), and the current distribution of planet masses suggest there will be many planets with a mass equal to that of Saturn or less (Marcy \\etal 2005), \\ie below current detection limits. These three observations imply many multiple planet systems will be detected in the future. Hence characterizing extrasolar planet interactions will be a critical aspect of the study of planet formation for the foreseeable future."
    },
    "0801/0801.1223_arXiv.txt": {
        "abstract": "CH$_4$ is proposed to be the starting point of a rich organic chemistry. Solid CH$_4$ abundances have previously been determined mostly toward high mass star forming regions. Spitzer/IRS now provides a unique opportunity to probe solid CH$_4$ toward low mass star forming regions as well. Infrared spectra from the Spitzer Space Telescope are presented to determine the solid CH$_4$ abundance toward a large sample of low mass young stellar objects. 25 out of 52 ice sources in the $c2d$ (cores to disks) legacy have an absorption feature at 7.7 $\\mu$m, attributed to the bending mode of solid CH$_4$. The solid CH$_4$ / H$_2$O abundances are 2-8\\%, except for three sources with abundances as high as 11--13\\%. These latter sources have relatively large uncertainties due to small total ice column densities. Toward sources with H$_2$O column densities above 2$\\times 10^{18}$ cm$^{-2}$, the CH$_4$ abundances (20 out of 25) are nearly constant at 4.7$\\pm$1.6\\%. Correlation plots with solid H$_2$O, CH$_3$OH, CO$_2$ and CO column densities and abundances relative to H$_2$O reveal a closer relationship of solid CH$_4$ with CO$_2$ and H$_2$O than with solid CO and CH$_3$OH. The inferred solid CH$_4$ abundances are consistent with models where CH$_4$ is formed through sequential hydrogenation of C on grain surfaces. Finally the equal or higher abundances toward low mass young stellar objects compared with high mass objects and the correlation studies support this formation pathway as well, but not the two competing theories: formation from CH$_3$OH and formation in gas phase with subsequent freeze-out. ",
        "introduction": "The presence and origin of complex organic molecules in protostellar regions and their possible incorporation in protoplanetary disks is an active topic of research. CH$_4$ is proposed to be a starting point of a rich chemistry, especially when UV photons are present \\citep{dartois05}. In particular CH$_4$ is believed to play a key role in the formation process of prebiotic molecules \\citep{markwick00}.  CH$_4$ is less well studied in interstellar and circumstellar media compared to other small organic molecules because CH$_4$ has no permanent dipole moment and therefore cannot be observed by pure rotational transitions at radio wavelengths. Solid CH$_4$ was first detected through its bending mode at 7.67 $\\mu$m from the ground by \\citet{lacy91}, and with the Infrared Space Observatory Short Wavelength Spectrometer (ISO-SWS) by \\citet{boogert96} toward a few high-mass sources. Tentative claims have been made toward some other objects including low-mass protostars, but are inconclusive because of the low S/N ratio of these data \\citep{cernicharo00, gurtler02, alexander03}. Solid CH$_4$ has also been detected from the ground through its stretching mode at 3.3 $\\mu$m, but only toward the brightest high mass sources due to problems in removing the many atmospheric lines in this spectral region \\citep{boogert04}.  Models predict CH$_4$ to form rapidly on cool grains through successive hydrogenation of atomic C; similarly H$_2$O is formed through hydrogenation of atomic O \\citep{vandehulst46, allen77, tielens82,brown88, hasegawa92,aikawa05}. Observations of CH$_4$ hence provide insight into the basic principles of grain surface chemistry. Compared to H$_2$O the observed gas- and solid-state CH$_4$ abundances are low; reported CH$_4$ abundances are typically a few percent with respect to H$_2$O \\citep{lacy91, boogert98}. This points to relatively low atomic C abundances at the time of CH$_4$ formation, with most C already locked up in CO as H readily reacts with C on surfaces \\citep{hiraoka98}. This is in agreement with the high CH$_3$OH abundances in several lines of sight, formed by hydrogenation of CO \\citep{dartois99,pontoppidan03}, and large CO$_2$ abundances, formed through oxidation of CO or hydrogenated CO. That these molecules are all formed through a similar process is corroborated by the profiles of solid CO$_2$ absorption bands, which usually show an intimate mixture of CO$_2$, CH$_3$OH and H$_2$O in  interstellar ices \\citep{gerakines99, boogert00, knez05}.  If CH$_4$ is formed efficiently through grain surface reactions as well, CH$_4$ should be similarly mixed with H$_2$O. Observations of solid CH$_4$ toward a few high mass young stellar objects (YSOs) show that the CH$_4$ absorption band profiles are broad and agree better with CH$_4$ in a hydrogen bonding ice, H$_2$O or CH$_3$OH, than with a pure CH$_4$ ice or CH$_4$ mixed with CO \\citep{boogert97}. This profile analysis does not, however, exclude CH$_4$ formation from photoprocessing of CH$_3$OH \\citep{allamandola88, gerakines96}. In addition, because of the small sample in previous studies, it is unclear if these broad profiles are a universal feature. Hence it cannot be excluded that CH$_4$ in some environments may form in the gas phase and subsequently freeze out.  Because formation pathway efficiency depends on environment, another method for testing formation routes is through exploring the distribution of CH$_4$ toward a large sample of objects of different ages, luminosities and ice column densities. In addition, correlations, or lack thereof, with other ice constituents may provide important clues to how the molecule is formed. If CH$_4$ is formed through hydrogenation on grain surfaces in quiescent clouds, the CH$_4$ abundance with respect to H$_2$O should be fairly source independent since this mechanism mainly depends on the initial conditions of the cloud before the star forms, which seem to vary little between different star forming regions. Because of the generality of this mechanism, the CH$_4$ and H$_2$O, and possibly CO$_2$, column densities should correlate over a large range of different environments. This is also the prediction of several models where the solid CH$_4$/H$_2$O and CH$_4$/CO$_2$ ratios in dark clouds vary little both as a function of time \\citep{hasegawa92} and distance into a cloud that has collapsed \\citep{aikawa05}. In contrast, CO, which is formed in the gas phase and subsequently frozen out, is predicted to only correlate with CH$_4$ during certain time intervals. If CH$_4$ instead forms in the gas phase, solid CH$_4$ should be better correlated with solid CO than with solid H$_2$O and CO$_2$, since pure CH$_4$ freezes out and desorbs at similar temperatures to CO \\citep[Fraser et al. submitted to $A\\&A$,][]{collings04}. Finally, if CH$_4$ forms by UV photoprocessing of CH$_3$OH more CH$_4$ would also be expected to form toward sources with stronger UV-fields i.e. higher mass objects.  The objective of this study is to determine the CH$_4$ abundances and distribution pattern toward a sample of low mass young stellar objects, varying in evolutionary stage and total ice column density. The distribution pattern and correlations with other ice constituents within the sample as well as comparison with high mass young stellar objects will be used to constrain the CH$_4$ formation mechanism.  This study is based on spectra acquired with the Spitzer Infrared Spectrometer (IRS) as part of our legacy program `From molecular cores to protoplanetary disks' ($c2d$), which provides a large sample (41 sources) of infrared spectra of low mass star formation regions \\citep{evans03}. In addition 11 sources are added from the GTO program 2 for which ground based 3-5 $\\mu$m already exist \\citep{pontoppidan03}. Overviews of the H$_2$O, CO$_2$, CH$_3$OH and other ice species in these data are found in Boogert et al. (ApJ submitted, Paper I) and Pontoppidan et al. (ApJ submitted, Paper II). The detection of solid CH$_4$ toward one of the sources, HH46 IRS, was published by \\citet{boogert04b}. We have detected an absorption feature, which is attributed to solid CH$_4$, toward 25 out of 52 low mass ice sources found in this $c2d$ sample. ",
        "conclusions": "We present Spitzer-IRS spectra of the solid CH$_4$ feature at 7.7 $\\mu$m toward a large sample of low mass young stellar objects. Our conclusions are as follows: \\begin{itemize} \\item 25 out of 52 low mass young stellar objects show a solid CH$_4$ feature at 7.7 $\\mu$m. \\item The solid CH$_4$ abundance with respect to H$_2$O is centered at 5.8\\% with a standard deviation of 2.7\\% in the sources with CH$_4$ detections. In the sources without detections the average upper limit is 15\\%, which is not significant compared with the detections. \\item The sources (two Ophiuchus and one Serpens) with more than 10\\% CH$_4$ all have H$_2$O column densities below 2$\\times 10^{18}$ cm$^{-2}$. Due to the low total column densities, two  of these three sources have uncertainities larger than 50\\%. Above 2$\\times 10^{18}$ cm$^{-2}$ the sources (20 out of 25) have a fairly constant CH$_4$ abundance of 4.7$\\pm$1.6\\%. \\item The 7.7$\\mu$m feature profiles are signficantly broader for all but one object than expected for pure solid CH$_4$ and toward most sources also broader than expected for CH$_4$ in H$_2$O dominated ices. Approximately 30\\% of the features have a blue wing, seen previously toward high mass YSOs and there attributed to solid SO$_2$ \\item The column densities of solid CH$_4$ and H$_2$O and CO$_2$ are clearly correlated, while CH$_4$ and CO and CH$_3$OH are only weakly correlated. \\item There is also no correlation between the CH$_4$ and CO abundances when both have been normalized to the H$_2$O abundance. \\item The Ophiuchus cloud has significantly higher CH$_4$ abundances compared to the rest of the sample, probably due to the low total column densities towards several of the sources. There are no significant differences between the remaining clouds. \\item The abundance variation is smaller for CH$_4$ compared to solid CH$_3$OH; CH$_4$ seems to belong to the class of molecules, also including H$_2$O and CO$_2$ that appear 'quiescent', i.e. their abundances are more or less constant, in contrast to highly variable ices like CH$_3$OH and OCN$^-$. If the Ophiuchus sources are included CH$_4$ is somewhere between the two classes. \\item Sample statistics and comparison with model predictions support CH$_4$ formation through hydrogenation of C on grain surfaces. \\end{itemize}"
    },
    "0801/0801.4406_arXiv.txt": {
        "abstract": "We calculate the contribution of cosmic strings arising from a phase transition in the early universe, or cosmic superstrings arising from brane inflation, to the cosmic 21 cm power spectrum at redshifts $z\\geq 30$. Future experiments can exploit this effect to constrain the cosmic string tension $G\\mu$ and probe virtually the entire brane inflation model space allowed by current observations. Although current experiments with a collecting area of $\\sim 1$ $\\rm{km}^2$ will not provide any useful constraints, future experiments with a collecting area of $10^4-10^6$ $\\rm{km}^2$ covering the cleanest $10\\%$ of the sky can in principle constrain cosmic strings with tension $G\\mu \\gtrsim 10^{-10}-10^{-12}$ (superstring/phase transition mass scale $>10^{13}$ GeV). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.1686_arXiv.txt": {
        "abstract": "We present and constrain a cosmological model which component is a pressureless fluid with bulk viscosity as an explanation for the present accelerated expansion of the universe. We study the particular model of a constant bulk viscosity coefficient $\\zeta_{{\\rm m}}$. The possible values of $\\zeta_{{\\rm m}}$ are constrained using the cosmological tests of SNe Ia Gold 2006 sample, the CMB shift parameter $R$ from the three-year WMAP observations, the Baryon Acoustic Oscillation (BAO) peak $A$ from the Sloan Digital Sky Survey (SDSS) and the Second Law of Thermodynamics (SLT). It was found that this model is in agreement with the SLT using only the SNe Ia test. However when the model is submitted to the three cosmological tests together (SNe+CMB+BAO) the results are:  1.- the model violates the SLT, 2.- predicts a value of $H_0 \\approx 53 \\; {\\rm km \\cdot sec^{-1} \\cdot Mpc^{-1}}$ for the Hubble constant, and 3.- we obtain a bad fit to data with a $\\chi^2_{{\\rm min}} \\approx 400$ ($\\chi^2_{{\\rm d.o.f.}} \\approx 2.2$). These results indicate that this model is ruled out by the observations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.4776_arXiv.txt": {
        "abstract": "We analyze the mid-infrared (MIR) spectra of ultraluminous infrared galaxies (ULIRGs) observed with the \\textit{Spitzer Space Telescope}'s Infrared Spectrograph. Dust emission dominates the MIR spectra of ULIRGs, and the reprocessed radiation that emerges is independent of the underlying heating spectrum. Instead, the resulting emission depends sensitively on the geometric distribution of the dust, which we diagnose with comparisons of numerical simulations of radiative transfer. Quantifying the silicate emission and absorption features that appear near 10 and 18\\um{} requires a reliable determination  of the continuum, and we demonstrate that including a  measurement of the  continuum at intermediate wavelength (between the features) produces accurate results at all optical depths. With high-quality spectra, we successfully use the silicate features to constrain the dust chemistry. The observations of the ULIRGs and local sightlines require dust that has a relatively high 18/10\\um{} absorption ratio of the silicate features (around 0.5). Specifically, the cold dust of Ossenkopf et al. (1992) % is consistent with the observations, while other dust models are not. We use the silicate feature strengths to identify two families of ULIRGs, in which the dust distributions are fundamentally different. Optical spectral classifications are related to these families. In ULIRGs that harbor an active galactic nucleus, the spectrally  broad lines are detected only when the nuclear surroundings are clumpy. In contrast, the sources of lower ionization optical spectra are deeply embedded in smooth distributions of optically thick dust. ",
        "introduction": "High infrared luminosity ($L_{IR} > 10^{12} L_\\sun$) characterizes ultraluminous infrared galaxies (ULIRGs). The underlying energy source may be accretion onto a supermassive black hole in an active galactic nucleus (AGN), intense bursts of star formation, or a combination of the two \\citep{San88a,Gen98}. Because the total luminosities are large, ULIRGs are critical sites to consider in obtaining a complete account of star formation, black hole  growth, and the relationship between these two phenomena over cosmic time. In all cases, dust is responsible for reprocessing the intrinsic hard radiation to emerge at longer wavelengths. This dust reprocessing erases the spectral signatures that would reveal the nature of the original source, but the resulting spectra serve as probes of the dust itself.  While most of the dust emission is continuum radiation, in the mid-infrared (MIR), two of its spectral features are observable,  arising near 10 and 18\\um. These are attributed to silicates, especially amorphous analogs of pyroxenes and olivines \\citep[and references therein]{Dra03r}. Both features have been measured in stellar observations for decades \\citep*{Gil68,Woo69,Low70}. Observations of the stronger 10\\um{} absorption feature in starburst galaxies and AGNs have a similarly long history \\citep*[\\textit{e.g.},][]{Gil75,Rie75a,Rie75b,Kle76}, but only recently has the high sensitivity of instruments in space made measurement of the 18\\um{} feature in large numbers of galaxies feasible.  Detecting silicate features in emission  can identify the dusty medium as optically thin, and deep absorption characterizes optically thick obscuration. In detail, however, quantifying the optical depth requires a complete radiative transfer calculation because the same material produces both the continuum and line features. Thus, the apparent optical depth that silicate absorption strength yields is not the true optical depth along the line of sight, although it is frequently treated as such a direct measurement \\citep[\\textit{e.g.},][]{Shi06,Gal04,Mai01}. Compared with complete calculations, however, high signal-to-noise spectra of both of these features together effectively reveal the geometric distribution of the dust, including the total line-of-sight optical depth. Moreover, while the continuum emission depends only weakly on the dust chemistry, comparison of the two silicate features provides one of the few accessible diagnostics of the dust composition.  Infrared studies that aim to deduce the nature of sources and surrounding dust geometry suffer from the degeneracy of the problem: a  large range of different configurations produces very similar overall spectral energy distributions. Reaching conclusions about geometry, let alone dust chemistry, based on the spectral shape is difficult if not impossible \\citep[\\textit{e.g.}][]{Vin03}. However, extremely deep  10\\um{} silicate absorption proves to be a good discriminant of clumpy and smooth dust distributions, as we show in \\citet{Lev07}. Here we extend this analysis to include both the 10 and 18\\um{} silicate features. We apply detailed models to MIR spectral measurements of ULIRGs to probe the dust chemistry and geometric distribution of material around the galaxies' powerful nuclear sources. While the MIR spectra of ULIRGs have been studied previously \\citep[\\textit{e.g.},][]{Arm04,Arm06,Arm07,Des07,Far07,Mar07,Spo04,Spo06}, here we investigate both silicate features in detail.  We consider the ULIRGs contained in the samples of \\citet{Spo07} and \\citet{Ima07}, which were all observed using the Infrared Spectrograph (IRS; \\citealt{Hou04}) on board the \\textit{Spitzer Space Telescope} \\citep{Wer04}. All the observations were obtained in low-resolution mode, which provides spectral resolution $R\\sim 100$ and complete coverage from 5--35\\um. The slits are up to 3.7 and 10.7\\arcsec{} wide in the short- and long-wavelength modules, respectively. Although these slits cover large areas of the galaxies (which have redshifts up to 0.93), the concentrated emission dominates the flux of the nuclear spectra we discuss. See \\citet{Spo07} for additional information about the data reduction. ",
        "conclusions": "The large luminosity of ULIRGs requires that these galaxies possess both an underlying source that is more energetic than star formation within ordinary galaxies and dust to reprocess the intrinsic radiation to long wavelengths. Observations in the IR offer the obvious advantage of directly detecting the light that emerges. Although they suffer the disadvantage that dust reprocessing erases the detailed signature of the nature of the energy source, as we have shown quantitatively, MIR spectroscopy does directly probe the dust itself. Analysis of the silicate features in particular provides useful diagnostics of the geometric distribution of dust and its optical properties, which ultimately depend on its mineralogy.  Comparing observations of ULIRGs with several models of dust chemistry, we find that numerical simulations of radiative transfer employing the observationally-motivated OHMc dust properties match the data well, and this dust is required when deep absorption is measured. The key difference between this model and others that are frequently employed is the relatively high silicate absorption peak near 18\\um{} relative to that near 10\\um; this strength ratio is around 0.5 in the OHMc dust. These properties are physically-motivated, based on laboratory measurements of amorphous silicates. Including the effects of grain oxidation results in the higher ratio of 18/10\\um\\ cross section, which is characteristic of bronzite \\citep[and references therein]{Oss92}. This dust chemistry is not peculiar to the ULIRGs or to the nuclei of galaxies, but also fits observations within the Galaxy and Magellanic Clouds. Moreover, local observations where the silicates appear in emission rule out the OHMw and  Draine dust at low optical depths.  Quantifying the strength of the silicate features requires determining  the underlying continuum accurately. We demonstrate that including measurements of the continuum at intermediate wavelengths between the features yields physically consistent results, and we recommend specific modifications to the idealized method when observational complications, such as photoionized emission or ice absorption bands, are present. We emphasize that because the same dust is responsible for both the feature and continuum, the apparent optical depth of silicate absorption does not reveal the line of sight optical depth, independent of the continuum-fitting procedure.  We identify two distinct families in the ULIRG measurements, which is a consequence of fundamental differences in obscuring geometry. One class shows deep absorption ($S_{10}< -1$). These energetic sources must be deeply embedded in a continuous medium that is geometrically and optically thick ($\\tau_V > 100$). The other class exhibits weak silicate features, which may appear in absorption or emission. Either a geometrically thin slab or a clumpy  dusty medium can account for these observations, but they do not represent the low optical depth realm of the continuous geometry. In addition, we find  that these two ULIRG families of geometry are related to the optical identification of the galaxy nuclei. All members of the second (low-strength) class are AGNs, while ULIRGs having spectra characteristic of LINERs and \\ion{H}{2} regions systematically exhibit the deepest absorption. Among the deeply embedded sources, the MIR emission and optical lines always emanate from different regions. Optically-identified LINERs may be a heterogeneous class, but the consistent finding of \\ion{H}{2} ULIRGs only in the deeply embedded group leads us to speculate that two different regions dominate the IR and optical emission line signatures in all such optical \\ion{H}{2} ULIRGs.  In spite of the general degeneracy of IR SEDs \\citep[\\textit{e.g.},][]{Vin03}, extreme absorption of the 10\\um{} feature alone proves to be a strong  indicator of smooth (as opposed to clumpy)  dust distributions \\citep{Lev07}. Here we find that the combination of 10 and 18\\um{} silicate features together constrains the geometry even more powerfully in a wider range of  situations. Furthermore, the analysis tool of the feature-feature diagram we introduce here provides a strong diagnostic of the dust chemistry. The optical properties fix the absolute zero points and the slopes of the linear trends of silicate strength that variations of the model parameters can produce. However, because analysis based on this diagram  is sensitive only to the feature strengths in the absorption cross section, it does not probe feature shape, peak wavelength, or properties such as mineral composition directly. We conclude that of the available models, the OHMc dust best describes cosmic silicate feature strengths, both in ULIRGs and local sources."
    },
    "0801/0801.0964_arXiv.txt": {
        "abstract": "We present the work of an international team at the International Space Science Institute (ISSI) in Bern that worked together to review the current observational and theoretical status of the non-virialised X-ray emission components in clusters of galaxies. The subject is important for the study of large-scale hierarchical structure formation and to shed light on the \"missing baryon\" problem. The topics of the team work include thermal emission and absorption from the warm-hot intergalactic medium, non-thermal X-ray emission in clusters of galaxies, physical processes and chemical enrichment of this medium and clusters of galaxies, and the relationship between all these processes. One of the main goals of the team is to write and discuss a series of review papers on this subject. These reviews are intended as introductory text and reference for scientists wishing to work actively in this field. The team consists of sixteen experts in observations, theory and numerical simulations. ",
        "introduction": "Clusters of galaxies are the largest gravitationally bound structures in the Universe. Their baryonic composition is dominated by hot gas that is in quasi-hydrostatic equilibrium within the dark matter dominated gravitational potential well of the cluster. The hot gas is visible through spatially extended thermal X-ray emission, and it has been studied extensively both for assessing its physical properties and also as a tracer of the large-scale structure of the Universe.  Clusters of galaxies are not isolated entities in the Universe: they are connected through a filamentary cosmic web. Theoretical predictions indicate the way this web is evolving. In the early Universe most of the gas in the web was relatively cool ($\\sim 10^4$~K) and visible through numerous absorption lines, designated as the so-called Ly$\\alpha$ forest. In the present Universe, however, about half of all the baryons are predicted to be in a warm phase ($10^5-10^7$~K), the Warm-Hot Intergalactic Medium (WHIM), with temperatures intermediate between the hot clusters and the cool absorbing gas causing the Ly$\\alpha$ forest.  The X-ray spectra of clusters are dominated by the thermal emission from the hot gas, but in some cases there appears to be evidence for hard X-ray tails or soft X-ray excesses. Hard X-ray tails are difficult to detect, and one of the topics for the team is a discussion on the significance of this detection (yet contradictory) in existing and future space experiments. Various models have been proposed to produce these hard X-ray tails, and our team reviews these processes in the context of the observational constraints in clusters.  While in some cases soft excesses in clusters can be explained as the low-energy extension of the non-thermal hard X-ray components mentioned above, there is evidence that a part may also be due to thermal emission from the WHIM. The signal seen near clusters then originates in the densest and hottest parts of the WHIM filaments, where the accelerating force of the clusters is highest and heating is strongest. A strong component of this emission is line radiation from highly ionised oxygen ions, and the role of this line emission and its observational evidence will be reviewed.  WHIM filaments not only can be observed because of their continuum or line emission, but also through absorption lines if a sufficiently strong continuum background source is present. The evidence for absorption in both UV and X-ray high-resolution spectra is discussed. Future space missions will be well adapted to study these absorption lines in more detail.  In particular in absorption lines the lower density parts between clusters become observable. In these low density regions of the WHIM not only collisional ionisation but also photo-ionisation is an important process. In general, the physics of the WHIM is challenging due to its complexity since there are many uncertain factors including the heating and cooling processes, the chemical enrichment, the role of supernova-driven bubbles or starburst winds, ram-pressure stripping, the role of shocks, magnetic fields, etc. More detailed (and sophisticated) hydrodynamical simulations with state-of-the-art spatial (and temporal) resolution are required in order to follow the impact of some (if not all) of these important processes. In particular chemical enrichment is an important process to consider as it leads to many observable predictions. We review the various physical processes relevant for the WHIM, the methods that are used to simulate this and the basic results from those models. ",
        "conclusions": ""
    },
    "0801/0801.2759.txt": {
        "abstract": "We conducted a high resolution mid-infrared spectroscopic investigation using {\\it Spitzer} of 32 late-type (Sbc or later) galaxies that show no definitive signatures of Active Galactic Nuclei (AGN) in their optical spectra in order to search for low luminosity and/or embedded AGN.  These observations reveal the presence of the high ionization [NeV] 14$\\mu$m and/or 24$\\mu$m line in 7 sources, providing strong evidence for AGNs in these galaxies.  Taking into account the variable sensitivity of our observations, we find that the AGN detection rate based on mid-infrared diagnostics in optically normal late-type galaxies is $\\sim$ 30\\%, implying an AGN detection rate in late-type galaxies that is possibly 4 times larger than what optical spectroscopic observations alone suggest. We demonstrate using photoionization models with both an input AGN and an extreme EUV-bright starburst ionizing radiation field that the observed mid-infrared line ratios in our 7 AGN candidates cannot be replicated unless an AGN contribution, in some cases as little as 10\\% of the total galaxy luminosity, is included. These models show that when the fraction of the total luminosity due to the AGN is low, optical diagnostics are insensitive to the presence of the AGN.  In this regime of parameter space, the mid-infrared diagnostics offer a powerful tool for uncovering AGN missed by optical spectroscopy.   The AGN bolometric luminosities in our sample inferred using our [NeV] line luminosities range from $\\sim$ 3$\\times$10$^{41}$ ergs s$^{-1}$ to $\\sim$ 2$\\times$10$^{43}$ ergs s$^{-1}$.  Assuming that the AGN is radiating at the Eddington limit, this range corresponds to a lower mass limit for the black hole that ranges from $\\sim$ 3$\\times$10$^3$M$_{\\odot}$ to as high as $\\sim$ 1.5$\\times$10$^5$M$_{\\odot}$.  These lower mass limits however do not put a strain on the well-known relationship between the black hole mass and the host galaxy's stellar velocity dispersion established in predominantly early-type galaxies.  Our findings add to the growing evidence that black holes do form and grow in low-bulge environments and that they are significantly more common than optical studies indicate. ",
        "introduction": "The vast majority of {\\it currently} known active galactic nuclei (AGN) in the local Universe reside in host galaxies with prominent bulges (e.g. Heckman 1980a; Keel 1983b; Terlevich, Melnick, \\& Moles 1987; Ho, Filippenko, \\&  Sargent 1997; Kauffmann et al. 2003). This result, together with the finding that the black hole mass, M$_{\\rm BH}$, and the host galaxy's stellar velocity dispersion, $\\sigma$, is strongly correlated (Gebhardt et al. 2000; Ferrarese \\&  Merritt 2000), has led to the general consensus that  black hole formation and growth is fundamentally connected to the build-up of galaxy bulges.  Indeed, it has been proposed that feedback from the AGN regulates the surrounding star formation in the host galaxy (e.g. Silk \\& Rees 1998; Kauffmann \\& Haehnelt 2000).  However, an important outstanding question remains unresolved: {\\it is a bulge in general a necessary ingredient for a black hole to form and grow}?  On the one hand, M33, the best studied nearby bulgeless galaxy, shows no evidence of a supermassive black hole (SBH), and the upper limit on the mass is significantly below that predicted by the M$_{\\rm BH}$-$\\sigma$ relation established in early-type galaxies (Gebhardt et al. 2001; Merritt et al. 2001).  On the other hand, both NGC 4395 (Filippenko \\& Ho 2003) and POX52 (Barth et al. 2004) show no evidence for a bulge and yet do contain AGN.  However, these two galaxies have remained isolated cases of bulgeless galaxies with accreting black holes, suggesting that they are anomalies.  Indeed in the extensive Palomar optical spectroscopic survey of 486 nearby galaxies (Ho, Filippenko, \\& Sargent1997; henceforth H97), there are only 9 optically identified Seyferts with Hubble type of Sc or later.  Only one galaxy of Hubble type of Scd or later in the entire survey is classified as a Seyfert (NGC 4395).  Greene \\& Ho (2004, 2007) recently conducted an extensive search for broad-line AGN with intermediate mass black holes in the Fourth Data Release of the Sloan Digital Sky Survey (2004, 2007). Of the 8435 broad line AGN, they found only 174 (2\\%) such intermediate mass objects, indicating that they are extremely rare.  Forty percent of these seem to be in late-type galaxies with colors consistent with morphological type Sab, although the Sloan images are of insufficient spatial resolution to extract precise morphological information.  Optical observations thus clearly suggest that AGN in late-type galaxies are uncommon.   However, determining if AGN reside in low bulge galaxies cannot be definitively answered using optical observations alone. The problem arises because a putative AGN in a galaxy with a minimal bulge is likely to be both energetically weak and deeply embedded in the center of a dusty late-type spiral.  As a result, optical emission lines will be dominated by the emission from star formation regions, severely limiting the diagnostic power of optical surveys in determining the incidence of accreting black holes in low-bulge systems. In our previous work, we demonstrated the power of mid-IR spectroscopy in detecting the low-power AGN in the Sd galaxy NGC 3621 (Satyapal et al. 2007; henceforth S07).  AGN show prominent high ionization fine structure line emission at mid-infrared wavelengths but starburst and normal galaxies are characterized by a lower ionization spectra characteristic of HII regions ionized by young stars (e.g. Genzel et al. 1998; Sturm et al. 2002; Satyapal et al. 2004).  In particular, the [NeV] 14 $\\mu$m  (ionization potential 97 eV) line is not generally produced in HII regions surrounding young stars, the dominant energy source in starburst galaxies, since even hot massive stars emit very few photons with energy sufficient for the production of Ne4+.  The detection of this line in any galaxy provides strong evidence for an AGN.  The goal of this paper is to answer the question: are AGN in late-type galaxies more common than previously thought?  If so, this would revise our understanding of the environments in which supermassive black holes (SMBHs) can form and grow and perhaps shed light on the nature of the M$_{\\rm BH}$-$\\sigma$ relation in low-bulge systems.  Toward this end, we conducted an archival investigation of 32 late-type galaxies observed by the {\\it Spitzer} high resolution spectrograph that show no definitive optical signatures of AGN to search for low-power and/or deeply embedded AGN.   This paper is structured as follows.  In Section 2, we summarize the properties of the {\\it Spitzer} archival sample presented in this paper.  In Section 3, we summarize the observational details and data analysis procedure, followed by a description of our results in Section 4.  In Section 5, we discuss the origin of the [NeV] emission , with each galaxy discussed individually, followed by a discussion of the implications of our discoveries in Section 6.  A summary of our major conclusions is given in Section 7. ",
        "conclusions": "We conducted a mid-infrared spectroscopic investigation of 32 late-type (Hubble type of Sbc or later) galaxies showing no definitive signatures of AGN in their optical spectra in order to search for low luminosity and/or embedded AGN.  The primary goal of our study was determine if AGN in low-bulge environments are more common than once thought.  Our high resolution {\\it Spitzer} spectroscopic observations reveal that the answer to this question is {\\it yes}.  Our main results are summarized below:  \\begin{enumerate}  \\item We detected the high ionization [NeV] 14.3$\\mu$m and/or 24.3$\\mu$m lines in 7 late-type galaxies, providing strong evidence for AGNs in these galaxies.  \\item We detected the high excitation [OIV] 25.9$\\mu$m and [NeIII] 15.5$\\mu$m lines in 5 out of the 7 of the galaxies with [NeV] emission.  Although these lines can be excited in star forming regions, our mapping observations (when available) suggest that the emission is centrally concentrated and likely to be dominated by the AGN.   \\item Taking into account the range of  sensitivities of our observations, our work suggests that the AGN detection rate based on mid-infrared diagnostics in late-type optically normal galaxies can be as much or more than $\\sim$ 30\\%.  This detection rate implies that the overall fraction of late-type galaxies hosting AGN is possibly more than 4 times larger than what optical spectroscopic observations alone suggest.  \\item Several of the galaxies with [NeV] detections have optical emission line ratios in the extreme ``starburst range'' indicating that there is absolutely no hint of an AGN based on their optical spectra. Three out of the 7 galaxies are classified based on their optical line ratios as ``transition objects'' but none show broad permitted optical lines.  \\item Amongst the 7 AGN candidates in our sample, 3 are Sbc, 3 are Sc, and 1 is of Hubble type Scd.  Since there are only 3 galaxies of Hubble type Sd, our limited sample size precludes us from making any definitive conclusions on the incidence of AGN in completely bulgeless galaxies.  The lowest central stellar velocity dispersion amongst the galaxies with published measurements is 40 km/s.  \\item We demonstrate using photoionization models with both an input AGN and an extreme EUV-bright starburst ionizing radiation field that the observed mid-infrared line ratios in our 7 AGN candidates cannot be replicated unless an AGN contribution is included. These models show that when the fraction of the total luminosity due to the AGN is low, the optical diagnostics are insensitive to the presence of the AGN.  In this regime of parameter space, the mid-infrared diagnostics offer a powerful tool in uncovering AGN missed by optical spectroscopy.  \\item Three of the galaxies, NGC 3556, NGC 3367, and NGC 4536, appear to be dominated by star-formation.  NGC 3556 could have as little as 10\\% of its luminosity coming from the AGN. NGC 4321, NGC 5055, NGC 3938, and NGC 4414 likely have a more dominant contribution to their luminosity from the AGN in addition to having some contribution to their emission line fluxes from shock-excited gas.  All of the galaxies  that H97 classifies as HII galaxies are well characterized by a low AGN contribution, while all the transition objects seem to require some shock component.  \\item The AGN bolometric luminosities inferred using our [NeV] line luminosities range from $\\sim$ 3$\\times$10$^{41}$ ergs s$^{-1}$ to $\\sim$ 2$\\times$10$^{43}$ ergs s$^{-1}$, with a median value of $\\sim$ 9$\\times$10$^{41}$ ergs s$^{-1}$. Assuming that the AGN is radiating at the Eddington limit, this corresponds a lower mass limit for the black hole that ranges from  $\\sim$ 3$\\times$10$^3$M$_{\\odot}$ to as high as $\\sim$ 1.5$\\times$10$^5$M$_{\\odot}$. These lower mass limits however do not put a strain on the well-known relationship between the black hole mass and the host galaxy\u0092s stellar velocity dispersion established in predominantly early-type galaxies.  \\end{enumerate}  The {\\it Spitzer} spectroscopic study presented here demonstrates that black holes do form and grow in low-bulge environments and that they are significantly more common than optical studies indicate.  In order to truly determine how common SBHs and AGN activity are in {\\it completely} bulgeless galaxies, a more extensive study with {\\it Spitzer} is crucial."
    },
    "0801/0801.0013_arXiv.txt": {
        "abstract": "The Vector Spectromagnetograph (VSM) instrument has been recording photospheric and chromospheric magnetograms daily since August 2003. Full-disk photospheric vector magnetograms are observed at least weekly and, since November 2006, area-scans of active regions daily. Quick-look vector magnetic images, plus X3D and FITS formated files, are now publicly available daily. In the near future, Milne-Eddington inversion parameter data will also be available and a typical observing day will include three full-disk photospheric vector magnetograms. Besides full-disk observations, the VSM is capable of high temporal cadence area-scans of both the photosphere and chromosphere. Carrington rotation and daily synoptic maps are also available from the photospheric magnetograms and coronal hole estimate images. ",
        "introduction": "The VSM currently operates at Kitt Peak, Arizona and is part of the Synoptic Optical Long-term Investigations of the Sun (SOLIS) project \\citep{Keller03}. The VSM provides a unique record of solar full-disk vector magnetograms along with high-sensitivity photospheric and chromospheric longitudinal magnetograms \\citep{Henney06}. The VSM began recording full-disk vector magnetograms in August 2003 at a temporary site in Tucson, Arizona. In April 2004, the VSM was relocated to Kitt Peak and resumed operations in May 2004. The VSM records the full Stokes profiles of the Fe I 630.15 and 630.25 nm lines with a current spatial and spectral sampling of 1.125 arcsec and 2.71 pm respectively. The 50-cm aperture VSM utilizes a Ritchey-Chr\\' etien optical design, where full-disk images are constructed from 2048 individual steps in declination of the projected solar image on the entrance slit (hereafter referred to as scan-lines). Observations with fewer scan-lines are referred to as area-scans and are typically recorded for active region areas. For example, 88 area-scans were recorded in the month of October 2003 during an exceptionally active flare period.  The VSM is available for user-requested programs and currently supports ongoing solar missions (e.g. MDI, STEREO, Hinode/SOT) and will support forthcoming missions (e.g. SDO/HMI). For example, VSM observations provide temporal and spatial context for the SOT observations. Though SOT/SP has exceptional spatial resolution, the VSM observes any solar region eight times faster than SOT/SP and ten times larger in one spatial direction. The VSM also provides full-disk coverage not available with SOT. In addition, the future SDO/HMI instrument will have high temporal cadence and spatial resolution, but will have very limited spectral resolution (with six filter images) that the VSM will complement with fully resolved Stokes profiles. \\begin{figure*}[!ht] \\plotone{henney_fig1_v1.ps} \\caption{An example VSM vector quick-look image of NOAA active region 10921, observed on November 3, 2006. The gray-scale illustrates the line-of-sight (LOS) magnetic flux and the arrows indicate the transverse field strength (the length scale is shown in the lower left-hand corner) and direction.} \\label{QL-AR} \\end{figure*}  \\subsection{VSM Upgrades} Starting in 2004, the VSM suffered from a slow degradation of the polarization modulator needed to measure full Stokes profiles. This complicated VSM vector calibrations and slowed efforts to release the vector magnetic data. In February of 2006, the vector modulator was replaced along with two modulators for the line-of-sight component of the magnetic field in the solar photosphere (630.2 nm) and the chromosphere (854.2 nm). During this maintenance period, a broken part was also replaced to reduce grating flexure. Though the new modulators are performing well, e.g. the strength of the 854.2 nm signal has significantly increased, the vector modulator has a polarization fringe pattern that is temperature dependent. The fringe pattern has recently been well parameterized to minimize the number of terms required for a good fit. The fringe fitting algorithm is expected to be incorporated in the SOLIS processing pipeline in early 2008. Since 2003, the VSM has incorporated two interim CMOS hybrid cameras made by Rockwell Scientific (now Teledyne Scientific \\& Imaging, LLC). During the first year of operation, several unexpected camera signals were revealed. Two of the pronounced artifacts, a variable dark level and signal cross-talk at the 0.5\\% level, required significant modification to the data reduction pipeline. The cameras were also found to have a 25\\% residual signal from previous frames. These interim cameras are scheduled to be replaced early in 2008 with cameras produced by Sarnoff that better match the spatial scale, sensitivity, low readout noise and desired frame rate for the VSM. ",
        "conclusions": ""
    },
    "0801/0801.2691_arXiv.txt": {
        "abstract": "{The observed clumpy structures in debris disks are commonly interpreted as particles trapped in mean-motion resonances with an unseen exo-planet. Populating the resonances requires a migrating process of either the particles (spiraling inward due to drag forces) or  the planet (moving outward). Because the drag time-scale  in resolved debris disks is generally long compared to the collisional time-scale, the planet migration scenario might be more likely, but this model has so far only been investigated for planets on circular orbits.} {We present a thorough study of the impact of a migrating planet on a planetesimal disk, by exploring a broad range of masses and eccentricities for the planet. We discuss the sensitivity of the structures generated in debris disks to the basic planet parameters.} {We perform many N-body numerical simulations, using the symplectic integrator SWIFT, taking into account the gravitational influence of the star and the planet on massless test particles. A constant migration rate is assumed for the planet.} {The effect of planetary migration on the trapping of particles in mean motion resonances is found to be very sensitive to the initial eccentricity of the planet and of the planetesimals. A planetary eccentricity as low as $0.05$ is enough to smear out all the resonant structures, except for the most massive planets. The planetesimals also initially have to be on orbits with a mean eccentricity of less than than $0.1$ in order to keep the resonant clumps visible.} {This numerical work extends previous analytical studies and provides a collection of disk images that may help in interpreting the observations of structures in debris disks. Overall, it shows that stringent conditions must be fulfilled to obtain observable resonant structures in debris disks. Theoretical models of the origin of planetary migration will therefore have to explain how planetary systems remain in a suitable configuration to reproduce the observed structures.} ",
        "introduction": "Since the first direct imaging of a debris disk around \\object{$\\beta$ Pictoris} by \\citet{1984Sci...226.1421S}, a dozen other optically thin dust disks have been spatially resolved around nearby main-sequence stars showing an infrared excess \\citep[][and references there in]{2007ApJ...661L..85K,2006ApJ...650..414S}.  The images often reveal asymmetric structures and clumps, interpreted as the signature of gravitational perturbations. A planet immersed in a debris disk usually produces structures such as a gap along  its orbit, by ejecting particles during close encounters, or density waves (e.g. a one-arm spiral), by modifying the precession rate of the dust particles \\citep{2005A&A...440..937W}. However, such structures cannot explain the observations of clumpy, non-axisymmetric disks \\citep{2004ASPC..321..305A,2007prpl.conf..573M}, and resonant mechanisms with unseen planets have been proposed to account for the observed asymmetries \\citep{2000ApJ...537L.147O,2002ApJ...578L.149Q,2003ApJ...588.1110K,2003ApJ...598.1321W}. A particle belongs to a mean motion resonance (MMR) when the particle to planet period ratio is a rational number, $m:n$ with $m$ and $n$ integers. An MMR is located at a semi-major axis $a$ given by $a/a_p=(m/n)^{2/3}$, where $a_p$ is the planet semi-major axis. In the Solar System,  for example, about $15\\%$ of the known Kuiper Belt objects, including Pluto, are trapped in the $3$:$2$ resonance with Neptune \\citep{2007prpl.conf..895C}. The interesting property of MMRs  for modeling asymmetric disks is that, as explained for example in \\citet{2000ssd..book.....M}, resonant objects are not uniformly distributed in azimuth around a star: rather they gather at specific longitudes relative to the perturbing planet and subsequently form clumps. This arises from properties specific to MMRs as a given particle trapped in a MMR with a planet undergoes conjunctions with the planet at specific locations along its orbit. The particles tend to gather around the most stable orbital configurations that ensure that the conjunctions occur at the maximum relative distance. The clumps, which are the result of the collective effects of resonant particles, generally corotate with the planet \\citep{2003ApJ...588.1110K}, while each of these resonant bodies has a different period from that of the planet (except for $1$:$1$ resonant planetesimals): hence the motion of these density waves differs from the orbital motion of the resonant particles. However, MMRs are very thin radial structures that usually trap a small number of particles in a given disk. Therefore, any structure due to MMRs has a high chance of being totally  hidden by the emission of the non-resonant particles, as illustrated in Fig. \\ref{withoutMigration}.  For clumps due to MMRs to be observed, the population of resonant particles must be significantly enhanced by an additional physical process. Two mechanisms can account for this: Poynting-Robertson (P-R) drag and planet migration. Dust particles that are too large to be ejected from the system by radiation pressure can spiral inward  into the star due to P-R drag and to some other minor forces like stellar wind drag \\citep[e.g.][]{2006A&A...455..987A}. In the course of its inward migration, a dust particle can be trapped into exterior MMRs with a planet, hence increasing the contrast of their asymmetric patterns \\citep{2003ApJ...588.1110K,2005ApJ...625..398D}. Planet migration, on the other hand, involves particles of all sizes, except those ejected by radiation pressure. Many particles can be trapped in MMRs by a planet migrating outward in the disk. Each non-resonant particle crossing an MMR has a chance trapped and subsequently migrating, following the resonance \\citep{2003ApJ...598.1321W}.  Several authors have studied either the effect of P-R drag, or of planet migration, on disk structures, using different methods (analytic, semi-analytic or numerical) and various planet parameters (mass and orbital eccentricity). A summary of the main previous studies is provided in Table \\ref{previousWorks}. The P-R drag scenario has been extensively  studied for a wide range of parameters, while  the migrating planet scenario has been investigated only for a planet on a circular orbit by  \\citet{2003ApJ...598.1321W}. It is thus important to better characterize the latter scenario in order to distinguish which of the two dominates the morphology of debris disks. Moreover, a number of studies \\citep{1996Icar..123..168L,2005A&A...433.1007W,2007A&A...462..199K} have shown that collisions may prevent MMRs from being populated by P-R drag since the collision timescale in massive debris disks might be much shorter than the P-R drag migration timescale.  Therefore,  we propose in this paper to extensively study the planet migration scenario, using numerical modeling, by generating a synthetic catalog similar to what has been done for the P-R drag scenario using analytical \\citep{2003ApJ...588.1110K} or numerical studies \\citep{2005ApJ...625..398D}. We extend the  pioneering work done by \\citet{2003ApJ...598.1321W} in studying the influence of the planet eccentricity on the visibility of the resonant patterns. In Section \\ref{lowEccOrb}, we discuss the case of a planet migrating on a circular or low-eccentricity orbit. In Section \\ref{highEccOrb}, we extend this study to planets on orbits with eccentricities up to $0.7$. This study is generalized to various migration rates and disk initial states in Section \\ref{Generalization}, and compared to previous studies in Section \\ref{Comparison}. The limitations of our approach are discussed in Section \\ref{Limitation}.  \\begin{table*} \\centering \\caption{\\label{previousWorks}Summary of recent papers on particles trapping in MMRs with a planet.} \\begin{tabularx}{\\textwidth}{c c c c c X} Authors & Method & \\multicolumn{2}{c}{Planet parameters} & Migration & Notes \\\\ && Mass ratio$^a$ & eccentricity \\\\ \\hline\\hline \\\\ \\citet{2003ApJ...598.1321W} & semi-analytic & $0.0003$ to $3$& $0$ & planet & Forced migration\\\\ \\citet{2006ApJ...639.1153W} & semi-analytic &&& none & This work extends the previous one to smaller particles sensitive to radiation pressure.\\\\ \\citet{2003ApJ...588.1110K} & analytic & $0.005$ to $15$ & $0$ to $0.6$ & particles & Only resonant particles trapped during migration due to P-R drag. \\\\ \\citet{2005ApJ...625..398D} & numerical & $0.01$ to $3$ & $0$ to $0.7$& particles & Particles migrate due to P-R drag and solar wind. \\\\ \\citet{2002AJ....124.2305M} & numerical &$0.05$&$0$& particles & Study of the Kuiper Belt. \\\\ This paper & numerical & $0.001$ to $3$ & $0$ to $0.7$ & planet & Only planetesimals disk, forced migration \\\\ \\hline \\end{tabularx} \\begin{list}{}{} \\item[$^a$]Planetary mass in Jovian mass divided by stellar mass in solar mass. \\end{list} \\end{table*}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{figures/withoutMigration.eps}} \\caption{\\label{withoutMigration} Example of a planetesimal disk without outward planet migration, nor inward P-R drag migration of the test particles, according to our numerical simulations as explained in Section \\ref{model}. The star and planet locations, projected onto the orbital plane of the planet, are represented by large red points, and the planet orbit by a thin green line. The initial planetesimal disk consists of $50\\,000$ planetesimals distributed between 40 and 75 AU, with the surface density distribution proportional to $r^{-1}$. Although some planetesimals are trapped in MMRs with the planet, they are not sufficiently numerous to generate spatial structures (besides the $1$:$1$ MMR). \\thanks{See the electronic edition of the Journal for a color version of this figure.}} \\end{figure} ",
        "conclusions": "We have studied the problem of the presence of observable structures in planetesimal disks due to mean motion resonance with an unseen planet migrating outward in the disk. Using numerical simulations, we have explored a large range of parameters for the planet (mass and orbital eccentricity) and the disk (initial distribution of planetesimal eccentricities). In the case of a planet on a circular orbit migrating inside a dynamically cold disk, our results are in agreement with previous analytical studies.  In the cases not already addressed, namely planets on eccentric orbits or dynamically warm disks, we have found that the observability of resonant structures demands very specific orbital configurations. The clumps produced by MMRs with a planet on a circular orbit are smoothed in the case of a planet on an even moderately eccentric orbit. An eccentricity as low as $0.05$ is enough to smooth all the resonant structures, except for the most massive planets. These results indicate that although trapping planetesimals in MMRs is an efficient mechanism to generate clumpy disks, stringent conditions must be fulfilled for this scenario to occur. Theoretical modeling of the origin of the planetary migration therefore will have to explain how planetary systems can remain under these conditions. Moreover, we only consider a planet migrating at a constant rate. A more realistic model with a variable, stochastic migration rate can reduce the population of resonances and thus their observability. A better model of planet migration  thus should be developed in future studies.  \\begin{table*} \\centering \\caption{\\label{normal} Summary of results for all simulations done in the present study with an initially unexcited disk (initial planetesimal eccentricities are equal to zero) and a standard migration rate of $0.5$ AU Myr$^{-1}$. For each simulation, we list the migration rate, the planet mass and eccentricity, the fraction of surviving planetesimals at the end of the simulation ($40$ Myr) and the resulting disk shape, following the convention of Fig. \\ref{resume}.} \\begin{tabular}{ccccc||ccccc} Mig. rate$^a$ & Mass$^b$ & Ecc. & Surv. planetesimals$^c$ & Disk shape$^d$& Mig. rate$^a$ & Mass$^b$ & Ecc. & Surv. planetesimals$^c$ & Disk shape$^d$\\\\ \\hline 0.5 & 0.0035 &0.0&$100\\%$&None&0.5 & 0.05 &0.0&$100\\%$&I\\\\ & &0.01&$100\\%$&None& & &0.01&$100\\%$&I\\\\ & &0.05&$100\\%$&None& & &0.05&$100\\%$&II\\\\ & &0.1&$100\\%$&None&  & &0.1&$100\\%$&None\\\\ & &0.2&$100\\%$&None& & &0.2&$50\\%$&None\\\\ & &0.3&$100\\%$&None& & &0.3&$10\\%$&None\\\\ & &0.4&$35\\%$&III& & &0.4&$5\\%$&None\\\\ & &0.5&$10\\%$&III& & &0.5&$10\\%$&None\\\\ & &0.6&$5\\%$&III& & &0.6&$5\\%$&None\\\\ & &0.7&$5\\%$&III& & &0.7&$5\\%$&None\\\\ 0.5 & 0.33 &0.0&$75\\%$&I&0.5 & 1 &0.0&$70\\%$&I\\\\ & &0.01&$70\\%$&I& & &0.01&$70\\%$&I\\\\ & &0.05&$85\\%$&II& & &0.05&$55\\%$&II\\\\ & &0.1&$75\\%$&None& & &0.1&$25\\%$&II\\\\ & &0.2&$25\\%$&None& & &0.2&$10\\%$&None\\\\ & &0.3&$10\\%$&None& & &0.3&$10\\%$&None\\\\ & &0.4&$10\\%$&None& & &0.4&$10\\%$&None\\\\ & &0.5&$10\\%$&None& & &0.5&$5\\%$&None\\\\ & &0.6&$5\\%$&None& & &0.6&$5\\%$&None\\\\ & &0.7&$5\\%$&None& & &0.7&$0\\%$&None\\\\ 0.5 & 2 &0.0&$65\\%$&I&0.5 & 3 &0.0&$55\\%$&I\\\\ & &0.01&$60\\%$&I& & &0.01&$50\\%$&I\\\\ & &0.05&$30\\%$&I& & &0.05&$20\\%$&I\\\\ & &0.1&$15\\%$&II& & &0.1&$10\\%$&I\\\\ & &0.2&$5\\%$&None& & &0.2&$0\\%$&None\\\\ & &0.3&$5\\%$&None& & &0.3&$0\\%$&None\\\\ & &0.4&$5\\%$&None& & &0.4&$0\\%$&None\\\\ & &0.5&$5\\%$&None& & &0.5&$0\\%$&None\\\\ & &0.6&$0\\%$&None& & &0.6&$0\\%$&None\\\\ & &0.7&$0\\%$&None& & &0.7&$0\\%$&None\\\\ \\hline \\multicolumn{10}{l}{$^a$In AU Myr$^{-1}$.}\\\\ \\multicolumn{10}{l}{$^b$In Jovian mass.}\\\\ \\multicolumn{10}{l}{$^c$Fraction of surviving planetesimals at the end of the simulation, i.e. $40$ Myr.}\\\\ \\multicolumn{10}{l}{$^d$As in Fig. \\ref{resume}.}\\\\ \\end{tabular} \\end{table*}  \\begin{table*} \\centering \\caption{\\label{fastslow} Same as Table \\ref{normal}, but for different migration rates.} \\begin{tabular}{ccccc||ccccc} Mig. rate$^a$ & Mass$^b$ & Ecc. & Surv. planetesimals$^c$ & Disk shape$^d$& Mig. rate$^a$ & Mass$^b$ & Ecc. & Surv. planetesimals$^c$ & Disk shape$^d$\\\\ \\hline 5 & 0.035 &0.0&$100\\%$&None&5 & 0.05 &0.0&$100\\%$&I\\\\ & &0.01&$100\\%$&None& & &0.01&$100\\%$&II\\\\ & &0.05&$100\\%$&None& & &0.05&$100\\%$&II\\\\ & &0.1&$100\\%$&None& & &0.1&$100\\%$&None\\\\ & &0.2&$100\\%$&None& & &0.2&$100\\%$&None\\\\ 5 & 0.33 &0.0&$100\\%$&I&5 & 1 &0.0&$100\\%$&I\\\\ & &0.01&$100\\%$&I& & &0.01&$100\\%$&I\\\\ & &0.05&$100\\%$&II& & &0.05&$100\\%$&II\\\\ & &0.1&$100\\%$&II& & &0.1&$95\\%$&II\\\\ & &0.2&$95\\%$&None& & &0.2&$75\\%$&None\\\\ 5 & 2 &0.0&$85\\%$&I&5 & 3 &0.0&$75\\%$&I\\\\ & &0.01&$85\\%$&I& & &0.01&$75\\%$&I\\\\ & &0.05&$80\\%$&I& & &0.05&$60\\%$&I\\\\ & &0.1&$65\\%$&II& & &0.1&$40\\%$&I\\\\ & &0.2&$50\\%$&None& & &0.2&$25\\%$&None\\\\ 0.05 & 0.035 &0.0&$100\\%$&II&0.05 & 0.05 &0.0&$90\\%$&I\\\\ & &0.01&$100\\%$&II& & &0.01&$85\\%$&I\\\\ & &0.05&$100\\%$&None& & &0.05&$85\\%$&II\\\\ & &0.1&$100\\%$&None& & &0.1&$75\\%$&II\\\\ & &0.2&$80\\%$&None& & &0.2&$10\\%$&None\\\\ 0.05 & 0.33 &0.0&$75\\%$&I&0.05 & 1 &0.0&$70\\%$&I\\\\ & &0.01&$75\\%$&I& & &0.01&$70\\%$&I\\\\ & &0.05&$65\\%$&II& & &0.05&$40\\%$&II\\\\ & &0.1&$40\\%$&II& & &0.1&$20\\%$&II\\\\ & &0.2&$5\\%$&None& & &0.2&$5\\%$&None\\\\ 0.05 & 2 &0.0&$60\\%$&I&0.05 & 3 &0.0&$55\\%$&I\\\\ & &0.01&$60\\%$&I& & &0.01&$55\\%$&I\\\\ & &0.05&$25\\%$&I& & &0.05&$15\\%$&I\\\\ & &0.1&$15\\%$&II& & &0.1&$10\\%$&None\\\\ & &0.2&$0\\%$&None& & &0.2&$0\\%$&None\\\\ \\hline \\multicolumn{10}{l}{$^a$In AU Myr$^{-1}$.}\\\\ \\multicolumn{10}{l}{$^b$In Jovian mass.}\\\\ \\multicolumn{10}{l}{$^c$Fraction of surviving planetesimals at the end of the simulation, i.e. $4$ Myr for $5$ AU Myr$^{-1}$ migration rate and $200$ Myr for $0.05$ AU Myr$^{-1}$.}\\\\ \\multicolumn{10}{l}{$^d$As in Fig. \\ref{resume}.}\\\\ \\end{tabular} \\end{table*}   \\begin{table*} \\centering \\caption{\\label{complete_hot} Same as Table \\ref{normal} but for simulations with initially excited disks. The maximum initial eccentricity of the planetesimals is mentioned for all simulations. The migration rate is $0.5$ AU Myr$^{-1}$ for all simulations.} \\begin{tabular}{ccccc||ccccc} Max. ecc.$^a$ & Mass$^b$ & Ecc. & Surv. planetesimals$^c$ & Disk shape$^d$ &Max. ecc.$^a$ & Mass$^b$ & Ecc. & Surv. planetesimals$^c$ & Disk shape$^d$ \\\\ \\hline 0.1 & 0.035 &0.0&$100\\%$&None&0.1 & 0.05 &0.0&$100\\%$&II\\\\ & &0.01&$100\\%$&None& & &0.01&$100\\%$&None\\\\ & &0.05&$100\\%$&None& & &0.05&$100\\%$&None\\\\ & &0.1&$100\\%$&None& & &0.1&$100\\%$&None\\\\ & &0.2&$100\\%$&None& & &0.2&$55\\%$&None\\\\ 0.1 & 0.33 &0.0&$70\\%$&II&0.1 & 1 &0.0&$60\\%$&I\\\\ & &0.01&$70\\%$&II& & &0.01&$55\\%$&II\\\\ & &0.05&$85\\%$&II& & &0.05&$50\\%$&II\\\\ & &0.1&$70\\%$&None& & &0.1&$25\\%$&II\\\\ & &0.2&$20\\%$&None& & &0.2&$5\\%$&None\\\\ 0.1 & 2 &0.0&$50\\%$&I&0.1 & 3 &0.0&$40\\%$&I\\\\ & &0.01&$45\\%$&I& & &0.01&$40\\%$&I\\\\ & &0.05&$25\\%$&I& & &0.05&$15\\%$&I\\\\ & &0.1&$15\\%$&II& & &0.1&$10\\%$&I\\\\ & &0.2&$5\\%$&None& & &0.2&$0\\%$&None\\\\ 0.2 & 2 &0.0&$35\\%$&I&0.2 & 3 &0.0&$30\\%$&I\\\\ & &0.01&$35\\%$&I& & &0.01&$25\\%$&I\\\\ & &0.05&$20\\%$&I& & &0.05&$10\\%$&I\\\\ & &0.1&$10\\%$&I& & &0.1&$5\\%$&I\\\\ & &0.2&$5\\%$&None& & &0.2&$0\\%$&None\\\\ \\hline \\multicolumn{10}{l}{$^a$For planetesimals.}\\\\ \\multicolumn{10}{l}{$^b$In Jovian mass.}\\\\ \\multicolumn{10}{l}{$^c$Fraction of surviving planetesimals at the end of the simulation, i.e. $40$ Myrs.}\\\\ \\multicolumn{10}{l}{$^d$As in Fig. \\ref{resume}.}\\\\ \\end{tabular} \\end{table*}"
    },
    "0801/0801.2372_arXiv.txt": {
        "abstract": "Two first order strongly hyperbolic formulations of scalar-tensor theories of gravity allowing nonminimal couplings (Jordan frame) are presented along the lines of the 3+1 decomposition of spacetime. One is based on the Bona-Mass\\'o formulation, while the other one employs a conformal decomposition similar to that of Baumgarte-Shapiro-Shibata-Nakamura. A modified Bona-Mass\\'o slicing condition adapted to the scalar-tensor theory is proposed for the analysis. This study confirms that the scalar-tensor theory has a well posed Cauchy problem even when formulated in the Jordan frame. ",
        "introduction": "\\label{sec:introduction}  Scalar-tensor theories of gravity (STT) are alternative theories of gravitation where a scalar field is coupled nonminimally to the curvature associated with the {\\em physical metric}\\/ (this is the so-called {\\em Jordan frame}\\/ representation).  The term ``physical metric'' refers to a situation where test particles follow the geodesics of that metric.  The variation of the action of the STT with respect to the physical metric gives rise to field equations which contain an effective energy-momentum tensor (EMT) involving second order derivatives in time and space of the scalar field.  Such EMT has the property that ``ordinary matter'', {\\em i.e.} matter associated with fields other that the scalar field, obeys the (weak) equivalence principle which mathematically translates into a conserved EMT for ordinary matter alone.  Since {\\em a priori}\\/ it was not clear how such second order derivatives could be eliminated (in terms of lower order derivatives), or managed so as to obtain a quasilinear system of hyperbolic equations for which the Cauchy problem was well-posed (in the Hadamard sense), many people decided to abandon this approach in favor of the so-called {\\em Einstein frame}\\/ representation where the nonminimal coupling (NMC) is absorbed into the curvature by means of a conformal transformation of the metric.  The new conformal metric is unphysical in the sense that (non-null) test particles do not follow the geodesics of that metric.  However, the mathematical advantage is that the field equations for the nonphysical metric resemble the standard Einstein field equations with an unphysical effective EMT which involves at most first order derivatives of a suitable transformed scalar field (this EMT is unphysical because the ``ordinary matter'' part is not separately conserved). In the Einstein frame one can show that by using standard gauges ({\\em e.g.} harmonic gauge) the field equations acquire the form required in the application of theorems that establish the well-posedness of the Cauchy problem.  In view of the apparent mathematical advantages and disadvantages of the Jordan and Einstein frames, several widespread misconceptions became common in the literature. One of these concerned the statement that the Cauchy problem is only well-posed in the Einstein frame~\\cite{Faraoni04}. In Ref.~\\cite{Salgado06}, however, one of us showed that this is not the case by following two different approaches.  One was in the spirit of a second order analysis consisting on reducing the set of field equations, both for the metric components and the scalar field, to a manifestly quasilinear diagonal second order hyperbolic form (the ``reduced'' field equations). This was achieved by manipulating the field equations in a way that allowed one to express the d'Alambertian on the scalar field in terms of at most first order derivatives, and also by implementing a modified harmonic gauge which was adapted to the STT. Such a modified gauge allowed one to eliminate the remaining second order derivatives of the scalar field which had previously prevented the applicability of Leray's theorem (see {\\em e.g.} Ref.~\\cite{Wald84} for the theorem).  The second approach followed in \\cite{Salgado06}, and which was directly related with the initial value problem, consisted in recasting the field equations of the STT in a 3+1 or Arnowitt-Deser-Misner (ADM)~\\cite{Arnowitt62} form written {\\em \\`a la} York~\\cite{York79} (hereafter referred to as the ADMY equations). To do so it was necessary to define suitable variables, followed by a manipulation of the equations in order to obtain well-defined constraints (independent of the choice of the lapse and shift), as well as first order in time evolution equations for the extrinsic curvature and the scalar-field variables.  In Ref.~\\cite{Salgado06} the main goal was to focus in the 3+1 approach rather than the second order one because of the subsequent numerical applications we had in mind. In fact, almost all the modern codes used in numerical relativity are based on first order in time formulations, and therefore we wanted to adapt them for the analysis of phenomena in STT.  Now, when taking the limit of pure general relativity (GR) ({\\em i.e.} no no-minimal coupling), the 3+1 equations presented in~\\cite{Salgado06} reduce to the standard ADMY equations with a minimally coupled scalar field plus ordinary matter sources. It is well-known that the ADMY equations are not strongly hyperbolic (see~\\cite{Alcubierre08a} for a detailed discussion), and therefore the corresponding equations for STT described in~\\cite{Salgado06} were expected not to be strongly hyperbolic either. As already mentioned in~\\cite{Salgado06}, such equations were only the first step towards a first order strongly hyperbolic system for which the well-posedness of the Cauchy problem could be established.  The aim of this paper is then to fill that gap, and to obtain a first order strongly hyperbolic system of partial differential equations based on the 3+1 equations of~\\cite{Salgado06}. Actually, we will show here two such systems: one that is related to the Bona-Mass\\'o-Seidel-Stela (BMSS) formulation~\\cite{Bona94b} (which in turn is based on the Bona-Mass\\'o (BM) formulation~\\cite{Bona92}), and a second one based on the Baumgarte-Shapiro-Shibata-Nakamura (BSSN) conformal decomposition~\\cite{Baumgarte:1998te,Shibata95}.  As was done in~\\cite{Salgado06}, a modified BM slicing condition will be used for the STT, while the shift vector will be taken as an {\\em a priori}\\/ given function of the coordinates.  A last comment about the Jordan vs. Einstein frame is in order. In~\\cite{Salgado06} it was a matter of principle to show that the Cauchy problem could be well formulated in the Jordan frame. It then became clear that there was no fundamental reason to continue using the Einstein frame. After all, the Jordan frame is the one associated with the (physical) quantities that are to be confronted with the observations. Moreover, the constraint and evolution equations are not more involved than those of usual GR, so from a numerical point of view the use of the Jordan frame does not add much complexity to the analysis. On the other hand, the direct use of the Jordan frame allows a better physical interpretation of the results and avoids the potential problems that might arise in the Einstein frame in cases when the conformal transformations back to the physical metric are not well-defined.  Finally, it has been argued that some of the energy densities associated with the Jordan frame effective EMT (JFEMT) are not positive definite and may therefore be unphysical, and also the corresponding ADM mass could then become negative. While it is certainly true that the JFEMT does not satisfy the energy conditions in general, it is very likely that any physical configuration consistent with observations will not carry (total) negative energy.  This is because the deviations of any alternative theory from GR are presumably very small in order to reproduce many of the current observations (cf. Ref.~\\cite{Will93}). Actually, in any viable STT cosmology it turns out that the total energy density of the Universe is always positive, even when taking into account the negative contributions due to the NMC~\\cite{Salgado96,Salgado97a,Salgado97b,Quevedo99}. Again, this is because the scalar-field contributions (positive or negative) to the total energy density should be in agreement with several observations related to the past and present history of the Universe. For instance, the expansion rate of the Universe modified by the NMC contributions has to be consistent with the one required to produce the correct abundance of primordial nucleosynthesis~\\cite{Salgado96}. In addition, the density perturbations in STT should also match the Cosmic Microwave Background data. In other words, a STT cosmology which produces a total effective negative energy density of the Universe at any given epoch will surely not be consistent with observations. By the same token, in a consistent STT cosmology the negative contributions (if any) are to be naturally suppressed by the positive ones, providing a consistent Universe.  In the case of astrophysical applications ({\\em e.g.} neutron star models) within a class of STT with a positive definite NMC function (which implies a positive definite effective gravitational ``constant''), the ADM mass turns to be always positive despite the negative contributions to the energy density due to the NMC~\\cite{Salgado98}. Moreover, the value of the NMC cannot be very high as otherwise the corresponding STT would put in jeopardy the agreement with the binary pulsar observations~\\cite{Damour96}.  Of course, the nice thing about matter fields satisfying the energy conditions is the applicability of several theorems ({\\em e.g.} positive mass theorems and singularity theorems; see~\\cite{Wald84} and references therein). However, the fact that the JFEMT does not satisfy the energy conditions in general simply indicates that the theorems cannot say anything about the positivity of mass or the formation of singularities in this case. This is not a problem of physical but of mathematical character. But again, the nonpositive definite terms arising from the NMC will be surely bounded if the STT in hand is to be consistent with the current observations and experiments, and therefore the nice features of an EMT respecting the energy conditions will also very likely appear in any observationally consistent STT. Of course, all these arguments are only sustained by numerical experiments when constructing viable phenomenological models (in cosmology and compact objects), and therefore do not constitute a theorem. It would then be quite interesting to explore the possibility of proving positive energy and singularity theorems in the Jordan frame if one can bound the non-positive definite contribution associated with the JFEMT.  This paper is organized as follows. In Sec.~\\ref{sec:ADM} we introduce the STT and the 3+1 equations described in~\\cite{Salgado06}. The notation of several variables will be slightly modified relative to Ref.~\\cite{Salgado06} in order to match with the one usually employed in numerical relativity.  In Sec.~\\ref{sec:hyper} a fully first order strongly hyperbolic system of the field equations of STT is obtained along the lines of the BMSS and BSSN formulations of GR. We conclude with a discussion of future numerical applications of the hyperbolic systems presented here. Finally, in Appendix A we include the complete set of equations for both formulations, and in Appendix B we introduce a simple example to complement the ideas of Sec.~\\ref{sec:hyper}. ",
        "conclusions": "\\label{sec:discussion}  In this paper we have constructed two novel first order strongly hyperbolic formulations of the STT in the Jordan frame along the lines of the BM and BSSN approaches. Such constructions show that both formulations have a well-posed Cauchy problem.  This analysis fills the gap of a previous study on the Cauchy problem of STT~\\cite{Salgado06} and confirms that the Jordan frame is mathematically adequate for treating the initial value problem.  One of the most interesting features of the formulations presented here is that a modified Bona-Mass\\'o slicing condition is required for the two new systems to be strongly hyperbolic while allowing several slicings ($f_{BM}>0$) which are natural generalizations of the slicings used in pure GR. In particular $\\Theta=1,f_{BM}$ in Eq.~(\\ref{STTBMlapse}) are two simple choices that lead to well behaved eigenfields.  In the absence of a scalar field, the equations of Sec.IIIA,B reduce (for $\\varsigma=-1$, $\\zeta=4$ and $\\xi=2$) to the known BM and BSSN formulations of GR (when the NMC function is trivial, {\\em i.e.}, $F(\\phi)\\equiv 1$, the ``gravitational'' and the scalar-field sectors decouple completely up to principal part).  What remains to be investigated is the usefulness and robustness of these formulations in actual numerical experiments, as well as the inclusion of a ``live shift''.  Actually, we plan to analyze the dynamical transition to the phenomenon of spontaneous scalarization in boson stars arising in STT \\cite{Whinnett00} and the subsequent gravitational collapse to a black hole with gravitational wave emission of scalar type, using one or both of the hyperbolic formulations presented here. Both phenomena (spontaneous scalarization and scalar gravitational waves) are even present in spherical symmetry due to the NMC, therefore by assuming such a symmetry one can simplify the equations without eliminating the interesting physical features. \\bigskip  An important consequence of the analysis presented here is that a slightly more general STT which includes a function $\\omega(\\phi)$ in the kinetic term of the scalar-field sector of the action~(\\ref{jordan}) ({\\em i.e.} the kinetic term has the form $\\omega(\\phi)(\\nabla \\phi)^2/2$) posses a well-posed Cauchy problem as well (except for some choices of $\\omega(\\phi)$; see Ref.~\\cite{Faraoni07}).  The fact is that all the terms with $\\omega(\\phi)$ do not contribute to the principal part of the equations associated with the {\\em metric}\\/ sector ({\\em i.e.} the equivalent of Eq.~(\\ref{Einst})), while the scalar-field sector (the equivalent of Eq.~(\\ref{KG})) preserves the quasilinear diagonal hyperbolic form (see Ref.~\\cite{Faraoni07} for the detailed equations). Thus, up to principal part such STT are identical to the ones analyzed here. The relevance of this generalization is that such STT can be mapped to the so-called {\\em modified theories of gravity}\\/ which are given by a Lagrangian density $f(R)$ ($R$ being the Ricci scalar)~\\cite{Faraoni07}. Some specific choices of $f(R)$ lead to gravity theories which have been recently analyzed in several contexts. Notably, in the cosmological setting such theories have been proposed as an alternative to dark energy, since they can produce an accelerating expansion of the Universe without any exotic form of matter~\\cite{Carroll04,Capozziello03}. However, it must be emphasized that some of these theories might violate the solar system tests~\\cite{Erickcek06,Chiba07}, and some modifications are required to circumvent such drawbacks (see Ref.~\\cite{Nojiri07} for a review).  The point we want to underline here is that, from the mathematical point of view, the viability of such theories relies heavily on the well-posedness of the Cauchy problem."
    },
    "0801/0801.2975_arXiv.txt": {
        "abstract": "{}{ Several photometric redshift (photo-$z$) codes are discussed in the literature and some are publicly available to be used by the community. We analyse the relative performance of different codes in blind applications to ground-based data. In particular, we study how the choice of the code-template combination, the depth of the data, and the filter set influences the photo-$z$ accuracy.}{We performed a blind test of different photo-$z$ codes on imaging datasets with different depths and filter coverages and compared the results to large spectroscopic catalogues. { We analysed the photo-$z$ error behaviour to select cleaner subsamples with more secure photo-$z$ estimates.} We consider \\emph{Hyperz}, \\emph{BPZ}, and the code used in the CADIS, COMBO-17, and HIROCS surveys.}  {{ The photo-$z$ error estimates of the three codes do not correlate tightly with the accuracy of the photo-$z$'s. While very large errors sometimes indicate a true catastrophic photo-$z$ failure, smaller errors are usually not meaningful.} For any given dataset, we find significant differences in redshift accuracy and outlier rates between the different codes { when compared to spectroscopic redshifts}.  However, different codes excel in different regimes. { % The agreement between different sets of photo-$z$'s is better for the subsample with secure spectroscopic redshifts than for the whole catalogue.  Outlier rates in the latter are typically larger by at least a factor of two.}} {{ Running today's photo-$z$ codes on well-calibrated ground-based data can lead to reasonable accuracy. The actual performance on a given dataset is largely dependent on the template choice and on realistic instrumental response curves. The photo-$z$ error estimation of today's codes from the probability density function is not reliable, and reported errors do not correlate tightly with accuracy. It would be desirable to improve this aspect for future applications so as to get a better handle on rejecting objects with grossly inaccurate photo-$z$'s. The secure spectroscopic subsamples commonly used for assessments of photo-$z$ accuracy may be biased toward objects for which the photo-$z$'s are easier to estimate than for a complete flux-limited sample, resulting in very optimistic estimates.}} ",
        "introduction": "\\label{sec:introduction} Photometric redshifts (hereafter, photo-$z$) have become a standard tool for the observing astronomer in the last years { \\citep{1986ApJ...303..154L, 1995AJ....110.2655C, 1999ASPC..191....3K, 1999A&A...343..399W, 2000ApJ...536..571B, 2000A&A...363..476B, 2001AJ....122.1151R, 2001AJ....122.1163B, 2001A&A...365..660W, 2003AJ....125..580C, 2003MNRAS.339.1195F, 2004PASP..116..345C, 2004MNRAS.353..654B, 2006A&A...457..841I, 2006MNRAS.372..565F}}.  Not only are large multi-colour imaging surveys planned and executed with the goal of estimating the redshift of as many galaxies as possible from their broad-band photometry, but also many smaller projects benefit from this technique by providing redshifts that are much cheaper, in terms of telescope time, than spectroscopic ones and may go deeper.  { Users of photo-$z$'s are often concerned with three main performance issues, which are the mean redshift error, the rate of catastrophic failures, and the validity of the probability density function (PDF) in a frequentist interpretation. The PDF may be correct in a Bayesian interpretation when including systematic uncertainties in the model fitting and correctly express a degree of uncertainty.  However, given the non-statistical nature of systematic uncertainties a frequentist PDF that correctly describes the redshift distribution in the real experiment is necessarily different, unless such systematics can be excluded.  The process depends on three ingredients: model, classifier, and data. A basic issue at the heart of problems with the PDF are the match between data and model, since best-fitting parameters and confidence intervals in $\\chi^2$-fitting are only reliable when the model is appropriate.  The importance in choosing the type of data is the need to break degeneracies between ambiguous model interpretations. Finally, the classifiers are expected to produce similar results, while they could produce them at dramatically different speed. Artificial Neural Nets, hereafter ANNs, are especially fast once training has been accomplished \\citep{2003MNRAS.339.1195F}.  There are many cases in the literature where the precision of photo-$z$'s has been improved after recalibrating the match between data and model \\citep[see e.g. ][]{2003AJ....125..580C,2004ApJS..150....1B,2006A&A...457..841I,2006AJ....132..926C,2007ApJS..172...99C}, although this process requires a large, representative training set of spectroscopic redshifts from the pool of data that is to be photo-$z$'ed.  If ANNs are trained with sufficiently large training samples they can achieve the highest accuracies within the training range as a mismatch between data and model is ruled out from the start.  The literature reports several different photo-$z$ estimators in use across the community, some of which use different template models and some of which allow implementation of user-defined template sets. Assuming a modular problem, where model (templates), classifier, and data can be interchanged, it is interesting to test how comparable the results of different combinations are. In this spirit, we have started the work presented in this paper, where we analyse photo-$z$ performance from real ground-based survey data, in dependence of magnitude, depth of data, filter coverage, redshift region, and choice of photo-$z$ code. We concentrate on the blind performance of photo-$z$'s which is the most important benchmark for any study that cannot rely on recalibration, e.g. in the absence of spectroscopic redshifts. We choose to focus on ground-based datasets because a lot of codes were tested on the Hubble-Deep-Field for which results can already be found in the literature \\citep[see e.g. ][]{1998AJ....115.1418H, 2000A&A...363..476B, 2000ApJ...536..571B}.  Meanwhile, a much larger initiative has formed to investigate all (even subtle) differences in workings and outcomes among codes and models. This initiative called PHAT\\footnote{\\url{http://www.strw.leidenuniv.nl/~hendrik/PHAT}} (PHoto-$z$ Accuracy Testing) engages a world-wide community of photo-$z$ developers and users and will hopefully develop our understanding of photo-$z$'s to a reliably predictive level.}  The paper is organised as follows. In Sect.~\\ref{sec:data} the imaging and spectroscopic datasets are presented. The photo-$z$ codes used for this study are described in Sect.~\\ref{sec:codes}. Sect.~\\ref{sec:strategy} presents our approach for describing photo-$z$ accuracy. The results are presented and discussed in Sect.~\\ref{sec:results}. { The different photo-$z$ estimates are compared to each other in Sect.~\\ref{sec:phz_vs_phz}.} A final summary and general conclusions are given in Sect.~\\ref{sec:conclusions}.  Throughout this paper we use Vega magnitudes if not otherwise mentioned. ",
        "conclusions": "\\label{sec:conclusions} We have shown that photo-$z$'s estimated with today's tools can produce a reasonable accuracy. The performance of a particular photo-$z$ code, however, cannot easily be characterised by a mere two numbers such as scatter and global outlier rate. The benchmarks are rather sensitive functions of filter set, depth, redshift range and code settings.  Moreover, there is at least a factor of two possible difference in performance between different codes which is again not stable for all setups but can vary considerably from one setup to another. There are, for example, redshift ranges where one code clearly beats another one in terms of accuracy only to loose at other redshifts. We give estimates of the performance for a number of codes in some practically relevant cases.  The estimation of photo-$z$'s from different ground-based datasets is not straightforward and results should not be expected to be identical to simulated photo-$z$ estimates. Rather, photo-$z$ simulations often seem to circumvent critical steps in ground-based photo-$z$ estimation. { Most importantly}, the match between observed colours and some template sets commonly used may be suboptimal.  In the preceding sections we have identified several aspects which are relevant to future optimisations of photo-$z$ codes. The photo-$z$ error estimation is one of the most unsatisfying aspects to date with error values often only very weakly correlated with real uncertainties. This is likely due to the insufficient inclusion of systematics since very low S/N objects, for which the errors should be dominated by photon shot-noise, show a tighter correlation. Chip-to-chip sensitivity variations, especially in the UV, could either be taken into account more accurately within the photo-$z$ codes or could be tackled by improved instrument design, survey strategy, and data reduction. The optimisation of template sets can be expected to be successfully done with ever larger spectroscopic catalogues becoming available.  { In general, biases can be removed by a recalibration which requires an extensive spectroscopic training set. Another proven successful route to better photo-$z$'s is improving the spectral resolution of the data, instead of their depth, as demonstrated by the COMBO-17 survey. This approach is also taken by the new ALHAMBRA survey \\citep[][ Ben\\'itez et al., 2007, A\\&A, submitted]{2005astro.ph..4545M} and COSMOS-21.  A general problem for all studies comparing photo-$z$'s to spectroscopic redshifts is our finding that secure spectroscopic samples can be biased. While surveys like VVDS are $>90\\%$ complete in obtaining spectra for galaxy samples the redshifts that are claimed to be $>90\\%$ secure only amount to $\\sim50\\%$. This subsample obviously consists of galaxies for which the photo-$z$ estimation works better than for the whole sample. In the future, it is desirable to put effort into spectroscopic surveys with secure redshift measurements for virtually every galaxy down to the same flux limit that is used for the analysis of photo-$z$ samples.  Several questions that are raised in this work will be tackled by the PHAT initiative mentioned above. PHAT aims to understand the issues presented here in a systematical and quantitative way in order to give guidance for better photo-$z$'s in the future.}"
    },
    "0801/0801.1147_arXiv.txt": {
        "abstract": "A two-dimensional electrodynamic model is used to study particle acceleration and non-thermal emission mechanisms in the pulsar magnetospheres. We solve distribution of the accelerating electric field with  the emission process and the pair-creation process in meridional plane, which includes the rotational axis and the magnetic axis. By solving the evolutions of the Lorentz factor,  and of the pitch angle, we calculate  spectrum in optical through $\\gamma$-ray bands with the curvature radiation, synchrotron radiation,   and inverse~-Compton process not only for  outgoing particles, but also for ingoing particles, which were ignored in  previous studies.  We apply the theory to the Vela pulsar. We find that the curvature radiation from the outgoing particles is the major emission process above 10~MeV bands. In soft $\\gamma$-ray to hard X-ray  bands, the synchrotron radiation from the ingoing primary particles in the gap dominates in the spectrum. Below hard X-ray bands, the synchrotron emissions from both outgoing and ingoing particles contribute to the calculated spectrum. The calculated spectrum is consistent with the observed phase-averaged spectrum of the Vela pulsar.  Taking into account the predicted dependency of the emission process and the emitting particles on the energy bands, we  compute the  expected pulse profile in X-ray and $\\gamma$-ray bands with a three-dimensional geometrical model. We show that  the observed five-peak pulse profile in the X-ray bands of  the Vela pulsar is reproduced by the inward and  outward emissions,  and the observed double-peak pulse profile in $\\gamma$-ray bands is explained by the outward emissions.  We also apply the theory to PSR B1706-44 and PSR B1951+32, for which X-ray emission properties have not been constrained observationally very well, to  predict the spectral features with  the present outer gap model. ",
        "introduction": "The Compton Gamma-Ray Observatory (CGRO) had  measured puled $\\gamma$-ray emissions from  younger pulsars (Thompson 2004). The multi-wavelength  observations of the $\\gamma$-ray pulsars have shown that the spectra of the non-thermal emissions extend in $\\gamma$-ray through  optical bands. The observations have also revealed the pulse profiles in optical through  $\\gamma$-ray bands,  and the polarization properties of the pulsed optical emissions from the $\\gamma$-ray pulsars (Kanbach et al. 2005; Mignani et al 2007). In the future, furthermore, the polarization of the X-rays and $\\gamma$-rays from the pulsars will probably be able to be  measured by ongoing projects (Kamae et al. 2007; Chang et al. 2007). The  multi-wave length observations on the  spectrum, the pulse profile and the polarization allow us to  perform a comprehensive theoretical discussion for the mechanisms of the particle acceleration and of the non-thermal emission in the pulsar magnetospheres.  The particle acceleration  and the non-thermal emission processes in the pulsar magnetosphere have  been mainly argued with the polar cap accelerator model (Sturrock 1971; Ruderman \\& Sutherland 1975) and the outer gap accelerator model (Cheng, Ho \\& Ruderman 1986a, 1986b). Also, the slot gap model (Muslimov and Harding 2004), which is an extension model of the polar cap model,  was  proposed. All models predict an acceleration of the particles by an  electric field parallel to the magnetic field. In the pulsar magnetosphere, the accelerating electric field arises in a charge depletion region from the so called Goldreich-Julian charge density (Goldreich \\& Julian 1969), and the  strong acceleration region in the magnetosphere  depends on th e model. The polar cap and the slot gap models predict a strong acceleration within several stellar radii on the polar cap, and the outer gap model predicts a strong acceleration beyond the null charge surface, on which the Goldreich-Julian charge density becomes zero.  As CGRO had observed,  most of the $\\gamma$-ray pulsars have a double peak structure in the pulse profile in $\\gamma$-ray bands. The outer gap model has been successful in explaining the observed double peak structure  (Romani \\& Yadigaroglu 1995). With the outer gap model, the double peak structure is naturally produced as an effect of the aberration and the time delay  of the emitted photons by the outgoing particles that move toward the light cylinder, which is defined by the positions, on which the rotating speed with the star is equal to speed of the light. The axial distance of the light cylinder is $R_{lc}=c/\\Omega$, where $c$ is the speed of the light, and $\\Omega$ is the angular velocity of the neutron star.  Because the Crab pulsar has the pulse profiles with the double peak structure in whole energy bands, the outer gap model  can explain the pulse profiles of the Crab pulsar in optical through $\\gamma$-ray bands (Takata \\& Chang 2007). Furthermore, the outer gap model can explain the observed  spectra, and the polarization characteristics in optical bands for the Crab pulsar (Cheng et al. 2000; Hirotani 2007; Jia, et al 2007; Takata et al, 2007; Takata \\& Chang 2007; Tang et al. 2007). Therefore, the outer gap accelerator model has successfully explained the results of the multi-wave length observations for the Crab pulsar.   For the Vela pulsar, multi-peak structure in pulse profiles in X-ray, UV and optical bands have been revealed (Harding et al. 2002; Romani et al. 2005), although the pulse profile in $\\gamma$-ray bands has two peaks in a single period (Fierro et al. 1998). For example, Harding et al. (2002) analysed RXTE data and revealed  the pulse profile, which  has at least five peaks in  a single period. Because the  two peaks in the five peaks are in phase with  the two peaks in the $\\gamma$-ray bands, the two peaks  will be explained by the traditional outer gap  model with the outward emissions from the outer gap. However, because other three peaks are not expected by the traditional outer gap model,  the origin of the three peaks has not been understood, so far.  The inward emissions emitted by particles accelerated toward stellar surface will give  one possibility to explain the unexpected three peaks of pulse profile in the X-ray bands of the Vela pulsar. Because the primary particles are produced near the inner boundary of the gap, the ingoing particles feel a small part of  whole potential drop in the gap before  escaping the gap from the inner boundary, although the outward moving particles can feel whole potential drop before escaping the gap from the outer boundary. In previous studies, therefore, the contribution of the inward emissions on the spectral calculation has been ignored by assuming that its flux is much smaller than that of the outward emissions. However, it will be true only in the $\\gamma$-ray bands with the curvature radiation. Because main emission mechanism in the X-ray bands will be the synchrotron radiation, the inward synchrotron emissions with a stronger magnetic field will be efficient enough to contribute on the observed emissions below $\\gamma$-ray bands.  Although  the inward emissions have been consistently dealt in the electrodynamic studies (Takata et al. 2006; Hirotani, 2007), only the outward emissions were taken into account for computing the spectrum.  In this paper, we calculate the spectrum  in optical through $\\gamma$-ray bands  by taking into account both  inward and outward emissions. Specifically, we solve the electrodynamics in the outer gap accelerator in the two-dimensional plane, which includes the rotational axis and the magnetic axis.  We calculate the spectrum of the curvature radiation, synchrotron radiation, and  inverse~-Compton process for the primary and the secondary particles. We also discuss the expected pulse profiles from optical through $\\gamma$-ray bands with a  three-dimensional model for comparison with the observations. Furthermore, we apply the theory to PSRs B1706-44 and B1951+32 to calculate the spectrum, because the properties of optical and X-ray emissions from the two pulsars have  not been understood well.  In section~2, we describe our two-dimensional electrodynamic model following Takata et al. (2004, 2006). In section~3, we apply  the theory to the Vela pulsar,  and we discuss the  spectrum and the pulse profile in optical through $\\gamma$-ray bands. We also show the expected spectra of  PSRs B1706-44 and B1951+32. ",
        "conclusions": "\\subsection{Phase resolved-spectrum} With the RXTE observations, Harding et al. (2002) proposed the spectral index $s$, which is defined by $I_{\\nu}\\propto \\nu^{-s}$, of  $s\\sim 1$ for  the phase-resolved spectrum of the peak (Pk~2-soft in figure~1 of Harding et al, 2002), which appears between the two peaks, whose phases are aligned with the peak positions in $\\gamma$-ray bands. As we discussed in section~\\ref{pulsep}, the  preset model explains emission mechanism of RXTE Pk~2-soft with the synchrotron emissions of the ingoing particles, which create a peak (Figure~\\ref{pulse}, IP1) in the pulse profile. Although the spectrum in Figure~\\ref{spe1} is calculated using two-dimensional model, we may be able to read the spectral index of the phase-resolved spectra of the peak IP1 from Figure~\\ref{spe1} ( from the slope of the thin dashed-dotted line in the right panel), which predicts the spectral index $s\\sim 0$ for the IP1 emissions. For the reason of this discrepancy between the spectral indexes of RXTE and the present model, we argue that the three-dimensional structure effects the phase-resolved spectrum, and/or that  RXTE observations might  not determine the spectral behavior in the X-ray bands very well due to a bright synchrotron X-ray nebula. The present model predicts the inward synchrotron emissions for the ingoing primary particles (Figure~\\ref{spe1}, the thin dashed line in the left panel) contribute to the spectrum in hard X-ray and soft $\\gamma$-ray bands. Therefore it is important to measure the phase-resolved spectra and the pulse profiles of the Vela pulsar in the hard X-ray and soft $\\gamma$-ray bands to see how  the phase-resolved spectra and the pulse profiles evolve.  \\subsection{Correlation between radio emission and the outer gap emission} Recently, the correlation between the arrival time of the radio emissions and the shape of the non-thermal X-ray pulse profile of the Vela pulsar was  discovered (Lommen et al. 2007). Although the radio emission mechanism has not been understood well, the polarization measurement indicates a correlation between the magnetic field geometry of the polar cap region and the swing of the position angle of the polarization of the radio emissions. Therefore, the polar cap accelerator model as the origin of the radio emissions from the pulsar has been widely accepted. A  magnetospheric model having both  the polar cap accelerator and the outer gap accelerator has been proposed upon request to balance between the energy loss rate and the angular momentum loss rate, which is equal to the energy loss rate divided by the angular velocity of the star, of the global magnetosphere (Shibata 1991; 1995). Although  the polar cap  accelerator and the outer gap accelerator can not exist on the same magnetic field lines because of the different current directions in the polar cap and outer gap accelerators, the polar cap accelerator affects the outer gap accelerator to obtain the torque balance. According to this model, therefore,  even though the  non-thermal X-rays originate from the outer gap accelerator, some correlations between X-ray emissions and the radio emission are expected. Furthermore, the present  model also predicts that the outer gap accelerator  affects the polar cap accelerator with the inward emissions, which pass near the stellar surface. The  inward emissions will affect the circumstance around polar cap region through  the pair-creation process and/or the scattering process.   \\subsection{Inward emissions from the Crab pulsar} \\label{crab} Unlike the Vela pulsar, the Crab pulsar has only two peaks in the pule profile in  optical through $\\gamma$-ray bands. This indicates that the observed emissions of the Crab pulsar are dominated by the outward emissions  in whole energy bands. For the Crab pulsar, the  photons above 1~GeV  emitted via the curvature radiation in the outer gap   are converted into the pairs outside the gap by the pair-creation process by the magnetospheric X-ray photons emitted by the secondary pairs. Outer gap model predicts the synchrotron and the inverse~-Compton processes of the secondary pairs explain the observed emissions in optical through $\\gamma$-ray bands. As we discussed in section~\\ref{spectrum} (Figure~\\ref{spe1},right), the outward synchrotron radiation by the outgoing secondary pairs created by the magnetospheric X-rays is much brighter than the inward synchrotron radiation, because the outgoing $\\gamma$-rays emitted by the curvature radiation is much more than the ingoing $\\gamma$-rays. For the Crab pulsar, therefore, we can expect that the outgoing secondary pairs produce  the observed spectrum with the synchrotron and the inverse~-Compton process.  Because the flux of the outward emissions of the secondary particles is much (say one or two order) larger than that of the inward emissions, the pulse profiles have strong two peaks in a single period in optical through $\\gamma$-ray bands for the Crab pulsar. Although we expect contribution of the inward emissions at off-pulse phase, it must be difficult to separate the tiny flux of the inward emissions from off-pulse outward emissions inside null charge surface and/or from the strong background  emissions from the synchrotron nebula."
    },
    "0801/0801.1692_arXiv.txt": {
        "abstract": "We have used archival \\chandra\\ and \\xmm\\ observations of quasars hosting intrinsic narrow UV absorption lines (intrinsic NALs) to carry out an exploratory survey of their X-ray properties. Our sample consists of three intrinsic-NAL quasars and one ``mini-BAL'' quasar, plus four quasars without intrinsic absorption lines for comparison. These were drawn in a systematic manner from an optical/UV-selected sample. The X-ray properties of intrinsic-NAL quasars are indistinguishable from those of ``normal'' quasars. We do not find any excess absorption in quasars with intrinsic NALs, with upper limits of $\\nh \\ls{\\rm a\\; few}\\times 10^{22}~\\cmm$.  We compare the X-ray and UV properties of our sample quasars by plotting the equivalent width and blueshift velocity of the intrinsic NALs and the X-ray spectral index against the ``optical-to-X-ray'' slope, \\aox. When BAL quasars and other AGNs with intrinsic NALs are included, the plots suggest that intrinsic-NAL quasars form an extension of the BAL sequences and tend to bridge the gap between BAL and ``normal'' quasars. Observations of larger samples of intrinsic-NAL quasars are needed to verify these conclusions. We also test two competing scenarios for the location of the NAL gas in an accretion-disk wind. Our results strongly support a location of the NAL gas at high latitudes above the disk, closer to the disk axis than the dense BAL wind.  We detect excess X-ray absorption only in Q0014+8118, which does not host intrinsic NALs.  The absorbing medium very likely corresponds to an intervening system at $z=1.1$, which also produces strong absorption lines in the rest-frame UV spectrum of this quasar. In the appendix we discuss the connection between UV and X-ray attenuation and its effect on \\aox. ",
        "introduction": "\\label{sec:intro}  The {\\it intrinsic} absorption lines found in the rest-frame UV spectra of quasars and active galactic nuclei (AGNs) are classified based on their line widths into broad absorption lines (BALs; FWHM$\\,>2000$~\\kms, found in $\\sim 10\\%$ of all quasars), narrow absorption lines (NALs; FWHM$\\,< 500$~\\kms, found in $\\sim 50\\%$ of quasars at $z\\sim 2$--4; these can be clearly separated from intervening lines via high-resolution spectroscopy), and mini-BALs (intermediate FWHM between BALs and NALs). They are typically blueshifted relative to the quasar and are thought to trace outflows from the quasar central engine. These outflows could be hydromagnetic accretion-disk winds (e.g., Blandford \\& Payne 1982; Emmering, Blandford \\& Shlosman 1992; K\\\"onigl \\& Kartje 1994; Everett 2005), or accretion-disk winds driven by radiation pressure (e.g., Murray et al. 1995; Arav, Li, \\& Begelman 1994, Proga, Stone, \\& Kallman 2000). Alternatively they could be outflows driven by thermal pressure and launched either from the accretion disk itself (e.g., Begelman, McKee, \\& Shields 1983) or from the obscuring torus invoked by AGN unification schemes and driven by thermal pressure (e.g., Balsara \\& Krolik 1993; Krolik \\& Kriss 1995, 2001; Chelouche \\& Netzer 2005). However, thermal accretion disk winds are too hot to produce absorption lines in the UV, while thermal winds from the obscuring torus are too slow to account for the observe velocities of NALs and BALs. Accretion disk winds, whatever their origin, are an integral part of the accretion flow, and may be the source of the broad emission lines that are characteristic of the optical and UV spectra of quasars and AGNs (e.g., Shields 1977; Chiang \\& Murray 1996; Murray \\& Chiang 1997). Moreover, the outflows may be important for cosmology since they deliver energy and momentum to the interstellar and intergalactic media (hereafter, ISM and IGM, respectively) and can affect galaxy evolution (e.g., Granato et al. 2004; Scannapieco \\& Oh 2004; Springel, Di Matteo \\& Hernquist 2005; Chartas et al. 2007a).  Most of our knowledge about intrinsic absorbers is derived from optical/UV observations. A combination of observational results and models suggest that intrinsic absorption lines of different widths represent either different lines of sight through the outflowing wind to the quasar continuum source (Ganguly et al. 2001; Elvis 2000), or different stages in the evolution of the absorbing gas parcels (e.g., Hamann \\& Sabra 2004; Misawa et al. 2005).  X-ray observations of BAL quasars have revealed large columns of nearly-neutral absorbing gas ($\\sim 10^{23-24}$~\\cmm; see, for example, Green \\& Mathur 1996; Gallagher et al. 2002a).  However, little is known about the X-ray properties of quasars with intrinsic NALs and mini-BALs. X-ray spectroscopy of the last two types of quasars can be extremely useful in many respects:  \\begin{itemize}  \\item In general terms, we can compare the X-ray properties of quasars that host intrinsic NALs with those that do not. This comparison allows us to assess whether there is a significant difference between the central engines of the two types of object. If no difference is found, and since intrinsic NALs are ubiquitous in quasar spectra (see Misawa et al 2007b and references therein), we will be led to a picture where intrinsic NALs, hence outflows, are a universal property of quasars.  \\item X-ray spectra can probe highly-ionized media that are not directly probed by optical/UV absorption lines. Moreover, they can yield more direct measurements of the total hydrogen column density of absorbers in a relatively low-ionization state (this cannot be easily done through UV absorption lines, unless the ionization state of the gas is very well constrained). Thus, we can use the X-ray spectra to investigate whether any hot gas (which may represent the bulk of the mass) co-exists (mixed or layered) with the UV absorbers. \\item The column densities determined from the X-ray spectra constitute a direct test of scenarios for the location of the NAL gas in the larger outflow. In particular, Elvis (2000) suggests that the NAL gas is located at very low latitudes above the accretion disk, implying very large column densities (comparable to those found in BAL quasars, i.e., $\\gs 10^{22}$~\\cmm). On the other hand, Ganguly et al. (2001) place the NAL gas at high latitudes above the accretion disk, implicitly suggesting lower column densities.  \\item Several examples of intrinsic NALs at high ejection velocities ($\\sim 0.2c$) are now known (see Misawa et al. 2007b and references therein). It is extremely interesting, therefore, to search quasars with intrinsic NALs for high-velocity X-ray absorption lines (see, for example, Chartas et al. 2002, 2003). Such observations can yield results relevant to cosmology since they can lead to estimates of the mass outflow rate and the kinetic power of the outflow (see, for example, Chartas et al. 2007b).  \\end{itemize}  With the above considerations in mind, we have used archival \\chandra\\ and \\xmm\\ data to carry out an exploratory survey of the X-ray properties of a small, but carefully-selected sample of quasars at $z\\approx 2.5$--3.8.  This sample consists of three quasars hosting intrinsic NALs, one quasar hosting a mini-BAL (HS1603+3820), and four quasars without intrinsic absorption lines for comparison (these were selected in a systematic way, as described in the next section). We compare the properties of the UV NALs (e.g., equivalent width, outflow velocity) with the parameters describing the X-ray spectrum (e.g., photon index, {\\it intrinsic} column density, and optical-to-X-ray slope, \\aox).  Thus, we place NAL quasars in the context of BAL quasars (see Gallagher et al. 2002b; Brandt, Laor, \\& Wills 2000) and investigate whether all types of objects follow the same trends in their properties. We also use the column densities determined from the X-ray spectra to carry out one of the tests outlined, above, namely, we attempt to distinguish between the two different suggestions for the location of the NAL gas in the larger outflow.  In \\S\\ref{sec:sample}, we describe how the sample was selected and summarize the properties of the constituent quasars.  In \\S\\ref{sec:obs} and \\S\\ref{sec:analysis}, we present the analysis of the data (observations, data screening, and model fits to the X-ray spectra). We compare the rest-frame UV and X-ray properties in \\S\\ref{sec:xuv}.  In \\S\\ref{sec:discussion} we summarize our findings, discuss our results, and consider prospects for future work.  We adopt a cosmological model with $H_{0}$ = 75~\\kms~Mpc$^{-1}$, $\\Omega_{m}$=0.3, and $\\Omega_{\\Lambda}$ = 0.7. Throughout this paper we give error bars corresponding to the 90\\% confidence level, unless noted otherwise. A significant part of our discussion makes use of the optical-to-X-ray slope of the spectral energy distribution, \\aox, which is affected by extinction. Thus we examine the effects of extinction on \\aox\\ in detail in the appendix. ",
        "conclusions": "\\label{sec:discussion}  The results presented above show that there are no differences in X-ray properties between quasars with and without intrinsic NALs in our sample. The spectra of all the quasars in our sample can be described by a very simple model, namely a power law modified by absorption by nearly-neutral matter at the redshift of the source.  In all but one quasar, we are only able to obtain upper limits to the column density of the intrinsic X-ray absorber; these limits are of order a few $\\times 10^{22}$~\\cmm.  The possibility that the intrinsic X-ray absorbers are ionized can be neither confirmed nor ruled out based on the data we have analyzed here. The only quasar in which we have found a measurable column density is Q0014+8118 ($\\nh\\approx 1\\times 10^{22}$~\\cmm), which does {\\it not} host any intrinsic NALs. We discuss this quasar further below. Our comparison with larger samples of quasars of comparable luminosity and redshift shows that quasars with intrinsic NALs do not differ from the general population in either their X-ray spectra or their spectral energy distributions as quantified by \\aox. The rest-frame equivalent widths of intrinsic NALs follow the same trend with \\aox\\ as nearby quasars and Seyfert galaxies, suggesting a relation between the medium responsible for the intrinsic NALs and the medium responsible for the X-ray continuum absorption. However, X-ray observations of a larger sample of quasars with intrinsic NALs is sorely needed to strengthen our conclusions.  A conclusion that follows immediately from the above results is that outflows that manifest themselves in the form of intrinsic NALs could be present in all quasars. We showed in our earlier work that intrinsic NALs are very common (occurring in 50\\% of all quasars at $z\\sim 2$--4; Misawa et al. 2007b) and the results presented here indicate that the typical spectral characteristics of quasars in general do not preclude their hosting intrinsic NALs.  \\begin{figure} \\centerline{\\includegraphics[width=3.4in]{f6.eps}} \\caption{Comparison of the properties of quasars in our sample with the properties of BAL quasars. Quasars from our sample with intrinsic NALs are represented by stars, while those without are represented by open stars. For the two quasars that were observed more than once (Q1422$+$2309 and Q1700$+$6416), we plot a single, representative value of \\aox\\ (see discussion in \\S\\ref{sec:xuv} of the text). Filled circles represent BAL quasars from the sample of Gallagher et al. (2006).  {\\it Top Panel:} Variation of the total rest-frame equivalent width (summed over all intrinsic NALs in the same quasar) with \\daox. The rest-frame equivalent width assigned to quasars without intrinsic NALs corresponds to the typical observed-frame detection limit of 0.056~\\AA\\ (see Misawa et al 2007b). {\\it Middle Panel:} Variation of X-ray photon index with \\daox.  Open triangles in this panel represent the extremely red quasars from Hall et al. (2006). {\\it Bottom Panel:} Variation of maximum blueshift velocity of intrinsic NALs with \\daox. One intrinsic NAL quasar from our sample (Q0130$-$4021; $\\vmax\\approx 65,000~\\kms$) is off scale and marked with upward arrow.\\label{fig:correlations}} \\end{figure}  In the remainder of this section, we compare the properties of quasars with intrinsic NALs from our sample with those of BAL quasars studied by other authors, we use our results to constrain scenarios for the geometry and location of the intrinsic NAL gas, and we discuss the properties of the absorber in Q0014+8118. We conclude by noting open questions and possible directions for future work.  \\subsection{Comparison with BAL quasars}\\label{sec:comp}  We compare the properties of the NAL quasars in our sample with those of BAL quasars in the sample of Gallagher et al. (2006) and very red quasars from Hall et al. (2006); their redshifts are between 1.4 and 2.8 and their luminosities are comparable to those of our quasars.  In Figure~\\ref{fig:daox}, we show the distribution of {\\it measured} values of \\daox\\ among BAL quasars on the same scale as the distribution of \\daox\\ among quasars in our sample. There is an obvious difference between the two distributions; a Kolmogorov-Smirnov test yields a chance probability of 0.003 that the intrinsic NAL quasars in our sample and the BAL and very red quasars from the bottom panel of Figure~\\ref{fig:daox} were drawn by chance from the same parent population.  This suggests that the X-ray source in intrinsic NAL quasars is not as heavily obscured as in BAL quasars (see the discussion of the \\daox\\ distribution in BAL quasars by Gallagher et al. 2006).  In Figure~\\ref{fig:correlations}, we plot properties of UV BALs as well as the X-ray photon index for the Gallagher et al. (2006) sample against \\daox. In the same plots we include the quasars from our own sample for comparison. The top two panels of Figure~\\ref{fig:correlations} suggest that quasars with intrinsic NALs fit in a progression between BALs and ``normal'' quasars. As the rest-frame equivalent width of the \\ion{C}{4} absorption line decreases, \\daox\\ increases, in a manner that is qualitatively consistent with a decrease in the column density of the absorber. This picture is supported by the photon index vs \\daox\\ plot, where quasars with intrinsic NALs occupy the low-column-density end of the BAL distribution. It is reassuring that the mini-BAL quasar in our sample, HS1603+3820, falls closer to the BAL locus in both of the top two panels of Figure~\\ref{fig:correlations} (as well as in the top panel of Fig.\\ref{fig:compaox}) than the intrinsic NAL quasars.  In contrast, in the plot of maximum blueshifted velocity vs \\daox\\ shown in the bottom panel of Figure~\\ref{fig:correlations}, quasars with intrinsic NALs do not fit into the BAL sequence, even though this sequence appears to connect smoothly to the ``normal'' quasars. This behavior can be reconciled with the behavior seen in the top two panels of the same figure in the context of an equatorial wind scenario. In the top two panels of Figure~\\ref{fig:correlations} the BAL quasar tracks are defined primarily by the column density of the absorbing medium since $\\log\\wrest$, $\\Gamma$, and \\daox\\ are all measures of the column density (the values of $\\Gamma$ of BAL quasars were inferred from the hardness ratio, thus they are directly affected by absorption and so are the 2~keV flux densities that are based on the same data). On the other hand, the orientation of the outflow relative to the line of sight and the acceleration mechanism play an important role in defining the BAL quasar track in the bottom panel of this figure, since they influence the value of \\vmax\\ (see for example, the discussion in Gallagher et al 2006). Therefore, the deviation of intrinsic NAL quasars from the BAL track suggests that the NAL gas is not part of the dense BAL flow thought to be located near the base of the wind, although it may still be associated with it. In other words, our line of sight through the NAL gas does not generally pass through the dense BAL gas. This conclusion also serves as a test of scenarios for the location of the NAL gas within the greater outflow, as we explain further in \\S\\ref{sec:geom}, below.   \\subsection{Geometry and Location of Absorbing Medium}\\label{sec:geom}  Using our observational results, we can test and constrain scenarios for the location of the intrinsic NAL gas in an (equatorial) accretion-disk wind whose dense, low-latitude parts are likely to give rise to BALs. The two competing scenarios are depicted in Figures~1--3 of Elvis (2000) and in Figure~13 of Ganguly et al. (2001). Elvis (2000) proposed that quasars with intrinsic NALs are viewed at a large inclination angle relative to the disk/wind axis (larger than in the case of BALs).  As a result the column density through the NAL line of sight is $\\nh\\gs 10^{22}$~\\cmm\\ and the apparent blueshift of the intrinsic NALs should be considerably smaller than that of BALs, ($\\voff\\sim10^3$~\\kms). On the other hand, in the Ganguly et al. (2001) picture the NAL gas is located {\\it above} the BAL gas, closer to the axis of the disk/wind. The column density is lower because the NAL gas is distributed in small parcels and the blueshifts can be as high as those found in BALs (or even higher).  Our results favor the Ganguly et al. (2001) scenario for a number of reasons. First, we do not detect large absorbing column densities; we find that $\\nh\\ls {\\rm a\\; few}\\times 10^{22}$~\\cmm\\ in all of our quasars. Second, the observed blueshifts of intrinsic NALs are comparable to or higher than those of BALs, suggesting that the absorbing gas is not associated with the base of the wind. Rather, this gas should have traveled an appreciable distance from its launch point and attained a speed close to its terminal speed (see, for example, Murray et al. 1995 and Hamann 1998). Third, the comparison of intrinsic NALs and BALs presented above (\\S\\ref{sec:comp} and bottom panel of Fig.~\\ref{fig:correlations}) shows that the intrinsic NALs do not connect smoothly to the BAL sequence in the \\vmax--\\daox\\ diagram, suggesting that the NAL gas is not a part of the dense BAL flow.  In contrast, the Ganguly et al. scenario passes these tests and also places the NAL gas far enough away from the central continuum source so that (a) it is not highly ionized, and (b) it can be accelerated to a high speed.   \\subsection{The Case of Q0014+8118}\\label{sec:case}  Q0014+8118 is the only quasar in which we have found excess absorption but it is also a quasar not known to have intrinsic UV NALs. We explore here whether these two observational results can be reconciled.  We have re-examined the rest-frame UV spectrum of this quasar presented in Misawa et al. (2007b) and found no strong absorption lines near the redshift of this quasar, which could be attributed to a neutral absorber with the column density comparable to that measured in the X-ray spectrum. We did, however, find an intervening absorption line system at $z=1.1$ with strong \\ion{Fe}{2}$\\;\\lambda$2600, \\ion{Mg}{2}$\\;\\lambda\\lambda$2796,2803 and \\ion{Mg}{1}$\\;\\lambda$2853 absorption lines with the following rest-frame equivalent widths: $W_{rest}(\\lambda 2600)=2.06\\;$\\AA, $W_{rest}(\\lambda 2796)=3.23\\;$\\AA, and $W_{rest}(\\lambda 2853)=0.67\\;$\\AA. This raises the possibility that the excess X-ray absorption is associated with the lower-redshift intervening system. We consider each of these two scenarios in turn.  To assess whether the X-ray absorber could be intrinsic to the quasar, we have carried out a number of photoionization simulations using the code Cloudy (Ferland et al. 1998) adopting the Mathews \\& Ferland (1987) model for the quasar spectral energy distribution (but modified to match the X-ray photon index of Q0014+8118), a Solar metallicity and abundance pattern, the total Hydrogen column density determined from the X-ray spectrum, and two different values of the density ($10^4$ and $10^8$~\\cmmm; the results turn out not be sensitive to the density). We computed four models with ionization parameters corresponding to $\\log\\xi_{\\rm XSTAR}=1.04$, 1.84, 2.24, and 2.54, which span the range of values inferred for the X-ray absorber (see \\S\\ref{sec:notes}). In all cases, the equivalent widths of the \\ion{C}{4}, \\ion{N}{5}, and \\ion{O}{6} absorption lines are large enough that these lines should have been detected in the high-resolution spectrum presented in Misawa et al. (2007b). Missing these lines would require a rather unlikely set of circumstances, namely: the ionization parameter should be at the high end of the allowed range so that the lines are weak and the outflow velocity should be such that the \\ion{C}{4} line is out of the range of the observed UV spectrum, the \\ion{N}{5} line should fall in the gap between echelle orders, and the \\ion{O}{6} line should be hidden in the Ly$\\alpha$ forest. Therefore, we disfavor the intrinsic interpretation of the absorber.  To assess the intervening absorber hypothesis we make use of the statistical observational results of Rao, Turnshek, \\& Nestor (2006). According to these authors, an intervening absorber with the observed \\ion{Fe}{2}, \\ion{Mg}{2}, and \\ion{Mg}{1} equivalent widths has a 65\\% probability of being a damped Ly$\\alpha$ system (DLA), with an average hydrogen column density of $4\\times 10^{20}$~\\cmm\\ and a column density of $1\\times 10^{21}$~\\cmm\\ being reasonably likely. Moreover, the observed relative strengths of the \\ion{Fe}{2} and \\ion{Mg}{2} lines suggest that the hydrogen column density can plausibly be as high as $5\\times 10^{21}$~\\cmm, while the observed \\ion{Mg}{1} strength suggests a 50\\% probability of $N_{\\rm H} > 1\\times 10^{21}$~\\cmm.  Thus we favor the intervening absorber hypothesis because the known intervening system at $z=1.1$ has a high probability of producing just the X-ray absorption we observe. In this context we may also understand the low attenuation of the rest-frame UV light of this quasar. By comparing the colors of large samples quasars with and without DLAs, Ellison, Hall, \\& Lira (2005) find that the dust content of DLAs is very small and place a limit of $E(B-V)<0.04$ on the color excess that they produce.  \\subsection{Open Questions and Future Prospects}\\label{sec:prospects}  There are a number of outstanding questions that our exploratory survey was not able to address.  \\begin{itemize}  \\item We have not detected an ionized absorber in any of the quasars with intrinsic NALs, even though such features are common in the X-ray spectra of Seyfert galaxies. Contributing factors to this were the low $S/N$ of the X-ray spectra and the high redshift of the quasars themselves (the signature of a warm absorber is most pronounced at energies that are redshifted out of the observable band). Detecting or placing limits on a warm absorbing medium is important because it provides a test of models for the structure of the outflow (especially the ionization structure at high latitudes above the disk; see, for example, Proga et al. 2000). Moreover, a warm medium has been invoked to interpret the variability of mini-BALs (see the discussion of HS1603+3820 by Misawa et al. 2007a).  \\item We have not detected high-velocity X-ray absorption lines in any of our spectra. This is not a surprise since such lines are rare. More specifically, unresolved lines have been found in about half a dozen Seyfert galaxies and quasars at $z<0.1$, including MCG-5-23-16 (Braito et al. 2006), IC~4329a (Markowitz et al. 2006), IRAS~13197-1627 (Dadina \\& Cappi 2004), Mrk~509 (Dadina et al. 2005), and PG~1211+143 (Pounds et al. 2007).  The observed outflow velocities are of order $0.1\\; c$, while the rest-frame equivalent widths typically range between 10 and 100~eV.  Broader, and somewhat stronger absorption lines have also been detected in a handful of high-redshift quasars ($z\\sim 2$--4), namely PG~1115+080 (Chartas et al. 2003, 2007a), H~1413+117 (Chartas et al 2007b), and APM~08279+5255 (Chartas et al. 2002; Hasinger et al. 2002).  In these cases the outflow speeds are of order a few tenths of the speed of light and the rest-frame equivalent widths range between a few hundred eV and a few keV. In addition to the cosmological implications of high-velocity X-ray absorption lines noted in \\S\\ref{sec:intro}, their detection will also help us constrain the ionization state of the NAL gas. The limits we have set for Q0014+8118 are rather stringent ($< 28$~eV), comparable to the weakest lines ever detected. In other quasars in our sample the limits are considerably higher, thus uninteresting, owing to the worse $S/N$ in their X-ray spectra. Nevertheless, it is interesting that the absence of high-velocity X-ray absorption lines in Q0014+8118 coincides with the absence of intrinsic UV NALs.  \\item Our results, especially Figure~\\ref{fig:correlations}, raise the issue of whether the UV and X-ray properties of NAL quasars connect smoothly to those of BAL quasars. Related to this issue is the question of where do mini-BALs fall in the grand scheme of things. Even though we have found some tantalizing trends, our sample of objects is too small for a firm conclusion on these issues. In this respect, further observations of mini-BALs would be particularly useful because they may fill in the gap between NALs and BALs in Figure~\\ref{fig:correlations}.  \\end{itemize}  The above questions can be addressed by employing two complementary strategies in future observations. Long exposures of selected objects can yield high-$S/N$ X-ray spectra that can be used for a sensitive search for warm absorbers or high-velocity X-ray NALs. At the same time, (relatively shallow) observations of larger samples of quasars hosting NALs and especially mini-BALs will allow us to explore connections between the corresponding quasar populations."
    },
    "0801/0801.1237_arXiv.txt": {
        "abstract": "Intrinsic stellar variability can hinder the detection of shallow transits, particularly in space-based data. Therefore, this variability has to be filtered out before running the transit search. Unfortunately, filtering out the low frequency signal of the stellar variability also modifies the transit shape. This results in errors in the measured transit depth and duration used to derive the planet radius, and orbital inclination. We present an evaluation of the magnitude of this effect based on 20 simulated light curves from the CoRoT blind exercise 2 (BT2). We then present an iterative filter which uses the strictly periodic nature of the transits to separate them from other forms of variability, so as to recover the original transit shape before deriving the planet parameters. On average with this filter, we improve the estimation of the transit depth and duration by 15\\% and 10\\% respectively. ",
        "introduction": "\\subsection{Planet parameters from observations and associated errors}   The radius ($R_{\\rm p}$, eq. \\ref{eq1}) and the mass ($M_{\\rm p}$, eq. \\ref{eq2}) of an exoplanet can be fully solved when measuring both the flux and the radial velocity variations of the parent star due to its orbiting planetary companion.  \\begin{equation} R_\\textrm{p}=R_{\\star} \\sqrt{\\frac{\\Delta F}{F}} \\label{eq1} \\end{equation}  \\begin{equation} M_\\textrm{p}=M_{\\star}^{\\frac{2}{3}}\\frac{K}{\\sin{i}}\\left(\\frac{P}{4\\pi G}\\right)^{\\frac{1}{3}} \\label{eq2} \\end{equation}  \\noindent where, $R_{\\star}$ and $M_{\\star}$ are the radius and the mass of the parent star, $\\frac{\\Delta F}{F}$ the flux variation due to the planet transiting the disc of its parent star, $K$ the amplitude of the radial velocity variation of the parent star due the gravitational influence of its orbiting planet, $i$ and $P$ the orbital inclination and period of the planet, and $G$ the gravitational constant. $\\frac{\\Delta F}{F}$, $P$, and  $i$ can be measured from the light curve. A common way to measure $R_{\\star}$ and $M_{\\star}$ is by comparing the stellar spectrum to stellar atmosphere models, allowing to derive the stellar parameters ($T_{eff}$,  $\\log{g}$, [Fe/H]) used to obtain the stellar mass and radius. $R_{\\star}$ can also be measured more precisely, and without the use of models, with interferometry, or with transit fitting (in the case of high precision light curves).  \\subsection{Planet parameters and planet evolution and formation models}  Improving the precision on observational planet masses and radii is important for both planet structure and planet formation models. The internal structure of a planet can be studied by comparing its mass and radius to model predictions of planets with different composition. Determining planet structure is important to derive observational statistics on planet types, which can then be compared to the predictions of planet formation models. \\cite{SKHal07} show that to determinate the composition of sub-Uranus planets, error bars of 2\\% on the planet parameters are required. The current uncertainties on planet masses and radii are of the order of 10\\%. Improving these measurements is thus vital to help confirm the models.  \\subsection{Sources of uncertainties on planet parameters}  The uncertainties on the planet mass, radius, and inclination depend on the uncertainties on the host star mass and radius, on the uncertainties on the transit parameters ($\\frac{\\Delta F}{F}$, total transit duration), and on the uncertainties on the radial velocity measurements. For large planets ($>$Jupiter), the uncertainties on the planet mass and radius are mainly due to the uncertainties on the stellar parameters. For sub-Uranus planets around active stars, the uncertainties on the planet mass and radius can be dominated by the uncertainties on the transit parameters. ",
        "conclusions": "Based on 20 simulated CoRoT light curves from the BT2 light curve sample, we have shown that \\cite{AI04} pre-detection filter, used to remove stellar variability prior to transit detection,  deforms the transit depth by 20\\% and the transit duration by 15\\%. To circumvent this, we have adapted \\cite{KBN05} iterative filtering method to the case of filtering stellar variability present in space-based light curves. The resulting post-detection iterative filter improves the estimation of the transit depth and duration by 15\\% and 10\\% respectively. \\\\ The two areas where we plan to focus future efforts are {\\it a)} to further automate the filtering process (some user interaction is currently needed to initialize the filter smoothing length), and {\\it b)} to evaluate the improvement in the planet parameters uncertainties resulting from the improvement in the transit parameters from this work. As {\\it b)} depends on the particulars of each system, we plan to derive these uncertainties using Monte Carlo simulations for known planets. Further tests will include using other pre-detection filters, including ones which can be applied to data with significant temporal gaps."
    },
    "0801/0801.2287_arXiv.txt": {
        "abstract": "{ Neuh\\\"auser et al. (2005) presented direct imaging evidence for a sub-stellar companion to the young T~Tauri star GQ Lup. Common proper motion was highly significant, but no orbital motion was detected. Faint luminosity, low gravity, and a late-M/early-L spectral type indicated that the companion is either a planet or a brown dwarf. } { We have monitored GQ Lup and its companion in order to detect orbital and parallactic motion and variability in its brightness. We also search for closer and fainter companions. } { We have taken six more images with the VLT Adaptive Optics instrument NACO from May 2005 to Feb 2007, always with the same calibration binary from Hipparcos for both astrometric and photometric calibration. By adding up all the images taken so far, we search for additional companions. } { The position of GQ Lup A and its companion compared to a nearby non-moving background object varies as expected for parallactic motion by about one pixel ($2 \\cdot \\pi$ with parallax $\\pi$). We could not find evidence for variability of the GQ Lup companion in the K$_{\\rm s}$-band (standard deviation being $\\pm 0.08$ mag), which may be due to large error bars. No additional companions are found with deep imaging. } { There is now exceedingly high significance for common proper motion of GQ Lup A and its companion. In addition, we see for the first time an indication for orbital motion ($\\sim 2$ to 3 mas/yr decrease in separation, but no significant change in the position angle), consistent with a near edge-on or highly eccentric orbit. We measured the parallax for GQ Lup A to be $\\pi = 6.4 \\pm 1.9$ mas (i.e. $156 \\pm 50$ pc) and for the GQ Lup companion to be $7.2 \\pm 2.1$ mas (i.e. $139 \\pm 45$ pc), both consistent with being in the Lupus I cloud and bound to each other. }  \\titlerunning{Astrometric and photometric monitoring of the GQ Lup system} ",
        "introduction": "Based on three epochs of imaging data spanning five years, Neuh\\\"auser et al. (2005, henceforth N05) presented evidence that the $\\le $ few Myr young T Tauri star GQ Lup has a co-moving companion with K$_{\\rm s} \\simeq 13.1$ mag about 0.7$^{\\prime \\prime}$ west ($\\sim 100$ AU at $\\sim 140$ pc) of the primary star. A low-resolution spectrum gave evidence that it has a late-M to early-L spectral type. Temperature and luminosity can be used to estimate the mass via theoretical evolutionary models. They are, however, very uncertain at young ages (up to at least $\\sim 10$ Myr; Chabrier et al. 2005). Using the Wuchterl \\& Tscharnuter (2003) model extended to planetary masses and few Myr of age, the GQ Lup companion could be a 1 to 3 M$_{\\rm Jup}$ object; whereas for the Tucson group models (Burrows et al. 1997), it is few to $\\sim 30$ M$_{\\rm Jup}$; and for the Lyon group models (Baraffe et al. 2002), it is few to $\\sim 40$ M$_{\\rm Jup}$. A comparison with the GAIA-1 model atmospheres indicated low gravity and, hence, a young age and very low mass (N05). Marois et al. (2007) re-analyzed archival HST and Subaru data to study the spectral energy distribution of GQ Lup A and its companion, showing possible evidence for an excess in the L- and R-bands (for the companion), possibly due to a disk and H$\\alpha$ emission; a fit of the data to the dusty GAIA model atmospheres confirmed the temperature and luminosity of the companion given in N05 and revised its radius to $0.38 \\pm 0.05$~R$_{\\odot}$, yielding a mass between 10 and 20 M$_{\\rm Jup}$. McElwain et al. (2007) then obtained higher-resolution spectra (R $\\simeq 2000$) of the GQ Lup companion in the J- and H-bands with the integral field spectrograph OSIRIS at Keck and found a slightly higher temperature (M6-L0) than in N05 (M9-L4), explained by the fact that H$_{2}$ collision-induced absorption is important for low gravity objects according to Kirkpatrick et al. (2006), but not considered in N05.  Higher-resolution spectra in the J-, H-, and K-bands taken with VLT/Sinfoni compared to GAIA-2 models could better constrain the parameters of the companion: the temperature is $2650 \\pm 100$ K, the gravity $\\log~g = 3.7 \\pm 0.5$ dex, and the radius $3.5^{+1.50} _{-1.03}$ R$_{\\rm Jup}$ (Seifahrt et al. 2007). Hence, its mass can be as low as few M$_{\\rm Jup}$, but also much higher. Comparing its parameters with 2M0535 A \\& B, where masses have been determined dynamically (an eclipsing double-lined binary brown dwarf in Orion, similar age as GQ Lup, Stassun et al. 2006), could give an upper mass limit of 35 M$_{\\rm Jup}$ (Seifahrt et al. 2007). Thus, the companion to GQ Lup can be regarded as a {\\em planet candidate} according to the best guess value of its mass, which is at or below the brown dwarf desert ($\\sim 30$ M$_{\\rm Jup}$; Grether \\& Lineweaver 2006), proposed as the deviding line between planets and brown dwarfs. Also, it is possible that the true mass is below 13 M$_{\\rm Jup}$, a more conservative upper mass limit for planets.   The large error in the luminosity of the companion, which is used for the mass estimate from evolutionary models, is mostly due to a large distance uncertainty, assuming that the object is in Lupus I ($\\pm 50$ pc, N05). Hence, a direct parallax measurement would yield a stronger contraint on luminosity and, hence, mass.  To finally confirm that the fainter object near GQ Lup is really a bound companion (rather than, e.g., another member of the Lupus I cloud not orbiting GQ Lup), one would need to see orbital motion.  For both measurements, parallax and orbital motion, we have monitored GQ Lup and its companion from 2005 to 2007 by taking six new images with Adaptive Optics. In Sects. 2 \\& 3, we explain the observations, data reduction and astrometric results.   We also use the data to monitor the brightness of GQ Lup A and its companion and search for photometric variability (see Sect. 4). We note that Seifahrt et al. (2007) found emission lines in the near-infrared spectra of the companion, indicative of accretion, so some variability is expected. We then add up all imaging data thus far available to obtain a very deep and high-dynamic-range image to search for additional, fainter and/or closer companions (see Sect. 5). We note that both Debes \\& Sigurdsson (2006) and Boss (2006) argued that if the GQ Lup companion is a planet, it should have been moved to its current large separation ($\\sim 100$ AU) from an originally closer orbit by an encounter with another more massive proto-planet, which may be detectable. We use the newly determined parameters to re-estimate some physical parameters of the GQ Lup companion in Sect. 6 and summarize our results at the end. ",
        "conclusions": "Given the newly constrained mean K$_{\\rm s}$-band magnitude of $13.39 \\pm 0.08$ mag and distance measured for the companion of $139 \\pm 45$ pc (this paper) and with B.C.$_{\\rm K}= 3.0 \\pm 0.1$ mag (following Golimowski et al. 2004) and the temperature of the companion newly constrained in Seifahrt et al. (2007), we can re-estimate the luminosity of the companion to be $\\log (L_{\\rm bol}/L_{\\odot}) = -2.375 \\pm 0.245$, similar to the value in N05, but with a smaller error bar.  With the temperature $2650 \\pm 100$~K and gravity $\\log g = 3.7 \\pm 0.5$ dex (g in $g/cm^{2}$) for the GQ Lup companion (Seifahrt et al. 2007), we can use luminosity and temperature to re-calculate its radius to be $3.0 \\pm 0.5$ R$_{\\rm Jup}$. With radius and gravity, we obtain a mass of $\\sim 20$ M$_{\\rm Jup}$ with a possible minimum (value $- 1 \\sigma$) being only few M$_{\\rm Jup}$, and the maximum being around the sub-stellar limit. However, the upper mass limit for the GQ Lup companion is still $36 \\pm 3$~M$_{\\rm Jup}$, because the GQ Lup companion is smaller, cooler, and fainter than both components in the eclipsing double-lined spectroscopic binary brown dwarf 2M0535 (Stassun et al. 2006), which has a similar age as GQ Lup, as already noticed by Seifahrt et al. (2007).  We can also use luminosity L, temperature T, gravity g, radius R, and the age of the young T Tauri star GQ Lup ($\\le 2$ Myr, N05, having strong IR excess) to estimate the mass of the companion from theoretical evolutionary models (as done in N05 with the original, less constrained parameters): from Burrows et al. (1997), we consistently obtain for all combinations of L, T, g, R, and age a mass of $\\sim 20$ M$_{\\rm Jup}$, and from Baraffe et al. (2002), we consistently get for all combinations L, T, and age a mass of $\\sim 20$ M$_{\\rm Jup}$. According to the calculations following Wuchterl \\& Tscharnuter (2003), as plotted in Fig. 4 in N05, the GQ Lup companion would have $\\sim 5$ M$_{\\rm Jup}$. It may be seen as intriguing that both the atmospheric and the conventional evolutionary models consistently give $\\sim 20$ M$_{\\rm Jup}$ as the best value. However, we note that the models by Burrows et al. and Baraffe et al. may not be valid for very young objects ($\\le 10$ Myr), as initial conditions are not taken into account, and none of the models used are tested positively for very low-mass objects or calibrated.  A better mass estimate can be obtained in the future by comparison with more very young objects with dynamically determined masses and/or atmospheric or evolutionary models that are calibrated.  We can summarize our results as follows. \\begin{enumerate} \\item With precise relative astrometry at six new epochs, we have detected small deviations in the separation between GQ Lup A and its companion, consistent with orbital motion of $\\sim 2$ to 3 mas/yr. \\item It remains unclear whether the slight decrease in separation observed is due to orbital motion around each other (bound) or due to different parallaxes (unbound). The latter case is, however, less likely. \\item By comparing the position of GQ Lup A and its companion to a third object in the small NACO FoV, a background star at 800 to 1700 pc, we could determine the parallaxes of GQ Lup A and its companion corresponding to a distance of $156 \\pm 50$ pc for A and $139 \\pm 45$ pc for the companion. \\item Apart from that background star, no additional companion candidates are detected. Outside of 115 mas (18 AU), we can exclude further companion candidates as bright or brighter than the known co-moving companion. \\item We could confirm photometric variability of GQ Lup A in the K$_{\\rm s}$-band (7 exposures over 3 years); variability of the companion is smaller than $\\pm 0.08$ mag. \\item From the newly constrained distance and the mean K$_{\\rm s}$-band magnitude, we re-estimate luminosity, radius, and mass of the companion, but we can still not decide whether the companion is below the deuterium burning mass limit (or below the radial velocity brown dwarf desert, which is at $\\sim 30$ M$_{\\rm Jup}$ according to Grether \\& Lineweaver 2006), i.e., whether the companion is of planetary mass or a brown dwarf. The large uncertainty in mass is due to uncertainties in the theoretical models and the gravity and distance measurements. \\end{enumerate}"
    },
    "0801/0801.2764_arXiv.txt": {
        "abstract": "We present high-resolution interferometric imaging of LH\\,850.02, the brightest 850- and 1200-\\micron\\ submillimetre (submm) galaxy in the Lockman Hole. Our observations were made at 890\\,\\micron\\ with the Submillimetre Array (SMA). Our high-resolution submm imaging detects LH\\,850.02 at $\\gsim$6$\\sigma$ as a single compact (size $\\lsim$1\\,arcsec or $\\lsim$8\\,kpc) point source and yields its absolute position to $\\sim$0.2-arcsec accuracy. LH\\,850.02 has two alternative radio counterparts within the SCUBA beam (LH850.02N \\& S), both of which are statistically very unlikely to be so close to the SCUBA source position by chance. However, the precise astrometry from the SMA shows that the submm emission arises entirely from LH850.02N, and is {\\it not} associated with LH850.02S (by far the brighter of the two alternative identifications at 24--\\micron\\ ). Fits to the optical-infrared multi-colour photometry of LH850.02N \\& S indicate that both lie at $z\\approx 3.3$, and are therefore likely to be physically associated. At these redshifts, the 24\\micronend--to--submm flux density ratios suggest that LH\\,850.02N has an Arp\\,220-type starburst-dominated far-IR SED, while LH\\,850.02S is more similar to Mrk\\,231, with less dust-enshrouded star-formation activity, but a significant contribution at 24--\\micron\\ (rest-frame ~$5-6$\\, \\micron ) from an active nucleus. This complex mix of star--formation and AGN activity in multi--component sources may be common in the high redshift ultraluminous galaxy population, and highlights the need for  precise astrometry  from high resolution interferometric imaging for a more complete understanding. ",
        "introduction": "\\begin{figure*} \\epsfig{figure=lh850-2N.ps,width=80mm} \\epsfig{figure=lh850-2S.ps,width=80mm} \\caption{Optical and near-IR photometry for LH\\,850.02N and LH\\,850.02S, along with fits to their spectral energy distributions (SEDs), and photometric redshifts \\citep{dye2008}.  Downward arrows indicate upper limits.} \\label{fig:photoz} \\end{figure*}  It has been well established that up to half of the far-infrared (far-IR) extragalactic background is produced by dusty starbursts and active galactic nuclei \\citep{hauser1998,dwek1998,fixsen1998,pei1999}. A significant fraction of this background was resolved at 850\\,\\micron\\ into discrete point sources \\citep{smail1997,hughes1998,barger1998} by the Submm Common-User Bolometric Array \\citep[SCUBA:][]{holland1999} on the 15-m James Clerk Maxwell Telescope (JCMT).  These so-called submm galaxies (SMGs) are thought to be high-redshift ultraluminous and hyperluminous IR galaxies \\citep[see][]{chapman2005} that represent massive systems in the process of formation \\citep{scott2002,blain2004}, and may dominate cosmic star formation for nearly the first half of the lifetime of the Universe \\citep[$z\\gsim 1$;][]{blain1999,blain2002}.  Since their discovery a decade ago, a series of surveys using SCUBA at 850\\,\\micron\\ \\citep{barger1999,eales1999,eales2000, chapman2002,cowie2002,scott2002,borys2002,webb2003,borys2003,serjeant2003,wang2004,knudsen2006,coppin2006,knudsen2008} and mm instruments at longer wavelengths \\citep{greve2004,dannerbauer2004, laurent2005,dscott2006,bertoldi2007,scott2008} have amassed catalogues of hundreds of SMGs. Unfortunately, the detailed study of these objects has been hindered somewhat by the poor angular resolution ($\\sim$11--18\\,arcsec) of current submm/mm telescopes. This problem was first addressed via deep radio continuum surveys, which exploited the radio--FIR correlation \\citep[see][for a review]{condon1992} in combination with statistical arguments \\citep[e.g.,][]{ivison2002, ivison2007} to associate 1.4-GHz sources with the submm emission.  Radio counterparts allowed SMGs to be localised with sub-arcsec precision and hence allowed a more detailed study of their properties. Optical spectroscopy of these radio-identified samples confirmed that SMGs lie preferentially at high redshift \\citep[median $z\\sim 2.3$;][]{chapman2005} and enabled CO spectroscopy which revealed them to be compact, massive, gas-rich, possibly merging systems \\citep[e.g.,][]{neri2003,sheth2004,kneib2005,greve2005,tacconi2006}.  Despite the undoubted success of deep radio imaging of SMGs, there is still a clear need for high-resolution {\\it submillimetre/millimetre} observations of at least a subset of the SMGs found in current, complete samples. Specifically, two classes of object require such observations to locate the source of the submillimetre emission with the required astronomical precision for optical/infrared follow-up.  First, because of the rapid dimming of the radio continuum with redshift \\citep[$I_\\nu\\sim (1+z)^{\\alpha-3}$, $\\alpha=-0.8$, where $I_{\\nu}\\propto\\nu^{\\alpha}$;][]{condon1992}, even the deepest existing radio imaging is relatively insensitive to SMGs at $z\\gsim 3$; as a result, typically only around two thirds of SMGs have been detected at radio wavelengths \\citep{ivison2002}.  Other techniques \\citep{ashby2006,pope2006} have been suggested, which make use of data from the IR Array Camera \\citep[IRAC:][]{fazio2004}, in combination with 24-\\micron\\ observations using the Multiband Imaging Photometer \\citep[MIPS:][]{rieke2004}, on board the {\\it Spitzer Space Telescope} to select counterparts.  However, they too may have hidden biases which are difficult to quantify, and the only unambiguous way to select the correct optical counterparts for these radio-unidentified SMGs is via time-intensive submm/mm high-resolution interferometric imaging. Interferometric observations at mm \\citep{downes1999,frayer2000,dannerbauer2002,genzel2003, greve2005,tacconi2006,dannerbauer2008} and submm \\citep[][see also Iono et al. 2008, in prep.]{iono2006,wang2007} wavelengths have successfully detected a growing catalogue of SMGs (in the process also confirming the reliability of the radio--submm association where a unambiguous radio counterpart has already been discovered). Most recently, \\citet{younger2007} followed up a flux-limited sample of seven mm-selected SMGs detected by the AzTEC camera \\citep{wilson2008} at the JCMT -- including five without reliable radio identifications -- at 890\\,\\micron\\ with the Submillimeter Array \\citep[SMA:][]{ho2004}, the counterparts of which suggested a population of very luminous SMGs at higher redshift than radio-identified samples (see also Yun et al. 2008, in prep.). That this bright (median 890-\\micron\\ flux density 12.0\\,mJy), AzTEC-selected sample contains a significant high-redshift tail of SMGs supports earlier evidence that the brightest SMGs may be the most distant \\citep[e.g.\\ Figure~9 of][]{ivison2002,dunlop2001}.  Second, within the radio-identified sub-samples of SMGs there remains some confusion. Specifically, there exists a statistically significant fraction of SMGs which possess more than one statistically robust radio counterpart \\citep[$\\sim20\\pm5\\%$;][]{ivison2007}. Monte--carlo simulations suggest that these associations are observed significantly more frequently than would be expected from chance associations. The nature of these multiply-identified SMGs remains somewhat uncertain, although the steepness of the SMG number counts suggests that the bright submillimetre sources are only rarely expected to arise from the blending/confusion of more moderate luminosity subcomponents. Interferometric imaging of the rest--frame far--IR continuum offers the best way to investigate the true nature of these interesting systems by precisely locating the source of the submm emission \\citep[see e.g., SMM J094303+4700 H6 and H7;][]{tacconi2006}.  In this paper, we present high-resolution SMA 890-\\micron\\ interferometric imaging of LH\\,850.02/LH\\,1200.04 (hereafter simply LH\\,850.02), the brightest source in both the MAMBO 1200-\\micron\\ \\citep{greve2004} and SCUBA 850-\\micron\\ \\citep{coppin2006} maps of the Lockman Hole.  In \\S~\\ref{sec:obs} we describe our observations and data reduction, and in \\S~\\ref{sec:discuss} we discuss some possible implications of our results.  Throughout this paper, we use the Vega magnitude system, and assume a flat concordance cosmology with $\\rm (\\Omega_m,\\Omega_\\Lambda,H_0) = (0.3,0.7,70$ km\\,s$^{-1}$\\,Mpc$^{-1}$). ",
        "conclusions": "We present high-resolution 890-\\micron\\ interferometric imaging of LH\\,850.02, a bright SMG in the LH, with the SMA. We detect LH\\,850.02 at $\\gsim 6\\sigma$ as a single compact (size $\\lsim$1\\,arcsec, or $\\lsim$8\\,kpc) point source, and determine a position accurate to $\\sim$0.2\\,arcsec.  From this, we identify multi-wavelength counterparts and find that only one (LH850.02N) of two candidate radio counterparts is associated with the submm emission. The nearby radio source with strong 24-\\micron\\ emission (LH850.02S) does {\\em not} contribute significantly to the submm flux. In this case, and by analogy other SMGs with multiple radio counterparts, the radio continuum more reliably locates the source of the submm -- and with it the properties of the associated starburst -- than does the mid--IR; similar to the results of \\citet{younger2007}. Since both radio sources have similar photometric redshifts, and therefore may be physically associated, their respective 24\\micronend--to--submm suggest that LH\\,850.02N has an Arp\\,220-type starburst-dominated far-IR SED, while LH\\,850.02S has a Mrk\\,231-type SED with a significant mid-infrared contribution from an active nucleus.  As a result of the relatively shallow, wide-field surveys conducted to date with the AzTEC camera -- wherein objects of this type were first found in significant numbers \\citep{younger2007} -- existing 1100-\\micron\\ samples are typically $\\gsim$2$\\times$ brighter than samples selected with SCUBA or MAMBO. We suggest that the recent prevalence of candidate high-redshift SMGs is more closely related to their high flux densities than to the long survey wavelength, a tendency noted by \\citet{ivison2002} and \\citet{wall2008}. It is therefore perhaps not surprising that we have here constrained the brightest source in the SHADES survey of the LH, LH\\,850.02, to lie at high redshift $z\\gsim 3$."
    },
    "0801/0801.0557_arXiv.txt": {
        "abstract": "I discuss the relationship between mass loss and nucleosynthesis on the Asymptotic Giant Branch (AGB). Because of thermal pulses and possibly other mixing processes, products of nucleosynthesis can be brought to the surface of AGB stars, increasingly so as the star becomes more luminous, cooler, and unstable against pulsation of its tenuous mantle. As a result, mass loss is at its most extreme when dredge-up is too. As the high rate of mass loss truncates AGB evolution, it determines the enrichment of interstellar space with the AGB nucleosynthesis products. The changing composition of the stellar atmosphere also affects the mass-loss process, most obviously in the formation of dust grains --- which play an important r\\^ole in driving the wind of AGB stars. ",
        "introduction": "Asymptotic Giant Branch (AGB) stars span a range in masses from $<1$ M$_\\odot$ to $\\sim8$ M$_\\odot$, which makes them important diagnostics of the star formation history and important players in galactic chemical enrichment from ages as young as $<100$ Myr up to (nearly) a Hubble time. They shed much, often most, of their initial mass in the form of dusty winds before leaving a white dwarf behind.  In this review I will first ask the question why we care about mass loss, with a bias towards those amongst us who are interested in nucleosynthesis. I then briefly discuss what determines the mass-loss rate, and comment on ways in which nucleosynthesis may, or may not, alter the conditions for the mass-loss mechanism. ",
        "conclusions": ""
    },
    "0801/0801.0885_arXiv.txt": {
        "abstract": "Model of supermassive black holes formation inside the clusters of primordial black holes is developed. Namely, it is supposed, that some mass fraction of the universe $\\sim10^{-3}$ is composed of the compact clusters of primordial (relic) black holes, produced during phase transitions in the early universe. These clusters are the centers of dark matter condensation. We model the formation of protogalaxies with masses about $2\\,10^{8}{\\rm M_{\\odot}}$ at the redshift $z=15$. These induced protogalaxies contain central black holes with mass $\\sim10^5{\\rm M_{\\odot}}$ and look like dwarf spheroidal galaxies with central density spike. The subsequent merging of induced protogalaxies and ordinary dark matter haloes corresponds to the standard hierarchical clustering scenario of large-scale structure formation. The coalescence  of primordial black holes results in formation of supermassive black holes in the galactic centers. As a result, the observed correlation between the masses of central black holes and velocity dispersion in the galactic bulges is reproduced. ",
        "introduction": "The problem of galaxy formation with supermassive central black hole (BH) becomes more and more intriguing and ambiguous in view of discovery of distant quasars at redshifts $z>6$ in Sloan Digital Sky Survey \\cite{z6}. The maximum observed red-shift $z=6.41$ belongs to the quasar with luminosity corresponding to the accretion onto BH with the mass $3\\,10^{9}{\\rm M_{\\odot}}$ \\cite{will03}. Such an early formation of BHs with masses $\\sim10^{9}{\\rm M_{\\odot}}$ is a serious difficulty for the standard astrophysical models of supermassive BH formation in galaxies supposing a fast dynamical evolution of the central stellar clusters in the galactic nuclei (see e.~.g. \\cite{spitzer,saslaw,lightshap78,rees84,dokrev} and references therein), a gravitational collapse of supermassive stars and massive gaseous disks in galactic centers (see e.~.g. \\cite{rees84,bintrem87,el2,haehreeesnar98}), the multiple coalescences of stellar mass BHs in galaxies (see e.~.g. \\cite{rees92,MouTan02,GebRicHo02,kawaguchi}) with the subsequent multiple merging of galactic nuclei in collisions of galaxies in clusters (see e.~g. \\cite{ValVal89,cattaneo05,komossa03,smirnova06,springel05,escala05,masjedi06}). All standard astrophysical (or galactic) scenarios of supermassive BH origin predict a rather late time of supermassive BH formation in the galactic nuclei. An other difficulty is that all these astrophysical scenarios are realized only in strongly evolved galactic nuclei. In view of these problems the cosmological scenarios of massive primordial BHs formation become attractive \\cite{zeld67,carr75,DolgovSilk93,quin,Ru1,carr05}. In cosmological scenarios the seeds of supermassive BHs are formed long before the formation of galaxies. These primordial black holes (PBH) can be the centers of baryonic \\cite{Ryan} and dark matter (DM) \\cite{DokEroPAZH} condensation into the growing protogalaxies. There are proposed two alternative possibilities: (i) a formation of initially massive primordial BHs and their successive growth up to $\\sim10^9{\\rm M_{\\odot}}$ due to accretion of ambient  matter or (ii) a formation of small-mass primordial BHs and their subsequent merging into the supermassive ones in the process of hierarchical clustering of protogalaxies.  An effective cosmological mechanism of massive primordial BH formation and their clusteri\\-zation was developed in works \\cite{Ru1,Ru2,KR04}. In this papers the properties of spherically symmetric primordial BH clusters were investigated. As a basic example, a scalar field with the tilted Mexican hat potential had been accepted. The properties of resulting primordial BH clusters appear to be strongly dependent on the value of initial phase. In addition, the properties of these clusters depend on the tilt value of the potential $\\Lambda $ and the scale of symmetry breaking $f$ at the beginning of inflation stage. As a result, the mass distribution of primordial BH clusters could vary in a wide range. In our previous paper \\cite{weqso} we considered the model parameters leading initially to large clusters with a rather heavy mass of the central primordial BH, $\\sim4\\,10^7{\\rm M_{\\odot}}$. These central heavy primordial BH can grow due to accretion up to $\\sim10^{9}{\\rm M_{\\odot}}$ and, therefore, may explain the observed early quasar activity.  The elaboration of a discussed mechanism of cosmological primordial BH formation is far from completion and detailed elaboration. For example, there are no now any physical substantiations for reconstruction of scalar field potential parameters and initial characteristics of primordial BH clusters. It is connected not only with the uncertainties of observational data but also with complexities of phase transition details. For example, the domain walls formed during the phase transition in the early universe has a topology of sphere but with a very complicated surface form. When these closed domain walls are turned out inside the horizon, they become self-gravitating. Inside the horizon domain walls tend to obtain a spherical form due to surface tension, but at the same time they strongly oscillate and generate gravitational and scalar waves. As a result their mass gradually diminish. An approximate consideration of this effect \\cite{Ru2} demonstrate that for a wide range of initial theoretical free parameters there are conditions for formation of cluster with a supermassive primordial BH. For this reason we will not fix here the definite values of free parameters in the discussed cosmological model of massive black hole formation. The influence of non-sphericity of the formed domain walls see in \\cite{Ru2,KR04}. An application of the mechanism found in \\cite{Ru1,Ru2,KR04} is not limited by a specific form of the scalar field potential. Below we demonstrate that substantial number of potentials, like e.~g. those used in hybrid inflation, also result in formation of massive BHs. Moreover, it is hard to avoid primordial BHs overproduction in the early Universe. In fact, any inflationary model using potential with two or more minima must take into account this mechanism of primordial BHs overproduction.  In this paper we choose parameters of the potential which lead to formation of relatively small primordial BH clusters. We suppose that relatively small primordial BH clusters provide the major contribution to initial density pertur\\-bations which afterwards evolve into protogalaxies. The hierarchical clustering of protogalaxies during the cosmological time leads to the observable large scale structure. We describe the gravitational dynamics of DM coupled with primordial BH clusters and demonstrate that a protogalaxy could be formed without any initial fluctuations in DM density. In this case the clusters of primordial BHs play the role of initial fluctuations. Two scenarios of supermassive BHs formation could coexist: (i) the most massive clusters of primordial BHs account for an early quasar activity \\cite{weqso}, but (ii) less massive (considered in this paper) clusters of primordial BHs produce more numerous supermassive BHs observed nowadays in almost all structured galaxies.  There are several stages of BHs and galaxies formation in the described scenario: (i) Formation of closed walls of scalar field just after the end of inflation with a subsequent collapse some of these walls with formation of massive primordial BH cluster with the most massive BH in the center after the horizon crossing according to \\cite{Ru1,Ru2}. (ii) Detachment of the central dense region of the primordial BH cluster from cosmological expansion and virialization. Numerous small-mass BHs merge with a central one. (iii) Detachment of the outer cluster region (where DM particles dominate) from cosmological expansion and a protogalaxy growth. Termination of a protogalaxy growth due to interaction with the surrounding standard DM fluctuations. (iv) Gas cooling and star formation accompanied by the merging of protogalaxies and final formation of modern galaxies. ",
        "conclusions": "We describe here a new model of protogalaxy formation with the cluster of primordial BHs as a source of initial density perturbation. The used mechanism of primordial BH formation \\cite{KR04,Ru2a} provide us with a set of primordial BH clusters of different total mass. This variety of initial conditions leads, therefore, to the variety of protogalaxies from the very beginning of their formation. In this paper we choose for numerical modeling only those BH clusters which produce the large number of small relatively protogalaxies. This model predicts the very early galaxy and quasar formation. An other inevitable consequence of this model it the existence of intermediate mass BHs beyond the dynamical centers of galaxies and in the intergalactic medium. May be one of these type intermediate mass BHs was already observed by the X-ray Chandra telescope in the galaxy M82 \\cite{Kaar00}.  More definitely in this model the protogalaxies are formed at redshift $z=15$. These induced protogalaxies have initially the following parameters: a constituent total mass of DM $M_{\\rm DM}=2.2\\,10^{8}{\\rm M_{\\odot}}$, a virial radius $1.8$~kpc, a mass of central BH $M_{\\mathrm{BH}}=7.2\\,10^{4}M_{\\odot}$. In the following cosmological and dynamical evolution, these protogalaxies are assembled by hierarchical clustering into the nowadays galaxies. The clustering process occurs in a stochastic manner and leads to the specific correlation between the central supermassive BH mass and galactic bulge velocity dispersion \\cite{DokEroPAZH}. An alternative proposed scenario is based on the initial large primordial BH clusters, when a resulting galaxy contains a single primordial BH growing due to accretion of ambient gas and stars and producing early quasar activity \\cite{weqso}.  It is worth to estimate in the framework of our model a probability to find a nowadays galaxy without supermassive BH. Induced galaxies (with a central cluster of primordial BHs) and ordinary small protogalaxies have mass $M_{\\mathrm{DM}}=10^{8}M_{\\odot}$, while the modern galaxies are much more massive, $M_{\\rm DM}=10^{12}{\\rm M}_{\\odot}$. The merging of induced galaxy with an ordinary protogalaxy produces a next generation protogalaxy with the massive central BH. Therefore, about $10^{4}$ collisions is required to form the nowadays galaxy. Suppose that an amount of induced galaxies is about 0.1\\% comparing with the ordinary ones. A corresponding probability to find a modern galaxy without supermassive BH is less than $0.999^{10000}\\simeq 4.5\\,10^{-5}$. Hence, even a very small fraction of induced galaxies is able to explain the observable abundance of AGN."
    },
    "0801/0801.3886_arXiv.txt": {
        "abstract": "{ This report is based on a rapporteur talk presented at the 30th International Cosmic Ray Conference held in Merida, Mexico (July 2007), and covers three of the OG sessions devoted to neutrino, gravitational wave, and $\\gamma$-ray detection.}  \\email{growell@physics.adelaide.edu.au}   \\begin{document} ",
        "introduction": "Summarised here are key results from papers and posters presented in the sessions OG~2.5 (neutrino detection), 2.6 (gravitational wave detection), and 2.7 ($\\gamma$-ray detection). The number of presentations in each session was 75 (OG~2.7), 19 (OG~2.5) and 2 (OG~2.6), with $\\gamma$-ray detection clearly dominating. OG~2.5 was devoted to neutrino/$\\gamma$-ray connections and related experimental and theoretical issues and contains overlaps with the HE~2 session. A more detailed summary of neutrino detectors and related astrophysical theory can be found in the HE~2 rapporteur by Tom Gaisser \\cite{Gaisser-RAPP}. Here, summaries of each sessions are ordered according to the number of contributions. I have kept to citing only ICRC presentations/posters since in most cases a detailed list of references may be found therein. ",
        "conclusions": "Some key conclusions on these sessions primarily devoted to the detection of $\\gamma$-rays and neutrino can be outlined as follows:  \\begin{itemize} \\item Ground-based $\\gamma$-ray astronomy is a now a mature field employing two established and complementary techniques (a) (Stereoscopic) Imaging Atmospheric Cherenkov Imaging; (b) Water Chereknov detection of air shower particles. All of the planned and proposed detectors appear to be making use of at least one of these two techniques.  \\item The next step beyond H.E.S.S., VERITAS, MAGIC-II etc.. is taking shape via continental-wide organisation efforts such as CTA (Europe) and the WhitePaper (USA), along with several other proposals. This is necessary in order to gather community-wide support for the required funding scales of order 100 Million dollar/Euros.  \\item GLAST is ready to go and its launch date is not far away. GLAST will no doubt provide a fresh and clearer look at the MeV to GeV Universe.  \\item Neutrino event rates in forthcoming detectors (IceCube, KM3NeT) appear to be a few per year from the strongest $\\gamma$-ray sources. This would give them a chance to realise discovery of extraterrestrial neutrino sources. In addition the first efforts in coordinated neutrino/$\\gamma$-ray observations have begun. \\end{itemize}  I wish to thank the organisers of the 30th ICRC for giving me the opportunity to summarise these topics. I also thank in particular Simon Swordy and Tom Gaisser for presenting my slides.   \\scriptsize"
    },
    "0801/0801.1551_arXiv.txt": {
        "abstract": "% We have studied the kinematics of the ionized gas and stellar component in Mrk334 using methods of panoramic  (3D) spectroscopy. The observations were performed at the prime focus of the SAO RAS  6-m telescope   with  the integral-field  spectrograph  MPFS  \\citep{mpfs} and with a scanning  Fabry-Perot interferometer (FPI),  installed on the multimode device SCORPIO \\citep{afan}. Based on these data, the  monochromatic  maps and velocity fields in  different emission lines  of the ionized gas were constructed. The diagnostic diagrams have been made based on the emission lines ratios. ",
        "introduction": "Mrk334 is a Sy1.8 galaxy, located at the distance of 88 Mpc (the scale is 430 pc/$''$). This galaxy demonstrates  merging features as tidal tail  and bright region `A' on the west from the nucleus. According \\citet{gonzadelga97},  the region \"A' can be a second nucleus or a remainder of a satellite galaxy, which was devoured by Mrk334. On the MPFS maps in different emission lines, three bright regions can be distinguish: nucleus, regions `A' and `B' (Fig.~\\ref{smir:fig1}). Spectra from these regions differ from each other and the gas ionization state vary from one region to another dramatically: on the diagnostic diagrams, the points belonging to the knot `B' lie in the region of ionization by a nonthermal radiation. All points that correspond to  the knot `A' fall into the photoionization region. Surprising, that the nucleus is located in the HII/LINER boundary i.e., the gas can also be partially ionized here by both shocks and radiation from young hot stars, but not by non-thermal source! Also, on the [SII] ratio map the region with the lowest density is the knot `B', whereas the density in the knot `A'  is much  higher.  \\begin{figure} \\plotone{smirnova1.eps} \\caption{ MFPS flux map (left) and velocity field  (middle) in the [OIII]$\\lambda$ 5007\\AA\\, emission line. The large-scale residual velocity field in the H$\\alpha$ line, isophotes in this line are overlapped (right).} \\label{smir:fig1}       % \\end{figure} ",
        "conclusions": ""
    },
    "0801/0801.4142_arXiv.txt": {
        "abstract": "We explore the noncommutative effect on single field inflation and compare with WMAP five-year data. First, we calculate the noncommutative effect from the potential and dynamical terms, and construct the general form of modified power spectrum. Second, we consider the leading order modification of slow-roll, DBI and K-inflation and unite the modification, which means the modification is nearly model independent at this level. Finally, comparing with the WMAP5 data, we find that the modified can be well realized as the origin of the relative large spectral index and the quite small running. ",
        "introduction": "Inflation \\cite{inflation} is a very successful paradigm for naturally understanding the puzzling aspects of hot big bang cosmology such as flatness, homogeneity and monople problems. As it causally explains the origin of the super-horizon density perturbations, inflation seeds the large structure of the universe we observe today. More like a paradigm than a theory because of the missing fundamental basis, inflation is well established by the observation of cosmic microwave background. Models based on pure de sitter spacetime predict a scale-invariant and adiabatic power spectrum, which can be characterized by $n_s-1=0$. However, this is not the history of our unverse. WMAP and other astrophysical observations predict nontrivial spectral index and running.  The WMAP team has reported the five-year data\\cite{WMAP5yr} and used them to constrain the physics of inflation. For the $\\Lambda $CDM model, WMAP5 data show that the index of power spectrum satisfies $n_{s}=0.963_{-0.015-0.028}^{+0.014+0.029}~~~(1\\sigma, 2\\sigma~ \\text{CL});  \\label{nsnumber1}$. Combining WMAP with SDSS and SNIa, the result is $n_{s}=0.960_{-0.013-0.027}^{+0.014+0.026}~~~(1\\sigma, 2\\sigma~ \\text{CL})$. Though the red power spectrum is still favored at the level of $3\\sigma $ CL, the result is a little bluer than that of WMAP3. And, the running of the spectral index is not favored anymore. With WMAP5 data only, the running is $\\alpha _{s}=\\frac{dn_{s}}{d\\ln k}=-0.037_{-0.028}^{+0.028}~~~(1\\sigma~\\text{CL})$. Combining with SDSS and SNIa data, the result is $\\alpha _{s}=\\frac{dn_{s}}{d\\ln k}=-0.032_{-0.020}^{+0.021}~~~(1\\sigma~\\text{CL})$. We find that the spectral index in WMAP5 is slightly larger than that in WMAP3 and the running is much smaller than that in WMAP3.   Theoretically, we can modify the inflation models, but it will increase the implicity from the aspect of a fundamental theory. In the series works of \\cite{ brandenberger1, easther1, kaloper}, running index emerges from non-Bunch-Davis vacuum and the new physics will imprint on CMB. However, the serious problem is we still know little about de sitter spacetime, which is noted as the ambiguity of vacuum selection.   In the papers \\cite{bean}, another method has been presented. The authors obtained the nontrivial spectral index and its running from the generalized slow-roll inflation with arbitrary sound speed. However, as a new challenge for most inflation models, we expect the spectral index and its running can provide the opportunity to observe the short distance physics and the first moment scenario of our universe.  As a candidate to probe the quantum gravity in string theory, noncommutativity is applied to describe D-brane physics. Considering the string effect and strong string interaction, short distance geometry becomes very different. In addition, while the usual concept of geometry is completely break down near the singularity geometry, noncommutativity is also a good description.  In the works of \\cite{brandenberger2, huang}, the authors introduced noncommutativity, which naturally emerges from string theory, to inflation physics. As the running of spectral index is large, they obtained nice results with WMAP3 experiment data. However, most recent, WMAP5 explores a quite small running index.  In this lecture, we use another method to calculate the modification. We directly calculate the modification of power spectrum from the noncommutative potential and dynamical terms while the quantum modes solution still remains classical. We study the modification of the single field inflation models including slow-roll, DBI and K-inflation. As our calculation includes all the potential and dynamics terms, the modification is exact. We find these three modification can be written in one unite form at leading order. This means that the modification can not be tested by different models. In addition, the leading order of modified power spectrum is $1/k^{2}$, which means that the modification is small. Our calculation shows that the noncommutative potential and dynamical terms can be realized as the possible origin of the spectral index and its running since it is favored by WMAP5 data.   This paper is organized as follows. In section 2, we review the inflation formalism with a general lagrangian. In section 3, we study the noncommutative effect on inflation including the modification of potential and dynamic and reconstruct the power spectrum. In section 4, we study concrete models and compare with the WMAP5 observation. Section 5 contains some conclusion and discussion. ",
        "conclusions": "As an important method to detect short distance physics, noncommutativity naturally emerges from string theory and is applied to inflation physics. Branderberger and Ho found noncommutativity can modify the power spectrum in a significant way. In addition, in papers\\cite{cai}, it has been used to regulate the eternal inflation.  Recent released WMAP5 data favor a red-tilted power spectrum of primordial fluctuations at the level of two standard deviations, which is the same as the WMAP3 result. However, qualitatively, the spectral index is slightly greater and the running is quite small than the three-year value.  Our calculation is quite different. First, we calculated the modification from the potential and dynamical terms while the solution of inflaton remains classic. Second, we calculated the modification of slow-roll, DBI and K-inflation. We found the modification is model independent, as we only considering the leading order. In addition, as the leading order modification of power spectrum is proportional to $1/k^{2}$, the modification of is quite small and can be favored by WMAP5. Third, as we used the string scale in our paper and the correction is proportional to $(\\theta^{12})^{2}$, it could be larger if we have a relative low noncommutative scale.  Finally, as WMAP five-year data provide stringent limits on deviation from the minimal, 6-parameter $\\Lambda $CDM model, we expect more years WMAP and Planck data give us more information about the anisotropy of CMB. In addition, future CMB missions such as ground-based polarization experiment of BICEP and ESA's Planck Surveyor, are also expected to help us faithfully understand the early universe."
    },
    "0801/0801.1767_arXiv.txt": {
        "abstract": "{We exploit a set of high signal-to-noise ($\\sim 70$), low-resolution ($R\\sim 800$) quasar spectra to search for the signature of the so-called proximity effect in the \\ion{H}{i} Ly$\\alpha$ forest. Our sample consists of 17 bright quasars in the redshift range $2.7 < z < 4.1$. Analysing the spectra with the flux transmission technique, we detect the proximity effect in the sample at high significance. We use this to estimate the average intensity of the metagalactic UV background, assuming it to be constant over this redshift range.  We obtain a value of $J = (9 \\pm 4) \\times 10^{-22}$ erg cm$^{-2}$ s$^{-1}$ Hz$^{-1}$ sr$^{-1}$, in good agreement with previous measurements at similar $z$. We then apply the same procedure to individual lines of sight, finding a significant deficit in the effective optical depth close to the emission redshift in every single object except one (which by a different line of evidence does nevertheless show a noticeable proximity effect). Thus, we clearly see the proximity effect as a universal phenomenon associated with individual quasars. Using extensive Monte-Carlo simulations to quantify the error budget, we assess the expected statistical scatter in the strength of the proximity effect due to shot noise (cosmic variance). The observed scatter is larger than the predicted one, so that additional sources of scatter are required. We rule out a dispersion of spectral slopes as a significant contributor. Possible effects are long time-scale variability of the quasars and/or gravitational clustering of Ly$\\alpha$ forest lines. We speculate on the possibility of using the proximity effect as a tool to constrain individual quasar ages, finding that ages between $\\sim 10^{6}$ and $\\sim 10^{8}$ yrs might produce a characteristic signature in the optical depth profile towards the QSO. We identify one possible candidate for this effect in our sample.} ",
        "introduction": "\\begin{figure*} \\sidecaption \\includegraphics*[width=12cm]{7088fig1.eps} \\caption{Example of a quasar spectrum taken from the dataset (HE 0940$-$1050, $z=3.088$). In the upper panel the flux density in units of erg cm$^{-2}$ s$^{-1}$ \\AA$^{-1}$ as function of wavelength is presented with the uncorrected (dotted line), corrected (long-dashed line) continuum level  and the noise of the spectrum (dashed-dotted line). In the lower panel the fitted transmission is shown in gray while the corrected continuum is plotted in black (See Sect.~\\ref{datared} and \\ref{syserrors} for details). In both panels the vertical dotted line represent the emission redshift. } \\label{spec} \\end{figure*}   The multitude of absorption lines seen in the spectra of high redshift quasars gives important information about the state of matter in the universe, tracing the physical conditions of the intergalactic medium (IGM) at various epochs. It is commonly believed that for column densities up to $N_{\\mbox{\\tiny\\ion{H}{i}}} \\approx 10^{17.5}\\,\\mathrm{cm^{-2}}$, the absorbers are in photoionisation equilibrium with a metagalactic ultraviolet background field (UVB), composed of the integral over all sources of UV radiation (essentially, star-forming galaxies and quasars). High-resolution spectra of the Lyman forest provide not only a rather detailed statistical characterisation of the absorber properties such as line number densities as well as temperature and density distributions \\citep[e.g.,][]{kim01}, but also physical parameters such as \\ion{H}{i} and \\ion{He}{ii} photoionisation rates \\citep{rauch97,fardal98}, which directly relate to the intensity of the UVB, as described recently by \\citet{bolton05} invoking hydrodynamical simulations. Independently, the UVB has been successfully synthesised by combining the observed quasar luminosity function and the UV emission from galaxies (although the latter is still very uncertain) with the propagation of diffuse radiation \\citep{haardt96}.  In the vicinity of strong UV sources such as bright quasars, the \\ion{H}{i} photoionisation rate should locally increase, further reducing the density of residual neutral hydrogen.  This enhances the transparency of the IGM to \\ion{H}{i} ionising radiation and should become observable as a weakening of the Lyman forest absorption near such sources. Such an effect has first been noted by \\citet{carswell82} and was later baptised `Inverse' \\citep{murdoch86} or `Proximity Effect' \\citet[][hereafter BDO88]{bajtlik88}. Its prime application has been so far the possibility to derive an independent estimate of the UVB intensity, by measuring the reduction of column densities (against the global evolution of absorption line density increasing with redshift) and combining this with the QSO luminosity at the Lyman limit, assumed to be known. The best constraints on the UVB using the proximity effect stem from the combined analysis of large quasar samples \\citep{cooke97,scott00,liske01}, yielding mostly values consistent with the above quoted other methods. However, the uncertainties are still substantial. Besides the problem that a limited number of lines of sight always suffers from `cosmic variance', there may also be systematic biases. In particular, if QSOs reside in intrinsically overdense environments then the signature of the proximity effect will appear weaker than predicted \\citep{loeb95,rollinde05}. Another uncertainty is the possibly limited lifetime of quasars. On the other hand, the proximity effect may also be used to derive constraints on this important, but largely unknown astrophysical quantity \\citep[e.g.,][]{pentericci02}.  In this paper we present an exploitation of new observational material in terms of the proximity effect (Sect.~\\ref{obs_set}). Rather than the traditional line counting we use the more sensitive flux transmission statistic to search for proximity effect signatures, augmented by extensive Monte-Carlo simulations to calibrate the systematic and statistical errors (Sect.~\\ref{lyaforest}). We deliver our results in Sections \\ref{combined} and \\ref{indiv_qsos}. Firstly we briefly present an analysis of the combined sample of 17 QSO spectra and derive an estimate of the UV background intensity. Secondly we demonstrate that the effect is measurable on single sightlines \\citep{williger94,lu96,savaglio97}, not only statistically in large samples \\citep{bajtlik88, scott00, liske01}.  The effect can actually be systematically detected in each single QSO of our sample (with one special case that is discussed separately). We speculate about the possibilities to detect signatures of finite quasar ages by virtue of the proximity effect.  Throughout this paper, we assume a flat $\\Lambda-$Universe with $\\Omega_\\mathrm{m}=0.3$ and $\\Omega_\\mathrm{\\Lambda}=0.7$ and $H_0=71\\,\\mathrm{km\\,s^{-1}\\,Mpc^{-1}}$.    \\begin{table*}[t] \\small\\centering \\caption[]{Log of observations.} \\label{Tab1} \\begin{tabular}{lcccccclc} \\hline\\hline\\noalign{\\smallskip} QSO &   $V$ mag  &$z_\\mathrm{em}$ &Exp. Time&Seeing & Airmass &Sky Condition$^\\mathrm{a}$ & Obs. Date&Ref. Mag\\\\ &&&(s)&(arcsec)&&& \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} \\object{CTQ 0247}        & 17.4 & 3.025 & 750  & 1.29 & 1.07 & CL, WI     & Dec. 8,  2002 & 4\\\\ \\object{CTQ 1005}       & 18.4 & 3.205 & 1500 & 1.13 & 1.21 & TN         & Jan. 9,  2003 & 3\\\\ \\object{CTQ 0460}        & 17.5 & 3.139 & 900  & 1.45 & 1.10 & PH         & Dec. 23, 2002 & 4\\\\ \\object{H 0055$-$2659}   & 17.5 & 3.665 & 600  & 1.34 & 1.02 & CL, WI     & Dec. 8,  2002 & 4\\\\ \\object{HE 0940$-$1050}  & 16.4 & 3.086 & 600  & 0.86 & 1.12 & TN, TK     & Nov. 26, 2002 & 3\\\\ \\object{HE 2243$-$6031}  & 16.4 & 3.010 & 600  & 1.02 & 1.24 & CL, PH     & Nov. 9,  2002 & 4\\\\ \\object{HE 2347$-$4342}  & 16.7 & 2.885 & 600  & 0.77 & 1.06 & PH         & Nov. 10, 2002 & 1\\\\ \\object{PKS 2126$-$15}   & 17.0 & 3.285 & 600  & 0.97 & 1.20 & PH         & Oct. 29, 2002 & 4\\\\ \\object{Q 0000$-$26}     & 18.0 & 4.098 & 600  & 1.31 & 1.03 & TN, CL     & Nov. 7,  2002 & 4\\\\ \\object{Q 0002$-$422}    & 17.2 & 2.767 & 600  & 1.39 & 1.45 & TK, TN     & Oct. 14, 2002 & 3\\\\ \\object{Q 0347$-$383}    & 17.7 & 3.220 & 800  & 1.45 & 1.15 & TN         & Jan. 9,  2003 & 1\\\\ \\object{Q 0420$-$388}    & 16.9 & 3.120 & 600  & 1.37 & 1.06 & CL, WI     & Dec. 8,  2002 & 1\\\\ \\object{Q 0913$+$0715}   & 17.8 & 2.787 & 800  & 1.45 & 1.18 & TN         & Jan. 9,  2003 & 2\\\\ \\object{Q 1151$+$0651}   & 18.1 & 2.758 & 900  & 0.94 & 1.17 & TN, CL     & Jan. 25, 2003 & 2\\\\ \\object{Q 1209$+$0919}   & 18.5 & 3.291 & 1500 & 0.78 & 1.80 & CL, TN, TK & Jan. 1,  2003 & 2\\\\ \\object{Q 1223$+$1753}   & 18.1 & 2.945 & 900  & 0.66 & 1.35 & TN, CL     & Jan. 25, 2003 & 2\\\\ \\object{Q 2139$-$4434}   & 17.7 & 3.214 & 800  & 0.85 & 1.33 & TN         & Apr. 30, 2003 & 3\\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} \\multicolumn{8}{l}{{}$^\\mathrm{a}$ Legend: PH-Photometric, CL-Clear, TN-Thin cirrus, TK-Thick cirrus, WI-Windy.} \\end{tabular} \\begin{list}{}{} \\item[1:]\\citet{worseck07}: PH conditions. \\item[2:]SDSS. \\item[3:]\\citet{veron06}. \\item[4:]Slit loss corrected only. \\end{list} \\end{table*} ",
        "conclusions": "We have presented new evidence of the line-of-sight proximity effect as a universal phenomenon occurring in the spectra of high-redshift quasars.  Even though our spectra are limited in spectral resolution, their high S/N and the power of the flux transmission method has enabled us to demonstrate the presence of the effect for every single line of sight, for the first time.  Our estimate of the mean UV background intensity for the redshift range $2.7<z<4$ is $\\log(J_\\nu)=-21.03^{+0.15}_{-0.22}$, in very good agreement with literature values for similar redshift ranges. We made a careful assessment of the error budget using extensive Monte-Carlo simulations. The errors are clearly dominated by cosmic variance, which implies that better spectral resolution would not necessarily have a dramatic impact on the measurement accuracy. Of course, high resolution spectra would be advantageous for a more detailed analysis of several other aspects of the proximity effect, such as the effects of gravitational clustering of absorbers near the QSO \\citep{rollinde05}, which may overestimate the value of the UV background by up to a factor of 3 \\citep{loeb95}.  We have quantified the strength of the proximity effect in individual spectra and find that this shows a higher dispersion than expected from only statistical shot noise errors. Among the most likely contributors for this additional dispersion are again gravitational clustering of absorbers near the QSO, or QSO variability over very long timescales; fluctuations of the UV background are also possible but unlikely to play a major role at this redshift range. We presented a speculative, but conceptually simple observational test to search for signatures of finite quasar ages in the optical depth profiles derived for a single quasar, and we even tentatively identified a candidate where such a pattern might be visible."
    },
    "0801/0801.4004_arXiv.txt": {
        "abstract": "{% Observations in polarized emission reveal the existence of large-scale coherent magnetic fields in a wide range of spiral galaxies. Radio-polarization data show that these fields are strongly inclined towards the radial direction, with pitch angles up to $35\\degr$ and thus cannot be explained by differential rotation alone. Global dynamo models describe the generation of the radial magnetic field from the underlying turbulence via the so called $\\alpha$-effect. However, these global models still rely on crude assumptions about the small-scale turbulence. To overcome these restrictions we perform fully dynamical MHD simulations of interstellar turbulence driven by supernova explosions. From our simulations we extract profiles of the contributing diagonal elements of the dynamo $\\alpha$-tensor as functions of galactic height. We also measure the coefficients describing vertical pumping and find that the ratio $\\hat{\\gamma}$ between these two effects has been overestimated in earlier analytical work, where dynamo action seemed impossible. In contradiction to these models based on isolated remnants we always find the pumping to be directed inward. In addition we observe that $\\hat{\\gamma}$ depends on whether clustering in terms of super-bubbles is taken into account. Finally, we apply a test field method to derive a quantitative measure of the turbulent magnetic diffusivity which we determine to be $\\sim 2\\,\\kpc\\kms$.} ",
        "introduction": "%  The last decade with its advances in observational technology has brought increasingly detailed maps of galactic magnetic fields. Radio-polarization data indicate large scale regular fields within the interstellar medium.  For typical spiral galaxies the ISM is highly turbulent \\citep{2004RvMP...76..125M}. The main drivers of the turbulence are thought to be winds of massive stars and explosions of supernovae (SNe) \\citep{2005A&A...436..585D,2006ApJ...653.1266J}, as well as (in less active regions) the magneto-rotational instability (MRI) \\citep*{2004A&A...423L..29D,2005ApJ...629..849P,2007ApJ...663..183P}. Furthermore, \\citet{2004ApJ...605L..33H,2006AN....327..469H} have shown that a Parker-type instability can be driven by cosmic rays, an idea first advocated by \\cite{1992ApJ...401..137P}. It, however, remains a matter of discussion, whether cosmic rays play a significant role in the interstellar dynamics \\citep{2006MNRAS.373..643S}.  The generation of a mean magnetic field from turbulent fluctuations can be explained via the so called $\\alpha$-effect which describes the correlations of the small-scale turbulent velocity and magnetic field giving rise to a mean electromotive force (EMF). In the case of the ISM there are three characteristic asymmetries leading to non-vanishing correlations: (i) the axis of rotation, (ii) the galactic shear gradient, and (iii) the (vertical) gradient in density and turbulence intensity \\citep{1993A&A...269..581R}. The mutual strength of the different terms puts constraints on the operability of a dynamo process. The ratio $\\hat{\\gamma}$ of the diamagnetic pumping over the $\\alpha$-effect has an influence on the efficiency of the dynamo. The strength of the $\\alpha$-effect compared to the shear puts limits on the pitch angle.  Until recently, a direct numerical simulation of the turbulent ISM has been infeasible. Therefore, early theoretical models like the SOCA approach by \\citet{1993A&A...269..581R} predicted the outcome of ISM-turbulence based on simplifying assumptions. In an alternative approach \\citet{1992ApJ...391..188F} derived the dynamo effect for isolated supernova remnants (SNRs) and super-bubbles (SBs). However, these early no-interaction models arrived at prohibitively high values for $\\hat{\\gamma}$. To test these findings first numerical simulations for single SNRs have been performed in 2D by \\citet*{1993A&A...274..757K} and in 3D by \\citet*{1996A&A...305..114Z}. Still the key issue of a dominating turbulent pumping remained.  The situation was somewhat improved by taking into account stratification \\citep[hereafter FER98]{1998A&A...335..488F} yielding a value $\\hat{\\gamma} \\approx 6$. The parameter range for $\\hat{\\gamma}$ permitting dynamo solutions has been explored by \\citet*{1994A&A...286...72S}.  The major limitation to the no-interaction models described above results from the fact that (even for SBs) the explosions are considered as isolated events taking place on a uniform background. But this is not the case for the ISM which due to thermal instability \\citep{1965ApJ...142..531F} is a highly heterogeneous medium. To overcome these limitations we perform direct numerical simulations of the local ISM and use a test field method \\citep{2005AN....326..245S,2007GAFD...101..81S} to obtain the dynamo parameters. ",
        "conclusions": "%  Since our vertical stratification is in radiative equilibrium, we do not observe an initial collapse of the disk. Turbulence builds up smoothly and after 100Myr reaches a quasi steady state. Fig.~\\ref{fig:slices} shows vertical and horizontal slices through the simulation box at $t=112\\Myr$. Most of the material is contained in cold clumps forming a $150-200\\pc$ wide disk. Close to the midplane the network of clumps and filaments is permeated by strong shocks from the SNe which are continously creating hot cavities. Looking at the lower panels (a) and (e) of Fig.~\\ref{fig:slices} one can see that for the region around the midplane there exists a significant correlation between the density and the magnetic field amplitude. The inferred slope of the correlation is somewhat steeper than the Chandrasekhar-Fermi value of $0.5$ but roughly consistent with $|B|\\sim\\rho^{2/3}$ as indicative of compressional amplification.  While single SNRs are largely confined to the midplane, super-bubbles break out of the central disk and drive moderate vertical flows. Their dense shells that are further compressed by shocks will form clouds that can efficiently cool and will in turn rain back into the gravitational potential thus forming what is termed the galactic fountain \\citep{1980ApJ...236..577B}. The Mach number of the flow is mildly supersonic below $1\\,\\kpc$ and becomes subsonic for the outer parts.  If we turn off the SN-clustering the morphology changes quite drastically. Instead of well confined SBs we see more disrupted features and chimney-like structures that channel strong vertical outflows. The velocity dispersion in the hot phase is twice as high as in the clustered case. These differences demonstrate the importance of a proper modeling of clustered explosions.  \\subsection{Thermal distribution and velocity dispersions}  \\begin{figure}\\begin{center} \\includegraphics[width=\\columnwidth]{fig/phs} \\caption{Phase space distribution of the SN-heated plasma. Adjacent plots show the mass- (thick) and volume- (thin line) distribution of density (upper panel) and pressure (right panel). We also plot isothermal contours (labeled in$\\K$), and the equilibrium cooling curve (dashed).} \\label{fig:phs_spc} \\end{center}\\end{figure}  In Fig.~\\ref{fig:phs_spc} we show the distribution of the SN-heated plasma as a function of density and thermal pressure at a time $t=112\\Myr$. The two stable branches of the equilibrium curve are densely populated but there also exists a considerable amount of gas in the radiatively unstable regimes.  Averaged velocity dispersions for the ISM phases are $5\\kms$ (cold), $11\\kms$ (cool), $25\\kms$ (warm), and $40-60\\kms$ (hot). While the values for the latter are consistent with the findings of \\citet{2005A&A...436..585D}, for the cold phases we fall short by a factor of two. We consider this a resolution issue but it might also be related to the different form of the driving. As a function of galactic height the time averaged rms-velocity rises steeply from its midplane value of $20\\kms$ to its peak value of $55\\kms$ which is reached at about $1\\,\\kpc$ and then falls off to about $30\\kms$ at $2\\,\\kpc$. Assuming that diamagnetic pumping follows the inverse gradient in turbulence intensity this implies an inward transport of magnetic flux for the central part of the disk.  \\subsection{Dynamo parameters}  In Fig.~\\ref{fig:alpha_prof} we plot the main $\\alpha$- and $\\eta$-coefficients for our standard run (mod. STD). To our knowledge this is the first time such profiles have been obtained from direct simulations of SN-turbulence. The corresponding vertically averaged amplitudes are listed in Table~\\ref{tab:results} -- for comparison we also list the peak values from FER98 (where $\\eta_t$ is distinguished into horizontal and vertical part). Dynamo numbers $C_{\\alpha}=\\alpha H/\\eta_t$ and $C_{\\Omega}=\\Omega H^2/\\eta_t$ are computed with $H=0.8\\,\\kpc$ and $\\eta_t$ values from the inner part of the disk. In accordance with analytical estimations based on these numbers, the runs without shear are still sub-critical, i.e., no field-amplification is observed. The $\\alpha^2\\Omega$-dynamo, however, does exponentially grow at a time scale of $\\sim 250\\Myr$.  During the course of the present simulations we do not yet reach the equipartition field strength. This implies that the values for the dynamo parameters (as well as the pitch angle) are representative of the unquenched regime. \\begin{figure}\\begin{center} \\includegraphics[width=0.9\\columnwidth]{fig/dyn} \\caption{Dynamo $\\alpha$- and $\\eta$-coefficients for the STD  model. Quantities indicated by the ordinate labels are plotted in dark ($\\alpha_{RR},\\dots$) respectively light ($\\alpha_{\\phi\\phi},\\dots$) colors. Shaded areas indicate the rms-fluctuation when applying the method to 4 temporal sub-intervals.} \\label{fig:alpha_prof} \\end{center}\\end{figure} Our values are still affected by fluctuations and a longer time-base is needed for a more accurate determination. However, we find some interesting trends in the data at hand: The antisymmetric part $\\gamma_z$ of the $\\alpha$-tensor is negative (positive) in the northern (southern) hemisphere, i.e., pumping is directed inward. This is contrary to the predictions by Ferri{\\`e}re for the case of non-interacting SNRs. In our simulations the inward pumping is opposed by an outward advection of the field with the mean flow. As can be seen from the middle panel of Fig.~\\ref{fig:alpha_prof} both terms have about the same magnitude for $z < 1\\,\\kpc$ while in the corona the wind is dominating. The balance in the inner region could be seen as an indication that the effects of turbulent pumping and advection are naturally linked. In consequence vertical transport processes might be of lesser importance than formerly believed.  The importance of modeling spatially coherent SBs can be seen from comparing the first two rows of Table~\\ref{tab:results}. In the case without clustered SNe (mod. NCL) we observe strong vertical streaming motions reflected in high values for $\\gamma_z$ and $\\bar{u}_z$. Also the turbulent diffusivity is high in the disk midplane which is not the case for the other runs where $\\eta_t$ increases with galactic height. From the model with type-II SNe only (mod. SN2) one can learn that the turbulent diffusivity predominantly arises from the more broadly distributed type-I SNe -- despite their lower rate by a factor of eight. The overall level of turbulent diffusion is about $3$--$6\\tms10^{26}\\,\\vis$. The off-diagonal components of the $\\eta$-tensor are somewhat smaller and cannot yet be determined accurately. Following \\citet{1996ApJ...457..798J} we crudely estimate that the measured amount of diffusion suffices to damp short wavelength MRI-modes for reasonably high $\\beta_{\\rm P}$ -- a definite conclusion, however, requires further investigations by means of combined direct simulations.  By comparing $\\alpha_{\\phi\\phi}$ with the SOCA results by \\citet{1993A&A...269..581R} we can estimate the coherence time $\\tau$ and the Coriolis number $\\Omega^*=2\\tau\\Omega$ of our turbulent flow. We find a value of $\\tau\\approx3.6\\Myr$ which is about a factor of three smaller than commonly assumed. The even lower value for our model NCL indicates that a higher coherence time might be achieved by a more realistic prescription for the modeling of SBs. In general the obtained profiles are consistent with their SOCA counterparts. Our obtained value for $\\hat{\\gamma}=|\\alpha_{\\phi R}|/|\\alpha_{\\phi\\phi}|$ agrees with the SOCA value of $\\hat{\\gamma}\\approx2.5$ in \\citet{2004maun.book.....R} and is smaller than the value of $\\hat{\\gamma}\\approx6$ in FER98.  In the lower panel of Fig.~\\ref{fig:alpha_prof} we overplot the scalar expression for the turbulent diffusivity which matches the profiles obtained from the simulations. The peak of the velocity dispersion close to the midplane is an artifact due to the static distribution of the SNe resulting in a disruption of the central disk. This issue has meanwhile been resolved.  \\begin{table*}\\begin{center}% \\begin{tabular}{|l|c|c|c|c|c|c|c|c|c|c|}\\hline &$|\\alpha_{RR}|$&$|\\alpha_{\\phi\\phi}|$&$|\\gamma_z|$&$\\hat{\\gamma}$ &\\multicolumn{2}{c|}{$\\eta_{t}$}&$\\tau$&$\\Omega^*$&$C_{\\alpha}$ \\tabularnewline & $[\\kms]$ & $[\\kms]$ &$[\\kms]$ & &\\multicolumn{2}{c|}{$[\\kpc\\kms]$} & $[\\Myr]$ & & \\tabularnewline\\hline\\hline STD & 3.1 & 3.8 & 7.8 & 2.1 & ~~0.9~~ & ~~2.7~~ & 3.6 & 0.7 & 3.3 \\tabularnewline\\hline NCL & 7.3 & 4.3 & 24.7  & 5.9 & 2.4 & 2.2 & 2.8 & 0.6 & 1.4 \\tabularnewline\\hline SN2 & 0.8 & 0.8 & 1.8 & 2.5 & 0.6 & 1.0 & 3.6 & 0.7 & 1.2 \\tabularnewline\\hline KIN & 3.2 & 2.0 & 8.3 & 4.1 & 0.6 & 1.2 & 3.4 & 0.7 & 2.6 \\tabularnewline\\hline\\hline FER98 & 6.0 &  2.6 & 16.0 &  6.2&  \\multicolumn{2}{c|}{3 -- 18} & -- & -- & 0.5 \\tabularnewline\\hline \\end{tabular} \\end{center} \\caption{Mean values of extracted parameters averaged over $\\approx150\\Myr$. Values for $\\eta_t$ apply to the inner ($|z|\\le 0.8\\,\\kpc$) respectively outer region of the disk. Coherence time $\\tau$ and Coriolis number $\\Omega^*$ are estimated from a comparison with SOCA-profiles for $\\alpha_{\\phi\\phi}$. The coefficients could not be obtained for model SHR. \\label{tab:results}} \\end{table*}  \\subsection{Pitch angles} In our models without shear we observe pitch angles ranging from $+10\\degr$ in the far outer regions to $-45\\degr$ near the midplane. This is hardly surprising since also the radial and azimuthal dynamo effects are found to be of equal strength in this case. If we include galactic shear with $q=-1$, pitch angles become solely negative reaching values up to $-30\\degr$ as consistent with observations. The minimum of $-10\\degr$ is found at the midplane. In comparison, \\cite{2006AN....327..469H} in their simulations of a cosmic ray driven galactic dynamo find the azimuthal field to be dominating by a factor of ten which corresponds to pitch angles of only $\\sim 5\\degr$."
    },
    "0801/0801.3412_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:intro} \\index{Introduction} Bars are a very common feature of disc galaxies. In a sample of $186$ spirals, $56\\%$ of the galaxies in the near infrared are strongly barred, while an additional $16\\%$ are weakly barred \\cite{esk00}. A large fraction of barred galaxies show two clearly defined spiral arms \\cite{elm82}, departing from the end of the bar. This is the case for instance in NGC~1300, NGC~1365 and NGC~7552. Spiral arms are believed to be density waves in a disc galaxy \\cite{lind63}. In \\cite{too69}, Toomre found that the spiral waves propagate towards the principal Lindblad resonances of the galaxy, where they damp down, and thus concludes that long-lived spirals need some replenishment. Danby argued that orbits in the gravitational potential of a bar play an important role in the formation of arms \\cite{dan65} and Kaufmann \\& Contopoulos argue that, in self-consistent models for three real barred spiral galaxies, spiral arms are supported also by chaotic orbits \\cite{kau96}.  The origin of rings has been studied by Schwarz who calculated the response of a gaseous disc galaxy to a bar perturbation \\cite{sch81,sch84,sch85}. He proposed that ring-like patterns are associated to the principal orbital resonances, namely ILR (Inner Lindblad Resonance), CR (Corotation Resonance), and OLR (Outer Lindblad Resonance). There are different types of outer rings. Buta classified them according to the relative orientation of the ring and bar major axes \\cite{but95}. If these two are perpendicular, the outer ring is classified as $R_1$. If they are parallel, the outer ring is classified as $R_2$. Finally, if both types of rings are present in the galaxy, the outer ring is classified as $R_1R_2$.  In Romero-G\\'omez {\\it et al.} \\cite{rom06,rom07}, we note that spiral arms and rings emanate from the ends of the bar and we propose that rings and spiral arms are the result of the orbital motion driven by the invariant manifolds associated to the Lyapunov periodic orbits around the unstable equilibrium points. In Romero-G\\'omez {\\it et al.} \\cite{rom06}, we fix a barred galaxy potential and we study the dynamics around the unstable equilibrium points. We give a detailed definition of the invariant manifolds associated to a Lyapunov periodic orbit. For the model considered, the invariant manifolds delineate well the loci of an $rR_1$ ring structure, i.e. a structure with an inner ring (r) and an outer ring of the type $R_1$. In Romero-G\\'omez {\\it et al.} \\cite{rom07}, we construct families of models based on simple, yet realistic, barred galaxy potentials. In each family, we vary one of the free parameters of the potential and keep the remaining fixed. For each model, we numerically compute the orbital structure associated to the invariant manifolds. In this way, we are able to study the influence of each model parameter on the global morphologies delineated by the invariant manifolds.  Voglis, Stavropoulos \\& Kalapotharakos study the chaotic motion present in self-consistent models of both rotating and non-rotating galaxies, concluding that rotating models are characterised by larger fractions of mass in chaotic motion \\cite{vog06a}. Patsis argues that the spiral arms of NGC~4314 are due to chaotic orbits and, to show it, he computes families of orbits with initial conditions near the unstable equilibrium points \\cite{pat06}. Voglis, Tsoutsis \\& Efthymiopoulos \\cite{vog06b} reproduce a spiral pattern found in a self-consistent simulation using the apocentric invariant manifolds of the short-period family of unstable periodic orbits. They give the angular position of the apocentres, which is where they state the stars spend a large part of their radial period, as a soliton-type solution of the Sine-Gordon equation.  In Sect.~\\ref{sec:mod}, we first present the galactic models used in the computations and the equations of motion. In Sect.~\\ref{sec:inv}, we give a brief description of the dynamics around the unstable equilibrium points and the role the invariant manifolds play in the transfer of matter around the galaxy. In Sect.~\\ref{sec:res}, we show the different morphologies that result from the computations. ",
        "conclusions": ""
    },
    "0801/0801.4232_arXiv.txt": {
        "abstract": "\\noindent We provide a new construction of the modes of the Poincar\\'e dodecahedral space $S^3/I^*$.  The construction uses the Hopf map, Maxwell's multipole vectors and orbifolds.  In particular, the *235-orbifold serves as a parameter space for the modes of $S^3/I^*$, shedding new light on the geometrical significance of the dimension of each space of $k$-modes, as well as on the modes themselves. \\\\  \\noindent Keywords:  Poincar\\'e dodecahedral space, spherical 3-manifold, eigenmodes of the Laplace operator, Hopf fibration, multipole vectors, orbifold ",
        "introduction": "\\label{SectionIntroduction}  Cosmological motivations \\cite{LaLu} have inspired recent progress in understanding the eigenmodes of the spherical spaces $S^3/\\Gamma^*$, \\ie, the quotients of the three-sphere $S^3$ by a binary polyhedral group $\\Gamma^*$. Such modes may be seen as the $\\Gamma ^*$-invariant solutions of the Helmoltz equation in the universal cover $S^3$. Their numeration and degeneracy were given by Ikeda \\cite{Ikeda}. Recent works \\cite{Caillerie,Gundermann,mlrEigenmodes} have provided  various means to calculate them.  Here we give a new point of view, using the Hopf map, multipole vectors and orbifolds to construct the modes of $S^3/\\Gamma^*$ and shed additional light on the geometrical significance of Ikeda's formula. Section~\\ref{SectionHopf} reviews the Hopf map and uses it to lift eigenmodes from $S^2$ to $S^3$. Section~\\ref{SectionEigenmodes} uses twist operators to extend the lifted modes to a full eigenbasis for $S^3$. Section~\\ref{SectionModesOfSphericalSpaces} generalizes the preceding results from the modes of $S^3$ to the modes of a spherical space $S^3/\\Gs$, showing that the latter all come from the lifts of the those eigenmodes of $S^2$ that are invariant under the corresponding (non-binary) polyhedral group $\\Gamma$. We then turn to a detailed study of the $\\Gamma$-invariant modes of $S^2$. Section~\\ref{SectionS2Modes} recalls Maxwell's multipole vector approach and uses it to associate each mode of $S^2/\\Gamma$ to a $\\Gamma$-invariant set of multipole directions.  Restricting attention to the case that $\\Gamma$ is the icosahedral group, Section~\\ref{SectionPDSModes} introduces the concept of an orbifold and re-interprets a $\\Gamma$-invariant set of multipole directions as a (much smaller) set of points in the *235-orbifold, which serves as the parameter space. Section~\\ref{SectionConclusion} pulls together the results of the preceding sections to summarize the construction of the modes of the Poincar\\'e dodecahedral space and state the dimension of the mode space for each $k$. ",
        "conclusions": "\\label{SectionConclusion}  Returning to the 3-dimensional Poincar\\'e dodecahedral space $S^3/I^*$, the results of the preceding sections may be summarized as follows.  Keep in mind that $S^3/I^*$ admits $k$-modes for even $k$ only;  odd $k$-modes cannot exist because $I^*$ contains the antipodal map. \\\\  \\noindent{\\bf Theorem~\\ref{SectionConclusion}.1}  To construct the modes of the Poincar\\'e dodecahedral space $S^3/I^*$, \\begin{itemize} \\item Each mode of $S^3/I^*$ corresponds to an $I^*$-invariant mode of $S^3$ (elementary). \\item Each $I^*$-invariant mode of $S^3$ is a sum of twists of $I^*$-invariant vertical modes of $S^3$ (Proposition~\\ref{SectionVerticalModesSuffice}.1). \\item Each $I^*$-invariant vertical $k$-mode of $S^3$ is the pull-back, under the Hopf map, of an $I$-invariant $\\ell$-mode of $S^2$, with $k = 2\\ell$ (Proposition~\\ref{SectionLiftingModes}.1). \\item The $I$-invariant $\\ell$-modes of $S^2$ are parameterized by $\\ell/60$ points on the *235-orbifold, possibly including fractional points (Section~\\ref{SectionOrbifold}). \\end{itemize}  \\noindent{\\bf Theorem \\ref{SectionConclusion}.2} The space of $k$-modes of the Poincar\\'e dodecahedral space $S^3/I^*$ has dimension \\begin{equation} (k + 1) ( 1 + \\lfloor\\frac{k/2}{2}\\rfloor + \\lfloor\\frac{k/2}{3}\\rfloor + \\lfloor\\frac{k/2}{5}\\rfloor - \\frac{k}{2}). \\end{equation} \\\\  \\noindent{\\it Proof.}  The space of $I$-invariant $k/2$-modes of the 2-sphere has dimension $ 1 + \\lfloor\\frac{k/2}{2}\\rfloor + \\lfloor\\frac{k/2}{3}\\rfloor + \\lfloor\\frac{k/2}{5}\\rfloor - \\frac{k}{2}$ (Proposition~\\ref{SectionImprovedDimensionFormula}.1) and thus the space of vertical $I^*$-invariant $k$-modes of the 3-sphere has this same dimension (Theorem~\\ref{SectionConclusion}.1).  The twist operators then take each vertical mode to a $(k+1)$-dimensional space of generic $I$-invariant modes (Table~\\ref{modeTable} and Proposition~\\ref{SectionVerticalModesSuffice}.1,), completing the proof. $\\blacksquare$\\\\"
    },
    "0801/0801.0642_arXiv.txt": {
        "abstract": "We cross-correlate large scale structure (LSS) observations from a number of surveys with cosmic microwave background (CMB) anisotropies from the Wilkinson Microwave Anisotropy Probe (WMAP) to investigate the Integrated Sachs-Wolfe (ISW) effect as a function of redshift, covering $z \\sim 0.1 - 2.5$. Our main goal is to go beyond reporting detections towards developing a reliable likelihood analysis that allows one to determine cosmological constraints from ISW observations. With this in mind we spend a considerable amount of effort in determining the redshift-dependent bias and redshift distribution ($b(z)\\times dN/dz$) of these samples by matching with spectroscopic observations where available, and analyzing auto-power spectra and cross-power spectra between the samples. Due to wide redshift distributions of some of the data sets we do not assume a constant bias model, in contrast to previous work on this subject. We only use the LSS data sets for which we can extract such information reliably and as a result the data sets we use are 2-Micron All Sky Survey (2MASS) samples, Sloan Digital Sky Survey (SDSS) photometric Luminous Red Galaxies, SDSS photometric quasars and NRAO VLA Sky Survey (NVSS) radio sources. We make a joint analysis of all samples constructing a full covariance matrix, which we subsequently use for cosmological parameter fitting. We report a 3.7$\\sigma$ detection of ISW combining all the datasets. We do not find significant evidence for an ISW signal at $z>1$, in agreement with theoretical expectation in $\\Lambda$CDM model. We combine the ISW likelihood function with weak lensing of CMB (hereafter Paper II \\cite{paperII}) and CMB power spectrum to constrain the equation of state of dark energy and the curvature of the Universe. While ISW does not significantly improve the constraints in the simplest 6-parameter flat $\\Lambda$CDM model, it improves constraints on 7-parameter models with curvature by a factor of 3.2 (relative to WMAP alone) to $\\Omega_K=-0.004^{+0.014}_{-0.020}$, and with dark energy equation of state by 15\\% to $w=-1.01^{+0.30}_{-0.40}$ [posterior median with ``$1\\sigma$'' (16th--84th percentile) range]. A software package for calculating the ISW likelihood function can be downloaded at {\\tt http://www.astro.princeton.edu/\\~{}shirley/ISW\\_WL.html}. ",
        "introduction": "Introduction}  The Cosmic Microwave Background (CMB) has provided us with a wealth of cosmological information. The large-scale anisotropies were first discovered by the Differential Microwave Radiometer (DMR) on Cosmic Background Explorer (COBE) satellite \\cite{smoot92}, and the smaller-scale CMB anisotropies were subsequently measured by various ground-based/balloon-borne experiments. More recently, the Wilkinson Microwave Anisotropy Probe (WMAP) satellite \\cite{bennett03, jarosik07} produced a cosmic variance limited map of CMB anisotropies down to $l\\sim 400$. The structure of the angular power spectrum when combined with other cosmological probes (such as Sloan Digital Sky Survey, \\cite{tegmark06}, Hubble Key Project \\cite{freedman94} and 2dF Galaxy Redshift Survey \\cite{cole05}), allows extremely precise measurements of the cosmological parameters of the $\\Lambda$CDM model. While most of the fluctuations seen by WMAP and other CMB experiments were generated at the last surface of scattering, structures formed at low redshift also leave imprints on the CMB.  These anisotropies, such as the thermal Sunyaev-Zeldovich (tSZ) \\cite{sunyaev80a} and kinetic Sunyaev Zeldovich effects (kSZ) \\cite{sunyaev80b}, the Integrated Sachs-Wolfe (ISW) effect \\cite{sachs67}, and gravitational lensing, contribute only slightly to the CMB power spectrum on scales measured by WMAP, but they can be detected by cross-correlating the CMB with suitable tracers of the large scale structure.  This is the first of two papers that measure the Integrated Sachs-Wolfe effect and gravitational lensing (Paper II) in cross-correlation. In this paper, we focus on large scale galaxy-temperature correlations and their large scale cosmological source, the Integrated Sachs-Wolfe (ISW) effect.  The ISW effect results from the red- or blue-shifting of the CMB photons as they propagate through gravitational potential wells. As the potential wells of the Universe (i.e., the spatial metric) evolve, the energy gained by photons falling into the potential well does not cancel out the energy loss as photons climb out of the well. This is important at late times when the Universe is not matter dominated and the gravitational potential is time dependent. It is only significant on large scales, since on small scales the amount of time spent by the photon in each coherence region of the gravitational potential is small and any small scale fluctuations will be smoothed out as the photon go through numerous potential wells along the line of sight.  To measure the above effect, we cross-correlate the CMB temperature anisotropies with maps of galaxies from the Two Micron All Sky Survey (2MASS), luminous red galaxies (LRGs) and quasars from the Sloan Digital Sky Survey, and radio sources from the NRAO VLA Sky Survey (NVSS). This incorporates most of the LSS tracers used by previous efforts \\cite{boughn98,fosalba03,scranton03,afshordi04,boughn04,fosalba04,nolta04,padmanabhan05ISW,gaztanaga05,cabre06,giannantonio06,vielva06, pietrobon06,mcewen07,rassat07} to detect the ISW effect. Our goal in this work extends this previous literature by going beyond detecting the ISW effect to measuring its redshift evolution and using that to constrain different cosmological models (e.g. the ISW effect due to spatial curvature occurs at significantly higher redshifts than that due to a cosmological constant). We therefore require a large redshift range ($z \\sim 0$ to $2.5$) but with sufficient redshift resolution to unambiguously discern any redshift evolution of the signal. In addition, to draw robust cosmological conclusions from an observed redshift evolution, we must constrain both the redshift distribution and evolution of the bias with redshift for each of the samples; the simple assumption of constant bias is in most cases no longer sufficient. These considerations drive our survey selections; we discuss these in more detail in Sec.~\\ref{sec:discussion}. Our final product is a likelihood code that can be applied to any cosmological model. In addition to providing complementary constraints on standard cosmological parameters, we expect it can be a strong discriminator of the modified gravity models, which have very distinctive ISW predictions \\cite{song07}.      We review the theory behind the ISW effect in Sec.~\\ref{sec:theory}. The CMB and LSS data sets used are described in Sec.~\\ref{sec:data}; the results of cross-correlating the two are in Sec.~\\ref{sec:cross}. Sec.~\\ref{sec:dndz} and ~\\ref{sec:sys} constrain the redshift distributions of the samples, and possible systematic contamination of the cross-correlations. Sec.~\\ref{sec:cosmo} presents the cosmological implications of these results, and Sec.~\\ref{sec:discussion} summarizes our conclusions. The companion paper (Paper II) uses the same data sets to detect the  weak lensing of the CMB. All of the theoretical predictions are made with WMAP 3 year parameters ($\\Omega_b h^2$=$0.0223$, $\\Omega_c h^2$ = $0.128$, $\\Omega_K=0$, $h=0.732$, $\\sigma_8 =0.761$) except in Section~\\ref{sec:dndz} or otherwise stated. ",
        "conclusions": "\\label{sec:discussion}  The main goal of this paper is to perform a full analysis of the integrated Sachs-Wolfe effect using the cross-correlations between WMAP CMB maps and maps of large scale structure. In contrast to previous work on this subject we place less emphasis on establishing a detection of ISW and more emphasis on developing a tool with which cosmological models can be compared to the data in a close to optimal fashion. For this reason we only select the data sets that can be reliably used towards this goal, as discussed in more detail below. The redshift range of the datasets we use is between 0 and 2.5. We use optimal weighting of the data both in angular space and in redshift space to extract the maximum amount of information, taking into account properly the correlations between them. Our final product is the likelihood function with which different cosmological models can be compared to each other.  As the ISW effect is both a probe of cosmological parameters and a consistency test of the standard $\\Lambda$CDM cosmology, there have been significant previous efforts made to detect it.  A number of different analysis methods have been used and the WMAP data have been cross-correlated with several samples.  These include the 2MASS XSC; several SDSS samples including magnitude-limited galaxy samples, LRGs, and quasars; the NVSS; and the HEAO hard X-ray map.  Most of these samples (or samples with similar spatial coverage and redshift range) are included in the present work, but not all. Here we compare our analysis with the previous work and comment on the reasons for our choice of data sets.  \\newcounter{oisw} \\begin{list}{\\arabic{oisw}. }{\\usecounter{oisw}} \\item {\\em Near-infrared galaxies (2MASS).}  The 2MASS galaxies are useful for ISW due to high sky coverage and the ability to see closer to the Galactic plane in the near-IR than in the optical.  However they can only probe the lowest redshifts ($z<0.2$). Afshordi et~al. \\cite{afshordi04} and Rassat et~al. \\cite{rassat07} have measured the ISW signal using the 2MASS sample and we deliberately cut our 2MASS sample into brightness bins such as theirs so that we can compare the results. We find that our measured signal from 2MASS is very similar. We do however derive cosmological constraints using a Markov chain (which fits all the cosmological parameters instead of just $\\Omega_\\Lambda$) from these samples. We also take into account (albeit in a crude way) the redshift dependence of the bias resulting from seeing all nearby galaxies but only the brightest and most biased galaxies at $z\\ge 0.1$.  \\item {\\em Optical galaxies (SDSS, APM).} Wide-angle multicolor galaxy surveys such as SDSS open almost limitless possibilities for constructing galaxy samples, and many of these samples have been used in previous ISW work.  Most work so far has been on either flux-limited samples \\cite{fosalba03, cabre06}, which have a broad redshift distribution, or photometric LRGs \\cite{fosalba03, scranton03, padmanabhan05ISW, cabre06}, which can be seen to larger distances and for which it is easier to construct reliable photo-$z$ cuts.  In SDSS, photometric LRG samples oversample the linear density field in the redshift range $0.2<z<0.6$ and the lower redshifts are covered by 2MASS, so the flux-limited galaxy samples would be redundant in terms of volume for our study; we therefore did not include them.  Our LRG samples cover the largest solid angle to date of any SDSS ISW analysis (6641 deg$^2$) and for the purposes of cosmological analysis are split into two photo-$z$ slices. Fosalba \\& Gazta\\~naga \\cite{fosalba04} have also used galaxies from the Automated Plate Measuring (APM) survey \\cite{maddox90}, which adds $\\sim 4300$ deg$^2$ in the Southern Hemisphere inaccessible to SDSS.  Their APM sample has a typical redshift $\\bar z\\approx 0.15$ and thus would add some information beyond the most distant of our 2MASS samples.  Considering that APM area is only 16\\% of 2MASS and that it only marginally extends the redshift range we have not used APM in our analysis. However adding a deeper galaxy survey in the South, comparable to or deeper than SDSS, would be valuable for improving ISW constraints. Overall signal to noise from SDSS LRG galaxies is about 3$\\sigma$, most of which comes from the higher redshift sample around $z \\sim 0.5$.  \\item {\\em Optical quasars (SDSS).}  Photometrically selected quasars can probe large-scale structure at much higher redshifts than ``normal'' galaxies because they are bright enough to be seen in wide-angle surveys (such as SDSS) even at $z\\sim 2$. The only ISW analysis with quasars so far has been that of Giannantonio et~al. \\cite{giannantonio06}, who cross-correlated WMAP with a sample of photometric quasars from the SDSS.  Our analysis uses similar selection criteria, but we have used photo-$z$ cuts to eliminate most of the lower-redshift objects, and used a combination of spectroscopic data and angular clustering to constrain $b*dN/dz$ taking into account the multimodal nature of the photo-$z$ failures.  We also slice our quasars into two photo-$z$ bins. Despite these improvements we find that the significance is only 1.3$\\sigma$ (1.24$\\sigma$ for the high redshift sample with $z>1$), and we therefore do not confirm that the 2.1$\\sigma$ signal seen in \\cite{giannantonio06} comes from $z>1$.  \\item {\\em Radio sources (NVSS).}  There have been several past WMAP$\\times$NVSS ISW analyses \\cite{boughn04, nolta04, pietrobon06, mcewen07}, taking advantage of the high redshift (compared to most optical samples) and wide sky coverage of the NVSS. We have used the angular power spectrum whereas the previous works have used correlation functions or wavelet coefficients. However, the most important difference between our analysis and the previous result is that we fit $b*dN/dz$ from cross-correlations rather than using the Dunlop \\& Peacock model \\cite{dunlop90} for the redshift distribution and assuming constant bias.  This is important as we find the fit $b*dN/dz$ looks very different (see Fig.~\\ref{fig:fnvss}).  All of these studies, including ours, have found positive cross correlations at the $\\sim 3\\sigma$ level. However, the interpretation of this result depends sensitively on one's ability to measure $b*dN/dz$ and this is where we believe our analysis is an improvement upon previous efforts.  \\item {\\em Hard X-ray background (HEAO).}  Boughn \\& Crittenden \\cite{boughn04} have used the HEAO hard X-ray map \\cite{boldt87} for ISW cross-correlation.  The background is due mainly to unresolved (by HEAO) active galactic nuclei and hence traces large scale structure at redshifts of order unity.  This, combined with the all-sky nature of HEAO, is beneficial for ISW projects. However, we decided not to add in HEAO sample to our analysis for several reasons. First is the difficulty in understanding the $b(z)*dN/dz$ of the sample (we use the general notation $dN/dz$ here even though for unresolved X-ray flux it would be more accurate to write $dF/dz$). Only $\\sim 75$\\% of the background is resolved by Chandra into sources with measured redshifts \\cite{cowie03, boughn04b}, and we have little guidance on where to place the other 25\\%.  Even if we knew $dN/dz$ perfectly, this does not tell us $b*dN/dz$: the modeled $dN/dz$ spans the range $0<z<3$ and it is unlikely that the bias would be even approximately constant over this range. An alternative is to fit for their bias and redshift distributions up to high $z$ using a cross-correlation method similar to that done for NVSS in Sec.~\\ref{sec:dndz}. Unfortunately HEAO has FWHM of $\\sim 2^\\circ$ and does not resolve individual sources, so we would have to fit the data to the model without small-scale information, which loses signal-to-noise on the cross-correlation very rapidly. A secondary reason is that there is considerable overlap between HEAO and NVSS, so it is likely that the two trace partly the same structure, and thus the improvement in ISW constraint is not as large as adding two independent data sets. We note that it may make sense to include the hard X-ray maps in parameter estimation in the future if a robust determination of $b*dN/dz$ becomes available. \\end{list}  In summary, we believe we used most of the available large scale structure data useful for ISW analysis. This not only updates previous work on ISW effect \\cite{boughn98,fosalba03,scranton03,afshordi04,boughn04,fosalba04,nolta04,padmanabhan05ISW,gaztanaga05,cabre06,giannantonio06,vielva06, pietrobon06,mcewen07,rassat07}, but is also the first one that attempts to do the tomography of ISW, in the sense of encompassing a wide redshift range via our mass tracers going from the local Universe to $z\\sim 2.5$, while reducing the amount of overlap in area and redshift as much as possible. We have argued that many of the previous measurements have a considerable overlap in redshift and area, which means that they cannot be combined independently and that the effective redshift of the sample is not necessarily the redshift from where most of the ISW signal is coming from. Our analysis, while attempting to minimize the overlap in the first place, takes the residual correlations into account explicitly via the construction of the full covariance matrix.  We note that Giannantonio et~al. \\footnote{T. Giannantonio et~al., in preparation.} are also pursuing an ISW tomography analysis, with somewhat different choices of LSS samples and cross-correlation methodologies.  We spend a significant fraction of our analysis obtaining the correct redshift distributions for all of the samples. To be more accurate, it is the $b*dN/dz$ that we constrain for all samples. The signature of ISW effect is highly affected by the redshift distribution of the tracer, and thus one would need to have an accurate idea of what the redshift distribution is in order to interpret the correlation. Apart from employing spectroscopic datasets that overlap in magnitude range and sky coverage, we correlate the tracer samples with one another so as to obtain the $b*dN/dz$ for some of the samples. This is mainly possible because LRGs have relatively good photometric redshifts and so we correlate the LRGs with other overlapping datasets to determine what are the $b*dN/dz$ at the redshift range that LRGs cover. In addition, we account for redshift-dependent bias in 2MASS and for the multimodal error distributions for the quasars. We also made the first determination of $b*dN/dz$ for NVSS sample which is not based simply on a theoretical model fitting the luminosity function.  Correlations of mass tracers with the CMB sky can be caused not only by the ISW effect, but also by other cosmological effect such as thermal SZ, Galactic foregrounds and extinction, and extragalactic point sources. We provide an estimate for all these effects and only include the scales deemed reliable, where the contamination is subdominant or negligible.  We report a detection of $3.7\\sigma$ of the ISW effect combining 2MASS, SDSS, and NVSS with WMAP data. We make a joint analysis of all samples by constructing a reliable covariance matrix including cross-correlations of different samples, which is needed for cosmological parameter fitting. We combine our ISW correlation functions with weak lensing of the CMB (Paper II) to derive cosmological constraints on three different cosmological models: (i) the ``vanilla'' $\\Lambda$CDM model, (ii) $\\Lambda$CDM+$\\Omega_K$, and (iii) $\\Lambda$CDM+$w$. We find a slight improvement of our measurement of $w$ in model (iii) over the measurement made by CMB alone: $w=-1.01^{+0.30}_{-0.40}$ instead of $-1.03^{+0.46}_{-0.43}$. The constraining power of our analysis is however most prominent in determining that curvature of the Universe: for CMB+ISW+WL we find $\\Omega_K=-0.004^{+0.014}_{-0.020}$ instead of $-0.050^{+0.044}_{-0.065}$ for CMB alone. These constraints are not as tight as that obtained by some other methods, such as combining the CMB with baryonic oscillations or with supernovae \\cite{eisenstein05,seljak06,spergel07}, but it is subject to very different systematics.  It is thus reassuring that all of them are consistent with each other. Even more importantly, there are other models where ISW can be crucial in distinguishing them from standard $\\Lambda$CDM, such as $f(R)$ models in which the growth of structure is not fixed by the background geometry \\cite{song07}. Some of these models may already be inconsistent with our ISW signal; we plan to present such constraints in a future paper. These constraints should improve further in the future with deeper galaxy surveys that should reach the cosmic variance limit out to $z\\sim 1-2$, and future CMB data that enables lower-noise lensing reconstruction.  Finally, we would like to note that we plan to release a package for calculating ISW likelihood function given the datasets and cosmological parameters. This will be described further in the documentation for the package."
    },
    "0801/0801.2192_arXiv.txt": {
        "abstract": "The non-thermal 3.6 cm radio continuum emission from the young stars S1 and DoAr21 in the core of Ophiuchus, has been observed with the Very Long Baseline Array (VLBA) at 6 and 7 epochs, respectively, between June 2005 and August 2006. The typical separation between successive observations was 2 to 3 months. Thanks to the remarkably accurate astrometry delivered by the VLBA, the trajectory described by both stars on the plane of the sky could be traced very precisely, and modeled as the superposition of their trigonometric parallax and a uniform proper motion. The best fits yield distances to S1 and DoAr21 of 116.9$^{+7.2}_{-6.4}$ pc and 121.9$^{+5.8}_{-5.3}$ pc, respectively.  Combining these results, we estimate the mean distance to the Ophiuchus core to be 120.0$^{+4.5}_{-4.2}$ pc, a value consistent with several recent indirect determinations, but with a significantly improved accuracy of 4\\%. Both S1 and DoAr21 happen to be members of tight binary systems, but our observations are not frequent enough to properly derive the corresponding orbital parameters. This could be done with additional data, however, and would result in a significantly improved accuracy on the distance determination. ",
        "introduction": "Ophiuchus is one of the most active regions of star-formation within a few hundred parsecs of the Sun (e.g.\\ Lada \\& Lada 2003). It has played an important role in the development of our understanding of star-formation, and remains an important benchmark for this field of research. Indeed, it has been one of the key targets of the Spitzer c2d legacy program (Padgett et al.\\ 2007); and has been observed in detail at numerous other wavelengths, including X-rays (Ozawa et al.\\ 2005, Gagn\\'e et al.\\ 2004), near-infrared (e.g.\\ Haisch et al.\\ 2002, Duch\\^ene et al.\\ 2004, and references therein), sub-millimeter (Motte et al.\\ 1998, Johnstone et al.\\ 2004), and radio (e.g.\\ Andr\\'e et al.\\ 1987, Leous et al.\\ 1991).  The detailed analysis of this wealth of observational data has been somewhat hampered by the relatively large uncertainty on the distance to the Ophiuchus complex. Traditionally assumed to be at 165 pc (Chini 1981), it has recently been suggested to be somewhat closer. For example, de Geus et al.\\ (1989) found a mean photometric distance of 125 $\\pm$ 25 pc. Knude \\& Hog (1998), who examined the reddening of stars in the direction of Ophiuchus as a function of their Hipparcos distances, also found a clear extinction jump at 120 pc. Using a similar method, M.\\ Lombardi et al.\\ (in prep.) also find a distance of about 120 pc for the Ophiuchus core. Finally, Mamajek (2007) identified reflection nebulae within 5$^\\circ$ of the center of Ophiuchus, and obtained the trigonometric parallax of the illuminating stars from the Hipparcos catalog. From the average of these Hipparcos parallaxes, he obtains a mean distance to Ophiuchus of 135 $\\pm$ 8 pc.  This latter result is based on parallax measurements, but considers a fairly large area around Ophiuchus. It could, therefore, include objects unrelated to Ophiuchus itself. The former results are restricted to regions more concentrated on Ophiuchus, but they are based on indirect distance determinations. Here, we will present measurements of the trigonometric parallax of two young stars (S1 and DoAr21) directly associated with the Ophiuchus core. This will allow us to estimate directly the distance to this important region of star-formation. ",
        "conclusions": "\\subsection{Properties of S1}  The mean 3.6 cm flux of S1 in our data is 4.8 mJy, and the dispersion about that mean is 1.2 mJy (see Fig.\\ 1). This shows that S1 is variable at the level of about 25\\% on timescales of months to years. This modest level of variability is certainly not unexpected for a non-thermal source associated with an active stellar magnetosphere (Feigelson \\& Montmerle 1999). As mentioned earlier, Andr\\'e et al.\\ (1991) reported a VLBI detection of S1 at 6 cm. They found --among many other things-- that the source was somewhat resolved in their observations, with a full width at half maximum extension of about 1.7 mas. The radio emission associated with S1 is also found to be resolved in {\\em all} six of our observations, with a deconvolved mean full width at half maximum of about 0.95 mas.  This is somewhat smaller than the figure reported by Andr\\'e et al.\\ (1991), but we note (i) that our observations and those of Andr\\'e et al.\\ (1991) were obtained at different wavelengths; and (ii) that at some of our epochs, the size of the emission reached 1.5 mas, whereas at other epochs, it was smaller than 0.5 mas. At the distance of S1 (see below), 0.95 mas corresponds to about 24 \\Rsun. The diameter of S1 is expected to be about 8.5 \\Rsun\\ (Andr\\'e et al. 1991), so its magnetosphere appears to be on average 3 times more extended than its photosphere.  The fact that S1 is resolved, and that its size varies from epoch to epoch likely produces small random shifts in the photocenter of the radio emission with an amplitude of a fraction of the size of the emitting region. The true uncertainties on the position of S1 are, therefore, likely to be somewhat larger than the figures quoted in Tab.\\ 1. Another factor that must be taken into account is that S1 is known to be a member of a binary system with a separation of about 20 mas (Richichi et al.\\ 1994). The companion is inferred to be about 4 times dimmer than the primary at K band, so it is likely to be significantly less massive (Richichi et al.\\ 1994). If we assume S1 to be a 6 \\Msun\\ star (as suggested by its B4 spectral type), we expect the orbital period to be about 0.7 yr, and the reflex motion of S1 to be about 1 to 2 mas if the companion is 10 to 20 times less massive than S1. Thus, the amplitude of the reflex motion is expected to be larger than the formal errors on the positions of S1 listed in Tab.\\ 1.  \\subsection{Properties of DoAr21}  The total radio flux of DoAr21 has long been known to be highly variable (Feigelson \\& Montmerle 1985). Our observations certainly confirm this strong variability since the ratio between the highest and the lowest measured flux exceeds 50 (Fig.\\ 1). In particular, the flux during our first two observations (10--20 mJy) is systematically about an order of magnitude higher than that (0.4--2 mJy) at any of the following 5 observations. Unfortunately, our time coverage is too coarse to decide whether these first two epochs correspond to two different flares, or to a single long-duration one.  The extreme variability of DoAr21, while at odds with the situation in S1, is reminiscent of the case of the spectroscopic binary V773 Tau (e.g.\\ Massi et al.\\ 2002). In the latter source, Massi et al.\\ (2002) showed that the variability had the same periodicity as the orbital motion, with the radio flux being highest at periastron. Interestingly, DoAr21 was found to be double during our second observation\\footnote{The position given in Tab.\\ 1 is that of the brightest of the two components. The other source is offset by more than 5 mas from its position of the steady component expected from the astrometry fits presented in \\S 4.3.}. This suggests that the same mechanism that enhances the radio emission when the two binary components are nearest, might be at work in both objects. The separation between the two components of DoAr21 in our second observation is about 5 mas.  This value, of course, corresponds to the projected separation; the actual distance between them must be somewhat larger. Moreover, if the mechanisms at work in DoAr21 and V773 Tau are similar, then DoAr21 must have been near periastron during our second epoch, and the orbit must be somewhat eccentric. As a consequence of these two effects, the semi-major axis of the orbit is likely to be a few times larger than the measured separation between the components at our second epoch, perhaps 10 to 15 mas. At the distance of DoAr21, this corresponds to 1.2 to 1.8 AU. For a mass of 2.2 \\Msun\\ (see Sect.\\ 1), the corresponding orbital period is 0.4 to 1.3 yr, and one would expect the source to oscillate with this kind of periodicity.  \\subsection{Astrometry}  The absolute positions of S1 and DoAr21 (listed in columns 3 and 4 of Tab.\\ 1) were determined using a 2D Gaussian fitting procedure (task JMFIT in AIPS). This task provides an estimate of the position errors (also given in columns 3 and 4 of Tab.\\ 1) based on the expected theoretical astrometric precision of an interferometer (Condon 1997). Systematic errors, however, usually limit the actual precision of VLBI astrometry to several times this theoretical value (e.g.\\ Pradel et al.\\ 2006, Loinard et al.\\ 2007).  Moreover, we have just seen that the extended magnetosphere of S1, and the reflex motions of both S1 and DoAr21 are likely to produce significant shifts in the positions of the source photocenters. While the effect of an extended magnetosphere might be to produce a random jitter, the reflex orbital motions ought to generate oscillations with a periodicity equal to that of the orbital motions.  Our observations, however, are currently insufficient to properly fit full Keplerian orbits. Instead, in the present paper, we represent the possible systematic calibration errors as well as the jitter due to extended magnetospheres and the oscillations due to reflex motions, by a constant error term (the value of which will be determined below) that we add quadratically to the errors given in Tab.\\ 1. The displacements of both S1 and DoAr21 on the celestial sphere are then modeled as a combination of their trigonometric parallaxes ($\\pi$) and their proper motions ($\\mu_\\alpha$ and $\\mu_\\delta$), assumed to be uniform and linear. The astrometric parameters were determined using a least-square fit based on a Singular Value Decomposition (SVD) scheme (see Loinard et al.\\ 2007 for details). The reference epoch was taken at the mean of each set of observations (JD 2453757.63 $\\equiv$ J2006.061 for S1, and JD 2453796.52 $\\equiv$ J2006.167 for DoAr21). The best fit for S1 (Fig.\\ 2a) yields the following astrometric parameters:  \\begin{eqnarray} \\alpha_{J2006.061} & = & \\mbox{ \\dechms{16}{26}{34}{174127} } ~ \\pm ~ \\mbox{ \\mmsec{0}{000026} } \\nonumber \\\\% \\delta_{J2006.061} & = & \\mbox{ \\decdms{-24}{23}{28}{44498} } ~ \\pm ~ \\mbox{ \\msec{0}{00028} } \\nonumber \\\\% \\mu_\\alpha \\cos \\delta & = & -3.88 ~ \\pm ~ 0.87 ~ \\mbox{mas yr$^{-1}$} \\nonumber \\\\% \\mu_\\delta & = & -31.55 ~ \\pm ~ 0.69 ~ \\mbox{mas yr$^{-1}$} \\nonumber \\\\% \\pi & = & 8.55 ~ \\pm ~ 0.50 ~ \\mbox{mas.} \\nonumber \\end{eqnarray}  \\noindent For DoAr21, on the other and, we get (Fig.\\ 2b):  \\begin{eqnarray} \\alpha_{J2006.167} & = & \\mbox{ \\dechms{16}{26}{03}{018535} } ~ \\pm ~ \\mbox{ \\mmsec{0}{000020} } \\nonumber \\\\% \\delta_{J2006.167} & = & \\mbox{ \\decdms{-24}{23}{36}{35830} } ~ \\pm ~ \\mbox{ \\msec{0}{00022} } \\nonumber \\\\% \\mu_\\alpha \\cos \\delta & = & -26.47 ~ \\pm ~ 0.92 ~ \\mbox{mas yr$^{-1}$} \\nonumber \\\\% \\mu_\\delta & = & -28.23 ~ \\pm ~ 0.73 ~ \\mbox{mas yr$^{-1}$} \\nonumber \\\\% \\pi & = & 8.20 ~ \\pm ~ 0.37 ~ \\mbox{mas.} \\nonumber \\end{eqnarray}   \\noindent To obtain a reduced $\\chi^2$ of 1 in both right ascension and declination, one must add quadratically 0.062 ms of time, and 0.67 mas to the statistical errors of S1 listed in Tab.\\ 1, and 0.053 ms of time, and 0.57 mas to the statistical errors of DoAr21. These figures include all the unmodeled sources of positional shifts mentioned earlier. Interestingly, the residuals of the fit to the S1 data (inset in Fig.\\ 2a) are not random, but seem to show a $\\sim$ 0.7 yr periodicity, as expected from the reflex motions (Sect.\\ 4.1). Similarly, the residuals from the fit to DoAr21 seem to show a periodicity of $\\sim$ 1.2 yr (Fig.\\ 2b, inset), within the range of expected orbital periods of that system (Sect.\\ 4.2). This suggests that the errors are largely dominated by the unmodeled binarity of both sources, and that additional observations designed to provide a better characterization of the orbits ought to improve significantly the precision on the trigonometric parallax determinations.  The distance to S1 deduced from the parallax calculated above is 116.9$^{+7.2}_{-6.4}$, while the distance deduced for DoAr21 is 121.9$^{+5.8}_{-5.3}$.  The weighted mean of these two parallaxes is 8.33 $\\pm$ 0.30, corresponding to a distance of 120.0$^{+4.5}_{-4.3}$. Since both S1 and DoAr21 are {\\it bone fide} members of the Ophiuchus core, this figure must represent a good estimate of the distance to this important region of star-formation. Note that it is in good agreement with several recent determinations (e.g.\\ de Geus et al.\\ 1989, Knude \\& Hog 1998, Lombardi et al.\\ ibid), but with a significantly improved relative error of 4\\%. This level of accuracy is likely to be further improved once additional observations of S1 and DoAr21 designed to characterize their orbital motions are available. Such observations are currently being collected at the VLBA. A significant improvement in the distance estimate will also be obtained once the parallax to other sources (also currently observed at the VLBA) are measured."
    },
    "0801/0801.2371.txt": {
        "abstract": "%{ that's for old abstracts % context heading (optional) {L-type ultra-cool dwarfs and brown dwarfs have cloudy atmospheres that could host weather-like phenomena.  The detection of photometric or spectral variability would provide insight into unresolved atmospheric {heterogeneities}, such as holes in a global cloud deck. Indeed, a number of ultra-cool dwarfs have indeed been reported to vary. Additional time-resolved spectral observations of brown dwarfs offer the opportunity to further constrain and characterize atmospheric variability. } % aims heading (mandatory) {It has been proposed that growth of {heterogeneities} in the global cloud deck may account for the L- to T-type transition as brown dwarf photospheres evolve from cloudy to clear conditions.  Such a mechanism is compatible with variability. We searched for variability in the spectra of five L6 to T6 brown dwarfs in order to test this hypothesis. } % methods heading (mandatory) {We obtained spectroscopic time series using the near-infrared spectrographs {ISAAC} on VLT--ANTU, over 0.99--1.13\\,$\\mu$m, and {SpeX} on the Infrared Telescope Facility for two of our targets, in $J, H$ and $K$ bands. %, as part of our {CLOUDS} survey for atmospheric variability of brown dwarfs. %for two brown dwarfs of type L6 and T2 in JHK using {\\sc SpeX} on IRTF. {We search for statistically variable lines and correlation between those.}\\, } % results heading (mandatory) {High spectral-frequency variations are seen in some objects, but these detections are marginal and need to be confirmed. We find no evidence for large amplitude variations in spectral morphology and we place firm upper limits of {2 to 3}\\,\\% on broad-band variability, depending on the targets and wavelengths, on the time scale of a few hours. In contrast to the rest of the sample, the T2 transition brown dwarf {SDSS\\,J1254$-$0122} shows numerous variable features, but a secure variability diagnosis would require further observations.} % conclusions heading (optional), leave it empty if necessary {Assuming that any variability arises from the rotation of patterns of large-scale clear and cloudy regions across the surface, we find that the typical physical scale of cloud cover disruption should be smaller than 5--8\\,\\% of the disk {area} for four of our targets, using simplistic heterogeneous atmospheric models. {The possible variations seen in  {SDSS\\,J1254$-$0122} are not strong enough to allow us to confirm the cloud breaking hypothesis.}\\, } ",
        "introduction": "L/T transition brown dwarfs are an informal class of sub-stellar objects comprising the latest L-type and the earliest T-type dwarfs (roughly {L8 to T4}). % , see e.g. Burgasser \\cite{Bur07}). Their near-infrared spectra may exhibit both CO and CH$_4$ absorption and they have $J-K$ colours intermediate between those of the red late-type L dwarfs and the blue mid-T dwarfs (e.g. $0<J-K<1.2$). The spectra of L~dwarfs are characterized by absorption due to metal hydrides (FeH, CrH) and {strong} alkali (K, Na, Rb, Cs) lines, absorption bands of water and CO, and the grey opacity of a cloud deck (Kirkpatrick et~al.~\\cite{Kir99}, Mart\\'\\i n et~al.~\\cite{Mar99}). These clouds are made of condensates of refractory elements such as Al, Ca, Fe, Mg, Si, as well as {Ti and V} whose oxides dominate the optical spectral morphology of M dwarfs, but which are depleted from cooler atmospheres as they {condense} into dust grains. The vertical distribution of the clouds is a complex issue, but  a balance between sedimentation (the equivalent of rain in the case of water) and upward mixing likely controls their vertical profile (e.g., Ackerman \\& Marley~\\cite{Ack01}). As the atmosphere becomes cooler in late-type L~dwarfs, the grey clouds gradually settle below the photosphere and the emergent spectrum %, with a strong wavelength dependence. becomes more strongly affected by molecular band opacities. In the next cooler spectral class, T, CO {completes its reduction} to CH$_4$ (Burgasser et~al.~\\cite{Bur02a}; Geballe et~al.~\\cite{Geb02}). The latter produces strong IR absorption, along with water and molecular hydrogen (collision-induced absorption). Synthetic spectra predicted by atmosphere models with a discrete cloud deck reproduce well the colours (Ackerman \\& Marley \\cite{Ack01}, Burgasser et~al.~\\cite{Bur02b}), the near-infrared spectra (Cushing et~al.~\\cite{Cus05}), and the mid-infrared spectra (Roellig et~al.~\\cite{Roe04}) of the L dwarfs. Likewise models with no grain opacity {fit well} the spectra of mid- to late-type T~dwarfs (Cushing et~al.~\\cite{Cus08}).  However the rapidity of the transition between the cloudy L~dwarf atmospheres and the clear T~dwarf atmospheres is not well understood.  Proposed mechanisms have included horizontal fragmentation of the global cloud deck (Burgasser et~al.~\\cite{Bur02b}), a rapid increase in the cloud particle sedimentation efficiency (Knapp et~al.~\\cite{Kna04}), or the sinking of a very thin cloud layer (Tsuji et~al.~\\cite{Tsu04}). The case for a dynamic mechanism, as opposed to a gradual sinking of a uniform cloud deck, was strengthened by recent measurements of geometric parallaxes for a significant number of brown dwarfs (Vrba et~al.~\\cite{Vrb04}), and the observations of brown dwarfs over a larger wavelength range of their spectra, particularly in the mid-infrared (Golimowski et~al.~\\cite{Gol04}). Effective temperatures derived independently of atmospheric modeling assumptions reveal that the L/T transition takes place over a small effective temperature range (about 200\\,K), {as was qualitatively first noticed by Kirkpatrick et~al.~(\\cite{Kir99}). This} is difficult to reconcile with a gradual settling of the clouds. It is also marked by a brightening in the $J$~band with increasing spectral type ({Dahn et~al.~\\cite{Dah02};} Knapp et~al.~\\cite{Kna04}, their Fig.~8 and 9). Burgasser et~al.~(\\cite{Bur02b}) interpreted this rapid brightening, and the corresponding $J-K$ colour change, as an indication of the fragmentation of the global cloud cover %and surface heterogeneities for {in the L/T transition brown dwarfs' atmospheres.} %This rapid brightening, originally reported for a smaller number of brown dwarfs, and the corresponding $J-K$ colour change, are interpreted by Burgasser et~al.~(\\cite{Bur02b}) as an indication of cloud cover breaking and heterogeneous surface of the transition brown dwarfs.  One possible way to test the horizontal fragmentation hypothesis would be to search for variability over a wide range of spectral types. Large-scale, patchy clouds, like those on Jupiter, would be expected to create photometric variations (Gelino \\& Marley~\\cite{Gel00}) as the structures rotate across the disk. If the signature of such phenomena is preferentially observed in the L/T transition objects this would strengthen the case for cloud fragmentation. If the typical length scale of horizontal cloud patches are small, however, little to no photometric or spectral variation would be expected. {This would also be the case for stable stripes along the brow dwarf parallels.} Thus detection of variability supports the cloud fragmentation hypothesis, but non-detection cannot rule it out. {Here we do not put forward any hypotheses on how such heterogeneities could arise.  We aim simply to  test if the surface is heterogeneous on large spatial scales. }  {Other explanations for the L/T transition enigma have been suggested.} Since we started this project, a number of L/T brown dwarfs have been resolved into binaries composed by a late-L {or early-T} primary, and a fainter T companion (e.g. {Reid et~al.~\\cite{Rei01}}, Burgasser et~al.~\\cite{Bur05}, Liu et~al.~\\cite{Liu06}). {Although the statistics are still not conclusive}, the seemingly higher rate of binaries among L/T~transition brown dwarfs compared to other brown dwarfs has prompted {Burgasser et~al.~(\\cite{Bur05,Bur06b}) and Liu et~al.~\\cite{Liu06})} to speculate that binarity could explain the peculiar relation {between effective temperature and spectral energy distribution} observed over the transition. {In the case where most L/T transition objects would be composed of earlier- and later-type components, we would not expect to find more variability among transition objects than among the populations from which the components are drawn, taking into account the binary effects. Similarly, our data do not allow to search confirmation of any mechanism that would not produce surface heterogeneities, including uniform and gradual sinking, or thinning (e.g., Knapp et~al.~\\cite{Kna04}, Tsuji et~al.~\\cite{Tsu04}) of the clouds.}  There have been a number of reports of brown dwarfs exhibiting variability. Tinney \\& Tolley~(\\cite{Tin99}), Bailer-Jones \\& Mundt~(\\cite{CBJ99,CBJ01}), Bailer-Jones \\& Lamm~(\\cite{CBJ03}), Clarke et~al.~(\\cite{Cla02a}) and Koen~(\\cite{Koe03}) studied samples of early to mid L~dwarfs in the $I$~band and reported variability detections  for a fraction of their targets. Gelino et~al.~(\\cite{Gel02}) conducted optical and near-infrared photometry follow-up of a dozen of brown dwarfs{, and found several variations in their targets' light curves}. {Enoch et~al.~(\\cite{Eno03}) conducted a $K_s$-only monitoring program of nine L and T~dwarfs, including three L/T transition dwarfs. At the 2-$\\sigma$ level, they found three variable dwarfs, including one T1 dwarf with a possible 2.96-h period, at the 89\\,\\% confidence level (C.L.), and did not  find a correlation between variability and spectral type. %At 99\\,\\% confidence level (C.L.), they find a variable L8 dwarf, but without a field star to be used as check. Another field star is also found variable at the same confidence level. Their average detection limit of sinusoidal variations is 12\\,\\% (99\\,\\% C.L.). } {Clarke et~al.~(\\cite{Cla02b}) reported {that Kelu-1 (L3, {Geballe et~al.~\\cite{Geb02}}) varies} photometrically in their 2000 optical observations, with a 1.8\\,h period, but this was not observed again in 2002 (Clarke et~al.~\\cite{Cla03}).} {Bailer-Jones \\& Lamm~(\\cite{CBJ03}) studied the effect of second-order extinction residuals in photometric variability observations. These residuals are due to the variable, wavelength-dependent {telluric} absorption of H$_2$O, CO$_2$ and O$_3$ and to the very different colours of the brown dwarfs and the comparison stars. They estimate these residuals to be of the percent level, and cannot generally be removed. Our spectroscopic observations are mostly not affected by this effect. Only deep, non-resolved lines in both the target spectrum {and the telluric absorption spectrum} could be affected.}  H$\\alpha$ has been monitored in two L3 dwarfs (Hall~\\cite{Hal02}, Clarke et~al.~\\cite{Cla03}) and an L5 dwarf (Liebert et~al.~\\cite{Lie03}), and the emission was found variable in all three objects. {Burgasser et~al.~(\\cite{Bur02c}) studied the peculiar, H$\\alpha$ emitting T6.5 brown dwarf 2MASS~J12373919+6526148 using optical spectroscopy and $J$-band photometry. A {tentative} shift in the H$\\alpha$ line was detected but no photometric variability, at the $\\approx 2.5$\\,\\% level. A series of eleven 20-min spectra of the L8 dwarf Gl~584C revealed ``subtle'' qualitative variations in the 9250--9600\\,\\AA\\  water absorption band, although a telluric origin cannot be excluded, and redward of the CrH 8600\\,\\AA\\  bandheaad (Kirkpatrick et~al.~\\cite{Kir01}). } Near-infrared spectroscopic data are more scarce. Nakajima et~al.~(\\cite{Nak00}) cautiously reported a variation in the water band at 1.53--1.58\\,$\\mu$m in the spectra of the T6 {SDSS\\,J1624+0029}, using the Subaru telescope, although over only 80\\,min of time. {The work most similar to this has recently been published by Bailer-Jones\\,(\\cite{CBJ08}). Two out of four ultracool dwarfs have been found to show correlated spectral numerous variations with 13-nm resolution, often (but not always) associated with one or more known features, particularly water and methane. }  Finally, Morales-Calder{\\'o}n et~al.~(\\cite{Mor06})  searched for photometric variability in three late L dwarfs  (DENIS-P~J0255$-$4700, 2MASS~J0908+5032, and 2MASS~J2244+2043) with the IRAC photometer on {\\em Spitzer Space Telescope}. Two out of the three objects studied exhibited some slight modulation in their light curves at 4.5, but not $8.8\\,\\rm \\mu m$.  In short, despite numerous variability searches with a great variety of instruments and technique, there has yet been no definitive detection of luminosity variability in a {field} brown dwarf. Here we report on a search for spectroscopic variability in a selection of L/T transition objects by the {\\sc Clouds} collaboration (Continuous Longitude Observations of Ultra-cool Dwarfs, Goldman et~al.~\\cite{Gol03})\\footnote{\\tt http://astronomy.nmsu.edu/CLOUDS/}. In the following we present our observations (Section\\,\\ref{obs}), and their reduction and how we search for variability (Section\\,\\ref{red}). In Section\\,\\ref{anal} we detail for each target the variability information we obtained from our time series, and present variability upper-limits. Finally in Section\\,\\ref{discussion} we interpret our results and link them to atmospheric physical processes.  %__________________________________________________________________ ",
        "conclusions": "We have obtained high SNR, low resolution spectroscopic time series over 0.99--1.13\\,$\\mu$m for five brown dwarfs of spectral types L6 to T6, sampling the L/T transition, and for a sub-set of two brown dwarfs of types L6 and T2 in the $J, H$ and $K$ bands. We cautiously report some high-frequency {(narrow-band)} variations, which we generally cannot tie to specific absorption lines. {An exception is FeH in the spectra of three targets.} {SDSS\\,J1254$-$0122} shows numerous variable features. All require confirmation by additional observations.  We place constraints on broad-band spectroscopic variability at the levels of 2\\,\\% ({2MASS\\,J1534$-$2952AB}) to 3\\,\\% ({2MASS\\,J0825+2115}, {2MASS\\,J1225$-$2739AB} and SDSS\\,J1624+0029), {with a 99\\,\\% confidence level.}  When comparing to the variations expected from crude atmospheric model interpolations between cloudy, L-type and clear, T-type atmospheres, we find that over the course of our observations, no significant variations (larger than $\\sim$5--8\\,\\% of the brown dwarf's disk) in the cloud coverage occurred in our sample. {Our data do not rule out smaller-size heterogeneity on the brown dwarf surface, nor does our analysis try to constrain other variability mechanisms.}"
    },
    "0801/0801.0704_arXiv.txt": {
        "abstract": "Neither standard model SEDs nor truncated standard model SEDs fit observed spectra of QU Carinae with acceptable accuracy over the range 900\\AA~to 3000\\AA. Non-standard model SEDs fit the observation set accurately. The non-standard accretion disk models have a hot region extending from the white dwarf to $R=1.36R_{\\rm wd}$, a narrow intermediate temperature annulus, and an isothermal remainder to the tidal cutoff boundary. The models include a range of $\\dot{M}$ values between $1.0{\\times}10^{-7}M_{\\odot}~{\\rm yr}^{-1}$ and $1.0{\\times}10^{-6}M_{\\odot}~{\\rm yr}^{-1}$ and limiting values of $M_{\\rm wd}$ between $0.6M_{\\odot}$ and $1.2M_{\\odot}$. A solution with $M_{\\rm wd}=1.2M_{\\odot}$ is consistent with an empirical mass-period relation. The set of models agree on a limited range of possible isothermal region $T_{\\rm eff}$ values between 14,000K and 18,000K. The model-to-model residuals are so similar that it is not possible to choose a best model. The Hipparcos distance, 610 pc, is representative of the model results. The orbital inclination is between $40\\arcdeg$ and $60\\arcdeg$. ",
        "introduction": "Cataclysmic variables (CVs) are semi-detached binary stars in which a late-type main sequence star loses mass onto a white dwarf (WD) by Roche lobe overflow \\citep{w1995}. A sub-class, the NL systems (discussed below) are of interest because they are photometrically stable. It is believed that an analytic expression describes the radial temperature profile on the accretion disk, leading to a corresponding spectral energy distribution (SED). Tests using stellar SEDs \\citep{wade1988} produced unsatisfactory results, but a theoretical model tailored to accretion disks \\citep{hubeny1990} is available and might be expected to lead to satisfactory results. In this paper we show that synthetic SEDs, built on the theoretically-prescribed model, do not fit the observed data, but non-standard empirical models fit the observations closely.  In non-magnetic systems the mass transfer stream produces an accretion disk with mass transport inward and angular momentum transport outward by viscous processes. The accretion disk may extend inward to the WD; the outer boundary extends to a tidal cutoff limit imposed by the secondary star in the steady state case. If the mass transfer rate is below a certain limit, the accretion disk is unstable and undergoes brightness cycles (outbursts), and if above the limit, the accretion disk is stable against outbursts \\citep{s1986,c1993,w1995,o1996}. The latter case objects are called nova-like (NL) systems; their accretion disks are nearly fully ionized to their outer (tidal cutoff) boundary and this condition suppresses dwarf nova outbursts. NL systems are of special interest because they are expected to have an accretion disk radial temperature profile\tgiven by an analytic expression \\citep[eq.5.41]{fr92}(hereafterFKR) thereby defining the so-called standard model.  Based on reports of QU Carinae as an emission-line object \\citep{st1968}, \\citet{gil1982} (hereafter GP) undertook a spectroscopic study to determine orbital characteristics of the system. They determined an orbital period of 10.9 hours. No secondary star spectrum was detected, indicating that the spectrum is dominated by light from an accretion disk or the primary star. Balmer line emission is relatively weak. This evidence in turn suggested a high rate of mass transfer and a correspondingly bright system. It also suggests that QU Car is a NL-type system. We discuss known parameters of QU Car in more detail in \\S4.  High speed photometry confirmed erratic variations of 0.1-0.2 mag. on time scales of minutes, as reported by \\citet{schild1969}. \\citet{kn1994} used $IUE$ spectra to study time variability of the UV resonance lines. They found that there are short-time-scale variations with a possible correlation with orbital period, but the data are insufficient to confirm the orbital nature of the variations.  \\citet{har2002}\tobtained $Hubble~Space~Telescope$ $STIS$ spectra of QU Car at three epochs. Blueshifted absorption in N V and C IV provides evidence of an outflow with a maximum velocity of $\\approx 2000$ km~${\\rm s}^{-1}$. Further analysis of these spectra is in \\citet{drew2003}. The latter authors suggest that the QU Car secondary is a carbon star, and that QU Car is the most luminous NL system known, possibly at a distance of ${\\approx}2000$ pc.  Our objective in this study is to use synthetic spectra, based on a physical model of QU Carinae, to constrain properties of this interesting system. ",
        "conclusions": "We have appropriated the term \"boundary layer`` to designate the region between the WD and the accretion disk $R=1.3611R_{\\rm wd}$. This region more generally is associated with a very hot region occupying the same space but predicted to emit hard X-rays if optically thin or soft X-rays if optically thick \\citep[sect. 2.5.4]{w1995}. This prediction is modified if a wind carries off part of the accretion mass \\citep{HD1991,HD1993}. In the latter case, predicted boundary layer $T_{\\rm eff}$'s are no higher than the values adopted in our models.\tBased both on its geometric location and the much higher local temperature than in the rest of the accretion disk, our appropriated term actually is in keeping with the usual sense of the term.  Adoption of the intermediate temperature layer obviously is a crude approximation. A temperature discontinuity at its inner and outer boundary, as in our present models, is not physically credible. However, this model is a much better fit to observation than with our earlier experimental (\\S6) models, which adopted a sharp, nonetheless more gradual temperature drop to the isothermal disk region than in this paper.  The outermost accretion disk region contibutes little to the system synthetic spectrum. The $T_{\\rm eff}$ values, for a limited region which would vary from one $\\dot{M}$ to another, could be closer to standard model values than the isothermal assumption without producing a detectable change in the system synthetic spectrum. On the other hand it is known that temperatures in the outer region of an accretion disk need an upward correction from the standard model to account for impact heating \\citep{las2001} and tidal heating \\citep{bm2001}.  The accretion disk temperature profiles found in this study differ from profiles found for other NL systems. In both SDSSJ0809 \\citep{linnell2007a}  and IX Vel \\citep{linnell2007b} the final model included an inner isothermal region with an outer region that followed a standard model. In the case of MV Lyr \\citep{l2005} in a high state, the final model had a standard model accretion disk extending from an inner truncation radius to an intermediate radius with an isothermal radius beyond. It is of interest that all of the NL systems included in these studies have accretion disks whose temperature profiles clearly differ from the standard model."
    },
    "0801/0801.1228_arXiv.txt": {
        "abstract": "The ultraviolet spectra of all \"weak emission line central stars of planetary nebulae\" (WELS) with available IUE data is analyzed. We found that the WELS can be divided in three different groups regarding their UV: (1) Strong P-Cygni profiles (mainly in C IV 1549); (2) Weak P-Cygni features and (3) Absence of P-Cygni profiles. We have measured wind terminal velocities for all objects presenting P-Cygni profiles in N V 1238 and/or C IV 1549. The results obtained were compared to the UV data of the two prototype stars of the [WC]-PG 1159 class, namely, A30 and A78. They indicate that WELS are distinct from the [WC]-PG 1159 stars, in contrast to previous claims in the literature. In order to gain a better understanding about the WELS, we clearly need to determine their physical parameters and chemical abundances. First non LTE expanding atmosphere models (using the CMFGEN code) for the UV and optical spectra of the star Hen 2-12 are presented. ",
        "introduction": "Hydrogen deficient central stars of planetary nebulae (CSPN) are generally divided in three main groups\\footnote{Here, we neglect other type of hydrogen deficient objects such as R Coronae Borealis (RCB) and O(He) stars}: [WR], PG 1159 and [WC]-PG 1159 stars. However, besides these classes, there is the \"weak emission line stars\" (WELS or [WELS]) group, which was introduced in the extensive observational study presented by Tylenda et al. (1993).  Although a considerably advance in the understanding of the origin and evolution of [WR], [WC]-PG 1159 and PG 1159 stars have been achieved in the last decade (see e.g. Werner \\& Herwig 2006), the evolutionary status of the WELS remains an open question. Motivated by the fact that only a few studies have been done focusing these objects, and that most of them were done in the optical part of the spectrum, we have investigated the UV spectra of all WELS with available IUE data. Moreover, we have applied the non LTE expanding atmosphere code CMFGEN to model the UV and optical spectra of one WELS, namely, Hen 2-12. ",
        "conclusions": "We have investigated the UV spectra of all WELS with available IUE data. We have found that they can be divided in three different groups: (1) Strong P-Cygni profiles (mainly in C IV 1549); (2) Weak P-Cygni features and (3) Absence of P-Cygni profiles. We have derived terminal velocities for all objects presenting P-Cygni profiles in N V 1238 and/or C IV 1549 in a homogeneous way. Both the $v_\\infty$ measurements and the UV characteristics observed indicate that the WELS are not [WC]-PG 1159 stars.  We have presented first non LTE expanding atmosphere models for the star Hen 2-12. Optical line fits are presented for the first time. Although interesting results could be obtained, in order to make an efficient comparison to the [WR] and PG 1159 classes, and to gain insight to the evolutionary status of the WELS, we clearly need an analysis of a large sample."
    },
    "0801/0801.1472_arXiv.txt": {
        "abstract": "We report on our attempts to achieve a nearly steady-state gas flow in hydrodynamical simulations of doubly barred galaxies. After exploring the parameter space, we construct two models, for which we evaluate the photometric and the kinematic integrals, present in the Tremaine-Weinberg method, in search of observational signatures of two rotating patterns. We show that such signatures are often present, but a direct fit to data points is likely to return incorrect pattern speeds. However, for a particular distribution of the tracer, presented here, the values of the pattern speeds can be retrieved reliably even with the direct fit. ",
        "introduction": "Bars within bars are common in disc galaxies -- up to 30\\% of early-type barred galaxies contain them (Erwin \\& Sparke 2002). It is generally expected that they play a role in the feeding of active galactic nuclei, though the direct theoretical confirmation is difficult, as the interaction between the bars considerably narrows the range of possible systems (Maciejewski \\& Sparke 2000, hereafter MS00; Maciejewski \\& Athanassoula 2008). Corsini, Debattista, \\& Aguerri (2003) demonstrated that in NGC 2903 the two bars rotate with two different pattern speeds. Maciejewski (2006) and Merrifield, Rand, \\& Meidt (2006) proposed extensions to the Tremaine-Weinberg method (Tremaine \\& Weinberg 1984) in order to derive multiple pattern speeds. There were also attempts to derive them from direct fitting of straight lines to the Tremaine-Weinberg (TW) integrals. Here we analyse in detail the behaviour of the TW integrals in hydrodynamical models of gas flow in double bars. Since only one dynamically plausible model of double bars has been constructed so far (MS00), we approach the problem from another direction here: by changing the parameters of the two bars, we search for steady-state flows, which may indicate preferred systems.  \\begin{figure}[t] \\plotone{singh1.ps} \\vspace{-137mm} \\caption{\\small Snapshots of gas density in two models of gas flow in double bars: Setup 1 at 955 Myr (left panel) and at 980 Myr (central panel); Setup 2 at 1445 Myr (right panel). Darker shading indicates higher density. Both bars rotate counterclockwise and the outer bar is vertical on the plots. Units on the axes are in kpc.} \\end{figure} ",
        "conclusions": "Correct pattern speeds in Setup 2 are recovered only at the late stages of evolution, because by then the majority of gas is accumulated in the region dominated by the inner bar. Earlier in the run, the slits close to the line of nodes sample both flows dominated by the inner and the outer bar. At the end of the run, the contribution from the outer bar is reduced, so that {\\it in the slits close to the line of nodes most of the tracer follows the inner bar.} This is the only gas distribution for which two pattern speeds can be derived with good confidence from a direct fit of straight lines to the TW integrals. Otherwise, the points indicating the TW integrals may gather on two lines, but the slope of the second line may show no relation to the pattern speed of the inner bar. Moreover, for many systems with two pattern speeds, points representing the TW integrals may not gather on any line, as in Setup 1. We conclude that if there is more than one rotating pattern in a galaxy, the direct fits to the TW integrals may provide reliable pattern speeds only for a very specific distribution of the tracer. Fully comprehensive work on this subject will be possible once we know the range of systems with two pattern speeds that are dynamically possible."
    },
    "0801/0801.2298_arXiv.txt": {
        "abstract": "Analyzing global starburst properties in various kinds of starburst and post-starburst galaxies and relating them to the properties of the star cluster populations they form, I explore the conditions for the formation of massive, compact, long-lived star clusters. The aim is to find out whether the relative amount of star formation that goes into star cluster formation as opposed to field star formation, and into the formation of massive long-lived clusters in particular, is universal or scales with star formation rate, burst strength, star formation efficiency, galaxy or gas mass, and whether or not there are special conditions or some threshold for the formation of star clusters that merit to be called globular clusters a few gigayears later. ",
        "introduction": "Star Cluster (SC) formation is a major or even dominant mode of all star formation (SF) and occurs in very different environments. This immediately raises the question whether young star clusters (YSCs) forming in different environments are similar or different.  SCs are not only interesting in their own right, but bear considerable power as tracers of SF in their parent galaxies. YSCs trace the spatial distribution of SF and its recent history within a galaxy, old globular clusters (GCs) trace violent SF phases in their parent galaxy -- all the way back to the very onset of SF, i.e. over a Hubble time. But SCs also fade and dissolve. In actively star forming galaxies, the youngest SCs may still be embedded in their natal dust clouds while part of the older SCs are already gone, dissolved and/or faded below detection. It is therefore important to take these processes into account when comparing SC populations in different galaxies. They depend on the initial properties of individual SCs, their masses, radii, abundances, stellar IMF, and of the SC population, i.e. the luminosity function, the mass function, distribution of radii, ages, etc.  That SC formation is an important mode of SF in starbursts was shown on the example of the Tadpole and Mice interacting galaxies. A pixel-by-pixel analysis of ACS data ($BVRI$) with GALEV evolutionary synthesis models showed that $\\sim 70$ \\% of the blue light is emitted by YSCs as opposed to only $\\sim 30$ \\% coming from field stars. We estimated that more than 35 \\% of all SF went into the formation of YSCs, not only in the main bodies of these two galaxies, but all along their extended tidal tails \\citep{deGrijs+03a}. Clearly, this analysis needs to be extended to different types of galaxies, starburst/non-starburst, dwarf/normal, gas-rich/gas-poor, interacting/non-interacting, in various stages of the interaction, etc. to explore the systematics. There are indications from \\cite*{Meurer+95}'s work that the contribution from YSCs to the UV light of a galaxy increases with increasing UV surface brightness, which itself is a measure of SF intensity. A burning question that we are currently exploring is whether the amount of SF that goes into massive long-lived SCs in relation to the amount of SF that goes into low-mass short-lived SCs and field stars, increases with increasing overall SFR, with bursts strength, or with star formation efficiency (SFE). If so, the immediate next question is whether this is a continuous increase or rather a threshold effect in the sense that only above a certain SF intensity or efficiency the massive long-lived SCs are formed that later are called GCs.  SCs form in very different environments, in normal star-forming galaxies, as e.g. M51 or NGC 5236 (cf. \\cite*{Larsen04}, \\cite*{Mora+07}, in dwarf galaxy starbursts like NGC 1569 \\citep{Anders+04b}, in interacting gas-rich galaxies like NGC 4038/39 (the Antennae) \\citep{Whitmore+95,Whitmore+05,Anders+07}. Slightly older and intermediate-age SCs are observed in post-starburst merger remnants like NGC 7252 \\citep{Whitmore+93,FB95,Miller+97,SchweizerSeitzer98} and dynamically young ellipticals like NGC 1316 \\citep{Goudfrooij+01b,Goudfrooij+04,Goudfrooij+07}, respectively. They can form all over the main body of a galaxy, as e.g. in the Antennae or NGC 7252, in and around a starburst nucleus, as e.g. in Arp 220 or NGC 6240 \\citep{Shioya+01,Pasquali+03}, all along some -- but not all -- extended tidal features \\citep{Knierman+03,deGrijs+03a,Trancho+07a}, as well as in group environments like Stephan's quintett \\citep{Gallagher+01}. These environments cover a huge range in terms of density, kinetic temperature, chemical abundances and it is by no means obvious whether or not all these SCs are similar or different, individually or as a population. Related questions are where, when, and how GC are formed and what a young GC looks like. Or how to tell apart YSCs into long-lived and short-lived ones -- by mass, concentration, mass function, ...?  Current cluster formation models require exceptionally high SF efficiencies ${\\rm SFE~:=~M_{\\ast}/M_{gas}~>30~\\%}$ as a prerequisite for the formation of massive strongly bound and long-term stable SCs, i.e. for the formation of young GCs \\citep{Brown+95,Burkert+96,ElmegreenEfremov97,Li+04}. On a global scale, SF efficiencies in normal spiral and irregular galaxies, as well as in starbursting dwarf galaxies are of order $0.1 - 3 \\% $ \\citep{Krueger+95}. On the smaller scale of molecular clouds in the Milky Way, the SF efficiency is of the same order of magnitude and so is the mass ratio between the molecular cloud core and the entire molecular cloud. No GC formation is therefore expected in spirals, irregulars or star-bursting dwarf galaxies by today. In giant gas-rich interacting galaxies, on the other hand, SF efficiencies of order $10 - 50 \\%$ are reported on global scales, and of order $30 - 90 \\%$ on nuclear scales of a few hundred pc up to $\\sim 1$ kpc. In those systems, GC formation should be possible. The fact that different submm lines (CO(1-0), HCN(1-0), CS(1-0)) trace molecular gas at different densities (${\\rm n \\geq 100, n~\\geq 3 \\cdot 10^4, n~\\sim 10^5~cm^{-3}}$), has allowed to see that while in the Milky Way and other nearby galaxies only a small fraction ($0.1 - 3 \\%$) of the mass of a molecular cloud makes up its high density core, the situation is drastically different in Ultraluminous Infrared galaxies (ULIRGs), which all are late stages of massive gas-rich mergers with strong nuclear (few 100 pc) starbursts. In those ULIRGs, almost all the molecular gas in the central starburst region is at the high densities of molecular cloud {\\em cores}, indicating that the molecular cloud structure must be very different from what we know in our Galaxy. The entire nuclear region is just one supergiant molecular cloud core, seriously raising the question whether the star and SC formation processes can be the same as in normal galaxies, not to mention the situation in extended, expanding, low-density tidal structures in the outskirts of other interacting galaxies.  In any case, before ALMA becomes operational, the YSCs forming in these different types of environments are our best proxy to the molecular cloud structure. In the Milky Way, molecular cloud cores, molecular clouds, and YSCs all feature power-law mass functions, suggesting scale-free self-similar evolution. Not even for the closest massive merger, the Antennae, can we presently determine the molecular cloud or cloud core mass functions (cf. \\cite*{Wilson+03}). The masses of YSCs and the shape of their mass function is all we can access \\citep{Wilson+06,Anders+07}. \\cite*{GaoSolomon04} and \\cite*{Solomon+92} have shown that for all galaxies -- from Blue Compact Dwarfs to spirals and ULIRGs -- there is a tight correlation between SFR, as derived from far-infrared luminosity, and the mass in molecular cloud cores, as derived from the HCN luminosity. They also find the SF efficiency to be proportional to the mass ratio of molecular gas at core and normal densities, i.e. to the ratio between HCN or CS luminosity and CO luminosity. The highest density molecular gas in all these environments is transformed into stars with almost 100 \\% efficiency and it is the amount of gas at those high densities that controls SF. The fraction of molecular gas at the highest densities therefore defines the SF efficiency.  The high ambient pressure building up in the course of massive gas-rich mergers can drive up SF efficiencies by $1-2$ orders of magnitude by compressing molecular clouds, increasing their masses and, in particular, their core mass fractions. \\cite*{JogDas92,JogDas96} have shown that the ISM pressure during mergers can easily become $3-4$ times higher than the typical internal molecular cloud pressure, raising the SF efficiency to $70-90 \\%$. This leads us to expect that the relative amount of SF that goes into the formation of massive, strongly bound young GCs in relation to the amount of SF that goes into field stars and low-mass, short-lived clusters is enhanced in massive gas-rich mergers. ",
        "conclusions": "{\\bf We so far conclude that starbursts in non-interacting dwarf galaxies do not form substantial populations of massive, long-lived YSCs that could evolve into GC, while massive gas-rich mergers do.}  NGC 7252 is not the only example of a merger that no doubt has produced a new generation of GCs. The $\\sim 1 - 3$ Gyr old merger remnants NGC 3921 \\citep{Schweizer+96}, NGC 34 \\citep{SchweizerSeitzer07}, and NGC 1316 \\citep{Goudfrooij+01a,Goudfrooij+01b,Goudfrooij+04,Goudfrooij+07} as well feature young GC populations formed during the mergers. The metallicities of these newly formed GCs agree well with expectations on the basis of spiral galaxy ISM properties. Evolutionary synthesis models predict these SCs to take on the optical colours of the red-peak GC widely observed in E/S0 galaxies by the time the tidal features indicative of the merger origin will have vanished. They also predict that they should readily be detectable against other populations of red GCs in combined optical and NIR observations \\citep{Fritze04} (see also R. Kotulla, this volume).  No example of a clearly merging gas-rich dwarf galaxy pair, nor of an accretion of a gas-rich dwarf by an elliptical or S0 has as yet been studied to check whether those would also give rise to new GC populations."
    },
    "0801/0801.0853_arXiv.txt": {
        "abstract": "We discuss a cosmology in which  cold dark matter begins to decay into relativistic particles at a recent epoch ($z < 1$).  We show that the large entropy production and associated bulk viscosity from such decays  leads to an accelerating cosmology as required by observations.  We investigate the effects of decaying cold dark matter in a $\\Lambda = 0$, flat, initially matter dominated  cosmology.  We show that this model satisfies the cosmological constraint from the redshift-distance relation for type Ia supernovae. The age in such models is also consistent with the constraints from the oldest stars and globular clusters. Possible candidates for this late decaying dark matter are suggested along with  additional observational tests of this cosmological  paradigm. ",
        "introduction": "Understanding the nature and  origin of both the dark energy  \\cite{garnavich} and cold dark matter \\cite{Feng06} constitutes a significant challenge to modern cosmology.  The simplest particle physics explanation for the cold dark matter is, perhaps, that of the lightest supersymmetric particle, an axion, or a heavy (e.g. \"sterile\") neutrino.   The dark energy, on the other hand is generally attributed to a cosmological constant, or possibly a vacuum energy in the form of a \"quintessence\" scalar field \\cite{wetterich,ratra} or $k$-essence \\cite{zlatev,steinhardt99}  which must be slowly evolving along an effective potential.  See \\cite{Barger06} for a review. In addition to these explanations, however, the simple coincidence that  both of these unknown entities currently contribute  comparable mass energy toward the closure of the universe begs the question as to whether they could be different manifestations of the same physical phenomenon.  Indeed, suggestions along this line have been made by many \\cite{Fabris,Chaplygin,Abramo04,Umezu}. See \\cite{Bean05} for  a recent review.  In Ref.~\\cite{Wilson07} yet another mechanism was considered by which a dark-matter particle could produce the cosmic acceleration.  In that work it was shown  that  entropy production and an associated bulk viscosity could result from  a decaying dark-matter particle.  Moreover,  that bulk viscosity would act as a negative pressure similar to a cosmological constant or quintessence.  However, in that paper it was shown the decay alone was not sufficient to produce the observed cosmic acceleration.  It was proposed, however,  \\cite{Wilson07} that some, but not all of the desired cosmic acceleration could be accounted for if the particle decay was delayed by proceeding thorough a cascade of long lived intermediate states before the final entropy-producing decay.  In this paper we expand on the hypothesis that the dark energy could be produced from a delayed decaying dark-matter particle.   Here, we show that a dark-matter particle which, though initially stable, begins to decay to relativistic particles near the time of the present  epoch will produce a cosmology consistent with the observed cosmic acceleration deduced from the type Ia supernova distance-redshift relation without  the need for a cosmological constant.  Hence, this  paradigm has the possibility to  account for the apparent dark energy without the well known fine tuning and smallness problems \\cite{Bean05} associated with a cosmological constant.  Also, for the model proposed herein, the apparent acceleration is a temporary phenomenon.  This avoids the difficulties \\cite{hellerman} in accommodating a cosmological constant in string theory.  The idea of delayed dark matter decay is not new.  It was previously introduced \\cite{Turner} as a means to provide an $\\Omega_M = 0.1-0.3$ without curvature ( $\\Omega_{tot} = 1$) by  hiding  matter in weakly interacting relativistic particles.  Here, we point out  that such a cosmology not only allows for a flat cosmology with low apparent cold dark matter matter content, but can produce  an accelerating cosmology consistent with observations.  We show that the bulk viscosity produced during the decay will briefly accelerate the cosmic expansion as matter is being converted from nonrelativistic to relativistic particles. This model thus shifts the dilemma in modern cosmology from that of explaining dark energy to one of explaining how an otherwise stable heavy particle might begin to decay at a late epoch.  In the next section we summarize the cosmology of late decaying dark matter and its associated bulk viscosity.   In the following section we discuss  candidate particles for such decays, and in Section IV we present fits to the  supernova  magnitude-redshift relation which  show that these data can be reasonably well fit  in a flat $k=0$, $\\Lambda = 0$ cosmology  with recent dark-matter decay.  We  summarize the constraints that the supernova data  places on the properties of the decay along with independent constraint from the ages of oldest stars and globular clusters. ",
        "conclusions": "We have considered models in which the present  cosmic acceleration derives from the temporary insertion of dissipative mass energy due to the bulk viscosity created by the recent decay of a cold dark-matter particle into  light (undetectable) relativistic species.  As  illustrative examples we have considered initially matter dominated flat ($\\Lambda=0$) cosmologies plus late-time particle decay. We find that models with bulk viscosity from late-time dark-matter decay are  consistent with observations of the supernova magnitude-redshift relation, and ages from the oldest stars and globular clusters.  We argue that it will likely satisfy other constraints as well.   Moreover, there is a difference in the SNIa magnitude-redshift relation for this cosmology compared to the standard $\\Lambda$CDM model.  This is because the deceleration is faster at high redshift due to a higher matter content during the matter-dominated epoch.  Thus,  as more data are accumulated at the highest redshifts,   it may be possible to distinguish  between this cosmology and a standard $\\Lambda$CDM model.  For now, however, there is sufficient success in the present model to motivate further work.  In a subsequent paper we plan  to consider  higher order terms in the transport equation in detail as well as the effects of such late decays on the observed power spectrum of the cosmic microwave background \\cite{WMAP}, and the growth of large scale structure.  Ultimately, of course, one must decide whether the dilemma of a cosmological constant with all of its difficulties is less palatable  than the dilemma of a bulk viscosity produced by the delayed onset of dark-matter decay. For now, however, our  purpose has simply been to establish that such a possibility exists and that it warrants further investigation."
    },
    "0801/0801.2584_arXiv.txt": {
        "abstract": "With the Navy Prototype Optical Interferometer (NPOI), the binary system $\\theta^{1}$ Orionis C, the most massive member of the Trapezium, was spatially resolved over a time period extending from February 2006 to March 2007. The data show significant orbital motion over the 14 months, and, after combining the NPOI data with previous measurements of the system from the literature, the observations span 10 years of the orbit. Our results indicate that the secondary did not experience an unusually close periastron passage this year, in contradiction to the prediction of a recently published, highly eccentric $\\sim$11 year orbit. Future observations of this source will be required to improve the orbital solution. Possible implications of the results in terms of system distance are discussed, although a main conclusion of this work is that a definitive orbit solution will require more time to obtain sufficient phase coverage, and that the interaction effects expected at periastron did not occur in 2007. ",
        "introduction": "Orion is the nearest example of a giant molecular cloud and is the site of both high-mass star formation and a prodigious number of recently-formed stars; the central 2.5 pc region of the Orion Nebula Cluster (ONC) alone contains $\\sim$3500 stars with a combined mass exceeding 900 $M_{\\odot}$ \\citep{hil97}. Although star-forming regions such as Taurus are closer, these regions are dark clouds associated with low-mass star formation and far fewer total stars. Initial star-count studies \\citep{lad91} implied that up to 95\\% of stars formed in clustered environments like Orion, while more recent {\\it Spitzer Space Telescope} results indicate that at least half of all stars originate in dense regions \\citep{meg04}. An analysis of binary star distributions \\citep{pat02} suggests that approximately 70\\% of field stars may have formed in a clustered environment. Thus, the study of Orion, the closest giant star forming cluster, is central to investigating the early history of the majority of stars. Furthermore, the variety of cluster morphologies investigated with {\\it Spitzer} observations suggest that OB stars play a significant role in the formation and evolution of star formation in clusters \\citep{meg04}.  The distance to Orion is a critical parameter that influences the interpretation of many of the properties of the region and its members. Measurements range from $\\sim$390 pc to $\\sim$480 pc using a variety of methods and assumptions including observations of water masers, radio sources, an eclipsing binary, and statistical analysis \\citep[and references therein]{gen81, sta04, san07, men07, jef07}. A closer distance would imply that the stars are less luminous, and members that are still contracting onto the main sequence are consequently older based on comparison with theoretical evolutionary tracks \\citep[e.g.,][]{sie00, pal99}. Older ages for the stars above the main sequence suggest a spread in ages.  Binary stars yield a model-independent distance with the combination of a spatially resolved orbit and a double-lined radial velocity orbit \\citep[e.g.,][]{tor97}. Given the distance to Orion, very high angular resolution is required to separate binary systems which exhibit significant orbital motion. The binary system $\\theta^{1}$ Ori C (HD 37022) in the Trapezium region of Orion has such a close separation that it was not detected until speckle observations barely resolved the pair \\citep{wei99} with a separation of less than the diffraction limit of the 6m telescope employed in the discovery. The separation has decreased over time and is now best monitored with interferometry; a recent orbit fit suggested the system might be just past periastron and completing an orbital cycle within a year \\citep{kra07}. In this letter, we present the results of interferometric observations with NPOI that add significantly to the binary orbit phase coverage and suggest the orbital period may be substantially longer than predicted.  The distance to $\\theta^{1}$ Ori C in particular and its physical parameters such as mass and age are of great importance since the O7 primary \\citep[e.g.,][]{gar82}, as the most massive member of the ONC, has the dominant role in shaping the properties of the surrounding nebula, and strongly impacts the circumstellar material around the ONC stars. The photoionizing radiation from $\\theta^{1}$ Ori C produces the brightening of many proplyds \\citep{ode93}, but also causes the material to escape. Observations of mass loss rates \\citep{joh98, hen99} and theoretical modeling of the results \\citep[e.g.,][]{sto99} imply that these structures cannot survive for $\\ga10^5$ yrs, substantially less than the $<$1$-$2 Myr age of the ONC \\citep{hil97}. Possible explanations for the apparent contradiction in disk lifetimes and stellar ages include radial orbits for the proplyds \\citep{sto99} or a very young age for $\\theta^{1}$ Ori C, such as has been proposed for $\\theta^{1}$ Ori B \\citep{pal01}. In contrast, recent models of the evolution of disk sizes in the ONC \\citep{cla07} suggest that disk survival times of 1-2 Myr in the uv field produced by $\\theta^{1}$ Ori C are possible -- consistent with the age of the stellar population. Key to estimating the age of $\\theta^{1}$ Ori C is placing the secondary accurately on the H-R diagram with a well-measured distance, luminosity, and mass given that the primary has already contracted onto the main sequence. We present new magnitude differences which will aid in the assessment of the luminosity, but concentrate on the orbital motion which is required to estimate the distance. ",
        "conclusions": ""
    },
    "0801/0801.0638_arXiv.txt": {
        "abstract": "Recent astronomical observations indicate that the Universe is in the phase of accelerated expansion. There are many cosmological models which explain this phenomenon, but should we prefer those models over the simplest one -- $\\Lambda$CDM model? According to the Occam's razor principle if all models describe the observations equally well we should prefer the simplest one. We consider the model comparison methods which involve such rules: the Akaike information criterion (AIC), Bayesian information criterion (BIC) and Bayesian evidence to compare the $\\Lambda$CDM model with its generalisation where the interaction between dark matter and dark energy is allowed. The analyses based on the AIC and Bayesian evidence indicate that there is only a weak evidence in favour of the $\\Lambda$CDM model over its generalisation, while those based on BIC quantity indicate the strong evidence in favour the simpler model. We also calculate some quantity which measure the effective number of model parameters that the given data can constrain. This value is used to compare the concordance LCDM model with its generalization basing on the extended interpretation of continuity condition-interacting $\\Lambda$CDM cosmology. We conclude that data set are not enough informative to constrain all parameters allows to vary in the models. ",
        "introduction": "Recent observations of type Ia supernova (SNIa) provide the main evidence that current Universe is in an accelerating phase of expansion \\cite{Riess:1998cb}. Cosmic microwave background (CMB) data indicate that the present Universe has also negligible space curvature \\cite{Spergel:2006hy}. Therefore if we assume the Friedmann-Robertson-Walker (FRW) model in which effects of nonhomogeneities are neglected, then the acceleration must be driven by dark energy component $X$ (matter fluid violating the strong energy condition $\\rho_{\\text{X}}+3p_{\\text{X}}\\geq 0)$. This kind of energy represents roughly $70\\%$ of the matter content of the current Universe. Because the nature as well as mechanism of the cosmological origin of the dark energy component are unknown some alternative theories try to eliminate the dark energy by modifying the theory of gravity itself. The main prototype of this kind of model is covariant brane models based on the Dvali-Gabadadze-Porrati (DGP) model \\cite{Dvali:2000hr} as generalized to cosmology by Deffayet \\cite{Deffayet:2000uy}. The simplest explanation of a dark energy component is the cosmological constant with effective equation of state $p=-\\rho$ but appears the problem of its smallness and hence its relatively recent dominance. Although the $\\Lambda$CDM model offers possibility of explanation of observational data it is only effective theory which contain the enigmatic theoretical term -- the cosmological constant $\\Lambda$. Other numerous candidates for dark energy description have also been proposed like to evolving scalar field \\cite{Peebles:1987ek} usually referred as quintessence, the phantom energy \\cite{Caldwell:1999ew, Dabrowski:2003jm}, the Chaplygin gas \\cite{Kamenshchik:2001cp} etc. Some authors believed that the dark energy problem belongs to the quantum gravity domain \\cite{Witten:2000zk}.  These theoretical models are consistent with the observations, they are able to explain the phenomenon of the accelerated expansion of the Universe. But should we really prefer such models over the $\\Lambda$CDM one? To answer this question we should use some model comparison methods to confront existing cosmological models having observations at hand.  Let us assume that we have $N$ pairs of measurements $(y_i,x_i)$ and that we want to find the relation between the $y$ and $x$ quantities. Suppose that we can postulate $k$ possible relations $y\\equiv f_i(x,\\bar{\\theta})$, where $\\bar{\\theta}$ is the vector of unknown model parameters and $i=1,\\dots,k$. With the assumption that our observations come with uncorrelated gaussian errors with mean $\\mu_i=0$ and standard deviation $\\sigma_i$ the goodness of fit for the theoretical model is measured by the $\\chi^2$ quantity given by \\begin{equation} \\chi^2=\\sum_{i=1}^{N} \\frac{(f_l(x_i,\\bar{\\theta}) - y_i)^2}{2\\sigma_i^2}=-2\\ln L, \\end{equation} where $L$ is the likelihood function. For the particular family of models $f_l$ the best one minimize the $\\chi^2$ quantity, which we denote $f_l(x,\\hat{\\bar{\\theta}})$. The best model from our set of $k$ models ${f_1(x,\\hat{\\bar{\\theta}}),\\dots,f_k(x,\\hat{\\bar{\\theta}})}$ could be the one with the smallest value of $\\chi^2$ quantity. But this method could give us misleading results. Generally speaking for more complex model the value of $\\chi^2$ is smaller, thus the most complex one will be choose as the best from our set under consideration.  A clue is given by so called Occam's razor principle: ``If two models describe the observations equally well, choose the simplest one.'' This principle has aesthetic as well as empirical justification. Let us quote simple example which illustrate this rule \\cite{MacKay:2003} and present it in Figure \\ref{Fig:1}. We see the black box and the white one behind it. One can postulate two models: 1. There is one box behind the black box, 2. There are two boxes of identical height and colour behind the black box. Both models explain our observations equally well. According to the Occam's razor principle we should accept the explanation which is simpler so that there is only one white box behind the black one. Is not it more probable that there is only one box than two boxes with the same height and colour?  \\begin{figure}[h] \\centering \\includegraphics[scale=0.5]{fig1.eps} \\caption{Illustration to the example explaining the Occam's razor principle.} \\label{Fig:1} \\end{figure}  We could not use this principle directly because the situations when two models explain the observations equally well are rare. But in the information theory as well as in the bayesian theory there are methods for model comparison which include such rule.  In the information theory there are no true models. There is only reality which can be approximated by models, which depend on some number of parameters. The best one from the set under consideration should be the best approximation to the truth. The information lost when truth is approximated by model under consideration is measured by so called Kullback-Leibler (KL) information so the best one should minimizes this quantity. It is impossible to compute the KL information directly because it depends on truth which is unknown. Akaike \\cite{Akaike:1974} found approximation to the KL quantity which is called the Akaike information criterion (AIC) and is given by \\begin{equation} \\text{AIC}=-2\\ln \\mathcal{L} +2d, \\end{equation} where $\\mathcal{L}$ is the maximum of the likelihood function and $d$ is the number of model parameters. Model which is the best approximation to the truth from the set under consideration has the smallest value of the AIC quantity. It is convenient to evaluate the differences between the AIC quantities computed for the rest of models from our set and the AIC for the best one. Those differences ($\\Delta_{\\text{AIC}}$) are easy to interpret and allow a quick `strength of evidence' for considered model with respect to the best one. The models with $0 \\le \\Delta_{\\text{AIC}}\\le 2$ have substantial support (evidence), those where $4<\\Delta_{\\text{AIC}}\\le 7$ have considerably less support, while models having $\\Delta_{\\text{AIC}} > 10 $ have essentially no support with respect to the best model.  It is worth noting that the complexity of the model is interpreted here as the number of its free parameters that can be adjusted to fit the model to the observations. If models under consideration fit the data equally well according to the Akaike rule the best one is with the smallest number of model parameters (the simplest one in such an approach).  In the Bayesian framework the best model (from the model set under consideration) is that which has the largest value of probability in the light of data (so called posterior probability) \\cite{Jeffreys:1961} \\begin{equation} P(M_{i}|D)=\\frac{P(D|M_{i})P(M_{i})}{P(D)}, \\end{equation} where $P(M_{i})$ is a prior probability for the model $M_{i}$, $P(D)$ is normalization constant [$P(D)= \\sum _{i=1}^{k} P(D|M_{i})P(M_{i})$], $D$ denotes data. $P(D|M_{i})$ is the marginal likelihood, also called evidence \\begin{equation} P(D|M_{i})=\\int P(D|\\bar{\\theta},M_{i})P(\\bar{\\theta}|M_{i}) d \\bar{\\theta} \\equiv E_{i}, \\end{equation} where $P(D|\\bar{\\theta},M_{i})$ is likelihood under model $i$, $P(\\bar{\\theta}|M_{i})$ is prior probability for ${\\bar{\\theta}}$ under model $i$.  Let us note that we can include the Occam's razor principle by assuming the greater prior probability for simpler model, but this is not necessary and rarely used in practice. Usually one assume that there is no evidence to favor one model over another which cause to equal value of prior for all models under consideration. It is convenient to evaluate the posterior ratio for models under consideration which in the case with flat prior for models is reduced to the evidence ratio (so called the Bayes factor $B$). The interpretation of the logarithm of Bayes factor values are as follows: $0<\\ln B\\leq 2$ as a weak, $2<\\ln B\\leq 6$ as a positive, $6<\\ln B\\leq 10$ as a strong and $\\ln B> 10$ as a very strong evidence in favor better model. This quantity involves the Occam's razor rule. Let us simplify the problem to illustrate how this principle works here \\cite{MacKay:2003,Trotta:2005ar}.  Assume that $\\bar{P}(\\bar{\\theta}|D,M)$ is the non normalized posterior probability for the vector $\\bar{\\theta}$ of model parameters. In this notation $E=\\int\\bar{P}(\\bar{\\theta}|D,M)d\\bar{\\theta}$. Suppose that posterior has a strong pick in the maximum: $\\bar{\\theta}_{\\text{MOD}}$. It is reasonable to approximate the logarithm of the posterior by its taylor expansion in the neighbour of $\\bar{\\theta}_{\\text{MOD}}$ so we finished with the expression \\begin{equation} \\bar{P}(\\bar{\\theta}|D,M)=\\bar{P}(\\bar{\\theta}_\\text{MOD}|D,M)\\exp \\left[-(\\bar{\\theta}-\\bar{\\theta}_{\\text{MOD}})^T\\text{C}^{-1}(\\bar{\\theta}-\\bar{\\theta}_{\\text{MOD}})\\right], \\end{equation} where $\\left[\\text{C}^{-1}\\right]_{ij}=-\\left[\\frac{\\partial^2\\ln\\bar{P}(\\bar{\\theta}|D,M)}{\\partial\\theta_i\\partial\\theta_j}\\right]_{\\bar{\\theta}=\\bar{\\theta}_\\text{MOD}}$. Posterior is approximated by the gaussian distribution with the covariance matrix $\\text{C}$ and the mean $\\bar{\\theta}_\\text{MOD}$. The expression for the evidence take a form $E=\\bar{P}(\\bar{\\theta}_{\\text{MOD}}|D,M)\\int\\exp \\left[-(\\bar{\\theta}-\\bar{\\theta}_{\\text{MOD}})^T\\text{C}^{-1}(\\bar{\\theta}-\\bar{\\theta}_{\\text{MOD}})\\right]d \\ \\bar{\\theta}.$ Because the posterior has a strong pick near the maximum, the most contribution to the integral comes from the neighbour close to $\\bar{\\theta}_\\text{MOD}$. Contribution from the other other region of $\\bar{\\theta}$ can be ignored, so we can expand the limit of the integral to whole $R^d$. With this assumptions one can obtain $E=(2\\pi)^{\\frac{d}{2}}\\sqrt{\\det\\text{C}}\\bar{P}(\\bar{\\theta}_{\\text{MOD}}|D,M)= (2\\pi)^{\\frac{d}{2}}\\sqrt{\\det\\text{C}}P(D|\\bar{\\theta}_\\text{MOD},M)P(\\bar{\\theta}_\\text{MOD}|M)$. Suppose that the likelihood function has sharp pick in $\\hat{\\bar{\\theta}}$ and the prior for $\\bar{\\theta}$ is nearly flat in the neighbour of $\\hat{\\bar{\\theta}}$. In this case  $\\hat{\\bar{\\theta}}=\\bar{\\theta}_{\\text{MOD}}$ and the expression for the evidence takes the form $E=\\mathcal{L}(2\\pi)^{\\frac{d}{2}}\\sqrt{\\det\\text{C}}P(\\hat{\\bar{\\theta}}|M)$. The $(2\\pi)^{\\frac{d}{2}}\\sqrt{\\det\\text{C}}P(\\hat{\\bar{\\theta}}|M)$ quantity is called the Occam factor (OF). When we consider the case with one model parameter with flat prior $P(\\theta|M)=\\frac{1}{\\Delta\\theta}$ the Occam factor OF$=\\frac{2\\pi\\sigma}{\\Delta\\theta}$ which can be interpreted as the ratio of the volume occupied by the posterior to the volume occupied by prior in the parameter space. The more parameter space wasted by the prior the smaller value of the evidence. It is worth noting that the evidence does not penalize parameters which are unconstrained by the data \\cite{Liddle:2006kn}.  As the evidence is hard to evaluation an approximation to this quantity was proposed by Schwarz \\cite{Schwarz:1987} so called Bayesian information criterion (BIC) and is given by \\begin{equation} \\text{BIC}=-2\\ln\\mathcal{L}+2d\\ln N, \\end{equation} where $N$ is the number of the data points. The best model from the set under consideration is this which minimize the BIC quantity. It is also convenient to analyse the difference between BIC quantities for the rest models from our set with the BIC for the best one. These differences could be interpreted in the following way: $0<\\Delta_{\\text{BIC}}\\leq 2$ as a weak, $2<\\Delta_{\\text{BIC}}\\leq 6$ as a positive, $6<\\Delta_{\\text{BIC}}\\leq 10$ as a strong and $\\Delta_{\\text{BIC}}> 10$ as a very strong evidence in favour of a better model.  One can notice the similarity between the AIC and BIC quantities though they come from different approaches to model selection problem. The dissimilarity is seen in the so called penalty term: $ad$, which penalize more complex models (complexity is identified here as the number of free model parameters). One can evaluated the factor by which the additional parameter must improve the goodness of fit to be included in the model. This factor must be greater than $a$ so equal to $2$ in the AIC case and equal to $\\ln N$ in the BIC case. Notice that the latter depends on the number of the data points.  It should be pointed out that presented model selection methods are widely used in context of cosmological model comparison \\cite{Hobson:2002de, Beltran:2005xd, Mukherjee:2005wg, Mukherjee:2005tr, Trotta:2005ar, Niarchou:2003hz, Liddle:2004nh, Saini:2003wq, John:2002gg, Parkinson:2004yx, John:2005bz, Serra:2007id, Biesiada:2007um, Liddle:2006kn, Godlowski:2005tw, Szydlowski:2006ay, Szydlowski:2005xv, Kunz:2006mc, Parkinson:2006ku, Trotta:2007hy, Liddle:2007fy, Kurek:2007tb}.  We should keep in mind that conclusions based on such quantities depend on the data at hand. Let us mention again the example with the black box. Suppose that we made a few steps toward this box that we can see the difference between the height of the left and right side of the white box. Our conclusion change now.  Let us quote example taking from \\cite{John:2005bz}. Assume that we want to compare the Newtonian and Einsteinian theories in the light of the data coming from laboratory experiment where general relativistic effects are negligible. In this situation Bayes factor between Newtonian and Einsteinian theories will be close to unity. Whereas comparing the general relativistic and Newtonian explanations of the deflection of a light ray that just grazes the sun's surface give the Bayes factor $\\sim 10^{10} $ in the favor of the first one (and even greater with more accurate data).  Having this in mind an interesting supplement to the above considerations seems to be quantity which measure the effective number of model parameters (C$_b$) that the given data set can constrain \\cite{Kunz:2006mc,Liddle:2007fy}. It can be computed using the relation \\begin{equation} \\text{C}_b = \\overline{\\chi^2(\\bar{\\theta})} - \\chi^2(\\hat{\\bar{\\theta}}), \\end{equation} where the mean is taken over the posterior pdf for $\\bar{\\theta}$ and $\\hat{\\bar{\\theta}}$ is the mode of the posterior pdf (different choices are also possible). As have been shown this quantity correspond to the number of parameters for which the width of the posterior probability distribution is significantly narrower than the width of the prior probability distribution \\cite{Kunz:2006mc}. These parameters can be considered to have been well measured by the data given our prior assumption in the model. It helps to determine if the data is informative enough to measure the parameters under consideration.  We share with George Efstathiou opinion \\cite{Efstathiou:2007gz,Chongchitnan:2007eb,Szydlowski:2004jv} that there is no sound theoretical basis for considering the dynamical dark energy, where as we are beginning to see an explanation for a small cosmological constant emerging from more fundamental theory. In our opinion the $\\Lambda$CDM model has the status of satisfactory effective theory. Estathiou argued why the cosmological constant should be given higher weight as a candidate for dark energy description than dynamical dark energy. In this argumentation Occam's razor is used to point out a more economical model explaining the observational data. On the other hand Biesiada advocated the use of Akaike information criterion which favour rather a dynamical model of dark energy (quintessence model) \\cite{Biesiada:2007um}. In our opinion quantification of Occam's razor by computing Bayesian evidence and Bayesian information criterion is more suitable as the AIC is unadequate for many statistical problems \\cite{Dowe:2007bn}.  The main aim of this paper is to compare the simplest cosmological model -- the $\\Lambda$CDM model -- with its generalisation where the interaction between dark energy and dark matter sector is allowed using methods described above. ",
        "conclusions": "We presented the methods of model comparison coming from information as well as bayesian theory. Both of them include the Occam's razor principle which states that if two models describe the observations equally well we should choose the simpler one. According to Akaike and Schwarz rule the model complexity is interpreted in the term of number of free model parameter while according to the Bayesian evidence more complex model wast greater volume of the parameter space. Finally we present the quantity which measures the effective number of model parameters which the data used in analysis can constrain. This quantity is also called the Bayesian complexity and reduce to the number of free model parameters when the data set is highly informative. We use those methods to answer the question if we should prefer the generalisation of the simplest $\\Lambda$CDM model where the interaction between dark matter and dark energy sector is allowed. The AIC and bayesian evidence give similar results: there is only weak evidence to favor the $\\Lambda$CDM model over the interacting $\\Lambda$CDM one. This is with contrary with the conclusion from the analysis of the BIC quantity: here there is a strong evidence to favor the $\\Lambda$CDM model. The analysis of the Bayesian complexity gives that data sets considered carry not enough information to constraint all parameters allows to vary in the models considered."
    },
    "0801/0801.0312_arXiv.txt": {
        "abstract": "We compare environmental effects in two analogous samples of galaxies, one from the Sloan Digital Sky Survey (SDSS) and the other from a semi-analytic model (SAM) based on the Millennium Simulation (MS), to test to what extent current SAMs of galaxy formation are reproducing environmental effects. We estimate the large-scale environment of each galaxy using a Bayesian density estimator based on distances to all ten nearest neighbors and compare broad-band photometric properties of the two samples as a function of environment. The feedbacks implemented in the semi-analytic model produce a qualitatively correct galaxy population with similar environmental dependence as that seen in SDSS galaxies. In detail, however, the colors of MS galaxies exhibit an exaggerated dependence on environment: the field contains too many blue galaxies while clusters contain too many red galaxies, compared to the SDSS sample. We also find that the MS contains a population of highly clustered, relatively faint red galaxies with velocity dispersions comparable to their Hubble flow. Such high-density galaxies, if they exist, would be overlooked in any low-redshift survey since their membership to a cluster cannot be determined due to the ``Fingers of God'' effect. ",
        "introduction": "Since the discovery of the morphology-environment relation \\citep{Dressler_1980}, it has been known that galaxy properties are correlated with their large-scale environment: the average morphology, color and luminosity of galaxies differ depending on how crowded their neighborhood is. On the face of it, it is not clear why or how the environment of a galaxy on Mpc scales should be related to the kpc-scale processes (star formation, supernova and AGN feedback) that determine the bulk properties of a galaxy. To further confuse matters, the strong correlation between morphology, color and luminosity \\citep[][and references therein]{Strateva_2001} makes it unclear which property is ultimately driven by environment let alone which physical processes are responsible. It is not even clear to what extent the environment of a galaxy effects it through nature (different formation conditions) rather than nurture (galaxy-galaxy interactions). A critical step towards answering such questions is to compare observed trends to those present in a simulated galaxy ensemble in which one knows all the processes at work.  The Sloan Digital Sky Survey \\citep[SDSS,][]{Adelman_2007} is a powerful tool for addressing questions of environmental effects. Its spectroscopic sample of galaxies is the largest such sample ever, ensuring that even relatively rare galaxy populations are well represented, and the survey's large contiguous footprint makes it easy to determine the large-scale environment for most of these galaxies. Previous researchers who have used the SDSS galaxy catalog to study environmental dependences have found three broad trends: the peaks of the bimodal color distribution of galaxies do not shift for different environments; blue and red galaxies are most common in low- and high-density environments, respectively; the luminosity of red galaxies increases with local density \\citep{Hogg_2003,Kauffmann_2004,Balogh_2004,Tanaka_2004,Blanton_2005,Zehavi_2005,Park_2007, Ball_2007}.  The Millennium Simulation \\citep[MS,][]{Springel_2005} is the largest ever cosmological simulation comprising some $10^{10}$ dark matter (DM) particles with a spatial resolution of $5h_{100}^{-1}$ kpc. At $z=0$, the simulation fills a cube $500h_{100}^{-1}$ Mpc per side. The MS does not explicitly model the gas, dust and stars which make up observable galaxies, but it produces a DM halo merger tree which serves as the backbone for a number of semi-analytic models (SAMs). Unlike N-Body/SPH simulations, SAMs do not simulate the gravitational and hydrodynamic forces involved in the formation and evolution of galaxies, but they do provide a computationally inexpensive way to explore the parameter space of sub-grid processes. The trade-off is that the parameters of a SAM must be tuned using observations (e.g. matching to the observed luminosity function), making truly independent comparisons between the model and reality more challenging. Numerous groups have developed SAMs which hierarchically form some $10^{7}$ galaxies from the MS merger tree  \\citep[eg:][and references therein]{Croton_2006, Bower_2006, DeLucia_2007}. Their models differ ---for example in their treatment of AGN feedback--- but all reproduce some of the empirical features of galaxy populations. The MS galaxies have a very realistic distribution of luminosities, thanks to judicious use of ``radio'' feedback. They also exhibit a bimodal color distribution as discovered in SDSS \\citep{Strateva_2001}. Finally, the power spectrum of the density fluctuations is in good agreement with the empirical data from 2dF and SDSS \\citep{Springel_2005}. Previous investigators have found that the brightest MS galaxies are red, dead ellipticals populating rich galaxy clusters \\citep{DeLucia_2006}, while the modeled galaxies in the very lowest-density environments have similar colors and star formation rates as analogous SDSS galaxies \\citep{Patiri_2006}.  In this work we compare the observed galaxy populations with those produced with SAMs. Our work differs from those listed above in the following ways: we use Bayesian number density as a proxy for local environment, rather than the commonly used surface density or two-point correlation function; we use the $u-r$ color, which has more leverage than the $g-r$ color; we use SDSS Data Release 5, rather than any of the previous (smaller) releases; and last but not least, we compare observed and modeled galaxies for the full range of galaxy environments and colors. ",
        "conclusions": "We have compared two analogous galaxy samples, one from the SDSS DR5 spectroscopic sample, and one from the SAMs of \\cite{DeLucia_2007}, after correcting for the observational effects present in the former. The density distribution and the luminosity function of the modeled galaxies qualitatively match those for the SDSS sample, but there are $50$\\% too many blue galaxies in the former. In detail, two additional discrepancies become apparent between the galaxy samples: MS galaxies are more blue in $u-r$ than SDSS galaxies; the colors of galaxies depend more strongly on environment in MS than in SDSS. The strong environmental dependence manifests itself as an over-representation of blue galaxies overall and suggests that the feedbacks implemented by \\cite{DeLucia_2007} exaggerate the role of galaxy environment. A population of relatively faint red galaxies in extremely high density environments is visible in the MS survey \\emph{sans} observational effects. Such high density environments would be imperceptible in SDSS due to velocity dispersions comparable to local Hubble flow."
    },
    "0801/0801.2121_arXiv.txt": {
        "abstract": "According to recent experimental data at GSI, the rate of the number of daughter ions ${^{140}}{\\rm Ce}^{58+}$, produced by the nuclear K--shell electron capture ($EC$) decay of the H--like ion ${^{140}}{\\rm Pr}^{58+}$, is modulated in time with a period $T_{EC} = 7.06(8)\\,{\\rm sec}$ and an amplitude $a_{EC} = 0.18(3)$.  We show that this phenomenon can be explained by neutrino mass differences and derive a value for the difference of squared masses $\\Delta m^2_{21} = m^2_2 - m^2_1 = 2.22(3)\\times 10^{-4}\\,{\\rm eV}^2$.  PACS: 12.15.Ff, 13.15.+g, 23.40.Bw, 26.65.+t ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.1122_arXiv.txt": {
        "abstract": "We present theoretical models for the formation and evolution of populations of low-mass X-ray binaries (LMXB) in the two elliptical galaxies NGC 3379 and NGC 4278. The models are calculated with the recently updated \\emph{StarTrack} code \\citep{Belczynski2006}, assuming only a primordial galactic field LMXB population. StarTrack is an advanced population synthesis code that has been tested and calibrated using detailed binary star calculations and incorporates all the important physical processes of binary evolution. The simulations are targeted to modeling and understanding the origin of the X-ray luminosity functions (XLF) of point sources in these galaxies. For the first time we explore the population XLF down to luminosities of $3\\times10^{36}\\, \\rm erg\\,s^{-1}$, as probed by the most recent observational results \\citep{Kim2006}. We consider models for the formation and evolution of LMXBs in galactic fields with different CE efficiencies, stellar wind prescriptions, magnetic braking laws and initial mass functions. We identify models that produce an XLF in excellent agreement with the observations both in shape and absolute normalization. We also find that the treatment of the outburst luminosity of transient systems remains a crucial factor for the determination of the XLF since the modeled populations are dominated by transient X-ray systems. ",
        "introduction": "A Low Mass X-ray binary (LMXB) is a Roche lobe overflowing, mass-transfering binary system with a compact object accretor, either a black hole (BH) or a neutron star (NS), and a low mass ($\\gtrsim 1\\,\\rm M_{\\odot}$) donor.  Since the late 80's it has been suggested that LMXBs should exist in early type galaxies, E and S0, and that they might even dominate the X-ray emission \\citep{TF1985, fabbiano1989, KFT1992}. The stellar populations in these galaxies are typically old and homogeneous. Massive stars have already evolved to compact objects and LMXBs are probably the only sources with X-ray luminosities above $10^{36}\\, \\rm erg\\,s^{-1}$.  Uncontroversial detection of LMXBs in early type galaxies became possible only this last decade with \\textit{Chandra's} increased angular resolution \\citep{Fabbiano2006, SIB2000}. The spectra of individual X-ray sources are consistent with those expected from LMXB models and the LMXBs observed in the Milky-way and M31 \\citep{HB2006, IAB2003}. For many galaxies observed with \\textit{Chandra}, the XLFs have been derived and they can usually be fitted with a single or a broken power law. The detections limit for these  surveys is usually a few times $10^{37}\\,\\rm erg\\,s^{-1}$. \\citet{KF2004} derived XLFs for 14 early type galaxies and they included completeness corrections. Each XLF is well fitted with a single power law with cumulative slope between -0.8 and -1.2. The composite XLF of these galaxies though is not consistent with a single power law. There is a prominent break at $(5\\pm1.6)\\times 10^{38}\\, \\rm erg\\,s^{-1}$, close to the Eddington luminosity ($L_{\\rm Edd}$) of a helium accreting NS-LMXB. This break might be hidden in the individual XLFs due to poor statistics \\citep[see also][]{SIB2000, KMZ2002, Jordan2004, Gilfanov2004}. Other recent studies by \\citet{JCBG2003},  \\citet{SSI2003} and \\citet{Jordan2004} suggested a break of the XLF at a higher luminosity ($\\sim 10^{39}\\, erg\\,s^{-1}$). The exact position and the nature of these breaks are still somewhat controversial, as the correct interpretation of the observed XLFs relies significantly on the proper completeness correction when looking at luminosities close to the detection limit and small number statistics at the high-end of the XLF.  Recent \\textit{Chandra} observations \\citep{Kim2006} have yielded the first low-luminosity XLFs of LMXBs for two typical old elliptical galaxies, NGC3379 and NGC4278.  The detection limit in these observations is $\\sim 3\\times 10^{36}\\rm \\, erg\\, s^{-1}$ which is about an order of magnitude lower than in most previous surveys of early type galaxies. The observed XLFs of the two ellipticals extend only up to $6\\times 10^{38}\\,\\rm erg\\, s^{-1}$ and are well represented by a single power law with a slope (in a differential form) of $1.9\\pm0.1$.  When \\textit{Chandra} observations are compared with optical images from \\textit{Hubble} or other ground based telescopes, it is generally found that a significant fraction of the LMXBs are inside globular clusters (GC). On average 4-5\\% of the GCs in a given galaxy are associated with a LMXB (with $L_x>\\sim 10^{37}\\rm \\,erg \\, s^{-1}$), while the fraction of LMXBs located in GC, varies from 10\\% to 70\\% depending on the type of the galaxy and its GC specific frequency. At present the origin and the properties of these systems, both in GCs and the field, are not yet well understood. It has been noted that XLFs at high luminosities for each sub-group (GCs and field) do not reveal any differences within the statistics of the samples considered \\citep{Fabbiano2006, KimE2006, KMZ2007, Jordan2004, Sarazin2003}. However, in more recent studies, \\citep{Fabbiano2007} and \\citet{VG2007} independently found that the two XLFs (GCs and field sources) show significant differences at low luminisities (below $10^{37}\\rm \\, erg\\, sec^{-1}$), pointing to a different LMXB formation mechanism in GCs.  The natural question that arises is whether  (\\textit{i}) all LMXBs were formed in GCs through dynamical interactions and some eventually escaped or some GCs dissolved in the field, or (\\textit{ii}) field LMXBs were born in situ through binary evolution of primordial binaries. The formation rates associated with these two possibilities are not understood well enough to give accurate predictions and use the relative numbers in the samples. \\citet{Juett2005} has shown that the observed relationship between the fraction of LMXBs found in GCs and the GC-specific frequency in early-type galaxies is consistent with the galactic field LMXB population being formed in situ. Similarly, \\citet{Irwin2005} compared the summed X-ray luminosity of the LMXBs to the number of GCs in a galaxy; in the case of all LMXBs having formed exclusively in GCs, the two should be directly proportional regardless of where the LMXBs currently reside. Instead he found that the proportionality includes an additive offset implying the existence of a LMXB population unrelated to GCs.     In the past, semi-analytical theoretical models have been introduced for the study  of the LMXB population in early galaxies. \\citet{White1998} studied the connection between the star formation rates of normal galaxies, i.e. galaxies without an active nucleus, and the formation rate of LMXBs and millisecond pulsars, assuming that all LMXBs are formed from primordial binaries. Considering a time-dependent star formation rate, they showed that the general relativity timescales relevant to the evolution of primordial binaries to LMXBs and to millisecond pulsars, lead to a significant time delay of the peak in the formation rate of these populations after the peak in the star formation rate. In a followup work \\citet{Ghosh2001}, using several updated star formation rate models, calculated the evolution of the X-ray luminosity of galaxies. They found that different star formation models lead to very different X-ray luminosity profiles, so the observed X-ray profiles can be used as probes of the star formation history. Finally they compared their models with the first \\textit{Chandra} deep imaging observations, to conclude that these first results were consistent with current star formation models. \\citet{Piro2002} argued that the majority of LMXBs in the field of elliptical galaxies have red giant donors feeding a thermally unstable disk and stay in this transient phase for at least 75\\% of their life. The very luminous X-ray sources ($L_x>10^{39} \\rm \\, erg\\, s^{-1}$) detected in \\textit{Chandra} surveys have been suggested to be X-ray binaries with highly super-Eddington mass inflow near the accreting component. In elliptical galaxies these objects have been suggested by King 2002 to be micro-quasar-like, as these galaxies contain no high mass X-ray binaries \\citep{King2002}. More recently  \\citet{Ivanova2006} also argued that this bright end of the XLF is most likely dominated by transient LMXBs with BH accretors during outburst and it can be used to derive constraints on the BH mass function in LMXBs; they also showed that the standard assumption of a constant transient duty cycle (DC) across the whole population seems to be inconsistent with current observations.  Semi-analytical population synthesis (PS) models of LMXBs have also been constructed for late type galaxies. \\citet{Wu2001} created a simple birth-death model, in which the lifetimes of the binaries are inversely proportional to their X-ray luminosity, and calculated the XLFs of spiral galaxies. His models reproduce some features, such as the luminosity break in the observed XLFs of spiral galaxies. The position of this break depends on the star formation history of the galaxy, and he suggested that it can be used as a probe of the galaxy's merger history.  The formation of LMXBs in GCs via dynamical interactions is less well studied, since apart from the binary stellar evolution, one has to also take into account the complex cluster dynamics.  \\citet{BD2004} considered a semi-analytical model for accretion from degenerate donors onto NSs in ultracompact binaries and showed that binaries with orbital periods of 8-10 minutes and He or C/O white dwarf (WD) donors of 0.06-0.08 $M_\\odot$ naturally provide the primary slope (-0.8 for cumulative form) typically derived from XLFs of elliptical galaxies. Ultra-compact systems are predicted to form in the dense GC environment and have relatively short persistent lifetimes ($<3\\times 10^6\\, \\rm yr$) but they form continuously through dynamical interactions. \\citet{Ivanova2007} presented PS studies of compact binaries containing NSs in dense GCs. They used \\emph{StarTrack} as their PS modeling tool in addition to a simplified treatment for the dynamical interactions. Their models produced a mixed population of LMXBs with red giant and MS donors, and  ultra-compact X-ray binaries; relative formation rates can be comparable but the different sub-populations have very different lifetimes.      In this paper we investigate the plausibility of an important contribution to the XLFs of these two galaxies from a primordial galactic field LMXB population using advanced PS simulations. In \\S 2 we describe briefly the physics included in our PS code and explain in detail the way we are constructing the modeled XLFs and the treatment of transient LMXBs.  We discuss the results of our simulations in \\S 3: the modeled XLFs from different models, a statistical comparison with the observed XLFs of the elliptical galaxies NGC3379 and NGC4278, and an analysis of the dependence of the modeled XLF properties on the PS parameters. Finally in \\S 4 we discuss the implication of our findings and the caveats of our methods. ",
        "conclusions": "The recent deep \\textit{Chandra} observations \\citep{Kim2006} of the two typical old elliptical galaxies: NGC3379 and NGC4278, led to the first observed low-luminosity XLFs of LMXBs, with the detection limit ($3\\times 10^{36}\\rm \\, erg\\, s^{-1}$) being about an order of magnitude lower than in most previous surveys. Motivated by this observational work, we developed PS simulations of LMXBs appropriate for these two galaxies. We considered formation of LMXBs only through the evolution of primordial binaries in the galactic fields and examined the possible contribution to the overall LMXB population. For our modeling we used the updated PS code \\emph{StarTrack} \\citep{Belczynski2006}.  Our main conclusions can be summarized as follows:  We found that some of our models produce XLFs in very good agreement with the observations, based on both the XLF shape and absolute normalization. There is no unique combination of PS parameters and modeling of transient sources (DC and outburst luminosity) that gives an XLF in agreement with the observation. We conclude that formation of LMXBs in the galactic field via evolution of primordial binaries can have a significant contribution to the total population of an elliptical galaxy, especially the ones with low GC specific frequency such as NGC3379 \\citep{Fabbiano2007}. Nevertheless, we are able to exclude the majority of our models, as inconsistent with the observations. Note that widely used, simple assumptions such as that all transients source in the outburst state are emitting X-rays at $L_{\\rm Edd}$, lead to XLFs clearly inconsistent with the observed ones. Our results appear to be robust since we do not have to fine-tune our code parameters in order to get a model that resembles the observed population.  As already suggested by \\citet{Piro2002}, the LMXB population has a significant contribution by transient systems (thermal disk instability) and with reasonable outburst  DCs they can even dominate the XLF. As a consequence the XLF shape is rather sensitive to the treatment of these transient systems. In Figure \\ref{xlf_dc} we show that keeping the same PS parameters and changing \\emph{only} the modeling of transient sources leads to completely different XLFs. We tried different methods of modeling the outburst characteristics of transient LMXBs and we found that: \\emph{(i)} When we assume the outburst luminosity of all transient LMXBs to be equal to $L_{\\rm Edd}$ (transient treatment E) or apply eq.(\\ref{Lx_P}) (transient treatment F) - which was empirically derived for Galactic BH-LMXBs - to the whole population, we get XLFs inconsistent with the observed ones regarding both their shape and the total number of sources predicted. \\emph{(ii)} A constant DC for all systems, although not physically motivated, can be sometimes a good first approximation. \\emph{(iii)} We get the best agreement with observations, when we consider a variable DC for NS-LMXB based on the theoretical study of \\citet{DLM2006} (see eqs. \\ref{dc} and \\ref{Lx_phys}), while for the BH-LMXBs we use the empirical correlation between orbital period and outburst luminosity, derived by \\citet{PDM2005} (see eq. \\ref{Lx_P}), and assuming a low constant DC ($\\sim 5\\%$).  The LMXB sub-populations that mainly contribute to the model XLFs are NS-LMXBs with red giant donors and BH-LMXBs with MS donors. A population of persistent ultra compact LMXBs with WD donors is also present in our models and in some cases has an important contribution too (see models $\\rm 9A^{IT}$, $\\rm 11A^{IT}$ and $\\rm 12A^{IT}$).  Of these sub-populations, the NS-LMXBs are the most dominant and primarily determine the XLF shape in the medium and low luminosity range (below $2 \\times 10^{38}\\rm \\, erg \\, s^{-1}$), while the BH-LMXBs have a significant contribution to the high-end of the XLF.   The normalization of the modeled XLFs is a less robust characteristic than its shape. We normalize the models so that the number of the primordial binaries we evolve, correspond to the known galaxy masses, given the IMF and the binary fraction. There are however uncertainties of the order of a few in the determination of the mass of the observed galaxies, due to uncertainties in their distance, the bolometric luminosity and the light-to-mass ratio. The majority of the models presented here  produce the observed number of LMXBs to within a factor of 3, consistent with the galaxy mass uncertainties. Exceptions are  models with transient treatment E or F (see Table \\ref{models_tran}) and models with high CE efficiencies (models 21-28, see \\ref{models}) which greatly overproduce hight luminosity LMXBs. Furthermore small changes in the CE efficiency can change the total number of sources produced by a model by a factor of two, without changing significantly the shape of the XLF (compare for example models $\\rm 10A^{IT}$, $\\rm 14A^{IT}$, $\\rm 18A^{IT}$ in the online supplemental material). In view of these uncertainties and our limited parameter space exploration for the models, we consider this normalization agreement satisfactory, but do not use it as an actual constraint on the models.  We do not claim that the work presented in this Paper is a complete PS study of field LMXBs in elliptical galaxies. It is meant to be a first effort in interpreting the recent deep \\emph{Chandra} observations of the two elliptical galaxies NGC3379 and NGC4278. Throughout the Paper we identify the caveats of our analysis, and we intend to address them in our future work. It turns out that the realistic treatment of the outburst properties of transient LMXBs is crucial for the modeling of XLFs of extragalactic populations. The derivation of an empirical correlation between the outburst luminosity and the period of Galactic NS-LMXBs, similar to the one derived by \\citet{PDM2005} for Galactic BH-LMXBs,  will provide a better understanding of the differences in the transient behavior of these two classes of X-ray sources. Another simplifying assumption we made, is the constant outburst luminosity of transient LMXBs. The use of model lightcurves  for the outburst phases of transient XLFs will possibly affect the shape of the modeled XLFs. We note that most of our models produce many systems with X-ray luminosity below the observational limit and down to $10^{35}\\rm \\, erg \\, s^{-1}$. The integrated diffuse X-ray emission from these galaxies can put a strong constraint on our models. The total luminosity of the observed diffuse emission will also include emission from gas and stellar coronae \\citep{PF1994, Revnivtsev2007} and thus  must be higher than the integrated luminosity of all the LMXBs in our models with luminosities below the observational detection limit.  \\clearpage \\newpage"
    },
    "0801/0801.1587_arXiv.txt": {
        "abstract": "We present a comparison between the 2001 XMM-Newton and 2005 Suzaku observations of the quasar, PG\\,1211+143 at $z=0.0809$. Variability is observed in the 7 keV iron K-shell absorption line (at 7.6 keV in the quasar frame), which is significantly weaker in 2005 than during the 2001 XMM-Newton observation. From a recombination timescale of $<4$ years, this implies an absorber density $n>4\\times10^{3}$\\,cm$^{-3}$, while the absorber column is $5\\times10^{22}<N_{\\rm H}<1\\times10^{24}$\\,cm$^{-2}$. Thus the sizescale of the absorber is too compact (pc scale) and the surface brightness of the dense gas too high (by 9-10 orders of magnitude) to arise from local hot gas, such as the local bubble, group or Warm/Hot Intergalactic Medium (WHIM), as suggested by McKernan et al. (2004, 2005). Instead the iron K-shell absorption must be associated with an AGN outflow with mildly relativistic velocities. Finally we show that the the association of the absorption in PG\\,1211+143 with local hot gas is simply a coincidence, the comparison between the recession and iron K absorber outflow velocities in other AGN does not reveal a one to one kinematic correlation. ",
        "introduction": "The cosmological requirement that half the baryons in the Universe are in a warm/hot intergalactic medium (WHIM) has motivated UV and X--ray absorption line studies to detect this otherwise invisible gas, using absorption against a bright AGN to probe the line of sight material (Bregman 2007). However, most AGN (except blazars) have intrinsic columns of warm absorbing gas in their nuclei, and the fact that this is generally connected to a nuclear outflow (Blustin et al 2005) complicates separating this intrinsic absorption from extrinsic line--of--sight material using velocity information. Nonetheless, various studies have looked for absorption lines (predominantly OVII/VIII) as signatures of the WHIM, though all current detections are more likely to be associated with material in our Galactic halo or Local Group (Bregman \\& Lloyd--Jones 2007). These have velocities within a few hundred km/s of the local standard of rest, and indicate columns of OVII/VIII $> 10^{16}$~cm$^{-2}$, equivalent to Hydrogen columns of $N_H\\sim 10^{19}$~cm$^{-2}$.  Much more controversial was the potential association of substantially higher columns of gas to such local material. McKernan et al (2004; 2005) noticed that several AGN with strong He-- or H--like Fe absorption features had these lines at energies which were approximately consistent with the local standard of rest.  For any single object the approximate match between the putative blueshifted outflow velocity and the galaxy redshift could be coincidental, but McKernan et al (2004; 2005) pointed out three AGN showing this trend, with redshifts spanning $\\sim 0.008-0.15$ (MCG--6--30--15, PG1211+143 and PDS~456). However, the derived columns of $N_H\\sim 10^{23}$~cm$^{-2}$ are high even for an intrinsic AGN outflow (Pounds et al 2003, hereafter P03; Reeves et al 2003). If these are instead of local origin then it requires a tremendously significant change to our understanding of the Galactic halo environment (McKernan et al 2004; 2005).  Here we use new Suzaku data to show that the highly ionised Fe absorption in PG~1211+143 is variable on a timescale of years. This conclusively demonstrates that it is intrinsic to the AGN, and not associated with our Galaxy or Local Group. We collate recent results to show that this is also the case for most other AGN with strong Fe K absorption lines, showing that powerful outflows are associated with luminous accretion flows. ",
        "conclusions": "We show that the observed variability in the absorption line in PG~1211+143 between the XMM-Newton and Suzaku data taken 4 years apart conclusively rules out a diffuse gas origin such as the local Galactic halo or WHIM. The gas must be associated with the AGN, as is further evidenced by its large column density in the XMM-Newton 2001 data which is far too high for local or intergalactic gas. The conclusive identification with the AGN, and the implausibility of any alternative line transition other than iron K$\\alpha$ means that the large outflow velocity of $\\sim 0.1$c is inescapable. Such material is predicted from the winds which are produced from luminous accretion discs in AGN (Proga \\& Kallman 2004).  The coincidence of the outflow velocity with the source redshift in these data is not significant. Other powerful AGN which show iron K absorption features clearly show a range of (intrinsic and observed) velocities, removing the apparent trend for the line energy to appear at the rest energy for these transitions. Thus it is clear that this material is a mildly relativistic outflow from the AGN. Its high velocity means that its kinetic energy can be comparable to the bolometric radiated luminosity of the AGN (King \\& Pounds 2003; Pounds et al 2003; Pounds \\& Reeves 2007), yet its high ionisation means that it is effectively invisible in all other wavebands. While such a major component of AGN energetics could go unnoticed in individual objects, there is increasing evidence that strong AGN feedback controls galaxy formation and evolution. Mildly relativistic winds provide more efficient heating than the jet due to their impact on a larger area, so these highly ionised disc winds may be the key to understanding the growth of structure in the Universe."
    },
    "0801/0801.1064_arXiv.txt": {
        "abstract": "We briefly review capabilities and requirements for future instrumentation in UV- and X-ray astronomy that can contribute to advancing our understanding of the diffuse, highly ionised intergalactic medium. ",
        "introduction": "\\label{Introduction}  A convergence of recent theoretical and observational work has generated rapidly growing interest in the physics of a highly ionised intergalactic medium (IGM). At low redshifts, this phase of the IGM likely holds the balance of the baryon mass density not accounted for by the mass densities in stars, diffuse gas in galaxies, the local Ly$\\alpha$ forest, the Intracluster Medium (ICM) in galaxy clusters and groups, etc. \\citep{FP04}. Model calculations suggest that in fact the baryon density in the IGM could indeed be comparable to the summed baryon densities of the known mass components, as would be required to reach the baryon density determined using independent arguments (big bang nucleosynthesis and the measured abundances of the light elements, and the measured fluctuation spectrum of the cosmic microwave background).  It is clear that the study of non-equilibrium phenomena in the low-redshift diffuse intergalactic medium, both within and outside bound structures, can produce a wealth of information on a large range of important astrophysical problems, ranging from the formation of galaxies to the generation of the first dynamically important magnetic fields. In this paper, we will attempt to summarise the requirements and prospects for significant progress on the observational study of these phenomena, considered in the light of planned or proposed instrumentation. We begin with a slightly abstract discussion of required and desired instrumental capabilities (spectral resolution, sensitivity, etc.), before we consider specific future or proposed instruments. ",
        "conclusions": ""
    },
    "0801/0801.1778_arXiv.txt": {
        "abstract": "{} {Investigation of the dense gas, the outflows and the continuum emission from the massive twin cores NGC6334I and I(N) at high spatial resolution.} {We imaged the region with the Australia Telescope Compact Array (ATCA) at 3.4\\,mm wavelength in continuum as well as CH$_3$CN$(5_K-4_K)$ and HCN(1--0) spectral line emission.} {While the continuum emission in NGC6334I mainly traces the UCH{\\sc ii} region, toward NGC6334I(N) we detect line emission from four of the previously identified dust continuum condensations that are of protostellar or pre-stellar nature. The CH$_3$CN$(5_K-4_K)$ lines are detected in all $K$-components up to energies of 128\\,K above ground toward two protostellar condensations in both regions. We find line-width increasing with increasing $K$ for all sources, which indicates a higher degree of internal motions of the hotter gas probed by these high K-transitions. Toward the main mm and CH$_3$CN source in NGC6334I we identify a velocity gradient approximately perpendicular to the large-scale molecular outflow. This may be interpreted as a signature of an accretion disk, although other scenarios, e.g., an unresolved double source, could produce a similar signature as well. No comparable signature is found toward any of the other sources.  HCN does not trace the dense gas well in this region but it is dominated by the molecular outflows. While the outflow in NGC6334I exhibits a normal Hubble-law like velocity structure, the data are consistent with a precessing outflow close to the plane of the sky for NGC6334I(N). Furthermore, we observe a wide ($\\sim$15.4\\,km\\,s$^{-1}$) HCN absorption line, much broader than the previously observed CH$_3$OH and NH$_3$ absorption lines. Several explanations for the difference are discussed.} {} ",
        "introduction": "The massive twin cores NGC6334I and I(N) at a distance of 1.7\\,kpc in the southern hemisphere \\citep{neckel1978,straw1989} have been subjected to investigations for more than two decades. The two regions are located at the north-eastern end of the much larger molecular cloud/H{\\sc ii} region complex NGC6334 (e.g., \\citealt{rodriguez1982,gezari1982,depree1995,kraemer1999b,sandell2000,carral2002}). The reason why NGC6334I (synonymous with NGC6334F) and NGC6334I(N) are so interesting from a comparison point of view is that they are only separated by approximately 1\\,parsec, hence they share a similar large-scale molecular environment, but they exhibit extremely different characteristics likely because they are at different evolutionary stages.  Both regions have been studied in much detail over the last decades; recent summaries of the past observations can be found, e.g., in \\citet{hunter2006}, \\citet{beuther2007b} or \\citet{rodriguez2007}. Here we just outline their main characteristics. NGC6334I is a prototypical hot molecular core right at the head of a cometary ultracompact H{\\sc ii} (UCH{\\sc ii}) region \\citep{depree1995,kraemer1995}.  It exhibits rich spectral line emission \\citep{mccutcheon2000,thorwirth2003,schilke2006}, a bipolar outflow \\citep{bachiller1990,leurini2006} and H$_2$O, OH, CH$_3$OH class {\\sc ii} and NH$_3$(3,3)/(6,6)/(8,6)/(11,9) maser emission \\citep{moran1980,forster1989,gaume1987,brooks2001,norris1993,caswell1997,walsh1998,beuther2007b,walsh2007}. In contrast to that, up to very recently NGC6334I(N) was considered a typical cold core since no mid-infrared and only faint near-infrared emission was detected \\citep{gezari1982,tapia1996,persi2005}. Furthermore, weak cm continuum and class {\\sc i} and {\\sc ii} CH$_3$OH maser emission was reported \\citep{carral2002,kogan1998,caswell1997,walsh1998}. The spectral line forest is considerably less dense compared to NGC6334I \\citep{thorwirth2003}, however, a few species are stronger toward NGC6334I(N) (\\citealt{sollins2004c}, Walsh et al.~in prep.). In addition, \\citet{megeath1999} report the detection of a molecular outflow in this region as well. In summary, both regions show signs of active star formation, however, the southern region NGC6334I appears to be in a more advanced evolutionary stage than the northern region NGC6334I(N).  To better characterize this intriguing pair of massive star-forming regions, we started a concerted campaign from cm to mm wavelengths with the Australia Telescope Compact Array (ATCA), the Submillimeter Array (SMA) and the Mopra single-dish telescope.  The previous ATCA NH$_3$(1,1) to (6,6) line observations revealed compact warm gas emission from both regions \\citep{beuther2005e,beuther2007b}, and temperatures estimated to exceed 100\\,K. While toward NGC6334I(N) the low energy NH$_3$ lines showed only extended emission, the high energy lines finally revealed compact gas components. The NH$_3$(6,6) line profile from NGC6334I(N) allowed speculation about a potential accretion disk. CH$_3$OH was strong in absorption toward the southern UCH{\\sc ii} region in NGC6334I, indicative of expanding gas.  In the mm continuum emission, \\citet{hunter2006} used the SMA to resolve several mm continuum sources toward both regions (4 in NGC6334I and 7 in NGC6334I(N)).  Furthermore, Hunter et al.~(in prep.) identified an additional SiO outflow in NGC6334I(N) that has its orientation in north-east south-west direction, approximately perpendicular to the one previously reported by \\citet{megeath1999}.  Here we present 3.4\\,mm continuum and HCN/CH$_3$CN spectral line observations obtained with the new 3\\,mm facility at the ATCA. These observations shed light on the outflow and dense gas properties of both regions as well as on the continuum emission from the embedded protostellar objects and the UCH{\\sc ii} region. ",
        "conclusions": "Our 3.4\\,mm continuum and CH$_3$CN/HCN spectral line study of the massive twin cores NGC6334I and I(N) reveals many new insights into that intriguing pair of massive star-forming regions. Both sets of spectral lines as well as the continuum emission are clearly detected toward both targets. While the continuum emission in NGC6334I mainly follows the UCH{\\sc ii} regions and the strongest protostellar 1.4\\,mm peak is detected only at a $\\sim 3\\sigma$ level, in NGC6334I(N), the 3.4\\,mm continuum emission traces four of the previously identified protostellar or pre-stellar condensations.  In both regions, the whole CH$_3$CN$(5_K-4_K)$ $K$-ladder from $K=0$ to 4 is detected toward the strongest protostellar condensations. While the emission is in most cases so optically thick that temperature estimates are prohibited, toward the secondary mm peak in NGC6334I(N) we can estimate a temperature of $170\\pm50$\\,K.  Toward all four detected CH$_3$CN emission sources, we find a correlation between increasing line-width and increasing excitation temperature of the $K$ components. Since increasing excitation temperatures are expected closer to the protostars, this implies more internal motions, e.g., outflow, infall or rotation, the closer one gets to the central protostar. Similar signatures are observed in the CH$_3$CN 2nd moment maps.  To investigate potential rotation, we produced 1st moment maps toward all CH$_3$CN peaks, and fitted the channel peak positions toward mm1 in NGC6334I to effectively increase the spatial resolution. We identify a velocity gradient toward mm1 in NGC6334I that is oriented approximately perpendicular to the known large-scale outflow. This may be interpreted as a signature of a rotating structure, maybe associated with a massive accretion disk. While early rotation-disk claims for that region were on scales of the molecular outflow \\citep{jackson1988}, we are now reaching spatial scales of the order a few hundred AU, much more reasonable for accretion disks (e.g., \\citealt{yorke2002,krumholz2006b}). However, we have to stress that the rotation signatures are not conclusive yet because, e.g., an unresolved double-source could produce similar signatures. Further investigation at higher angular resolution are required to resolve this issue. While higher angular resolution images with the ATCA at 3\\,mm wavelengths are difficult because of decreasing phase stability with increasing baseline length, one may tackle that problem in some highly excited NH$_3$ lines in the 12\\,mm band. Furthermore, ALMA will allow the investigation of this source in much greater depth. Toward the previously suggested disk candidate in NGC6334I(N) we cannot identify a similar velocity gradient.  In contrast to conventional wisdom that HCN traces the dense gas cores, we find it most prominently in the molecular outflows of both massive star-forming regions. The velocity structure of the outflow in NGC6334I is relatively normal and follows the well-known Hubble-law for molecular outflows. In addition to that, we find the broadest HCN line-width toward the main mm continuum peak mm1. In contrast to the previously found elongation of NH$_3$(6,6) maser emission indicating that mm2 could drive the molecular outflow, these data suggest that mm1 harbors the outflow driving source. However, it is also possible that there are two molecular outflows with different properties that are preferentially detected in different tracers (HCN in this work and CO for the previous larger-scale outflow detection at a slightly different position-angle).  The velocity structure of the NGC6334I(N) outflow is more peculiar.  There we find a change between blue- and red-shifted outflow emission on one side of the outflow.  Taking into account the additionally bended morphology of that outflow, a possible explanation for this velocity structure is a precessing outflow close to the plane of the sky.  Furthermore, HCN exhibits a broad absorption feature with a line-width of $\\sim 15.4$\\,km\\,s$^{-1}$ toward the UCH{\\sc ii} region in NGC6334I. Comparing the line-width of the previously observed absorption features of NH$_3$ and CH$_3$OH, which are of the order 2\\,km\\,s$^{-1}$, with the HCN line-width as well as the line-width observed in the ionized gas of $\\sim$32\\,km\\,s$^{-1}$, two explanations are possible to explain the different line-widths. The velocity gradient identified in the ionized gas indicates that it may be influenced by the molecular outflow close by. If the opacity of HCN were very large it traced only the outer gas layers around the UCH{\\sc ii} region and could also be influenced by the outflow. On the other hand, NH$_3$ is optically thin based on the absent absorption in the hyperfine satellite lines. If HCN were optically thin as well, another possible explanation would be based on the different critical densities of the various molecules: In the picture of an expanding UCH{\\sc ii} region, the densities close to the UCH{\\sc ii} region surface should be higher than those further outside.  Hence HCN may trace the gas closer to the expanding UCH{\\sc ii} region and is then much stronger affected by the expansion process that the lower-density regions further out that are traced by NH$_3$ and CH$_3$OH. To differentiate between both models, observations of a rarer HCN isotopologues are required to derive its optical depth."
    },
    "0801/0801.3243_arXiv.txt": {
        "abstract": "The Alpha Magnetic Spectrometer (AMS), whose final version AMS-02 is to be installed on the International Space Station (ISS) for at least 3 years, is a detector designed to measure charged cosmic ray spectra with energies up to the TeV region and with high energy photon detection capability up to a few hundred GeV, using state-of-the art particle identification techniques.  Among several detector subsystems, AMS includes a proximity focusing RICH enabling precise measurements of particle electric charge and velocity. The combination of both these measurements together with the particle rigidity measured on the silicon tracker endows a reliable measurement of the particle mass.  The main topics of the AMS-02 physics program include detailed measurements of the nuclear component of the cosmic-ray spectrum and the search for indirect signatures of dark matter. Mass separation of singly charged particles, and in particular the separation of deuterons and antideuterons from massive backgrounds of protons and antiprotons respectively, is essential in this context. Detailed Monte Carlo simulations of AMS-02 have been used to evaluate the detector's performance for mass separation at different energies.  The obtained results and physics prospects are presented. ",
        "introduction": "The Alpha Magnetic Spectrometer (AMS)\\cite{bib:ams}, whose final version AMS-02 is to be installed on the International Space Station (ISS) for at least 3 years, is a detector designed to study the cosmic ray flux by direct detection of particles above the Earth's atmosphere using state-of-the-art particle identification techniques. AMS-02 is equipped with a superconducting magnet cooled by superfluid helium. The spectrometer is composed of several subdetectors: a Transition Radiation Detector (TRD), a Time-of-Flight (TOF) detector, a Silicon Tracker, Anticoincidence Counters (ACC), a Ring Imaging \\CK\\ (RICH) detector and an Electromagnetic Calorimeter (ECAL). Fig.~\\ref{amsdet} shows a schematic view of the full AMS-02 detector. A preliminary version of the detector, AMS-01, was successfully flown aboard the US space shuttle Discovery in June 1998\\cite{bib:ams01res}.   \\begin{figure}[htb]  \\center  \\vspace{0.2cm}  \\mbox{\\epsfig{file=ams2.eps,width=0.48\\textwidth,clip=}}   \\caption{Exploded view of the AMS-02 detector.\\label{amsdet}}  \\vspace{-0.5cm}  \\end{figure}   The main goals of the AMS-02 experiment are:  \\begin{itemize}  \\item A precise measurement of charged cosmic ray spectra in the rigidity region between \\mbox{$\\sim$ 0.5 GV} and \\mbox{$\\sim$ 2 TV}, and the detection of photons with energies up to a few hundred GeV;  \\item A search for heavy antinuclei ($Z \\ge$ 2), which if discovered would signal the existence of cosmological antimatter;  \\item A search for dark matter constituents by examining possible signatures of their presence in the cosmic ray spectrum.  \\end{itemize}  The long exposure time and large acceptance (0.5 m${}^2\\cdot$sr) of AMS-02 will enable it to collect an unprecedented statistics of more than $10^{10}$ nuclei. ",
        "conclusions": "AMS-02 will provide a major improvement on the current knowledge of cosmic rays. A total statistics of more than 10${}^{10}$ events will be collected during its operation. Detailed simulations have been performed to evaluate the detector's particle identification capabilities, in particular those of the RICH. Simulation results show that the separation of light isotopes is feasible. Using a set of simple cuts based on event data, relative mass resolutions of $\\sim$~2 \\% and rejection factors higher than 10${}^4$ have been attained in D/p separation at energies of a few GeV/nucleon. The separation procedure presented here might be crucial for the identification of an antideuteron flux resulting from neutralino annihilation."
    },
    "0801/0801.1070_arXiv.txt": {
        "abstract": "We use an independent new sample of cataclysmic variables (CVs), constructed by selecting objects for H$\\alpha$ emission, to constrain the properties of the intrinsic CV population.  This sample is restricted to systems that are likely to be non-magnetic and unevolved; it consists of 17 CVs, of which at least 10 have orbital periods above 3~h.  We find that even very generous allowance for selection effects is not sufficient to reconcile the large ratio of short- to long-period CVs predicted by standard CV evolution theory with the observed sample, possibly implying that short-period systems evolve faster than predicted by the disrupted magnetic braking model.  This would require that an angular momentum loss mechanism, besides gravitational radiation, acts on CVs with orbital periods below the period gap.  To bring the model into agreement with observations, the rate of angular momentum loss below the period gap must be increased by a factor of at least 3, unless the model also over estimates the angular momentum loss rate of long-period CVs. ",
        "introduction": "Despite its wide relevance, the evolution of cataclysmic variables (CVs) is still not very well understood.  Much of the problem stems from the fact that strong selection effects act on observed samples of CVs, making it difficult to constrain theory observationally.  The angular momentum loss rate ($-\\dot{J}$) is the crucial ingredient of CV evolution theory.  Angular momentum loss leads to mass transfer from the secondary to the white dwarf.  The relation between the mass transfer rate ($\\dot{M}$) and $\\dot{J}$ depends on the structure of the secondary star, and the reaction of the secondary to mass loss determines the orbital period ($P_{orb}$) evolution of the system.  Population synthesis methods combine the changing $\\dot{M}$ and $P_{orb}$ with a model of the birthrate of CVs to predict the present-day distribution of CVs as a function of $M_1$, $M_2$ (the mass of the white dwarf and secondary, respectively), $\\dot{M}$, and $P_{orb}$.  Only one of these parameters, $P_{orb}$, can practically be measured for a large number of CVs (but see \\citealt{Patterson98} for an indirect method to measure mass ratios).  The orbital period distribution of known CVs is therefore one of the few observational properties of the CV population that can be used to constrain theory.  The other obvious constraint is the observed CV space density.  In comparing both the size and period distribution of the known CV population to theoretical predictions, observational bias must be accounted for.  This requires that observed CV samples be restricted to objects selected in well-defined (and preferably homogeneous) ways.  A long-standing problem is that theory predicts a large population of short-period CVs, consisting mostly of period bouncers (systems that have evolved beyond the period minimum at about 76~min), which is not observed.  The models of \\cite{Kolb93} and \\cite{HowellRappaportPolitano97} both predict that only $\\simeq 1$\\% of CVs are long-period systems, while roughly 70\\% are period bouncers.  Existing observations have already been used to argue that the short-period CV population cannot be as large as predicted (e.g. \\citealt{Patterson98}; \\citealt{PretoriusKniggeKolb07}).  We will do the same here, using a new and independent CV sample.  The CV samples that have been available to date are heavily biased against the intrinsically faint, short-period CVs, mainly because most have bright flux limits.  The new sample considered here is also limited to apparently bright systems.  However, almost all surveys incorporate a second selection cut (most commonly a blue cut) that also discriminates against the discovery of short-period CVs.  The CV sample that we have constructed differs from most existing samples in that the only selection criterion (other than a flux limit) is based on line emission.  The spectra of the majority of CVs show Balmer emission lines, originating mainly in the accretion flow.  The luminosity of CVs is anti-correlated to the equivalent widths (EWs) of their emission lines (\\citealt{Patterson84}; \\citealt{WithamKniggeGansicke06}; we will take EWs of emission lines as positive throughout).  The explanation for this anti-correlation is that intrinsically faint CVs are low-$\\dot{M}$ systems with low density discs, in which recombination line cooling is very efficient.  Therefore, in contrast to other selection criteria (such as blue optical colours and variability), an emission line EW-based selection cut should favour the discovery of intrinsically faint, short-period systems.  We have used the AAO/UKST SuperCOSMOS H$\\alpha$ Survey (SHS) to define a homogeneous sample of CVs, selected on the basis of H$\\alpha$ emission.  Observations of the new CVs were presented in an earlier paper (\\citealt{HalphaI}; hereafter Paper I).  Here we describe the construction and completeness of the sample, examine the observational biases affecting it, and compare it to the predictions of theory. ",
        "conclusions": "We have compared a homogeneous CV sample, selected for H$\\alpha$ emission, to a model CV population based on standard CV evolution theory.  The magnitude limit and Galactic latitude range of the observed sample was modelled in some detail, while conservative assumptions were made to account for the effects of variability and the $\\mathrm{EW}(\\mathrm{H}\\alpha)$-based selection cut.  The model population is inconsistent with the observed sample.  Specifically, the model predicts relatively too many short-period CVs.  This confirms earlier results, based on independent observations.  The reason for the mismatch between the predicted and observed ratio of short- to long-period CVs may be that the theoretical evolutionary time-scale for CVs below the period gap is too long.  A (very simplistic) consideration of the relative numbers of long- and short-period CVs included in the sample indicates that the disrupted magnetic braking model underestimates $-\\dot{J}$ of short-period CVs by a factor of at least 3, assuming that the model is correct for long-period CVs.  Although surveys now in progress will in future provide much better observational constraints on CV evolution theory than can be derived at the moment, it is already clear that the standard magnetic braking model is in need of revision.  Furthermore, it seems that the correct approach to take is to investigate angular momentum loss rates in excess of the gravitational radiation rate in CVs below the period gap."
    },
    "0801/0801.1841_arXiv.txt": {
        "abstract": "We present refined values for the physical parameters of transiting exoplanets, based on a self-consistent and uniform analysis of transit light curves and the observable properties of the host stars. Previously it has been difficult to interpret the ensemble properties of transiting exoplanets, because of the widely different methodologies that have been applied in individual cases. Furthermore, previous studies often ignored an important constraint on the mean stellar density that can be derived directly from the light curve. The main contributions of this work are \\emph{i}) a critical compilation and error assessment of all reported values for the effective temperature and metallicity of the host stars; \\emph{ii}) the application of a consistent methodology and treatment of errors in modeling the transit light curves; and \\emph{iii}) more accurate estimates of the stellar mass and radius based on stellar evolution models, incorporating the photometric constraint on the stellar density.  We use our results to revisit some previously proposed patterns and correlations within the ensemble.  We confirm the mass-period correlation, and we find evidence for a new pattern within the scatter about this correlation: planets around metal-poor stars are more massive than those around metal-rich stars at a given orbital period.  Likewise, we confirm the proposed dichotomy of planets according to their Safronov number, and we find evidence that the systems with small Safronov numbers are more metal-rich on average. Finally, we confirm the trend that led to the suggestion that higher-metallicity stars harbor planets with a greater heavy-element content. ",
        "introduction": "\\label{sec:introduction}  The transiting exoplanets are only a small subset of all the known planets orbiting other stars, but they hold tremendous promise for deepening our understanding of planetary formation, structure, and evolution. Observations of transits and occultations\\footnote{The word \\emph{transit} is sometimes assumed in the field of exoplanet research to be synonymous with \\emph{eclipse}. In reality, it has a more restricted meaning and has long been used in the eclipsing binary field to describe an eclipse of the larger object by the smaller one. The term \\emph{occultation} is used to refer to the passage of the smaller object (in this case, the planet) behind the larger one (the star) \\citep[see, e.g.,][]{Popper:76}. To avoid confusion, we advocate that \\emph{occultation} or \\emph{secondary eclipse} are preferable to neologisms such as ``secondary transit'' or ``anti-transit''.}  (along with the spectroscopic orbit of the host star) not only allow one to measure the mass and radius of the planet, but also provide opportunities to measure the stellar spin-orbit alignment \\citep{Queloz:00, Winn:07a}, the planetary brightness temperature \\citep{Charbonneau:05, Deming:05}, the planetary day-night temperature difference \\citep{Knutson:07}, and even absorption lines of planetary atmospheric constituents \\citep{Charbonneau:02, Vidal-Madjar:04, Tinetti:07}. These and other observations have been accompanied by theoretical progress in modeling the physical processes in the planetary interiors and atmospheres, as well as the planets' interactions with their parent stars. This rapid progress has been stimulated in no small measure by the remarkable diversity of planet characteristics that has been found among the members of the transiting ensemble.  The accuracy and precision with which the properties of the planet can be derived from transit data depend strongly on whatever measurements and assumptions are made regarding the host star. For example, for a given value of the transit depth, the inferred planetary radius scales in proportion to the assumed stellar radius; and for a given spectroscopic orbit of the host star, the inferred planetary mass scales as the two-thirds power of the stellar mass.  In the literature on transiting planets, a wide variety of methods have been used to estimate the radius and mass of the parent star, ranging from simply looking them up in a table of average stellar properties as a function of spectral type, all the way to fitting detailed stellar evolutionary models constrained by the luminosity, effective temperature, and other observations that may be available for the star. As a result, the ensemble of planet properties at our disposal is inhomogeneous, and in many cases the uncertainty of those determinations is dominated by systematic errors in the stellar parameters that are treated differently by different investigators.  This situation is unfortunate because it hinders our ability to gauge the reliability of any patterns that are discerned among the ensemble properties of transiting exoplanets. With 23 systems that have been reported in the literature, this subfield should be poised for the transition from a handful of results to a large and diverse enough sample for meaningful general conclusions to be drawn, but the heterogeneity of reported results is clearly an obstacle.  It may be surprising that we are limited in many cases by our knowledge of the properties of the parent stars; one would think that stellar physics is a ``solved problem'' in comparison to exoplanetary physics. However it must be remembered that the host stars are usually isolated (and therefore no dynamical mass measurement is possible), and that many of the host stars are distant enough that they do not even have measured parallaxes.  Of the 23 cases in the literature, five have \\hip\\ parallaxes \\citep{Perryman:97}.  For those few it is fairly straightforward to estimate the stellar properties, but for the other systems, more indirect methods have been used.  These indirect methods often rely on the value of the stellar surface gravity that is derived by measuring the depths and shapes of gravity-sensitive absorption lines in the stellar spectrum.  This is a notoriously difficult measurement and the result is often strongly correlated with other parameters that affect the spectrum.  Recently, however, \\cite{Sozzetti:07}, Holman et al.~(2007), and others demonstrated that it is possible to do better by using the information about the mean stellar density that is encoded in the transit light curve. This information was typically overlooked prior to these studies.  The study presented in this paper was motivated by the desire for a homogeneous analysis, and by the desire to take advantage of the photometric estimates of the stellar mean density.  We have revisited the determination of the stellar parameters for all of the transiting planets that have been reported in the literature.  We have taken the opportunity to merge all existing measurements of the atmospheric parameters (mainly the effective temperature and metallicity) with the goal of presenting the best possible values.  We have chosen a uniform method for analyzing photometric data and have re-analyzed existing light curves where necessary to provide homogeneity.  Our hope was that by applying these procedures across the board, we and other investigators could view the ensemble properties with greater clarity and uncover any interesting clues the transiting planets might provide us about the origin, structure, and evolution of exoplanets.  This paper is organized as follows. \\S\\,\\ref{sec:stellar} describes the procedures by which we estimated the stellar properties, using the available spectroscopic and photometric datasets, and a particular set of theoretical stellar evolution models. As a check on the models, \\S\\,\\ref{sec:modchecks} compares the results of a subset of our calculations with those derived from a different set of evolutionary models. \\S\\,\\ref{sec:obschecks} investigates alternate ways of estimating the stellar properties. \\S\\,\\ref{sec:gj436} deals with GJ~436, which needs special treatment because the host star has such a lower mass than the other host stars. \\S\\,\\ref{sec:results} presents the final results for the planetary parameters. \\S\\,\\ref{sec:discussion} uses the new results to check on some of the previously proposed correlations among the properties of the transiting ensemble, and \\S\\,\\ref{sec:remarks} provides final remarks. The Appendix lists the data sets and other issues that are particular to each system. ",
        "conclusions": "\\label{sec:discussion}  Many investigators have sought and claimed possible correlations between various stellar and planetary parameters of transiting systems.  In principle such correlations could lead to important insights into the formation, structure, and evolution of exoplanets. A primary motivation for presenting a more complete, accurate, and homogeneous set of these parameters in this work was to facilitate such studies. The relatively large array of properties now available offers the opportunity to find new correlations, or to revisit old ones incorporating additional variables. While it is beyond the scope of the present work to investigate all possible correlations with statistical rigor, in this section we check on three of the most intriguing and potentially important relations that have been proposed.  \\subsection{Planetary mass versus orbital period} \\label{sec:massperiod}  \\cite{Mazeh:05} were the first to point out the apparent correlation between $M_p$ and $P$ for transiting planets \\citep[see also][]{Gaudi:05}. The original suggestion was based on only 6 systems, but additional discoveries have generally supported the trend of decreasing mass with longer periods, although the scatter has also become larger. This is shown in Figure~\\ref{fig:mp}a. HAT-P-2b would be an extreme outlier in this plot; we have excluded it because it is so much more massive than the other planets and may belong to a different category of planet (see \\S\\,\\ref{sec:safronov}). Similarly, we have excluded GJ~436b because it is so much less massive than the others, and may be a rocky or rock-ice planet rather than a gas giant. A simple linear fit is shown for reference (dashed line).   \\begin{figure}        % \\vskip 0.0in \\epsscale{1.3} {\\hskip -0.2in\\plotone{f9xv.ps}} \\vskip 0.0in  \\figcaption[]{(a) $M_p$ as a function of period for all transiting planets except HAT-P-2b ($M_p = 8.7$~M$_{\\rm Jup}$, $P = 5.63$~d) and GJ~436b ($M_p = 0.073$~M$_{\\rm Jup}$, $P = 2.64$~d); see text. The dashed line is a linear fit. (b) $O\\!-\\!C$ residuals $\\Delta M_p$ from the linear fit in the top panel shown as a function of the metallicity of the host star. The dashed line is a linear fit. (c) Same as (a), with the dependence on [Fe/H] removed based on the fit in (b). The fitted line has the expression $M_p = (+1.70 \\pm 0.13) - (0.281 \\pm 0.044) \\times P$. \\label{fig:mp}}  \\end{figure}  We also investigated the scatter in this relation, seeking any ``third variable'' that might correlate with the residuals.  We seem to have found such a third variable: the metallicity of the host star. Figure~\\ref{fig:mp}b displays the $O\\!-\\!C$ residuals from the top panel as a function of [Fe/H], indicating a rather clear correlation (dashed line): $\\Delta M_p = (+0.152 \\pm 0.050) - (1.17 \\pm 0.23) \\times {\\rm [Fe/H]}$. After removal of this trend, the relation between $M_p$ and period becomes tighter (Figure~\\ref{fig:mp}c).  The scatter in the mass-period relation is reduced from 0.26~M$_{\\rm Jup}$ in the top panel to 0.17~M$_{\\rm Jup}$ in the bottom panel. It seems unlikely that this is a statistical fluke. However, the scatter is still larger than the formal observational uncertainties, suggesting that these three variables are not completely determinative.  What might be the implications of this metallicity dependence?  It has been proposed that the mass-period relation is related to the process by which close-in exoplanets migrated inward from their formation sites, or more specifically, to the mechanism that halts migration at orbital periods of a few days. The trend of larger masses at shorter orbital periods could suggest that the halting mechanism depends on mass, and larger planets are able to migrate further in. The dependence on metallicity may then be interpreted to indicate that planets in metal-poor systems need to be more massive in order to migrate inward to the same orbital period as more metal-rich planets. In this context, the trend in Figure~\\ref{fig:mp}b could be interpreted as evidence that the efficiency of the migration (or halting) mechanism is affected to some degree by the chemical composition.  A dependence of migration on metallicity is in fact predicted by some theories, and could arise, as pointed out by \\cite{Sozzetti:06}, either from slower migration rates in metal-poor protoplanetary disks \\citep{Livio:03, Boss:05} or through longer timescales for giant planet formation around metal-poor stars, which would effectively reduce the efficiency of migration before the disk dissipates \\citep{Ida:04, Alibert:05}. However, the above processes are more aimed at addressing the apparent lack of short-period planets among very metal-poor stars claimed by some authors \\citep[see][]{Sozzetti:06}, whereas among the transiting planets it is not the lack of more metal-poor examples we are concerned with, but rather their different properties (such as mass) compared to metal-rich planets \\emph{at the same orbital period}.  To give a quantitative example, we find that for a period of 2.5 days (near the average for known transiting systems), the mass of a planet with [Fe/H] $= -0.2$ is $\\sim$40\\% larger than one with average metallicity ([Fe/H] = +0.13, excluding HAT-P-2b and GJ~436b), while the mass of a planet with [Fe/H] $= +0.4$ is about 30\\% smaller. This range of metallicities, from [Fe/H] $= -0.2$ to +0.4 (covering a factor of 4 in metal enhancement), is approximately the full range observed. We note that the strength of the metallicity effect is modest rather than overwhelming, and this too is in agreement with theoretical expectations \\citep[e.g.,][]{Livio:03}. The massive planet HAT-P-2b obviously does not conform to this trend of $M_p$ versus period, which may indicate some fundamental difference either in its formation or migration.  An alternative interpretation, also proposed by \\cite{Mazeh:05}, is that the $M_p$ versus $P$ relation is more a reflection of survival requirements in close proximity to the star, due to thermal evaporation from the extreme UV flux. Close-in planets must be more massive to avoid ablation to the point of undetectability. The role of metallicity in this case would be through the difference in the internal structure \\citep{Santos:06}. Metal-rich planets have been suggested to be more likely to develop rocky cores \\citep[][see also \\S\\,\\ref{sec:cores}]{Pollack:96, Guillot:06, Burrows:07}. If the presence of such a core somehow slows down or prevents complete evaporation, as has been proposed \\citep{Baraffe:04, Lecavelier:04}, survival at a given period would then have a dependence on metallicity.  \\begin{figure} \\vskip -0.35in \\epsscale{1.3} {\\hskip -0.2in\\plotone{f10.eps}} \\vskip -0.3in  \\figcaption[]{Surface gravity versus orbital period for all transiting planets except HAT-P-2, which is off the scale ($g_p \\sim 234$ m~s$^{-2}$). A linear fit is shown.\\label{fig:loggperiod}}  \\end{figure}  A diagram related to the one considered above is that of planetary surface gravity versus orbital period, first presented by \\cite{Southworth:07}.  Because $g_p$ does not require knowledge of the mass or radius of the star, it is in a sense a cleaner quantity that should be free from systematic errors in $M_{\\star}$ and $R_{\\star}$ and is nearly independent of stellar evolution models. This not true of other bulk properties such as the mean planet density. An updated version of the $g_p$ versus $P$ relation is shown in Figure~\\ref{fig:loggperiod}, along with a linear fit.  We examined the residuals from this linear fit for a possible correlation with stellar metallicity, as we did earlier for the case of the $M_p$ versus $P$ diagram, but we found none.  Given that $g_p \\propto M_p/R_p^2$, we note that the apparent influence of metallicity on the planetary radii must be contributing significantly to the scatter in the $g_p$ versus period relation. We discuss this further in \\S\\,\\ref{sec:cores}.  \\subsection{Safronov number versus equilibrium temperature} \\label{sec:safronov}  Recently, \\cite{Hansen:07} proposed a distinction between two classes of hot Jupiters, based on a consideration of the Safronov number and the zero-albedo equilibrium temperature. The Safronov number is a measure of the ability of a planet to gravitationally scatter other bodies \\citep{Safronov:72}, and is defined as $\\Theta = \\frac{1}{2}(V_{\\rm esc}/V_{\\rm orb})^2 = (a/ R_p)(M_p/ M_{\\star})$, the ratio between the escape velocity and the orbital velocity squared.  We assume that the zero-albedo equilibrium temperature scales as $T_{\\rm eq} = T_{\\rm eff}(R_{\\star}/2a)^{1/2}$ (i.e., we assume that the heat redistribution factor $f$ is common to all planets, in the absence of more complete knowledge).  We list $\\Theta$ and $T_{\\rm eq}$ for all transiting planets in Table~\\ref{tab:planetary}.  \\cite{Hansen:07} pointed out a gap in the distribution of Safronov numbers, and defined Class~I planets as those with $\\Theta \\sim 0.07 \\pm 0.01$ and Class~II as $\\Theta \\sim 0.04 \\pm 0.01$.  They tentatively proposed also that these two categories have other distinguishing characteristics, such as a difference in the average temperature of the host stars, or the orbital separations. Upon the discovery of HAT-P-5b \\citep{Bakos:07c} and HAT-P-6b \\citep{Noyes:07}, those authors pointed out that the distinction now seems less clear, as these two planets tend to fill the gap between Class~I and Class~II.  The larger sample now available, and especially the more accurate and homogeneous set of properties presented here, offers the opportunity to revisit the issue. Figure~\\ref{fig:safronov} shows an updated version of the $\\Theta$--$T_{\\rm eq}$ diagram for transiting planets, which has some significant differences compared to the original version. Following \\cite{Hansen:07} we have excluded the massive planet HAT-P-2b, with a Safronov number so much larger than all the others ($\\Theta = 0.94$) that it would seem to be in a different class altogether, as well as GJ~436b, which is of much lower mass and a presumably different composition.  \\begin{figure} \\vskip -0.35in \\epsscale{1.3} {\\hskip -0.2in\\plotone{f11.eps}} \\vskip -0.3in  \\figcaption[]{Diagram of Safronov number ($\\Theta$) versus equilibrium temperature ($T_{\\rm eq}$) for transiting planets. Planet classes are labeled following \\cite{Hansen:07}. A tentative dividing line at $\\Theta = 0.05$ is indicated. \\label{fig:safronov}}  \\end{figure}  In our updated diagram the separation between Class~I and Class~II is still quite striking. The clustering of the Class~II objects around the value $\\Theta \\sim 0.04$ has tightened, if anything, and HAT-P-5b (\\#21) and HAT-P-6b (\\#22) do not encroach on the gap in Safronov numbers with Class~I. (The Class~II object with the lowest value of $\\Theta$ is OGLE-TR-111b [\\#6].) Thus, the dichotomy remains very suggestive.  \\cite{Hansen:07} have argued that the principal distinction between the two classes is based on mass (planetary and/or stellar), and that planets of Class~II are, on average, less massive than those in Class~I, and orbit stars that are typically more massive.  While we agree with the latter part of this statement regarding the \\emph{stellar} masses, we find that the distribution of \\emph{planetary} masses is indistinguishable between the two groups. We do confirm, however, that an updated diagram of $M_p$ versus $T_{\\rm eq}$ (see Figure~\\ref{fig:masstemp}) shows planets in Class~II to be systematically less massive \\emph{for the same equilibrium temperature} (level of irradiation), as was also found by \\cite{Hansen:07}, so in this sense their general claim that the distinction has something to do with mass seems to be supported by the observations.  \\begin{figure} \\vskip -0.35in \\epsscale{1.3} {\\hskip -0.2in\\plotone{f12.eps}} \\vskip -0.3in  \\figcaption[]{Planetary mass as a function of equilibrium temperature for transiting planets. Class~II planets ($\\Theta < 0.05$) are represented with squares. A tentative dividing line is indicated. \\label{fig:masstemp}}  \\end{figure}  An important characteristic that seems to be different in the two groups is the metallicity of the parent stars. This is illustrated in Figure~\\ref{fig:saffeh}. Parent stars of Class~II planets tend to be slightly more metal-rich. A Kolmogorov-Smirnov test of the two metallicity distributions, which appear to be centered around [Fe/H] $\\sim 0.0$ for Class~I and [Fe/H] $\\sim +0.2$ for Class~II, indicates only a 1.7\\% probability that they are drawn from the same parent population. There is perhaps a hint that the Safronov numbers for Class~I planets show a decreasing trend with metallicity in Figure~\\ref{fig:saffeh}, whereas the $\\Theta$ values for Class~II planets are independent of [Fe/H].  \\begin{figure} \\vskip -0.3in \\epsscale{1.3} {\\hskip -0.2in\\plotone{f13.eps}} \\vskip -0.3in  \\figcaption[]{Safronov number as a function of stellar metallicity for transiting planets. Class~II planets ($\\Theta < 0.05$) are represented with squares.\\label{fig:saffeh}}  \\end{figure}  \\subsection{Heavy element content versus stellar metallicity} \\label{sec:cores}  Transiting planet discoveries have challenged our theoretical understanding of these objects from the very beginning. The inflated radius of HD~209458b, argued to be due to an overlooked internal heat source in the planet \\citep[see, e.g.,][and references therein]{Fabrycky:07}, still poses a problem for modelers, and this first example is now joined by several other oversized planets indicating it is not an exception (see \\S\\,\\ref{sec:results}).  On the other end of the scale, HD~149026b was the first transiting giant planet found to have a radius significantly \\emph{smaller} than predicted by standard theories \\citep{Sato:05}. The implication is that it must have a substantial fraction of heavy elements of perhaps $\\sim$70~M$_{\\earth}$ (2/3 of its total mass), which is often assumed to be in a core.  More recently HAT-P-3b \\citep{Torres:07a} has also been found to be enhanced in heavy elements on the basis of its small size for the measured planet mass. In this case metals make up $\\sim$1/3 of the total mass.  \\cite{Guillot:06} and also \\cite{Burrows:07} have found evidence that the heavy element content correlates with the metallicity of the parent star, a trend that was not anticipated by theory. These studies were based, respectively, on the 9 and 14 transiting planets known at the time. The larger sample now available warrants a second look at this possible correlation, for its potential importance for our understanding of planet formation.  Here we have estimated the heavy element content $M_{\\rm Z}$ of each planet using the recent models by \\cite{Fortney:07}, which include the effects of irradiation from the central star. The mass in heavy elements was calculated from the measured mass of the planet, its radius, the orbital semimajor axis, and the age of the system as derived above from stellar evolution models. The uncertainty in these estimates is often very large, due mostly to errors in the measured values of $R_p$ and especially the age. In ten cases the models indicate no metal enhancement at all, and some of these are in fact ``inflated'' hot Jupiters. In a few other cases only an upper limit can be placed. For HAT-P-2b the \\cite{Fortney:07} models as published do not provide a large enough range of core masses.  The results for $M_{\\rm Z}$ are listed in Table~\\ref{tab:planetary}.  A comparison with values from \\cite{Guillot:06} and \\cite{Burrows:07} for the subset of planets in common indicates significant differences in some cases, either in $M_{\\rm Z}$ or its error.  \\begin{figure} \\vskip -0.35in \\epsscale{1.3} {\\hskip -0.2in\\plotone{f14.eps}} \\vskip -0.3in  \\figcaption[]{Heavy element content for transiting giant planets as a function of the metallicity of the host star. HAT-P-2 is excluded (see text).\\label{fig:cores}}  \\end{figure}  The mass in heavy elements is plotted against the stellar metallicity in Figure~\\ref{fig:cores}. The overall trend is similar to that pointed out by previous authors, in the sense that the upper envelope of the distribution appears to increase with [Fe/H]. It is natural to expect that a higher stellar metallicity implies a higher-metallicity protoplanetary disk. In the context of the core-accretion model of planet formation, one would also naturally expect a more metal-rich disk to lead to more metal-rich planets, and in this sense the observed trend is in accordance with the core-accretion theory."
    },
    "0801/0801.3596_arXiv.txt": {
        "abstract": "This article investigates the predictions of an inflationary phase starting from a homogeneous and anisotropic universe of the Bianchi~$I$ type. After discussing the evolution of the background spacetime, focusing on the number of $e$-folds and the isotropization, we solve the perturbation equations and predict the power spectra of the curvature perturbations and gravity waves at the end of inflation.  The main features of the early anisotropic phase is (1) a dependence of the spectra on the direction of the modes, (2) a coupling between curvature perturbations and gravity waves, and (3) the fact that the two gravity waves polarisations do not share the same spectrum on large scales. All these effects are significant only on  large scales and die out on small scales where isotropy is recovered. They depend on a characteristic scale that can, but a priori must not, be tuned to some observable scale.  To fix the initial conditions, we propose a procedure that generalises the one standardly used in inflation but that takes into account the fact that the WKB regime is violated at early times when the shear dominates. We stress that there exist modes that do not satisfy the WKB condition during the shear-dominated regime and for which the amplitude at the end of inflation depends on unknown initial conditions. On such scales, inflation loses its predictability.  This study paves the way to the determination of the cosmological signature of a primordial shear, whatever the Bianchi~$I$ spacetime. It thus stresses the importance of the WKB regime to draw inflationary predictions and demonstrates that when the number of $e$-folds is large enough, the predictions converge toward those of inflation in a Friedmann-Lema\\^{\\i}tre spacetime but that they are less robust in the case of an inflationary era with a small number of $e$-folds. ",
        "introduction": "Inflation~\\cite{lindebook,pubook} (see Ref.~\\cite{linde2007} for a recent review of its status) is now one of the cornerstones of the standard cosmological model. In its simplest form, inflation has very definite predictions: the existence of adiabatic initial scalar perturbations and gravitational waves, both with Gaussian statistics and an almost scale invariant power spectrum~\\cite{ref:inf,mbf}. Various extensions, which in general involve more fields, allow e.g. for isocurvature perturbations~\\cite{iso}, non-Gaussianity~\\cite{ng}, and modulated fluctuations~\\cite{bku}. All these features let us hope that future data will shed some light on the details (and physics) of this primordial phase of the universe.  Almost the entire literature on inflation assumes that the universe is homogeneous and isotropic while homogeneity and isotropy are what inflation is supposed to explain. Indeed, the dynamics of anisotropic inflationary universes has been widely discussed~\\cite{olive}. It was demonstrated that under a large variety of conditions, inflation occurs even if the spacetime is initially anisotropic~\\cite{infanisogen}, regardless of whether it is dominated by a pure cosmological constant or a slow-rolling scalar field. The isotropization of the universe was even generalized to Bianchi braneworld models~\\cite{braneinf} recently. It should be emphasized, however, that a deviation from isotropy~\\cite{infanisogen} or flatness~\\cite{inflK} may have a strong effect on the dynamics of inflation, and in particular on the number of $e$-folds.\\\\  The study of the perturbations during the isotropization phase has been overlooked, mainly because in the past decades one was mostly focused on large field inflationary models, which generically give very long inflationary phases. This was also backed up by the ideas of chaotic inflation and eternal inflation~\\cite{linde2007}. In such cases, it is thus an excellent approximation to describe the universe by a Friedmann-Lema\\^{\\i}tre (FL) spacetime when focusing on the modes observable today since they exited the Hubble radius approximatly in the last 60 $e$-folds. In this case, the origin of the density perturbations is understood as the amplification of vacuum quantum fluctuations of the inflaton. In particular, the degrees of freedom that should be quantized deep in the inflationary phase, and known as the Mukhanov-Sasaki variables~\\cite{MSvar}, were identified~\\cite{mbf}, which completely fixes the initial conditions and makes inflation a very predictive theory.  In the context of string theory, constructing a string compactification whose low energy effective Lagrangian is able to produce inflation is challenging (see Ref.~\\cite{Allister}). In particular, it has proved to be difficult to build large field models~\\cite{bauman} (in which the inflaton moves over large distances compared to the Planck scale in field space). This has led to the idea that, in this framework, an inflationary phase with a small number of $e$-folds is favored (see however Ref.~\\cite{roulette}). If so, the predictions of inflation are expected to be sensitive to the initial conditions, and in particular the classical inhomogeneities are not expected to be exponentially suppressed, which makes the search for large scale deviations from homogeneity and isotropy much more motivated, as well as the possibility that the inflaton has not reached the inflationary attractor. More important, it is far from obvious (as we shall discuss in details later) that all observable modes can be assumed to be in their Bunch-Davies vacuum initially, and more puzzling, that for a given mode modulus the possibility of setting the initial conditions will depend on its direction. If so, then the initial conditions for these modes would have to be set in the stringy phase, an open issue at the time.  The theory of cosmological perturbations in a Bianchi universe was roughed out in Ref.~\\cite{TomitaDen,Dunsby,pitrou07} (see also Ref.~\\cite{NohHwang} and Ref.~\\cite{Abbotetal} for the case of higher-dimensional Kaluza-Klein models and Ref.~\\cite{Qaniso} for the quantization of test fields and particle production in an anisotropic spacetime). Recently, we performed a full analysis of the cosmological perturbations in an arbitrary Bianchi~$I$ universe~\\cite{ppu1}. It was soon followed by an analysis~\\cite{GCP} that focused on Bianchi~$I$ universe with a planar symmetry. As we shall see in this work, the case of a planar symmetric spacetime is not generic (both for the dynamics of the background and the evolution of the perturbations).\\\\  From a more observationally oriented perspective, the primordial anisotropy can imprint a preferred direction in the primordial power spectra. This could be related to the possible large scale statistical anomalies~\\cite{lowQ} of the cosmic microwave background (CMB) anisotropies. Many possible explanations have been proposed, including foregrounds~\\cite{prunet}, non-trivial spatial topology~\\cite{topologie} (which implies a violation of global isotropy~\\cite{modetopo}), the breakdown of local isotropy due to multiple scalar fields~\\cite{picon}, the presence of spinors~\\cite{spinor} or dynamical vectors \\cite{dulaney}, the effect of the spatial gradient of the inflaton~\\cite{donoghue} or a late time violation of isotropy~\\cite{Jaffe2005}. In the case of universes with planar symmetry, the signatures on the CMB were derived by many authors~\\cite{GCP,CMB_bianchi,acker}. In this case, a large primordial shear is indeed necessary, which is not in contradiction with the constraints obtained from the CMB~\\cite{LimitShear} or from the big-bang nucleosynthesis~\\cite{bbn}. This brings a secondary motivation to our analysis: can the CMB anisotropy be related to an anisotropic primordial phase or, on the other hand, can it constrain the primordial shear and what are the exact predictions of a primordial  anisotropic phase? \\\\  This is however not our primary motivation. In the first place, we are interested in understanding the genericity of the predictions of inflation, and in particular with respect to the symmetries of the background spacetime. As we shall see, the simple extension considered in this article drives a lot of questions, concerning both the initial conditions in inflation and, more generally, quantum field theory in curved spacetime.  In this article, we build on our previous work~\\cite{ppu1} to investigate the dynamics and predictions of one field inflation starting from a generic Bianchi~$I$ universe. After discussing the dynamics of the background in \\S~\\ref{sec1}, we focus on the evolution of the perturbations in \\S~\\ref{sec2}. In particular, we will need to understand the procedure for the quantization during inflation. As we shall see in~\\S~\\ref{sec2.2}, this procedure deviates from the one standardly used in a Friedmann-Lema\\^{\\i}tre spacetime, and even in a planar symmetric spacetime as discussed in Ref.~\\cite{GCP}. The main reason for such an extension lies in the fact that there always exist modes that were not in a WKB regime during the shear-dominated inflationary era. It follows that, while our procedure leads to similar predictions to the standard one on small scales, it appears that there is a lack of predictibility on large scales. Indeed, we do not want to push our description beyond the Planck or string scale where extensions of general relativity have to be considered. This may give a description of both the early phase of the inflationary era and also a procedure to fix the initial conditions without any ambiguity. In \\S~\\ref{sec3}, we explicitly compute the primordial spectra, for both scalar modes and gravity waves. We will describe in details the effect of the anisotropy and show how the isotropic predictions are recovered on small scales. ",
        "conclusions": "\\label{sec4}  In this article, we have worked out the predictions of an anisotropic inflationary era for a generic Bianchi~$I$ spacetime. We have discussed the isotropization both at the background level and at the linear order in perturbation theory, both for scalar modes and gravity waves.  Generically these spacetimes always enjoy a bouncing direction, apart from the particular case $\\alpha=\\pi/2$ considered in Ref.~\\cite{GCP}. (Note that the predictions for this case are in fact singular among the predictions since they do not converge uniformly when $\\alpha\\rightarrow\\pi/2$).  Since at early time, the modes are not in a WKB regime, we had to extend the standard procedure to fix the initial conditions (see \\S~\\ref{sec2.2.c}). We showed that the modes larger than $\\kref$ always enter a WKB regime before they become super-Hubble during inflation. As we discussed, our procedure reproduces the standard one on small scales.  In the particular case where we tune the initial shear so that these initial conditions can be set unambiguously while still having an imprint of the CMB anisotropy, we presented the imprint of the primordial anisotropy in the power spectra of the gravity waves and curvature perturbation at the end of inflation. Two examples were studied but we can provide predictions for any Bianchi~$I$ universe. Note that these predictions were drawn by assuming that the slow-roll attractors were reached before the time the modes of observational relevance had exited the horizon. In the case of inflation with small number of $e$-folds before that time, it is not clear that this is a realistic hypothesis. If so, one would have to make the predictions in terms of trajectories, that is predictions that will depend on the initial conditions of the scalar field. This is not specific to Bianchi universes but to all models in which the inflationary period is short~\\cite{levprivate}.\\\\  Concerning the initial conditions, two problems arise (see discussion in \\S~\\ref{sec2.2.e}). First, there exists an early shear-dominated phase where the WKB approximation is violated. This forbids us to set the initial conditions as in a Friedmann-Lema\\^{\\i}tre universe. As we showed, the sub-Hubble modes at the onset of the accelerating phase can be quantized because quantum fluctuations act at all times and can always source oscillatory solutions, when they exist. On the other hand, there always exist non-oscillating modes. These modes are expected not to be much smaller than the Hubble radius today (since there is no trivial imprint of the anisotropy on the CMB). This implies that there exists a cut-off scale above which we cannot predict the spectra from first principles, at least in the theoretical set-up we are considering. Consequently, we conclude that above this scale we can only measure the power spectra or postulate their functional form, as was actually done to set the initial conditions before the invention of inflation. We have shown that even if unknown pre-WKB initial conditions can grow, this growth is at most of order unity. Therefore we can safely assume that the perturbations at the end of inflation reflect only those modes that have been seeded during the WKB regime.  Second, for these modes, one cannot assume that they are independent. It implies that we have to consider 3 interacting fields, which also complicates the quantization procedure. Such an issue was addressed perturbatively in the interaction picture in the case of self-interacting field in order to estimate the non-Gaussianity~\\cite{interacting} but no general formalism has been designed when no interaction-free regime can be exhibited.\\\\  In our analysis we have assumed the validity of general relativity up to the singularity. Indeed, we do not take this model for more than what it actually is, e.g. we cannot extrapolate it beyond the Planck or the string scale. There, more degrees of freedom have to be included and can change the dynamics. This could introduce an early chaotic phase~\\cite{billiard} since Bianchi~$I$ models coupled to $p$-form fields have never ending oscillatory behaviour exhibited by generic string theory. Examples of an early dynamics have also been given in terms of Kalb-Ramond axion fields~\\cite{kaloper}, non-commutative geometry~\\cite{ncg} and recently loop quantum gravity~\\cite{lcq}. Any of these developments may give a description of both the early phase and a procedure to fix the initial conditions.\\\\  Coming back to inflation, our study demonstrates to which extent its predictions are sensitive to initial (classical) large scale anisotropies and that in the presence of a non-vanishing shear, it is impossible to define a Bunch-Davies vacuum in the standard way (see also Ref.~\\cite{picon2} for a discussion of this issue in the standard picture and the arbitrariness on the choice of the initial state and its influence of the prediction of inflation). Indeed, if we tune the initial conditions such that the number of $e$-folds is large, none of the problems we address here will affect the observable modes, simply because only modes such that $k/\\kref\\gg1$ are observable and we have shown that in this limit we recover the standard inflationary predictions. If this is the case, one would have no observational imprint of the primordial shear. But our analysis and conclusions may be of some relevance for inflationary model building in the framework of string theory if the feeling that no large field model (and thus no large number of $e$-folds) can be constructed persists, an issue far beyond the scope of this article.  Our analysis also shows the importance (and peculiarity) of the Friedmann-Lema\\^{\\i}tre background in our theoretical predictions and on the quantization procedure, and it gives as well an explicit construction of the difficulties encountered when these symmetries do not exist. We showed that if the number of $e$-folds is large the inflationary predictions converge toward the isotropic predictions, hence demonstrating that they are robust in that regime. In the case of a small number of $e$-folds, they are very sensitive to the initial shear but also the theoretical construction is less under control.  If the indication of the breakdown of statistical isotropy from the CMB were to be confirmed and related to such an early anisotropic phase, then a new coincidence will appear in the cosmological models since one would need to understand why the characteristic scale is of the order of the Hubble radius today, $\\kref\\sim k_0$.  From a more pragmatic attitude, the present work allows us to draw the CMB signatures of such an anisotropic early phase, for any Bianchi~$I$ universe. We stress that the general form of the initial power spectra are more general than those heuristically considered in the analysis performed in Refs.~\\cite{acker,kam} and go beyond the one derived for the singular case $\\alpha=\\pi/2$ studied in Ref.~\\cite{GCP}. The existence of correlation between gravity waves and curvature perturbation, as well as the fact that the two gravity wave polarisations do not share the same spectrum, may also lead to specific signatures to be investigated."
    },
    "0801/0801.3075_arXiv.txt": {
        "abstract": "Circumnuclear star forming regions, also called hotspots, are often found in the inner regions of some spiral galaxies where intense processes of star formation are taking place. In the UV, massive stars dominate the observed circumnuclear emission even in the presence of an active nucleus, contributing between 30 and 50 \\% to the H$\\beta$ total emission of the nuclear zone. Spectrophotometric data of moderate resolution ( 3000 $<$ R $<$ 11000) are presented from which the physical properties of the ionized gas: electron density, oxygen abundances, ionization structure etc. have been derived. ",
        "introduction": "The inner ($\\sim$ 1Kpc) zones of some spiral galaxies show high star formation rates frequently arranged in a ring or pseudo-ring pattern around their nuclei. In general, Circumnuclear Star Forming Regions (CNSFR) -- also referred to as ``hotspots'' -- and luminous and large disk HII regions are very much alike, but the first look more compact and show higher peak surface brightness  (Kennicutt et al. 1989).  In many cases they contribute substantially to the emission of the entire nuclear region. Their large H$\\alpha$ luminosities, typically higher than 10$^{39}$ erg s$^{-1}$, point to relatively massive star clusters as their ionisation source. % These regions then constitute excellent places to study how star formation proceeds in high metallicity, high density circumnuclear environments.  The importance of an accurate determination of the abundances of high metallicity HII regions cannot be overestimated since they constitute most of the HII regions in early spiral galaxies (Sa to Sbc) and the inner regions of most late  type ones (Sc to Sd) without  which our description of the metallicity distribution in galaxies cannot be  complete.  The question of how high is the highest oxygen abundance in the gaseous phase of galaxies is still standing and extrapolation of known radial abundance gradients would point to CNSFR as the most probable sites for these high metallicities. Accurate measurements of elemental abundances of high metallicity regions are therefore crucial to obtain reliable calibrations of empirical abundance estimators, widely used but poorly constraint, whose choice can severely bias results obtained for quantities of the highest relevance for the study of galactic evolution like the luminosity-metallicity (L-Z) relation for galaxies. CNSFR are also ideal cases for studying the behaviour of abundance estimators in the high metallicity regime. ",
        "conclusions": "Electron densities for each observed region have been derived from the  [SII] $\\lambda\\lambda$ 6717, 6731 \\AA\\ line ratio, following standard methods. They were found to be, in all cases, $\\le$ 600 cm$^{-3}$, higher than those usually derived in disk HII regions, but still below the critical value for collisional de-excitation. The low excitation of the regions, as evidenced by the weakness of the [OIII] $\\lambda$ 5007 \\AA\\ line (see the blue spectrum in Figure 1, precludes the detection and measurement of the auroral [OIII] $\\lambda$ 4363 \\AA\\ necessary for the derivation of the electron temperature. We have therefore used a semi-empirical procedure for the derivation of abundances. Firstly we have produced a calibration of the [SIII] temperature as a function of the abundance parameter SO$_{23}$ defined as: \\[ SO_{23} = \\frac{I([OII]\\lambda 3727,29)+I([OIII]\\lambda 4959,5007)}{I([SII]\\lambda 6716,31)+I([SIII]\\lambda 9069,9532)} \\] This parameter is similar to the S$_3$O$_3$ proposed by Stasi\\'nska (2006) but is, at first order, independent from geometrical (ionization parameter) effects. To perform this calibration we have compiled all the data so far available with sulfur emission line detections of both the auroral and nebular lines at $\\lambda$ 6312 \\AA\\ and $\\lambda\\lambda$ 9069,9532 \\AA\\ respectively (see D\\'\\i az et al. 2007). The calibration is shown in Figure 2 together with the quadratic fit to the high metallicity HII region data: \\[ t_e([SIII]) = 0.596 - 0.283 log SO_{23} + 0.199 (log SO_{23})^2 \\] For one of the observed regions, R1+R2, the [SIII] $\\lambda$ 6312 \\AA\\  line has been detected and measured yielding a value T$_e$([SIII])= 8400$^{+ 4650}_{-1250}$K, represented as a solid black circle in Figure 2 (left). We have used this calibration to derive t$_{e}$([SIII]) for our observed CNSFR. In all cases the values of log$O_{23}$ are inside the range used in performing the calibration, thus requiring no extrapolation of the fit.These temperatures, in turn, have been used to derive the S$^+$/H$^+$ and S$^{++}$/H$^+$ ionic ratios. Once the sulphur ionic abundances have been derived, we have estimated the corresponding oxygen abundances, assuming that sulphur and oxygen temperatures follow the relation given by Garnett (1992) and t$_e$ ([OII]) $\\simeq$ t$_e$([SIII]).  Finally, we have derived the N$^{+}$/O$^{+}$ ratio assuming that t$_e$([OII]) $\\simeq$ t$_e$([NII]) $\\simeq$ t$_e$([SIII]).   \\begin{figure}[!h] \\plottwo{TeSIII_cal_new.eps}{eta-prima-plot.eps} \\caption{{\\itshape Left:\\/}Empirical calibration of the [SIII] electron temperature as a function of the abundance parameter SO$_{23}$. The solid line represents a quadratic fit to the high metallicity HII region data. {\\itshape Right:\\/} The [OII]/[OIII] vs [SII]/[SIII] logarithmic ratios for different ionised regions. References for the data in the two figures can be found in D\\'\\i az et al. (2007)} \\end{figure}   The observed CNSFR being of high metallicity show however marked differences with respect to high metallicity disk HII regions. Even though their derived oxygen and sulphur abundances are similar, they show values of the O$_{23}$ and the N2 parameters whose distributions are shifted to lower and higher values respectively with respect to the high metallicity disk sample. Hence, if pure empirical methods were used to estimate the oxygen abundances for these regions, higher values would in principle be obtained. This would seem to be in agreement with the fact that CNSFR, when compared to the disk high metallicity regions, show the highest [NII]/[OII] ratios.  CNSFR also show lower ionization parameters than their disk counterparts, as derived from the [SII]/[SIII]  ratio. Their ionization structure also seems to be different with CNSFR showing radiation field properties more similar to HII galaxies than to disk high metallicity HII regions. This can be seen in Figure 2 (right), a diagram of the emission line ratios [OII]/[OIII] vs [SII]/[SIII], that works as a diagnostics for the nature and temperature of the radiation field. In this plot, diagonal lines of slope unity show the locus of ionized regions with constant ionization temperature. CNSFRs are seen to segregate from disk HII regions. The former cluster around a value of T$_{ion}\\sim$ 40,000 K, while the latter cluster around T$_{ion}\\sim$ 35,000 K. Also shown are the data corresponding to HII galaxies. Indeed, CNSFRs seem to share more the locus of HII galaxies than that of disk HII regions.  One possible concern about these CNSFR is that, given their proximity to the galactic nuclei, they could be affected by hard radiation coming from a low luminosity AGN. Alternatively, the spectra of  these regions harbouring massive clusters of young stars might be affected by the presence of shocked gas. Diagnostic diagramas of the kind presented by Baldwin, Phillips, \\& Terlevich (1981) can be used to investigate the possible contribution by either a hidden AGN or the presence of shocks to the emission line spectra of the observed CNSFR. When plotted on one of these diagnostic diagramas, log ([NII]/H$\\alpha$) vs log ([OIII])/H$\\beta$, some of our CNSFR are found close to the transition zone between HII region and LINER spectra but only one region in NGC~3504 may show a hint of a slight contamination by shocks."
    },
    "0801/0801.3305_arXiv.txt": {
        "abstract": "We investigate the accretion of angular momentum onto a protoplanet system using three-dimensional hydrodynamical simulations. We consider a local region around a protoplanet in a protoplanetary disk with sufficient spatial resolution. We describe the structure of the gas flow onto and around the protoplanet in detail. We find that the gas flows onto the protoplanet system in the vertical direction crossing the shock front near the Hill radius of the protoplanet, which is qualitatively different from the picture established by two-dimensional simulations. The specific angular momentum of the gas accreted by the protoplanet system increases with the protoplanet mass. At Jovian orbit, when the protoplanet mass $M_{\\rm p}$ is $M_{\\rm p}\\lesssim 1\\mj$, where $\\mj$ is Jovian mass, the specific angular momentum increases as $j\\propto M_{\\rm p}$. On the other hand, it increases as $j\\propto M_{\\rm p}^{2/3}$ when the protoplanet mass is $M_{\\rm p}\\gtrsim 1\\mj$. The stronger dependence of the specific angular momentum on the protoplanet mass for $M_{\\rm p}\\lesssim 1\\mj$ is due to thermal pressure of the gas. The estimated total angular momentum of a system of a gas giant planet and a circumplanetary disk is two-orders of magnitude larger than those of the present gas giant planets in the solar system. A large fraction of the total angular momentum contributes to the formation of the circumplanetary disk. We also discuss the satellite formation from the circumplanetary disk. ",
        "introduction": "Until now, more than 200 extrasolar planets (or exoplanets) have been detected mainly by measuring the radial motion of their parent star along the line of sight. Almost all exoplanets observed by this method are giant planets, like Jupiter and Saturn in our solar system, because massive planets are preferentially observed. Although these planets are supposed to be formed in the disk surrounding the central star (i.e., circumstellar disk or protoplanetary disk), their formation process has not been fully understood yet. In the core accretion scenario \\citep{hayashi85}, a solid core or protoplanet with $\\simeq 10\\me$, where $\\me$ is the Earth mass, captures a massive gas envelope from the protoplanetary disk by self-gravity to become a gas giant planet.  The evolution of the gaseous protoplanet has been studied with the approximation of  spherical symmetry  including radiative transfer \\citep[e.g.,][]{mizuno80,bodenheimer86,pollack96,ikoma00}. \\citet{ikoma00} showed that rapid gas accretion is triggered when the solid core mass exceeds $\\simeq 5-20\\me$, and the protoplanet quickly increases their mass by gas accretion. However, the angular momentum of the accreting gas was ignored in these studies, because they assumed spherical symmetry. Since the gas accretes onto the solid core with a certain amount of the angular momentum, a circumplanetary disk forms around the protoplanet, analogously to the formation of a protoplanetary disk around a protostar. The difference in the disk formation between the protostar and protoplanet is the region from which the central object acquires the angular momentum. The protostar acquires the angular momentum from a parent cloud, while the protoplanet acquires it from the shearing motion in the protoplanetary (circumstellar) disk. In addition, the gravitational sphere of the protostar spreads almost infinitely, while the gravitational sphere (i.e., the Hill sphere) of the protoplanet is limited in the region around the protoplanet because the gravity of the central star exceeds that of the protoplanet outside the Hill sphere.  The numerical simulations are useful to investigate the gas accretion onto a protoplanet and its circumplanetary disk (hereafter we just call them as a protoplanet system). \\citet{korycansky91} studied giant planet formation using one-dimensional quasi-spherical approximation with angular momentum transfer. They showed that as the protoplanet contracts, outer layers of the envelope containing sufficient specific angular momentum remain in bound orbit, and form a circumplanetary disk. However, owing to the spherical symmetry, the accretion flow from the protoplanetary disk to the protoplanet could not be investigated in their study. To study the accretion flow onto a protoplanet and the acquisition process of the angular momentum in detail, a multidimensional simulation is necessary. \\citet{sekiya87} investigated the gas flow around a protoplanet with relatively low resolution, and found that the spin rotation vector of the protoplanet becomes parallel to the orbital rotation vector. Recently, the flow pattern was carefully investigated by many authors \\citep{miyoshi99, lubow99, kley01, dangelo02, dangelo03}. However, since the main purpose of these studies was to clarify the planet migration process in a large scale (i.e., outside the Hill radius), they did not investigate the region near the protoplanet (i.e., inside the Hill radius) with sufficient resolution. Thus, in their simulations, we cannot study the gas stream inside the Hill radius.  To investigate the accretion flow onto the protoplanet system, we need to cover a large spatial scale from the region far from the Hill sphere to that in the proximity to the protoplanet. For Jupiter, since the Hill radius is $\\rh = 744\\,\\rp$, where $\\rp$ is Jovian radius, we have to resolve at least  $\\sim$1000 times different scales. To cover a large dynamical range in scales, a few authors used the nested-grid method. \\citet{dangelo02,dangelo03} investigated the relation between the spiral patterns within the Hill radius and migration rate using three-dimensional nested-grid code. Although they resolved the region inside the Hill radius, they did not investigate the structure in the proximity to the protoplanet because they adopted a sink cell at $0.1r_{\\rm H}$ that corresponds to $\\sim70\\,\\rp$ at Jovian orbit (5.2 AU). Thus, we cannot observe a circumplanetary disk at $r\\ll\\rh$ in their calculation. \\citet{tanigawa02a} also investigated the gas flow around a protoplanet using two-dimensional nested-grid code. They resolved the region from $12\\,\\rh$ to $0.005\\,\\rh$. They found that a circumplanetary disk with $100\\me$ is formed around the protoplanet. \\citet{tanigawa02b} and \\citet{machida06c} investigated the evolution of the protoplanet system using three-dimensional nested-grid code. They found that the gas flow pattern in three dimensions is qualitatively different from that in two dimensions: the gas is flowing into the protoplanet system only in the vertical direction in three-dimensional simulations.    In the present study, we calculated the evolution of the protoplanet system using three-dimensional nested-grid code. We found that after the flow around the protoplanet reaches a steady state, the angular momentum accreting onto the protoplanet system is well converged regardless of both the cell width and the size of the sink cell region, while the mass accretion rate is not well converged. Although we calculated the evolution of the protoplanet system with spatial resolution much higher than previous studies, we still need further higher spatial resolution to determine the mass accretion rate onto the protoplanet. Thus, in this paper, we focus on the gas flow onto and around the protoplanet system and the accretion process of the angular momentum, and do not deal with the mass accretion rate (we plan to investigate the mass accretion rate with higher spatial resolution using a higher-performance computer in a subsequent paper). Note that the specific angular momentum accreting onto the protoplanet system with fixed protoplanet mass does not strongly depend on the mass accretion rate as described in the following sections.    The structure of the paper is as follows. The frameworks of our models are given in \\S 2 and the numerical method is described in \\S3. The numerical results are presented in \\S 4. \\S 5 is devoted for discussions. We summarize our conclusions in \\S6. ",
        "conclusions": "\\subsection{Effect of Thermal  Pressure} \\label{sec:thermal} In this paper, we investigated the evolution of the protoplanet system under the isothermal approximation, which might not be problematic in the circumplanetary disk, but may not be valid in the very proximity to the protoplanet. To investigate the thermal effect around the protoplanet, we adopted the sink cell in some models. As shown in \\S\\ref{sec:sink}, the cell width in the fiducial model having $l_{\\rm max}=8$ is about 4 Jovian radius. Thus,  gas inside $r < 4 \\rj$  has the same thermal energy. If the protoplanet has almost the same size as the present gas giant planet, the thermal energy around the protoplanet may be overestimated in model without the sink cell, because an actual protoplanet is embedded in the small part of the innermost cell. On the other hand, when we adopt the sink cell, the thermal energy around the protoplanet is underestimated, because the thermal energy is artificially removed from the sink cell. However, as shown in \\S\\ref{sec:sink}, when the sink radius is much smaller than the Hill radius ($r_{\\rm sink} \\ll 1/100 \\rh$), there are little differences in the angular momentum acquired by the protoplanet system between models with and without the sink cell. This is because a large part of the angular momentum of the protoplanet system is distributed around the Hill radius ($0.5\\,\\rh\\lesssim r \\lesssim \\rh$) as shown in \\S\\ref{sec:res-ang}. The Jovian radius is much smaller than the Hill radius ($\\rh = 744\\, \\rj$). In addition, the angular momentum flowing into the protoplanet system is determined by the shearing motion in the region of  $r \\simeq \\rh$ (see \\ref{sec:dis-evo}). Therefore, under the assumption that the protoplanet is smaller than the innermost cell width, it is expected that the thermal effect from the protoplanet is sufficiently small for the acquisition process of the angular momentum when we are using  sufficient smaller cells than the Hill radius.   \\citet{mizuno80}, \\citet{bodenheimer86}, and \\citet{ikoma00} suggested that gaseous protoplanet has a large envelope with high temperature. For example, \\citet{mizuno80} showed that, for Jovian case, gas distributed in the range of $r\\gtrsim 2\\times 10^{11}$\\,cm behaves isothermally, while that distributed in the range of $r \\lesssim 2\\times 10^{11}$\\,cm behaves adiabatically in their spherically symmetric calculations (see Fig.4 of \\citealt{mizuno80}). Thus,  gas in the range of $r \\lesssim 2\\times 10^{11}$\\,cm ($\\sim30\\,\\rj$) has higher temperature than the ambient medium. However, it is expected that the gaseous envelope hardly affect the angular momentum of the protoplanet system, because the size of the envelope ($\\sim30\\,\\rj$) is much smaller than the Hill radius ($\\rh = 744\\, \\rj$). A large part of the angular momentum is distributed in the range of $0.5\\rh \\lesssim r \\lesssim \\rh$, as shown in \\S\\ref{sec:res-ang}. On the other hand, the size of the circumplanetary disk may be affected by the thermal envelope. Thus, when we investigate the formation of the circumplanetary disk, we have to include the realistic thermal evolution around the protoplanet. However, it is difficult to study the thermal evolution around the protoplanet, because we need to solve the radiation hydrodynamics in three dimensions. In a subsequent paper, we will investigate the effect of the thermal envelope under simple assumptions.      \\subsection{Angular Momentum for Jupiter and Saturn} \\label{sec:jupiter-saturn} In \\S\\ref{sec:dis-evo},  we fixed the physical quantities as those at $\\ro=5.2$\\,AU. However,  in our calculation, since we use the dimensionless quantities, we can rescale those at any  orbits $\\ro$. Under the standard model \\citep{hayashi85}, we can generalize equation~(\\ref{eq:jlm}) as \\begin{equation} j_{\\rm lm} = 7.8 \\times 10^{15} \\left(\\dfrac{M_{\\rm p}}{M_{\\rm J}} \\right) \\left(\\dfrac{\\ro}{1\\,{\\rm AU}}  \\right)^{7/4} \\jcm = 2.5 \\times 10^{13} \\left( \\dfrac{M_{\\rm p}}{\\me} \\right)   \\left( \\dfrac{\\ro}{1\\, {\\rm AU}} \\right)^{7/4} \\jcm. \\label{eq:fit3} \\end{equation} We assume that the protoplanet system acquires the gas with the average specific angular momentum of equation~(\\ref{eq:fit3}) when the protoplanet has a mass of $M_{\\rm p}<1\\mj$. Thus, to estimate the angular momentum of the protoplanet system, we need to integrate equation~(\\ref{eq:fit3}) by the mass until the present values of gas giant planets. For example, the angular momentum in our model for Jupiter is \\begin{equation} J_{\\rm J} = \\int^{\\mj} j_{\\rm lm} (5.2\\,{\\rm AU})\\, dM = 1.3 \\times 10^{47} {\\rm g\\jcm}, \\end{equation} which is about 30 times larger than that of present Jupiter ($4.14\\times 10^{45}$g$\\jcm$). On the other hand, the angular momentum in our model for Saturn is \\begin{equation} J_{\\rm S} = \\int^{M_{\\rm s}} j_{\\rm lm} (9.6\\,{\\rm AU}) \\, dM = 3.6\\times 10^{46} \\ \\ {\\rm g\\,\\jcm}, \\end{equation} which is 50 times larger than that of present Saturn ($7.2 \\times 10^{44}$g$\\jcm$), where $M_{\\rm s} = 95.16\\me$ is the Saturnian mass. Note that the orbital angular momenta of the {\\em present} Jovian and Saturnian satellites can be ignored in the protoplanet system because they are considerably smaller  than the Jovian and Saturnian spin angular momenta.    The above estimation corresponds to the angular momentum of proto planet-disk system, which includes the central planet and protoplanetary disk. We showed that the angular momentum of the protoplanet system is $30-50$ times larger than  present spin of the gaseous planet. The angular momenta flowing into the Hill sphere are distributed into the spin of the protoplanet and orbital motion of the circumplanetary disk. It is expected that a large fraction of the total angular momentum contributes to the formation of the circumplanetary disk and the residual contributes to the spin of the planet. Although the fraction of the angular momentum of the circumplanetary disk to the spin of protoplanet is not correctly estimated in our calculation, a part of the angular momentum is certainly distributed into the disk. Thus, the angular momentum transfer and dissipation mechanism for the circumplanetary disk are necessary to follow further evolution of the protoplanet system. \\citet{takata96} proposed the despin mechanism of the protoplanet by planetary dipole magnetic field, in which an initially rapidly rotating protoplanet can be spun down to the present value by the magnetic interaction between the protoplanet and circumplanetary disk.     \\subsection{Disk formation and Implication for Satellite Formation} \\label{sec:dis-disk} When the circumplanetary disk is formed around the protoplanet, it is possible to form satellites in the disk. While there are many scenarios for the satellite formation \\citep[e.g.,][]{stevenson86},  regular satellites around gas giant planets are supposed to be formed in the gaseous disk as for the planet formation in the protoplanetary disk \\citep[e.g.,][]{korycansky91, canup02}. Observations showed that the regular satellites around Jupiter and Saturn are distributed only in the close vicinity of the planet ($r\\lesssim 50\\,\\rp$), and are on prograde orbits near the equatorial plane. Thus, it is expected that these regular satellites  formed in the circumplanetary disk.    We discussed the angular momentum of the protoplanet system in \\S\\ref{sec:dis-evo}. As the angular momentum flowing into the Hill sphere is brought into both the protoplanet and circumplanetary disk, we cannot estimate the fraction of angular momentum for the circumplanetary  disk. When we assume that the centrifugal force is balanced with the gravity of the protoplanet, we can derive the centrifugal radius $r_{\\rm cf}$ as \\begin{equation} r_{\\rm cf} = \\dfrac{j^2}{GM_{\\rm p}}. \\label{eq:cent} \\end{equation} To quantify the disk size, we adopt the specific angular momentum in equation~(\\ref{eq:cent}) as those in equation~(\\ref{eq:fit3}). The centrifugal radius for the proto-Jovian disk with $M_{\\rm p} = 1 \\mj$ and $\\ro =5.2$\\,AU is \\begin{equation} r_{\\rm cf, J}  = 1.5 \\times10^{11} \\ \\ {\\rm cm}, \\label{eq:r-est} \\end{equation} which is 22 times as large as Jovian radius ($r_{\\rm Jup} = 7.1\\times 10^9\\cm$). The Galilean satellites, which are regular  satellites, are distributed in the range of $6\\,r_{\\rm Jup}\\lesssim r \\lesssim 27\\,r_{\\rm Jup}$, which is consistent with the centrifugal radius derived from our calculation. Note that the disk size is expected to be larger than equation~(\\ref{eq:r-est}) owing to the thermal effect around the protoplanet. In the same way, we estimate the centrifugal radius of proto-Saturnian disk as \\begin{equation} r_{\\rm cf, S}  = 4.1\\times10^{11} \\ \\ {\\rm cm}, \\end{equation} which is 68 times as large as Saturnian radius ($r_{\\rm Sat}=6.0\\times10^9$\\,cm). The Saturnian representative regular satellites (Mimas, Enceladus, Tethys, Dione, Rhea, Titan, Hyperion, and Iapetus) are distributed in the range of $3\\, r_{\\rm Sat} \\lesssim  r \\lesssim 60\\, r_{\\rm Sat}$, which is also corresponding to the centrifugal radius of our result.    Figure~\\ref{fig:13} shows the structure around the protoplanet at the end of the calculation ($\\deft \\simeq 20$) for model M1, in which the protoplanet has mass of $1\\mj$. Each right panel is 4 times magnification of each left panel. Figures~\\ref{fig:13}{\\it a}, {\\it b}, and {\\it c} show the shock fronts outside the Hill radius, and a thick disk between the shock front and the Hill sphere. In the proximity to the protoplanet, a concave structure is seen in  Figures~\\ref{fig:13}{\\it e} and {\\it f}. The contours in these figures rapidly drop around the rotation axis (i.e., $z$-axis). Thus, even inside the Hill sphere, except for the central region, the disk has a thick torus-like structure. Figure~\\ref{fig:13}{\\it g}, {\\it h}, and {\\it i} show the structure in the very proximity to the protoplanet. On the midplane, the elliptical structure is seen in a large scale (Fig.~\\ref{fig:13}{\\it d}), while an almost round shape is seen in the proximity to the protoplanet (Fig.~\\ref{fig:13}{\\it g}) because the force acting on gas in this region is dominated by the gravity of planet. The contours in Figures~\\ref{fig:13}{\\it h} and {\\it i} show  a very thin disk in the region of $\\tl{r} \\lesssim 0.1$ (or $r \\lesssim 50\\,\\rj$). Thus, the regular satellites might be formed in these circumplanetary disks. However, to study the satellite formation in more detail, we need more realistic calculations."
    },
    "0801/0801.3133_arXiv.txt": {
        "abstract": "{I review briefly different aspects of the MOND paradigm, with emphasis on phenomenology, epitomized here by many MOND laws of galactic motion--analogous to Kepler's laws of planetary motion. I then comment on the possible roots of MOND in cosmology, possibly the deepest and most far reaching aspect of MOND. This is followed by a succinct account of existing underlying theories. I also reflect on the implications of MOND's successes for the dark matter (DM) paradigm: MOND predictions imply that baryons alone accurately determine the full field of each and every individual galactic object. This conflicts with the expectations in the DM paradigm because of the haphazard formation and evolution of galactic objects and the very different influences that baryons and DM are subject to during the evolution, as evidenced, e.g., by the very small baryon-to-DM fraction in galaxies (compared with the cosmic value). All this should disabuse DM advocates of the thought that DM will someday be able to reproduce MOND: it is inconceivable that the modicum of baryons left over in galaxies can be made to determine everything if a much heavier DM component is present.} ",
        "introduction": "\\label{section1} MOND is an alternative paradigm to Newtonian dynamics, whose original motivation was to explain the mass discrepancies in galactic systems without invoking dark matter (DM) (Milgrom 1983a). It constitutes a modification of dynamics in the limit of low accelerations that rests on the following basic assumptions: (i) There appears in physics a new constant, $\\az$, with the dimensions of acceleration. (ii) Taking the formal limit $\\az\\rar 0$ in all the equations of physics restores the equations of classical (pre-MOND) dynamics. (iii) For purely gravitational systems, the opposite, deep-MOND limit, $\\az\\rar \\infty$, gives limiting equations of motion that can be written in a form where the constants $\\az$ and $G$, and all masses in the problem, $m_i$, appear only in the product $m_iG\\az=m_i/\\mz$, where $\\mz\\equiv(G\\az)^{-1}$ (Milgrom 2005)\\footnote{By this I mean that starting with equations that involve (including in derivatives) $\\vr$, $t$, $G$, $\\az$, $m_i$, and gravitational field degrees of freedom, we can rewrite them, possibly by redefining the gravitational field, so that only $\\vr$, $t$, $m_i/\\mz$, and the gravitational field appear.}. This last fiat reproduces the desired MOND phenomenology for purely gravitational systems. A MOND theory is one that incorporates the above tenets in the nonrelativistic regime. \\par Since all our knowledge of MOND comes, at present, from the study of purely gravitational systems (galactic systems, the solar system, etc.) it is still an open question how exactly to extend the third MOND tenet to systems involving arbitrary interactions. One possibility is to require that for $\\az\\rar\\infty$, the limiting equations of motion can be brought to a form where $G$, $\\az$, and $m_i$ appear only as $G\\az^2=\\az/\\mz$ and $m_i/\\az$. This requirement will automatically cause $\\az$, $G$, and $m_i$ to appear as $m_i/\\mz$ in the deep MOND limit of purely gravitational systems, as in such cases $G$ and $m_i$ always appear in equations as $Gm_i$. Such a general requirement would also replace Newton's second law $\\vF=m\\va$ by $\\vF=m\\vQ/\\az$ in the deep MOND regime, where $\\vQ$ is some functional with dimensions of acceleration squared that does not depend on $\\az$. \\par Detailed reviews of phenomenological and theoretical aspects of the paradigm can be found, for instance, in Sanders and McGaugh (2002), and, more recently, in Scarpa (2006) and in Bekenstein (2006).   It follows from the third tenet that an underlying MOND theory must be nonlinear in the sense that an acceleration of a test particle due to a combination of several fields is not simply the sum of the accelerations produced by the individual fields. Take, as an example, the purely gravitational case where we modify the equations for the gravitational field. Linearity would mean that in the nonrelativistic limit of the theory the acceleration of a test particle at position $\\vr$ in the field of $N$ masses $m_i$ at positions $\\vr_i$ is given by \\beq \\va=\\sum_{i=1}^N m_i q_i(G, \\az, \\vr, \\vr_1,...,\\vr_N). \\eeqno{lino} The third assumption says that in the deep MOND case $q_i\\propto \\mz^{-1}$, but this is dimensionally impossible. Clearly, the acceleration produced even by a single point  cannot be linear in its mass, as we shall see explicitly below. \\par It also follows from the third tenet, and the assumption that $\\az$ is the only new dimensioned constant, that the deep MOND limit of any theory must satisfy the following scaling laws for purely gravitational systems: In general, on dimensional grounds, all physics must remain the same under a change of units of length $\\ell\\rar\\lambda\\ell$, of time $t\\rar\\lambda t$, and no change in mass unit $m\\rar m$. Under these we have $\\az\\rar \\lambda^{-1}\\az$ and $G\\rar \\lambda G$; so, the constant $\\mz$, which alone appears in the limiting theory, is invariant under the scaling\\footnote{Other dimensioned quantities, such as the gravitational potential, have to be scaled appropriately. Note that velocities are invariant.}. This tells us that we are, in fact, exempt from scaling the constants of the theory when we scale $(\\vr,t)\\rar \\lambda(\\vr,t)$. The theory is thus invariant under this scaling; namely, if a certain configuration is a solution of the equations, so is the scaled configuration. Specific theories may have even higher symmetries; for example, the above scaling property is only part of the conformal invariance of the deep MOND limit for the particular MOND theory of Bekenstein \\& Milgrom (1984), as found in Milgrom (1997).  \\par As a corollary of the scaling invariance we have: if $\\vr(t)$ is a trajectory of a point body in a configuration of masses $m_i$ at positions $\\vr_i(t)$ (which can be taken as fixed, for example), then $\\hat\\vr(t)=\\lambda\\vr(t/\\lambda)$ is a trajectory for the configuration where $m_i$ are at $\\lambda\\vr_i(t/\\lambda)$, and the velocities on that trajectory are $\\hat\\vV(t)=\\vV(t/\\lambda)$. (An extended mass changes its size and density such that the total mass remains the same. A point mass remains a point mass of the same value.) \\par Since $\\mz$ has dimensions of $mt^4/\\ell^4$, another scaling under which it is invariant is $m_i\\rar\\lambda m_i$, $\\vr_i\\rar\\vr_i$, $t\\rar\\lambda^{-1/4} t$. (The most general scaling allowed is $m\\rar\\lambda m$, $\\vr_i\\rar\\kappa\\vr_i$, $t\\rar\\kappa\\lambda^{-1/4}t$.) This means that for a purely gravitational system deep in the MOND regime, scaling all the masses leaves all trajectory paths the same but the bodies traverse them with all velocities scaling as $m^{1/4}$; accelerations then scale as $m^{1/2}$.  \\par It is enlightening to draw an analogy between the role of $\\az$ in MOND with the role of $\\hbar$ in quantum physics, or that of $c$ in relativity. These constants, each in its own realm, mark the boundary between the classical and modified regimes; so formally pushing these boundaries to the appropriate limits ($c\\rar\\infty$, $\\hbar\\rar0$, $\\az\\rar 0$) one restores the corresponding classical theory (for quantum theory, in some weak sense). In addition, these constants enter strongly the physics in the modified regime where they feature in various phenomenological relations. For example, $\\hbar$ appears in the black body spectrum, the photoelectric effect, atomic spectra, and in the quantum Hall effect. The speed of light appears in the Doppler effect, in the mass vs. velocity relation, and in the radius of the Schwarzschild horizon. Without the respective, underlying theory, these disparate phenomena would appear totally unrelated, and the appearance of the same numerical constant in all of them would constitute a great mystery. The MOND paradigm similarly predicts a number of laws related to galactic motion, some of which are qualitative, but many of which are quantitative and involve $\\az$. Since they appear to be obeyed by nature it should indeed be a great mystery why that should be so without the underlying MOND paradigm (i.e. with Newtonian dynamics plus DM). I now discuss some of these predictions in some detail. ",
        "conclusions": ""
    },
    "0801/0801.2303_arXiv.txt": {
        "abstract": "We present here some of the first results we have obtained on the study of the optical spectra of {\\em Spitzer/MIPS} 24 $\\rm \\mu m$-selected galaxies  in the COSMOS field.  This is part of a series of studies we are conducting to analyse the optical spectral properties of mid-infrared (mid-IR) galaxies with different IR luminosities up to high redshifts. The results shown here correspond to the brightest $S_{24 \\, \\rm \\mu m}>2 \\, \\rm mJy$ IR galaxy population at $z<1$. ",
        "introduction": "Since the discovery of an extragalactic component for the IR background a decade ago~\\cite{ref:pug}, numerous studies have demonstrated the key importance of IR galaxies to reconstruct the history of galaxy star formation and stellar mass assembly. Recent surveys conducted with the {\\em Spitzer Space Telescope} have allowed for the resolution of more than 70\\% of the mid and far-IR background light~\\cite{ref:dole} and a quite detailed understanding of the evolution of IR galaxies with cosmic epoch, e.g.~\\cite{ref:lefl,ref:cap06a,ref:cap06b,ref:cap07}.  The IR galaxy luminosity function has a strong positive luminosity evolution from redshifts $z=0$ to 1 (see~\\cite{ref:lefl,ref:cap07}), implying that more luminous IR galaxies dominate the IR output with increasing redshift. Beyond $z\\sim1$, basically all the IR light is produced by luminous and ultra-luminous IR galaxies (LIRGs and ULIRGs, respectively) with bolometric IR luminosities $L_{\\rm bol.}> 10^{11} \\, \\rm L_\\odot$. Active galactic nuclei (AGN) have some  significant  (but not yet well-determined) contribution to the IR background and govern the shape of the bright-end of the mid-IR LF at $z\\sim2$~\\cite{ref:cap07}.  The study of the optical spectra of IR galaxies can reveal several aspects of the physical conditions in which star formation and AGN activity take place. What is the evolutionary stage of the hosts of IR activity? What is the relationship between dust emission and the observed  reddening in optical bands? Does the IR phase accompany all the process of star formation or is it set at a particular time? The large sample of zCOSMOS spectra~\\cite{ref:lil07} being taken over the 2 deg$^2$  COSMOS field offers a perfect opportunity to address  these questions. ",
        "conclusions": ""
    },
    "0801/0801.2245_arXiv.txt": {
        "abstract": "{} { We have carried out an accurate investigation of the Ca isotopic composition and stratification in the atmospheres of 23 magnetic chemically peculiar (Ap) stars of different temperature and magnetic field strength. } { With the UVES spectrograph at the 8\\,m ESO VLT, we have obtained high-resolution spectra of Ap stars in the wavelength range 3000--10000\\,\\AA.  Using a detailed spectrum synthesis calculations, we have reproduced a variety of Ca lines in the optical and ultraviolet spectral regions, inferring the overall vertical distribution of Ca abundance, then we have deduced the relative isotopic composition and its dependence on height using the profile of the the IR-triplet \\ion{Ca}{ii} line at $\\lambda$~8498\\,\\AA. } { In 22 out of 23 studied stars, we found that Ca is strongly stratified, being usually overabundant by 1.0--1.5\\,dex below $\\log\\tau_{5000}\\approx -1$, and strongly depleted above $\\log\\tau_{5000}=-1.5$.  The IR-triplet \\ion{Ca}{ii} line at $\\lambda$~8498\\,\\AA\\ reveals a significant contribution of the heavy isotopes $^{46}$Ca and $^{48}$Ca, which represent less than 1\\,\\% of the terrestrial Ca isotopic mixture. We confirm our previous finding that the presence of heavy Ca isotopes is generally anticorrelated with the magnetic field strength. Moreover, we discover that in Ap stars with relatively small surface magnetic fields ($\\le$\\,4--5 kG), the light isotope $^{40}$Ca is concentrated close to the photosphere, while the heavy isotopes are dominant in the outer atmospheric layers. This vertical isotopic separation, observed for the first time for any metal in a stellar atmosphere, disappears in stars with magnetic field strength above 6--7\\,kG. } { We suggest that the overall Ca stratification and depth-dependent isotopic anomaly observed in Ap stars may be attributed to a combined action of the radiatively-driven diffusion and light-induced drift. } ",
        "introduction": "\\label{intro}  After the pioneering work by Michaud (\\cite{Michaud70}), atomic diffusion in stellar envelopes and atmospheres has been recognized as the main process responsible for the atmospheric abundance anomalies in the peculiar stars of the upper main sequence. Due to their unique characteristics, such as an extremely slow rotation, strong, global magnetic fields, and the absence of convective mixing, magnetic chemically peculiar (Ap and Bp) stars  exhibit the most clear manifestation of diffusion effects and thus represent privileged laboratories for investigation of chemical transport processes and magnetohydrodynamics.  Detailed diffusion calculations performed for a set of chemical elements in the atmospheres of magnetic peculiar stars predict separation of chemical elements over the stellar surface and with height in stellar atmosphere (abundance stratification). These theoretical predictions can be directly tested through the comparison with empirical maps of chemical elements inferred from observations. For a small number of elements, including Ca, an effect of the vertically stratified element distribution on the spectral line profiles was demonstrated in early studies (Borsenberger at al. \\cite{BPM81}). However, the absence of high-resolution, high signal-to-noise spectroscopic observations did not allow a robust comparison between observations and theoretical diffusion modelling. This step was carried out by Babel (\\cite{Babel92}), who calculated the Ca abundance distribution in the atmosphere of magnetic Ap star 53~Cam and showed that the unusual shape of \\ion{Ca}{ii} K line -- a combination of the wide wings and extremely narrow core (Babel \\cite{Babel94}; Cowley et al. \\cite{CHK06}) -- is a result of a step-like Ca distribution with abundance decrease at $\\log\\tau_{5000}\\approx-1$. Following Babel, the step-function approximation for the abundance distributions was employed in many stratification studies based on the observed profiles of spectral lines (\\txtbf{Savanov et al. \\cite{SKV01}}; Bagnulo et al. \\cite{bagetal01}; Wade et al. \\cite{WLRK03}; Ryabchikova et al. \\cite{RPK02,RLK05,RRKB06}; \\txtbf{Glagolevskii et al. \\cite{GRC05}; Cowley et al. \\cite{CHC07}}).  Ca was found to be stratified in nearly the same way as in 53 Cam (enhanced concentration of Ca below $\\log\\tau_{5000}\\approx-1$ and its depletion above this level) in all stars for which stratification analysis have been performed: $\\beta$~CrB (Wade et al. \\cite{WLRK03}), $\\gamma$~Equ (Ryabchikova et al. \\cite{RPK02}), HD~204411 (Ryabchikova et al. \\cite{RLK05}), HD~133792 (Kochukhov et al. \\cite{KTRM06}) and HD~144897 (Ryabchikova et al. \\cite{RRKB06}).  In addition to remarkable diffusion signatures, recently another Ca anomaly was detected, first in the spectra of the Hg-Mn stars by Castelli \\& Hubrig (\\cite{CastH04}) and then in Ap stars by Cowley \\& Hubrig (\\cite{CH05}). These authors found a displacement of the lines of \\ion{Ca}{ii} IR triplet due to significant contribution of the heavy Ca isotopes. Ryabchikova, Kochukhov \\& Bagnulo (see review paper by Ryabchikova (\\cite{MONS05})) were the first to apply spectrum synthesis calculations to investigate effects of Ca isotopes on the calcium line profile shape in magnetic CP stars. In particular, we have demonstrated the general anticorrelation between the presence of heavy Ca and magnetic field strength: in Ap stars the contribution of heavy Ca isotopes decreases with the increase of the field modulus, and disappears when the field strength exceeds $\\sim$\\,3\\,kG. Cowley et al. (\\cite{CHC07}) have come to a similar conclusion, but with reservations.  In the present paper we summarise our detailed analysis of the vertical stratification of Ca abundance in the atmospheres of magnetic Ap stars of different temperatures and magnetic field strengths. We combine this vertical Ca mapping with the analysis of the Ca isotopic anomaly and its dependence on height, based on the spectrum synthesis modelling of the IR triplet \\ion{Ca}{ii} line at $\\lambda$~8498\\,\\AA. Our study is the first to present a homogeneous and systematic determination of the vertical stratification of a given element in the atmospheres of a large number of stars, enabling us to study the signatures of atmospheric atomic diffusion in the presence of strong magnetic field as a function of stellar parameters and magnetic field intensity.  Our paper is organized as follows. In Sect.~\\ref{obs} we describe spectroscopic observations and  the data reduction. Determination of the stellar atmospheric parameters is detailed in Sect.~\\ref{parameters}. Sects.~\\ref{synthesis}--\\ref{isot} present our methodological approach to magnetic spectrum synthesis, determination of Ca stratification and the study of Ca isotopic composition. Our findings are summarised in Sect.~\\ref{results} and are discussed in Sect.~\\ref{discus}. ",
        "conclusions": "\\label{discus}  In this paper we have investigated the vertical stratification and isotopic anomaly of Ca in a large sample of cool magnetic Ap stars. Our study is the first to address the interesting problem of the presence of heavy Ca isotopes in chemically peculiar stars with detailed polarized radiative transfer calculations, which take into account the effects of the magnetic field and chemical separation in stellar atmospheres. We derive stratification of Ca for 23 Ap stars using a sample of \\ion{Ca}{i} and \\ion{Ca}{ii} lines distributed over a broad spectral range. All but one program stars clearly show signatures of the Ca stratification, whereas comparison stars reveal no inhomogeneities in the vertical Ca distribution when analysed with the same techniques and atomic data. Although our stratification modelling was based on a simplified step-function approximation and adopted a simplified homogeneous model for the magnetic field geometry, the inferred parameters of the vertical Ca distributions impose important observational constraints for theoretical models of the radiative diffusion processes in stellar atmospheres.  Analysis of the IR Ca triplet line \\ion{Ca}{ii} 8498\\,\\AA\\ provided information on the relative contribution of different Ca isotopes. Spectrum synthesis calculations show that significant isotopic shifts observed in the core of \\ion{Ca}{ii} 8498\\,\\AA\\ cannot be attributed to the overabundance of heavy Ca isotopes throughout the whole stellar atmosphere. Instead, \\txtbf{we show that} the Ca line profile shape \\txtbf{is consistent with} the vertical separation of different Ca isotopes, with heavy Ca located in a cloud above the most abundant isotope $^{40}$Ca. Even though the presence of heavy Ca is prominent in the line core, the normal Ca isotope dominates the line wings and is more abundant in the lower atmosphere where the total Ca abundance is also much larger. Thus, our \\txtbf{tentative} model calls for a less extreme heavy Ca enrichment than was suspected in previous investigations limited to the centroid measurements of the \\ion{Ca}{ii} 8498\\,\\AA\\ line core.  Recently different stratification of He isotopes was found by Bohlender (\\cite{B05}) in the analysis of $^{3}$He and related stars. That paper used a similar approach to modelling chemical stratification, but has only addressed He vertical distribution in hot non-magnetic chemically peculiar stars.  We have successfully used the model with Ca stratification and vertical isotopic separation to explain the appearance of the \\ion{Ca}{ii} 8498\\,\\AA\\ in all stars showing excess of heavy Ca isotopes. The prominent anticorrelation between the presence of heavy Ca and magnetic field strength, first reported by Ryabchikova (\\cite{MONS05}), is confirmed and strengthened in the present paper. We find that only stars with sufficiently weak field show traces of heavy Ca. According to our knowledge, this interesting relation is the only case when definite dependence of the chemical abundance characteristic on the magnetic field strength is found for Ap stars. Furthermore, we find that in stars with stronger fields the heavy Ca isotopes tend to accumulate higher in the atmosphere.  According to the results of our study, pulsating (roAp) and non-pulsating Ap stars are not distinguished by the characteristics of their Ca isotopic anomaly. Both groups of stars are equally likely to show an excess of heavy Ca and follow the same trend of the heavy isotope contribution versus the magnetic field strength and the \\ion{Ca}{ii} 8498\\,\\AA\\ line intensity.  If the overall distribution of Ca abundance in the atmospheres of Ap stars follows the predictions of the radiatively driven diffusion, our results on the isotopic separation favour the light-induced drift (LID) as the main process responsible for  this separation. According to Atutov \\& Shalagin (\\cite{AS88}), LID arises when the radiation field is anisotropic inside the line profile. Such an anisotropy takes place for a line of the trace isotope, for instance $^{46}$Ca and $^{48}$Ca in the terrestrial calcium mixture, which is located in the wing of a strong line of the main isotope $^{40}$Ca. The main isotope induces the drift velocity for other isotopes. If the trace isotope's line is located in the red wing of the line due to main isotope, the drift velocity is directed towards the upper atmosphere and the trace isotopes are pushed upwards. This is in agreement with observed vertical distribution of Ca isotopes. The Zeeman splitting changes the line shape and decreases the flux anisotropy for the line of trace isotope. When magnetic field becomes strong enough, $\\sim$5--6~kG, the flux anisotropy disappears and the isotopic separation is ceasing. Therefore, the observed Ca isotopic anomaly in magnetic stars may be qualitatively explained by the combined action of the radiatively-driven diffusion and the light-induced drift. Theoretic study of a combination of the radiatively-driven diffusion and LID was presented by LeBlanc \\& Michaud (\\cite{LM93}) for He. These authors showed that LID accelerates separation of $^{3}$He from $^{4}$He in hotter CP stars. Detailed theoretical chemical diffusion calculations (LeBlanc \\& Monin \\cite{LM04}) should incorporate LID in order to test our hypothesis that this effect may be important for the chemical transport processes in cool Ap-star atmospheres."
    },
    "0801/0801.4769_arXiv.txt": {
        "abstract": "{We present results of a study of large-scale neutral hydrogen (\\HI) gas in nearby radio galaxies. We find that the early-type host galaxies of different types of radio sources (compact, \\FRI\\ and \\FRII) appear to contain fundamentally different large-scale \\HI\\ properties: enormous regular rotating disks and rings are present around the host galaxies of a significant fraction of low power compact radio sources, while no large-scale \\HI\\ is detected in low power, edge-darkened \\FRI\\ radio galaxies. Preliminary results of a study of nearby powerful, edge-brightened \\FRII\\ radio galaxies show that these systems generally contain significant amounts of large-scale \\HI, often distributed in tail- or bridge-like structures, indicative of a recent galaxy merger or collision. Our results suggest that different types of radio galaxies may have a different formation history, which could be related to a difference in the triggering mechanism of the radio source. If confirmed by larger studies with the next generation radio telescopes, this would be in agreement with previous optical studies that suggest that powerful \\FRII\\ radio sources are likely triggered by galaxy mergers and collisions, while the lower power \\FRI\\ sources are fed in other ways (e.g. through the accretion of hot IGM). The giant \\HI\\ disks/rings associated with some compact sources could - at least in some cases - be the relics of much more advanced mergers.}  \\FullConference{From planets to dark energy: the modern radio universe\\\\ October 1-5 2007\\\\ University of Manchester, Manchester, UK}  \\begin{document} ",
        "introduction": "Radio sources are generally hosted by early-type galaxies. A significant fraction of these radio galaxies, in particular the more powerful ones, show optical peculiarities (tail, bridges, shells, etc.) that are indicative of galaxy mergers or interactions \\citep{hec86,bau92}. This has led to the suggestion that galaxy mergers or collisions could be the trigger for the radio-AGN activity in these systems.  Here we present the large-scale neutral hydrogen (\\HI) properties of different types of nearby radio galaxies in order to investigate whether these systems have been involved in a recent or past gas-rich merger or interaction. Simulation show that in a major merger between gas-rich galaxies, large-scale gaseous tidal structures can be expelled from the merging system, which after several rotations ($\\gtrsim$ Gyr) can partially fall back onto the host galaxy and settle in a low-density, regular rotating disk or ring \\citep{bar02}. \\HI\\ observations -- in particular when combined with a stellar population analysis of the host galaxy \\citep{tad05} -- therefore provide an excellent tool to trace and date galaxy mergers on relatively long time-scales, which can be compared with the age of the radio-loud AGN in these galaxies in order to investigate a possible connection with the triggering of the radio source. ",
        "conclusions": "\\label{sec:discussion}  Our results of large-scale \\HI\\ around nearby radio galaxies imply that there could be {\\sl a fundamental difference in large-scale \\HI\\ content between the various types of radio galaxies}, which is visualized in more detail in Fig.~\\ref{fig:HImassplot}. While the host galaxies of several low power compact radio sources contain a large-scale regular rotating \\HI\\ disk or ring (possibly as the result of an advanced merger), the extended \\FRI\\ radio galaxies in the same sample lack any large-scale \\HI\\ structures outside the optical body of the host galaxy up to a conservative detection limit of a few $\\times 10^{8} M_{\\odot}$. On the other hand, preliminary results for a small sample of the more powerful \\FRII\\ radio sources show that most of their host galaxies do contain significant amounts of \\HI, often distributed in tail- or bridge-like structures. If confirmed by future studies with greater statistical significance and lower detection limits, this could mean that there is a fundamental difference in the formation history of different types of radio galaxies and, related, the triggering of their radio source. This would be in agreement with optical studies that suggest that powerful \\FRII\\ sources are generally fueled by major mergers, while \\FRI\\ source are likely powered in another way, for example through the accretion of gas from the hot inter-galactic medium \\citep{hec86,bau92}.  Future instruments, with eventually the Square Kilometre Array, will be essential to verify our results by observing \\HI\\ emission in large statistical samples of radio galaxies.  \\begin{figure}[t] \\centering \\includegraphics[width=10cm]{Proc_SKA_HImass_v3_lowres.eps} \\caption{Large-scale \\HI\\ mass plotted against the total linear extent of the radio source. Black circles and arrows represent the sources from our B2 sample (compact + \\FRI), while red squares and arrows are preliminary results from our \\FRII\\ sample. In case of non-detection a conservative upper limit (3$\\sigma$ across 400 \\kms) was estimated.} \\label{fig:HImassplot} \\end{figure}"
    },
    "0801/0801.3463_arXiv.txt": {
        "abstract": "The 21-cm anisotropies from the neutral hydrogen distribution prior to the era of reionization is a sensitive probe of primordial non-Gaussianity. Unlike the case with cosmic microwave background, 21-cm anisotropies provide multi-redshift information with frequency selection and is not damped at arcminute angular scales. We discuss the angular trispectrum of the 21-cm background anisotropies and discuss how the trispectrum signal generated by the primordial non-Gaussianity can be measured with  the three-to-one correlator and the corresponding angular power spectrum. We also discuss the separation of primordial non-Gaussian information in the trispectrum with that generated by the subsequent non-linear gravitational evolution of the density field. While with the angular bispectrum of 21-cm anisotropies one can limit the second order corrections to the primordial fluctuations  below $f_{\\rm NL}\\sim 1$, using the trispectrum information we suggest that the third order  coupling term, $f_2$ or $g_{\\rm NL}$, can be constrained to be arounde 10 with future 21-cm observations over the redshift interval of 50 to 100. ",
        "introduction": "The cosmic 21-cm background involving spin-flip line emission or absorption of neutral hydrogen contains unique signatures  on how the neutral gas evolved from last scattering at $z \\sim 1100$ to complete reionization at $z < 10$ \\cite{Furlanetto}. Subsequent to recombination, the temperature of neutral gas is coupled to that of the cosmic microwave background (CMB). At redshifts below $\\sim$ 200 the gas cools adiabatically, its temperature drops below that of the CMB, and neutral hydrogen resonantly absorbs CMB flux through the spin-flip transition \\cite{field,loeb,Bharadwaj}. The inhomogeneous neutral hydrogen density distribution generates anisotropies in the brightness temperature measured relative to the blackbody CMB \\cite{zaldarriaga}. The large cosmological and astrophysical information content in 21-cm background is well understood in the literature \\cite{loeb,Iliev,Santos,Bowman,McQuinn,Pen,sigurdson}.  Parallel to the large effort in analytical and numerical calculations of the 21-cm properties, there are now several first generation 21-cm experiments underway focusing on the 21-cm signal during the era of reionization. At the low redshifts probed by these first generation interferometers, the 21-cm signal is modified by the  astrophysics of first  sources and the associated UV photon background. There, one naturally expects fluctuations to be dominated by inhomogeneities in source properties \\cite{Zahn,Santos2}. The associated non-Gaussianity in 21-cm anisotropies leads to a measurable three-point correlation function or a bispectrum \\cite{Coo05,Bharadwaj2,morales2}. Such a non-Gaussianity is also expected to dominate the signature in 21-cm anisotropies generated by the primordial non-Gaussian density field. However, the 21-cm background generated by neutral hydrogen at redshifts of 50 to 100 prior to the onset of reionization and the appearance of first stars is expected to provide a cleaner probe of the primordial density perturbations in the same manner CMB observations are used to study primordial fluctuations \\cite{loeb}. If the primordial fluctuations are non-Gaussian, then the 21-cm anisotropies at these high redshifts will naturally contain a signature associated with that non-Gaussianity \\cite{Cooray,pillepich}  The primordial non-Gaussianity in the density field can be studied with the three-point and higher-order correlation functions of the 21-cm background. In particular, the second order corrections to the density perturbations generated by primordial non-Gaussianity lead to a bispectrum with a dependence on the second-order correction to the curvature perturbations, $f_{\\rm NL}$  \\cite{Komatsu,Bartolo3}. With future low frequency data out to $z \\sim 100$,  21-cm background anisotropies could potentially limit $f_{\\rm NL} < 0.1$ \\cite{Cooray}. In comparison, the expected non-Gaussianity under standard inflation is of order $|n_s-1|$ and with the scalar spectral index $n_s \\sim 0.98$ \\cite{Spergel}, $f_{\\rm NL}$ is expected to be well below unity \\cite{Maldacena}. The primordial non-Gaussianity parameter at the second order $f_{\\rm NL}$, however, has a correction associated with evolution of second and higher-order perturbations after inflation \\cite{Bartolo}. For standard slow-roll inflation then $f_{\\rm NL} = -5/12(n_s-1) + 5/6 +3/10f(k)$ where the last term is momentum dependent. In this case $f_{\\rm NL}$ is at the level of a few tenths and could be as high as 1. The ability for 21-cm anisotropies to probe $f_{\\rm NL}$ as low as 0.1 is important since even a perfect CMB experiment limited by cosmic variance alone can only restrict $f_{\\rm NL} > 3$ \\cite{Komatsu,Babich,Smith} while there is no significant improvement when using low redshift large-scale structure \\cite{Scoc,Dalal}.  The possibility to make primordial non-Gaussianity measurements with the 21-cm background is important since compared to most other probes of inflationary parameters, $f_{\\rm NL}$ is one of the few parameters for which we have limited number of probes sensitive to the low amplitude non-Gaussianity expected under standard slow-roll inflation. While CMB as a probe of non-Gaussianity is well known, when compared to CMB temperature and polarization anisotropies, 21-cm background has two distinct advantages:  (1) the ability to probe multiple redshifts based on frequency selection and (2) the lack of a damping tail in the 21-cm anisotropy  spectrum, unlike damping of CMB anisotropies at a multipole around 2000.  While the 21-cm background as a potentially interesting and a useful probe of $f_{\\rm NL}$ is now known, different scenarios for primordial fluctuations may produce a small $f_{\\rm NL}$ but a large third-order correction. In certain alternative models of inflation higher order terms may be significant even if the second-order term is small and such scenarios include the new ekpyrotic cosmology \\cite{Buchbinder} and, under certain conditions, the curvaton model \\cite{Sasaki}. Thus, beyond the non-Gaussianity at the three-point level with the bispectrum, it is also useful to study the non-Gaussianities of the 21-cm background at the four point-level involving the trispectrum.  Here we show that the angular trispectrum of 21-cm anisotropies contains a measurable non-Gaussianity from primordial fluctuations if the scale-independent cubic corrections to the gravitational potential captured with an amplitude parameter $f_2$ (or $g_{\\rm NL}$ as described in Ref.~\\cite{Amico}) has a value of order $\\sim 10$, even if the scale-independent quadratic corrections to the gravitational potential captured by $f_{\\rm NL}$ has a value around $\\sim$ 1. While $f_{\\rm NL}$ is currently constrained with WMAP data \\cite{Spergel}, there is no real constraint on this third order non-Gaussianity parameter. There are, however, theoretical expectations: for slow-roll inflation, a trispectrum of the form $T_4=1/2\\tau_{\\rm NL} [P(k_1)P(k_2)P(k_3) + ...]$ is expected for curvature perturbations \\cite{Seery} with an expectation value of $\\tau \\lesssim r/50$ where $r$ is the tensor-to-scalar ratio ($r \\lesssim 0.6$ is recent CMB data \\cite{Spergel}). In such a model, there is also a direct connection between $\\tau_{\\rm NL}$ and $f_{\\rm NL}$ such that $\\tau_{\\rm NL}=(6f_{\\rm NL}/5)^2$ (see the discussion in Ref.~\\cite{Kogo:2006} for connections between coupling terms of various models for primordial trispectrum). While such a relation is generally assumed when constraining $\\tau_{\\rm NL}$ \\cite{Lyth}, in future it may be that one can directly test the above relation between $f_{\\rm NL}$ and $\\tau_{\\rm NL}$ with data.  To measure the non-Gaussianity at the four-point level, we introduce the three-to-one correlator statistic, extending the two-to-one correlator of Ref.~\\cite{Coo01} and applied to WMAP data in Ref.~\\cite{Szapudi}. We optimize the angular power spectrum of the 3-1 correlator to detect the primordial trispectrum by appropriately filtering 21-cm anisotropy data to remove the non-Gaussian confusion  generated by the non-linear evolution of gravitational perturbations. To do this properly, one requires prior knowledge on the configuration dependence  of the primordial 21-cm trispectrum, but its amplitude is a free variable to be determined from the data.  This paper is organized as following: we first discuss the bispectrum in 21-cm anisotropies associated with primordial perturbations resulting from quadratic corrections to the primordial potential. We discuss ways to measure this bispectrum in the presence of other non-Gaussian signals and determine the extent to which $f_{\\rm NL}$ can be measured from 21-cm background data. In the numerical calculations described later, we take a fiducial flat-$\\Lambda$CDM cosmological model with $\\Omega_b=0.0418$, $\\Omega_m=0.24$, $h=0.73$, $\\tau=0.092$, $n_s=0.958$, and $A(k_0 = 0.05\\ \\mathrm{Mpc}^{-1})=2.3\\times10^{-9}$. This model is consistent with recent measurements from WMAP \\cite{Spergel}. ",
        "conclusions": "\\label{sec:results}  In Fig.~1 we summarize the power-spectrum of 21-cm anisotropies generated by the neutral hydrogen distribution at a redshift of 100 with a bandwidth for observations of 1 MHz. Here, we also plot $X_l^{\\rm prim}$, $X_l^{\\rm grav}$, and $N_l^{\\rm tot}$ for the same redshift with the optimal filter applied with $f_{\\rm NL}=10$ and $f_2=1$ as the non-Gaussian scale-independent amplitude of the primordial second- and third-order curvature perturbations, respectively. As shown, $N_l^{\\rm tot} > X_l^{\\rm grav}$, suggesting that the noise term is dominated by the Gaussian variance (second term in equation~\\ref{eqn:nl}). This statement is independent of $f_{\\rm NL}$ and $f_2$ and thus the non-Gaussian detection is dominated by the Gaussian term in $N_l$ regardless of what is assumed about non-Gaussianity. Note that we have estimated $N_l^{\\rm tot}$ in Fig.~1 under the assumption that observations are limited only by the cosmic variance and not accounting for any instrumental noise variance, which will also lead to a cut-off in $l$ out to which we can make measurements. Using the cosmic variance alone allows us to establish the potentially achievable limit and compare directly with cosmic variance limit with CMB data. When calculating $X_l^{\\rm prim}$ and $X_l^{\\rm grav}$ in equation~(\\ref{eqn:xl}) we set maximum value of $L$ in the sum associated with $T^{{l_1l_2}_{l_3l}}(L)$ to be $L_{\\rm max}=100$ We tested our calculation for $L_{\\rm max}=150$ and found results to be within a percent, but such a higher value slows the numerical calculation significantly.  Finally, due to computational limitations of the numerical calculation, we restrict estimate of $X_l$ to $l=10^3$.  In Fig.~1 when calculating $X_l^{\\rm grav}$, following the derivation in the Appendix and the discussion there, we use the exact analytical result for the 2-2 trispectrum of the density field with the mode coupling captured by $F_2(\\bk_1,\\bk_2)$ \\cite{Goroff,fry84}. For the 3-1 trispectrum of the density field under non-linear gravitational evolution, given that an analytical result for the angular trispectrum is cumbersome, we use the angular averaged value for $F_3(\\bk_1,\\bk_2,\\bk_3)$. We refer the reader to the Appendix for details.  In Fig.~2 we summarize the estimate related to signal-to-noise ratio for a detection of $X_l^{\\rm prim}$ as a function  of $l$. The typical signal-to-noise ratio, when measurements are out to a multipole of $10^3$, is at the level of $\\sim$ 0.5 if one assumes that the coupling parameters $f_{\\rm NL}=10$ and $f_2=3$ for 21-cm observations centered at a redshift of 100 over a bandwidth of 1 MHz. The signal-to-noise ratio for the case with $f_{\\rm NL}=0$ and $f_2=3$ is $\\sim$ 0.03. Since $S/N \\propto f_2$ when $f_{\\rm NL}=0$, out to $l=10^3$, a signal-to-noise ratio of 1 is achieved if $f_2 \\sim 10^2$. While there are neither strong theoretical motivations on the expected value of $f_2$ nor a real bound on its value from existing data, it is likely that with 21-cm data one can constrain $f_2$ down to a level well below this value for a single redshift. This is due to the fact that 21-cm observations lead to measurements at multiple redshifts, though one cannot make arbitrarily small bandwidths to improve the detection since at scales below a few Mpc, anisotropies in one redshift bin will be correlated with those in adjacent bins \\cite{Santos3}. For example, if 21-cm observations are separated to 30 independent bins over the redshift interval 50 to 100 (as can be achieved with 1 MHz bandwidths), then an approximate estimate of the cumulative signal-to-noise ratio, ${\\rm S}/{\\rm N} = \\sqrt{\\sum_z [{\\rm S}/{\\rm N}(z)]^2}$, is $\\sim 0.15 (f_{2}/3)$ if $f_{\\rm NL}=0$. In return, one can potentially probe $f_{\\rm NL}$ values as low as 20 roughly out to $l_{\\rm max}\\sim 10^3$.  With  the first-generation  radio interferometers, we would at most survey 1\\% of the sky. Assuming instrument noise is dominating at multipoles above $10^3$ between 30 MHz and 60 MHz at 1 MHz bandwidths (corresponding to redshifts 30 to 100), we find a signal-to-noise ratio of $5 \\times 10^{-3} f_{2}$, which could lead to a limit on $f_2$ of order 200.  Above discussion on the application of $X_l^{\\rm prim}$ assumes that $f_{\\rm NL}=0$. Since $X_l^{\\rm  prim} \\propto f_{\\rm NL}^2$, if $f_{\\rm NL}$ is greater than one, the overall signal-to-noise ratio for the detection of the three-to-one angular power spectrum is increased and the dominant contribution to $X_l^{\\rm prim}$ comes from the coupling associated with $f_{\\rm NL}$ and not $f_2$.  In the case when $f_2=0$, the $f_{\\rm NL}$ one probes with the trispectrum can be related to $\\tau_{\\rm NL}$ of Refs.~\\cite{Seery,Lyth} for the primordial trispectrum. In Fig.~2, we show the signal-to-noise ratio with $f_{\\rm NL}=10$ and $f_2=3$. Since in this case $f_{\\rm NL}$ term dominates, this provides an approximate estimate of the signal-to-noise ratio for $f_{\\rm NL}$ with the trispectrum. Using a single redshift bin out to $l_{\\rm max} \\sim 10^3$, 21-cm observations  achieve a signal-to-noise ratio of 1 if $f_{\\rm NL} \\sim 15$. This in return constraints $tau_{\\rm NL} \\lesssim 300$. Assuming 30 redshift bins over the redshift interval 50 to 100, assuming $f_2=0$, we find that one can constrain $\\tau_{\\rm NL} < 50$. This result is only out to $l_{\\rm max} =10^3$, but since 21-cm observations are not damped as in the case of CMB observations, higher resolution data can improve limits on both $f_2$ and $\\tau_{\\rm NL}$ significantly especially if observations can be pushed to $l > 10^4$.  While a detection of the CMB angular trispectrum has been motivated as a way to constrain $f_{\\rm NL}$ or $\\tau_{\\rm NL}$ \\cite{Kogo:2006}, this is probably not necessary with 21-cm data. Once radio interferometers start probing the redshift interval of 50 to 100, the angular bispectrum, which can be probed with the two-to-one angular power spectrum \\cite{Cooray}, can limit $f_{\\rm NL} < 0.1$. This will facilitate a separation of the contribution to the three-to-one power spectrum from $f_{\\rm NL}$ and $f_2$. While we have assumed that $f_{\\rm NL}$ and $f_2$ are momentum independent, it is likely that for specific models of inflation, these coupling terms are momentum dependent \\cite{Bartolo2} and then the ability to separate $f_{\\rm NL}$ and $f_2$ by combining 21-cm bispectrum and trispectrum information with the two-to-one and three-to-one correlator respectively will strongly depend on the exact momentum dependence of the coupling factors. We leave such a study for future research. While measuring $f_2$ is well motivated, one can also test the consistency between $f_{\\rm NL}$ probed by the bispectrum and the $\\tau_{\\rm NL}$ from the trispectrum related to the slow-roll inflation predictions for the non-Gaussianity. While we have not discussed the extent to which this consistency relation can be established with upcoming experiments after taking into account of instrumental noise it will also be useful to return to such  a calculation in the future once 21-cm observations begin to probe the universe at $z > 10$.  \\begin{figure}[t] \\includegraphics[scale=0.5,angle=-90]{Fig2.eps} \\caption{ Signal-to-noise ratio for a detection of $X_l^{\\rm prim}$ with $f_{\\rm NL}=10$ and $f_2=3$ (solid line) and $f_{\\rm NL}=0$ and $f_2=3$ (dot-dashed line). For reference, we show the signal-to-noise ratio associated with the detection of the full 21-cm bispectrum with a dashed line. } \\label{sn} \\end{figure}     The 21-cm anisotropies from the neutral hydrogen distribution prior to the era of reionization is expected to be more sensitive to primordial non-Gaussianity than the cosmic microwave background due to both the three-dimensional nature of the 21-cm signal and the lack of a damping tail at arcminute angular scales. Previous calculations have discussed the extent to which 21-cm bispectrum can be used as a probe of  primordial non-Gaussianity at the two-point level with an non-Gaussianity parameter $f_{\\rm NL}$ \\cite{Cooray,pillepich}.  Here, we extend these calculations to discuss the possibility to use a four point statistic of the 21-cm background as a probe of the primordial non-Gaussianity associated with the trispectrum.  We have calculated the angular trispectrum of the 21-cm background anisotropies and have introduced th three-to-one correlator and the corresponding angular power spectrum as a probe of the primordial trispectrum captured by both $f_{\\rm NL}$ and the third order non-Gaussianity parameter $f_2$ (described in some publications as $g_{\\rm NL}$). Since the primordial non-Gaussianity is confused with a non-Gaussian signal in the 21-cm background generated by the non-linear evolution of the density perturbations under gravitational evolution, we have discussed  way to separate the two using an optimal filter. While with the angular bispectrum of 21-cm anisotropies one can limit the second order corrections to the primordial fluctuations as low as $f_{\\rm NL}\\sim 0.1$ below the value of $\\sim 1$ expected for inflationary models, using the trispectrum information we suggest that one can constrain third order coupling term $f_2$ to about few tens. If $f_{\\rm NL}$ is large, it could potentially be possible to test the consistency between $f_{\\rm NL}$ from the bispectrum and the slow-roll non-Gaussianity $\\tau_{\\rm NL}$ at the four-point level with the relation $\\tau_{\\rm NL}=(6f_{\\rm NL}/5)^2$. We hope to return to such a detailed study in the future."
    },
    "0801/0801.4106.txt": {
        "abstract": "} }% %BeginExpansion \\begin{abstract} %EndExpansion The ratio of the Bondi and Jeans lengths is used to develop a cloud-accretion model that describes both an inner Bondi-type regime where gas pressure is balanced by the gravity of a central star and an outer Jeans-type regime where gas pressure is balanced by gas self-gravity. The gas density profile provided by this model makes a smooth transition from a wind-type inner solution to a Bonnor-Ebert type outer solution. It is shown that high-velocity dust impinging on this cloud will tend to pile-up due to having a different velocity profile than gas so that the dust-to-gas ratio is substantially enriched above the 1\\% ISM\\ level. %TCIMACRO{\\TeXButton{End abstract}{ ",
        "introduction": "Laboratory-scale plasma jets have been produced by a magnetohydrodynamic mechanism believed analogous to the mechanism responsible for driving astrophysical jets %TCIMACRO{\\TeXButton{\\citet{Hsu2002,Bellan2005}}{\\citep{Hsu2002,Bellan2005}}}% %BeginExpansion \\citep{Hsu2002,Bellan2005}% %EndExpansion . The laboratory jets are driven by capacitor bank power supplies that provide poloidal and toroidal magnetic fields and the jet acceleration mechanism can be considered as due to the pressure of the toroidal magnetic field inflating flux surfaces associated with the poloidal magnetic field. If these jets are indeed related to astrophysical jets, the obvious question arises as to what constitutes the power supply responsible for the toroidal and poloidal magnetic fields in the astrophysical situation. Existing models of astrophysical jets typically assume the poloidal field is simply given and that the toroidal field results from the rotation of an accretion disk twisting up the assumed poloidal field. The author has developed an alternate model which postulates that the toroidal and poloidal field result instead from a dusty plasma dynamo mechanism that converts the gravitational energy of infalling dust grains into an electrical power source that drives poloidal and toroidal electric currents creating the respective toroidal and poloidal fields. A brief outline of how infalling charged dust can drive poloidal currents has been presented in %TCIMACRO{\\TeXButton{\\citet{Bellan2007}}{\\citet{Bellan2007}}}% %BeginExpansion \\citet{Bellan2007}% %EndExpansion .  An important requirement of the model is that there should be sufficient infalling dust to provide the jet power. It has been well established that the dust-to-gas mass ratio in the Interstellar Medium (ISM) is 1\\%. If one assumes, as has been traditional, that this ratio holds throughout the accretion process, then the gravitational energy available from infalling dust would be insufficient. However, %TCIMACRO{\\TeXButton{\\citet{Fukue2001}}{\\citet{Fukue2001}} }% %BeginExpansion \\citet{Fukue2001} %EndExpansion has recently shown via numerical solution of coupled dust and gas equations of motion that the dust-to-gas ratio can become substantially enriched during Bondi-type accretion.  Star formation has also been previously examined from a molecular cloud physics point of view which differs significantly from the Bondi accretion point of view in the treatment of self-gravity and inflows. Molecular cloud physics is clearly important because observations indicate that stars form within the dense cores of molecular clouds (e.g., see %TCIMACRO{\\TeXButton{\\citet{Evans2001}}{\\citet{Evans2001}}}% %BeginExpansion \\citet{Evans2001}% %EndExpansion ). Analysis of the force balance in molecular clouds shows that the cloud radial density profile can be characterized by the Bonnor-Ebert sphere solution %TCIMACRO{\\TeXButton{\\citet{Ebert1955,Bonnor1956}}{\\citep{Ebert1955,Bonnor1956}% %}}% %BeginExpansion \\citep{Ebert1955,Bonnor1956}% %EndExpansion . This solution does not take into account flows associated with accretion. Dust dynamics is also not taken into account; instead it is typically assumed that the dust-to-gas mass ratio is fixed at 1\\% for all radii.  To summarize the above discussion, we note that the Bondi-type analysis reported by %TCIMACRO{\\TeXButton{\\citet{Fukue2001}}{\\citet{Fukue2001}} }% %BeginExpansion \\citet{Fukue2001} %EndExpansion emphasizes inflow physics and dust-gas coupling but does not take into account self-gravity whereas molecular cloud analysis such as used by %TCIMACRO{\\TeXButton{\\citet{Evans2001}}{\\citet{Evans2001}} }% %BeginExpansion \\citet{Evans2001} %EndExpansion emphasizes self-gravity, but does not take into account inflow or dust-gas coupling.  This paper will address the physics of coupled dust and gas accretion using a methodology similar in concept to that presented by %TCIMACRO{\\TeXButton{\\citet{Fukue2001}}{\\citet{Fukue2001}}}% %BeginExpansion \\citet{Fukue2001}% %EndExpansion , but extended to bridge the gap between Bondi accretion models and molecular cloud force balance models. \\ The enrichment mechanism observed by %TCIMACRO{\\TeXButton{\\citet{Fukue2001}}{\\citet{Fukue2001}} }% %BeginExpansion \\citet{Fukue2001} %EndExpansion will be examined in detail and will be shown to result from the inherently non-uniform nature of the Bondi/molecular cloud/ISM system. Properties of dust and gas for radii ranging from the cloud-ISM interface to the inner Bondi region will be considered by examining a sequence of successively smaller concentric regions. The basic character of each region and its scale, defined in terms of the nominal radial distance from a central object star, is as follows:  \\begin{description} \\item \\textit{ISM scale:} The outermost scale is that of the Interstellar Medium (ISM). The ISM\\ has a gas density $\\sim$10$^{7}$ m$^{-3}$, a dust-to-gas mass ratio of 1 percent, a gas temperature $T_{g}^{ISM}\\sim100$ K$,$ and is optically thin. The ISM is assumed to be spatially uniform and to bound a molecular cloud having radius $r_{edge}$.  \\item \\textit{Molecular cloud scale:}\\ The molecular cloud scale has much higher density than the ISM and is characterized by force balance between gas self-gravity\\ and gas pressure. The molecular cloud scale is sub-divided into a large, radially non-uniform low-density outer region and a small, approximately uniform, high-density inner core region. Clouds have a characteristic scale given by the Jeans length $r_{J}.$ The radial dependence of density is provided by the Bonnor-Ebert sphere solution which acts as the outer boundary of the Bondi accretion scale.  \\item \\textit{Bondi accretion scale:} The Bondi accretion scale %TCIMACRO{\\TeXButton{\\citep{Bondi1952}}{\\citep{Bondi1952}} }% %BeginExpansion \\citep{Bondi1952} %EndExpansion is $\\sim r_{B}$ which is sufficiently small that gas self-gravity no longer matters so equilibrium is instead obtained by force balance between gas pressure and the gravity of a central object assumed to be a star having mass $M_{\\ast}$. The Bondi scale is sub-divided into three concentric radial regions: an outermost region where the gas flow is subsonic, a critical transition radius at exactly $r_{B}$ where the flow is sonic, and an innermost region where the gas flow is free-falling and supersonic.  \\item \\textit{Collisionless dusty plasma scale (to be considered in a future publication): } Free-falling dust grains collide with each other in one of the above scales and coagulate to form large-radius grains which are collisionless and optically thin. The optically thin dust absorbs UV photons from the star, photo-emits electrons and becomes electrically charged. The charged dust grains are subject to electromagnetic forces in addition to gravity. Motions of charged dust grains relative to electrons result in electric currents with associated poloidal and toroidal magnetic fields [see preliminary discussion in %TCIMACRO{\\TeXButton{\\citet{Bellan2007}}{\\citet{Bellan2007}}}% %BeginExpansion \\citet{Bellan2007}% %EndExpansion ]. \\ \\  \\item \\textit{Jet scale (to be considered in a future publication): }The electric currents interact with the magnetic fields to produce magnetohydrodynamic forces which drive astrophysical jets in a manner analogous to that reported in \\ %TCIMACRO{\\TeXButton{\\citet{Hsu2002}}{\\citet{Hsu2002}} }% %BeginExpansion \\citet{Hsu2002} %EndExpansion and %TCIMACRO{\\TeXButton{\\citet{Bellan2005}}{\\citet{Bellan2005}}}% %BeginExpansion \\citet{Bellan2005}% %EndExpansion . \\end{description}  \\bigskip The separation-of-scales requirement $r_{J}>>r_{B}$ implies existence of a small parameter% \\begin{equation} \\varepsilon_{BJ}\\ =\\frac{r_{B}}{r_{J}}\\ \\label{eps}% \\end{equation} quantifying the separation between the Bondi and Jeans scales. For purposes of relating the Bondi and Jeans scales to each other it is convenient to introduce a geometric-mean scale with characteristic length given by \\begin{equation} r_{gm}=\\sqrt{r_{B}r_{J}}=\\sqrt{\\varepsilon_{BJ}}r_{J}\\ \\label{def rgm}% \\end{equation} in which case% \\begin{equation} r_{B}\\ll r_{gm}\\ll r_{J}. \\label{separation}% \\end{equation}   Gas and dust must be considered separately in each scale. We assume that gas motion influences dust motion but not vice-versa so that the dust dynamics can be ascertained after gas dynamics has been determined. This assumption is appropriate so long as dust mass and energy densities are small compared to corresponding gas densities or if the dust is decoupled from the gas. The sequence \\ of scales is summarized in Tables \\ref{gas table} and \\ref{dust table}. The numerical values of table elements with asterisks will be predicted by the model to be presented here while table elements with filled-in numbers represent\\ prescribed physical boundary conditions. Here $m_{g}$ is the gas molecular mass, $n_{g}$ is the gas density, $T_{g}$ is the gas temperature, $u_{g,d}$ are the radially inward gas and dust fluid velocities, and $\\ $ $\\rho_{g,d}$ are the gas and dust mass densities. Quantities with superscripts `ISM' are evaluated in the ISM, un-superscripted quantities refer to Bondi or molecular cloud regions.  Our methodology will be conceptually similar to %TCIMACRO{\\TeXButton{\\citet{Fukue2001}}{\\citet{Fukue2001}} }% %BeginExpansion \\citet{Fukue2001} %EndExpansion but will also differ in important ways. %TCIMACRO{\\TeXButton{\\citet{Fukue2001}}{\\citet{Fukue2001}} }% %BeginExpansion \\citet{Fukue2001} %EndExpansion used a wind equation scheme to model Bondi-type physics and, unlike the analysis to be presented here, assumed the Bondi accretion region was directly bounded by the ISM, i.e., no molecular cloud region with Jeans-type scaling was taken into account. In addition, Fukue assumed that (i) the gas was heated by the combined effects of adiabatic compression and friction due to dust-gas collisions, and (ii) the entire region was optically thin so that the dust was subject to radiation pressure. Because of the complexity introduced by the heating of the gas, Fukue obtained results via numerical solution of three coupled differential equations (gas momentum, dust momentum, gas heating). Our approach will differ by assuming that (i) the gas is isothermal, (ii) a molecular cloud region lies between the Bondi accretion region and the ISM, and (iii) the system is optically thick outside the inner part of the Bondi region so that dust is shielded from optical radiation until it penetrates to the inner part of the Bondi region. Furthermore, rather than using numerical solutions, we will attempt analytic solutions as much as possible, a goal made feasible by the isothermal assumption. We believe our assumptions (i)-(iii) reasonably correspond to observations that gas is nearly isothermal and that stars form inside the cores of optically thick molecular clouds.  \\bigskip% %TCIMACRO{\\TeXButton{B}{\\begin{table}[tbp] \\centering}}% %BeginExpansion \\begin{table}[tbp] \\centering %EndExpansion% \\begin{tabular} [c]{lllll}\\hline\\hline region & location & $n_{g}$ & $T_{g}$ & $u_{g}$\\\\\\hline\\hline ISM & $r>r_{edge}$ & 10$^{7}$ m$^{-3}$ & 100 K & *\\\\ B-E sphere & $r_{gm}<r<r_{edge}$ & * & 10 K & *\\\\ Bondi-Jeans interface & $r_{gm}=\\sqrt{r_{B}r_{J}}$ & * & \\multicolumn{1}{c}{\"} & *\\\\ Bondi subsonic & $r_{B}<r<r_{gm}$ & * & \\multicolumn{1}{c}{\"} & $<\\sqrt{\\kappa T_{g}/m_{g}}$\\\\ Bondi transonic & $r=r_{B}$ & * & \\multicolumn{1}{c}{\"} & $\\sqrt{\\kappa T_{g}/m_{g}}$\\\\ Bondi supersonic & $r<r_{B}$ & * & \\multicolumn{1}{c}{\"} & $>\\sqrt{\\kappa T_{g}/m_{g}}$% \\end{tabular} \\caption{Sequence of regions for gas, * indicates quantity to be discussed/calculated in text}\\label{gas table}% %TCIMACRO{\\TeXButton{E}{\\end{table}}}% %BeginExpansion \\end{table}% %EndExpansion %  %TCIMACRO{\\TeXButton{B}{\\begin{table}[tbp] \\centering}}% %BeginExpansion \\begin{table}[tbp] \\centering %EndExpansion% \\begin{tabular} [c]{llll}\\hline\\hline region & location & $\\rho_{d}/\\rho_{g}$ & $u_{d}$\\\\\\hline\\hline ISM & $r>r_{edge}$ & 0.01 & 3 km/s (turbulent acceleration)\\\\ B-E sphere & $\\sqrt{r_{B}r_{J}}<r<r_{edge}$ & * & slowing down by gas\\\\ Bondi-Jeans interface & $r_{gm}=\\sqrt{r_{B}r_{J}}$ & * & entrained with gas\\\\ Bondi subsonic & $r_{B}<r<r_{gm}$ & * & \\multicolumn{1}{c}{\"}\\\\ Bondi transonic & $r=r_{B}$ & * & \\multicolumn{1}{c}{\"}\\\\ Bondi supersonic & $r<r_{B}$ & * & \\multicolumn{1}{c}{\"}% \\end{tabular} \\caption{Sequence of regions for dust physics, * indicates to be discussed/calculated in text}\\label{dust table}% %TCIMACRO{\\TeXButton{E}{\\end{table}}}% %BeginExpansion \\end{table}% %EndExpansion ",
        "conclusions": "The logical requirement that the Bondi length must be smaller than the Jeans length introduces the dimensionless parameter $\\varepsilon_{BJ}$ given by the ratio of these lengths. An intermediate distance $r_{gm}$ can be defined as the geometric mean of the Bondi and Jeans lengths so, if the two lengths are well separated, $r_{gm}$ is much larger than the Bondi length but much smaller than the Jeans length. Thus, $r_{gm}$ constitutes the large $r$ limit of the solution to the Bondi problem and the small $r$ limit of the solution to the Jeans problem. Matching the Bondi and Jeans solutions at $r_{gm}$ provides a solution valid over scales ranging from smaller than the Bondi length to larger than the Jeans length. This allows accounting for transonic flow, mass accretion, and gravitational self-confinement in one self-consistent solution to the gas equation.  By assuming that the dust mass density and energy density are not larger than the corresponding gas densities, the gas equations can be evaluated ignoring interaction with dust. Once the gas behavior has been worked out, the dust can be considered as moving through a pre-determined gas profile that retards the dust due to frictional drag. \\ The dust slows down greatly whereas the gas velocity changes relatively little, a distinction that enriches the dust-to-gas mass ratio. Plausible values of $\\varepsilon_{BJ}$ suggest that the dust-to-gas mass ratio will always be enriched to be somewhat less than unity, because at unity the assumption that the gas is unaffected by dust fails. \\ The high dust density at small radius means that dust-dust collisions will become important and cause coagulation of dust. This coagulation will cause the dust to become optically thin and collisionless again. Because the coagulated dust is optically thin, it will absorb photons from the star and become electrically charged due to photo-emission of electrons. Two-dimensional motion of collisionless, charged dust in a gravitational field resulting in a poloidal field dynamo has been discussed in %TCIMACRO{\\TeXButton{\\citet{Bellan2007}}{\\citet{Bellan2007}}}% %BeginExpansion \\citet{Bellan2007}% %EndExpansion ; three dimensional motion and resulting poloidal/toroidal field dynamo action will be discussed in a future publication.  The author wishes to thank an anonymous referee for his/her thoughtful comments.  \\bigskip"
    },
    "0801/0801.3525_arXiv.txt": {
        "abstract": "Low-resolution spectra from 3000--9000 \\AA\\ of young low-mass stars and brown dwarfs were obtained with LRIS on {\\it Keck I}.   The excess UV and optical emission arising in the Balmer and Paschen continua  yields mass accretion rates ranging from $2\\times10^{-12}$ to $10^{-8}$ \\mdotyr.  These results are compared with {\\it HST}/STIS spectra of roughly solar-mass accretors with accretion rates that range from $2\\times10^{-10}$ to $5\\times 10^{-8}$ $M_\\odot$ yr$^{-1}$.  The weak photospheric emission from M-dwarfs at $<4000$ \\AA\\ leads to a higher contrast between the accretion and photospheric emission relative to higher-mass counterparts. The mass accretion rates measured here are systematically $\\sim 4-7$ times larger than those from H$\\alpha$ emission line profiles, with a difference that is consistent with but unlikely to be explained by the uncertainty in both methods. The accretion luminosity correlates well with many line luminosities, including high Balmer and many \\ion{He}{1} lines.  Correlations of the accretion rate with H$\\alpha$ 10\\% width and line fluxes show a large amount of scatter.  Our results and previous accretion rate measurements suggest that $\\dot{M}\\propto M^{1.87\\pm0.26}$ for accretors in the Taurus Molecular Cloud. ",
        "introduction": "In the magnetospheric accretion model for accretion onto young stars, the stellar magnetosphere truncates the disk at a few stellar radii.  Gas from the disk accretes onto the star along magnetic field lines and shocks at the stellar surface.  The $\\sim 10^4$ K optically-thick post-shock and optically-thin pre-shock gas produce emission in the Balmer and Paschen continuum and in many lines, including the Balmer and Paschen series and the \\ion{Ca}{2} IR triplet \\citep{Har94,Cal98}.  The accretion luminosity may be estimated from any of these accretion tracers and subsequently converted into a mass accretion rate.   Historically, mass accretion rates onto young stars have been measured directly using optical veiling from the Paschen continuum and stronger Balmer continuum emission shortward of the Balmer jump.  Modelling of the Paschen and Balmer continuum emission with a generic isothermal, plane-parallel slab yields $\\dot{M}\\sim 10^{-7}-10^{-10}$ $M_\\odot$ yr$^{-1}$ for accretors with masses of 0.4-1.0 $M_\\odot$  \\citep{Val93,Gul98}. \\citet{Cal98} obtained similar accretion rates by applying more realistic shock geometries with a range of temperatures to explain the observed accretion continuum.  The accretion luminosity can be measured instead at longer wavelengths with optical veiling measurements from high-dispersion spectra \\citep[e.g.][]{Bas90,Har91,Har95,Har03,Whi04}, albeit with a larger uncertainty in bolometric correction.  Several deep surveys have since identified large populations of young very low-mass stars and brown dwarfs with masses $<0.2$ $M_\\odot$ \\citep[e.g.][]{Bri98,Luh04,Sle06,Gui07}.  Excess IR emission indicates the presence of disks around some of these very low-mass objects$^2$ \\citep[e.g.][]{Muz03,Jay03,Nat04}. The excess Balmer continuum emission produced by accretion onto very low-mass stars and brown dwarfs has not been observed previously. Instead, accretion has been been measured based on magnetospheric models of the H$\\alpha$ line profile (Muzerolle et al. 2000, 2003, 2005), by extending relationships between $\\dot{M}$ and line luminosities established at higher masses to brown dwarf masses (Natta et al. 2004, 2006; Mohanty et al. 2005), or in a few cases from measurements of the Paschen continuum \\citep{Whi03}. \\footnotetext[2]{The objects studied in this paper span the mass range from 0.024--0.2 $M_\\odot$; hereafter all are referred to as ``stars'' for ease of terminology.  We refer to $\\sim 0.3-1$ $M_\\odot$ stars as ``higher-mass stars.''}   In this paper, we present the first observations of the Balmer jump, and hence direct measurements of the accretion luminosity, for young very low-mass stars and brown dwarfs using data obtained with LRIS on {\\it Keck I}.  We compare the Balmer and Paschen continuum emission from very low-mass stars to that from higher-mass accretors that were observed with {\\it HST}. The weak U-band photospheric emission of M-dwarfs leads to a high contrast between accretion and photospheric emission at $\\lambda<3700$ \\AA.  The properties of the Balmer continuum and jump for very low-mass stars and brown dwarfs are mostly similar to those of the higher-mass accretors, although the sensitivity to low $\\dot{M}$ improves to smaller stellar mass.  The Balmer jump tends to be larger for smaller mass accretion rates, indicating lower optical depth for the emitting gas. The luminosity in many lines, including the Balmer series and \\ion{He}{1} lines, correlate well with the accretion luminosity. Our $\\dot{M}$ estimates are systematically $\\sim 4-7$ times higher than those estimated from modelling the H$\\alpha$ line profile.  We find that $\\dot{M}\\propto M^{1.9}$ for accretors in Taurus and discuss biases in the selection of those accretors that may affect this calculated relationship. ",
        "conclusions": "We analyzed blue spectra of 18 accreting very low-mass stars and brown dwarfs obtained with LRIS on {\\it Keck I}.  Ours are the first data on the Balmer continuum and high Balmer line emission in this mass range.   Most of this sample was selected based on the presence of accretion identified with other diagnostics.  Our observations are compared with archival {\\it HST}/STIS spectra of the Balmer continuum from several higher-mass stellar accretors.  We obtain the following results:  1.  We detect excess hydrogen Balmer continuum emission from 16 of the 18 mid-late M-dwarfs in our sample.  11 of these objects also show hydrogen Paschen continuum emission.  This continuum is attributed to accretion for most of our objects.  Weak Balmer continuum emission onto earlier-type stars, including the K6 dwarf V836 Tau, can be attributed to accretion because non-accretors show no evidence for chromospheric Balmer continuum emission.  However, we cannot rule out a chromospheric origin for weak excess Balmer emission from mid-M dwarfs.  2.  Measurements of the Balmer and Paschen continua emission yield $\\dot{M}=10^{-12}-10^{-8}$ for very-low mass stars and brown dwarfs with masses from 0.024--0.17 $M_\\odot$.  Most of these accretion rates are smaller than those typically measured for accreting $\\sim 0.3-1.0$ $M_\\odot$ stars.  Analysis of the random and systematic errors indicate that these $\\dot{M}$ may be underestimated by a factor of 1.2 (or 2.4 if the hydrogen continuum emission incident upon the star is reprocessed into photospheric emission), with a relative uncertainty of $\\sim 0.6$ dex.  Comparing these $\\dot{M}$ with $\\dot{M}$ from other studies requires accounting for different assumptions regarding radiative transfer.  3. The Balmer jump tends to be larger and the Balmer continuum tends to be redder for stars with lower mass accretion rates.  These properties indicate lower opacity (with $\\log n_e=13.5$ for 2M1207, compared with $\\log n_e\\sim14.3$ for stars with higher $\\dot{M}$ in the plane-parallel slab models) and lower temperature (with $T\\sim7000$ K for 2M1207, compared with $\\sim 9500$ for stars with higher $\\dot{M}$).  4.  The $\\dot{M}$ calculated here are well correlated with but $\\sim 4-7$ times larger than the $\\dot{M}$ calculated from published models of the H$\\alpha$ line profiles.   We suggest that this offset is methodological, although the difference is within the range of the combined errors for both methods.  The $\\dot{M}$ measured here also correlates with H$\\alpha$ 10\\% width but introduces 0.9 dex uncertainty in $\\dot{M}$ predicted from H$\\alpha$ 10\\% width alone.  5.   We provide empirical relationships between $L_{acc}$ and $L_{line}$, which are often well correlated.  In particular, estimating $L_{acc}$ from the luminosity in the high Balmer and many \\ion{He}{1} lines introduces an uncertainty of only 0.2--0.35 dex.  On the other hand, estimating the $L_{acc}$ from line surface fluxes alone introduces an uncertainty of $\\sim 0.7-1.1$ dex in $\\dot{M}$ estimates.   6.  We find that $\\log \\dot{M}\\propto M^{1.87}$ for Taurus stars with $A_V<2.5$, which is consistent with previous estimates.  This relationship may be biased if the higher sensitivity of ${\\dot{M}}/M$ at lower $M$ causes the percentage of accretors that are detected to also vary with $M$."
    },
    "0801/0801.2465_arXiv.txt": {
        "abstract": "{ The magnetar model involves an isolated neutron star with a very high magnetic field ($B\\sim 10^{14-15}$ G), and is invoked to explain the emission processes of two classes of sources, the Anomalous X-ray Pulsars (AXPs) and the Soft Gamma-Ray Repeaters (SGRs). Five of them have been recently identified to be persistent sources in the hard X-ray band (20--200 keV). AXPs, in particular, present the hardest known persistent spectra in the hard X/soft $\\gamma$-ray energy range. The broad band modeling of their spectra still suffers from the non-simultaneity of the observations and from a lack of sensitivity above 20 keV. We present the Simbol X simulated observations of these objects and show that that this mission could surely help to disentangle the contribution of the different spectral components, and to understand how they contribute to the secular flux variations observed in these sources. ",
        "introduction": "Most of the known neutron stars (NSs) are either isolated rotation powered pulsars, or accretion powered neutron stars hosted in binary systems. A dozen (plus a few candidates) of sources, dubbed magnetars, do not fit in either of these categories, since their dominant source of energy is believed to be the magnetic one. In fact, in the magnetar model \\citep{dt92} it is the decay of the huge magnetic field ($B\\sim 10^{14-15}$ G) of isolated neutron stars that powers their electromagnetic emission.  Magnetars are historically divided in two categories, due to the fact that they were discovered in different ways. We will briefly summarize their properties here. For a complete review of this kind of objects see \\citet{woodsrew}.  \\subsection{Soft Gamma-Ray Repeaters}  SGRs were originally identified in the late '70s a subclass of Gamma Ray Bursts, due to the fact that these bursts had softer spectra and originated repeatingly from fixed positions across the sky. Four confirmed SGRs are known to date, three of them (1806--20, 1900+14, and 1627--41) are located in our Galaxy at several kpc, while one, 0525--66, is located in the Large Magellanic Cloud. They sporadically emit short ($\\sim$0.1 s) and bright (L$\\sim$ 10$^{39}$-10$^{42}$ ergs s$^{-1}$) bursts of soft gamma-ray radiation during active periods, alternating with quiescent periods that can last several years. In rare occasions they emit so-called {\\em giant flares}, which are characterized by very brights spikes (L$\\sim$ 10$^{44}$-10$^{46}$ ergs) lasting a fraction of a second, followed by a pulsating tail lasting a few hundred seconds. Up to know three out of four SGRs have emitted a giant flare, and the pulsations found in their tails pushed forward the idea that these objects were associated with neutron stars, with an extremely high magnetic field. In fact, the energy required to confine a fireball for the duration of the giant flare pulsating tail strongly points towards a magnetic field of the order of 10$^{15}$ G.  Three of the confirmed SGRs have quiescent pulsating X-ray counterparts for which the period and its derivative could be measured. Assuming that this radiation is due to magnetic braking in a dipolar surface magnetic field, these values confirm the high magnetic field, and the absence of Doppler modulations in the observed light curves exclude the presence of companion stars.   \\subsection{Anomalous X-ray Pulsars}  AXPs have been originally identified as a class due to their common X-ray and timing properties \\citep{merestella}. Their rotational periods cluster in the 5-12 s range, their period derivatives in the 0.05-4$\\times$10$^{-11}$ s s$^{-1}$. Their X-ray luminosity is L$_{X}\\sim$10$^{34-36}$ erg s$^{-1}$ . From their timing porterties one can derive the dipole magnetic field at the surface of the NS as follows,  \\begin{equation} B=\\left(\\frac{3Ic^{3}P\\dot P}{2\\pi^{2}R^{6}}\\right)\\simeq3.2\\times10^{19}(P\\dot P)^{1/2} \\; \\rm G, \\end{equation}  where $I$ ($\\simeq$10$^{45}$g cm$^{2}$) is the NS moment of inertia, and $R$ ($\\simeq$10$^{6}$ cm) is the NS radius. The derived field values for AXPs exceed the quantum critical value of $B_{Q}\\equiv m_{e}^{2}c^{3}/(e\\hbar)=4.4\\times10^{13}$ G. This fact together with the lack of an observable companion star and with the fact that $L_{X}$ largely exceeds the rotational energy loss ($I\\omega\\dot\\omega$), points towards the inclusion of the AXPs in the magnetar class. In addition the detection of Soft Gamma-Ray Repeater-like bursts from five AXPs \\citep[e.g.][]{gavrill} has strengthened the association between these objects and the SGRs. ",
        "conclusions": "One of the major issues today in the magnetar's field is that no clear physical model has been developed yet, in order to explain the broad band spectrum of AXPs and SGRs. An important reason for this is that with the current instrumentation it is not possible to obtain good quality broad band (0.1--200 keV) spectra. For every astrophysical object the broad band coverage is nowadays an important issue, but the unexpected discovery of very hard tails in AXPs is a unique feature, and deserves particular attention in this sense.  We have shown that with Simbol-X we will be able to monitor for the first time all the spectral components of these variable objects at once. In particular we will be able to pinpoint the differences between AXPs and SGRs, that before the IBIS/ISGRI hard X-ray observations looked very similar for what concerns their persistent emission. We will also be able to understand what is the driver of spectral changes in magnetars, i.e. if the changes below 10 keV are in reality caused by the variation of an underlying high-energy component.  Concerning the timing, pulsations in the hard X-ray band have been detected for a few AXPs and possibly for \\zerosei~ (G\\\"otz et al., 2007, in preparation). Their detection needs long integration times, and the combination of different instruments on different satellites, which introduces unavoidable uncertainties in the estimation of the pulsed fraction. Simbol -X, being a focusing telescope with low background at hard X-rays, will surely reduce by an order of magnitude our uncertainties and the exposure times required. In addition it will allow us to do broad band phase resolved spectroscopy.  The detection of resonant cyclotron absorption lines would be the direct confirmation of the presence of  a huge magnetic field in magnetars. Up to now, despite the deep searches below 10 keV, no firm detection of these features has been reported. One interesting feature around 6--15 keV has been reported by \\citet{iwasawa} using {\\it Ginga} data, during the outburst of  the AXP 1E 2259+586. Being this feature close to the upper energy boundary of the instrument, its detection is questionable. Simbol-X could clearly settle an issue like that.  A few of the AXPs are transients: this is a new subclass of magnetars that undergo large outbursts lasting years where their flux changes by more than 3 orders of magnitude \\citep[e.g.][]{cxo}. The mechanisms of these long outbursts is not known, and in addition due to several constraints, these sources could not be observed efficiently during their outbursts with hard X-ray telescopes. Such an outburst (but also the quiescent phase) observed at hard X-rays with Simbol-X could surely help to shed some light of the mechanisms that drive these long lived outbursts."
    },
    "0801/0801.2186_arXiv.txt": {
        "abstract": "{New eclipse minimum timings of the M4.5/M4.5 binary CM Dra were obtained  between the years 2000 and 2007. In combination with published timings going back to 1977, a clear non-linearity in observed-minus-calculated (O-C) times has become apparent. Several models are applied to explain the observed timing behavior.} {Revealing the processes that cause the observed O-C behavior, and testing the evidence for a third body around the CM Dra system.} {The O-C times of the system were fitted against several functions, representing different physical origins of the timing variations.} {An analysis using model-selection statistics gives about equal weight to a parabolic and to a sinusoidal fitting function. Attraction from a third body, either at large distance in a quasi-constant constellation across the years of observations or from a body on a shorter orbit generating periodicities in O-C times is the most likely source of the observed O-C times. The white dwarf GJ 630.1B, a proper motion companion of CM Dra, can however be rejected as the responsible third body. Also, no further evidence of the short-periodic planet candidate described by Deeg et al. (2000) is found, whereas other mechanisms, such as period changes from stellar winds or Applegate's mechanism can be rejected.} {A third body, being either a few-Jupiter-mass object with a period of 18.5$\\pm$4.5 years or an object in the mass range of $1.5M_\\mathrm{jup}$ to $0.1M_{\\sun}$ with periods of hundreds to thousands of years is the most likely origin of the observed minimum timing behavior.} ",
        "introduction": "\\object{CM Dra} (\\object{LP 101.15}, \\object{G225-067}, \\object{GJ 630.1}) is a  detached spectroscopic eclipsing M4.5/M4.5 binary with one of the lowest known total masses, of $0.44 M_{\\sun}$. With its nearly edge-on inclination of 89.59$\\degr$(see \\citealt{koz+04} for the most recent orbital and physical elements) it was chosen as the target of the first photometric search for planetary transits. Performed by the 'TEP' project, with an intense observing campaign during the years 1994 - 1999 \\citep{ddk+98, ddk+2000}, over 1000 hours of coverage of that system were obtained with several 1m-class telescopes. This lightcurve was initially searched for the presence of transits from planets in circumbinary 'P-type' orbits with 5 - 60 day periods, with a negative result \\citep{ddk+2000}. The same lightcurve provided, however, a further possibility to detect the presence of third bodies, from their possible light-time effects on the binary's eclipse minimum times.  To date, no unambiguously circumbinary planets have been detected\\footnote{A possible circumbinary planet is \\object{HD202206c}, orbiting around HD202206a and b. This depends however on the classification of the 17.4M$_\\mathrm{jup}$ object \\object{HD202206b} as being a planet or a brown dwarf \\citep{correia+05}}, and their discovery would constitute  a new class of planets. Motivation for this work also arises from previous successes of precise timing measurements to detect the presence of planets. The first known extrasolar planets were detected through light-time effects in the signals of the Pulsar PSR1257+12 \\citep{wolyfrail92} and recently, sinusoidal residuals in the pulsation frequency of the sdB pulsating star V391 Peg have been explained through the presence of a giant planet \\citep{silvotti+07}, leading to the first detection of a planet orbiting a post-red-Giant star. In both cases, timing measurements have led to detections of planets that would have been difficult or impossible to find with other planet-detection methods, a situation that is similar to the detection of planets around eclipsing binaries -- unless they exhibit transits.  A first analysis of 41 eclipse minima times for CM Dra, presented in \\citet{deeg+00} (hereafter DDK00) gave a low-confidence indication for the presence of a planet of 1.5 - 3 Jupiter masses, with a period of 750 - 1050 days. This result was based on a power-spectral analysis of the minimum timings' residuals against a linear ephemeris, indicating a periodic signal with an amplitude of about 3 seconds. Motivated by the result of DDK00, we continued surveying the CM Dra system with occasional eclipse observations in the following years. During this time, it became increasingly clear that a simple linear ephemeris would not provide a sufficient description of the general trend of the eclipse times any longer. This led to the objective of this paper -- a thorough and systematic discussion of the possible processes acting on this interesting system, and an evaluation of the presence of a third body of planetary mass or heavier.  In the following sections, we present the data for this analysis (Sect.~\\ref{sec:data}), evaluate the effect of several physical processes on the eclipse timings (Sect.~\\ref{sec:processes}), and compare the significance of several numerical fits for the explanation of the observed trend (Sect.~\\ref{sec:stats}). This is followed in Sect.~\\ref{sec:disc} by a discussion of the physical implications of these fits, with conclusions in Sect.~\\ref{sec:conclusion}. ",
        "conclusions": "} The O-C timing shows a clear indication of a non-linear trend. After a review of potential causes (previous section), the remaining explanations are given by the presence of a third body. The nearly equal statistical weight of the parabolic and the sinusoidal fits currently prevents setting a clear preference for the one or the other model. For the period of the third body there exist two distinct possibilities: A Jupiter-type planet with $M \\sin i$ of 1-2 $M_\\mathrm{jup}$ with a period of $18.5\\pm4.5$ yr, or an object such as a giant planet or heavier, with a period of hundreds to thousands of years. Intermediate-length periods of about 25 - 100 years are less likely due to poor compatibility with either the sinusoidal or the polynomial fits. Regarding shorter periodic bodies, the power-spectrum shown in Section~\\ref{sec:powerspectra} excludes O-C modulations with amplitudes of larger than $\\approx 3.5$ seconds with periods of $\\la$10 years. The sensitivity of O-C timing detections against orbiting third bodies decreases with their period, and hence Jupiter-mass objects on such shorter periods cannot be excluded. A low-significance candidate for such an object with a period of about 900 days was presented in DDK00. That candidate was based on a single peak in a power-spectrum with an amplitude of 3 seconds, whereas the newer data show several peaks of amplitudes up to 3.5 seconds. The newer data therefore cannot refute that candidate, however it has become most likely to have been the result of a fortuitous combination of O-C timing values.  For long-period bodies at a quasi-constant distance and position during the observed time-span, Eq.~\\ref{eq:minmass} allows us to set a maximum lateral separation from CM Dra for any given third-body mass. We also assume that any nearby body larger than about $\\approx 0.1M_{\\sun}$ would have become apparent in existing images of the CM Dra field, for which a large collection of CCD images exist from the TEP project \\citep{ddk+98, ddk+2000}, or in 2Mass images in the IR. A maximum lateral distance of 6 arcsec, or 95AU may therefore be set for the presence of an as yet undiscovered third body of $\\la  0.1M_{\\sun}$. The corresponding maximum distances for undiscovered brown dwarfs or planets that could explain CM Dra's O-C timing behavior are 5 and 2 arcsec, respectively. While the setting of third-body minimum masses (and implied brightnesses) for a given lateral distance aids in the definition and interpretation of observing projects, we note however that the observed acceleration term may be caused by relatively small third-body masses. If we take 100 years as the shortest circular period that may mimic the quasi-stationary case, such an object's orbital distance would be 16.4 AU. If it is aligned such that $|\\mathbf{r}| \\approx r_{\\parallel}$ (e.g. $i$ close to $90\\deg$), then solving Eq.~\\ref{eq:accel} indicates a mass of about $1.5 M_\\mathrm{jup}$. This mass may be considered the absolute minimum mass for a third body at a quasi-stationary distance, with objects at larger distances requiring larger masses.  In conclusion, two possibilities for the source of CM Dra's timing variations remain valid: A mass of a few Jupiters on a two decade-long orbit, or an object on a century-to-millenium long orbit, with masses between 1.5 Jupiters and that of a very low mass star. Continued observations of the timing of CM Dra's eclipses over the next 5 - 10 years should, however, be decisive regarding the continued viability of the sinusoidal-fit model, and hence, about the validity of a jovian-type planet in a circumbinary orbit around the CM Dra system."
    },
    "0801/0801.2379_arXiv.txt": {
        "abstract": "{Optical nova lightcurves often have structures, such as rapid declines and recoveries, due to nebular or dusty phases of the ejecta. Nova Cygni 2006 (V2362 Cyg) underwent an unusual brightening after an early rapid decline.  The shape of the lightcurve can be compared to that of V1493 Aql, but the whole event in that case was not as bright and only lasted a couple of weeks.  V2362 Cyg had a moderately fast decline of $t_2 \\,=\\,9.0\\pm0.5$ days before rebrightening, which lasted 250 days after maximum.} {We present an analysis of our own spectroscopic investigations in combination with American Association of Variable Star Observers (AAVSO) photometric data covering the whole rebrightening phase until the return to the final decline.} {We used the medium resolution spectroscopy obtained in ten nights over a period of 79 nights to investigate the change of the velocity structure of the ejecta. The publicly available AAVSO photometry was used to analyze the overall properties and the energy of the brightening.} {Although the behavior of the main outburst (velocity, outburst magnitude, and decline timescales) resembles a ``normal'' classical nova, the shell clearly underwent a second fast mass ejecting phase, causing the unusual properties. The integrated flux during this event contributes $\\approx$40\\% to the total radiation energy of the outburst. The evolution of the H$_\\alpha$ profile during the bump event is obtained by subtracting the emission of the detached shells of the main eruption by a simple optically thin model. A distance of $D \\approx 7.5^{+3.0}_{-2.5}$~kpc and an interstellar extinction $E(B-V) = 0\\fm6\\pm0\\fm1$ was also derived.} {} ",
        "introduction": "NOVA Cygni 2006 (V2362 Cyg) was discovered April 2, 2006 by H. Nishimura as a 10\\fm5 object \\citep{discover}. The early spectra, obtained $\\approx 0\\fd6$ prior to the maximum in $V$, showed a rather flat continuum, on which the H$_\\alpha$ line shows a clear P-Cyg profile. The FWHM of the emission was 500 km s$^{-1}$ and the absorption minimum was 700 km s$^{-1}$ blueshifted \\citep{IAUC8698_1}. Within a day ($\\approx -0\\fd2$), the profiles showed two absorption features at --880 and --1\\,730~km~s$^{-1}$ \\citep{early_spec}. Based on spectra obtained on April 13, 2006 ($\\approx 6\\fd7$), \\citet{siviero} classified the nova as 'Fe-II' and found a very broad and structured profile with FWZI of 3\\,750~km~s$^{-1}$ and FWHM of 1\\,800~km~s$^{-1}$ from the Balmer emission lines.  With the exact position found by \\citet{position}, \\citet{progenitor} identified the progenitor as an emission line star in the IPHAS H$_\\alpha$ survey. With progenitor magnitudes of $r'=20\\fm3\\pm0\\fm05$ and $i'=19\\fm76\\pm0\\fm07$ the outburst amplitude was about 12 magnitudes.  In July 2006 ($\\approx 100^{\\rm d}$) the first unusual behavior was reported by \\citet{no_ir_dust}. The nova, after having already declined by 4$^{\\rm m}$, still showed only low excitation lines. The target brightened \\citep{munari_c} from August 2006 ($\\approx 130^{\\rm d}$) until beginning of December 2006 ($\\approx 240^{\\rm d}$) . \\citet{goranskij_1} found an optical period of 0\\fd2070 and an increase of the emission line width during this phase. They also pointed out similarities  to the lightcurve of V1493~Aql (Nova Aql 1999a). The second maximum in V1493~Aql had a shorter duration of 2 weeks and a smaller amplitude of only $\\approx$0\\fm75. The duration in V2362 Cyg was 6 months with an amplitude of 2\\fm1. At the beginning of December 2006 the brightness dropped by 2\\fm5 over a period of two weeks to a value expected from the extrapolation of the early decay. The spectra also changed rapidly \\citep{nebular_phase} and reached a typical nebular phase at the end of December 2006 \\citep{munari_d}. The formation of very hot dust, $T_{\\rm d} \\approx 1\\,410~K$, began before December 12, 2006 ($\\approx 250^{\\rm d}$) \\citep{ir_dec,ir_late}. In May 2007 ($\\approx 396^{\\rm d}$), \\citet{ir_may07} reported continuing dust emission but at a lower temperature, $T_{\\rm d} < 520\\,$K.  \\citet{x-ray} also detected the source in X-rays on October 14, 2006 ($\\approx 191^{\\rm d}$), and reported that the spectrum is harder than expected from supersoft X-ray binaries. On May 5, 2007 ($\\approx 394^{\\rm d}$), \\citet{XMM} reported that the XMM-Newton spectra were poorly fitted with a single component absorbed model. They fitted a two-temperature thermal plasma model with a low and a high temperature of 0.2~keV and 2.3~keV, respectively.  Here we report the results of a spectral monitoring throughout the main brightening phase and further on until the target reached its final decline. ",
        "conclusions": "The derived ratio of $t_2/t_3$ and the decline rate of the first 60\\,\\,days and 4\\fm0 in brightness fit perfectly the normal decline described in \\citet{kato_06,kato_07}.  In addition, this ratio also resembles the prediction from a linear fit through the novae with $t_2 < 20\\fd0$ in Tab. 5.2 of \\citet{warner}. The correlation of the nova speed with the expansion velocities \\citep{mclaughlin,warner} also resembles that of a normal system. Thus, it seems reasonable to use the samples in \\citet{kato_06,kato_07} to estimate the white dwarf mass $M_{\\rm WD} = 1.2\\pm0.1 M_\\odot$ with $t_3$. Using the curves of \\citet{kato_1} leads to an even higher estimated mass. Although the samples mentioned above have a very large scatter of individual objects, they also give estimates for such mass for the accretor. On the other hand, even if the early decline looks very much like that of a normal object, one  has to be dubious whether one is allowed to use such models for a nova with such a massive secondary maximum. Thus, the results here have to be estimates. For the proposed orbital period of 0\\fd2070 \\citep{goranskij_1} and the mass of the accreting white dwarf of $M_{\\rm WD} \\approx 1.2$~M$_\\odot$ the secondary has to be a late type $M_{\\rm SE}\\le 0.75~$M$_\\odot$ K to M star.  We used \\citet{Cox} to derive a temperature $T$ from stellar colors and the evolution of the luminosity $L$ from the lightcurve and the bolometric corrections. During the bump the temperature increased slightly from 4\\,250 to 4\\,900~K and decreased again at maximum to 4\\,250~K (see ($V-I_{\\rm C}$) in Fig. \\ref{lightcurve}). The temperature derived from the early spectra \\citep{early_spec} was used as the temperature at maximum. Thereafter the evolution of the radius of the ``quasi'' photosphere derived by $R \\propto T^{-2}\\,L^{0.5}$ was obtained. The shrinking of the photosphere stopped at day 60 and stayed constant at 15~\\% of the radius of maximum until day 170. In case of a homologous expansion this is only possible with a significant increase of density. The fast ejecta reached the photosphere around day 170. With a speed of 2\\,200\\,km\\,s$^{-1}$ -- \\ derived from our spectra and consistent with the value measured in November 2006 by \\citet{munari_b} -- and the estimated radius above, this  bump material  was accelerated about 45 to 60 days after maximum, assuming a negligibly small radius of the accelerating region. Subtracting the extrapolated power law decline and assuming that the bolometric correction derived by the color is constant at $\\approx$ 0\\fm12 to 0\\fm2, leads to an estimate of the radiative energy in the event. About 40\\% of the total radiated energy of the object during the whole outburst is contained in the bump. The high velocity of the fast ejecta during this epoch also significantly increased the mechanical energy. Hence we presume that the bump event contained a significant fraction of the total outburst energy of the system.  The onset of massive dust formation, marked by the rapid decline, the fast rise of the color (see Fig. \\ref{lightcurve}), and the IR excess \\citep{ir_late} around day 250 occurred most likely in the inner region only. Thus, the spectral components of the outer optically-thin regions continued with a smooth decline, while the fast component was obscured within a few days (see Fig. \\ref{h-alpha2}). The forbidden [OI] lines, which showed only the narrow component throughout the whole event, were generated in the outer region only. Although the obscured inner region, even seen later in the IR excess \\citep{ir_may07}, cannot directly heat the outer shell, an effective energy transport changed the excitation state quickly -- [OI] $\\rightarrow$ [OII] $\\rightarrow$ [OIII] within 2 weeks.  The distance of $D \\approx 7.5^{+3.0}_{-2.5}$~kpc places V2362~Cyg at a galactocentric radius of about 11~kpc and a distance to the plane of about 310~pc. As the direction marks the maximum of the Galactic warp above the plane \\citep{warp}, the real distance to the disk is about twice this value. This coincides with the spiral arm +II beyond the Perseus arm \\citep{perseus}.  Further investigations of the ejecta of this unusual object and collection of all the data of various observers from the early decline until now is required to build the basis for a more sophisticated theoretical modeling. We would like to encourage modeling groups and would like to enliven a scientific discussion about the nature of the second outburst of V2362~Cyg and its older twin V1493~Aql."
    },
    "0801/0801.1860_arXiv.txt": {
        "abstract": "Observations with the {\\it Spitzer Space Telescope} have recently revealed a significant population of high-redshift ($z\\sim2$) dust-obscured galaxies (DOGs) with large (rest-frame) mid-infrared to ultraviolet luminosity ratios. Due to their optical faintness, these galaxies have been previously missed in traditional optical studies of the distant universe.  We present a simple method for selecting this high-redshift population based solely on the ratio of the observed mid-infrared 24$\\mu$m to optical $R$-band flux density. We apply this method to observations of the $\\approx 8.6~{\\rm deg}^2 $ Bo\\\"otes Field of the NOAO Deep Wide-Field Survey, and uncover $\\approx$2,600 DOG candidates (i.e., a surface density of 0.089~arcmin$^{-2}$) with 24$\\mu$m flux densities $F_{\\rm 24\\mu m} \\ge 0.3$mJy and $(R-[24])\\ge 14$ (i.e., $F_\\nu({\\rm 24\\mu m})/F_\\nu(R) \\simgt 1000$). These galaxies have no counterparts in the local universe. They become a larger fraction of the population at fainter 24$\\mu$m flux densities, increasing from 7$\\pm$0.6\\% of sources at $F_{24\\rm \\mu m}\\ge 1$~mJy to $\\approx 13\\pm1$\\% of the population at $\\approx$ 0.3~mJy.  These galaxies exhibit evidence of both star-formation and AGN activity, with the brighter 24$\\mu$m sources being more AGN-dominated. Their mid-infrared spectral energy distributions range from power-laws (likely AGN-dominated at mid-IR wavelengths) to systems showing a ``bump'', the latter likely resulting from the redshifted 1.6$\\mu$m peak characteristic of most stellar populations. Using primarily the W. M. Keck Observatory and {\\it Spitzer}, we have obtained spectroscopic redshifts for \\nz\\ objects within this sample, and find a broad redshift distribution which can be modeled as a Gaussian centered at $\\bar{z}\\approx1.99\\pm0.05$ and $\\sigma(z)\\approx0.45\\pm0.05$. The space density of this population is $\\Sigma_{\\rm DOG} (F_{\\rm 24\\mu m}\\ge 0.3~{\\rm mJy}) = (2.82\\pm 0.05)\\times 10^{-5} h_{70}^3~ {\\rm Mpc^{-3}}$, similar to that of bright sub-millimeter-selected or UV-selected galaxies at comparable redshifts. These redshifts also imply very large luminosities, with a sample median $\\nu L_\\nu(8\\mu{\\rm m})\\approx 4\\times 10^{11}\\Lsun$, implying ${\\rm 8\\mu m - 1mm}$ luminosities of $L_{\\rm IR}\\gtsim 10^{12-14}\\Lsun$ for the population. The infrared luminosity density contributed by this relatively rare DOG population is ${\\rm log}(L_{\\rm IR})\\approx 8.23^{+0.18}_{-0.30}$. This is $\\approx60^{+40}_{-15}$\\% of that contributed by $z\\sim2$ ultraluminous infrared galaxies (ULIRGs, with $L_{\\rm IR}>10^{12}\\Lsun$), and suggests that our simple selection criterion effectively identifies a significant fraction of $z\\sim 2$ ULIRGs. This IRLD is also $\\approx26\\pm14$\\% of the total contributed by all $z\\sim2$ galaxies, and comparable to that contributed by the luminous UV-bright star-forming galaxy populations at $z\\approx 2$. We suggest that these DOGs are the progenitors of luminous ($\\sim 4L^*$) present-day galaxies and are undergoing an extremely luminous, short-lived phase of both bulge and black hole growth. They may represent a brief evolutionary phase between sub-millimeter-selected galaxies and less obscured quasars or galaxies. ",
        "introduction": "Infrared-luminous sources (with $L_{\\rm IR}>10^{11}{\\rm L_\\odot}$) dominate the bright end of the bolometric luminosity function of galaxies in the local universe \\citep[]{soi1986}.  Their role in the story of galaxy evolution is not yet understood, but it is likely that their prodigious luminosities are the product of an extremely active phase during which these systems form stars and/or grow their central black holes at a rapid rate \\citep[]{san1988}.  Did all large galaxies exhibit such a phase in their early formative stages, or is this active phase characteristic of only some of the most massive systems?  To address these questions, we need to find and study the primary high-redshift galaxy populations contributing to the mid-IR and far-IR counts and understand their relation to the local galaxy population.  We know a great deal about the local population of IR-luminous galaxies, due in large part to the pioneering observations made with the {\\it Infrared Astronomical Satellite} \\citep[{\\it IRAS}; e.g.,][]{soi1987a,san1988,san1996}. Their mid-IR number counts show strong evolution \\citep[e.g.,][]{elb1999,chary2001,pap2004,chary2004,marl2004}. Although rare locally, infrared luminous galaxies become an increasing fraction of the galaxy population at high redshift, dominating the IR energy density at redshifts $z\\sim1$ \\citep[e.g.,][]{fra2001,lag2004,gru2005,lef2005,per2005}.  Far less is known about their higher redshift counterparts despite several ground-breaking studies.  Prior to the launch of the {\\it Spitzer Space Telescope}, our knowledge of high-$z$ IR-luminous galaxies was derived from studies of a handful of the most extreme {\\it IRAS} sources, constraints on the redshift evolution out to $z\\sim 1$ \\citep[largely due to {\\it ISO}; e.g.,][]{fra2001,chary2001,lag2003}, small samples of galaxies discovered at sub-millimeter wavelengths \\citep[e.g.,][]{sma2002,bla2002,cha2004b,cha2005}, and several galaxies discovered as a result of their extreme optical-to-near-infrared colors \\citep[the dusty `Extremely Red Objects', or EROs, with $R-K>6$; e.g.,][]{hu1994,dey1999,mcc2004a}.  The launch of the {\\it Spitzer Space Telescope} has revolutionized studies of the dusty galaxy population. It has provided the first mid-IR spectroscopy of dust emission in high-redshift star-forming galaxies and AGN, and revealed populations of galaxies that are extremely faint at (observed) visible wavelengths \\citep[e.g.,][Higdon et al. and Desai et al., both in preparation]{hou2005,yan2005,yan2007}. Many of these galaxies lie at redshifts $z\\simgt 1$, and their mid-IR spectra reveal either strong silicate absorption features, suggestive of an AGN, or emission from polycyclic aromatic hydrocarbons (PAHs) which generally arise in photodissociation regions associated with star-forming molecular clouds. A significant fraction of these active high-redshift galaxies are under-luminous at rest-frame UV wavelengths, suggesting that the prodigious star-formation and / or AGN accretion luminosity is hidden by dust that reprocesses and radiates the bulk of the energy at far-IR wavelengths \\citep[e.g.,][]{san1996}.  In this paper, we present a simple and economical method for selecting these high-redshift dusty populations, based solely on the ratio of the observed 24$\\mu$m to optical flux ratio. This selection, applied to a wide area survey, results in a large, significant population of extremely luminous high-redshift galaxies for which there are no local counterparts. These are systems with unusually red spectral energy distributions, which we interpret as being the result of extinction by large columns of dust. These galaxies appear to be largely absent from (rest-frame) UV-selected high-redshift galaxy samples. We have measured a large number of spectroscopic redshifts for this population, and confirm that they lie at redshifts $z\\sim2$. Given their large mid-infrared luminosities and implied far-infrared luminosities, these dust obscured galaxies may be undergoing an intense burst of star formation, accretion activity (onto a nuclear black hole) or plausibly both. We speculate that this population of objects, which may be selected simply on the basis of optical to mid-infrared flux ratio, are good candidates for systems in the throes of formation of their stellar spheroids and nuclear black holes. Despite their rarity, this population contributes a significant fraction of the infrared luminosity density at $z\\sim 2$.  This paper is structured as follows. In \\S2, we briefly describe the observations from which the sample is drawn and discuss the selection criteria we use to isolate the candidate high-redshift dust-obscured galaxy (hereinafter DOG) population. In \\S3, we present the main results: the optical-to-mid-infrared color distribution of the DOGs, their spectral energy distributions, redshift distribution and rest-frame UV and IR luminosities. In \\S4, we discuss the nature of the DOG population, derive their space density and their contribution to the $z\\approx 2$ infrared luminosity density, and discuss their possible relationship to other $z\\approx 2$ galaxy populations. We summarize our findings in \\S5.  Throughout this paper we adopt a cosmology with ${\\rm H_0 = 70~km\\ s^{-1}\\ Mpc^{-1}}$, $\\Omega_m = 0.3$, $\\Omega_\\lambda = 0.7$. Magnitudes quoted in this paper are in Vega units, except where explicitly defined otherwise (e.g., in the discussion of rest-frame 2200\\AA\\ luminosities). ",
        "conclusions": "\\subsection{The Nature of the Extreme Red Population}  What are the extreme red objects that become an increasing fraction of the population at fainter 24$\\mu$m flux densities? The spectroscopically measured redshifts for this population suggest prodigious luminosites ($10^{12-14}\\Lsun$) and redder colors than any local ULIRGs.   We therefore suggest that the rise in this red fraction is predominantly an evolutionary effect due to the appearance of a more enshrouded population of luminous objects at higher redshift. The space density of this population is low; if the comoving space density remains constant with time, they may evolve into the fairly luminous local ellipticals. We speculate that we have uncovered a population of young galaxies undergoing very rapid bulge formation and / or AGN accretion, that evolve into fairly luminous galaxies ($\\sim 4L^*$) in the present-day universe.  The optical-mid-IR SEDs of this population show a range of color. However, by selection, these objects are extreme in their optical-to-mid-IR properties and, as demonstrated by Figure~\\ref{color}b, have colors redder than any local dusty galaxy or ULIRG. Since we have only limited spectral information, the dust reddening cannot be determined unambiguously. However, the existing mid-infrared spectroscopy shows that these galaxies have spectra that rise into the mid-infrared, suggesting emission from warm dust, and exhibit either deep silicate absorption \\citep[e.g.,][]{hou2005,wee2006} or PAH emission \\citep[][; Desai et al. 2008, in prep]{yan2007}. Near-infrared spectroscopy by our group for a handful of galaxies reveals, where measurable, large Balmer H$\\alpha$/H$\\beta$ decrements in both the narrow and broad components of the Balmer lines \\citep[]{bra2007}. This suggests that the galaxies contain significant quantities of dust distributed on scales of, at the very least, the extended (i.e., kiloparsec scale) narrow-line emitting region. Since these galaxies lie at high redshift and are therefore very luminous, they must be undergoing significant AGN accretion and / or star-formation activity. This would suggest that the large rest-frame near-infrared to ultraviolet luminosity ratios result from heavy extinction of the most luminous emitting regions.  It is difficult to estimate the extinction in these galaxies from the limited photometric information, but we can derive some general constraints by considering the extinction necessary to create such red SEDs given a young stellar population. Figure~\\ref{AV} shows the $V$-band extinction (in magnitudes) necessary to create an $(R-[24])=14$ color for template galaxies with ages less than 1~Gyr. This is based on the assumption of an extinction law derived using the functional form of \\citet[]{cal2000} for $\\lambda \\le 2.2\\mu$m and values of the extinction as tabulated in Table~4 of \\citet[]{dra2003} at longer wavelengths. Although the assumption of a monolithic dust screen is unrealistic, it does provide a lower limit on the extinction required in these galaxies. In addition, the fact that the entire galaxy is so underluminous  at rest-frame UV wavelengths suggests that either the stellar population in these systems is old (i.e., $\\simgt 1$~Gyr) and intrinsically faint at UV wavelengths, or that the dust is distributed on large enough scales to shroud the stellar component as well. If these galaxies are young and dominated by a star-forming phase, then they are enveloped on kiloparsec scales by prodigious amounts of dust.  Given the redshifts ($z\\sim2$) and spectral energy distributions of these galaxies, it is very unlikely that they are dominated by old stellar populations, and we therefore hypothesize that these objects are dust-obscured galaxies (DOGs), where most of the luminous regions of the galaxy are shrouded in dust. Such a scenario would suggest an overall SED for these galaxies containing three (perhaps related) components: a reddened stellar component peaking at a rest-frame wavelength of 1.6$\\mu$m, cold dust enshrouding the star-forming regions reradiating the absorbed UV emission at long wavelengths (corresponding to, say, a rest-frame wavelength of 60$\\mu$m), and warm dust enshrouding the AGN dominating the emission in the rest-frame mid-infrared (i.e., at rest-frame wavelengths of 5-20$\\mu$m).  If the luminosity from the sources selected by $(R-[24])\\ge14$ is largely the result of star formation activity, then the implied star-formation rates are very large \\citep[$\\approx 1700 \\Msun$/yr for an $L_{\\rm IR}=10^{13}\\Lsun$ object;][]{ken1998b}. Such a violent starburst is unlikely to be confined to a small region, and the (observed) very large infrared-to-ultraviolet luminosity ratio must result from dust distributed on very large (kiloparsec) scales.  In addition, star-formation rates this large cannot be sustained for long periods, suggesting that these galaxies are viewed during a short-lived phase during which the bulk of the stars is produced.  If, instead, the sources are dominated by AGN, then the dust could be mostly confined to the nuclear regions of the galaxy. However, based on the measurement of large Balmer decrements in both broad- and narrow-line components in a few galaxies, \\citet[]{bra2007} argue that the dust distribution must be patchy and over large volumes. Even if the galaxy luminosity is dominated by AGN activity, this phase of evolution must be fairly short lived. The bolometric luminosity of an AGN may be written (in terms of its Eddington Luminosity) as $L_{\\rm bol}^{AGN}\\equiv \\epsilon L_{\\rm Edd}\\approx 0.33\\times 10^{13} (\\epsilon/0.1)(M_{\\rm BH}/10^9\\Msun) \\Lsun$, or in terms of its accretion rate as $L_{\\rm bol}^{AGN}=\\eta \\Mdot c^2 \\approx 0.15\\times 10^{13}(\\eta/0.1)(\\Mdot/1\\Msun yr^{-1})\\Lsun$. Large AGN luminosities therefore generally imply high accretion rates onto massive black holes. The present data on the DOGs do not allow any useful constraints on either the radiative efficiency, the accretion efficiency, or the black hole mass, but their luminosities of $10^{13}\\Lsun$ suggest that they must be radiating at close to the Eddington luminosity even if they harbor a very massive black hole. Even at high accretion efficiencies (i.e., $\\eta = 0.1$), the timescales to build up a $10^9\\Msun$ black hole would be short ($\\approx$0.15~Gyr). A flux density limited sample of galaxies (such as ours) will preferentially identify the galaxies in their most luminous phase, suggesting that we may be witnessing the last throes of the AGN formation in these obscured galaxies.  \\begin{figure}[t] \\epsscale{1.0} \\plotone{f13.eps} \\caption{The minimum extinction required to redden the spectral energy distributions of young stellar populations and AGN to produce an optical-to-mid-infrared color of $(R-[24])=14$. The template SEDs are models from the \\citet[]{bc2003} population synthesis library with ages of 5~Myr (solid line), 25~Myr (dotted line), 100~Myr (short-dashed line) and 900~Myr (dot-dashed line), and the median QSO template (long-dashed red line) from \\citet[]{elv1994}. The reddening law used is a combination of that from \\citet[]{cal2000} and longer wavelength estimates from \\citet[]{dra2003}, and assumes the unrealistic case of a dust screen in front of the emitting source in order to derive a firm lower limit on $A_V$. \\label{AV}} \\end{figure}  What is the relevance of this rather extreme population to the overall galaxy population at $z\\approx 2$? In the following subsections, we estimate the space density of the DOGs and their contribution to the infrared luminosity density in an attempt to address this question. In contrast to previous work, we do not discriminate between AGN and star-forming galaxies within our sample, but estimate their joint contribution to the space densities and contribution to the infrared luminosity density.  \\subsection{Space Density of DOGs}   Based on the observed redshift distribution, we now estimate the space density of this population. We assume that the comoving volume sampled by the $(R-[24])\\ge 14$ and $F_{\\rm 24\\mu m}\\ge 0.3$ selection criteria is simply the fraction of the volume between $0.5\\le z \\le 3.5$ weighted by the redshift distribution. In other words, the effective co-moving volume sampled is given by: \\begin{equation} V_c^{\\rm eff}({\\rm DOGs}) = \\Delta\\Omega \\int{g(z) dV_c} \\end{equation} where $\\Delta\\Omega$ is the solid angle covered by the survey, \\begin{equation} g(z) = {{1}\\over{\\sigma_z\\sqrt{2\\pi}}} {\\rm exp}\\left[-{{(z-\\bar{z})}\\over{2\\sigma_z^2}}\\right] \\end{equation} (with $\\bar{z}=1.98$ and $\\sigma_z=0.53$) is the (normalized) Gaussian fit to the redshift distribution, and $dV_c$ is the comoving volume element per unit solid angle per unit redshift interval \\citep[e.g.,][]{pee1993,hog1999}. For the Bo\\\"otes field, the volume sampled by the redshift distribution is roughly a third of the comoving volume between $0.5\\le z \\le 3.5$. The resulting average space density for this population in the redshift range $0.5\\le z \\le 3.5$ is \\begin{equation} \\Sigma_{\\rm DOG} (F_{\\rm 24\\mu m}\\ge 0.3~{\\rm mJy}) = (2.82\\pm 0.05)\\times 10^{-5} h_{70}^3~ {\\rm Mpc^{-3}}. \\end{equation}   This space density of DOGs is comparable to that of highly UV-luminous $z\\sim 2-3$ star-forming galaxies selected using the Lyman-break or near-infrared selection techniques \\citep[e.g.,][and references therein]{red2007,vand2006}. For example, the dust-obscured galaxy population has a space density comparable to that of $M_{\\rm AB}({\\rm 1700\\AA})-5{\\rm log}h_{0.7}\\approx -22.2$ UV-bright galaxies at $z\\sim 1.9-3.4$, based on the luminosity functions derived by \\citet[]{red2007}. Although the space densities are comparable, it is important to note that each DOG contributes significantly more IR (and bolometric) luminosity per galaxy to the universe at $z\\approx 2$. The population of ``massive galaxies'' selected as ``Distant Red Galaxies'' (DRGs) at near-infrared wavelengths has a space density of $(2.2\\pm0.6)\\times 10^{-4} h_{70}^3~{\\rm galaxies~ Mpc^{-3}}$ \\citep[]{vand2006}, roughly ten times more numerous than DOGs.  As mentioned above, the DRG and DOG populations also overlap in their rest-frame UV luminosity distributions, and it is therefore possible that a fraction of the star-forming DRGs are selected by our $(R-[24])\\ge 14$ criterion.  A recent measurement of the luminosity function by \\citet[]{wak2006} finds that luminous (i.e., $>4L^*$) red galaxies in the local Universe have a space density of $\\approx 2.5\\times 10^{-5} h_{70}^3 {\\rm Mpc^{-3}}$, very comparable to that of the $z\\approx 2$ DOGs. It is tempting to suggest that DOGs may therefore represent a common phase of evolution for most massive galaxies. Indeed, if this phase is short lived (i.e., if the duty cycle of the enshrouded, luminous phase when galaxies may be selected as DOGs is short), then the DOGs may be the progenitors of a large fraction of present-day $>L^*$ galaxies.  More detailed comparisons will require better determinations of the redshift distribution as a function of apparent magnitude and 24$\\mu$m flux density.  \\subsection{Infrared Luminosity Density of DOGs}  In order to estimate the infrared luminosity density contributed by the DOGs, we begin by assuming, as before, that all the DOGs with no measured spectroscopic redshifts have a redshift probability distrbution described by the Gaussian approximation to the observed redshift distribution. For the DOGs with spectroscopic redshifts, we computed their $L_8$ values and converted them to $L_{\\rm IR}$ estimates based on the procedure described above. For the DOGs without measured redshifts, we assigned random redshifts (drawn from the probability distribution) to each object and computed their luminosities. As in section \\S3.3.2, we computed luminosities using the conversion of \\citet[]{cap2007}, and derive a range based on the two extremes $L_{\\rm IR}=[5-15]\\times L_8$.  Since DOGs with fainter $F_{\\rm 24\\mu m}$ are also typically fainter in the IRAC bands, we adopt the mean $\\bar{\\alpha}=2.296$ as the spectral index for all sources.  The infrared luminosity density (IRLD) of the entire sample was estimated using 1000 Monte Carlo simulations of the redshift distribution of the galaxies. We find the total infrared luminosity density contributed by the DOG population to be \\begin{equation} {\\rm log[IRLD_{\\rm DOG}(F_{\\rm 24\\mu m}\\ge 0.3mJy)]} = 8.228^{+0.176}_{-0.302}, \\end{equation} where IRLD is in units of $\\Lsun~{\\rm Mpc^{-3}}$ and the limits are not 1$\\sigma$ uncertainties, but the likely range of the IRLD based on the uncertainties in the $L_8$ to $L_{\\rm IR}$ conversion (i.e., the lower and upper limits corresponding to $L_{\\rm IR}/L_8= 5$ and 15 respectively).  DOGs thus provide a significant contribution to the IRLD at $z\\approx 2$. According to \\citet[]{cap2007}, the total IRLD from all 24$\\mu$m-selected $z\\sim2$ galaxies is ${\\rm log(IRLD_{\\rm Total})}\\approx 8.819^{+0.073}_{-0.079}$, with the contribution from ULIRGs alone being $42^{+15}_{-22}\\%$ of the total.  Although DOGs presumably represent only a fraction of the $z\\sim 2$ population (i.e., the reddest subset), they appear to contribute $\\approx26\\pm 14\\%$ of the total IRLD from the 24$\\mu$m-selected $z\\sim 2$ galaxy population, and $\\approx 60^{+40}_{-15}\\%$ of the ${\\rm IRLD_{\\rm ULIRG}}$.  The \\citet[]{cap2007} $z\\sim2$ sample is identified (in the GOODS-N and -S fields) using photometric redshift estimates and should, in principle, include DOGs. What is noteworthy, however, is that the DOGs (selected solely on the basis of extreme $R-[24]$ color) comprise such a significant component of the IRLD$_{\\rm ULIRG}$, which suggests that the DOG selection criterion effectively selects out the bulk of ULIRGs at $z\\sim 2$.  Not all $z\\sim2$ galaxies may be selected as 24$\\mu$m sources. Indeed, one of the most successful methods of selecting high-redshift galaxy populations is based on the UV emission of young, star-forming galaxies \\citep[e.g.,][and references therein]{red2007}. Recently, \\citet[]{red2007} have estimated the contribution of UV-selected galaxies to the IRLD at $z\\sim2$. The IRLD contributed by DOGs is 158\\%, 60\\%, 35\\% and 56\\% of the IRLD contributed by the $L_{\\rm IR}=10^{9-10}$, $10^{10-11}$, $10^{11-12}$, and $>10^{12}\\Lsun$ UV-bright populations respectively.  Do the bright and faint DOGs contribute comparably? To investigate this, we divided the DOG samples at $F_{\\rm 24\\mu m}=0.6$~mJy, resulting in 2,149 galaxies fainter than the cut with a median flux density of 0.37~mJy, and 452 galaxies brighter than the cut with a median flux density of 0.85mJy. The redshift distributions on either side of this cut are slightly different ($\\bar{z}\\approx1.9$/2.1 for the fainter/brighter galaxies with spectroscopically measured redshifts), but there are only 22 redshifts below the cut, so the distribution is not well constrained. If we assume that the distributions are the same at all flux densities, we find the IRLDs for the bright and faint subsamples are ${\\rm log[IRLD_{\\rm DOG}(F_{\\rm 24\\mu m}\\ge 0.6)]=7.78\\pm0.02}$ and  ${\\rm log[IRLD_{\\rm DOG}(0.3\\le F_{\\rm 24\\mu m}< 0.6)]=7.91\\pm0.01}$ respectively. The faint and bright sources thus contribute comparable amounts to the IRLD, with the increasing number making up for the lower luminosity of the fainter source population.  In actuality, there is a large range in optical-to-infrared color, and we have chosen the $(R-[24])\\ge 14$ criterion simply to isolate the higher redshift galaxies. Since the redshift distribution is broad, selecting objects at a bluer threshold would still result in high-redshift population, but with a larger low-redshift contamination. Moreover, we recall that the number density of DOGs rises toward fainter 24$\\mu$m flux densities, suggesting that $F_{\\rm 24}<0.3$mJy sources may contribute significantly to the IRLD as well. In addition, we have not corrected our IRLD estimates for the sample incompleteness at the faint 24$\\mu$m limit of our data. For all these reasons, the estimates of space density and luminosity density reported here for the DOGs are likely to be lower limits, corresponding to the contribution of only the most extreme (in $F_{\\rm 24\\mu m}$ and $(R-[24])$ color) subset of the population. It is therefore remarkable that even this relatively rare population can contribute such a significant fraction of the overall infrared luminosity density.   \\subsubsection{Comparison to Other $z\\approx 2$ Galaxy Populations}  The galaxies selected by virtue of their extreme mid-infrared to optical flux density ratios seem to be under-represented in most optically selected samples of high-redshift galaxies. Indeed, our 24$\\mu$m flux density criterion for DOGs (i.e., $F_{\\rm 24\\mu m}\\ge 0.3$~mJy, imposed by the existing MIPS observations of the Bo\\\"otes Field) selects a relatively rare population that has only a handful of members in most existing small-field, deep samples. In this subsection, we attempt to compare our sample with other known $z\\approx 2$ galaxy populations, most of which have been selected using observations at optical or near-infrared (i.e., rest-frame UV or optical) wavelengths.  As we mentioned in \\S3.3.3 (and shown in Figures~\\ref{RbandHist} and ~\\ref{2200Hist}), the DOG population is systematically fainter than the UV-bright populations selected at $z\\approx 2$ by, say, the ``BX/BM'' selection criteria of \\citet[]{red2007}. This is not surprising, since we are selecting galaxies to have a very large mid-infrared to UV flux ratio. What is interesting is that our selection results in a sample that lies at high redshift, with no bright, low-redshift members. Moreover, it is important to note that the DOG population, while being faint at optical wavelengths, is actually bolometrically very luminous, and thereby contributes significantly to (and perhaps {\\it dominates}) the top end of the bolometric luminosity function (i.e., at  $L_{\\rm IR}\\simgt 10^{12}\\Lsun$).   A comparison with the $z\\approx 2-3$ ``Distant Red Galaxy\" (DRG) samples selected using near-infrared selection techniques \\citep[typically, $J-K>2.3$; e.g.,][]{franx2003,vand2004} is more difficult, since the Bo\\\"otes NDWFS field lacks very deep near-infrared imaging data. However, the DRG samples appear to have median $R_{\\rm AB}\\approx 25.9$ with 25(75)\\%-ile values of the distribution at $R_{\\rm AB}\\approx 25.1(26.7)$ \\citep[]{vand2006}. The median rest-frame luminosity of DRGs is $M_{\\rm 2200\\AA} \\approx -18.3$~AB mag for $K<21$ (Vega) samples selected from the Hubble Deep Field South and MUSYC Deep fields \\citep[G. Rudnick, personal communication;][]{gaw2006}. These magnitudes, magnitude ranges and rest-frame luminosities are comparable to those measured for our $(R-[24])\\ge14$ population, suggesting that these populations may overlap. It is now understood that the DRG selection results in a mixed sample of passively evolving systems, star-forming galaxies and AGN \\citep[e.g.,][]{kriek2006,kriek2007}. Since our selection criteria require galaxies to be bright at 24$\\mu$m, it is likely that our selection does not include the passively evolving galaxies.   Recently, much work has been done using the $BzK$ selection technique pioneered by \\citet[]{dad2004} to select star-forming and passive galaxy populations at high redshift (i.e., $z\\approx 2$). We examined the $(R-24)$ colors of the $BzK$ selected star-forming galaxies in the GOODS-N field \\citep[]{dad2007a,dad2007b}, and find that only 4 of the 187 $BzK$ galaxies in GOODS-N satisfy our selection criteria. The vast majority of the $BzK$ galaxies are less extreme in their $(R-[24])$ colors, and are fainter in their 24$\\mu$m flux densities. This may be due to the lower AGN contribution (by selection) to the mid-infrared luminosities of the $BzK$ star-forming population.  Given the surface density of bright DOGs (i.e., with $F_{\\rm 24\\mu m}\\ge 0.3$mJy, we would have expected roughly 14 galaxies in the GOODS-N field. Indeed, there are 21 galaxies with these selection criteria in the GOODS-N field, 3 of which are brighter than $F_{\\rm 24\\mu m}=1.0$~mJy (E. Daddi and M. Dickinson, personal communication). % Fifteen of these 21 galaxies are also selected by the $BzK$ techniques, but all but 4 are rejected from the $BzK$ `star-forming' galaxy sample because they are hard X-ray sources (E. Daddi, personal communication).  The $(R-[24])\\ge 14$ and $BzK$ selections therefore appear to be comparable, with the star-forming systems at bright 24$\\mu$m flux densities effectively selected by both techniques. The simple $(R-[24])$ selection results in a slightly larger sample with perhaps a slightly broader redshift distribution. It is worth noting that the $BzK$ technique can be fairly expensive in telescope time, requiring deep imaging in three bands. The simple technique described here of using the $R$-[24] color selects similar star-forming galaxy and AGN populations more economically (90~sec with {\\it Spitzer}/MIPS and 6000~sec of $R$-band per field).  Recently, a similar population of dust-obscured galaxies was identified by \\citet[]{yan2007} using the following selection criteria: $F_{\\rm 24\\mu m}>0.9$~mJy; $(\\nu F_\\nu({\\rm 24\\mu m})/\\nu F_\\nu({\\rm 8\\mu m}) > 3.16; (\\nu F_\\nu({\\rm 24\\mu m})/\\nu F_\\nu({\\rm 0.7\\mu m}) > 10$. Our selection criterion for DOGs corresponds to \\\\ $(\\nu F_\\nu({\\rm 24\\mu m}) / \\nu F_\\nu({\\rm 0.7\\mu m}) > 28.6$, and thus selects the redder, and a higher redshift subset of the population. The \\citet[]{yan2007} galaxies with $(R-[24])\\ge 14$(15) all lie at $z > 1$(1.6). The $(\\nu F_\\nu({\\rm 24\\mu m})/\\nu F_\\nu({\\rm 8\\mu m})$ distribution for DOGs shows that the selection criteria used by \\citet[]{yan2007} would reject 36\\% of the DOGs.  Finally, although dust-obscured galaxies have been found in other studies using comparable selection criteria, we emphasize the simplicity and robustness of the current approach. In particular, the simple selection criteria will make it easier, in principle, to quantify the selection function once a large sample of spectroscopic redshifts become available. One significant difference in the way we select the sample of DOGs is that we do not discriminate between AGN-dominated and star-formation-dominated systems. The mid-infrared colors of many of the DOGs suggest that they may indeed be dominated by AGN at bright 24$\\mu$m flux densities (i.e., $F_{\\rm 24\\mu m}\\simgt 0.8$). Although the majority of the fainter population does show stellar ``bumps'', PAH emission features in their mid-infrared spectra, and other evidence of starlight in their spectral energy distributions, their mid-infrared,  far-infrared, or bolometric luminosity may yet contain a significant contribution from an AGN. The combination of AGN and star-formation signatures in the DOG population (and in some cases in the very same objects) suggests that the population may be undergoing both rapid black-hole growth and rapid star-formation.  \\subsection{The Possible Role of DOGs in Galaxy Evolution}  We have demonstrated that the simple selection criterion of large mid-infrared to optical flux density ratio efficiently isolates a significant population of $z\\approx 2$ galaxies. These systems are extremely luminous at rest-frame mid-infrared wavelengths, presumably due to warm dust heated by some combination of AGN and hot, young stars. What is the role of these galaxies in galaxy evolution?  In a seminal paper, \\citet[]{san1988} proposed that local ULIRGs represented a brief dust-enshrouded stage in the formation of quasars. Despite the extreme colors of the DOGs and the lack of local galaxies (even ULIRGs) with similar colors, there are some similarities between the two populations. They are both at the extreme luminous end of the galaxy population at their epochs, and they both appear to be heavily enshrouded systems exhibiting evidence for both intense star formation and AGN accretion activity.  We therefore speculate that the DOGs represent a fairly short-lived phase in the evolution of massive (i.e., $\\sim4L^*$) galaxies, in between the phases represented by the sub-millimeter galaxy population (SMGs) and the more ``passive'' galaxies (e.g., selected as passive $BzK$s or $DRG$s). We envision a scenario in which gas accumulates in deep potentials, triggering an early episode of star-formation which quickly results in the formation of large quantities of dust. This results in a system which is luminous at sub-mm wavelengths (as an SMG), with cold dust temperatures. The mid-infrared spectra of these systems would show PAH features, characteristic of star-forming systems, as has been observed \\citep[e.g.,][]{men2007,pop2008}. At some point, as the accretion and star-formation proceed, an AGN is triggered which heats the dust to warmer temperatures. It is at this point that the system would have a warmer characteristic dust temperature and be selected as a DOG. If the AGN is able to destroy or expel most of the dust (or the star-formation proceeds to the point where the dust and gas are sufficiently consumed), the DOG would evolve into an optically visible AGN. Whether or not the DOG is transformed into an optically-luminous AGN will depend on the relative timescales of dust destruction / consumption and AGN accretion. It has been suggested recently that AGN feedback and star-burst driven winds can provide two important mechanisms for terminating the star-formation in galaxies \\citep[e.g.,][and references therein]{hop2007b}. In such a scenario, the galaxy evolutionary phase represented by DOGs would follow the bulge growth phase (i.e., the SMGs?), but precede the phase when the galaxy may be visible as an optical quasar or, eventually, as a red, passively-evolving galaxy (i.e., the DRGs or passive $BzK$ galaxies?). If, instead, the AGN phase is more short lived than the UV-bright star-forming phase, then DOGs would evolve into a UV-bright star-forming population.  There is some circumstantial support for such a scenario from the measured space density and redshift distribution of the DOGs. The DOG surface density on the sky (of 0.089~arcmin$^{-2}$) is similar to that of the luminous sub-mm galaxy population with 850$\\mu$m flux densities $F_{\\rm 850\\mu m} > 6$mJy \\citep[]{cop2006}, and they have a comparable redshift distribution \\citep[]{cha2005}. The 850$\\mu$m sub-mm flux densities of the DOG population have not yet been measured, but the upcoming SCUBA-II Legacy survey should provide useful constraints. In the scenario described above, the DOGs are expected to have warmer dust temperatures than the SMGs, and it is indeed interesting that the detection rate of SMGs at MIPS wavelengths (24$\\mu$m and 70$\\mu$m) is low, with only $\\simlt 40$\\% of the SMGs in the GOODS-N field having 24$\\mu$m flux densities $F_{\\rm 24\\mu m}\\ge 0.3$mJy \\citep[]{pop2008}. The space density of the brightest DOGs (with $F_{\\rm 24\\mu m}\\ge 1$~mJy) is comparable to that of unobscured QSOs at comparable redshifts \\citep[e.g.,][]{bro2006}. If these brighter sources are powered primarily by AGN, they represent a significant component of the accretion history of supermassive black holes, missed by traditional QSO surveys.  Clustering analyses of the population may provide a clue to the possible evolutionary products of DOGs and the lifetime of the DOG phase relative to the SMG phase, and we will investigate this in a future paper."
    },
    "0801/0801.1051_arXiv.txt": {
        "abstract": "We present spectra of accretion discs around white-dwarfs calculated with an improved and updated version of the Shaviv \\& Wehrse (1991) model. The new version includes line opacities and convective energy transport and can be used to calculate spectra of hot discs in bright systems (nova--like variables or dwarf novae in outburst) as well as spectra of cold accretion discs in quiescent dwarf novae. ",
        "introduction": "\\label{sec:intro}   In weakly magnetized cataclysmic variable stars (CVs) matter lost by a Roche-lobe filling low-mass companion forms an accretion disc around the white-dwarf primary. In bright systems such as nova-like binaries and dwarf-novae in outburst the disc is the dominant source of luminosity but even in quiescence the disc emission might provide crucial information about the physics of accretion, in particular about the mechanisms transporting angular momentum and releasing gravitational energy.  Until now one could calculate only spectra for hot accretion discs \\citep[see][and references therein]{wh98}. In such an approach the disc is divided into rings whose vertical structure is calculated using the program TLUSDISK \\citet{hubeny90,hubeny91}. A different way of solving radiative transfer equations in accretion discs was presented by \\citet[][hereafter SW]{sw91}. Also, in their approach one could apply it to hot discs only. The preliminary attempt to apply the SW code to cold quiescent discs in \\citet{Idan1} was not conclusive. The main obstacle in solving the vertical structure of cold accretion discs is the presence of convection-dominated zones. \\citep[The description of the quiescent disc of SS Cyg by][neglects convection]{kromer}.  The absence of spectral models for cold quiescent discs is the reason why fitting models to observations can be a frustrating exercise\\citep[see][and references therein]{urbsion}. Using the models of hot stationary discs of \\citet{wh98} to describe the spectra of cold (and non-stationary) quiescent discs cannot be expected to give relevant results. In fact since the observational data were the UV spectra obtained by the IUE, it was not very surprising that the best fits were obtained with no disc emission at all: quiescent discs are not supposed to be UV emitters so that the only source of UV radiation in such systems should be the white dwarf.  The description of the quiescent state of the dwarf-nova cycle is the Achilles heel of the widely accepted disc instability model (DIM) according to dwarf-nova outbursts are due to a thermal-viscous instability \\citep[see][for a review]{l01}. For example while the DIM predicts the quiescent disc to be optically thick (especially in its outer regions), the observations of the quiescent dwarf-nova IP Peg seem to suggest the opposite \\citep{little01,ribeiro07}. Correct modeling of the emission from quiescent discs is therefore an obvious prerequisite to solving this thorny problem .  In general a full description of dwarf-nova outburst cycle should include the radiation transfer equations. In practice this much too costly and therefore unfeasible. However, instead of including one can combine the DIM with the solutions of the radiative transfer equations for the same disc parameters. The \\citet[][hereafter HMDLH]{hmdl98} version of the DIM in which the time-dependent radial evolution equations use as input a pre-calculated grid of hydrostatic vertical structures (a 1+1D scheme) is well suited to such enterprize. The vertical structures are calculated using the standard equations of stellar structure or the grey-atmosphere approximation in the optically thin case. The angular-momentum viscosity transport mechanism is described by the $\\alpha$-ansatz of \\citet{ss73}. For a given $\\alpha$, $M/R^3$ and effective temperature $\\Teff$ there exist a unique solution describing the disc vertical structure. Such solutions which are very handy for solving the time-dependent equations of the disc evolution cannot be used to produce emission spectra. These can be, however, calculated by using the same input parameters ($\\alpha$, $M/R^3$ and $\\Teff$) in a radiative-transfer code on the condition that the vertical solutions calculated by the two methods are the same. In this way one can reproduce spectra of the whole cycle of a dwarf-nova outburst.  The outline of the present article is as follows. In section \\ref{sec:model} we discuss the model used with special stress put on marking the differences with original SW code on which it has been based. In section \\ref{sec:results} we present solutions obtained for various types of physical set-ups. In subsection \\ref{subsec:sc} we compare our solutions for vertical structure with those obtained with HMDLH code. Then in \\ref{subsec:wh} we compare our disc spectra with the UV spectra obtained by \\citet{wh98}. In the following subsection {subsec:non-stat} we present spectra of cold non-equilibrium discs which in the DIM represent the quiescent state of the dwarf-nova outburst cycle. Section \\ref{sec:concl} contains the discussion, and conclusion. ",
        "conclusions": ""
    },
    "0801/0801.4674_arXiv.txt": {
        "abstract": "We provide measurements of the neutral hydrogen fraction $x_{\\rm HI}$ at $z\\sim 6$, by comparing semi-analytical models of the Ly$\\alpha$ forest with observations of high-$z$ quasars and Gamma Ray Bursts absorption spectra.\\\\ We analyze the transmitted flux in a sample of 17 QSOs spectra at $5.74\\leq z_{em}\\leq 6.42$ studying separately the narrow transmission windows (peaks) and the wide dark portions (gaps) in the observed absorption spectra. By comparing the statistics of these spectral features with our models, we conclude that $x_{\\rm{HI}}$ evolves smoothly from $10^{-4.4}$ at $z=5.3$ to $10^{-4.2}$ at $z=5.6$, with a robust upper limit $x_{\\rm{HI}} < 0.36$ at $z=6.3$. We show the results of the first-ever detected transverse proximity effect in the HI Ly$\\alpha$ forest, produced by the HII region of the faint quasar RD J1148+5253 at $z=5.70$ intervening along the LOS of SDSS J1148+5251 at $z=6.42$.\\\\ Moreover, we propose a novel method to study cosmic reionization using absorption line spectra of high-redshift GRBs afterglows. We show that the time evolution and the statistics of gaps in the observed spectra represent exquisite tools to discriminate among different reionization models.  By applying our methods to GRB~050904 detected at $z=6.29$, we show that the observation of this burst provides strong indications of a highly ionized intergalactic medium at $z\\sim 6$, with an estimated mean neutral hydrogen fraction $x_{\\rm HI}=6.4\\pm 0.3 \\times 10^{-5}$ along that line of sight.\\\\ \\\\ PACS 98.62.Ra - Intergalactic matter; quasar absorption systems; Lyman forest \\\\ PACS 98.54.Aj - Quasars\\\\ PACS 98.70.Rz - $\\gamma$-ray bursts\\\\ ",
        "introduction": "Although observations of cosmic epochs closer to the present have indisputably shown that the InterGalactic Medium (IGM) is in an ionized state, it is yet unclear when the phase transition from the neutral state to the ionized one started. Thus, the redshift of reionization, $z_{rei}$, is still very uncertain.  In the last few years, our knowledge of the reionization process has been enormously increased mainly owing to the observation of $z\\sim 6$ QSOs by the SDSS survey \\cite{fan06} and CMB data \\cite{page07}. Long gamma ray bursts (GRB) may constitute a complementary way to study the reionization process possibly probing even larger redshifts, the current recold holder being GRB~050904 at $z=6.3$ \\cite{tagliaf05} \\cite{haislip06}.  We provide measurements of the neutral hydrogen fraction $x_{\\rm HI}$ at epochs approaching reionization, by comparing semi-analytical models of the Ly$\\alpha$ forest with observations of the highest-$z$ QSOs and GRBs absorption spectra. ",
        "conclusions": "We measure the neutral hydrogen fraction $x_{\\rm HI}$ at epochs approaching the reionization, by comparing semi-analytical models of the Ly$\\alpha$ forest with observations of high-$z$ quasars and Gamma Ray Bursts absorption spectra. We consider an Early Reionization Model (ERM), characterized by a highly ionized Universe at $z\\sim 6$ and a Late Reionization Model (LRM) in which reionization occurs at $z\\sim 6$.\\\\ By comparing statistical analysis of the transmitted flux in a sample of 17 QSOs spectra at $5.74\\leq z_{em}\\leq 6.42$ with our models, we find that both ERM and LRM provide good fits to the observed LGW distribution, favoring a scenario in which $x_{\\rm HI}$ smoothly evolves from $10^{-4.4}$ at $z\\approx 5.3$ to $10^{-4.2}$ at $z\\approx 5.6$, with a robust upper limit $x_{\\rm{HI}} < 0.36$ at $z=6.3$. Discriminating among the two reionization scenarios would require a sample of QSO at even higher redshifts.\\\\ We show the results of the first-ever detected transverse proximity effect in the HI Ly$\\alpha$ forest, produced by the HII region of the faint quasar RD J1148+5253 at $z=5.70$ intervening along the LOS of SDSS J1148+5251 at $z=6.42$.\\\\ Moreover, we show that the time evolution of gaps in GRBs absorption spectra represent exquisite tools to discriminate among different reionization models.  By applying our methods to GRB~050904 detected at $z=6.29$, we show that the observation of this burst provides strong indications of a highly ionized intergalactic medium at $z\\sim 6$, with an estimated mean neutral hydrogen fraction $x_{\\rm HI}=6.4\\pm 0.3 \\times 10^{-5}$ along that line of sight."
    },
    "0801/0801.1926_arXiv.txt": {
        "abstract": "Mean-motion resonances (MMRs) are likely to play an important role both during and after the lifetime of a protostellar gas disk. We study the dynamical evolution and stability of planetary systems containing two giant planets on circular orbits near a 2:1 resonance and closer.  We find that by having the outer planet migrate inward, the two planets can capture into either the 2:1, 5:3, or 3:2 MMR. We use direct numerical integrations of $\\sim 1000$ systems in which the planets are initially locked into one of these resonances and allowed to evolve for up to $\\sim 10^7$~yr.  We find that the final eccentricity distribution in systems which ultimately become unstable gives a good fit to observed exoplanets.  Next, we integrate $\\sim 500$ two-planet systems in which the outer planet is driven to continuously migrate inward, resonantly capturing the inner; the systems are evolved until either instability sets in or the planets reach the star.  We find that although the 5:3 resonance rapidly becomes unstable under migration, the 2:1 and 3:2 are very stable.  Thus the lack of observed exoplanets in resonances closer than 2:1, if it continues to hold up, may be a primordial signature of the planet formation process. ",
        "introduction": "The discovery of extrasolar planets around sun-like stars \\citep{mayorqueloz95,marcybutler95,marcybutler98} has revealed that early large-scale migration likely plays an important role in the formation of planetary systems. For example, observations have revealed a nontrivial number (roughly 6\\% overall) of close orbiting gas giants, called `hot Jupiters,' planets that would have great difficulty forming in their present locations \\citep{bodenheimeretal00}. In fact, disk-planet interaction theory predicts disturbingly short migration time scales that seem to threaten the survival of everything ranging from planetary embryos to gas giants \\citep{ward97}. Additionally, there are currently at least eight planetary systems that have two planets in MMR \\citep{udryetal07}. These include systems such as GJ 876 \\citep{marcyetal01} and HD 82943 \\citep{leeetal06} with two planets in a 2:1 resonance, and HD12661 \\citep{fischeretal03}, which contains two planets in what may be a 6:1 resonance. The origins of such resonant systems are thought to be due to disk-planet interactions that induce angular momentum transfer and differential migration  \\citep{snellgroveetal01,leepeale02,papa03,kleyetal04}.  There are currently two competing theories to explain planet formation. The ``core accretion\" model starts with the sedimentation and collisional growth of dust grains and smaller planetesimals in the protoplanetary disk \\citep[see][and references therein]{lissauer93}. These rocky cores (protoplanets) continue to build up until their gravitational influence allows for the accretion of the surrounding gas, forming gas giants. The other theory invokes direct gravitational instabilities in the disk \\citep[e.g.,][]{cameron78,boss00}.  In these theories the dynamical relaxation time of a planetary system can, in principle, be longer than the time scale for planet formation. Therefore, long term stability is not guaranteed for a planetary system's initial configuration. Gravitational interactions between massive protoplanets can lead to orbit crossing and instability, resulting in massive planets being thrown closer to their parent stars \\citep{rasioford96,weidenschillingmarzari96}. In the case of one of the planets being ejected, the surviving planet is left with high eccentricity (some with $e > 0.90$). Additionally, migration occurs within the disk due to the exchange of angular momentum between the planet and the surrounding disk material through planet-disk interactions \\citep{goldtrem79,goldtrem80,linpapa79,linpapa93,papalin84,ward86}. In a laminar disk, there are two basic modes of migration: For an embedded body, imbalance between the torques from the inner and outer parts of the disk is thought to lead to orbital decay; this is commonly referred to as type I migration \\citep[e.g.,][]{ward97}.  If a body is massive enough, it locks itself to the disk by opening a deep annular gap, and is thus carried along as the disk accretes onto the star \\citep{linpapa86,ward97}; this is called type II migration.\\footnote{There is also a proposed type III migration, which involves disk material flowing through the co-orbital region of the planet, which will not be considered here \\citep[see][]{papaetal07}.} When orbital migration occurs for different planets at different rates, convergent migration and locking into mean-motion commensurabilities can occur.  For the two planet case, the regions of dynamical stability (i.e., the region where close encounters are impossible) can be determined analytically by studying how the value of the Jacobi integral relates to the topological (Hill) stability of the system \\citep{gladman93}. However, this criterion does not take into account the fact that planets may be in resonance. It has been shown that resonant systems have additional regions of stability \\citep{barnesgreen07}, which current analytic criteria predict to be unstable. In fact, \\cite{barnesgreen07} have shown that nearly all observed resonant systems lie in these extended regions.  Overall, planet-planet-disk interactions are likely key in determining the final configuration of a planetary system. To date, each type of dynamical process (planet-disk interactions, dynamical instabilities, and resonances) has usually been considered individually. The motivation of this study is to explore physical situations where all of these processes occur simultaneously. In particular, we consider gap-opening planets located close to particular resonances within a protoplanetary disk. This results in convergent migration and resonance capture, but possibly also eventual instability due to close encounters.  There are two objectives to this study. The first is to better quantify the regions of stability for particular resonances. More explicitly, we will consider the 3:2, 5:3, and 2:1 resonance for two gap-opening planets of varying mass. We place an upper bound on the possible masses for planets in these three resonances. This provides an interesting problem in dynamics by allowing us to more generally map the parameter space for these particular resonances in and out of the protoplanetary disk, compared to \\citet{barnesgreen07} who studied particular observed extrasolar systems. Additionally, in the case where there are instabilities, this study allows us to see how the final configurations (e.g., distribution of eccentricity) differ, if at all, from non-resonant planet scattering \\citep[see, e.g.,][]{fordetal01,chatterjeeetal07}.  Our second objective is to study the dynamics of gas giants in the region of the protoplanetary disk where we believe these planets form, between 5 and 10 AU. The three above resonances are typical resonances that planets can resonantly capture into, since they require little to no eccentricity to give a high probability of capture when planets migrate at rates consistent with Type II migration. Here we are assuming that planets are born on nearly circular orbits. We then study the evolution as the coupled system of two planets migrates with the disk. To date, the closest resonance to be observed---that is, the resonance with the smallest ratio of outer to inner period---is the 2:1. We aim to better quantify why we have not found any systems in the 3:2 or 5:3 resonances and measure how common we should expect each of these resonances to be in extrasolar planetary systems.  This paper is organized as follows. In \\S 2, we explain the numerical treatment used in both the resonance stability and disk simulations. In \\S 3, we give further details of the theory used in our models for studying the stability of the aforementioned resonances. We then present the results of those simulations. In \\S 4, we develop the theory used to model the planet-disk interactions in the protoplanetary disk as well as the results of those runs. We conclude with a summary and discussion in \\S 5. ",
        "conclusions": "In this paper, we have examined the stability of particular first and second-order resonances  with and without disk interactions. We have shown that the two-planet stability criterion (computed to third-order) is invalid for these resonances (see Figures \\ref{fig:21Stab}, \\ref{fig:32Stab}, and \\ref{fig:53Stab}). The 2:1 resonance has regions of stability that extend up to mass ranges comparable to brown dwarfs, while the 3:2 resonance has a region of stability with combined planet masses up to twice the above criterion. The 5:3 resonance has a reduced region of stability compared to analytic criteria.  Next, we have examined the results of instability. As shown in Figure \\ref{fig:ResCumHisto}, combining an equal number of unstable outcomes originating in each of the above three resonances (with appropriate relative fractions of ejections and collisions) yields an eccentricity distribution which fits, at least, as well as previous studies of dynamical instability in multi-planet systems (in particular \\citealt{chatterjeeetal07} for the three-planet case, and \\citealt{jurictremaine07} for larger ensembles).  This is an intriguing result which raises the possibility of a link between these MMRs and exoplanet eccentricities.  However, significant caution is warranted in interpreting it.  Our ``outcomes\" represent snapshots at the time of instability, with no attempt made to model the subsequent effect of the disk on the remaining single planet.  Also, the way in which we assemble resonances in \\S \\ref{sec:initial conditions} is rather artificial, at least in comparison to convergent migration of \\S \\ref{sec:ppdinteract}.  With the addition of migration in \\S \\ref{sec:ppdinteract}, plus the eccentricity damping expected from disk interactions, we find the 2:1 resonance to be the most stable, with 100\\% of the runs remaining stable even without eccentricity damping (see Table \\ref{tab:resoutcome}). Additionally, the 3:2 resonance remains completely stable with our adopted damping prescription. A second-order resonance, the 5:3, has a high capture efficiency but goes unstable at very low eccentricities due to close encounters.  This means that in order for planets to become locked into the 3:2 resonance, their primordial period separation must generally lie between 5:3 and 3:2.  In this study, we selected the initial semi-major axes of the planets to coincide with the location where it is believed planets form \\citep[see, e.g.,][]{kokuboida02,thommesduncanlevison03}, rather than allowing the planets to migrate from further out in the disk. Under this assumption that planets form between 5 and 10 AU, if planets were able to form at any location with equal probability, the probability that a planet forms between the 3:2 and 5:3 resonances of a planet located at 5 AU is  \\begin{equation} \\frac{\\pi\\left(\\left[(5/3)^{2/3} - (3/2)^{2/3} \\right]\\cdot 5\\right)^2}{\\pi\\cdot5^2} \\simeq 0.009. \\end{equation} That is, the probability of forming such a configuration is about 1\\%.  To date, we have not discovered any planets in a closer than 2:1 resonance.  The lack of 5:3 planets is readily accounted for by this study:  Since this resonance is unstable at low eccentricity, it is unstable to survive all the way to a mature planetary system.  The fact that 3:2 planets are {\\it also} not seen, notwithstanding the resonance's robustness, may well be telling us something fundamental about how planets form:  It may simply be that neighboring giant planets are seldom born with primordial period separations of {\\it less} than 5:3.  The planets in HD 12661 are believed to be in a 6:1 resonance, and there is much debate as to how they could have become locked in such a high-order resonance. This study suggests a mechanism. We have shown that an instability in one of the lower-order resonances, which needs little to no initial eccentricity for planets to become trapped in it, can result in the planets scattering apart, to then be brought together again and re-locked in a more distant, higher-order resonance (See Table \\ref{tab:rescap}). Further work is necessary to include the adjustment of the disk edge after the planets have been locked into a more distant MMR.  However, migration of the outward-scattered planet would likely be somewhat more gentle in a self-consistent disk---generally we would expect an annulus of gas to initially exist between the two planets---in which case resonant re-capture would actually tend to be {\\it more} likely than in our current simple implementation.  Aside from the physical findings, this study highlights an important point on orders of accuracy. Since the resonance stability tests relied on two planets of unequal mass, we chose to use up to the third order term in Equation \\ref{eqn:glad}, because this was the first term that was not symmetric in the masses (that is, a mapping $M_i \\leftrightarrow M_o$ would not be an identity mapping).  Furthermore, as shown by Table 1, additional terms can change the result up to twice  the first-order value. If we were to truncate Equation \\ref{eqn:glad} to just the first term, we would have concluded that the equation gives a better representation of the region of stability for the resonant case than it actually does.  This paper represents a first step in our study of resonances in protoplanetary discs. Here we have constructed systems where resonance encounters were impossible to avoid: requiring that we assumed the planets were coplanar and existed near these resonances without going unstable. A natural question this paper does not address is whether the dynamics that might occur due to relative inclinations affects the overall stability of the resonant systems. It has been shown \\citep{thommeslissauer03} that it is possible for planets starting in a 2:1 resonance to later enter an inclination resonance, which generates large mutual inclinations even in an initially almost coplanar system.  It is not clear \\textit{a priori}  whether this would promote dynamical instability or protect against it. Additionally, we have modeled the eccentricity damping time scale as a constant proportional to the migration time scale. Recent studies \\citep{ogilvielubow03,moorheadadams07} have shown that damping may not be uniform and depends on the local properties of the disk as well as the overall geometry of the planets. In particular, $de/dt$ may even change sign depending on the value of the eccentricity. Therefore, it is important to better understand the nature of eccentricity growth and damping in order to understand the long-term stability of these resonances in the protoplanetary disk.  Our improved distribution shown in Figure \\ref{fig:ResCumHisto} was derived by using an equal number of cases from each of the three resonances studied. However, Table \\ref{tab:resoutcome} suggests that we should find more planets in the 2:1 resonance compared to the 3:2 resonance during the disk's lifetime when we assume that all the planets migrated from a region farther out than the 2:1 resonance. In the case of gravitational instability, it may be assumed that planets could form in a more equally distributed manner before undergoing migration and locking into these resonances. Complete three dimensional integrations would allow us to gain an accurate representation of the distribution of resonance captures due to migration alone. Further stability tests will involve allowing one planets to migrate from greater distances, where near-resonance interactions cannot affect the overall stability of the system until the planets (given they remain stable) can migrate close enough together to encounter a particular resonance.  Additionally, gap-forming planets on the edge of a disk in general ought to still be accreting some mass across their gap, and therefore mass growth will be taken into consideration. In this case, planets could lock into resonance under conditions we have shown to be stable, and through mass accretion enter a regime that analytic and numerical criteria determine to be unstable. A comparison of whether this region of stability is similar to what we have found here for constant-mass planets would give us further important insights into the possible range of planetary system configurations.  We have gained a clearer picture of the range of possible dynamics within a protoplanetary disk. In doing so, we have seen how resonances can strongly affect planetary systems, long before the surrounding disk material has dissipated. This suggests that what happens before the disk disappears is key in determining a planetary system's ultimate dynamical architecture. One possible outcome is a system which, after the disk is gone, is left with planets in resonance.  In order for this to happen, the planets must enter these resonances relatively late in the disk's lifetime, so that they ultimately survive the coupled migration.  After the disk has dissipated, some resonances may eventually become unstable. Another possibility we have demonstrated is that resonances can be broken by dynamical instability that occurs while the disk is still present.  Our results allow us to make a connection between the planet formation process and observations of mature, resonant planetary systems, none of which have yet been observed in a closer than 2:1 resonance:  For planets that form further apart than a period ratio of 2:1 and migrate convergently, the very stable 2:1 constitutes a formidable barrier.  Between 2:1 and 5:3, the initial outcome is likely to be capture into 5:3, followed by a one-way trip to instability.  The 3:2 resonance is quite stable (unless eccentricity damping by the disk is absent), so provided that an appreciable number of planets form with a primordial period separation of less than the 5:3, we would expect to find some planets in a 3:2 resonance.  Thus, if future observations continue to reveal no such systems, then this may suggest a lower limit to orbital separations with which neighboring planets first form."
    },
    "0801/0801.4732_arXiv.txt": {
        "abstract": "{This is the third of a series of papers presenting the first attempt to analyze the growth of the bar instability in a consistent cosmological scenario. In the previous two articles we explored the role of the cosmology on stellar disks, and the impact of the gaseous component on a disk embedded in a cosmological dark matter halo.} {The aim of this paper is to point out the impact of the star formation on the bar instability inside disks having different gas fractions.} {We perform  cosmological simulations of the same  disk-to-halo mass systems as in the previous works where the star formation was not triggered. We compare the results of the new simulations with the previous ones to investigate the effect of the star formation by analysing the morphology of the stellar  components, the  bar strength, the behaviour of the pattern speed. We follow the gas and the central mass concentration during the evolution and their impact on the bar strength.} {In all our cosmological simulations a stellar bar, lasting 10 Gyr, is still living at z=0.\\\\ The central mass concentration of gas and of the new stars has a mild action on the ellipticity of the bar but is not able to destroy it; at z=0 the stellar bar strength is enhanced by the star formation. The bar pattern speed is decreasing with the disk evolution.} {} ",
        "introduction": "The connection between the bar feature and the star formation  process has already been pointed out in the past: \\citet{mart77} observed that non interacting galaxies  displaying the highest star forming activity have strong and long bars. Conversely not all strong and long bars are intensively creating stars. \\citet{Ma01}  performed SPH simulations to investigate the dependence of the star formation in disks on the geometry and dynamical state of the DM halo. They showed that the star formation lengthens the life time of the bar. However in  their works the evolution of the disk+halo system arises in a isolated framework, outside the cosmological scenario.\\\\ The treatment of the star formation in galaxies involves a large number of physical processes  which arise with different time scales  and  spatial lengths. Therefore its inclusions in simulations of galaxy formation is complex (\\citet{MaCu03} and references therein) and many sub-grid parametrisations (i.e. representing all the astrophysical processes that cannot be resolved due to resolution or computing power limitations) of the star formation process, exist in literature. To mimic all these processes with a numerical model of a realistic disk galaxy is difficult too. Thus the numerical work on this subject focuses either on a detailed analysis of the interstellar medium (ISM) and a smaller simulated area \\citep{Wa01, avil2000} or on the study of the global instabilities and of the star formation at the cost of simplified ISM models \\citep{Spri04,  Sem02, Comb02}. Models that have tackled  both  topics have either been in two dimensions \\citep{WaNo01} or restricted to a box size a few hundred parsecs across \\citep{Wa01}. \\citet{Ros95} worked in two dimensions allowing the ISM to evolve self-consistently but treated the stars as collisional rather than collisionless fluid. \\citet{task06} presented the first three dimensional simulations of a global disk without the need to simplify the structure of the ISM. Their paper is devoted to understand the fundamental processes  of star formation and feedback in a disk galaxy.\\\\ Here we  performed cosmological simulations to analyse the effects of the star formation process inside stellar-gaseous disks embedded in a DM cosmological halo evolving in a fully consistent scenario. We employ the sub-grid star formation model by \\citet{Spri03} which is able to produce a self-regulated star formation in galactic disks. We point out that our model is not a general galaxy evolution model, since the gradual formation and growth of the stellar disk has not been taken into account. However our approach,  developed in two previous papers \\citep{Cu06, Cu07}, allows us  to investigate the effects of different parameters like the disk-to-halo mass ratio and the gas fraction inside the disk, on the growth of the bar instability, its coupling with the star formation rate  and  its dependence on such parameters in a self-consistent cosmological framework. We compare the results of  such a new set of simulations with those of  our previous sets with the star formation switched off (\\citet{Cu06}, hereafter Paper 1, \\citet{Cu07}, hereafter Paper 2). In Paper 1 we presented simulations  of purely stellar disks with the same disk-to-halo mass ratios, in the same cosmological scenario, and with the same initial conditions as in this new paper. In Paper 2 we investigated the growth and the evolution of the bar instability in disks with the same disk-to-halo mass ratios as in Paper 1, and different gas fractions, without star formation. Thus, we will compare the results of simulations here performed, with those corresponding to the same disk-to-halo mass ratio as in Paper 1 and to the same gas fraction as in Paper 2. The comparison between the parameters characterising the bar can be done only evaluating that of the old stars component, since the new stars are not present in Paper 1  and  2. In both such Papers it was shown that in DM dominated disks a bar feature is triggered and maintained by the Cosmology whereas, in the more massive disks, a gas fraction 0.2 is able to destroy the bar. The focus of the present work is to determine if the  star formation changes such a result.  The plan of the paper is the following: in Section 2, we  summarise our recipe for the initial $disk+halo$ system fully described in Paper 1,  and present our star formation recipe. In Section 3  there are our cosmological simulations, in Section 4 we point out our  results. The parameters related to the bars formed in the new and the old stars and to the global, old+new, stellar populations are given in such a section. Section 5 is devoted to our discussion and conclusions. It contains a Table including the parameters characterising the old stellar bar at z$=0$ in the  three Papers to allow a more easy comparison between the results of  our Papers (Table 2). ",
        "conclusions": "We presented six  cosmological simulations  with the same  disk--to--halo mass ratios   as  in  Paper 1 and Paper 2. In order to study the impact of the forming stars from the gaseous   component, here we included and varied its percentages inside disks of different disk-to-halo mass ratios, as in  Paper 2 where the star formation was switched off.  In Table \\ref{fin_table} we show the final data for ellipticities and semimajor axes resuming the results of our three papers.  The old star component show  a long lasting bar, 10 Gyr old, in all the simulations  of this work. \\begin{table*} \\caption{Cosmological disks simulations: global results} \\label{fin_table} \\centering \\begin{tabular}{c c c c c c } \\hline\\hline Disk mass& gas fraction & {$   $} no star formation & {}& star formation & {}\\\\ {}& {}  & ellipticity& $a_{max}$ & ellipticity & $a_{max}$ \\\\ \\hline 0.33 & 0. &  0.52 & 5 & - & - \\\\ 0.33&  0.1 &  0.68 & 8.4 & 0.65 & 8 \\\\ 0.33& 0.2 &  0.1 &  no bar  & 0.55 & 11 \\\\ 0.33& 0.4 & 0.07 &  no bar & 0.6 & 8.4 \\\\ 0.1& 0. &  0.3 & 6.5 & - & - \\\\ 0.1&0.1 &  0.58 & 5.8 & 0.39 & 3 \\\\ 0.1& 0.2 &  0.6 & 5.4  & 0.45 & 3 \\\\ 0.1&  0.6 & 0.42 & 5.8 & 0.5 & 3 \\\\ \\hline \\end{tabular} \\end{table*} Moreover, in all such simulations,  except for simulation c6, the  bar is stronger than that developed in the pure stellar case with the same disk-to-halo mass ratio but weaker than that formed in the case of the same gas fraction without star formation.  We find that the star formation, reducing the central gaseous mass concentration, allows the bar to survive until the end of the evolution also in the more massive disks, at variance with the results in Paper 2 (where the gas was not allowed to form stars):  in such a case a gas fraction  0.2  was able to destroy the bar. Even if  some details of the disk morphologies  obviously depend on the star formation prescription, the  bars arising here  in the c2 and c3 simulations are   due to the reduced central gas concentration and thus to the weakening of its effect on the bar itself. We verified this point  by rerunning simulation c2 changing the star formation prescription.  Instead of the GADGET-2 effective model, we used a simple algorithm in which, when a gas particle reaches a given density threshold  with a temperature lower than a minimum temperature threshold, it is converted into a star particle, as in  \\citet{Katz96}. Fig. \\ref{katz} shows the morphology of the old stellar component for such a  different star formation recipe: the survival of the bar is due to the star formation, independently on its details.\\\\ In all the simulations of more massive disks, the new stellar component shows a barred shape coupled with that of old stars but with semimajor axis and ellipticity values smaller than those. \\begin{figure} \\begin{center} \\includegraphics[width= 7cm]{isoold_xyz_levfixG179_12lev_rot35_Katz.ps} \\end{center} \\caption{ Isodensity contours of the old stars at z=0, for the  simulation c2, rerunned with the Katz'96  star formation recipe} \\label{katz} \\end{figure} In  all the simulations  of DM dominated disks, a  bar feature is maintained in  the old stellar component at the disk center, and the new stars form a spheroidal bulge that  hides the  bar of the old stars.  This effect could be of interest as far as the  observations about the Milky Way are concerned where, in addition to the main bar, a small nuclear bar consisting of old stars seems to be included in the bulge (see e. g. \\citet{ala01}).  The final  strenght of these old star bars  increases by  increasing gas fraction  and their pattern speed  is quickly decreasing before z$=1$.   The classical results obtained outside the cosmological scenario are no longer applicable. This conclusion  remarks the results of Paper 1, where it was shown that in the DM dominated disks the bar feature is triggered and maintained by the cosmological  properties of the halo (namely its triaxiality and its dynamical state).  From the models presented here we suggest that a very low pattern speed (few $ Km\\, s^{-1}\\, Kpc^{-1}$) could be a  signature of  a dominating halo and not a classical product of the disk instability.  The whole set of cosmological simulations we  presented here is suggesting that the star formation works in favour of maintaining the bar feature in self gravitating disks,  against the effect of increasing gas fractions,  whereas in  DM  dominated disks, it slightly reduces the bar strength. Moreover, inside such  disks, the new  stars  are arranged in non barred systems which are decoupled from the  bar  of the old star component.   {\\bf Acknowledgements} Simulations have been performed on the CINECA IBM CLX Cluster, thanks to the INAF-CINECA grants cnato43a/inato003 ``Evolution of disks in cosmological contests: effect of the star formation inside the disk'', and on the Linux PC Cluster of the Osservatorio Astronomico di Torino. We wish to thank   V. Springel for kindly providing us with his code GADGET, and Martina Giovalli which implemented the simple star formation prescription in it.  We thank the referee for his suggestions, useful to improve the paper."
    },
    "0801/0801.2578_arXiv.txt": {
        "abstract": "The frequency of microlensing planet detections, particularly in difficult-to-model high-magnification events, is increasing. Their analysis can require tens of thousands of processor hours or more, primarily because of the high density and high precision of measurements whose modeling requires time-consuming finite-source calculations.  I show that a large fraction of these measurements, those that lie at least one source diameter from a caustic or the extension from a cusp, can be modeled using a very simple hexadecapole approximation, which is one to several orders of magnitude faster than full-fledged finite-source calculations.  Moreover, by restricting the regions that actually require finite-source calculations to a few isolated ``caustic features'', the hexadecapole approximation will, for the first time, permit the powerful ``magnification-map'' approach to be applied to events for which the planet's orbital motion is important. ",
        "introduction": "\\label{sec:intro}}  Microlensing planets are being discovered at an accelerating rate, with one being reported in 2003 \\citep{ob03235}, three in 2005 \\citep{ob05071,ob05390,ob05169}, two in 2006 \\citep{gaudi08}, and perhaps as many as six in 2007.  It has been a huge challenge for modelers to keep up with these discoveries, in large part because the computing requirements are often daunting: the parameter space is large, the $\\chi^2$ surface is complex and generally contains multiple minima, and the magnification calculations are computationally intensive. Proper treatment of individual events can require tens of thousands processor hours, or more.  Indeed, some potential planetary events have still not been fully modeled because of computational challenges.  Planetary microlensing events are recognized through two broad channels, one in which the lightcurve perturbation is generated by the so-called ``planetary caustic'' that is directly associated with the planet \\citep{gouldloeb92} and the other in which it is generated by the ``central caustic'' that is associated with the host star \\citep{griestsafi}. The former events are relatively easy to analyze and indeed the event parameters can be estimated reasonably well by inspection of the lightcurve. The latter events are generally much more difficult.  There are several interrelated reasons for this.  First, planets anywhere in the system can perturb the central caustic.  This is why these events are avidly monitored, but by the same token it is often not obvious without an exhaustive search which planetary geometry or geometries are responsible for the perturbation. Second, this very fact implies that several members of a multi-planet system can be detected in central-caustic events \\citep{gaudi98}. Multiple planets create a larger, more complicated parameter space, which can increase the computation time by a large factor. Even if there are no obvious perturbations caused by a second planet, an exhaustive search should be conducted to at least place upper limits on their presence.  Third, if the source probes the central caustic, it is ipso facto highly magnified.  Such events are brighter and more intensively monitored than typical events and so have more and higher-quality data.  While such excellent data are of course a boon to planet searches, they also require more and more-accurate computations, which requires more computing time.  Fourth, central-caustic planetary events are basically detectable in proportion to the size of the caustic, which roughly scales $\\propto q/|b-1|$, where $q$ is the planet/star mass ratio and $b$ is the planet-star separation in units of the Einstein ring. Thus, these events are heavily biased toward planets with characteristics that make the caustic big.  Such big caustics can undergo subtle changes as the planet orbits its host, and in principle these effects can be measured, thus constraining the planet's properties.  Exploration of these subtle variations requires substantial additional computing time. Finally, there is also a bias toward long events, simply because these unfold more slowly and so increase the chance that they will be recognized in time to monitor them intensively over the peak. Such long events often display lightcurve distortions in their wings due to the Earth's orbital motion which, if measured, can further constrain the planet properties.  However, probing this effect (called ``microlens parallax'') requires yet another expansion of parameter space. Moreover, the ``parallax'' signal must be distinguished from ``xallarap'' (effects of the source orbiting a companion), whose description requires a yet larger expansion of parameter space.  There are two broad classes of binary-lens (or triple-lens) magnification calculations: point-source and finite-source.  The former can be used whenever the magnification is essentially constant (or more precisely, well-characterized by a linear gradient) over face of the source, while the latter must be used when this condition fails.  Point-source magnifications can be derived from the solution of a 5th (or 10th) order complex polynomial equation and are computationally very fast.  When the point-source is well separated from the caustics (as it must be to satisfy the linear-gradient condition) then this calculation is also extremely robust and accurate.  The main computational challenge in modeling planetary events comes from the finite-source calculations.  Almost all integration schemes use inverse-ray shooting, which avoids all the pathologies of the caustics by performing an integration over the {\\it image} plane (where the rays behave smoothly) and simply asks which of the rays land on the source.  The problem is that a large number of rays must be ``shot'' to obtain an accurate estimate of the magnification, which implies that high-quality data demand proportionately longer computations.  Of course, the higher the magnification, the larger the images, and so the more rays are required.  There are various schemes to expedite inverse ray-shooting, including clever algorithms for identifying the regions that must be ``shot'' and pixelation of the source plane.  The bottom line is, however, that the overwhelming majority of computation time is spent on finite-source calculations.  Here I present a third class of binary-lens (or triple-lens) computation that is intermediate between the two classes just described: the hexadecapole approximation.  \\citet{pejcha07} expand finite-source magnification to hexadecapole order and illustrate that the quadrupole term by itself can give quite satisfactory numerical results. In this paper, I develop a simple prescription for evaluating this expansion.  While this algorithm is about 10 times slower than point-source calculations, it is one to several orders of magnitude faster than finite-source calculations.  The method can be applied whenever the source center is at least two source radii from a caustic or the extension from a cusp. Typically, well over half the non-point-source points satisfy this condition, meaning that the method can reduce computation times by a factor of several.  Moreover, by isolating the small regions of the lightcurve where finite-source calculations must be used, the method opens up the possibility that the finite-source computations themselves can be radically expedited for the special, but very interesting, class of events in which planetary orbital motion is measured. ",
        "conclusions": "\\label{sec:conclude}}  I have identified an intermediate regime between the ones where finite-source effects are dominant and negligible, respectively. In this regime, magnifications can be evaluated with very high precision using a simple hexadecapole approximation, for which I give a specific prescription.  Outside of small (few source-diameter crossing times) regions associated with caustic crossings and cusp approaches, the approximation has a fractional error of well under 0.1\\%. Some events can now be analyzed without any traditional finite-source calculations, while for others these calculations will be drastically reduced.  In particular, by restricting the regions that do absolutely require finite-source calculation to a few isolated zones, this approximation opens the possibility of applying ``magnification maps'' to planetary systems experiencing significant orbital motion."
    },
    "0801/0801.4865_arXiv.txt": {
        "abstract": "We construct self-consistent dynamical models for disk galaxies with triaxial, cuspy halos.  We begin with an equilibrium, axisymmetric, disk-bulge-halo system and apply an artificial acceleration to the halo particles.  By design, this acceleration conserves energy and thereby preserving the system's differential energy distribution even as its phase space distribution function is altered.  The halo becomes triaxial but its spherically-averaged density profile remains largely unchanged.  The final system is in equilibrium, to a very good approximation, so long as the halo's shape changes adiabatically.  The disk and bulge are ``live'' while the halo is being deformed; they respond to the changing gravitational potential but also influence the deformation of the halo.  We test the hypothesis that halo triaxiality can explain the rotation curves of low surface brightness galaxies by modelling the galaxy F568-3. ",
        "introduction": "Dark matter halos -- at least the ones found in cosmological simulations -- have a number of universal traits.  Most famously, their density profiles have a shape that is nearly independent of mass, formation epoch, and cosmological model \\citep{nfw96}.  Their angular momentum distribution \\citep{bullock01}, phase space density \\citep{taylor01}, and velocity anisotropy \\citep{hansen06} profiles also appear to follow universal forms.  In other respects, halos are rather diverse.  Simulated halos are typically triaxial with axis ratios that range from $0.6$ to $1$ (See \\citet{dubinski91, warren92} and more recently, \\citet{novak06}).  The shapes of real halos are more difficult to determine but promising observational approaches do exist.  Probes of the Galactic halo include flaring of the gas disk \\citep{olling} and tidal streams of satellite galaxies \\citep{johnston99}.  The shapes of halos in other galaxies can be determined, at least statistically, by weak gravitational lensing surveys \\citep{hoekstra, mandelbaum, parker}.  On the other hand, triaxiality can bias attempts to determine a halo's density profile from the rotation curve of the disk that sits within it \\citep{hayashi06}.  Our main goal in writing this paper is to introduce a novel scheme to generate self-consistent dynamical models for disk galaxies with triaxial halos.  Our models can be tailored to fit observational data for specific galaxies and therefore provide a testing ground to study the disk-halo connection.  We consider the effects of halo triaxiality on the rotation curves of low surface brightness galaxies (LSBs) and briefly discuss other applications of the models.  Halos in simulations of a cold dark matter (CDM) universe have central cusps with $\\rho\\propto r^{-\\gamma}$ where $\\gamma\\simeq 1$ \\citep{nfw96}.  In dark matter-dominated galaxies, this density profile would seem to imply a rotation curve where $v\\propto r^{1/2}$ as $r\\to 0$.  By contrast, a halo with a constant density core implies $v\\propto r$ as $r\\to 0$.  LSBs, which are believed to be dark-matter dominated at small radii, have rotation curves that generally favor a constant density core over a $\\gamma=-1$ cusp.  This result represents one of the most serious challenges to the CDM scenario \\citep{moore94, flores94, mcg98} and has inspired some rather exotic alternatives. De Blok \\& McGaugh (1998), for example, suggested that LSB rotation curves could be explained by Modified Newtonian Gravity while \\citet{firmani00} and \\citet{mo00}) invoked dark matter self-interactions to flatten the central cusp of the halo.  The connection described above between a galaxy's rotation curve and the intrinsic density profile of its halo assumes that the halo is spherically symmetric.  However, if a galaxy's halo is triaxial, then gas in the disk will move on non-circular orbits and, under certain conditions, the observed rotation curve will rise approximately linearly even if the intrinsic halo density profile has a steep cusp. \\citet{hayashi06} and \\citet{hayashi07} presented this argument as a means of reconciling LSB rotation curves with the predictions of the CDM model.  \\citet{hayashi06} and \\citet{hayashi07} derived model rotation curves by calculating closed orbits in the potential generated by a triaxial halo.  In this paper, we derive rotation curves by making pseudo-observations of a disk that is embedded in the halo. Deviations from axial symmetry in disk and halo are generated concurrently and self-consistently.  A number of methods exist for constructing models of, and embedding disks in, triaxial halos.  For example, \\citet{moore04} show that the remnant of a major merger between two equilibrium spherical halos is triaxial.  \\citet{bailin07} describe how to set up an equilibrium disk in a combined halo-disk potential.  Our approach produces an N-body galaxy complete with disk, bulge, and triaxial halo.  (Central black holes may also be included, as in \\citet{wid05}.)  It is inspired by the method outlined in \\citet{holley01}.  In that scheme, dubbed ``adiabatic squeezing'', particles of an equilibrium halo are subjected to an artificial drag by modifying the force law of a standard N-body code and evolving the system forward in time.  A triaxial halo is created if the drag has a different strength along three orthogonal directions, that is, along what become the three principle axes of the halo.  The final model will be in equilibrium, to a good approximation, so long as the timescale for the halo's shape to change is slow as compared to the typical orbital timescale of the system.  Adiabatic squeezing causes a halo to shrink in size.  For an isolated halo, this shrinking can be reversed by simply rescaling the positions and velocities of the particles.  Obviously, this method is unsuitable for disk-bulge-halo systems since the disk and bulge would be disrupted in an unphysical way.  We propose a modification of this method in which drag is applied along one axis and ``negative drag'' is applied along the other two or vice versa depending on whether one wants a prolate or oblate halo.  We require that for each particle, the change in energy due to the artificial drag force is zero.  In this way, we change the phase space distribution of the particles but not their energy distribution.  As noted in \\citet{BT}, {\\it if two systems have the same energy distribution, their spherically-averaged density profiles will be very similar even if their phase space distribution functions are different.}  Our starting point is the equilibrium model of \\citet{wid07} which comprises a Sersic bulge, cuspy dark halo, and exponential disk.  The model is described in terms of a phase space distribution function (DF) which, in turn, is a function of the integrals of motion.  In the current version of the model, the halo component of the DF depends only on the energy.  In the absence of a disk, the halo is spherically symmetric.  With the disk included, the halo is flattened slightly but is still axisymmetric.  Adiabatic deformation allows us to extend our disk-bulge-halo model to systems with triaxial halos.  In Section 2, we describe our method and construct an example of an isolated triaxial halo.  We then consider the LSB galaxy F568-3.  In Section 3, we present axisymmetric, equilibrium models for this galaxy based on its published surface brightness profile and circular speed curve.  In Section 4, we show how transforming the axisymmetric halo in one of these models into a triaxial halo changes the shape of the rotation curve.  We conclude in Section 5, by summarizing our results and briefly discussing further applications of the method. ",
        "conclusions": "The adiabatic squeezing method \\citet{holley01} produces triaxial halos that have shrunk in size and therefore requires that the positions and velocities of the particles be rescaled.  This awkward step precludes the technique from being applied to compound systems. Our approach avoids this problem by using an {\\it energy-conserving} artificial force to deform the halos.  Our analysis of the LSB galaxy F568-3 begins with a discussion of axisymmetric models.  We attempt to fit both photometric and kinematic observations using Bayesian statistics and the MCMC method.  Our excellent fit of the surface brightness profile requires a bulge and disk truncation, neither of which were included in previous studies. As for the rotation curve, we find that constant density cores do better than density cusps in agreement with earlier studies of LSBs.  The second stage of our analysis is to deform the halo of a compound system.  In agreement with \\citet{hayashi06}, we show that the rotation curve of F568-3 may indicate the presence of a triaxial halo rather than a problem with the standard CDM model of structure formation.  \\citet{hayashi06} and \\citet{hayashi07} construct rotation curves by finding closed orbits in the gravitational potential of a triaxial halo.  We calculate the rotation curves by making pseudo-observations of a disk that is self-consistently embedded in a dark halo.  There are two improvements that will add a further level of realism to the analysis: the inclusion of a gas disk in the galactic models and an iterative scheme whereby the initial model and artificial acceleration parameters are adjusted so that the final model fits the data in detail.  These improvements will be considered in a future publication.  Our triaxial models have a wide range of applications.  For example, they can be used to study the effect a non-spherical halo has on the morphology of tidal streams from satellite galaxies and flaring and warping of the gas disk.  The method can also be applied to bulges where departures from axial symmetry are thought to be important."
    },
    "0801/0801.3959_arXiv.txt": {
        "abstract": "We demonstrate the ability to measure precise stellar barycentric radial velocities with the dispersed fixed-delay interferometer technique using the Exoplanet Tracker (ET), an instrument primarily designed for precision differential Doppler velocity measurements using this technique. Our barycentric radial velocities, derived from observations taken at the KPNO 2.1 meter telescope, differ from those of Nidever et al. by 0.047 km s$^{-1}$ (rms) when simultaneous iodine calibration is used, and by 0.120 km s$^{-1}$ (rms) without simultaneous iodine calibration. Our results effectively show that a Michelson interferometer coupled to a spectrograph allows precise measurements of barycentric radial velocities even at a modest spectral resolution of R $\\sim 5100$. A multi-object version of the ET instrument capable of observing $\\sim$500 stars per night is being used at the Sloan 2.5 m telescope at Apache Point Observatory for the Multi-object APO Radial Velocity Exoplanet Large-area Survey (MARVELS), a wide-field radial velocity survey for extrasolar planets around TYCHO-2 stars in the magnitude range $7.6<V<12$. In addition to precise differential velocities, this survey will also yield precise barycentric radial velocities for many thousands of stars using the data analysis techniques reported here. Such a large kinematic survey at high velocity precision will be useful in identifying the signature of accretion events in the Milky Way and understanding local stellar kinematics in addition to discovering exoplanets, brown dwarfs and spectroscopic binaries. ",
        "introduction": "The discovery of various tidal streams in our galaxy (Ibata et al. 1994; Yanny et al. 2006; Belokurov et al. 2006) challenges the monolithic collapse formation model (Eggen 1962), but the streams discovered so far form only a small fraction of the mass of the galaxy. Large kinematic surveys are necessary to determine if such merger and accretion events are more common. Radial velocity surveys using multi-object spectrographs, like RAVE (Steinmetz et al. 2006), are already producing barycentric radial velocities for many thousands of stars at a precision of $2-3$ km s$^{-1}$, enabling studies of galactic dynamics (e.g. Veltz et al. 2008). In the northern hemisphere the magnitude limited Geneva-Copenhagen Survey (Nordstrom et al. 2004) has provided radial velocities for over 14000 bright nearby F and G dwarfs ($V < 8.6$) at a precision of $100-300$ m s$^{-1}$. Coupled with gravity and chemical composition derived from the observed spectra, such datasets of precise barycentric radial velocities are very valuable in identifying past accretion events and substructure in the galactic disk (Helmi et al. 2006), finding members of moving groups (Jones 1970) and determining cluster membership.  The ability to determine precise differential radial velocities has also improved dramatically in the last few years, and planet search surveys are currently at a precision of $1-3$ m s$^{-1}$ (Rupprecht et al. 2004; Butler et al. 1996). Such planet search surveys, however, use an arbitrary zero point for the stellar velocities since only differential velocities are measured. The true radial velocity of the center-of-mass of a star is a difficult quantity to measure accurately using spectroscopy since it involves additional modelling (or assumptions) of effects like convective blue shift and gravitational redshift (Lindegren \\& Dravins 2003). A quantity that is easier to obtain, in terms of observables and known physical constants, is the apparent Doppler shift of the lines in the stellar spectra, as measured in the solar system barycenter, with the Sun setting the velocity zero point. This Doppler shift can be converted into a velocity that we define as the barycentric radial velocity of the star. Very precise barycentric radial velocities have been measured using high resolution echelle spectrographs by Udry et al. 1999 (UD99 hereafter) for a number of stable F, G, K dwarfs, as well as by Nidever et al. 2002 (ND02 hereafter) for a larger sample.  In this paper we demonstrate the ability of a dispersed fixed-delay interferometer (DFDI) to measure precise barycentric radial velocities. The Exoplanet Tracker (ET) is a prototype of this class of instruments, which use a Michelson interferometer followed by a medium resolution spectrograph working in the first grating order. Fixed delay interferometers have long been used by solar astrophysicists to measure solar oscillations (Gorskii \\& Lebedev 1977; Kozhevatov 1983; Harvey et al. 1995). In 1997 D. J. Erskine proposed adding a spectrograph as a post disperser for increasing the measurement precision of differential radial velocities (Erskine \\& Ge 2000; Ge, Erskine \\& Rushford 2000). When using this technique sinusoidal interference fringes are formed in the slit direction at the position of the stellar absorption lines. A shift in the dispersion direction of the underlying spectra due to a Doppler shift manifests itself as the phase shift of the sinusoidal fringes in the slit direction (Ge 2002).  Precise radial velocities can be measured with DFDI even at low wavelength dispersions since information is also being encoded as sinusoidal fringes in the slit direction. This results in each stellar spectrum taking up only a small fraction of the available number of pixels on the CCD detector used to record the spectrum. This feature of the DFDI technique is a great advantage since it has the potential to enable a multi-object instrument (similar to ET) to simultaneously obtain precise radial velocities for a large number of stars (Ge 2002; Ge et al. 2002) without the need for an expensive cross-dispersed high spectral resolution instrument. One such instrument, the W. M. Keck Exoplanet Tracker, is being used at the Sloan 2.5 m telescope for the magnitude limited Multi-object APO Radial Velocity Exoplanet Large-area Survey (MARVELS), a survey for extrasolar planets around TYCHO-2 stars in the magnitude range $7.6<V<12$ that is being conducted as part of SDSS III (Weinberg et al. 2007). This instrument can observe 59 stars simultaneously, and approximately 500 stars per night. Designed primarily as a wide-field precision radial velocity survey to discover short and intermediate period planets, MARVELS will also provide a dataset of precise barycentric radial velocities for thousands of late F, G, and K TYCHO-2 stars. ",
        "conclusions": "We have demonstrated the ability to measure barycentric radial velocities at a high level of precision (rms = $50 -120$ m s$^{-1}$) using a dispersed fixed delay interferometer instrument. Our results also show that the radial velocities derived using a very different technique and an instrument at substantially lower resolution are not very different in precision from those derived using high resolution spectroscopy and cross-correlation with stellar (ND02; UD99) or Fe I line templates (Gullberg \\& Lindegren 2002). Simultaneous calibration with the iodine cell yields more precise velocities, but the data analysis using only the templates (Section 3.1, 4.1) may be more appropriate for faint stars where the low signal to noise may introduce significant errors when extracting velocities from the star+iodine data. We have corrected the observed radial velocities to the barycenter of the solar system. Using the corrected $\\eta$ Cas radial velocity from ND02 to set the velocity scale for our radial velocities then automatically corrects all velocities for the gravitational redshift and convective blueshift of the Sun (used by ND02 as the zero point). By adopting the radial velocity of $\\eta$ Cas from ND02 we have automatically adopted ND02 as our velocity reference for the absolute calibration of a zero-point and performed a secondary measurement to determine the velocities of all our other target stars. No attempt is made to further correct the stellar radial velocities for spectral type and luminosity class dependent effects. This results in our derived velocities having typical accuracies of $0.2-0.3$ km s$^{-1}$ (derived from ND02, our velocity reference) for late F, G and K dwarfs, and precisions of $\\sim 50-100$ m s$^{-1}$ when comparing stars of similar spectral types. With the ET instrument scheduled to be a facility instrument at the KPNO 2.1m, this velocity measurement capability will be available to the wider astronomical community. These techniques will also be applied to data from the MARVELS survey that will observe $\\sim$ 10,000 stars as part of SDSS III, leading to precise barycentric radial velocities being measured for the majority of the stars in the sample.  Future work with these instruments will focus on using the day sky spectra of the Sun as a stellar template to determine the velocity zero point, allowing independent comparison with the radial velocities of ND02 and UD99 without having to adopt either of these works as the velocity reference, and a better understanding of any zero point differences that may arise due to the much lower spectral resolution (Dravins \\& Norlund 1990) of ET compared to high resolution echelle spectrographs."
    },
    "0801/0801.1995_arXiv.txt": {
        "abstract": "From the Sloan Digital Sky Survey (SDSS) Data Release 5 (DR5), we extract a sample of 4594 galaxies at redshifts $0.02<z<0.03$, complete down to a stellar mass of $M=10^{10}~M_{\\odot}$. We quantify their structure (S\\'ersic index), morphology (S\\'ersic index + ``Bumpiness''), and local environment.  We show that morphology and structure are intrinsically different galaxy properties, and we demonstrate that this is a physically relevant distinction by showing that these properties depend differently on galaxy mass and environment.  Structure mainly depends on galaxy mass whereas morphology mainly depends on environment. This is driven by variations in star formation activity, as traced by color, which only weakly affects the structure of a galaxy but strongly affects its morphological appearance.  The implication of our results is that the existence of the morphology-density relation is intrinsic and not just due to a combination of more fundamental, underlying relations. Our findings have consequences for high-redshift studies, which often use some measure of structure as a proxy for morphology. A direct comparison with local samples selected through visually classified morphologies may lead to biases in the inferred evolution of the morphological mix of the galaxy population, and misinterpretations in terms of how galaxy evolution depends on mass and environment. ",
        "introduction": "The morphology-density relation \\citep[MDR,][]{dressler80} implies that the environment affects the star formation history, color, and structure of galaxies.  Many possible mechanisms have been suggested to explain the suppressed or quenched star formation activity of galaxies located in dense environments, most notably ram pressure stripping \\citep{gunn72}, harassment \\citep{farouki81,moore96}, strangulation \\citep[e.g.,][]{larson80,kauffmann93,diaferio01}, and tidal interactions and merging \\citep[e.g.,][]{park07b}.  At the same time, it is now also well established that strong correlations exist between galaxy mass and, for example, color \\citep[e.g.,][]{baldry06} and star-formation rate \\citep[e.g.,][]{brinchmann04}. Several suggestions have been made to explain the dependence of star formation history on galaxy mass, all of which are related to feedback mechanisms via either supernovae \\citep[e.g.,][]{white91}, AGNs \\citep[e.g.,][]{croton06}, or shock heating of infalling gas \\citep{dekel06}.  Galaxy structure as measured by, e.g., concentration or S\\'ersic index, behaves differently from color and star formation in the sense that structure depends strongly on galaxy mass but only weakly on environment \\citep{hogg04,kauffmann04}.  In the context of the MDR this may be surprising as it is clear that concentration and morphology are closely related quantities \\citep[e.g.,][]{bell03}. One possibility is that morphology and concentration only correlate with environment through the underlying correlation between environment and galaxy mass.  Another possibility is that morphology and concentration are intrinsically different galaxy properties. For example, it has been suggested \\citep[by, e.g.,][]{kauffmann04} that morphology is strongly related to star formation, more so than concentration.  \\citet{blakeslee06} developed an automated, quantitative scheme, the $B$-$n$ method ($B$ for ``Bumpiness'', and $n$ for \\citet{sersic68} index, see Sec. 2), to distinguish E+S0 galaxies from later types in \\textit{Hubble Space Telescope} imaging of distant clusters.  In \\citet{vanderwel07b} (hereafter, vdW07) we generalize the $B$-$n$ method to classify galaxies both at high redshift and in the Sloan Digital Sky Survey (SDSS).  In both cases the classifications from the $B$-$n$ method agree very well with visually determined morphologies as long as only two types are considered. In the current Letter we study, for a local sample extracted from the SDSS, the relation between structure (in this Letter measured by S\\'ersic index $n$), morphology (measured by $n$ and $B$), color, mass, and environment. This allows us to decide which factors determine the morphological appearance of a galaxy, and provides insight into whether or not the mechanisms that are responsible for morphological transformations are identical to the processes that suppress star formation.  We use the Fifth Data Release (DR5) from the SDSS \\citep{adelman07} and adopt the cosmological parameters $(\\Omega_{\\rm{M}},~\\Omega_{\\rm{\\Lambda}},~h) = (0.3,~0.7,~0.7)$.  The stellar masses used in this Letter are calculated for a ``diet'' \\citet{salpeter55} IMF \\citep{bell03}.  \\begin{figure} \\epsscale{1.2} \\plotone{f1.eps} \\caption{ Rest-frame $g-r$ color vs. S\\'ersic index $n$.  Magenta points are early-type (E+S0) galaxies, and green circles are late-type (Sa+later) galaxies, both according to the $B$-$n$ method. There are significant number of late-type galaxies with high S\\'ersic indexes.} \\label{n_col} \\end{figure}  \\begin{figure} \\epsscale{1.2} \\plotone{f2.eps} \\caption{ Rest-frame $g-r$ color vs. stellar mass. \\textit{Top:} Magenta points are early-type galaxies, and green circles indicate late-type galaxies, both according to the $B$-$n$ method. \\textit{Bottom:} Black points are galaxies with $n>2.5$, and orange circles are galaxies with $n<2.5$. The dashed lines show the least-squares linear fit to the early-type galaxies, i.e., the magenta points in the top panel, iteratively rejecting $3\\sigma$ outliers. The solid lines are the same as the dashed lines but shifted blue-ward by $2\\sigma$.} \\label{M_gr} \\end{figure} ",
        "conclusions": "We have shown that structure (in this Letter quantified by the S\\'ersic parameter $n$) and morphology (quantified by $n$ and the bumpiness parameter $B$) are distinct galaxy properties. There is a significant number of galaxies with high S\\'ersic indexes but late-type morphologies (Fig. \\ref{n_col}). The physical significance of the difference between structure and morphology becomes apparent when their behavior as a function of galaxy mass and environment is analyzed. Structure mainly depends on galaxy mass whereas morphology, at fixed galaxy mass, depends on environment (Fig. \\ref{struct}).  The different behavior is linked to star formation activity, which only weakly affects the structure of a galaxy, but strongly affects its morphological appearance (Fig. \\ref{colorB}). This implies that the existence of the MDR is intrinsic, and is explained by the environmental dependence of star formation activity. This seems trivial; however, it means that the MDR is not simply the result of underlying correlations, i.e., the strong dependence of structure and star formation on galaxy mass, and the environmental dependence of the mass function.  The author gratefully acknowledges financial support from NASA grant NAG5-7697, and is indebted to support from John Blakeslee and Holland Ford. Stimulating discussions with Andrew Zirm and Marijn Franx helped shape this Letter. Attentive reading by and valuable comments from Andrew Zirm, Roderik Overzier, Alison Crocker, John Blakeslee, Bradford Holden, Ricardo Demarco, and Marc Postman have greatly improved the readability."
    },
    "0801/0801.3397_arXiv.txt": {
        "abstract": "Very massive primordial stars ($140\\ M_{\\odot} < M < 260\\ M_{\\odot}$) are supposed to end their lives as PISN. Such an event can be traced by a typical chemical signature in low metallicity stars, but at the present time, this signature is lacking in the extremely metal-poor stars we are able to observe. Does it mean that those very massive objects were not formed, contrarily to the primordial star formation scenarios ? Could it be possible that they avoided this tragical fate ?  We explore the effects of rotation, anisotropical mass loss and magnetic field on the core size of very massive Population III models. We find that magnetic fields provide the strong coupling that is lacking in standard evolution metal-free models and our $150\\ M_{\\odot}$ Population III model avoids indeed the pair-instability explosion. ",
        "introduction": "According to \\citet{HFWLH03}, the fate of single stars depends on their He-core mass ($M_{\\alpha}$) at the end of the evolution. They have shown that at very low metallicity, the stars having $64\\ M_{\\odot} < M_{\\alpha} < 133\\ M_{\\odot}$ will undergo pair-instability and be entirely disrupted by the subsequent supernova. This mass range in $M_{\\alpha}$ has been related to the initial mass the star must have on the main sequence (MS) through standard evolution models: $140\\ M_{\\odot} < M_{\\rm ini} < 260\\ M_{\\odot}$. Since we will present here a non-standard evolution, we will rather keep in mind the $M_{\\alpha}$ range. ",
        "conclusions": "The model we presented here is exploratory. We cannot draw general conclusions from it. But our model shows that heavy mass loss is possible even at $Z=0$, and the answer to our title's question is: \\emph{yes, under certain conditions, very massive stars can indeed avoid PISN}.  Some aspects need yet to be clarified. Is the WR mass loss rate we have used really valid? Can the CNO lines alone really drive a WR wind? In a more general perspective, we still lack a good mass loss recipe for the strict $Z=0$ case. A word of caution must also be cast on the inclusion of magnetic fields. The validity of the Tayler-Spruit dynamo is still under debate, and more work need to be done before we may confidently rely on results that have been obtained with the actual treatment. Moreover, there is not a clear consensus whether magnetic fields were present in the early Universe or not. In the Taylor-Spruit dynamo, the mechanism amplifies a pre-existing field, so if none were present at the start, we cannot use it.  Anyway, the physics used in the present model is today's ``state of the art'' and it is interesting to study what can be achieved with it. Our result is encouraging, because the computation has been accomplished with reasonable assumptions: the initial rotation rate used here was fast but not extreme, and the mechanisms called upon (anisotropy of the winds and magnetic fields) are not exotic ones, but two natural effects which arise when one treats properly the case of rotation.  After this first step, more work is needed. First, we have to refine the mass loss treatment at $\\Omega\\Gamma$-limit and check if the mass loss stays as strong as here. If it does so, we have to check if our result are valid in the whole PISN mass domain. Higher mass models should experience higher mass loss, but it is necessary to check if it would still be sufficient to help them avoid pair-instabilities."
    },
    "0801/0801.3674_arXiv.txt": {
        "abstract": "We study the spectrum of cosmological fluctuations in the D3/D7 brane inflationary universe with particular attention to the parametric excitation of entropy modes during the reheating stage. The same tachyonic instability which renders reheating in this model very rapid leads to an exponential growth of entropy fluctuations during the preheating stage which in turn may induce a large contribution to the large-scale curvature fluctuations. We take into account the effects of long wavelength quantum fluctuations in the matter fields. As part of this work, we perform an analytical analysis of the reheating process. We find that the initial stage of preheating proceeds by the tachyonic instability channel. An upper bound on the time it takes for the energy initially stored in the inflaton field to convert into fluctuations is obtained by neglecting the local fluctuations produced during the period of tachyonic decay and analyzing the decay of the residual homogeneous field oscillations, which  proceeds by parametric resonance. We show that in spite of the fact that the resonance is of narrow-band type, it is sufficiently efficient to rapidly convert most of the energy of the background fields into matter fluctuations. ",
        "introduction": "In recent years a lot of interest has focused on the interface area between superstring theory and cosmology. The reasons are twofold. Firstly, cosmology can provide a possible arena to test string theory observationally. Secondly, superstring theory may be able to resolve the conceptual problems from which the current realizations of scalar field-driven inflation suffer (see e.g. \\cite{RHBrev} for a discussion). In particular, a lot of attention has recently been devoted to attempts to obtain periods of inflationary expansion of space in compactifications of superstring theories to four space-time dimensions. Since such compactifications contain a large number of scalar fields and since some of those remain massless before supersymmetry breaking, string theory gives rise to promising candidates for the inflaton, the scalar field driving the inflationary expansion of space (see e.g. \\cite{Linderev,Cliffrev,JCrev,EvaLiam} for recent reviews of attempts to obtain inflation from string theory compactifications).  D-branes have played a particularly important role in the recent approaches to obtaining inflation from string theory \\cite{Dvali,Stephon,Shafi,Cliff,JuanGB,KKLMMT}. A particularly promising model is the D3/D7 brane inflation model, in which the inflaton is the separation between the two branes in the directions transverse to both branes \\cite{DHHK}. This model features a shift symmetry \\cite{shift} for the inflaton field which ensures that at the classical level, the potential is flat along the inflaton direction (this shift symmetry is broken under quantum corrections \\cite{kors}). Inflation is induced by supersymmetry breaking terms. Various aspects of this model have been studied previously \\cite{herd,D37strings,previous}.  The large number of light moduli fields in string inflation models \\footnote{These moduli in general could be the complex or K\\\"ahler structure moduli of the internal space. They could also be the moduli of the branes in our theory. Once the K\\\"ahler and complex structure moduli are fixed by fluxes \\cite{GVW,DRS,GKP}, they could be the remaining brane moduli. In this paper, since we are not fixing any of the moduli, the light fields could come from all the above moduli.}, however, leads to a new danger: light fields which are not the inflaton are potential entropy modes and can give rise to entropy fluctuations. The ``curvaton'' \\cite{Mollerach,SY,LM,Moroi,LW,Sloth} and ``modulated reheating'' \\cite{DGZ,Lev,MR,Uzan,Vernizzi} scenarios are examples of this effect. Non-parametric generation of fluctuations from a global symmetry breaking (also involving an key isocurvature field) has been investigated in \\cite{KRV,Matsuda}.  As pointed out in \\cite{BaVi,FB2}, these entropy modes may undergo parametric instability during the initial stages of reheating. In particular, this instability occurs on super-Hubble (but sub-horizon) scales. This is possible \\cite{FB1}since the background fields carry the causal information to scales larger than the Hubble radius. The entropy fluctuations, in turn, will seed a curvature mode (which we will call ``secondary'' mode). There are two field scalar field toy models in which this secondary curvature mode dominates over the primary mode, the pure adiabatic linear perturbation theory mode. In fact \\cite{Zibin} the secondary mode can grow to be larger than unity before back-reaction shuts off the parametric instability which drives the growth of the entropy fluctuation. In this case, the model is ruled out by the observed (small) amplitude of curvature fluctuations on large cosmological scales.  The D3/D7 brane inflation model is a prototypical example in which there is an entropy mode which can undergo resonance during the initial stages of reheating, the preheating \\cite{JB1,KLS1,JB2,KLS2} phase. In this paper, we study the growth of this entropy mode and calculate the magnitude of the induced secondary curvature fluctuations. It turns out that for reasonable particle physics parameters, the secondary fluctuations remain in the linear regime at the end of the reheating phase. However, their amplitude can be (in the absence of back-reaction effects) larger than the amplitude of the primary mode. This would necessitate changing the parameters of the model in order to reproduce the observed magnitude of the large-scale cosmological perturbations.  Note that a similar conclusion was recently found \\cite{Larissa} in another model of brane inflation, the KKLMMT model \\cite{KKLMMT}. In that model, it had already been observed that due to their large phase space enhancement, second order fluctuations may give a dominant contribution \\cite{BC2}. We should emphasize that the effect we are calculating in this paper is an effect which arises in linear cosmological perturbation theory. Both the linear effects described here and the quadratic processes discussed in \\cite{BC2} and elsewhere (e.g. \\cite{Losic,Patrick}) can be operational at the same time. In fact, the enhanced growth of the linear perturbations which is the focus of this work will increase the strength of the quadratic processes. We expect that effects similar to those discussed in \\cite{BC2} will also arise in the D3/D7 system.  The structure of this paper is as follows: we first revisit the calculation of the amplitude of the spectrum of cosmological perturbations in the D3/D7 inflationary model based on the magnitude of the primary adiabatic mode alone. Next, we give an approximate but analytical study of reheating in the D3/D7 brane inflation scenario. The initial stages of reheating are governed by a tachyonic decay channel. The same tachyonic resonance channel leads to exponential growth of super-Hubble entropy modes, which then in turn induce a secondary curvature mode. The last major section of this paper focuses on the determination of the amplitude of this mode.  Our analysis of the reheating phase provides results which are interesting in their own right. A new aspects of our analysis is that we allow for an initial offset in the value of the ``waterfall\" field. Such an offset could arise as a consequence of the back-reaction of long wavelength \\footnote{Long means longer than the Hubble radius.} fluctuations of the matter fields. Such an offset will suppress the formation of domains on small scales. If the offset were sufficiently large, it could prevent the onset of the tachyonic instability. However, if we take the offset to be generated by the abovementioned fluctuations, we find that for most of the realistic space of parameters of the D3/D7 inflationary model, the resulting values of the field lie in the configuration space region for which the tachyonic resonance channel \\cite{tachyonic} is effective. Thus, the initial stages of the transfer of the energy in the inflaton field to matter fluctuations will proceed via the tachyonic instability in a time interval which is less than the characteristic time it takes for the matter fields to oscillate about their minima. A substantial part of the initial inflaton energy is converted to localized nonlinear fluctuations during the period of tachyonic resonance. Neglecting the back-reaction of these fluctuations on the homogeneous background, we see that a fraction of order unity of the initial inflaton energy density still remains in the background when the tachyonic instability shuts off. We are interested in demonstrating analytically that the transfer of energy from the inflaton to fluctuations is effective in rapidly draining almost all of the energy from the homogeneous fields. The dominant energy conversion processes are those related to the interactions of the nonlinear fluctuations. These have been studied numerically in great detail in \\cite{tachyonic}. Since our aim is to demonstrate analytically that within less than a Hubble time the energy can be effectively drained from the homogeneous background, we will neglect the nonlinear interactions and focus on the evolution of the homogeneous component of the fields once the tachyonic resonance shuts off (due to our approximation scheme, we will thus only obtain an upper bound for the decay time). We show that the decay of the residual homogeneous component of the inflaton field occurs via a parametric resonance instability \\cite{JB1,KLS1,JB2,KLS2}. In spite of the fact that the resonance turns out to be of narrow-band type, it leads to an efficient conversion of most of the energy from the background inflaton field to matter fluctuations.  One of the advantages of the D3/D7 brane inflation model is that the matter fields are directly associated with the branes whose separation provides the inflaton field\\footnote{Although not directly related to the main theme of the paper, we would like to point out that some recent papers \\cite{beasly} have shown how to get standard model with chiral fermions in a D3/D7 set-up from F-theory. Of course one need to change our background manifold to a certain Del-Pezzo surface to accomodate the standard model. It would be interesting to realise D3/D7 inflationary dynamics with this manifold.}. As a consequence, there is a direct coupling between the inflaton field and the matter fields, rendering it easy, as we will see, to obtain efficient reheating. This is a feature the D3/D7 brane inflation model shares with the D4/D6 model in which reheating was analyzed in \\cite{Easson} and with the brane-antibrane inflation models in which reheating was studied in \\cite{CFM,BC} (see also \\cite{TM}), but contrasts with the situation in inflationary models in which the inflaton dynamics and the standard model fields live in different throats, in which case efficient reheating is more difficult to obtain (see e.g. \\cite{BBC,KY,DSU,FMM,CT,PL}).  The outline of this paper is as follows. In the following section, we review the D3/D7 brane inflationary model, with particular emphasis on the form of the potential for the inflaton field. In Section 3 we discuss the vacuum manifold of the model, and in Section 4 we revisit the calculation of the spectrum of cosmological perturbations. Section 5 focuses on the dynamics of the inflaton and of the other relevant scalar fields of the model before the onset of reheating, introduces the conditions for effectiveness of tachyonic resonance and demonstrates that for most of the relevant parameter space of the theory, the tachyonic resonance condition is satisfied. However, the tachyonic instability turns off before most of the energy of the inflaton field is released as matter fluctuations. To obtain an upper bound on the time it takes for the inflaton energy stored in the residual oscillations to further dissipate, we study the later dynamics of the homogeneous field components which proceeds via the less efficient parametric resonance channel (Section 6). Section 7 is devoted to the computation of the secondary curvature fluctuation mode, the mode induced by the entropy fluctuations. ",
        "conclusions": "We have studied reheating and structure formation in the D3/D7 brane inflation model of \\cite{DHHK}, with particular emphasis on the tachyonic instability of super-Hubble scale entropy fluctuations. These entropy perturbations induce a curvature fluctuations which we call ``secondary''. We have found that   under certain conditions these secondary fluctuations are larger than the primordial adiabatic ones which have been considered in the past. They do, however, remain in the regime of applicability of linear perturbation theory.   This would imply that the parameters of the model have to be changed compared to what is usually assumed in order to obtain agreement with the observed amplitude of the large-scale curvature fluctuations.  Along the way, we have given an extensive analytical analysis of the reheating mechanism in the D3/D7 brane inflation model\\footnote{Our analytical analysis is complementary to the numerical work of \\cite{tachyonic}. We need to make certain approximations and neglect some effects, as discussed in the text. On the other hand, numerical work can only handle a limited range of scales. We are interested in large cosmological scales, whereas the basic physical scale of the system is microphysical. It is not clear that numerical work which is sensitive to the microphysical scale can make reliable predictions for questions involving cosmological scales.}. In the low energy field theory limit, the dynamics of this system is a special case of supersymmetric hybrid inflation. We have taken into account the quantum fluctuations in the ``waterfall'' field $\\phi_+$. These play an important role for both the reheating dynamics and for the generation of entropy fluctuations. For COBE-normalized values of the string theory parameters and for reasonable values of the string coupling constant we have found that, including the effects of the above mentioned quantum fluctuations, the initial stages of reheating occur via a tachyonic instability. This instability, however, shuts off quite early, leaving some of the initial inflaton energy in the background fields. To obtain an upper bound on the time it takes to drain most of the residual energy from the background homogeneous fields, we have neglected the interactions of the nonlinear fluctuations produced during the initial phase of tachyonic decay, and focused on the residual homogeneous field dynamics. This later dynamics proceeds by the parametric resonance instability of \\cite{JB1,KLS1,JB2,KLS2}. In fact, except for at the onset of the reheating process, the system is in the narrow-band region of parameters. Nevertheless, the reheating is sufficiently efficient to convert a substantial fraction of the inflaton energy into matter fluctuations.  We note that efficient reheating in the D3/D7 brane inflation model is easier to achieve than in multi-throat inflation models since the matter fields are directly coupled to the inflaton in the form of strings stretching between the two branes whose separation constitutes the inflaton.  In our work, we have neglected the issue of moduli stabilization. This is a very important caveat to our analysis. Moduli stabilization in our model has been considered in the first reference of \\cite{previous} and more recently in \\cite{BBDD}. We plan to study the implications of the corrections to the potential induced by moduli stabilization on the dynamics of reheating in a followup paper.  We have also simplified the dynamics of the background fields. Namely, we have considered in-phase oscillations of the two background fields $S$ and $\\phi_+$, and we have neglected their phases. These phases provide extra low mass entropy modes. It would be interesting to consider their excitation during reheating."
    },
    "0801/0801.4438_arXiv.txt": {
        "abstract": "The proposed design for PILOT is a general-purpose, wide-field ($1\\dg$) 2.4m, f/10 Ritchey-Chr\\'etien telescope, with fast tip-tilt guiding, for use $0.5-25\\mum$. The design allows both wide-field and diffraction-limited use at these wavelengths. The expected overall image quality, including median seeing, is $0.28-0.3''$ FWHM from $0.8-2.4\\mum$. Point source sensitivities are estimated. ",
        "introduction": "The free atmospheric conditions at Dome C are known to be exceptional. The seeing (above $\\sim$ 30m), coherence time, isoplanatic angle, infrared sky emission, water vapour absorption and telescope thermal emission are all better than at the best mid-latitude sites (Lawrence {\\em et al.\\/} \\cite{Lawrence}, Agabi {\\em et al.\\/} \\cite{Agabi}, Walden {\\em et al.\\/} \\cite{Walden}).  PILOT (the Pathfinder for an International Large Optical Telescope) is intended to show that we can fully utilise these conditions for optical/infra-red astronomy. It is also intended to demonstrate that large optical telescopes can be built and operated in Antarctica; to fully characterise the possibilities for adaptive optics; and to perform cutting-edge science in its own right. A design study by the AAO and UNSW is currently underway, with completion mid-2008. As part of that study, we are actively seeking partners in the PILOT project, and input into the design and the scientific requirements  \\begin{figure} \\begin{center} \\includegraphics[width=9.5cm,angle=-90]{saunders1_fig1.ps} \\end{center} \\caption{Optical layout for PILOT, for wide-field optical ($r$-band) use with $1\\dg$ field (above) and wide-field infra-red ($K_{dark}$-band) use with $15'$ field and cold stop (below). The spot diagrams (on right) include the Airy disc and a scale bar of 18$\\mum$/$0.15''$ (above) and 72$\\mum$/$0.62''$ (below). The optical design is diffraction limited above 500nm.} \\end{figure} ",
        "conclusions": ""
    },
    "0801/0801.1501_arXiv.txt": {
        "abstract": "Using {\\it Spitzer} archival data from the SAGE (Surveying the Agents of a Galaxy's Evolution) program, we derive the Cepheid period-luminosity (P-L) relation at $3.6$, $4.5$, $5.8$ and $8.0$ microns for Large Magellanic Cloud (LMC) Cepheids. These P-L relations can be used, for example, in future extragalactic distance scale studies carried out with the {\\it James Webb Space Telescope}. We also derive Cepheid period-color (P-C) relations in these bands and find that the slopes of the P-C relations are relatively flat. We test the nonlinearity of these P-L relations with the $F$ statistical test, and find that the $3.6\\mu \\mathrm{m}$, $4.5\\mu \\mathrm{m}$ and $5.8\\mu \\mathrm{m}$ P-L relations are consistent with linearity. However the $8.0\\mu \\mathrm{m}$ P-L relation presents possible but inconclusive evidence of nonlinearity. ",
        "introduction": "The Cepheid period-luminosity (P-L) relation is an important component in extragalactic distance scale and cosmological studies. The most widely used P-L relations in the literature are obtained from Large Magellanic Cloud (LMC) Cepheids in optical $BVI$ \\citep[e.g., see][]{mad91,tan99,uda99a,san04,kan06} and near infra-red (NIR) $JHK$ \\citep[e.g., see][]{mad91,gie98,gro00,nik04,per04,nge05} bands. In addition, there are also some LMC P-L relations obtained for MACHO $V_{MACHO}R_{MACHO}$ \\citep{nik04,nge05} and EROS $V_{EROS}R_{EROS}$ \\citep{bau99} bands. Recently, \\citet{nge07} applied a semiempirical approach to derive the LMC P-L relation in Sloan $ugriz$ bands: still within optical and/or NIR regimes. Therefore, Cepheid P-L relations are well-developed for wavelengths ranging from optical to NIR.  In contrast, there are currently no P-L relations available for wavelengths longer than the $K$ band in the literature. The main motivation for having a P-L relation at these longer wavelengths is in order to apply it in future extragalactic distance scale studies. The {\\it Near Infrared Camera (NIRCam)} and the {\\it Mid-Infrared Instrument (MIRI)}, that are scheduled to be installed on the {\\it James Webb Space Telescope (JWST)}, will operate in the NIR and mid-infrared: a wavelength range of 0.6-5 microns and 5-27 microns, respectively\\footnote{See the links given in {\\tt http://www.stsci.edu/jwst/instruments/}}. It is possible that extragalactic Cepheids will be discovered and/or (re-)observed (for those galaxies that were observed by the $HST\\ H_0$ Key Project) by the {\\it JWST}. Hence, P-L relations at longer wavelengths are needed in order to use {\\it JWST} to derive a Cepheid distance to these galaxies. In this Paper, we derive the LMC P-L relations at $3.6$, $4.5$, $5.8$ and $8.0$ microns from the {\\it Spitzer} archival data. As far as we are aware, this is the first time such a relation has been derived.  In addition to the P-L relations, we can also derive period-color (P-C) relations and construct color-color plots and the color-magnitude diagrams (CMD) for LMC Cepheids in these {\\it Spitzer} bands. In Section 2 we discuss our data selection. In Section 3 we present our analysis and results for the P-L relations. Section 4 we show the P-C relations, the CMD and the color-color plot for the Cepheids in our sample. Our conclusion is given in Section 5. Extinction is ignored in this Paper because it is expected to be negligible in the {\\it Spitzer's IRAC} bands (hereafter {\\it IRAC} band). ",
        "conclusions": "In this Paper, we derive P-L relations for LMC Cepheids in {\\it IRAC} $3.6$, $4.5$, $5.8$ and $8.0$ microns bands. These P-L relations can be potentially applied to future extragalactic distance scale studies with, for example, the {\\it JWST}. The data are taken from the {\\it Spitzer's} archival database from the SAGE program. After properly removing the outliers, the fitted P-L relations are presented in Table \\ref{tab1}. We have tested the P-L relations with various period cuts and found that our results are insensitive to period cuts up to $\\log P_{cut} \\sim 0.65$. We also argue that the random phase corrections may not be important for {\\it IRAC} band P-L relations. When comparing P-L relations from $B$ to $8.0\\mu \\mathrm{m}$ bands, the slope of the P-L relation appears to be the steepest around $K$ band to $3.6\\mu \\mathrm{m}$ band, while the dispersion of the P-L relation reaches a minimum between those bands. The shallower slopes and larger P-L dispersions in the $5.8\\mu \\mathrm{m}$ and $8.0\\mu \\mathrm{m}$ band are in contrast to the theoretical expectation. This could be due to the smaller number of Cepheids and larger measurement errors toward the faint end in these two bands. We also test the nonlinearity of the P-L relations in the {\\it IRAC} band using the $F$ statistical test. As expected, the $F$-test results show that the P-L relations are linear in $3.6$, $4.5$ and $5.8$ microns band, but the $8.0\\mu \\mathrm{m}$ P-L relation is found to be nonlinear. However, the nature of the nonlinear $8.0\\mu \\mathrm{m}$ P-L relation is still inconclusive. For the P-C relations, it was found that the slopes of the P-C relation are relatively flat in the {\\it IRAC} bands. Finally, the LMC Cepheids show a well-defined instability strip in the CMD and clustered in a small region in the color-color plot. This can potentially be used to identify Cepheids observed in the {\\it IRAC} bands. Even though there may be some associated problems for the $8.0\\mu \\mathrm{m}$ P-L relation, the $3.6\\mu \\mathrm{m}$, $4.5\\mu \\mathrm{m}$ and perhaps the $5.8\\mu \\mathrm{m}$ P-L relations can still be used in future distance scale studies."
    },
    "0801/0801.3732_arXiv.txt": {
        "abstract": "{} {We report on the production of a large area, shallow, sky survey, from \\xmm slews. The great collecting area of the mirrors coupled with the high quantum efficiency of the EPIC detectors have made \\xmm the most sensitive X-ray observatory flown to date. We use data taken with the EPIC-pn camera during slewing manoeuvres to perform an X-ray survey of the sky.  } { Data from 218 slews have been subdivided into small images and source searched. This has been done in three distinct energy bands; a soft (0.2-2 keV) band, a hard (2-12 keV) band and a total \\xmm band (0.2-12 keV). Detected sources, have been quality controlled to remove artifacts and a catalogue has been drawn from the remaining sources.} {A 'full' catalogue, containing 4710 detections and a 'clean' catalogue containing 2692 sources have been produced, from 14\\% of the sky. In the hard X-ray band (2-12 keV) 257 sources are detected in the clean catalogue to a flux limit of $4\\times10^{-12}$ ergs s$^{-1}$ cm$^{-2}$. The flux limit for the soft (0.2-2 keV) band is $6\\times10^{-13}$ ergs s$^{-1}$ cm$^{-2}$ and for the total (0.2-12 keV) band is $1.2\\times10^{-12}$ ergs s$^{-1}$ cm$^{-2}$. The source positions are shown to have an uncertainty of 8\\arcsec (1~$\\sigma$ confidence).} {} ",
        "introduction": "There is a strong tradition in X-ray astronomy of using data taken during slewing manoeuvres to perform shallow surveys of the sky. The Einstein (\\cite{elvis}), Exosat (\\cite{reynolds}) and RXTE (\\cite{revnivtsev}) slew surveys all provide a useful complement to dedicated all-sky surveys such as ROSAT (\\cite{voges}) and HEAO-1 (\\cite{piccinotti})  and the smaller area, medium sensitivity ASCA survey (\\cite{ueda}) and pencil-beam \\xmm (\\cite{hasinger}) and Chandra (\\cite{brandt}) deep looks. It was long recognised that XMM-Newton (\\cite{jansen}) with its great collecting area, efficient CCDs, wide energy band and tight point-spread function (PSF) has the potential to make an important contribution to our knowledge of the local universe from its slew data. Early estimates (\\cite{Lumb98}; \\cite{jonesandlumb}), based on an expected slewing speed of 30 degrees per hour, and a slightly lower background level than that actually encountered in orbit, predicted a 0.5--2 keV flux limit of $2\\times10^{-13}$ erg cm$^{-2}$s$^{-1}$. While the chosen in-orbit slew speed of 90 degrees per hour reduces the sensitivity, initial assessments of the data showed that the quality of the data is good and that many sources are detected (\\cite{freyberg}). A review of properties shows that the \\xmm slew survey compares favourably with other large area surveys in terms of depth and positional accuracy (Table 1).  During slews all the three imaging EPIC cameras take data in the observing mode set in the previous pointed observation and with the Medium filter in place. The slew speed of 90 degrees per hour combined with the slow readout time of the MOS detectors (2.6s; \\cite{turner}) means that sources appear as long streaks in the MOS cameras but are well formed in the fast observing modes of the pn camera (\\cite{struder}). For this reason, only the EPIC-pn data have been analysed.  In this paper we present a catalogue drawn from slews taken between revolutions 314 and 978 covering a sky region of 6240 square degrees. The main properties of the slew survey discussed in this paper are given in Table 2.  \\begin{table} \\caption{Properties of large area X-ray surveys} \\label{table:SurveySumm}      % \\begin{center} \\begin{tabular}{l c c c c}        % \\hline\\hline                 % Satellite & Energy range & Coverage$^{a}$ & Flux lim. & Position \\\\    % & (keV) &    \\% of sky   &  $^{b}$      &  error \\\\    % \\hline                        % RASS & 0.2-2.4 & 92 & 0.03 & 12\\arcsec \\\\      % Einstein slew & 0.2-3.5 & 50 & 0.3    & 1.2\\arcmin \\\\ {\\bf XMM slew (soft)} & {\\bf 0.2-2} & {\\bf 14} & {\\bf 0.06}  & {\\bf 8\\arcsec}  \\\\ \\hline                        % EXOSAT slew & 1-8 & 98 & 3  & 20\\arcmin \\\\ HEAO-1/A2 & 2-10 & 100 & 3 & 60\\arcmin \\\\      % RXTE slew & 3-20 & 95 & 1.8  & 60\\arcmin \\\\ {\\bf XMM slew (hard)} & {\\bf 2-12} & {\\bf 14} & {\\bf 0.4}  & {\\bf 8\\arcsec} \\\\ \\hline                                   % \\end{tabular} \\\\ \\end{center} $^{a}$ The \\xmm slew sky coverage has been computed by adding the area contained in all of the images used in source searching with an exposure time greater than 1 second. \\\\ $^{b}$ Flux limit, units of $10^{-11}$ ergs $s^{-1}$ cm$^{-2}$ \\\\ \\end{table}  \\begin{table} \\caption{Properties of the \\xmm slew survey} \\label{table:slewprops}      % \\begin{center} \\begin{tabular}{l c c c}        % \\hline\\hline                 % Property & \\multicolumn{3}{c}{Range (keV)}\\\\    % & 0.2--2 & 2--12 & 0.2--12  \\\\    % \\hline                        % Observing time (s) & $5.4\\times10^{5}$ & $5.4\\times10^{5}$ & $5.4\\times10^{5}$ \\\\ Mean exposure time$^{a}$ & 6.2 s & 6.0 s & 6.1 s\\\\ Total number of photons & $1.6\\times10^{6}$ & $2.2\\times10^{6}$ & $3.8\\times10^{6}$ \\\\ Mean background$^{b}$ & 0.09 & 0.14 & 0.23\\\\ Median source count rate$^{c}$ & 0.68 & 0.90 & 0.81\\\\ Limiting source flux$^{d}$ & $6.0\\times10^{-13}$ & $4.0\\times10^{-12}$ & $1.2\\times10^{-12}$\\\\ Num. sources (full cat)$^{e}$ & 2606 & 692 & 3863\\\\ Num. sources (clean cat)$^{e}$ & 1874 & 257 & 2364\\\\  \\hline                                   % \\end{tabular} \\\\ \\end{center} $^{a}$ The mean exposure time, after correcting for the energy-dependent vignetting.\\\\ $^{b}$ cnts/arcmin$^{2}$. \\\\ $^{c}$ cnts/second. \\\\ $^{d}$ ergs/s/cm$^{2}$. Based on a detection of 4 photons from a source passing near the detector centre (exp. time = 10s), with a power-law spectrum of slope 1.7 and galactic absorption of 3$\\times10^{20}$ atoms cm$^{-2}$. \\\\ $^{e}$ The number of sources flagged as good in the full and clean catalogues (see text). \\\\ \\end{table}  Slews have been source searched down to a likelihood threshold of 8, which after manual rejection of false detections gives 4710 candidate sources. Using simulations we have been able to identify a subset of high significance sources, with a likelihood threshold dependent upon the background conditions, that gives 2692 sources with a spurious fraction of $\\sim4\\%$ (the \"clean\" catalogue). Of these, 2621 are from unique sources. In the hard (2--12 keV) band the clean catalogue contains 257 sources (253 unique) of which $\\sim9\\%$ are expected to be due to statistical fluctuations.  The slew catalogue and accompanying images and exposure maps have been made available through the XMM Science Archive (XSA) as a queryable database and as FITS files\\footnote{The catalogue was initially released on May 3 2006. In this paper we discuss the updated version released in October 2006.}. A summary of scientific highlights from the slew survey has been published in \\cite{Read06}. ",
        "conclusions": "The \\xmm slew data represent the deepest near all-sky hard band X-ray survey made to date, while the soft band survey is comparable with the RASS. The source density, from the clean catalogue, is $\\sim0.45$ per square degree and $\\sim70\\%$ have plausible identifications. The \\xmm slew survey catalogue will continue to grow as the mission continues and it is expected that a sky coverage in excess of 50\\% will eventually be achieved. With the excellent attitude reconstruction this will leave a powerful legacy for future variability studies in the soft and hard X-ray bands."
    },
    "0801/0801.0605.txt": {
        "abstract": "% context heading (optional) % {} leave it empty if necessary {Over the last decade, several groups of young (mainly low-mass) stars have been discovered in the solar neighbourhood (closer than $\\sim$100 pc), thanks to cross-correlation between X-ray, optical spectroscopy and kinematic data. These young local associations -- including an important fraction whose members are Hipparcos stars -- offer insights into the star formation process in low-density environments, shed light on the substellar domain, and could have played an important role in the recent history of the local interstellar medium.} % aims heading (mandatory) {To study the kinematic evolution of young local associations and their relation to other young stellar groups and structures in the local interstellar medium, thus casting new light on recent star formation processes in the solar neighbourhood.} % methods heading (mandatory) {We compiled the data published in the literature for young local associations. Using a realistic Galactic potential we integrated the orbits for these associations and the Sco-Cen complex back in time.} % results heading (mandatory) {Combining these data with the spatial structure of the Local Bubble and the spiral structure of the Galaxy, we propose a recent history of star formation in the solar neighbourhood. We suggest that both the Sco-Cen complex and young local associations originated as a result of the impact of the inner spiral arm shock wave against a giant molecular cloud. The core of the giant molecular cloud formed the Sco-Cen complex, and some small cloudlets in a halo around the giant molecular cloud formed young local associations several million years later. We also propose a supernova in young local associations a few million years ago as the most likely candidate to have reheated the Local Bubble to its present temperature.} % conclusions heading (optional), leave it empty if necessary {} ",
        "introduction": "The first moving groups of stars were discovered early on in the history of Galactic dynamics. As early as 1869, R.A. Proctor published a work in which he identified a group of stars around the Hyades open cluster that was moving together through the Galaxy. He also found another 5 comoving stars in the Ursa Major constellation. In the 1960s, O.J. Eggen suggested the existence of a {\\it Local Association} (see for example Eggen \\cite{Eggen61}, \\cite{Eggen65a}, \\cite{Eggen65b}) formed by a group of young stars with approximately the same spatial velocity (also referred as the {\\it Pleiades moving group}). Eggen's Local Association included the nearest bright B-type stars, stars in the Pleiades, the $\\alpha$ Perseus and IC2602 clusters, and the stars belonging to the Sco-Cen complex. However, it was difficult to defend a unitary view of the group due to the wide range of ages ($\\la$100 Myr) and widespread spatial distribution ($\\sim$300 pc) of its constituent stars.  At the same time as Eggen was studying his Local Association, the first X-ray detectors were installed in rockets and launched into space.  This led to the discovery of the diffuse soft X-ray background (SXRB). The anticorrelation observed between the SXRB and the HI column density was rapidly interpreted as evidence of a local cavity in the interstellar medium (ISM) filled by an X-ray emitting plasma, which became known as the {\\it Local Bubble} (LB). The launch of the ROSAT satellite in 1990 allowed the LB to be studied in more detail. The opening of the X-ray window in the sky and observations of stars belonging to clusters with known ages also led to the realisation that X-ray emission persists in young stars for a period of the order of 100 Myr. An important fraction of these young, X-ray emitting stars are T Tauri stars. These very young ($\\la$10 Myr) stars also exhibit excess IR emission (due to the presence of nearby heated dust particles forming disks or envelopes) as well as UV line and continuum emission produced by the accretion of the surrounding gas and dust. When the stars reach the age of $\\sim$10 Myr, both IR emissions and optical activity decline considerably. X-ray activity then becomes the basis for ascertaining youthfulness.  The ROSAT All-Sky Survey (RASS) detected more than 100,000 X-ray sources at keV energies (Voges et al. \\cite{Voges00}). Using a kinematic approach and searching for groups of comoving stars allows us to determine which of them are T Tauri stars. This greatly reduces the number of candidate stars that must be observed spectroscopically to confirm their youthfulness.  In this way, several young stellar associations have been discovered within 100 pc of Earth during the last decade (see Jayawardhana \\cite{Jayawardhana00}, and Zuckerman \\& Song \\cite{Zuckerman04b} for a recent review). The stars belonging to these young local associations (YLA) have ages ranging from a few million to several tens of millions of years. Due to their proximity to us, these stars are spread over a large area of the sky (up to several hundred square degrees) making it difficult to identify the associations as single entities. Furthermore, they are far away from molecular clouds or star forming regions (SFR). This is why they remained unnoticed until recently.  These young stars offer insights into the star formation process in low-density environments, which is different from the dominant mechanism observed in denser SFR. The discovery of these YLA has also contributed to substellar astrophysics, since dozens of dwarfs have been identified in them. For instance, it has been confirmed that isolated brown dwarfs can form in these low-density environments (see for example Webb et al. \\cite{Webb99}; Lowrance et al. \\cite{Lowrance99}; Chauvin et al. \\cite{Chauvin03}). YLA also provide important clues regarding recent star formation in the solar neighbourhood (since they contain the youngest stars near the Sun) and its effect on the local ISM.  In Sect. \\ref{sect.LB+YLA} of this paper we review current understanding of the LB and the Sco-Cen complex and present published data for the nearly 300 stars belonging to YLA. In Sect. \\ref{sect.orbits} we present our method for orbit integration, which makes use of a Galactic potential that includes the general axisymmetric potential and the perturbations due to the Galaxy's spiral structure and the central bar. This leads us to study, in Sect. \\ref{sect.origin}, the origin and evolution of the local structures and, in Sect. \\ref{sect.sce}, to propose a scenario for recent star formation in the solar neighbourhood. This scenario explains the origin of the Sco-Cen complex and YLA through the collision of a parent giant molecular cloud (GMC) with the Sagittarius-Carina spiral arm. Finally, the conclusions of our work are summarised in Sect. \\ref{sect.conc}.   % %__________________________________________________________________ ",
        "conclusions": "\\label{sect.conc}  This paper studies the kinematic evolution of the Sco-Cen complex and the so-called young local associations (YLA). It makes use of most of the astrophysical data published in the literature for all the known members of YLA (more than 200 stellar systems). This information appears in Appendix \\ref{sec.app} and can also be accessed via a webpage\\footnote{http://www.am.ub.es/$\\sim$dfernand/YLA/}.  The study of the orbits integrated back in time for all these associations allows us to propose a scenario for recent star formation in the solar neighbourhood (Sect. \\ref{sect.sce}) and a possible link between these associations and the origin and/or evolution of the LB (Sect. \\ref{sect.YLA+LB}). In our scenario, the oldest Sco-Cen associations were brought about by the impact of the inner spiral arm against a giant molecular cloud. YLA were born later (and outside Sco-Cen) due to the shock fronts of the most massive supernovae belonging to the LCC or UCL associations. Our results suggest that a YLA is the most likely place to have harboured the supernova that reheated the LB a few million years ago.  As seen in this paper, the recent discovery of a set of YLA, together with the use of appropriate tools (mainly, the integration back in time of orbits), sheds light on several apparently unrelated topics in astrophysics. These topics range from star formation mechanisms for low-mass stars distant from star forming regions, to the recent history of the LB, and include the origin of the Sco-Cen stellar complex and its possible independence from the Gould Belt. The possible discovery in the coming years of new members of these associations, or even of new associations, will be very useful in confirming the results obtained in this work."
    },
    "0801/0801.2391_arXiv.txt": {
        "abstract": "We observed the first known very high energy (VHE) $\\gamma$-ray emitting unidentified source, TeV J2032+4130, for 94 hours with the MAGIC telescope. The source was detected with a significance of 5.6$\\sigma$. The flux, position, and angular extension are compatible with the previous ones measured by the HEGRA telescope system five years ago. The integral flux amounts to (4.5$\\pm$0.3$_{\\rm stat}\\pm$0.35$_{\\rm sys}$)$\\times$10$^{-13}$ ph~cm$^{-2}$~s$^{-1}$ above 1~TeV. The source energy spectrum, obtained with the lowest energy threshold to date, is compatible with a single power law with a hard photon index of $\\Gamma$=-2.0$\\pm$0.3$_{\\rm stat}\\pm$0.2$_{\\rm sys}$. ",
        "introduction": "The TeV source J2032+4130 (Aharonian et al. 2002) was the first unidentified very high energy (VHE) $\\gamma$-ray source, and also the first discovered extended TeV source, likely to be Galactic.  Intensive observational campaigns at different wavelengths have been carried out on TeV J2032+4130. Butt et al. (2003) presented an analysis of the CO, HI, and infrared emissions, together with first observations by {\\it Chandra} (5 ksec) and a reanalysis of VLA data. These observations showed that the TeV source region is positionally coincident with an outlying group of stars (from the Cygnus OB2 core), although they failed to identify a counterpart. Mukherjee et al. (2003) analyzed the same {\\it Chandra} data and provided optical follow-up observations of several of the brightest X-ray sources, confirming that most were either O stars or foreground late-type stars. A deeper {\\it Chandra} observation (50 ksec, Butt et al. 2006), found hundreds of star-like sources and yet no diffuse X-ray counterpart emission.  A deep ($\\sim$ 50~ksec) \\textit{XMM-Newton} exposure has also been obtained (Horns et al. 2007). After the subtraction of the contribution of known sources from the data, an extended X-ray emission region with a FWHM size of $\\sim$~12 arcmin was reported. The centroid of the emission is co-located with the position of TeV J2032+4130 and was proposed as the counterpart of the TeV source. The question whether the result reported by Horns et al. can be interpreted as a truly diffuse background, or it could be a result of unresolved X-ray sources, remains disputable.   Paredes et al. (2007) and Mart\\'i et al. (2007) have provided deep radio observations covering the TeV J2032+4130 vicinity using the Giant Metrewave Radio Telescope and discovered a population of radio sources, some in coincidence with X-ray detections by Butt et al. (2006) and with optical/IR counterparts. At least three of these sources are non-thermal, and one has a hard X-ray energy spectrum. They found extended non-thermal diffuse emission in the radio band apparently connecting with one or two radio sources. It is yet to be determined if one or more of these sources is similar to some of the known $\\gamma$-ray binaries (e.g., Aharonian et al. 2006, Albert et al. 2006a).  Several theoretical explanations for the TeV emission from J2032+4130 have been given. Among them, those related with extragalactic counterparts, e.g., a radiogalaxy (Butt et al. 2006) or a proton blazar (Mukherjee et al. 2003), face the difficulty of explaining the extended appearance of the source. Gamma-ray production in hypothetical jet termination lobes of Cyg X-3 was explored (Aharonian et al. 2002), but the putative northern lobe of Cyg X-3 (now considered a mere thermal HII region, Mart\\'i et al. 2006) is far from the location of the TeV source. A yet unknown pulsar wind nebula (PWN) was proposed by Bednarek (2003), although no clear PWN signal was observed. A distant microquasar was proposed by Paredes et al. (2007), perhaps related with one of the X-ray/radio sources they discovered. If such an association is accepted, the extension of the source could be explained by the diffusion of accelerated particles into a hypothetical nearby molecular enhancement (see Bosch-Ramon et al. 2005).  Torres et al. (2004) and Domingo-Santamar\\'ia \\& Torres (2006) studied the relationship between the TeV emission and the known massive stars in the area, through the interaction of relativistic protons with wind ions. The distribution of stars in the neighborhood favors this interpretation (Butt et al. 2006). An explanation involving the excitation of giant dipole resonances of relativistic heavy nuclei in radiation dominated environments has also been suggested (Anchordoqui et al. 2007). ",
        "conclusions": "MAGIC observations confirm the location of TeV J2032+4130 found by HEGRA. The MAGIC observation shows an extended source with a significance of 5.6$\\sigma$. We find a steady flux with no significant variability within the three year span of the observations (with the flux being at a similar level of the HEGRA data of the period 2002-2005). We also present the source energy spectrum obtained with the lowest energy threshold to date, which, within errors, is compatible with a single power law."
    },
    "0801/0801.3324.txt": {
        "abstract": "A numerical approach is considered for spherically symmetric spacetimes that generalize Lema\\^\\i tre--Tolman--Bondi dust solutions to nonzero pressure (``LTB spacetimes''). We introduce quasi--local (QL) variables that are covariant LTB objects satisfying evolution equations of Friedman--Lema\\^\\i tre--Robertson--Walker (FLRW) cosmologies. We prove rigorously that relative deviations of the local covariant scalars from the QL scalars are non--linear, gauge invariant and covariant perturbations on a FLRW formal ``background'' given by the QL scalars. The dynamics of LTB spacetimes is completely determined by the QL scalars and these exact perturbations. Since LTB spacetimes are compatible with a wide variety of ``equations of state'', either single fluids or mixtures, a large number of known solutions with dark matter and dark energy sources in a FLRW framework (or with linear perturbations) can be readily examined under idealized but non--trivial inhomogeneous conditions.  Coordinate choices and initial conditions are derived for a numerical treatment of the perturbation equations, allowing us to study non--linear effects in a variety of phenomena, such as gravitational collapse, non--local effects, void formation, dark matter and dark energy couplings and particle creation.  In particular, the embedding of inhomogeneous regions can be performed by a smooth matching with a suitable FLRW solution, thus generalizing the Newtonian ``top hat'' models that are widely used in astrophysical literature. As examples of the application of the formalism,  we examine numerically the formation of a black hole in an expanding Chaplygin gas FLRW universe, as well as the evolution of density clumps and voids in an interactive mixture of cold dark matter and dark energy. ",
        "introduction": "Recent observations suggest that present cosmic dynamics is dominated by elusive sources called ``dark matter'' and ``dark energy'', the former clustered in galactic halos and the latter possibly associated with a large scale repulsive (yet unknown) interaction. A large body of phenomenological and theoretical models have been proposed to describe these sources (see \\cite{review} for a comprehensive review). The dominant model of dark matter is a collision--less gas of supersymmetric particles: ``cold'' dark matter~\\cite{review}, whereas dark energy has been described by a  cosmological constant \\cite{Lambda}, as well as by quintessence scalar fields, Phantom fields, tachyons, branes (see \\cite{review}). A unified perspective is given by the Chaplygin gas~\\cite{review,chaplygin1,chaplygin2}, a single source that exhibits the expected dynamical behavior of dark matter and dark energy. There is also a large number of empiric and phenomenological descriptions of dark matter interacting with dark energy \\cite{intmix1,intmix2} (see \\cite{intmix_new} for a recent appraisal), as well as empiric dark energy ``equations of state'' that fit observations~\\cite{EOS1,EOS2}.  While dark matter at the galactic scale is considered inhomogeneous, due to its association with structure formation, most research work on dark energy sources has been conducted in homogeneous Friedman--Lema\\^\\i tre--Robertson--Walker (FLRW) spacetimes and/or linear perturbations, since dark energy is assumed to have a significant effect only at a larger cosmic scale in which the universe appears to be homogeneous.  However, the possibility of anisotropic or inhomogeneous dark energy (or combined dark matter and energy sources) is beginning to be discussed in the literature \\cite{chapinhom,chapinhom2,chapstar,intmix3,Mota_anis}.  Given our ignorance on the fundamental properties of these sources, there is no theoretically binding reason to assume, {\\it a priori}, that no valuable new information would result from studying their interaction under inhomogeneous conditions, at least in scales associated with structure formation and gravitational clustering.  Alternative proposals to the dark energy paradigm also consider a full relativistic treatment of cosmological inhomogeneities~\\cite{Inh_rew1,Inh_rew2}. Numerous articles describe the possibility that cosmic acceleration (or at least part of it) could result from the presence of inhomogeneites in photon trajectories in observations from high red--shift objects~\\cite{InhObs1, InhObs2}. Following a more theoretical perspective, various averaging formalisms~\\cite{ave1,ave2} have examined the occurrence of an effective acceleration from the so--called ``back--reaction'', associated with non--local effects that emerge as inhomogeneous sources are averaged (see \\cite{theo1,theo2} for further discussion and \\cite{ltbave,condsBR} for scalar averages in the spherical dust case). Also, an interesting  connection  between cosmic acceleration, back--reaction and non--trivial gradients of quasi--local energy has been proposed~\\cite{wiltshire}, which could clarify important theoretical  issues (the connection between the material presented in this paper and back--reaction issues is discussed in \\cite{condsBR,QLvars}). Possible observational consequences of these theoretical alternatives to dark energy are certainly worth serious consideration~\\cite{Obs_BR} (see also \\cite{InhObs1,InhObs2} ).  Even if dark energy prevails over these alternative proposals, there is no harm done (and perhaps valuable new information) in probing its behavior under inhomogeneous conditions.   Ideally, the study of inhomogeneous sources should be conducted on  spacetimes not restricted by simplifying symmetries, but this would require fully 3--dimensional codes of high complexity. As a compromise between linear perturbations and this type of ``realistic'' generality, we present in this article a class of inhomogeneous spherically symmetric models that can be well handled with relatively simple numerical methods. Although these models are idealized, they are non--trivial, exhibit non--linear behavior, and are general enough for testing a wide variety of physical assumptions on modern cosmological sources.  By assuming a spherically symmetric Lema\\^\\i tre--Tolman--Bondi metric, we derive in section \\ref{genLTB} a class of spacetimes (``LTB spacetimes'') that generalize to nonzero pressure (and nonzero pressure gradients) the well know solutions for inhomogeneous dust associated with this metric \\cite{kras}. The most general source for this metric (in the comoving frame) is a fluid with anisotropic stresses, which will end up being interpreted as fluctuations of the isotropic pressure. However, regardless of the interpretation, we remark that the existence of pressure anisotropy is far from a drawback or shortcoming, not only because it can be associated to a number of well motivated physical effects~\\cite{anisP}, but because an inhomogenous source with only isotropic pressure is a far more idealized situation than one with anisotropic pressure.  In section \\ref{1plus3} we show that LTB spacetimes can be fully characterized by a set of local covariant scalars, hence their dynamics becomes completely determined by solving the evolution equations for these scalars in the covariant fluid flow (or ``1+3'') representation~\\cite{ellisbruni89,BDE,LRS,1plus3}. By looking at the relation between energy density and the Misner--Sharp quasi--local mass--energy invariant~\\cite{MSQLM,kodama,szab,hayward1,hayward2}, we introduce in section \\ref{QLdefs} ``quasi--local'' (QL)  scalar functions, which are LTB objects that satisfy FLRW dynamics. Expressing all local covariant scalars as perturbations of the QL scalars, we rewrite in section \\ref{eveqs} the 1+3 fluid flow equations of section \\ref{1plus3} as evolution equations for the QL scalars and these perturbations, which on the basis of known criteria that define a perturbation formalism~\\cite{ellisbruni89, BDE,LRS,1plus3,bardeen}, are shown in section \\ref{perturb} to be covariant, gauge--invariant, non--linear perturbations on a FLRW formal ``background'' given by the QL variables. These evolution equations are equivalent to ODE's in which radial dependency enters as a parameter (see Appendix C of \\cite{suss08})).  As shown in section \\ref{EOS}, once we choose an ``equation of state'' (EOS) between the QL energy density and pressure (which are ``background'' variables), the QL evolution equations become determined. We point out that, by choosing such an EOS, correlation terms appear in the relations between local scalars, in a similar way as virial corrections to the ideal gas EOS appear in Newtonian self--gravitating systems~\\cite{saslaw1,saslaw2}. These terms can be associated with the long range nature of gravity~\\cite{Padma_GTD}. Since the fundamental physics of dark matter and dark energy sources is still unknown, we cannot rule out the possibility that this type of non--local effects could play an important dyamical role when these sources are studied under non--trivial and non--linear inhomogeneity.  In section \\ref{mixtures} we apply the perturbation formalism to a fluid mixture, which can be interactive or with each component separately conserved (decoupled). This mixture description can be readily used to generalize to inhomogeneous conditions similar mixture models derived in a FRLW context in the many references quoted in \\cite{intmix1,intmix2,intmix_new}. We show in this section that scalar fields can only be compatible with LTB spacetimes in mixture sources, playing the role of the homogeneous dark energy component (coupled with inhomogeneous dark matter~\\cite{intmix3})  Coordinate choices appropriate for the numerical treatment for the evolution equations of sections \\ref{eveqs} and \\ref{mixtures} are discussed in \\ref{FLRWlike}, while in section \\ref{FLRWsubs} we show that perturbation functions are well defined as long as shell crossing singularities are absent. A possible (but not compulsory) way in which LTB spacetimes can be embedded into a FLRW background is by smoothly matching an inhomogeneous LTB section with a section of a suitable FLRW spacetime. We provide in section \\ref{tophats} the conditions for such a matching, leading to smooth and fully relativistic generalizations of the Newtonian ``top hat'' models that are widely used in the astrophysical literature~\\cite{tophats}.  We examine in section \\ref{applications} two examples of the application of the formalism presented in the previous sections. In the first example, we solve numerically the evolution equations for a smooth and fully relativistic Chaplygin gas ``top hat'' model, in which a black hole is formed in an inhomogeneous comoving section of the LTB  Chaplygin gas, smoothly embedded in an expanding Chaplygin gas FLRW universe. In the second example, we examine a mixture of interactive dark matter and dark energy of the type examined in \\cite{intmix1,intmix2,intmix_new}, but give the interaction a phenomenological interpretation in terms of particle creation, with dark matter particles decaying into dark energy~\\cite{intmix_new,lima}. We show how the radial profiles of dark matter density evolve from initial clumps into deep voids, thus providing a toy model for a void formation scenario in the context of this type of interactions. In the appendices we provide detailed information on how initial conditions and coordinate choices can be set up for the numerical integration of the evolution equations. ",
        "conclusions": "We have presented a formalism, based on quasi--local (QL) variables, describing a large class of spherically symmetric models (LTB spacetimes) as non--linear, gauge invariant and covariant perturbations of a FLRW formal background. The integration of the resulting evolution equations require numerical work, but this task can be handled with relatively simple numerical techniques. A summary of how this formalism was motivated, introduced, developed and applied is furnished in the Introduction (section II).  Since LTB spacetimes are compatible with a wide variety of ``equations of state'' (EOS) and physical assumptions, these models and the formalism developed for them are useful theoretical tools to examine a large amount of cosmological sources of interest that have been studied only in a FLRW context (or linear perturbations)~\\cite{review,Lambda,chaplygin1,chaplygin2,intmix1,intmix2,intmix_new,EOS1,EOS2}. While the inhomogeneity of LTB spacetimes is certainly idealized, the resulting models are non--trivial and still exhibit non--linear behavior, and can be useful to examine important phenomena (gravitational collapse, void formation) that cannot be studied with FLRW models or their linear perturbations.  As we argued in section \\ref{EOS}, the EOS that determines the evolution equations is given between the QL energy density and pressure, so that local density and pressure and pressure are related by means of correlation terms that could contain non--local information that could be important for self--gravitating sources. Since the fundamental nature of dark matter and dark energy is not known, we cannot rule the possibility of non--local effects playing an important role when we examine these sources under inhomogeneous conditions, at least in a scale smaller than 300 Mpc. We believe that LTB models can be useful to examine these and other non--linear effects, not only for dark matter or dark energy, but also in theoretical proposals that seek to explain cosmic acceleration without dark energy~\\cite{Inh_rew1,Inh_rew2,InhObs1,InhObs2,ave1,ave2,ave3,theo1,theo2,ltbave,condsBR,wiltshire}. As shown elsewhere~\\cite{condsBR,QLvars}, the formalism presented here can be applied to understand the issue of ``back--reaction'' that appears in various averaging formalisms~\\cite{ave1,ave2,ave3} aiming to explore this type of alternative explanations.  An important result of this article is the smooth and fully relativistic generalization of the Newtonian ``top hat'' models (section \\ref{tophats}), which are useful toy models for structure formation and are widely used in Astrophysical literature. One of the numerical examples of application of the formalism that we provided (section \\ref{chaplygin}) is a Chaplygin gas ``top hat'' model that describes the formation of a local black hole smoothly embedded in an expanding Chaplygin gas FLRW universe.  The second numerical application example is that of an interactive mixture of cold dark matter (CDM) and dark energy. By interpreting the interaction term in terms of particle creation, we obtained a model in which initial clump (or overdensity) radial profiles of CDM evolve into void profiles as CDM decays into dark energy. While this model is not ``realistic'', it serves to illustrate how LTB spacetimes can illustrate important non--linear features of dark energy/matter sources that cannot be studied within a homogeneous context (nor with linear perturbations). As shown in \\cite{intmix2,intmix_new}, mixtures of this type can be studied by means of a dynamical systems approach. This type of approach is compatible with the evolution equations of LTB spacetimes (autonomous equations), and so a generalization can be readily made of these dynamical system studies in a FLRW context. In fact, a dynamical systems study of LTB dust solutions has been achieved using the same QL variables~\\cite{suss08}. The extension of this work for LTB spacetimes with nonzero pressure is being currently undertaken."
    },
    "0801/0801.2450_arXiv.txt": {
        "abstract": "{The X-ray reflection features of irradiated accretion disks around black holes enable us to probe the effects of strong gravity. We investigate to which precision the reflection signs, i.e. the iron K-line and the Comptonized hump, can be observed with Simbol-X for nearby Seyfert galaxies. The simulations presented include accurate computations of the local reprocessed spectra and modifications due to general relativistic effects in the vicinity of the black hole. We discuss the impact of  global black hole parameters and of the irradiation pattern of the disk on the resulting spectra as they will be detected by the Simbol-X mission. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.0039_arXiv.txt": {
        "abstract": "Anomalous X-ray pulsars (AXPs) are thought to be magnetars, which are neutron stars with ultra strong magnetic field of $10^{14}$--$10^{15}$ G. Their energy spectra below $\\sim$10 keV are modeled well by two components consisting of a blackbody (BB) ($\\sim$0.4 keV) and rather steep power-law (POW) function (photon index $\\sim$2-4). Kuiper et al.(2004) discovered hard X-ray component above $\\sim$10 keV from some AXPs. Here, we present the Suzaku observation of the AXP 1E 1841-045 at the center of supernova remnant Kes 73. By this observation, we could analyze the spectrum from 0.4 to 50 keV at the same time. Then, we could test whether the spectral model above was valid or not in this wide energy range. We found that there were residual in the spectral fits when fit by the model of BB + POW. Fits were improved by adding another BB or POW component. But the meaning of each component became ambiguous in the phase-resolved spectroscopy. Alternatively we found that NPEX model fit well for both phase-averaged spectrum and phase-resolved spectra. In this case, the photon indices were constant during all phase, and spectral variation seemed to be very clear. This fact suggests somewhat fundamental meaning for the emission from magnetars. ",
        "introduction": "Anomalous X-ray pulsars (AXPs) are thought to be magnetars, which are strongly magnetized ($\\sim 10^{14}$ -- $10^{15}$ G) neutron stars with emissions powered by the dissipation of the magnetic energy (see \\cite{Woods Thompson 2006} for a review). The spectra of the AXP had been modelled well by the compound of a blackbody and a power-law function below $\\sim 10$ keV region. Besides, two temperature blackbody spectrum was also fit well for some AXPs \\cite{Morii et al 2003, Naik et al 2007}. On the other hand, separate hard X-ray emission was discovered for some AXPs above $\\sim 10$ keV \\cite{Kuiper Hermsen Mendez 2004, Kuiper et al 2006}.  AXP 1E 1841-045 is located on the center of the supernova remnant (SNR) Kes 73 (G 27.4+0.0) with diameter of about 4$^\\prime$. The kinematic distance toward the SNR was estimated to be between 6 and 7.5 kpc \\cite{Sanbonmatsu Helfand 1992}. The pulse period of 1E 1841-045 was 11.8 s \\cite{Vasisht Gotthelf 1997}. Morii et al. (2003) \\cite{Morii et al 2003} reported that the spectrum was fitted well with the model consisting of the blackbody ($kT = 0.44 \\pm 0.002$ keV) and the power-law function with the hardest photon index among AXP ($\\Gamma = 2.0 \\pm 0.3$), using Chandra data (0.6 -- 7.0 keV). In this analysis, they also showed two blackbody model fits well. Kuiper, Hermsen, \\& Mendez (2004) \\cite{Kuiper Hermsen Mendez 2004} discovered hard X-ray emission up to $\\sim 150$ keV by using RXTE. ",
        "conclusions": "It is interesting to compare the NPEX parameters between this AXP and accretion-powered pulsars. $\\alpha$ and $-\\beta$ were $\\sim0$ and $\\sim -2$ for accretion-powered pulsars, while those are $3.54 \\pm 0.04$ and $0.92 \\pm 0.06 $ for this AXP (Table \\ref{tab: a}). Both of this AXP are larger than those of accretion powered pulsars by about three. For accretion powered pulsars, power-law component with $\\alpha \\sim 0$ supplies seed photons for Comptonization. On the other hand, for this AXP, the soft power-law component with $\\alpha = 3.54 \\pm 0.04$ would supply seed photons. This steep power-law may be a modified blackbody, which is originated from magnetar's surface and deformed by strong magnetic field of magnetar \\cite{Ozel 2001}. The $\\beta$ parameters can be interpreted as follows. The Comptonization is parameterized by $y$ parameter \\cite{Rybicki Lightman 1979}. For $y >> 1$, $\\beta$ becomes 2, meaning that the spectrum becomes saturated. For accretion-powered pulsars, the emission comes from the accretion column where the plasma density is large, so $y$ become large. On the other hand, the plasma density of this AXP is smaller than those of accretion-powered pulsars. The $y$ parameter become small, then $\\beta$ becomes small. The $y$ must be $\\Gtsim 1$ in this AXP, so that Comptonization process can be effective. This fact suggests that there are corona around magnetosphere of this magnetar \\cite{Beloborodov Thompson 2007}."
    },
    "0801/0801.4795_arXiv.txt": {
        "abstract": "{The recently-discovered lack of close binaries, among extreme horizontal branch (EHB) stars in Galactic globular clusters, has thus far constituted a major puzzle, in view of the fact that blue subdwarf stars~-- the field counterparts of cluster EHB stars~-- are well-known to present a high binary fraction.} {In this {\\em Letter}, we provide new results that confirm the lack of close EHB binaries in globular clusters, and present a first scenario to explain the difference between field and cluster EHB stars.} {First, in order to confirm that the lack of EHB binaries in globular clusters is a statistically robust result, we undertook a new analysis of 145 horizontal branch stars in NGC\\,6752, out of which forty-one belong to the EHB. To search for radial-velocity variations as a function of time, we repeated high-resolution ($R=18\\,500$) spectroscopy of all stars, four times during a single night of observations.} {We detected a single, hot (25\\,000~K), radial-velocity variable star as a close-binary candidate. From these results, we estimate an upper-limit for the close (period $P \\leq 5$~day) binary fraction $f$ among NGC\\,6752 EHB stars of 16\\% (95\\% confidence level), with the most probable value being $f=4\\%$. Thus our results clearly confirm the lack of close binaries among the hot HB stars in this cluster.} {We suggest that the confirmed discrepancy between the binary fractions for field and cluster EHB stars is the consequence of an $f$-age relation, with close binaries being more likely in the case of younger systems. We analyze theoretical and observational results available in the literature, which support this scenario. If so, an age difference between the EHB progenitors in the field and in clusters, the former being younger (on average) by up to several Gyr, would naturally account for the startling differences in binary fraction between the two populations. } ",
        "introduction": "\\label{capintro}  The presence of a large population of binaries among field B-type subdwarf (sdB) stars, also referred to as Extreme Horizontal Branch (EHB) stars, is well-established in the literature \\citep{Ferguson84,Allard94,Ulla98,Aznar01,Maxted01, Williams01,Reed04,Napiwotzki04}. The measurement of binary fraction varies from one survey to another, probably because of the different survey selection effects. It is generally agreed however, that the binary fraction must be large. Moreover, extensive analysis of orbital parameters \\citep{Moran99,Saffer98,Heber02,MoralesRueda04,MoralesRueda06} reveal that {\\em close} binaries ($P \\leq 10$~d) comprise a major fraction of the field EHB stars, with short-period systems---including either a degenerate or a low-mass main sequence (MS) companion---constituting about half of the entire field sdB population. Indeed, theoretical results indicate that these stars can be naturally explained within the context of binary star evolution \\citep{Han02,Han03}. Moreover, EHB stars of binary origin may also account for the UV flux excess (``UV upturn'') observed in elliptical galaxies \\citep{Han07}, although it should be noted that single-star scenarios have also been suggested \\citep[see][for a recent review]{Catelan07}. One way or another, the link between field sdB stars and binary systems is very well-established, both on observational and theoretical grounds, and close, short-period systems are the most frequently found amongst them.  It therefore came as a great surprise that the first radial-velocity (RV) surveys among EHB stars in globular clusters (GCs) revealed a remarkable lack of close binary systems \\citep[see][for a review]{Moni07b}.  In this {\\em Letter} we present, in light of new observational results, a first scenario to account for this discrepancy.   \\begin{figure*}[t] \\begin{center} \\includegraphics[width=17cm]{res6752.ps} \\caption{Results of our RV variability search. The maximum RV variation (${\\rm RV}_{\\rm max}-{\\rm RV}_{\\rm min}$) observed for each star is plotted at its effective temperature. For the sake of clarity, we plot only the 3$\\sigma$ error bars, which indicate the statistical significance of the observed variations. The data for the only star with a clearly-detected RV variation, is plotted using a larger, solid circle symbol, and thicker lines to indicate its errorbar.} \\label{figresults} \\end{center} \\end{figure*} ",
        "conclusions": "Using an independent dataset, a sample more than twice as large as in \\citet{Moni06a}, and a resolution in RV variations that is higher by almost a factor of two, for the first time we were able to find a good binary candidate among the EHB stars in \\object{NGC\\,6752}. That notwithstanding, our results confirm that the corresponding (close) binary fraction $f$ is very small in \\object{NGC\\,6752}, with a most likely value of $f = 4\\%$ and an upper limit of $f = 16\\%$ at the 95\\% confidence level.  There are hints that a small $f$ is not a peculiarity of \\object{NGC\\,6752}, but could also be a characteristic of other globular clusters \\citep{Moni06b}. This is in sharp contrast with the situation for field sdB stars, where close binaries are at least a factor of ten more frequent, comprising up to 70\\% of the entire sdB population \\citep[see \\S4 in][for a recent review]{Catelan07}. There is however no knowledge about cluster EHB stars in long-period binaries, or with a close low-mass companion, because no survey has investigated their role yet. These kind of systems are known to exist among field sdBs, but are just a minor population, and their presence in GCs would not alleviate the striking contrast with field results.  What is the origin of this startling difference between field and cluster EHB stars? We believe that it may not be completely unexpected: there should be a relation between the close binary fraction and the mean age of a sdB population, as a consequence of the different efficiency of binary channels responsible for EHB star formation.  Theoretical arguments strongly suggest that sdB stars in close binary systems should have undergone at least one common envelope (CE) phase. Although \\citet{Han02} found an upper limit for the initial mass of the sdB progenitor in this scenario, they also pointed out that within the permitted values a higher initial mass favors the CE channel, and leads to sdB binaries with shorter periods. In fact, a higher mass implies a more tightly-bound envelope, which requires a greater amount of (orbital) energy to be released. On the other hand, \\citet{Han02} explored the stable Roche Lobe Overflow (RLOF) scenario, and found that a higher initial progenitor mass makes it harder for the RLOF to be stable (see their Table~3), because the minimum mass of the companion increases with increasing progenitor mass. For higher values, fewer MS and white dwarf (WD) secondaries are sufficiently massive for the activation of this channel (sdB's with neutron star companions are indeed very rare). The progeny of systems that underwent stable RLOF, that is wide binaries with very long periods, is therefore generated primarily by progenitors of lower initial mass.  In light of these results, a relation between $f$ and the mean initial mass of the sdB progenitors could be naturally expected, hence implying a relationship between $f$ and (mean) age. More specifically, one may naturally expect that field sdB's formed from progenitors with a wide spectrum of initial masses (up to about $2\\,M_{\\sun}$), whereas in an old population (such as in globular clusters) only the progeny of less massive stars are currently found on the EHB, those of more massive ones having long evolved away from the He-burning phase. The stable RLOF is an efficient channel for sdB formation in the old case, while the CE one is not---and the CE itself would be released at earlier stages, before the orbits shrink substantially. Therefore, in old populations we should expect to find predominantly wide binaries, or/and single EHB stars formed through other channels \\citep[see][for a recent discussion]{Catelan07}. In fact, WD mergers, the third binary channel studied by \\citet{Han02}, can form (single) sdB stars, and is expected to be particularly efficient for old populations (see below). Moreover, it may be expected to play an important role within the dense environment of a globular cluster, where stellar encounters can harden close binaries \\citep{Heggie75} and thus enhance channel efficiency, as proposed earlier by \\citet{Bailyn89}. However, dynamical effects are not required for the formation of EHB stars \\citep{Whitney98}, as indicated by the lack of radial gradients among the EHB stars in \\object{NGC\\,2808} \\citep{Bedin00} and $\\omega$~Centauri \\citep[\\object{NGC\\,5139};][]{DCruz00}.  There are other results in the recent literature that provide support to our proposed scenario. In fact, \\citet{Napiwotzki04} already invoked a possible $f$-age or $f$-metallicity relation to explain their results. Among field sdB stars, they found a close binary fraction lower than in previous surveys, and noted that their sample was on average fainter, possibly including more thick disk or halo (hence older/metal poorer) stars. This suggested difference between the data used has never been fully investigated, and a further check of the populations of these stars would be invaluable. Even stronger support is provided by the recent simulations by \\citet{Han07}, modeling the UV upturn of elliptical galaxies from a binary sdB population. Their Figure~7, though aimed at reproducing the spectral energy distribution of elliptical galaxies, shows how the relative contribution of sdB's formed through different channels evolves with time. One finds that the contribution to the UV flux from RLOF sdB's is essentially constant with time, whereas the one from sdB's that experienced the CE phase first becomes important for a population of age 1.5~Gyr and then slowly fades, being more and more marginal for increasing ages. On the other hand, for populations older than 5~Gyr these authors find that sdB's formed through the merger channel dominate---but these are not binaries any longer. We suggest that these theoretical results implicitly hint at a solution to the heretofore puzzling lack of close EHB binaries in globular clusters, on the one hand, and their large numbers among field stars, on the other.  As an alternative to the binary scenario, it is worth noting that many single-star evolutionary channels have been invoked to explain EHB star formation in GCs, including interactions with a close planet \\citep[][see also \\citealt{Silvotti07}]{Soker98}, He mixing driven by internal rotation \\citep{Sweigart79,Sweigart97} or by stellar encounters \\citep{Suda07}, dredge-up induced by H-shell instabilities (\\citealt{vonRudloff88}, but see also \\citealt{Denissenkov03}), close encounters with a central, intermediate-mass black hole \\citep{Miocchi07}, and a sub-population of stars with high helium abundance \\citep[e.g.,][]{DAntona05}.  We note that any EHB formation mechanism involving single progenitors should be particularly inefficient among field stars (but possibly not so among GCs), since the EHB progenitors are on average younger, and thus more massive, in the field than in GCs. As a consequence, a much higher amount of RGB mass loss is needed to successfully produce a EHB star in the case of a field progenitor, whereas the ancient GC red giants have a much smaller envelope to be removed before they become EHB stars. This notwithstanding, a small component of single-star progeny may still be needed among field stars to fully explain the available observations \\citep{Lisker05}."
    },
    "0801/0801.2877_arXiv.txt": {
        "abstract": "Observations show the increase of high-frequency wave power near magnetic network cores and active regions in the solar lower atmosphere. This phenomenon can be explained by the interaction of acoustic waves with a magnetic field. We consider small-scale, bipolar, magnetic field canopy structure near the network cores and active regions overlying field-free cylindrical cavities of the photosphere. Solving the plasma equations we get the analytical dispersion relation of acoustic oscillations in the field-free cavity area. We found that the $m=1$ mode, where $m$ is azimuthal wave number, cannot be trapped under the canopy due to energy leakage upwards. However, higher ($m \\geq 2$) harmonics can be easily trapped leading to the observed acoustic power halos under the canopy. ",
        "introduction": "Waves play an important role in the dynamics of the solar atmosphere. Observations show an increase of high-frequency power ($\\nu > 5\\;$mHz) in the surroundings of active regions in velocity power maps sometimes called as photospheric power halos \\citep{braun1992,brown1992,hindbro1998,jainheb2002}. The halos were not found in Doppler power maps at lower frequencies ($3\\;$ mHz). Observations also show a lack of power halos in continuum intensity power  maps \\citep{hindbro1998,jainheb2002, Mug2005}.  On the other hand, recent observations reveal the decrease of the acoustic high frequency power in the chromosphere and its increase in the photosphere near active regions \\citep{Mug2003,Mug2005}. It has also been  shown that the quiet-Sun chromospheric magnetic network elements are surrounded by \"magnetic shadows\", which lack the oscillatory power at higher frequency range \\citep{maci2001,kru2001,Vecchio2007}. Therefore, both the photospheric power halos and the chromospheric magnetic shadows probably reflect the same physical process of acoustic wave interaction with overlying magnetic field \\citep[while in subsurface regions the rotational and the meridional non-uniform flows are supposed to have an impact on the formation of the acoustic wave power spectra.][]{sherg2005}.  The properties of propagating acoustic waves are closely related to the magnetic field structure. The numerical calculations show that the propagation of acoustic disturbances in the solar atmosphere is strongly determined by the overlying magnetic canopy \\citep{Rosenthal2002,Bogdan2003}.  The canopy has been usually modeled with purely horizontal magnetic field \\citep{Evans1990}, but recent high-resolution observations reveal more complex small-scale structure of the field \\citep{dew,centeno2007}. It has been suggested that the magnetic field has small-scale closed loop structure in the vicinity of network cores \\citep{maci2001,schr2003}. The inclined magnetic field may channel low-frequency photospheric oscillations in the chromosphere/corona \\citep{DePont2004}.  Here we use a model of small-scale bipolar magnetic canopy near a chromospheric network core and/or an active region. We suggest that granular cells may form field-free cylindrical cavities under the magnetic canopy due to the transport of magnetic flux towards boundaries. These cavities may trap high-frequency acoustic oscillations, while the lower-frequency harmonics may propagate upwards in form of magneto-acoustic waves.  Sect. 2 gives the analytical approach, obtained dispersion relation and resulting oscillation spectrum. Sect. 3 includes discussion and comparison of theoretical findings to observations. Sect. 4 briefly summarizes the results. ",
        "conclusions": "We have studied the spectrum of acoustic oscillations in the cylindrical field-free cavity regions under the small-scale magnetic canopy near the magnetic network cores and active regions. We found that the first harmonic of acoustic oscillations cannot be trapped in the cavity due to the energy radiation by fast magneto-acoustic waves in the canopy. However, the higher ($m \\geq 2$) harmonics can be trapped there, leading to the observed enhancement of high-frequency acoustic power in the photosphere. Future detailed study of the proposed mechanism including gravitational stratification and 3D models is necessary."
    },
    "0801/0801.0043_arXiv.txt": {
        "abstract": "The Local Supercluster is an ideal laboratory to study distribution of luminous and dark matter in the nearby Universe. The 1100 small groups have been selected using algorithm based on assumption that a total energy of physical pair of galaxies must be negative. The properties of the groups have been considered. ",
        "introduction": "We used an selection algorithm based on assumption that a total energy of physical pair of galaxies must be negative \\cite{MK00}. Individual mass of the galaxies are estimated from their K-luminosity using appropriate value of M/L. It is chosen to reproduce the well known nearby groups of galaxies, like the Centaurus~A, M~81, M~83 and IC~342.   \\begin{figure}[b] \\includegraphics[width=0.32\\textwidth]{lscmass.eps} \\includegraphics[width=0.64\\textwidth]{skymap.eps} \\caption{The mass distribution in the Local Supercluster. The left panel presents distribution of surface mass density in the Local Supercluster plane. The right panel shows an all sky map of surface mass density in the Supergalactic coordinates.} \\label{fig:sky} \\end{figure} ",
        "conclusions": ""
    },
    "0801/0801.0569_arXiv.txt": {
        "abstract": "The rapidly oscillating Ap stars are of importance for studying the atmospheric structure of stars where the process of chemical element diffusion is significant. We have performed a survey for rapid oscillations in a sample of 9 luminous Ap stars, selected from their location in the colour-magnitude diagram as more evolved main-sequence Ap stars that are inside the instability strip for rapidly oscillating Ap (roAp) stars.  Until recently this region was devoid of stars with observed rapid pulsations.  We used the VLT UV-Visual Echelle Spectrograph (UVES) to obtain high time resolution spectroscopy to make the first systematic spectroscopic search for rapid oscillations in this region of the roAp instability strip.  We report 9 null-detections with upper limits for radial-velocity amplitudes of $20-65$\\,\\ms\\ and precisions of $\\sigma=7-20$\\,\\ms\\ for combinations of Nd and Pr lines.  Cross-correlations confirm these null-results.  At least six stars are magnetic and we provide magnetic field measurements for four of them, of which three are newly discovered magnetic stars. It is found that four stars have magnetic fields smaller than $\\sim2$\\,kG, which according to theoretical predictions might be insufficient for suppressing envelope convection around the magnetic poles for more evolved Ap stars. Suppression of convection is expected to be essential for the opacity mechanism acting in the hydrogen ionisation zone to drive the high-overtone roAp pulsations efficiently. Our null-results suggest that the more evolved roAp stars may require particularly strong magnetic fields to pulsate. Three of the studied stars do, however, have magnetic fields stronger than 5\\,kG. ",
        "introduction": "Why do some stars oscillate while others do not? This question is particularly relevant for the class of rapidly oscillating Ap (roAp) stars for which the discovery of high-frequency oscillations came as a surprise \\citep{kurtz82}. The $\\delta$\\,Scuti pulsations of stars in this area of the colour-magnitude diagram (CMD) were not theoretically predicted to produce the observed high-frequency, high-overtone modes. Subsequently, considerable observational and theoretical efforts have been able to describe most pulsation properties of roAp stars through i) the oblique pulsator model \\citep{kurtz82}, ii) the probable driving mechanism of the pulsations \\citep{balmforthetal01} and iii) their link with the magnetic fields present in Ap stars (see, e.g., \\citealt{cunha06}; \\citealt{saio05}).  Ap stars are chemically peculiar stars that range from early-B to early-F spectral types with the majority having detectable magnetic fields with strengths of a few $10^2$ to a few $10^4$\\,G (\\citealt{bychkovetal03}, according to whom $\\sim55$ per cent have mean longitudinal fields above 400\\,G).  For the cool Ap stars (also called CP2 stars), abundance anomalies are typically visible in their spectra as abnormally strong absorption lines, particularly of rare-earth elements (REEs). Candidate roAp stars can therefore be selected photometrically due to the influence of these spectral features on narrow-band photometric indices such as those of the Str\\\"omgren filter system (\\citealt{martinez93}, see also Sect.\\,\\ref{sec:selection}).  The high-overtone roAp oscillations are thought to be driven by the opacity-mechanism acting in the partial hydrogen ionisation zone \\citep{dziembowskietal96,balmforthetal01,cunha02,saio05}.  The frequency range of the excited modes is related to the presence of strong magnetic fields in roAp stars that i) directly stabilise the low-frequency, low-order $\\delta$\\,Scuti oscillations, excited by the $\\kappa$-mechanism acting in the \\ion{He}{ii} partial ionisation zone \\citep{saio05}, and ii) indirectly enhance the driving of high-order oscillations by the $\\kappa$-mechanism acting in the partial hydrogen ionisation zone \\citep{balmforthetal01}. In particular, the stabilisation of the $\\delta$\\,Scuti modes can be explained by the dissipation of slow Alfv\\'en waves \\citep{saio05} and by the suppression of turbulent motion in the outer layers of the magnetic pole regions of roAp stars, which speeds up gravitational settling of helium and drains it from the \\ion{He}{ii} ionisation zone \\citep{theado05}.  An implication of the high radial order of the observed modes of roAp stars is that the observable pulsation amplitudes become depth dependent.  This pulsational structure provides a unique opportunity for studying the magneto-acoustic structure of the pulsations in 3D because of the element stratification in the atmospheres of these stars. Different elements, such as Fe, Pr, Nd, Eu and H, act as tracers for the pulsation structure due to their different properties regarding formation depth and extent of formation regions. Examples of depth-dependent measurements of pulsations are now numerous, such as \\citet{kurtzetal05, elkinetal2005b, ryabchikovaetal07}.  The roAp stars are thus unique objects for studying the interactions among stellar pulsation, magnetic fields and atomic diffusion. The latter is important for, e.g., estimates of globular cluster ages where He settling is significant, pulsation driving mechanisms in sdB and $\\beta$\\,Cephei stars where radiative levitation of Fe is needed for instability, and the Standard Solar Model where both He settling and radiative levitation of some metals are included.  Only 37 roAp stars are known at present (\\citealt{kurtzetal06b,tiwarietal07}) despite several searches for rapid pulsation in Ap stars, such as those by \\citet{nelsonetal93, martinezetal94, handleretal99, ashokaetal00, weissetal00, dorokhovaetal05,joshietal06}.  The selection of roAp candidates is mostly based on photometric indices pointing to chemical peculiarities in the spectra.  There is at present no clear correlation between radial velocity (RV) amplitude and photometric amplitude in roAp stars (see Table~1 of \\citealt{kurtzetal06b}).  For a (non-exhaustive) list of spectroscopic studies of roAp stars, see \\citet{kurtzetal06a}. \\citet{cunha02} calculated a theoretical instability strip for roAp stars and compared it to locations of 16 known roAp stars.  The region around the terminal age main-sequence is remarkably devoid of pulsators which, as Cunha points out, could be an observational bias. She predicted that more luminous and evolved Ap stars may pulsate with lower frequencies ($0.67 - 0.83$\\,mHz) which makes it harder to detect them in the typically short high-speed observing runs.  Note that this frequency range overlaps with the highest $\\delta$\\,Sct frequencies, such as for HD\\,34282 (0.92\\,mHz, \\citealt{amadoetal04}). Alternatively, she proposed that this absence of known evolved roAp stars could reflect a real deficiency in the fraction of pulsators in that region of the HR diagram, resulting from the fact that the magnetic field is less likely to suppress envelope convection in more evolved stars.  The photometric survey by \\citet{martinezetal94} was indeed less sensitive to low-amplitude roAp periods longer than 15\\,min and got a null-result for HD\\,116114 for which a 21-min oscillation was recently detected spectroscopically by \\citet{elkinetal05}.  A dedicated survey of luminous Ap stars therefore seems pertinent.  Ap stars inside the roAp instability strip that do not exhibit any detectable variability are called non-oscillating Ap (noAp) stars. As seen in, e.g., the astrometric HR-diagram by \\citeauthor{hubrigetal05} (\\citeyear{hubrigetal05}, their figure 2) for roAp and Ap stars, the apparent noAp stars occupy essentially the same regions as the roAp stars. However, the noAp stars appear to be systematically more evolved than the roAp stars \\citep{north97,handleretal99,hubrig00}. To fully understand the mechanism responsible for driving oscillations in roAp stars, it is essential to confirm and understand why so many stars with similar characteristics are apparently stable against pulsations. The apparent absence of oscillations in noAp stars does not, however, necessarily mean that oscillations are suppressed.  For instance, detection of oscillations may depend on lifetimes of excited modes (may be days for some roAp stars, \\citealt{handler04}), on beating between multiple modes, on the geometry of the magnetic field's orientation and the aspect in which the stellar surface is observed together with the surface distribution of the chemical elements used to detect the oscillations.  Further, the signal-to-noise ratio ($S/N$) of the obtained photometry or spectroscopy may simply be too low to detect the typically low-amplitude roAp pulsations. Finding high precision null-results for a statistically significant sample of roAp candidates would be a safe basis for concluding whether noAp stars exist or whether the current lack of known luminous roAp stars is an observational bias. From a theoretical point of view, stars located inside the theoretical instability strip may be noAp stars if the magnetic field is too weak to effectively suppress convection in the stellar envelope. A good test sample should, therefore, also include stars with strong magnetic fields.  To try to answer the question of whether the absence of observed variability among the luminous Ap stars is an observational selection effect only, or in fact due to intrinsic properties, we have studied 9 such stars at high radial-velocity precision. Our immediate goal with UVES was to search these luminous Ap stars for roAp oscillations at high radial velocity precision. Table\\,\\ref{tab:obslog} lists the known properties of the 9 targets and gives an observing log for the collected spectra. We find no pulsation in any of the 9 stars; Tables\\,\\ref{tab:cogpow} and \\ref{tab:ccpow} below present the null-results of the frequency analyses.  In the following sections, we describe selection and observation of the targets along with the data reduction in Sect.\\,2, then follows the data analyses including estimation of physical parameters and the radial-velocity analysis in Sect.\\,3. We show that our analyses reach the same precision as other radial-velocity studies of roAp stars in the literature and we then give the results for each target on a star-by-star basis. Finally we discuss the null-results in Sect.\\,4.  \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{\\label{tab:obslog}Observing log indicating number of observations and the mid-series heliocentric Julian Date for each series of spectra. Exposure times were 40\\,s, except for the fainter stars HD\\,110072 and HD\\,170565 where 80\\,s was used; readout and setup times were $24-27$\\,s.  The last two columns indicate $S/N$ measured in the continuum based on the standard deviation in the residual flux from two consecutive spectra. {\\it Lower} and {\\it Upper} respectively refer to wavelengths below or above the gap between the two CCDs at 6000\\,\\AA, while the ranges in parentheses correspond to the $S/N$ range for all spectra in that series. The $S/N$ estimates assume random noise only.} \\begin{tabular}{@{}lrccccll@{}} \\hline Star        &\\#spec& $\\alpha_{2000.0}$&$\\delta_{2000.0}$ &HJD\\,obs&Time&S/N Lower& $S/N$ Upper \\\\ &     & (h:m:s)    & (\\degr:\\,\\arcmin\\,:\\,\\arcsec\\,) &(d)&(h)&  &   \\\\ \\hline HD\\,107107 & 111 & 12:19:04.8 & -40:09:47 & 2453509.504 &1.94 & 79\\phantom{1} (77--82)   & 62\\phantom{1}  (56--77) \\\\ HD\\,110072 &  69 & 12:39:50.2 & -34:22:30 & 2453509.591 &2.06 & 50\\phantom{1} (43--55)   & 38\\phantom{1}  (32--49) \\\\ HD\\,131750 & 111 & 14:56:20.7 & -30:52:36 & 2453509.681 &1.98 & 80\\phantom{1} (74--87)   & 66\\phantom{1}  (57--74) \\\\ HD\\,132322 & 111 & 15:01:36.0 & -63:55:35 & 2453510.519 &2.03 &151            (147--155) & 141            (132--152)\\\\ HD\\,151301 & 111 & 16:49:28.3 & -54:26:47 & 2453510.715 &1.95 & 82\\phantom{1} (76--84)   & 67\\phantom{1}  (60--79) \\\\ HD\\,170565 &  85 & 18:30:08.2 & -02:35:26 & 2453509.778 &2.50 & 89\\phantom{1} (83--94)   & 77\\phantom{1}  (69--86) \\\\ HD\\,197417 & 125 & 20:48:48.1 & -72:12:44 & 2453509.879 &2.20 &107            (94--113)  & 93\\phantom{1}  (79--110)\\\\ HD\\,204367 & 111 & 21:28:41.2 & -25:38:39 & 2453510.798 &1.96 &113            (108--120) & 101            (93--115)\\\\ HD\\,208217 & 138 & 21:56:56.4 & -61:50:44 & 2453510.896 &2.52 &156            (142--165) & 148            (128--159)\\\\ \\hline \\hline \\end{tabular} \\end{minipage} \\end{table*} ",
        "conclusions": "The class of roAp stars is notoriously difficult to supplement with new members, as demonstrated by several photometric and spectroscopic studies before this one.  But motivated by the recent discovery of the luminous roAp star HD\\,116114 \\citep{elkinetal05}, we have spectroscopically tested a sample of 9 luminous Ap stars for rapid pulsations. Using lines known to show pulsations in roAp stars, we reach typical upper amplitude limits in radial velocity of: $40-75$\\,\\ms\\ ($\\sigma=13-24$\\,\\ms) for the line core of \\halpha, $20-65$\\,\\ms\\ ($\\sigma=7-20$\\,\\ms) when combining all measured Nd and Pr lines, and $20-40$\\,\\ms\\ ($\\sigma=7-11$\\,\\ms) when combining all measured lines.  With cross-correlations, using large wavelength regions, we typically reach upper amplitude limits of $4-10$\\,\\ms\\ ($\\sigma=1-4$\\,\\ms). In spite of a clear theoretical prediction \\citep{cunha02} and empirical (HD\\,116114) evidence for roAp pulsations in this part of the Hertzsprung--Russell diagram, we end up with 9 null-results, or noAp stars.  A number of questions are therefore pertinent to discuss.  \\vspace{2mm} \\noindent{\\em How well does our test sample resemble known roAp stars?} \\vspace{1mm}  \\noindent All studied stars have strong REE lines, except for HD\\,204367 and possibly also HD\\,197417. They have the core-wing anomaly typical for roAp stars, and {\\it Hipparcos} luminosities with our temperature estimates place them inside the predicted roAp instability strip. Several appear to have spotted surface distributions of REEs (such as HD\\,170565, HD\\,151301, HD\\,132322 and perhaps also HD\\,197417) which is typically associated with the strong magnetic fields common in known roAp stars.  Indeed most of the stars are magnetic, and cover a range in magnetic field strengths of $0.4-8.0$\\,kG, comparable to that of known roAp stars \\citep{kurtzetal06b}. In the case of the sharp-lined HD\\,110072, we compared its spectrum in detail with those of two known roAp stars, and found remarkable similarities for the REEs.  However, HD\\,110072 and HD\\,208217 are double-lined and single-lined binaries, respectively, which might indirectly influence their stability to high-frequency pulsations by reducing the magnetic field intensity (see \\citealt{cunha02} and references therein). Still, the orbits are probably too wide in both cases for tidal interaction to occur and HD\\,208217 has a known strong magnetic field ($\\left<B\\right>=8$ kG). The known roAp stars have cases of wide binaries, such as $\\beta\\,$CrB (spectroscopic binary), HR\\,3831, $\\alpha$\\,Cir, $\\gamma$\\,Equ and HD\\,99563 (visual binaries).  In these regards, the studied sample has the characteristics of roAp stars.  \\vspace{2mm} \\noindent{\\em Could pulsations have been overlooked?} \\vspace{1mm}  \\noindent\\citet{kurtzetal06b} published radial velocity amplitudes for the \\halpha\\ cores of 16 roAp stars.  Of these, only 3 have amplitudes below 75\\,\\ms\\ (HD\\,116114, HD\\,154708 and HD\\,166473, of which two have amplitudes above 3\\,$\\sigma$), while the rest range from $148-2528$\\,\\ms.  Further, radial velocity series for Pr and Nd line measurements in UVES spectra of these 16 roAp stars (Kurtz et al., partly unpublished), show that about 75 per cent of the measured and significant (3\\,$\\sigma$) amplitudes are within the range $350-1600$\\,\\ms.  Five of these stars have very small Nd and Pr amplitudes ($60-90$\\,\\ms): HD\\,166473, HD\\,116114, $\\beta$\\,CrB, 33\\,Lib and HD\\,154708. In such difficult cases, other lines or combinations of several lines makes detection of pulsations possible. We successfully tested our procedures on the two latter roAp stars in Sect.\\,\\ref{sect-rvshift}, and also used combinations of several lines, including of different elements. More of the other 19 known roAp stars have low amplitudes, such as 10\\,Aql, but our tests and analyses show that we reach these amplitude levels and should have detected such rapid pulsations if present in the studied sample.  A complication for our analyses is typical \\vsini\\,$\\sim10-30$\\kms combined with double lines due to either spots on the stellar surface and/or magnetic splitting, which results in considerable line blending and makes line identification and analysis more difficult.  However, our radial-velocity analysis was based partly on cross-correlations that are more robust than line measurements (as shown by our tests for two roAp stars) and this method similarly results in flat amplitude spectra that exclude rapid oscillations to relatively small roAp-amplitude levels.  It also seems improbable that, e.g., unfavourable viewing angles of the global pulsations or short mode lifetimes could explain a momentary lapse of detectable pulsation amplitudes simultaneously in all nine stars.  In fact, we know independently that HD\\,208217 was observed near its magnetic negative extremum where roAp pulsations are expected to have maximum amplitudes. Future surveys like this one may benefit from being repeated at different rotation phases.  We also note that the near-normal REE abundances of HD\\,197417 and HD\\,204367 reduce the probability that they are roAp stars, given the strong peculiarity of all known roAp stars.  \\vspace{2mm} \\noindent {\\em Do these stars really not pulsate?} \\vspace{1mm}  \\noindent In pulsators such as roAp stars oscillations are intrinsically unstable. Their excitation depends on the balance between the driving and damping of the oscillations over each pulsation cycle.  In roAp stars this balance is thought to be particularly delicate. On one hand, the amount of energy input through the opacity mechanism acting on the hydrogen ionisation region depends strongly on the interaction between the magnetic field and envelope convection, being maximal in the regions where envelope convection is suppressed \\citep{balmforthetal01,cunha02}.  On the other hand, the direct effect of the magnetic field on pulsations can introduce significant energy losses, through slow Alfv\\' en waves in the interior and through acoustic waves in the atmosphere, both resulting from mode conversion in the magnetic boundary layer \\citep{cunha00,saio05}.  Due to this delicate balance, it is not too surprising that roAp and noAp stars occupy the same locus in the HR diagram. Despite the developments in theoretical studies of linear non-adiabatic pulsations in models of roAp stars, we still lack a theoretical study that takes into account all these phenomena simultaneously and, thus, cannot firmly predict the conditions under which pulsations should be expected in roAp stars. In fact, both studies of \\cite{cunha02} and \\cite{saio05} considered the extreme case in which envelope convection is fully suppressed. Moreover, the first of these studies did not consider the direct effect of the magnetic field on pulsations and neglected the energy losses as a result of mode conversion, and the second study, while considering mode conversion, assumed the waves are fully reflected at the surface, hence neglecting energy losses through acoustic running waves in the atmosphere.  As discussed by \\cite{cunha02}, the condition for suppression of envelope convection, which seems necessary to make the high frequency modes unstable, is in principle harder to fulfil in evolved stars due to the increase with age of the absolute value of the buoyancy frequency in the region where hydrogen is ionised. Hence, it is likely that in evolved stars oscillations are excited only if the magnetic field is relatively strong.  Unfortunately, the complexity of the interaction between magnetic field and convection makes it impossible to derive a global convective stability criterion, even if local criteria for convective stability may be established \\citep{Gough66,Moss69} \\citep[see also][for an extensive discussion on this subject]{theado05}. Thus, the magnetic intensity needed to suppress convection at a given age, for a given mass, is very hard to establish.  The more evolved roAp star HD\\,116114, in which a relatively low frequency oscillation was found well in agreement with theoretical predictions, has a magnetic field modulus of $\\approx 6$~kG. In contrast with this, most stars in our sample have estimated mean magnetic field moduli around or below 2~kG. The clear exceptions are HD\\,107107, HD\\,131750 and HD\\,208217. Of these three, the latter is clearly an important test case to check the theoretical predictions. It has the strongest confirmed magnetic field in our sample and is strongly peculiar. However, we observed HD\\,208217 when one of its magnetic poles were almost visible, so pulsation should have been near its maximum amplitude.  From the observational point of view, one way to investigate the conditions under which roAp star oscillations are excited, and thus test theoretical models, is by identifying systematic differences between roAp and noAp stars.  This study would have been able to detect pulsations in all the known roAp stars, and any missed rapid pulsations must have amplitudes lower than these.  Hence we conclude that based on the obtained data, most stars in our sample are indeed noAp stars, and also excellent roAp candidates.  Despite this conclusion, it is premature to state that we can confirm the evidence that noAp stars are in average more luminous and more evolved than roAp stars, as indicated by earlier studies based on photometric surveys for pulsations in roAp candidates \\citep{north97,handleretal99,hubrig00}. To conclude that, we would also need to search for rapid pulsations in a control group of less evolved roAp stars with lower luminosity and compare the frequency of null-results in the two cases.  Such a survey has been started and results are expected in the near future.  The next step is to analyse the noAp stellar atmospheres in detail, taking temperature gradients and the abundance stratification into account.  New spectra are needed of HD\\,132322 to clarify the origin of its double line structures, and of HD\\,110072 to verify its secondary spectrum and to test for the `spurious' absorption lines in the recent FEROS spectrum. Polarimetry of HD\\,110072 and HD\\,204367 is needed to confirm the detected magnetic fields."
    },
    "0801/0801.2270_arXiv.txt": {
        "abstract": "{} {The present work provides the results of the first six years of operation of the systematic night-sky monitoring at ESO-Paranal (Chile).} {The $UBVRI$ night-sky brightness was estimated on about 10,000 VLT-FORS1 archival images, obtained on more than 650 separate nights, distributed over 6 years and covering the descent from maximum to minimum of sunspot cycle n. 23. Additionally, a set of about 1,000 low resolution, optical night-sky spectra have been extracted and analyzed.} {The unprecedented database discussed in this paper has led to the detection of a clear seasonal variation of the broad band night sky brightness in the $VRI$ passbands, similar to the well known semi-annual oscillation of the Na~I~D doublet. The spectroscopic data demonstrate that this seasonality is common to all spectral features, with the remarkable exception of the OH rotational-vibrational bands. A clear dependency on the solar activity is detected in all passbands and it is particularly pronounced in the $U$ band, where the sky brightness decreased by $\\sim$0.6 mag arcsec$^{-2}$ from maximum to minimum of solar cycle n.~23. No correlation is found between solar activity and the intensity of the Na~I~D doublet and the OH bands. A strong correlation between the intensity of N~I 5200\\AA\\/ and [OI]6300,6364\\AA\\/ is reported here for the first time. The paper addresses also the determination of the correlation timescales with solar activity and the possible connection with the flux of charged particles emitted by the Sun.} {} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.2100_arXiv.txt": {
        "abstract": "We investigate the evolution of accretion luminosity $L_{\\rm acc}$ and stellar luminosity ${L_\\ast}$ in pre-mainsequence stars. We make the assumption that when the star appears as a Class II object, the major phase of accretion is long past, and the accretion disc has entered its asymptotic phase. We use an approximate stellar evolution scheme for accreting pre-mainsequence stars based on Hartmann, Cassen \\& Kenyon, 1997. We show that the observed range of values $k = L_{\\rm acc}/L_\\ast$ between 0.01 and 1 can be reproduced if the values of the disc mass fraction $M_{\\rm disc}/M_*$ at the start of the T~Tauri phase lie in the range 0.01 -- 0.2, independent of stellar mass. We also show that the observed upper bound of $L_{\\rm acc} \\sim L_\\ast$ is a generic feature of such disc accretion. We conclude that as long as the data uniformly fills the region between this upper bound and observational detection thresholds, then the degeneracies between age, mass and accretion history severely limit the use of this data for constraining possible scalings between disc properties and stellar mass. ",
        "introduction": "The claimed relationship between accretion rate, $\\dot{M}$, and mass, $M_*$, in pre-mainsequence stars (e.g. Natta, Testi \\& Randich 2006 and references therein) originates from the more direct observational result that in Class II pre-mainsequence stars the accretion luminosity, $L_{\\rm acc}$, is similarly found to correlate with stellar luminosity, $L_\\ast$. The accretion luminosity is deduced from the luminosity in the emission lines of hydrogen (in Natta et al., 2006, these are Pa$\\beta$ and Br$\\gamma$), together with modeling of the emission process (Natta et al. 2004; Calvet et al. 2004). The modeling appears to be relatively robust. In order to deduce the quantities $\\dot{M}$ and $M_*$ from the quantities $L_{\\rm acc}$ and $L_\\ast$, it is necessary to be able to deduce the stellar properties. This is done by placing the object in a Hertzsprung-Russell diagram and making use of theoretical pre-mainsequence tracks for non-accreting stars (e.g. D'Antona \\& Mazzitelli 1997).  The result of this exercise indicated a correlation between $\\dot{M}$ and $M_*$ of the form $\\dot{M} \\propto M_*^{\\alpha}$, where $\\alpha$ is close to $2$. The simplest theoretical expectation (in the case of disc accretion where neither the fraction of material in the disc nor the disc's viscous timescale scale systematically with stellar mass) is instead that $\\dot{M} \\propto M_*$. The claimed  steeper than linear relationship thus motivated a variety of theoretical ideas about specific scalings of disc parameters with stellar mass (Alexander \\& Armitage 2006, Dullemond et al 2006).  \\begin{figure} \\resizebox{\\colwidth}{!}{\\includegraphics[angle=0]{fig1_till.eps}} \\caption{The distribution of Classical T Tauri stars in Ophiuchus in the $L_\\ast$, $L_{\\rm acc}$ plane. Detections (filled squares) and upper limits (crosses) deduced from emission line data. Adapted from Natta et al 2006.} \\label{hereifany2} \\end{figure}  Clarke \\& Pringle (2006) however suggested that the claimed correlation between $\\dot{M}$ and $M_*$ could be understood as follows. They noted that the distribution of datapoints in the plane of $L_\\ast$ versus $L_{\\rm acc}$ (Figure 1) more or less fills a region that is bounded, at high $L_{\\rm acc}$ by the condition $k = L_{\\rm acc}/L_\\ast \\sim 1$ and, at low $L_{\\rm acc}$,  by observational detection thresholds, which roughly follow a relation of the form $L_{\\rm acc} \\propto L_\\ast^{1.6}$. They noted that when this relation between $L_{\\rm acc}$ and $ L_\\ast$ is combined with specific assumptions about the relationship between $L_\\ast$ and $M_*$, and between $L_\\ast$ and stellar radius $R_\\ast$ (based on placing the stars on pre-mainsequence tracks at a particular age, or narrow spread of ages) then the resulting relationship between $\\dot{M}$ and $M_*$ is indeed $\\dot{M} \\propto M_*^2$. They thus claimed that it is currently impossible to reject the possibility that the claimed, steeper than linear, relationship is an artefact of detection biases.   Clarke \\& Pringle (2006) also pointed out that the distribution of detections in the $L_{\\rm acc}$ -- $L_\\ast$ plane may nevertheless  be used to tell us something about conventional accretion disc evolutionary models. In this paper we explore the extent to which this might be achieved. The first question which needs to be addressed is why there is a spread of values of $L_{\\rm acc}$ at a given value of $L_\\ast$. Secondly we need to understand {\\it why} the upper locus of detections corresponds to $L_{\\rm acc} \\sim L_\\ast$, since there seems no obvious reason why this should correspond to a detection threshold or selection effect.   We begin here  (Section 2) by setting out a simple argument why -- as a generic property of accretion from a disc with mass much less than the stellar mass -- one would expect the distribution to be bounded by the condition $L_{\\rm acc} \\sim L_\\ast$. In Section 3 we use the ideas of Hartmann, Cassen \\& Kenyon (1997) to develop an approximate evolutionary code for accreting pre-mainsequence stars and demonstrate its applicability by comparison with the tracks of pre-mainsequence stars computed by Tout, Livio \\& Bonnell (1999). In Section 4 we apply the code to pre-mainsequence stars accreting from an accretion disc with a declining accretion rate and investigate the range of disc parameters required to obtain the observed range of $L_{\\rm acc}/L_\\ast$. We draw our conclusions in Section 5.   \\section {The maximum ratio of accretion luminosity to stellar luminosity}  From a theoretical point of view there are three relevant timescales in the case of star gaining mass by accretion from a disc. The first is  \\begin{equation} t_M = M_*/\\dot{M}, \\end{equation} which is the current timescale on which the stellar mass is increasing. The second is  \\begin{equation} t_{\\rm disc} = M_{\\rm disc}/\\dot{M}, \\end{equation} which is the current timescale on which the disc mass is decreasing. The third is the Kelvin-Helmholtz timescale  \\begin{equation} t_{\\rm KH} = \\frac{GM_*^2}{R_\\ast L_\\ast}, \\end{equation} which is the timescale on which the star can radiate its thermal energy. It is also the timescale on which a star can come into thermal equilibrium. On the pre-main sequence, when the luminosity of the star is mainly due to its contraction under gravity, $t_{\\rm KH}$ is also the evolutionary timescale.  We see immediately  (since $L_{\\rm acc} \\sim G M_* \\dot M/R$) that  \\begin{equation} \\frac{L_{\\rm acc}}{L_\\ast} = \\frac{t_{\\rm KH}}{t_M}. \\end{equation}   The pre-mainsequence stars in which we are interested, Class~II (the Classical T~Tauri stars), all have the following properties:  (i) They have passed through the major phase of mass accretion which occurs during the embedded states Class 0 and Class I.  This implies typically that they have been evolving for some time with their own gravitational energy as the major energy source. Deuterium burning interrupts this briefly but only delays contraction by a factor of 2 or so.  Thus for these stars we may expect that they have an age $T \\approx t_{\\rm KH}$ roughly.  (ii) The masses of their accretion discs are now small, $M_{\\rm disc} < M_*$. Further, accretion disc models at such late stages -- i.e. the so called asymptotic stage when most of the mass that was originally contained in them has been accreted on to the central object --  exhibit a  power law decline in $\\dot {M}$ with time and thus have  the generic property that the age of the disc is about $t_{\\rm disc}$. If most of the star's lifetime has been spent accreting in this asymptotic regime, we may roughly equate the age of the disc and the age of the star and write $T \\approx t_{\\rm disc}$.  Given this, we see immediately that, for these stars, we expect $t_M = M_*/\\dot{M} > M_{\\rm disc}/\\dot{M} = t_{\\rm disc} \\approx t_{\\rm KH}$ and therefore that $L_{\\rm acc} <  L_\\ast$.  In the following section we test this argument by evolving a suite of model star-disc systems in the $L_\\ast - L_{\\rm acc}$ plane. ",
        "conclusions": "We have shown, using simplified stellar evolution calculations for stars subject to a time dependent accretion history, that a plausible morphology of the $L_* -  L_{\\rm acc}$ plane for T~Tauri stars has the following features,  (i) an upper locus corresponding to $L_* \\approx L_{\\rm acc}$ and  (ii) diagonal tracks for falling accretion rates as stars descend Hayashi tracks.   We have argued that (i) is a generic feature of stars which are accreting from a disc reservoir with mass less than the central star coupled with the assumption that the mass depletion timescale of the disc is comparable to the age  of the star. \\footnote {We may in fact argue that the existence of such an upper limit in the data is a good signature of temporally declining accretion through a disc, since if Classical T Tauri stars instead accreted external gas at a constant rate (e.g. Padoan  et al 2005), no such feature would be apparent in the data}. The first is necessary for stars to be classified as Classical T~Tauri stars while the second results from the gradual decay of the accretion rate.   We show that the slope of of the diagonal tracks (ii) is related to the index of the power law decline of accretion rate with time at late times.  This itself depends on the scaling of viscosity with radius in the disc. For reference, if $\\nu \\propto R$, then $L_{\\rm acc} \\propto L_\\ast^{1.7}$ along such tracks, whereas if $\\nu \\propto R^{1.5}$ then $L_{\\rm acc} \\propto L_\\ast^{2.4}$.  We however argue that, as long as observational datapoints appear to fill the whole region of the $L_\\ast - L_{\\rm acc}$ plane that is bounded by the upper limit (i) and luminosity dependent sensitivity thresholds, it will remain impossible to deduce what  viscosity law or what dependence of disc parameters on stellar mass will be  required in order to populate this region. This is because of the severe degeneracies that exist between stellar mass, age and accretion history (our results show that in order to populate the desired region, it is necessary that there is, at the least,  a spread in both stellar mass and initial disc properties). It is not, however,   necessary to posit any particular scaling of  disc properties with stellar mass in order to populate the plane, in contrast to the hypotheses of Alexander \\& Armitage 2006 and Dullemond et al 2006. Likewise, although, as noted above, the trajectories followed by individual stars in this plane are a sensitive diagnostic of the disc viscosity law, this information is washed out in the case of an ensemble of stars which simply fill this  plane.  We therefore conclude that such data could only yield useful insights if regions of the plane turned out to be unpopulated (or very sparsely populated). There is a shadow of a suggestion in Figure 1 that at low $L_\\ast$  the data falls further below the locus $L_{\\rm acc} = L_\\ast$ than at higher luminosities although, given the sample size, this difference is not statistically significant (Clarke \\& Pringle 2006). Any such edge in the distribution could in principle re-open the possibility of constraining disc models. We however emphasise that this low luminosity regime can in any case not be explored by the simple evolutionary models that we employ here, and that -- if it becomes necessary to investigate this regime --  it will be essential to use full stellar evolution calculations.  Finally, we note that we have chosen to relate our study to observational data in the $L_{\\rm acc} -L_*$ plane because the ratio of these quantities is straightforwardly derivable from emission line equivalent widths. Thus this comparison is more direct than the alternative, i.e. the  comparison  with observational data that has been transformed into the plane of $\\dot M$ versus $M_*$ via the use of isochrone fitting and the application of theoretical mass-radius relations. It is nevertheless worth  noting that the upper locus of $L_{acc} \\sim L_*$ corresponds to a line with $\\dot M \\propto M_*$. Although it is often taken as self-evident that this is the expected ratio (in the case that the maximum initial ratio of disc mass to stellar mass is independent of stellar mass), this actually only corresponds to $\\dot M \\propto M_*$ in the case that the disc's initial viscous time is independent of stellar mass. Given the possible variety of viscosity values and radii of protostellar discs it is not at all evident that these should conspire together to give the same viscous timescale --  until, that is, that one realises that all discs in the asymptotic regime have viscous timescale which is equal to the disc age. The conclusion that the upper locus should follow the relation $\\dot M \\propto M_*$ is thus a very general one.   In summary, we conclude, as in the study of  Clarke \\& Pringle (2006), that the observed population of the $L_{\\rm acc} - L_*$ plane, and the consequent relationship between $\\dot {M}$ and $M_*$, is strongly driven by luminosity dependent sensitivity thresholds and an upper locus at $L_{\\rm acc} \\approx L_\\ast$.  The main new ingredient injected by our numerical calculations and plausibility arguments is that we now understand that this observed upper locus can be understood as a consequence of plausible accretion histories, given the current properties of the relevant pre-mainsequence stars. In addition, our computations indicate the range of initial disc properties that are required in order to populate the region of parameter space occupied by observational datapoints: we are able to reproduce the observed range of values of $L_{\\rm acc}/L_\\ast$ at a given value of $L_\\ast$ by assuming once the T~Tauri phase has been reached, the disc mass fractions lie roughly in the range $0.01 \\le M_{\\rm disc}/M_* \\le 0.2$, independent of stellar mass."
    },
    "0801/0801.0333_arXiv.txt": {
        "abstract": "We consider a gauge inflation model in the simplest orbifold $M_4 \\times S^1/\\mathbb{Z}_2$ with the minimal non-Abelian $SU(2)$ hidden sector gauge symmetry. The inflaton potential is fully radiatively generated solely by gauge self-interactions. Following the virtue of gauge inflation idea, the inflaton, a part of the five dimensional gauge boson, is automatically protected by the gauge symmetry and its potential is stable against quantum corrections. We show that the model perfectly fits the recent cosmological observations, including the recent WMAP 5-year data, in a wide range of the model parameters. In the perturbative regime of gauge interactions ($\\gfd \\lesssim 1/(2\\pi R\\mpl)$) with the moderately compactified radius ($10 \\lesssim R\\mpl \\lesssim 100$) the anticipated magnitude of the curvature perturbation power spectrum and the value of the corresponding spectral index are in perfect agreement with the recent observations.  The model also predicts a large fraction of the gravitational waves, negligible non-Gaussianity, and high enough reheating temperature. ",
        "introduction": "It is now widely accepted that an early period of accelerated expansion of the universe, or inflation~\\cite{inf}, can resolve many cosmological problems such as horizon problem and can provide the desired initial conditions for the subsequent hot big bang evolution of the observed universe~\\cite{books}. Many observational facts such as the flatness of the universe and the isotropy of the cosmic microwave background (CMB) are natural consequences of inflation, and hence they strongly support the existence of such a period of acceleration in the very early universe. In particle physics point of view, inflation occurs when one or more scalar fields, the inflaton fields, dominate the energy density of the universe with their potential being overwhelming~\\cite{Lyth:1998xn}. Under such a condition, dubbed slow-roll condition, curvature perturbation $\\mathcal{R}$ is produced which is nearly scale invariant and is heavily constrained by the measurements of the anisotropies of the CMB and the observations of the large scale structure~\\cite{obs}, including the recent Wilkinson Microwave Anisotropy Probe (WMAP) 5-year data set~\\cite{Komatsu:2008hk}. The slow-roll condition says that the inflaton potential should be very flat, i.e. the effective mass of the inflaton should be very small compared with the inflationary Hubble parameter. This is, however, quite difficult to be achieved: for example, in supergravity, any generic scalar field is expected to have an effective mass of $\\mathcal{O}(H)$~\\cite{sugrainf}, which completely spoils the desired slow-roll condition. It is thus natural to consider some symmetry principle which protects the small inflaton mass from large radiative corrections or supergravity effects.   Shift symmetry, under which the field is invariant with respect to the transformation \\begin{equation} \\phi \\to \\phi + a \\, , \\end{equation} with $a$ being an arbitrary constant, is one of such symmetries and the corresponding fields remain completely massless, i.e. their potential is exactly flat as long as shift symmetry is unbroken. When the symmetry is explicitly broken, they become pseudo Nambu-Goldstone bosons (pNGBs) and the potential acquires a tilt depending on the symmetry breaking scale $f$ and is of the form \\begin{equation}\\label{naturalinf_V} V(\\phi) = \\Lambda^4 \\left[ 1 \\pm \\cos\\left( \\frac{\\phi}{f} \\right) \\right] \\, . \\end{equation} The model of natural inflation~\\cite{naturalinf} makes use of such a pNGB with large $f$, which renders the potential very flat. This largeness, however, requires $f \\gg \\mpl \\equiv G^{-1/2} \\sim 10^{19} \\mathrm{GeV}$ and hence the valid region of successful natural inflation lies in the regime where we completely lose theoretical control. This has been regarded as a significant drawback of natural inflation.   An idea to evade this problem was suggested in Ref.~\\cite{Arkani-Hamed:2003wu}\\footnote{The idea of identifying the higher dimensional gauge field as the inflaton is also suggested in Ref.~\\cite{Kaplan:2003aj}.}. In this scenario, called gauge inflation, the inflaton field is essentially coming from a part of the fifth component of the gauge boson in five dimensional bulk $A_5$, by which the gauge invariant Wilson line is defined as \\begin{equation} e^{i\\theta} \\equiv \\exp \\left[i g \\oint A_5 dy \\right] \\, , \\end{equation} where $g$ is the gauge coupling constant. By the one-loop interactions with the charged particles, the potential of the canonically normalized field $\\phi \\equiv f\\theta$ has basically the same form as Eq.~(\\ref{naturalinf_V}) while $f=1/(g_4 L)$ has extra dimensional nature with $L$ being the size of the fifth dimension. Hence the potential can be trusted even when $f$ is larger than $\\mpl$ in the perturbative regime of gauge interaction, $g_4 \\lesssim 1/(\\mpl L)$. An interesting possibility to unify the grand unification theory (GUT) and cosmic inflation in the framework of the gauge inflation was suggested in Ref.~\\cite{Park:2007sp} by one of the authors of the present paper: starting from the higher dimensional $SU(5)$ GUT, the standard model is recovered by orbifold projection by $S^1/{\\mathbb{Z}_2}$ and a massless scalar boson from the fifth component of the gauge boson is shown to play the role of the inflaton with a fully radiatively generated flat potential.   In this paper, we take a different approach. Here we try to build the minimal model of gauge inflation using the simplest orbifold $S^1/\\mathbb{Z}_2$ and the minimal non-Abelian gauge group $SU(2)$ for the hidden sector. The obvious advantages of this approach is two-fold. First, by taking the hidden sector gauge interaction, the theory is less constrained than the visible sector gauge theory, e.g. GUT~\\cite{Park:2007sp}. This is important since the required coupling constant is typically very small, $g_4 \\sim 10^{-2}$, so it is hard to reconcile this smallness with the coupling constant unification. Second, the minimality can apply to the particle contents. Different from the Abelian case, the non-Abelian gauge theory accommodates the gauge self-interactions by which the potential for the inflaton ($\\sim A_5$) can be generated even without any additional charged matter fields. So the theory remains minimal in particle contents and its predictions are quite robust and free from parameters such as the masses and other quantum numbers of newly introduced matter fields. Combining the minimal choice of the orbifold and the gauge group, no exotic particles appear in the theory. The model is clean.   The structure of the paper is as follows. In the next section, we give the detail of the higher dimensional gauge theory and outline the subsequent cosmological scenario. Section~\\ref{sec_cos} is devoted to the cosmological evolution of the model and we analytically calculate the observable quantities produced during inflation and study their relations to the parameters of the model. We also address the issue of reheating and estimate the reheating temperature $T_\\mathrm{RH}$. In Section~\\ref{sec_con} we then conclude. ",
        "conclusions": "\\label{sec_con}     In this paper, we have presented a cosmological scenario from the hidden sector $SU(2)$ gauge symmetry in the five dimensional orbifold $M_4\\times S^1/\\mathbb{Z}_2$. The model is minimal in several aspects: the minimal non-Abelian gauge group and the minimal orbifold compactification with the minimal number of extra dimensions. Thanks to the non-Abelian nature, the bulk gauge boson, the fifth component $A_5$ in particular, could have a one-loop induced effective potential without introducing any exotic field in the model. This makes sure the minimality of the model. The advantage of this minimal setup is as follows: the inflaton field is a built-in ingredient of the theory and is automatically free from quantum gravitational effects because of its higher dimensional locality and the gauge symmetry. Fully radiatively generated one-loop potential is naturally able to support a long enough period of slow-roll inflation provided that the theory is weakly coupled, i.e. $\\gfd \\ll 1$, during the inflationary epoch. In very good numerical precision, the minimal model essentially provides a realization of natural inflation \\begin{eqnarray} V(\\phi) \\approx \\Lambda^4 \\left[ 1-\\cos \\left( \\frac{\\phi}{\\feff} \\right) \\right] \\, , \\end{eqnarray} with $\\Lambda^4 = 9R^{-4}/(2\\pi)^6$ and $\\feff = (2\\pi\\gfd R)^{-1}$. For $10 \\lesssim R M_{\\rm Pl} \\lesssim 100$ and  $1 \\lesssim \\feff/\\mpl \\lesssim 100$, the model predicts the observable cosmological quantities \\begin{equation} 1.2 \\times 10^{-5} \\lesssim \\mathcal{P}_\\mathcal{R} \\lesssim 4.9 \\times 10^{-5} \\, , \\end{equation} \\begin{equation} 0.952 \\lesssim n_\\mathcal{R} \\lesssim 0.966 \\, , \\end{equation} \\begin{equation} 0.03 \\lesssim r \\lesssim 0.13 \\, . \\end{equation} The power spectrum of the curvature perturbation $\\mathcal{P}_\\mathcal{R}$ and the corresponding spectral index $n_\\mathcal{R}$ are in good agreement with the current observations. While $|f_\\mathrm{NL}|$ is always far smaller than 1 and no detectable non-Gaussianity is expected, very interestingly the predicted tensor-to-scalar ratio $r$ is quite close to sensitivity of the near future cosmological experiments. This would be the first test of our minimal cosmological model. The reheating temperature $T_\\mathrm{RH}$ is estimated to be high enough to successfully follow the standard hot big bang evolution.     \\subsection*{Acknowledgements}   We are grateful to Misao Sasaki and the Yukawa Institute for Theoretical Physics at Kyoto University where some part of this work was carried out during ``Scientific Program on Gravity and Cosmology'' (YITP-T-07-01) and ``KIAS-YITP Joint Workshop: String Phenomenology and Cosmology'' (YITP-T-07-10). JG thanks Daniel Chung, L. Sriramkumar and Ewan Stewart for helpful conversations, and is partly supported by the Korea Research Foundation Grant KRF-2007-357-C00014 funded by the Korean Government. SCP appreciates Yasunori Nomura for his comments on low energy constraints of $\\gfd$ and also thanks C. S. Lim for encouragement to publish this paper."
    },
    "0801/0801.1429_arXiv.txt": {
        "abstract": "We perform a combined X-ray and strong lensing analysis of RX~J1347.5-1145, one of the most luminous galaxy clusters at X-ray wavelengths. {  We show that evidence from strong lensing alone, based on published VLT and new HST data, strongly argues in favor of a complex structure.} The analysis takes into account arc positions, shapes and orientations and is done thoroughly in the image plane. The cluster inner regions are well fitted by a bimodal mass distribution,  with a total projected mass of $M_{tot} = (9.9 \\pm 0.3)\\times 10^{14} M_\\odot/h$ within a radius of $360~\\mathrm{kpc}/h$ ($1.5'$). Such a complex structure could be a signature of a recent major merger as further supported by X-ray data. A temperature map of the cluster, based on deep Chandra observations, reveals a hot front located between the first main component and an X-ray emitting South Eastern sub-clump. The map also unveils a filament of cold gas in the innermost regions of the cluster, most probably a cooling wake caused by the motion of the cD inside the cool core region. A merger scenario in the plane of the sky between two dark matter sub-clumps is consistent with both our lensing and X-ray analyses, and can explain previous discrepancies with mass estimates based on the virial theorem. ",
        "introduction": "\\label{intro} Accurate determination of total mass of galaxy clusters is important to understand properties and evolution of these systems, as well as for many cosmological applications. Gravitational lensing, through multiple image systems (strong lensing) as well as from distortions of background sources (weak lensing), provides a reliable method to determine the cluster  mass,   which is  independent of the equilibrium properties of the cluster \\citep{mel99}. The lensing  mass determination can be compared to estimates based on measured X-ray surface brightness and temperature, which is instead based on the assumption of hydrostatic equilibrium  \\citep{sar88} or to dynamical estimates, which rely on the assumption of virialized systems. Combining optical, X-ray and radio observations of galaxy clusters is a major tool to investigate their intrinsic properties. In particular, the comparison of lensing and X-ray studies can give fundamental insights on the dynamical state of the galaxy clusters (see for example \\citet{allen2002, def+al04}), on the validity of the equilibrium hypothesis and on their 3-dimensional structure \\citep{def+al05, ser+al06, ser2007}.  RX J1347.5-1145 ($z=0.451$) is one of the most X-ray luminous and massive galaxy cluster known. This cluster has been the subject of numerous  X-ray \\citep{schindler95,sch+al97,allen2002,gi+sc04}, optical (Sahu et al.1998) and Sunyaev-Zeldovich (SZE) effect studies \\citep{kom+al01,kit+al04}. Formerly believed to be a well relaxed cluster, with a good agreement between weak-lensing \\citep{fi+ty97, kli+al05}, strong lensing \\citep{sahu98} and X-ray mass estimates \\citep{sch+al97}, more recent investigations revealed a more complex dynamical structure. { In particular}, a region of enhanced emission in the South-Eastern quadrant was first detected by SZE effect observations \\citep{kom+al01} and later confirmed by X-ray observations that also measured an hotter temperature for the excess component \\citep{allen2002}. This feature  has been interpreted as an indication of a recent merger event \\citep{allen2002,kit+al04}.  {Furthermore}, a spectroscopic survey on the cluster members   found  a velocity dispersion of $910 \\pm 130\\ \\mathrm{km~s^{-1}}$, which is  significantly smaller than that derived from weak lensing, $1500 \\pm 160\\ \\mathrm{km~s^{-1}}$ (Fisher \\& Tyson 1997), strong lensing, roughly 1300 $\\mathrm{km~s^{-1}}$ and X-ray analyses $1320 \\pm 100\\ \\mathrm{km~s^{-1}}$ \\citep[see their table~4]{co+kn02}. A major merger in the   plane  of the sky was proposed as a likely scenario to reconcile all measurements \\citep{co+kn02}.  In this paper, we further investigate the merger hypothesis  by performing a combined strong lensing and X-ray analysis of archive data. We perform a strong lensing investigation based on a family of multiple arc candidates first proposed in Brada{\\v c} et al. (2005)   using  deep VLT observations. Differently from previous studies, we take care of performing the statistical analysis in the lens plane, which is a more reliable approach than the source-plane investigation when working with only one multiple image system. We further refine our analysis by taking into account   not only the image positions, but also the shape and orientation of the arcs.   In addition, we exploit {\\it Chandra} observations to gain additional insights into the dynamical status of the cluster, through  spectral and morphological analyses of the X-ray halo, and to discriminate between different evolutionary scenarios.  The paper is organised as follows. Section~\\ref{sec:data} discusses the strong lensing image candidates selection from archive VLT and {  HST} data. Section~\\ref{pdsa} describes the statistical method used in the lensing analysis whereas Sec.~\\ref{Sec:xray} is devoted to the X-ray data analysis. Section~\\ref{sec:dis} discusses the merger hypothesis. Summary and conclusions are presented in Sect.~\\ref{sec:summary}. Throughout this paper we use a flat model of universe with a cosmological constant  with $\\Omega_m=0.3$ and $H_0=70\\ \\mathrm{km\\ s}^{-1} \\mathrm{Mpc}^{-1}$. This implies a linear scale of $5.77\\ {\\rm kpc}/\\arcsec$ at the cluster redshift. ",
        "conclusions": "\\label{sec:summary}  We  have analysed both the lensing and X-ray properties of RX~J1347.5-11.45, one of the most luminous and massive X-ray cluster known. Based on the analysis of an arc family, photometrically selected, we have estimated the total cluster  mass distribution, within a radius of $\\sim 500\\ \\rm{kpc}$ from the cluster centre. We performed a $\\chi^2$ analysis in the lens plane and scrupulously modelled the arc configuration. A model with two smooth dark matter components of similar mass   accurately  reproduces the observations  and yields a mass estimate in agreement with previous strong and weak lensing and X-ray studies. Our strong lensing analysis suggests a major merger between two sub-clumps of similar mass located within the central $300~\\mathrm{kpc}$.  X-ray observations further strengthen our view of a complex structure of the inner regions, revealing a hot front in the South-Eastern area, most probably a remnant of the occurred merger. This merging framework, which arises naturally from our strong lensing model, can also reconcile the observed discrepancy between dynamical mass estimates and X-ray, lensing and SZE  ones, which instead give consistent results. Agreement between X-ray and lensing mass estimates further indicates that the gas might have had time to virialize. Spectroscopic measurements additionally suggest that the merger is taking place in the plane of the sky.  Whereas the presence of a merger is confirmed on several grounds, its properties are though still unclear.  The detailed features of our model are strictly related to the selection of the members of the multiple image system, in particular to the inclusion of the arc AC.  Despite  all candidates should be confirmed spectroscopically,  a photometric analysis suggests that A4 and A5 almost undoubtedly belong to the same system. Remarkably, based on the shapes and orientations of A4 and A5 alone, without any constraint on AC, we can exclude the presence of a single mass component. A  spectroscopic confirmation of further arc candidates is the required step  to confirm and accurately define   the merging  scenario  and to provide a final, more accurate description of the dynamical state of RX J1347.5-1145."
    },
    "0801/0801.3276_arXiv.txt": {
        "abstract": "When  constraining the primordial non-Gaussianity parameter $f_{\\rm NL}$ with cosmic microwave background anisotropy maps, the bias resulting from the covariance between primordial non-Gaussianity and secondary non-Gaussianities to the estimator of $f_{\\rm NL}$ is generally assumed to be negligible. We show that this assumption may not hold when attempting to measure the primordial non-Gaussianity out to angular scales below a few tens arcminutes with an experiment like Planck, especially if the primordial non-Gaussianity parameter  is around the minimum detectability level with $f_{\\rm NL}$ between 5 and 10. In future, it will be necessary to jointly estimate the combined primordial and secondary contributions to the CMB bispectrum and establish $f_{\\rm NL}$ by properly accounting for the confusion from secondary non-Gaussianities. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.1645.txt": {
        "abstract": " ",
        "introduction": "Models with warped extra dimensions~\\cite{Randall:1999ee} have arisen in the last few years as strong candidates for a natural theory of electroweak symmetry breaking (EWSB).  The original solution to the hierarchy problem has been supplemented by the addition of a natural flavor structure~\\cite{flavor}, a custodial symmetry to protect the $T$ parameter~\\cite{Agashe:2003zs} and the $Z \\bar{b}_L b_L$ coupling~\\cite{Agashe:2006at}, and the realization of the Higgs as the pseudo Goldstone Boson of a broken global symmetry~\\cite{Contino:2003ve}--\\cite{Agashe:2005dk}.  These developments have produced calculable models that successfully address most of the mysteries related to the electroweak (EW) scale.  On the other hand, dark matter (DM), that is often considered one of the strongest --albeit indirect-- hints of physics beyond the Standard Model (SM), has so far lacked a generic implementation in models with warped extra dimensions.  The main reason is the inherent asymmetry in warped backgrounds, that do not posses a natural KK parity as the one present in Universal Extra Dimensions (UED)~\\cite{Appelquist:2000nn}.  In this article we explore a generic procedure to introduce an exact discrete exchange symmetry that results in new stable states, without introducing new parameters.  This is done via a doubling of part of the field content.  The exchange symmetry we advocate has been first introduced in \\cite{Panico:2006em} (where it was dubbed ``mirror symmetry'') and \\cite{Regis:2006hc} to alleviate the fine--tuning and to get a viable DM candidate in Gauge-Higgs Unification (GHU) models in flat space.  As already anticipated in \\cite{Panico:2006em}, it can be extended straightforwardly to warped models.  Given a bulk field, $\\phi$, satisfying certain boundary conditions (b.c.), the procedure consists of replacing $\\phi$ by a pair of fields, $\\phi_{1}$ and $\\phi_{2}$, and imposing the symmetry $\\phi_1 \\leftrightarrow \\phi_2$.~\\footnote{A similar discrete symmetry has recently been used in \\cite{Bai:2008cf} to get a DM candidate in the context of little Higgs theories.} The even linear combination $\\phi_+ \\equiv (\\phi_1 + \\phi_2)/\\sqrt{2}$ is identified with the original field (in particular, it inherits the b.c. obeyed by $\\phi$, as well as its couplings).  The couplings of the orthogonal combination, $\\phi_- \\equiv (\\phi_1 - \\phi_2)/\\sqrt{2}$ are determined by those of $\\phi_{+}$.  Under the exchange symmetry one has $\\phi_{\\pm}\\rightarrow \\pm \\phi_{\\pm}$, so that one can assign a multiplicative charge $+1$ to $\\phi_+$ and $-1$ to $\\phi_-$.  Provided the discrete exchange symmetry is an exact symmetry at the quantum level, the lightest Kaluza--Klein (KK) resonance among all $\\Z_2$-odd states in the model is absolutely stable, and will be referred to as the LOP (Lightest Odd Particle).  We argue that the above symmetry is indeed exact, by showing that possible 5D Chern--Simons (CS) terms, in general needed to restore gauge invariance in 5D theories, do not violate it~\\cite{Hill:2007zv}.  The choice of which fields to double must be guided by phenomenological considerations.  For example, DM direct and indirect searches impose stringent constraints on the possible couplings of the DM candidate to SM fields.  It is therefore natural to look for charge and color neutral fields that can lead to viable DM candidates.  In fact, the models we will consider always contain $U(1)$ factors, and it will be natural for the DM candidate to be a $U(1)$ massive gauge field, $X_{-}$.  This is similar to the 5D UED case and the GHU model of \\cite{Regis:2006hc} in flat space, in which the DM can be identified with the first KK mode of the hypercharge gauge field.  If the LOP were the only $\\Z_2$-odd particle, it would couple to SM fields only via non-renormalizable interactions.  For the typical scales involved, its annihilation rate would then be extremely small, and the resulting thermal relic density unacceptably large.  Hence, it is necessary in general to apply the previously described ``doubling'' construction to additional fields, making sure that $X_-$ remains the lightest among the $\\Z_{2}$-odd particles.  We can choose b.c. for the LOP so that it has a mass about one order of magnitude smaller than the mass of the first $\\Z_2$-even gauge resonances.  Thus, our DM candidate has typically sub--TeV masses, which, as we will see, can range from $\\sim 100$ GeV in Higgsless models up to $\\sim 700$ GeV in Randall--Sundrum (RS) models.  Together with its mass, the coupling of the LOP to SM matter is the other crucial parameter governing its relic density.  In all the cases we will consider, the $U(1)$ symmetry is related to the SM $U(1)_Y$ symmetry, with a coupling constant of electroweak size, and the $\\Z_{2}$-odd fields will consist of $X_-$ and a subset of fermion fields, typically associated with the top or bottom quarks. It is clear that our DM candidate is a weakly interacting massive particle (WIMP) with approximately the correct mass and couplings to give rise to the observed DM relic density in the universe.  The construction outlined above turns out to be particularly natural in the specific class of GHU scenarios in warped space.\\footnote{It is known that such models can also be seen, thanks to the AdS/CFT dictionary \\cite{Ads-cft}, as strongly coupled 4D composite Higgs models.  However, since calculability requires the 5D picture, we will mostly use the 5D language, adopting only occasionally the 4D dual language.  } In this case, the additional set of fields introduced by our procedure is not only minimal, but is also singled out by the role the doubled sector plays in generating the EW scale dynamically. Although it is not easy to use the discrete symmetry to protect EW observables from large tree-level corrections, as in R-parity preserving supersymmetric scenarios, Little Higgs theories with T-parity, or UED with KK parity, we will see that the physics of EWSB in the GHU model with the discrete symmetry naturally leads to a region in parameter space where the constraints due to precision measurements are relaxed.  Furthermore, this same region of parameter space, in which coannihilation effects can be relevant, predicts a dark matter abundance in accord with observation.  Thus, the DM sector is tightly connected to the physics of EWSB, and plays an indirect role in leading to agreement with precision constraints.\\footnote{See \\cite{Hambye:2007vf} for a model in which EWSB and DM are related.  In this reference, however, the hierarchy problem is not addressed.} Also, the model is rather predictive, with several fermionic resonances nearly degenerate with the DM particle that should lead to an exciting collider phenomenology.  We also consider some non-perturbative effects --the formation of bound states and the effect of QCD Coulomb-like forces-- which might invalidate the usual perturbative computation of the relic density. We argue that the problem of the formation of bound states does not occur and that the effect of Coulomb--like forces on the perturbative cross-sections is negligible.  We also discuss in some detail the degree of\u00cafine-tuning involved in the EWSB pattern, and compare to supersymmetric extensions of the standard model.  It turns out that in the GHU framework (with or without DM) the fine-tuning is somewhat worse than expected by naive considerations.  We point out, however, that here the fine-tuning seems to be associated with accommodating a top mass of order the EW scale, rather than with an intrinsic tension in the Higgs sector. When restricted to the region of parameter space with a fixed top mass, the low-energy properties of the model turn out to be essentially insensitive to the microscopic parameters of the model, and therefore to the detailed properties of the new physics, which is a very interesting feature.  Although we find the application of our construction to the GHU framework particularly appealing, we stress that it is of more general applicability.  We also briefly discuss the implementation of the discrete exchange symmetry in two other models with warped extra dimensions: the simplest RS model with the SM fermions and gauge bosons in the bulk, and a Higgsless model.  Hence, our construction leads to viable DM candidates in a variety of scenarios, without the introduction of new parameters.  The organization of the paper is as follows.  In section~\\ref{sect:Z2} we review the generic properties of the exchange $\\Z_2$ symmetry that gives rise to a dark matter candidate.  In section~\\ref{sec:GHU} we discuss in detail its implementation in a model of GHU. In particular, we study the interplay between EWSB, the EW constraints and the DM relic abundance, the (ir)relevance of the above mentioned non-perturbative effects, and discuss the fine-tuning.  In section~\\ref{sec:othermodels} we describe the implementation of the discrete exchange symmetry in other models. Section~\\ref{sec:phenomenology} is devoted to particular details of DM collider phenomenology and DM direct detection in our construction. We comment on the issue of anomalies in section~\\ref{sect:anomalies}, and conclude in section~\\ref{sect:conclusions}.  We relegate some technical details to the Appendices.   %----------------------------------------------------------------------- ",
        "conclusions": "}   In this paper we have proposed a generic construction that allows to endow given models with a DM candidate.  This is achieved through an extension in which the model acquires a $\\Z_{2}$ exchange symmetry. The lightest $\\Z_{2}$-odd particle is then absolutely stable.  We have considered several models with warped extra dimensions. Although these scenarios are well-motivated extensions of the SM, explaining both the Planck-weak scale hierarchy as well as the flavor structure of the SM, they do not contain, generically, stable particles that can account for the observed DM component of the universe.  We have shown that our mechanism can easily solve this deficiency.  We have payed special attention to a class of warped scenarios that is particularly appealing: GHU/composite Higgs scenarios.  In this case, the $\\Z_{2}$ structure responsible for the stability of the DM candidate is tightly connected to the physics that leads to the dynamical breaking of the EW symmetry.  As in supersymmetry, the dark matter mass and couplings are intimately connected to the EW scale. In fact, not only does our construction not introduce new parameters, but there is a further sense in which it can be considered minimal. As was emphasized in the main text, the physics of EWSB in such models crucially depends on certain fields that have two properties: they are fermionic fields without zero-modes and, through their strong connection to the top sector, they give a significant --in fact, crucial-- contribution to the Higgs potential.  This is precisely the sector that gives rise to the $\\Z_{2}$-odd particles that can lead to a realistic dark matter candidate.  As a result, the $\\Z_{2}$-odd sector, through its contribution to the Higgs potential, plays a key role in the realization of an EWSB vacuum with the desired physical properties.  As explained in detail in the main text, the effects due to the $\\Z_{2}$-odd sector go in the direction of alleviating the constraints imposed by low-energy measurements.  The picture that emerges is rather compelling: the observed DM relic abundance and the very precise measurements of the EW observables at LEP and SLC point to a \\textit{common} region in parameter space.  This region is characterized by several fermionic resonances nearly degenerate with the DM candidate (a spin-1 particle), which in turn is predicted to have a mass in the $300-500~{\\rm GeV}$ range.  We have also pointed out that although the GHU scenarios (with or without DM) present a sensitivity of order one percent or worse with respect to microscopic parameters (larger than what naive considerations would indicate), this sensitivity disappears once the top mass measurement is imposed.  Thus, the moderate sensitivity of order a percent, that is present in virtually every extension of the SM that has any relation to the physics of EWSB acquires a new twist: it indicates an intrinsic difficulty in accommodating a heavy top.  By contrast, once the top mass is fixed to a value of order the EW scale, the Higgs mass easily satisfies the LEP bound, while typically lying below $170~{\\rm GeV}$ or so.  Regarding the phenomenological implications of our construction, the low scale predicted for the new $\\Z_2-$odd sector and, in particular, new vector-like quarks, guarantees large production cross-sections. Due to the particular features of the NLOP spectrum, which is very degenerate with the LOP, the most common signature will be jets plus missing energy accompanied in some cases by the semileptonic decays of a pair of off-shell $W^\\ast$.  These are challenging signatures that will require dedicated analysis.  The large production cross-sections and the information coming from the easier channels in the $\\Z_2-$even sector should however be sufficient to guarantee a discovery of the DM sector at colliders.  We have also seen that DM direct detection is not very promising in present or near future experiments, due to the fact that couplings to light valence quarks are suppressed by intergenerational mixing, leading to very small cross-sections.    \\vspace{5mm} \\noindent %  %-----------------------------------------------------------------------"
    },
    "0801/0801.4456_arXiv.txt": {
        "abstract": "{Hyper-velocity stars move so fast that only a supermassive black hole (SMBH) seems to be capable to accelerate them. Hence the Galactic centre (GC) is their only suggested place of origin. Edelmann et al.~(2005) found the early B-type star HE\\,0437$-$5439 to be too short-lived to have reached its current position in the Galactic halo if ejected from the GC, except if being a blue straggler star. Its proximity to the Large Magellanic Cloud (LMC) suggested an origin from this galaxy.} {The chemical signatures of stars at the GC are significantly different from those in the LMC. Hence, an accurate measurement of the abundance pattern of HE\\,0437$-$5439 will yield a new tight constraint on the place of birth of this hyper-velocity star.} {High-resolution spectra obtained with UVES on the VLT are analysed using state-of-the-art non-LTE modelling techniques.} {We measured abundances of individual elements to very high accuracy in HE\\,0437$-$5439 as well as in two reference stars, from the LMC and the solar neighbourhood, respectively. The abundance pattern is not consistent at all with that observed in stars near the GC, ruling our an origin from the GC. However, there is a high degree of consistency with the LMC abundance pattern. Our abundance results cannot rule out an origin in the outskirts of the Galactic disk. However, we find the life time of HE\\,0437$-$5439 to be more than~three times shorter than the time of flight to the edge of the disk, rendering a Galactic origin unlikely.} {Only one SMBH is known to be present in Galaxy and none in the LMC. Hence the exclusion of an GC origin challenges the SMBH paradigm. We conclude that there must be other mechanism(s) to accelerate stars to hyper-velocity speed than the SMBH. We draw attention to dynamical ejection from dense massive clusters, that has recently been proposed by Gvaramadze et al.~(2008).} ",
        "introduction": "Stars moving faster than the Galactic escape velocity were~first predicted to exist by \\citet{hills88}. The first such hyper-velocity stars~(HVSs) were discovered serendipitously only recently \\citep[][E05: \\object{HE\\,0437$-$5439} at 723\\,km\\,s$^{-1}$ heliocentric radial velocity]{brown05,hirsch05,edelmann05}. A systematic search for such objects has resulted in the discovery of seven additional HVS up to now \\citep[see][]{brown07}.  \\cite{hills88} predicted that the tidal disruption of a binary by a supermassive black hole (SMBH) could lead to the ejection of stars with velocities exceeding the escape velocity of our Galaxy. The Galactic centre (GC) is the suspected place of origin of the HVSs as it hosts a SMBH.  Both, the moderate rotational velocity ($v\\sin i$\\,$=$\\,54\\,\\kms) and the chemical composition (consistent with solar abundances within a factor of a few from a LTE analysis of a spectrum with S/N\\,$\\approx$\\,5) were considered evidence for a main sequence nature of HE\\,0437$-$5439 (E05). This put the star at a distance of $\\approx$60\\,kpc.  Numerical kinematical experiments were carried out to trace the trajectory of HE\\,0437$-$5439 from the GC to its present location in the Galactic halo. However, its travel time (100\\,Myrs) was found to be much longer than its main sequence lifetime ($\\approx$\\,25--35\\,Myrs), rendering a GC origin unlikely, though possible if HE\\,0437$-$5439 would be a blue straggler star. Alternatively the star could have originated from the Large Magellanic Cloud (LMC) as it is much closer to this galaxy (18\\,kpc) than to the~GC.  An accurate determination of its elemental abundance pattern should allow to distinguish between an origin in the LMC or in the GC because their chemical compositions differ distinctively. HE\\,0437$-$5439 is by far the brightest HVS known \\citep[$V$\\,$=$\\,16.36,][]{bonanos08} and therefore suited best for a detailed analysis. Hence we obtained high-resolution, high S/N spectra of HE\\,0437$-$5439 at the ESO VLT with UVES and performed a quantitative spectral analysis using state-of-the-art non-LTE modelling techniques. ",
        "conclusions": "We have performed a quantitative spectral analysis of the HVS HE\\,0437$-$3954 and reference objects in the LMC and the solar neighbourhood. Use of state-of-the-art non-LTE modelling techniques allowed precise elemental abundances to be obtained.  From the results of the spectroscopic analysis we can rule out a GC origin for HE\\,0437$-$3954 because its abundance pattern indicates neither high metallicity nor $\\alpha$-enrichment. On the other hand a high degree of consistency with the LMC pattern is found when accounting for an abundance scatter within the LMC\\footnote{More comprehensive investigations than presently available, using improved analysis techniques on larger samples of young stars covering the entire region of the LMC are needed to draw definite conclusions on the degree of chemical (in)homogeneity in that galaxy.}. The observed abundance pattern could also be consistent with an origin from the outskirts of the Galactic disk, in the case HE\\,0437$-$3954 being a blue straggler. However, a comparison of the time-of-flight with the life time of the blue straggler practically rules out this scenario. Hence, HE\\,0437$-$3954 is unlikely of Galactic origin.  \\citet{edelmann05} suggested that HE\\,0437$-$3954 was ejected from the LMC. Our spectroscopic analysis lends strong support to this scenario. The abundance pattern is consistent with that of our LMC reference star. An ejection velocity of about 1000\\,{\\kms} is required for the star to have reached its position from the LMC. Accurate proper motion measurements are needed to finally confirm the LMC origin.  Because no SMBH is known to exist in the LMC, the SMBH slingshot scenario or other ejection models \\citep[e.g.][]{baumgardt06} invoking a SMBH are ruled out. Hence there must be an additional physical mechanism capable of accelerating a massive star to a space velocity of 1000\\,{\\kms}.  Two such alternative scenarios have been proposed. {\\sc i}) A close encounter of a binary with an intermediate-mass black hole (IMBH) more massive than 10$^3$\\,$M_{\\sun}$ has been proposed as a viable ejection mechanism by \\citet{GPZ07}, who suggest the populous dense clusters \\object{NGC\\,2004} or \\object{NGC\\,2100}\\, in the LMC could harbour such an~IMBH. {\\sc ii}) Dynamical ejection by interaction of massive binaries in the cores of dense clusters is plausible \\citep{GGPZ08}. This process is more likely to occur in the LMC than in the Galaxy because the LMC hosts a significant number of sufficiently massive and dense clusters. From the list of \\citet{mackay03}, eight young clusters can be identified to be of the right age to qualify as candidate place of birth of HE\\,0437$-$3954. Primary candidates are NGC\\,2100 and NGC\\,2004, followed by \\object{NGC\\,1850} and \\object{NGC\\,1847}.  In consequence, hyper-velocity stars are not such simple probes for the shape of the Galactic halo as suggested e.g. by \\citet{gnedin05}. The fundamental assumption of the exclusive origin of HVSs in the GC has to be dropped, as acceleration may occur in clusters throughout the Galactic disk as~well."
    },
    "0801/0801.1363_arXiv.txt": {
        "abstract": "We develop a framework for studying collective three-flavor neutrino oscillations based on the density matrix formalism. We show how techniques proven useful for collective two-flavor  neutrino oscillations such as corotating frames can be applied  readily to  three-flavor mixing. Applying two simple assumptions and the conservation of two ``lepton numbers'' we use this framework to demonstrate how the adiabatic/precession solution emerges. We illustrate with a numerical example how two stepwise spectral swaps appear naturally if the flavor evolution of the neutrino gas can be described by such a solution. For the special case where mu and tau flavor neutrinos are equally mixed and are produced with identical energy spectra and total numbers, we find that one of the spectral swaps in the three-flavor scenario agrees with that in the two-flavor scenario when appropriate mixing parameters are used. Using the corotating frame technique we show how the adiabatic/precession solution can obtain even in the presence of a dominant ordinary matter background. With this solution we can explain why neutrino spectral swapping can be sensitive to  deviations from maximal 23-mixing when the ``mu-tau'' matter term is significant. ",
        "introduction": "It has long been recognized that, in addition to the conventional Mikheyev-Smirnov-Wolfenstein (MSW) effect \\cite{Wolfenstein:1977ue,Wolfenstein:1979ni,Mikheyev:1985aa}, neutrino self-coupling can be important for neutrino flavor evolution when neutrino number densities are large \\cite{Fuller:1987aa,Notzold:1988kx,Pantaleone:1992xh,Sigl:1992fn,% Fuller:1992aa,Qian:1993dg,Samuel:1993uw,Kostelecky:1994dt,Pastor:2001iu,% Pastor:2002we,Balantekin:2004ug,Fuller:2005ae}. Recently two-flavor neutrino oscillations in the core-collapse supernova environment have been intensively investigated \\cite{Duan:2005cp,Duan:2006an,Duan:2006jv,Hannestad:2006nj,% Raffelt:2007yz,EstebanPretel:2007ec,Duan:2007mv,Duan:2007bt,Fogli:2007bk}. These studies show that supernova neutrinos can indeed experience collective flavor evolution because of neutrino self-coupling, even when the neutrino self-coupling is subdominant compared to the MSW potential \\cite{Duan:2005cp}.  An important result is that collective two-flavor neutrino oscillations can exhibit ``stepwise spectral swaps'' or ``spectral splits'' in the final neutrino energy spectra when the neutrino number density slowly decreases from a high value, where neutrinos experience synchronized oscillations \\cite{Pastor:2001iu}, towards zero (see, e.g., Ref.~\\cite{Duan:2006jv}). When this occurs, $\\nu_e$'s appear to swap their energy spectra with $\\nu_x$'s at energies below or above (depending on the neutrino mass hierarchy) a transition energy $\\Ec$ (where the superscript ``s'' can stand either for ``swapping point'' or ``splitting point''). Here $\\nu_x$ is some linear combination of $\\nu_\\mu$ and $\\nu_\\tau$. The phenomenon of spectral swapping is present in both the ``single-angle approximation'', where flavor evolution along various neutrino  trajectories is assumed to be the same as that along a representative trajectory (e.g., the radial trajectory), and  the ``multi-angle approximation'', where flavor evolution along different trajectories is independently followed \\cite{Duan:2006an}. For the inverted neutrino mass hierarchy, it also has been found that stepwise neutrino spectral swapping is essentially independent of the $2\\times2$ effective vacuum mixing angle when this angle is small (see, e.g., Ref.~\\cite{Duan:2007bt}).  Collective two-flavor neutrino oscillations are best understood with the help of neutrino flavor polarization vectors \\cite{Sigl:1992fn} or neutrino flavor isospins \\cite{Duan:2005cp}. Using the spin analogy, one can represent the flavor content of a neutrino mode by a spin vector in  flavor isospace. In this analogy, the matter effect is described as a spin-field coupling, and the neutrino self-coupling is described as spin-spin coupling. The phenomenon of spectral swapping is a result of collective precession of all neutrino flavor isospins with a common angular velocity $\\wpr$ at any given neutrino number density \\cite{Duan:2006an,Duan:2007mv}. This collective precession of neutrino flavor isospins is described the two-flavor adiabatic/precession solution \\cite{Raffelt:2007cb,Raffelt:2007xt}. In this solution, all neutrino flavor isospins stay aligned or antialigned with a total effective field in a reference frame that rotates with angular velocity $\\wpr$. Numerical simulations have shown that, in the supernova environment, neutrinos can first experience collective MSW-like flavor transformation (in which the MSW effect is enhanced by neutrino self-coupling) and then subsequently the adiabatic/precession solution \\cite{Duan:2007fw}.  In the real world, however, there are three active neutrino flavors. Some limited progress has been made on understanding collective three-flavor neutrino oscillations. The first fully-coupled simulation of three-flavor neutrino oscillations in the supernova environment showed a spectral swapping phenomenon similar to the two-flavor scenario except possibly with two swaps \\cite{Duan:2007sh}. Another simplified numerical study with a single neutrino energy bin \\cite{EstebanPretel:2007yq} reveals that collective neutrino oscillations can be sensitive to deviations from maximal 23-mixing when there is a dominant ``mu-tau'' term arising from a high order contribution from virtual $\\mu$'s and $\\tau$'s (e.g., Refs.~\\cite{Fuller:1987aa,Botella:1987aa,Roulet:1995qb}). Ref.~\\cite{Dasgupta:2007ws} has extended the neutrino flavor polarization vector notation to the three-flavor scenario and discussed the collective three-flavor oscillations as  ``factorization'' of two two-flavor oscillations.   In this paper we develop a framework for studying collective three-flavor neutrino oscillations based on the density matrix formalism. Using this framework we find a generalized three-flavor adiabatic/precession solution and show how stepwise spectral swapping can appear as a natural result of such a solution.  The rest of this paper is organized as follows. In Sec.~\\ref{sec:eom} we develop a framework centered around a $3\\times3$ reduced neutrino flavor density matrix. This density matrix is equivalent to  an 8-component neutrino flavor vector, a generalization of neutrino flavor isospin and similar to the Bloch vector used in Ref.~\\cite{Dasgupta:2007ws}. We show how techniques important in studying collective two-flavor oscillations such as corotating frames can be applied in the framework. In Sec.~\\ref{sec:sol} we  demonstrate how the three-flavor adiabatic/precession solution can be found using two simple assumptions and the conservation of two ``lepton numbers''. We  also illustrate with a numerical example how two spectral swaps can form from the adiabatic/precession solution when the total neutrino number density vanishes. In Sec.~\\ref{sec:matt} we  employ the corotating frame technique and show that the adiabatic/precession solution obtains even in the presence of a dominant matter background. In particular, we show that, in the presence of a large mu-tau term, neutrino spectral swapping becomes sensitive to deviations from maximal 23-mixing. In Sec.~\\ref{sec:conclusions} we  give our conclusions. ",
        "conclusions": "}  We have developed a framework for studying collective three-flavor neutrino oscillations. Important techniques in studying collective two-flavor oscillations such as corotating frames can be applied readily  to three-flavor scenarios in this framework. We have shown that the three-flavor adiabatic/precession solution obtains when both the precession ansatz and the adiabatic ansatz are satisfied. If the flavor evolution of a neutrino gas is described by the adiabatic/precession solution, the final neutrino energy spectra will exhibit the stepwise swapping phenomenon. We have shown that stepwise spectral swapping appears in a numerical example in which neutrinos are directly emitted from the neutrino sphere of a bare neutron star into vacuum. For this special example, because the neutrinos in $\\mu$ and $\\tau$ flavors are equally mixed and have identical energy spectra initially, the adiabatic/precession solutions for both the three-flavor and the two-flavor scenarios produce the same spectral swapping at the atmospheric neutrino mass-squared-difference scale. In more general cases, however, this may not be the case, and the full $3\\times3$ mixing framework should be employed.  Strictly speaking, the adiabatic/precession solution obtains  only when the $CP$-violating phase $\\delta=0$. This is because, if $\\delta\\neq0$, the unitary transformation matrix $U$ connecting the flavor states and vacuum mass eigenstates is not real, and \\begin{equation} U\\rho_\\bfp^*U^\\dagger\\neq(U\\rho_\\bfp U^\\dagger)^* \\text{ and } U\\bar\\rho_\\bfp^*U^\\dagger\\neq(U\\bar\\rho_\\bfp U^\\dagger)^*. \\end{equation} As a result, Eqs.~\\eqref{eq:eom-nu} and \\eqref{eq:eom-anu} become invalid in the vacuum mass basis, and all the following derivations are invalid. In practice, however, the adiabatic/precession solution may still be a good approximation even when $\\delta$ is large. This is because $\\theta_{13}\\simeq0$ and the transformation matrix $U$ is almost real.  In fact, numerical simulations in Ref.~\\cite{Duan:2007sh} show that, at least for the parameters employed in those simulations, varying the $CP$ phase $\\delta$ has little effect on the final neutrino energy spectra except for changing the relative mixing of $\\mu$ and $\\tau$ neutrino flavors. The possible effects of $CP$ violation in stellar collapse have been discussed in a different context (see, e.g., Ref.~\\cite{Balantekin:2007es}).  We also have  demonstrated that the adiabatic/precession solution obtains even in the presence of a dominant matter background and a large mu-tau term. When the matter term is much larger than the vacuum mixing term, the presence of the ordinary matter only reshuffles the neutrino states in which the spectral swapping  has the most dramatic manifestation. For the supernova environment, this means that the regime where collective neutrino oscillations occur is solely determined by the neutrino fluxes and does not depend sensitively on the matter density profile. This is in agreement with the previous analysis in two-flavor scenarios \\cite{Duan:2005cp}. The supernova neutrino signals observed on earth will depend of course  on the matter profile in the supernova. In part, this is because the MSW effect will modify the neutrino energy spectra subsequent to the collective oscillations discussed here. This dependence of neutrino signals on the matter profile can provide important information on the conditions deep in the supernova envelope \\cite{Schirato:2002tg,Duan:2007sh,Lunardini:2007vn}."
    },
    "0801/0801.1680_arXiv.txt": {
        "abstract": "{Rare types of variable star may give unique insight into short-lived stages of stellar evolution. The systematic monitoring of millions of stars and advanced light curve analysis techniques of microlensing surveys make them ideal for discovering also such rare variable stars. One example is the R Coronae Borealis (RCB) stars, a rare type of evolved carbon-rich supergiant.} {We have conducted a systematic search of the EROS-2 database for the Galactic catalogue Bulge and spiral arms to find Galactic RCB stars.} {The light curves of $\\sim$100 million stars, monitored for 6.7 years (from July 1996 to February 2003), have been analysed to search for the main signature of RCB stars, large and rapid drops in luminosity. Follow-up spectroscopy has been used to confirm the photometric candidates.} {We have discovered 14 new RCB stars, all in the direction of the Galactic Bulge, bringing the total number of confirmed Galactic RCB stars to about 51.} {After reddening correction, the colours and absolute magnitudes of at least 9 of the stars are similar to those of Magellanic RCB stars. This suggests that these stars are in fact located in the Galactic Bulge, making them the first RCB stars discovered in the Bulge. The localisation of the 5 remaining RCBs is more uncertain: 4 are either located behind the Bulge at an estimated maximum distance of 14 kpc or have an unusual thick circumstellar shell; the other is a DY Per RCB which may be located in the Bulge, even if it is fainter than the known Magellanic DY Per. From the small scale height found using the 9 new Bulge RCBs, $61<h^{RCB}_{Bulge}<246$ pc (95\\% C.L.), we conclude that the RCB stars follow a disk-like distribution inside the Bulge.} ",
        "introduction": "}  The R Coronae Borealis (RCB) stars constitute a rare type of hydrogen deficient, carbon-rich supergiant star. Only 38 examples are currently known in the Galaxy \\citep{1996PASP..108..225C, 2005AJ....130.2293Z}, 17 in the Large Magellanic Cloud \\citep{2001ApJ...554..298A} and 6 in the Small Magellanic Cloud \\citep{2004A&A...424..245T, 2005ApJ...631L.147K}. This rarity presumably indicates a brief phase of stellar evolution. A detailed review of their characteristics has been written by \\citet{1996PASP..108..225C}. RCBs exhibit spectacular, unpredictable, rapid declines in brightness (up to 9 magnitudes in optical wavelengths over a few days or weeks) thought to be due to the photosphere being obscured by newly formed dust clouds along the line of sight. As the dust clouds disperse, the original brightness gradually recovers in months. Several of these clouds have been directly observed in various directions around the RCB star RY Sgr, see \\citet{2004A&A...428L..13D, 2007A&A...466L...1L}.  Up to now, there has been no reliable distance estimate to any Galactic RCB star, so their absolute magnitudes is only known from the RCBs discovered in the Magellanic Clouds (MC). A direct relation between effective temperature and absolute magnitude was observed by \\citet{2001ApJ...554..298A}. Hipparcos gave lower limits to the distances to several of the brightest Galactic RCB stars \\citep{1997PASP..109.1089T, 1998PASA...15..179C}. Such limits imply that the RCB stars studied must be brighter than $M_V = -3$.  The true spatial distribution of RCB stars is still uncertain. This is due to the small number of known RCB stars, which is also biased toward higher Galactic latitudes as reddened stars in the Galactic plane can be missed in magnitude limited surveys \\citep{1990MNRAS.247...91L, 1990ASPC...11..566L}. Different views are found in the literature. First, \\citet{1985ApJS...58..661I} reports a scale height of $\\sim400$ pc for known RCB stars assuming $M_{Bol} = -5$. So RCB stars may be part of an old disk-like distribution. This idea is also supported by \\citet{2005AJ....130.2293Z}, who mention a possible thick-disk distribution. However, \\citet{1998PASA...15..179C} noted that the Hipparcos velocity dispersion of RCB stars is quite similar to those of other cool hydrogen-deficient carbon stars and extreme Helium (eHe) stars, suggesting that RCB stars might have a Bulge-like distribution.  Two major evolutionary scenarios are suggested to explain their origin : the Double Degenerate scenario (DD) and the final Helium Shell Flash (FF) scenario \\citep{1996ApJ...456..750I, 1990ASPC...11..549R}. The DD model involves the merger of a CO- and a He- white dwarf and was recently strongly supported by the observations of $^{18}$O over-abundance in seven H-deficient carbon and RCB stars, that is not expected in the FF model \\citep{2007ApJ...662.1220C}. The birthrates of RCB stars would then help us to better understand the rates of mergers for objects that could be Supernovae type Ia progenitors \\citep{1984ApJ...277..355W,2005ApJ...629..915B}. The FF model involves the expansion of a star, on the verge of becoming a white dwarf, to a supergiant size. This outburst phenomena has already been observed in three stars (Sakurai's object, V605 Aql and FG Sge), which transformed themselves into cool giant stars with spectral properties similar to those of RCB stars \\citep{1997AJ....114.2679C, 1999A&A...343..507A,1998ApJS..114..133G}. However, \\citet{2006ApJ...646L..69C} note that differences between FF star and RCB star light curve variations and abundance patterns (such as $^{12}$C$/^{13}$C ratio) indicate that the former are unlikely to be the evolutionary precursors of the majority of the latter.  One way to make progress in the comprehension of these stars is to enlarge the small sample of known objects. Microlensing surveys (OGLE, MACHO, MOA, EROS and Supermacho) are ideal for increasing the sample because they monitor millions of stars and thus could observe the characteristic RCB signature of rapid fading. The EROS-2 (Exp\\'erience de Recherche d'Objets Sombres) experiment is currently the survey that has monitored the largest sky area ($\\sim106$ deg$^2$ in our Galaxy) during a period of $\\sim$6.7 years. EROS can thus be used to significantly increase the number of known Galactic RCBs.  This article reports a sample of 14 new RCB stars. The photometric and spectroscopic data used are presented in Section~\\ref{sec_obs}. A discussion of the previously known RCB stars in and near the area monitored by EROS-2 is given in Section~\\ref{knownRCB}. Following the discovery of the first RCB stars in the SMC \\citep{2004A&A...424..245T}, we tried to increase our detection efficiency to prevent a selection effect due to reddening. The modified detection techniques, and the spectroscopic observations, are described in Section~\\ref{mining.sect}. Finally, the general characteristics of the new RCBs discovered are discussed in Sections~\\ref{sec_NewRCB} and \\ref{summary}.    \\begin{table*} \\caption{Known Galactic RCB stars located in the EROS-2 fields \\label{tab.KnownRCB.GenInfo}} \\medskip \\centering \\begin{tabular}{llllc} \\hline \\hline Star name & Coordinates ($J_{2000}$) & True MACHO Id & Situation on EROS-2 & Large variations\\\\ & & & reference image & seen by EROS-2?\\\\ \\hline \\object{V1783 Sgr}   &  18:04:49.74 -32:43:13.6  & & saturated, cg0842m & no\\\\ \\object{V739 Sgr}   &  18:13:10.54 -30:16:14.7  & \\object{MACHO-123.24377.7} & saturated, cg6245k & yes\\\\ \\object{V3795 Sgr}  &  18:13:23.58 -25:46:40.8  & \\object{MACHO-161.24445.6} & saturated, cg6273l & no\\\\ \\object{VZ Sgr}      &  18:15:08.58 -29:42:29.4  & \\object{MACHO-117.25166.4693} & saturated, cg0290n & yes \\\\ \\object{MACHO-135.27132.51} &  18:19:33.87 -28:35:57.8  & & too faint, cg0353k & yes\\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "\\label{summary}  Our search for new RCB stars in the EROS-2 Galactic photometry resulted in the discovery of 14 new RCB stars, one being a DY Per star. The total number of confirmed Galactic RCB stars is therefore now 51. No RCB was found in the 29 fields monitored in the spiral arms. Overall, we estimate that our detection of RCB stars should be $\\sim$60\\% complete for the majority of our survey.  MACHO-118.18666.100 has been misclassified as an RCB star. The new spectrum presented shows that it is an M giant star with light curve variations that resemble those of sequence-D variable stars.  If no interstellar reddening correction is applied, the new RCB sample would consist of the coolest RCB stars ever found. An extinction correction that assumes that the stars are located in the Bulge give intrinsic magnitudes and colours for 9 of the new RCBs that overlap those of the LMC and SMC. This indicates that they are indeed in the Galactic Bulge -- the first RCB stars identified in the Bulge.  Concerning the remaining classical RCBs, 4 have the lowest colour temperatures and faintest luminosities (EROS2-CG-RCB-5, -9, -12 and -14), assuming that the stars are located in the Bulge. However, these unusual properties may be explained either by an unusual thick circumstellar carbon shell or by a location beyond the Bulge at a maximum distance of $\\sim14$ kpc. The supplementary distance would thus give rise to an extra reddening. From the study of the Ca II triplet strength, we suggest that EROS2-CG-RCB-5, -9, and -14 have typical RCB temperatures ($T_{eff}\\sim5000$K) and that EROS2-CG-RCB-12 is the coolest of the classical RCBs found.   Finally, we note that the density of RCB stars seems to increase strongly as the galactic plane is approached. We measured a small scale height of $109^{+137}_{-48}$pc at the distance of the Galactic Centre using the 9 RCBs that are more likely to be located inside the Bulge. This small value is not consistent with measured Bulge scale height values and suggests a disk-like distribution inside the Bulge. Our derived distribution suggests that the RCB star density should be higher than one per deg$^2$ for $\\vert b \\vert <2\\degr$ towards the Galactic Bulge. Hence, we suggest that the next generation of infrared variable star surveys of the Bulge (e.g. VISTA\\footnote{Visible and Infrared Survey Telescope for Astronomy, URL: http://www.vista.ac.uk/}) should be able to discover many RCB stars near the Galactic Centre.  \\begin{figure*} \\includegraphics[scale=0.45]{./cg6075k_26836.eps} \\includegraphics[scale=0.45]{./Eros2-M135.27132.51.eps} \\includegraphics[scale=0.45]{./Eros2-V739Sgr.eps} \\includegraphics[scale=0.45]{./Eros2-VZSgr.eps} \\caption{EROS-2 light curves of MACHO-118.18666.100 (re-classified as a M giant star, see Sect.~\\ref{knownRCB}) and three known Galactic RCB stars, MACHO-135.27132.51, V739 Sgr and VZ Sgr. The dashed vertical lines indicate the 2MASS epochs. The photometry of the last three RCB stars has been reprocessed as these stars were either saturated or too faint on the original EROS-2 reference images.} \\label{lc_KnownRCB} \\end{figure*}  \\begin{figure*} \\includegraphics[scale=0.45]{./cg0030l_3740.eps} \\includegraphics[scale=0.45]{./cg0030m_154.eps} \\includegraphics[scale=0.45]{./cg0062l_18324.eps} \\includegraphics[scale=0.45]{./cg0711m_12518.eps} \\caption{Light curves of the 14 new RCBs star, one being a DY Per star (EROS2-RCB-CG-2). The arrows represent our detection limit. The dashed vertical lines indicate the 2MASS epochs. Top: $B_E - R_E$ colour vs. time; Middle: $R_E$ light curve ; Bottom: $B_E$ light curve.} \\label{lc} \\end{figure*}   \\begin{figure*} \\includegraphics[scale=0.45]{./cg0715n_14430.eps} \\includegraphics[scale=0.45]{./cg1024m_18795.eps} \\includegraphics[scale=0.45]{./cg1026l_13692.eps} \\includegraphics[scale=0.45]{./cg1066m_15039.eps} \\caption{Light curves of the new RCBs stars (continued).} \\end{figure*}  \\begin{figure*} \\includegraphics[scale=0.45]{./cg1127l_13076.eps} \\includegraphics[scale=0.45]{./cg1133k_13881.eps} \\includegraphics[scale=0.45]{./cg1187l_16577.eps} \\includegraphics[scale=0.45]{./cg1235m_14181.eps} \\caption{Light curves of the new RCBs stars (continued).} \\end{figure*}  \\begin{figure*} \\includegraphics[scale=0.45]{./cg6110k_17996.eps} \\includegraphics[scale=0.45]{./cg6267k_28863.eps} \\caption{Light curves of the new RCBs stars (continued).} \\label{lc_end} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[scale=0.82]{./All_Charts.eps} \\caption{Charts of the new Galactic RCB stars (2'x2'). North is up, East is to the left.} \\label{charts} \\end{figure*}"
    },
    "0801/0801.1649_arXiv.txt": {
        "abstract": "We have simulated the 3D emergence and interaction of two twisted flux tubes, which rise from the interior into the outer atmosphere of the Sun. We present evidence for the multiple formation and eruption of flux ropes inside the emerging flux systems and hot arcade-like structures in between them. Their formation is due to internal reconnection, occurring between oppositely directed, highly stretched and sheared fieldlines at photospheric heights. Most of the eruptions escape into the corona, but some are confined and fade away without leaving the low atmosphere. As these flux ropes erupt, new reconnected fieldlines accumulate around the main axis of the initial magnetic flux systems. We also show the complex 3D fieldline geometry and the structure of the multiple current sheets, which form as a result of the reconnection between the emerging flux systems. ",
        "introduction": "Partial eruption of a flux rope can be produced by favourable, \\textit{imposed} shearing motions at the photosphere. This shearing can occur as sunspots with opposite polarity move apart along the neutral line of an active region. Then, as the rising magnetic system expands, the sheared magnetic fieldlines are stretched and undergo internal reconnection, at low heights, which leads to partial eruption of magnetic flux \\citep{Manchester+ea04}. This is consistent with the `tether-cutting' mechanism (\\citealt{Sturrock89, Moore01}), in which, effective reconnection of magnetic fieldlines at the core of  a magnetic bipole releases the energy stored in the sheared field, resulting in an eruption. An important factor in determining whether an emerging flux rope will erupt is the amount of twist of the fieldlines around the central axis of the tube (\\citealt{Murray06}). If the winding of the fieldlines is high enough, a kink instability occurs that eventually leads to either a confined or an ejective eruption (\\citealt{Torok05, Gibson+ea06}).  Observations (\\citealt{Sterling04, Sterling05, Jiang07}) have shown active-region filament eruptions when new flux emerges near the location of a preexisting emerging flux region. It is likely that the new emergence is important in the eruption onset. Several numerical models, requiring multipolar magnetic configurations for eruptions, have reproduced the above and similar-like observations. In the breakout model (\\citealt{Antiochos98}), a sheared bipole rises into an overlying bipole and \\lq external\\rq\\ reconnection occurs at a magnetic null point at their interface. When the reconnection becomes fast enough the rising bipole breaks through the ambient field and erupts.  Another possibility, shown in magnetofrictional simulations, is that flux ropes are formed and erupt due to slow diffusion of decaying flux from active regions, which interact through converging motions on the photosphere (\\citealt{Duncan06, Adrian07}). Also, \\citet{Chen00} simulated the interaction between a flux system with a detached flux rope and an emerging flux system that appears to one side of the primary field. Reconnection between the flux systems leads to a vertical current sheet below the detached flux rope, which eventually experiences a fast ejection.  In this Letter, we consider the 3D extension of the model by \\citet{Archontis07} (hereafter called AHB model) in which one (leading) flux tube rises through a stratified atmosphere to create a coronal field for the second (following) tube to rise into. The 3D model simulates: 1) the intricate reconnection of the two flux systems 2) the multiple eruptions of flux ropes, which are initiated by the new emergence of flux and 3) the formation of current sheets and hot arcade structures. ",
        "conclusions": "The initial evolution of the system is similar to the evolution in the AHB model. When the leading tube reaches the photosphere a bipolar region is formed with a north-south orientation due to the initial strong twist of the fieldlines. Eventually, the bipolar region moves towards an east-west direction as more internal magnetic layers rise to the photosphere and the inclination of the anchored flanks of the tube becomes more vertical. After $t=35$, the leading tube rises into the non-magnetized corona where it expands and \\textit{creates an ambient field} for the following tube to rise into. The direction of the fieldlines and the field strength vary away from the center of the expanding volume of the leading system and, thus, the newly formed ambient field is more complex than in the previous simulations of the interaction between an emerging field and a uniform coronal field.  As the newly emerged bipolar magnetic field is moving in opposite directions along the neutral line, shearing of the field occurs so that the magnetic field lines lose their strong azimuthal nature and run nearly parallel to the neutral line. Observations have shown that most bipolar sheared magnetic fields in active regions can produce ejective explosions such as flares and CMEs \\citep{Falconer01}. Also, in our simulations the twisted and sheared magnetic field adopts an overall \\textit{sigmoidal shape} with two oppositely curved \\textit{elbow-like} regions at the two (east and west) ends of the neutral line (as illustrated in Fig.~\\ref{fig1}, left panel). The neutral line is stretched along the middle of the sigmoidal structure while the elbows are found on either side. During the evolution of the system, the curved region of the elbows consist of expanding loops of fieldlines that rise in opposite (north-east and south-west) directions. The parts of the elbows close to the middle of the sigmoid, shear past each other and are ready to reconnect if they come into contact. This magnetic field configuration with the elbow-like structures has been reported in observations of solar eruptive events by \\citet{Moore01}.  When the second flux system emerges  at $t\\approx 40$, a thin, curved, sheet-like current structure (Fig.~\\ref{fig1}, middle panel) is formed just above the top layers of the following tube and parallel to the interface of the magnetic flux systems, along the y-direction. An interesting result is that already at this early stage of the interaction, \\textit{asymmetry} is introduced into the system (Fig.~\\ref{fig1}, left panel) by the 3D nature of the fan-like expansion of the field. On the west side of the emerging regions, the fieldlines of the rising systems make a contact angle which is larger than on the east side, leading eventually to a higher current concentration and more efficient reconnection at the west side of their interface.  In a similar manner to the 2.5D AHB model, soon after the build up of the interface current concentration \\textit{plasmoid-like} structures develop inside the sheet. Due to the asymmetry, they are formed firstly towards the west side of the current sheet. The cool and dense plasmoids grow and are ejected upwards and sideways, as in the AHB model, allowing reconnection in the diffusion regime to occur at a \\textit{faster rate}.  Two new flux systems are formed by the reconnected fieldlines (Fig.~\\ref{fig1}, middle panel). At the \\textit{upper end} of the current sheet, reconnection forms an overlying field that joins the two emerging systems externally, mainly from the north-west side of the leading system to the south-east side of the following system while it spreads out mostly above the west sides of the emerging systems. At the start of the reconnection, the overlying field links the west-side flank of the leading tube with the east-side flank of the following tube. At the \\textit{lower end} of the current sheet, reconnection forms a series of magnetic loops in between the two emerging systems. The configuration of the low-lying magnetic loops is an \\textit{arcade-like} structure joining the leading and following tubes and is reminiscent of \\textit{post-flare loops}. The arcade is stretched and curved along the y-direction and grows in size as reconnection proceeds and more fieldlines are added to the top of the arcade. The temperature here could be as high as $2$ MK. The arcade brightens when high velocity reconnection outflows, which carry hot material ejected from the current sheet, compress the plasma at the upper part of the arcade and increase the temperature.  During this time, due to the emergence of the leading tube the inner fieldlines are sheared horizontally and stretched vertically. Hence, oppositely directed fieldlines, which are rooted to the sigmoidal neutral line, are forced into contact. This internal reconnection of fieldlines is similar to the `tether-cutting' mechanism \\citep{Moore92}. In experiments with a single emerging tube, with the same initial properties, we find that internal `tether-cutting' reconnection occurs at $t\\approx 180$. In the experiment with two emerging bipoles it occurs much  earlier at $t\\approx 135$.    In the model with the two tubes, reconnection occurs first (externally) at X-type points at the ends of the interface current sheet. At the lower end of the current sheet, the hot reconnected fieldlines form an arcade structure. The magnetic pressure inside this new arcade increases as reconnection proceeds and eventually exceeds the magnetic pressure of the leading flux system. Thus, the arcade expands sideways and pushes the south side of the leading system towards its neutral line. It is the combination of shearing along the neutral line, stretching due to expansion and the increasing magnetic pressure force of the arcade structure that leads to internal reconnection of the fieldlines just above the neutral line of the leading emerging system.  As the two emerging flux tubes continue to rise and interact, internal reconnection occurs at \\textit{many positions along the neutral line} of the leading magnetic flux system. Due to internal reconnection, \\textit{flux-rope-like structures} are formed inside the magnetized volume of the leading tube. Eventually, the ropes rise and erupt into the outer atmosphere. The eruption is due to the vigorous sideway pressure by the fast expanding arcade and the release of the magnetic tension force of the overlying field, after the ejection of the plasmoids out of the interface current sheet.  A striking result is that the eruptions produced in our simulations \\textit{do not appear simultaneously}. Also, multiple eruptions can be triggered from the same location inside the expanding volume of the leading system. As an example, the first formation and eruption of a \\textit{flux rope} occurs on the south-west part of the neutral line. This part of the leading tube is closer, due to the asymmetry mentioned above, to the following tube. Thus, the dynamical rearrangment of the magnetic field due to reconnection occurs faster at this region of the leading emerging tube. In general, the eruptions occured in our experiment follow the direction of the overal 3D expansion of the emerging systems, which for the south-west side is preferentially the direction towards the following tube and for the north-east side is the opposite direction. Thus, the eruption of the first flux rope is directed towards the south side (Fig.~\\ref{fig1}, right panel). It is a confined eruption, being trapped between the envelope field of the leading tube and the expanding field of the upcoming tube that halts the further rise of the flux rope. Other eruptions, which occur at later stages of the evolution of the emergence, are more dynamic and seem to fully escape from the corona, in a \\textit{CME-like} manner. These eruptions are mainly directed towards the north side of the leading magnetic flux system.  The time evolution of the flux ropes formed \\textit{inside} the leading region, can be followed by locating the center (or {\\it O-point}) of the ropes at various  $x\\, z$ cross-sections (different values of $y$) as functions of time. The height-time relation for 6 values of $y$ is shown in the top panel of Fig.~\\ref{fig2}. The bottom panel is the projection of their motion onto an $x\\, z$ slice. The confined eruption of the first flux rope is shown by the lines at $y=-20$ and $y=-10$. The center of the rope for $y=-20$ is lower than at $y=-10$ at all times.   There are \\textit{two phases} during the eruption of the ropes. In the first, the flux ropes are moving with moderate speeds, while in the second they are accelerated towards the outer atmosphere. These curves are qualitatively similar to the time profiles of filament heights shown in the observations by \\citet{Sterling04, Sterling05}. We also estimate the temperature and current density in the area below the central part of the erupting flux ropes where a vertical current sheet is formed due to internal reconnection. We find that there is a slow increase in the first phase, which is followed by an intensity increase over the period of the fast eruption. The peak of the temperature and current density occurs when the fast rise is well underway. Since both quantities follow the same time evolution, it is likely that the reason for the change in the rise phase of the eruption is that reconnection is slower at the beginning but becomes faster, and operates with a higher growth rate, as the ropes rise. These results are also consistent with the observations mentioned above in which soft X-ray emission is found to occur during the slow-rise phase of a filament's eruption and hard X-ray bursts occured after the start of the fast eruption.  As shown above, the first formation of a flux rope and its eruption occurs in the south-west part of the leading tube. Figure ~\\ref{fig2} (bottom panel) shows that a secondary formation and eruption occurs at the same location (y=-10, dashed-triple dotted line), starting at t=145. Thus, \\textit{at least two eruptions are triggered from the same location} within the leading emerging system until the end of the simulation. Further eruptions may be triggered inside the expanding magnetized volumes since the system of the interacting bipoles does not reach an equilibrium but evolves dynamically. As the flux ropes erupt, they grow in radial extent as additional fieldlines reconnect at the vertical current sheet beneath them. The magnetized plasma inside the erupting ropes is a mixture of cold and hot plasma. The cold and dense plasma comes from the low atmosphere and is carried upwards by the tether-cutting reconnection forming concave upwards fieldlines. The hot plasma (up to $2$ MK) is emitted inside the envelope of the rope by the reconnection jet, which blasts from the upper end of the vertical current sheet below the rope. The velocity at the center of the rope is $\\approx $10-15$ Km/sec$. However, the maximum upward velocity inside the complete volume of the rope reaches values up to $150$ Km/sec, mainly due to the emission of the reconnection jet. The rising motion of the flux rope can result in a CME-like eruption.  Figure ~\\ref{fig3} (left panel) shows the temperature distribution in a vertical plane and magnetic fieldlines, which are traced from near the axis of the flux rope. The cold plasma is mostly concentrated near the axis of the flux rope while the hot plasma is overlying the axis on the right side of the eruption. The fieldlines near the axis of the rope come from \\textit{both of the emerging flux systems}. More precisely, fieldlines from the following tube joins the leading tube, through the arcade structure between them before forming the erupting flux rope. These fieldlines undergo multiple reconnection events along their lengths and, thus, may end up in either of the flanks at the east-side boundary. The middle panel in Fig.~\\ref{fig3} shows isosurfaces of current density, which are overplotted on the same set of fieldlines. A striking result is that long and twisted sheets are found close to the core of the erupting flux rope. In addition, arch-like current sheets that are less twisted and have different orientation to each other, are formed in the region between the emerging systems (see `A' and `B' in the figure). These structures, with enhanced current density, may be quasi seperatrix layers. The right panel in Figure ~\\ref{fig3} is a blow-up of the region at the vertical current sheet below the erupting flux rope shown in the left panel. The fieldlines shown in the figure have been traced from the lower end of the vertical sheet. These are reconnected fieldlines that surround the leading flux tube's original axis. Thus, the field below the vertical current sheet forms a tightly wound flux tube with its axis lying parallel to the neutral line of the leading bipolar region. The temperature rises up to $0.2$ MK at the bottom of the current sheet, as hot reconnected plasma moves downward from the sheet. These fieldlines are also linked: a) to fieldlines that envelope the erupting flux rope from the north side (two fieldlines on left side of figure) and b) to fieldlines that form the hot arcade structure between the emerging flux systems.  Our 3D simulations show that the interaction of emerging flux tubes may lead to multiple eruptions of flux ropes. A network of twisted current sheets, with an overall S-shaped configuraton, and hot cusp-like, arcade structures are also formed during the evolution of the system."
    },
    "0801/0801.3150_arXiv.txt": {
        "abstract": "% We present the results of the first multiwavelength observing campaign on the high-mass X-ray binary \\lsi\\ comprising observations at the TeV regime with the MAGIC telescope, along with X-ray observations with {\\it Chandra}, and radio interferometric observations with the MERLIN, EVN and VLBA arrays, in October and November 2006. From our MERLIN observations, we can exclude the existence of large scale ($\\sim 100$ mas) persistent radio-jets. Our 5.0 GHz VLBA observations display morphological similarities to previous 8.4 GHz VLBA observations carried out at the same orbital phase, suggesting a high level of periodicity and stability of the processes behind the radio emission. This makes it unlikely that variability of the radio emission is due to the interaction of an outflow with variable wind clumps. If the radio emission is produced by a milliarcsecond scale jet, it should also show a stable, periodic behavior. It is then difficult to reconcile the absence of a large scale jet ($\\sim$100 mas) in our observations with the evidence of a persistent relativistic jet reported previously. We find a possible hint of temporal correlation between the X-ray and TeV emissions and evidence for radio/TeV non-correlation, which points to the existence of one population of particles producing the radio emission and a different one producing the X-ray and TeV emissions. Finally, we present a quasi-simultaneous energy spectrum including radio, X-ray and TeV bands. ",
        "introduction": "%  \\lsi\\ is a high-mass X-ray binary consisting of a low-mass [$M\\sim (1-4)\\,\\msun$] compact object orbiting around an early type B0\\,Ve star along an eccentric ($e=0.7$) orbit \\citep[and references therein]{casares05}. The modulation of both, the radio \\citep{gre78,gre02} and X-ray \\citep{tay96,paredes97} emissions, display a period of $P_\\textrm{orb} =26.496$~d, attributed to the orbital motion. The position of the maximum of the radio emission along the orbit, as well as its intensity, are modulated with a superorbital period of $P_\\textrm{sup} = 1667 \\pm 8$ days \\citep{gre02}. \\lsi\\ is positionally coincident with an EGRET $\\gamma$-ray source \\citep{kni97}. Moreover, variable emission at TeV energies has been recently detected with the MAGIC telescope \\citep{alb06}. These authors found that the peak flux at TeV energies occurs at orbital phase $\\phi_\\textrm{orb} \\approx 0.65$, while no high-energy emission is detected around periastron passage ($\\phi_\\textrm{orb} \\approx 0.23$). From $\\sim$50 mas resolution radio images of \\lsi\\ obtained with MERLIN, \\cite{mas04} suggested the existence of a precessing relativistic ($\\beta\\approx0.6$) jet up to angular scales of $\\sim$0.1 arcsec, which led them to interpret \\lsi\\ within the framework of the microquasar scenario \\citep{valenti06,gustavo07}. However, recent VLBA imaging obtained by \\cite{dha06} over a full orbit of \\lsi\\ has shown the radio emission to come from angular scales smaller than about 7~mas (projected size 14 AU at an assumed distance of 2~kpc). This radio emission appeared cometary-like, interpreted to be pointing away from the high mass star within a particular scenario, and no relativistic motion, nor halos, nor larger-scale structures were detected at any phase of the orbit. Based on these findings, \\cite{dha06} concluded that the radio and TeV emissions from \\lsi\\ are originated by the interaction of the wind of a young pulsar with that of the stellar companion \\citep{maraschi81,dub06}.  In this article, we present and discuss the results of a multi-wavelength campaign including radio, X-ray, and TeV $\\gamma$-ray observations on \\lsi, aimed at shedding light on the physical processes going on in the system, as well as yielding useful input for a detailed, time-dependent modeling of this relevant system. For this, we study the correlations between the different detected emissions in terms of their morphological, temporal and spectral features, both in the intra-day and day-to-day timescales. ",
        "conclusions": "%  The comparison between VHE $\\gamma$-ray and radio data during the October campaign (see Figure~\\ref{fig:obs}) shows a detection at TeV energies for $\\phi_\\textrm{orb} = 0.66$ with a flux level of $\\sim 15\\%$ of the Crab Nebula flux, during a period when the radio emission is constant at 35 mJy. During the November campaign, however, there is no detection at TeV energies, while the radio data show a peak flux twice as high as in October. \\citet{alb06} reported on radio and TeV peaks detected almost simultaneously, while for our campaing we see the TeV peak for a flat and low radio flux and a radio peak for no significant TeV emission. Therefore, we exclude a general TeV-radio correlation. A plausible explanation is that the emissions are produced by different particle populations. On the other hand, the detections at X-ray and TeV energies, both at particularly high flux values and one day apart, might point to a correlation between X-ray and TeV fluxes, and hence to the fact that both radiations are produced by the same population of particles. Yet, our data are too scarce to make any solid conclusion at present.  The results obtained by radio imaging at different angular scales show that the size of the radio emitting region of \\lsi\\ is constrained below $\\sim$6~mas ($\\sim$12 projected AU), and the presence of persistent jets above this scale is therefore excluded.  We resolve a radio-emitting region at 5.0 GHz with VLBA, extending east-southeast from the brighter, unresolved emitting core. The outflow velocity implied by our observations is at least of $\\sim$0.1~c, in agreement to what previously suggested by \\cite{che06}, for an over-pressured pulsar wind, although recent hydrodynamical simulations \\citep{bogo08} show that the shocked pulsar wind could be relativistic.  The comparison of our VLBA image from 25 October with an image at 8.4 GHz obtained at a similar orbital phase but 10 orbital cycles apart shows a high level of similarity between them. It is worth noting that \\cite{dha06} obtained high-resolution radio images of the source along a complete orbital cycle, and found an extended feature whose orientation with respect to the central core varied by $360^\\circ$ along the orbit. Therefore, the similarity in the morphology and orientation between both images suggests a significant level of periodicity and stability of the physical processes involved in the radio emission. This result points to the fact that the extended radio emission is produced by the interaction of steady flows (from a relativistic pulsar wind, jet and/or stellar wind) rather than by the interaction of such an outflow with wind clumps. We note that if the radio emission is produced by a milliarcsecond scale jet, the required stability and periodic behavior of such a jet is difficult to reconcile with the non-persistent nature of a large scale ($\\sim$100 mas), putative relativistic jet, as deduced from our MERLIN observations in combination with those obtained by \\cite{mas04}.   Finally, we combine our X-ray data from 25 October, the TeV data from 26 October and the average VLBA in both nights to produce a quasi-simultaneous multi-wavelength spectrum, including radio, X-ray and VHE $\\gamma$-ray observations, which we display in Figure~\\ref{sed}. However, the simultaneous data cover only a small range of the orbital phase of \\lsi. Given the high variability of the physical conditions of the system along the orbit, more simultaneous, multiwavelength data, and particularly involving longer exposure times, orbital phase coverage and redundancy, will shed further light in our understanding of this peculiar object."
    },
    "0801/0801.4546_arXiv.txt": {
        "abstract": "The Roma International Conference on Astroparticle Physics covered gamma-ray astronomy, air shower experiments and neutrino astronomy on three successive days.  I organize my brief summary comments into four topics that cut across these three techniques. They are detector calibration, galactic sources, extra-galactic sources and cosmology. ",
        "introduction": "A main theme of the conference was the multi-messenger approach to the origin of cosmic rays.  The conference had a local flavor that illustrates the strong Italian contributions to major experimental activities in all three fields, gamma-ray astronomy, cosmic-ray experiments and neutrino astronomy.  In these brief remarks I will point out some common techniques and approaches and focus on a few important questions being addressed by current experiments in particle astrophysics.  I make no attempt to give a balanced and systematic review of the field here~\\cite{ICRC2007}. ",
        "conclusions": ""
    },
    "0801/0801.4894_arXiv.txt": {
        "abstract": "We present a new public archive of light-motion curves in Sloan Digital Sky Survey (SDSS) Stripe 82, covering 99$^{\\circ}$ in right ascension from $\\alpha = 20.7^{\\mbox{\\small h}}$ to 3.3$^{\\mbox{\\small h}}$ and spanning 2\\fdg52 in declination from $\\delta = -$1\\fdg26 to 1\\fdg26, for a total sky area of $\\sim$249~deg$^{2}$. Stripe 82 has been repeatedly monitored in the $u$, $g$, $r$, $i$ and $z$ bands over a seven-year baseline. Objects are cross-matched between runs, taking into account the effects of any proper motion. The resulting catalogue contains almost 4~million light-motion curves of stellar objects and galaxies. The photometry are recalibrated to correct for varying photometric zeropoints, achieving $\\sim$20~mmag and $\\sim$30~mmag root-mean-square (RMS) accuracy down to 18~mag in the $g$, $r$, $i$ and $z$ bands for point sources and extended sources, respectively. The astrometry are recalibrated to correct for inherent systematic errors in the SDSS astrometric solutions, achieving $\\sim$32~mas and $\\sim$35~mas RMS accuracy down to 18~mag for point sources and extended sources, respectively.  For each light-motion curve, 229 photometric and astrometric quantities are derived and stored in a higher-level catalogue. On the photometric side, these include mean exponential and PSF magnitudes along with uncertainties, RMS scatter, $\\chi^2$ per degree of freedom, various magnitude distribution percentiles, object type (stellar or galaxy), and eclipse, Stetson and Vidrih variability indices. On the astrometric side, these quantities include mean positions, proper motions as well as their uncertainties and $\\chi^2$ per degree of freedom. The here presented light-motion curve catalogue is complete down to $r~\\sim~21.5$ and is at present the deepest large-area photometric and astrometric variability catalogue available. ",
        "introduction": "\\begin{table*} \\centering \\caption{The list of SDSS imaging runs included in the LMCC. All runs were processed with version 40 of the SDSS {\\tt frames} pipeline except those runs marked with an asterisk, which were processed with version 41. Runs 4203 and 5823 (underlined) are the reference runs used for the astrometric calibrations (see Section~2.4).} {\\scriptsize \\begin{tabular}{@{}llll} \\hline Year & Month & North Strip Runs & South Strip Runs \\\\ \\hline 1998 & Sep & 94 & 125 \\\\ & & & \\\\ 1999 & Oct & 1033 & 1056 \\\\ & & & \\\\ 2000 & Sep & 1752 & - \\\\ 2000 & Oct & - & 1755 \\\\ 2000 & Nov & - & 1894 \\\\ & & & \\\\ 2001 & Jun & 2385 & - \\\\ 2001 & Sep & 2570, 2578, 2589 & 2579, 2583, 2585 \\\\ 2001 & Oct & 2649, 2650, 2659, 2662, 2677 & - \\\\ 2001 & Nov & 2700, 2708, 2728, 2738  & 2709 \\\\ 2001 & Dec & 2768, 2820 & - \\\\ 2002 & Jan & 2855, 2861, 2873 & 2886 \\\\ & & & \\\\ 2002 & Sep & - & 3325$^{*}$ \\\\ 2002 & Oct & 3362, 3384, 3437 & 3355, 3360$^{*}$, 3388, 3427, 3430, 3434, 3438 \\\\ 2002 & Nov & 3461 & 3460, 3465 \\\\ & & & \\\\ 2003 & Sep & 4128, 4153, 4157 & 4136, 4145 \\\\ 2003 & Oct & 4184, 4188, 4198, 4207 & 4187, 4192, \\underline{4203} \\\\ 2003 & Nov & 4253 & 4247, 4263, 4288 \\\\ & & & \\\\ 2004 & Aug & 4797 & - \\\\ 2004 & Sep & 4858 & - \\\\ 2004 & Oct & 4868, 4874, 4895, 4905, 4917 & - \\\\ 2004 & Nov & 4933, 4948 & 4930 \\\\ 2004 & Dec & - & 5042, 5052 \\\\ & & & \\\\ 2005 & Sep & 5566, 5603, 5610, 5622, 5633, 5642, 5658 & 5582, 5597, 5607, 5619, 5628, 5637, 5646, 5666 \\\\ 2005 & Oct & 5709, 5719, 5731, 5744, 5759, 5765, 5770, 5777, 5781, 5792, 5800 & 5675, 5681, 5713, 5729, 5730, 5732, 5745, 5760, 5763, 5771, 5776, 5782, 5786, 5797 \\\\ 2005 & Nov & 5813, \\underline{5823}, 5842, 5865, 5866, 5872, 5878, 5898, 5902, 5918 & 5807, 5820, 5836, 5847, 5853, 5870, 5871, 5882, 5889, 5895, 5905, 5924 \\\\ \\hline \\end{tabular}} \\label{tab:runs} \\end{table*}  One meaning of the verb ``to vary'' is to change in amount or level, especially from one occasion to another. In astronomy, there are many types of temporal variability. Stars may change in brightness, in which case they are termed ``variables'', or they may change position on the sky, in which case they have a proper motion. Over a sufficiently long period of time, the shape of constellations change. Galaxies exhibit variability across the electromagnetic spectrum since their emission is made up of the radiation from billions of sources, although their most obvious source of variation comes from active galactic nuclei or supernovae. Transient events, such as gamma-ray bursts, or microlensing events, are instrinsically variable. Solar system objects from planets to asteroids drift slowly across the sky, waxing and waning on various timescales. In fact, at some level, everything in the sky is variable.  The introduction of CCD detectors to astronomy greatly enhanced the ability to conduct variability surveys. The extension of CCD cameras to mosaic and wide-field formats along with the exponential progression of computing power have allowed the subsequent development of more ambitious surveys reaching to deeper magnitudes, higher cadences and larger sky areas. For more details we direct the reader to \\citet{bec2004} who present a clear summary of modern variability surveys. In this work we concentrate on optical photometric and astrometric variability (hence ``light and motion'') over a $\\sim$249~deg$^{2}$ patch of sky.  Large sky surveys such as the Sloan Digital Sky Survey~(SDSS; \\citealt{yor2000}) have in many ways revolutionised our knowledge of the Universe. SDSS has imaged approximately a quarter of the sky in five photometric wave bands. The exploitation of this impressive dataset has resulted in hundreds of publications covering a wide range of astronomical topics, from the structure of the Milky Way to the mapping of a large fraction of the Universe. The bulk of this data, however, contain only single measurements of objects from the north Galactic cap with no information on possible photometric variability or astrometric motion. Substantial efforts have been made by \\citet{mun2004} (see also \\citet{gou2004}) to measure proper motions by matching SDSS data from the north Galactic cap with the USNO-B catalogue (\\citealt{mon2003}). The resultant proper motion catalogue is 90\\% complete down to $g~=~19.7$ with the magnitude limit being set by the USNO-B catalogue faint magnitude limits.  One of the primary goals of the SDSS is the study of the variable sky (\\citealt{ade2007}) of which our knowledge is still very incomplete (\\citealt{pac2000}). To this end, the SDSS has repeatedly imaged a 300 square degree area, the so called Stripe 82, during the later half of each year since 1998. In 2005, the SDSS-II Supernova Survey (\\citealt{fri2008}) started with the aim of detecting Type-I supernovae in Stripe 82, greatly improving the cadence of measurements within the stripe. By averaging a subset of the repeated observations of unresolved sources in Stripe 82, \\citet{ive2007} built a standard star catalogue containing $\\sim$1 million nonvariable sources with $r$ band magnitudes in the range 14-22, by far the deepest and most numerous set of photometric standards available. Using these same multi-epoch photometric data, \\citet{ses2007} analysed the photometric variability for $\\sim$1.4~million unresolved sources in the stripe, drawing interesting conclusions on the spatial distribution of RR Lyrae stars and the variability of quasars.  Here we present a new public archive of light-motion curves in SDSS Stripe 82. The archive has been constructed from the set of high-precision multi-epoch photometric and astrometric measurements made in the stripe since the first SDSS runs in 1998 until the end of 2005. In constructing the catalogue, we only use measurements of objects that are cleanly detected in individual SDSS runs. The catalogue contains almost 4~million objects, galaxies and stars, and is complete down to magnitude 21.5 in $u$, $g$, $r$ and $i$, and to magnitude 20.5 in $z$. Each object has its proper motion calculated based only on the multi-epoch SDSS J2000 astrometric measurements. The catalogue reaches almost two magnitudes deeper than the SDSS/USNO-B catalogue, making it the deepest large-area photometric and astrometric catalogue available.  The catalogue comes in two flavours, the Light-Motion Curve Catalogue (LMCC), which contains the set of individual light-motion curves, where measured quantities for each object are listed as a function of wave band and epoch, and the Higher-Level Catalogue (HLC), which presents a set of derived quantities for each light-motion curve. For many purposes, it is more convenient to work with the HLC, especially for selecting subsets of interesting objects. The construction, calibration and format of the LMCC are discussed in Section~2, and the HLC is described in Section~3. Between the two subcatalogues, there is all the necessary information available to explore the photometric and astrometric variability of $\\sim$249 square degrees of equatorial sky. In Section~4, we investigate the quality of the photometric and astrometric properties of our catalogues by comparing them against suitable external catalogues, and we analyse the behaviour of our proper motion uncertainties. ",
        "conclusions": "The Light-Motion Curve Catalogue (LMCC) contains almost 4~million light-motion curves for stars and galaxies covering $\\sim$249~deg$^{2}$ in the SDSS Stripe 82. A light-motion curve provides wave band and time dependent photometric and astrometric quantities, where $\\sim$76 per cent of light-motion curves in the LMCC have at least 20 epochs of measurements. The LMCC is complete to magnitude 21.5 in $u$, $g$, $r$ and $i$, and to magnitude 20.5 in $z$, making it the deepest large-area catalogue of its kind. The photometric RMS accuracy for stars is $\\sim$20~mmag at $r\\sim$18~mag and for galaxies it is $\\sim$30~mmag at $r\\sim$18~mag. In both the RA and Dec coordinates, an RMS accuracy of $\\sim$32~mas and $\\sim$35~mas at $r\\sim$18~mag is achieved for pre-2005 data for stars and galaxies, respectively, and an RMS accuracy of $\\sim$35~mas and $\\sim$46~mas at $r\\sim$18~mag is achieved for 2005 data for stars and galaxies, respectively.  The LMCC is thus an ideal tool for studying the variable sky, and in order to aid in its exploitation, we have created the Higher-Level Catalogue (HLC). The HLC consists of 229 derived photometric and astrometric quantities for each light-motion curve, and it is stored in a set of FITS binary tables, a format that is widely used by the astronomical community. The photometry presented in the HLC is fully consistent with the IV07 standard star catalogue and, since it is based on many more epochs, the random uncertainties in the mean magnitudes are smaller. Reassuringly, we show that only a very small percentage of standard stars ($\\sim$0.065\\% or $\\sim$650 objects) from IV07 are actually photometrically variable. The HLC proper motions of 30546 stars are consistent to within uncertainties with those derived by \\citet{gou2004} by combining SDSS~DR1 and USNO-B proper motions.  The power in using the HLC is well illustrated by the work of \\citet{vid2007} who construct a reduced proper motion diagram for Stripe 82 in order to find rare white dwarf populations, and consequently they identify 8 new candidate ultracool white dwarfs and 10 new candidate halo white dwarfs. Also, \\citet{bec2008} report the discovery of an eclipsing M-dwarf binary system 2MASS~J01542930+0053266, having employed the corresponding light-motion curve along with radial velocity measurements to determine the masses and radii of the stellar components. We therefore encourage the astronomical community to actively use and explore these catalogues that are so ample in their content."
    },
    "0801/0801.1109.txt": {
        "abstract": "% E+A galaxies have been interpreted as post-starburst galaxies based on % the presence of strong Balmer absorption lines combined with the % absence of major emission lines ([OII] nor H$\\alpha$). As a population % of galaxies in the midst of the transformation, E+A galaxies has been % subject to an intense research activity. %It has been, however, % difficult to investigate E+A galaxies statistically since E+A galaxies % are an extremely rare population of galaxies ($<$1\\% of all galaxies).   We have performed a two-dimensional spectroscopy of three nearby E+A (post-starburst) galaxies with the Kyoto3DII integral field spectrograph. In all the cases, H$\\delta$ absorption is stronger at the centre of the galaxies, but significantly extended in a few kpc scale. For one galaxy (J1656), we found a close companion galaxy at the same redshift. The galaxy turned out to be a star-forming galaxy with a strong emission in H$\\gamma$. For the other two galaxies, we have found that the central post-starburst regions possibly extend toward the direction of the tidal tails. Our results are consistent with the merger/interaction origin of E+A galaxies, where the infalling-gas possibly caused by a galaxy-galaxy merging creates a central-starburst, succeeded by a post-starburst (E+A) phase once the gas is depleted. ",
        "introduction": "Galaxies with strong Balmer absorption lines without any emission in [OII] nor H$\\alpha$ are called E+A galaxies. The existence of strong Balmer absorption lines shows that E+A galaxies have experienced starburst recently \\citep[within a gigayear;][]{2004A&A...427..125G}.  However, these galaxies do not show any sign of on-going star formation as non-detection in the [OII] emission line indicates. Therefore, E+A galaxies have been interpreted as post-starburst galaxies, that is,  a galaxy which truncated starburst suddenly \\citep{1983ApJ...270....7D,1992ApJS...78....1D,1987MNRAS.229..423C,1988MNRAS.230..249M,1990ApJ...350..585N,1991ApJ...381...33F,1996ApJ...471..694A}. Recent study found that E+A galaxies have $\\alpha$-element excess \\citep{2007MNRAS.377.1222G}, which also supported the post-starburst interpretation of E+A galaxies. However, the reason why they started starburst, and why they abruptly stopped starburst remain one of the mysteries in the galaxy evolution.  At first, E+A galaxies are found in cluster regions, especially at higher redshift \\citep{1985MNRAS.212..687S,1986ApJ...304L...5L,1987MNRAS.229..423C,1988MNRAS.235..827B,1991ApJ...381...33F,1995A&A...297...61B,1996MNRAS.279....1B,1998ApJ...498..195F,1998ApJ...507...84M,1998ApJ...497..188C,1999ApJS..122...51D,1999ApJ...518..576P,2003ApJ...599..865T,2004ApJ...617..867D,2004ApJ...609..683T}. Therefore a  cluster specific phenomenon such as the ram-pressure stripping was thought to be responsible for the  violent star formation history of E+A galaxies \\citep{1951ApJ...113..413S,1972ApJ...176....1G,1980ApJ...241..928F,1981ApJ...245..805K,1983ApJ...270....7D,1999MNRAS.308..947A,1999PASJ...51L...1F,2000Sci...288.1617Q,2003ApJ...584..190F,2003ApJ...596L..13B,2004PASJ...56..621F}.  However, \\citet{2004MNRAS.355..713B} found that low redshift E+A galaxies are located predominantly in the field environment, suggesting that a physical mechanism that works in the field region is at least partly responsible for these E+A galaxies. Recently, \\citet{2005MNRAS.357..937G} has shown that E+A galaxies have more close companion galaxies than average galaxies, showing that the dynamical merger/interaction could be the physical origin of field E+A galaxies. Dynamically disturbed morphologies of E+A galaxies also support this scenario \\citep{2006astro.ph.12053L,2005MNRAS.359.1557Y} To reconcile the situation, independent evidence on the origin of E+A galaxies has been waited.  Previous work mentioned above has been focused on the investigation of the global/external properties of E+A galaxies, such as the environment of E+A galaxies\\citep{2005MNRAS.357..937G}, and the integrated spectra of the E+A galaxies \\citep{2004ApJ...617..867D}. However, if the physical origin of E+A galaxies is merger/interaction or gas-stripping, these mechanisms should leave traces to the spatial distribution of stellar-populations in E+A galaxies. For example, a centrally-concentrated post-starburst region is  expected to be found in case of the merger/interaction origin\\citep[e.g.,][]{1992ARA&A..30..705B}. In contrast, the gas-stripping would create a more uniform, galaxy-wide post-starburst region. Thus, the spatial-distribution of the post-starburst region inside the E+A galaxy contains important and independent clues on the physical origin of E+A galaxies \\citep[e.g.,][]{2005MNRAS.359.1421P,2005ApJ...622..260S}. In this work, we try to obtain such independent hints on the origin of E+A galaxies by revealing internal structure of E+A galaxies using the Kyoto3DII integrated field spectrograph (IFS). We perform spatially-resolved spectroscopy of three nearby E+A galaxies with the Kyoto-3DII IFS. By revealing the spatial distribution of H$\\delta$ absorption, we aim to obtain an independent evidence to shed light on the physical origin of E+A galaxies.  Unless otherwise stated, we adopt the WMAP cosmology: $(h,\\Omega_m,\\Omega_L) = (0.71,0.27,0.73)$ \\citep{2003ApJS..148....1B}. ",
        "conclusions": "We have observed three nearby E+A galaxies with the Kyoto3DII IFS. The spatially-resolved spectra show that the strong Balmer absorption characteristic to E+A galaxies is concentrated around the galaxy centre, but is spatially-extended to a few kpc scale. %Both H$\\delta$ and H$\\gamma$ absorption are significantly weaker in the outskirt of the galaxy. These results support a scenario where the infalling gas caused by galaxy-galaxy merging may create the post-starburst phase in the centre of the galaxy. We have found that the extensions of the post-starburst regions coincide with the direction of the tidal tails for J2102 and J2337, possibly supporting the merger/interaction scenario. We have confirmed that a nearby object to J1656 was at the same redshift, and thus is physically-associated. However interestingly, this nearby companion is a star-forming galaxy, and not an E+A galaxy. This implies that the galaxy-galaxy merging may create E+A galaxies at a certain condition, but not all  merging galaxies will evolve into an E+A galaxy. We have found that the H$\\delta$ and H$\\gamma$ absorption lines show different radial trends, which were difficult to be interpreted with the age/metallicity mixture.  %\\begin{figure*} %\\begin{center} %\\includegraphics[bb= 250 200 350 300, clip,scale=2]{j2102_contour_rainbow.ps} %\\includegraphics[bb= 250 150 350 300, clip,scale=2]{j1656_hd_rainbow.ps} %\\includegraphics[bb= 250 150 350 300, clip,scale=2]{j1656_hg_rainbow.ps} %\\includegraphics[bb= 250 100 450 300, clip,scale=1.5]{j2337_contour_rainbow.ps} %\\end{center} %\\caption{Panels (d) of Figures \\ref{fig:J2102}-\\ref{fig:J2337} are colour coded for clarity. %}\\label{fig:color} %\\end{figure*}"
    },
    "0801/0801.3827_arXiv.txt": {
        "abstract": "The 7 year data set of the  Milagro TeV observatory contains $2.2\\times10^{11}$ events of which most are due to hadronic cosmic rays.  This data is searched for evidence of intermediate scale structure. Excess emission on angular scales of $\\sim10^\\circ$ has been found in two localized regions of unknown origin with greater than $12\\sigma$ significance.  Both regions are inconsistent with pure gamma-ray emission with high confidence. One of the regions has a different energy spectrum than the isotropic cosmic-ray flux at a level of $4.6\\sigma$, and it is consistent with hard spectrum protons with an exponential cutoff, with the most significant excess at $\\sim10$ TeV. Potential causes of these excesses are explored, but no compelling explanations are found. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.2156_arXiv.txt": {
        "abstract": "{ We introduce on/off intermittency into a mean field dynamo model by imposing stochastic fluctuations in either the alpha effect or through the inclusion of a fluctuating electromotive force.  Sufficiently strong small scale fluctuations with time scales of the order of 0.3-3 years can produce long term variations in the system on time scales of the order of hundreds of years. However, global suppression of magnetic activity in both hemispheres at once was not observed. The variation of the magnetic field does not resemble that of the sunspot number, but is more reminiscent of the $^{10}$Be record. The interpretation of our results focuses attention on the connection between the level of magnetic activity and the sunspot number, an issue that must be elucidated if long term solar effects are to be well understood. ",
        "introduction": "The sun shows variability on a broad range of time scales, from milliseconds to millennia and on to even  longer scales. Here we are interested in the time scales associated with the solar cycle and its grand minima. A salient manifestation of the cyclic behavior is seen in the sunspot number, whose annual mean oscillates on a scale of eleven years (or twenty-two years, if one goes by magnetic polarity variations). The term ``cycle'' is used to describe this oscillation in the sense that the sunspot number qualitatively performs the same kind of oscillation approximately every eleven years, with an amplitude that varies  in a way that is reminiscent of some chaotic oscillators.  However, it has been found that we do not have sufficient data to decide whether  the global solar magnetic variation is  chaotic, in the sense that its temporal behavior may be that of a low-order deterministic dynamical system (Spiegel \\& Wolf 1987).   Nevertheless, a chaotic oscillator offers a natural way to model various irregularities of the solar cycle (Spiegel 1977; Tavakol 1978; Ruzmaikin 1981).  Like certain simple chaotic systems, the sun repeats itself magnetically over and over again, but never quite the same way twice, much in the manner of simple chaotic oscillators. On the other hand, the simplest chaotic dynamos (Allan 1962; Robbins 1979) do not exhibit the kind of strong intermittency such as was seen in the Maunder Minimum (Eddy 1978) when the amplitude of the oscillation in sunspot numbers went nearly to zero for seventy-five years in the time of Newton. This behavior is suggestive of the possibility that a strongly intermittent chaotic oscillator needs to be invoked in trying to understand the solar oscillation. Oscillators that behave this way are known (Spiegel 1981; Fautrelle \\& Childress 1982) but, when they are in the active phase, their variations do not normally resemble those of the sunspot number.  In modeling the solar magnetic variability with simple oscillators, it is not obvious how to connect the output of the models with the sunspot number. It may therefore not be damning if the variations produced by a model do not reproduce even qualitatively the variations seen in the sunspot number. Nevertheless, in terms of nonlinear lumped models, whose behavior is purely temporal, it has been possible to produce variations of the cyclic behavior that resemble those seen in the grand minima qualitatively either by modulation (Weiss et al.\\ 1984) or intermittency (Platt et al.\\ 1993b). But the solar variability is manifestly spatio-temporal and lumped models can at best provide clues to the actual processes involved in the solar cycle. Here we describe an attempt to go beyond models with purely temporal variations to models showing spatio-temporal variations that also produce analogues of the grand solar minima.  An earlier version of this paper was composed in the early nineties and some added remarks were inspired by discussions at the Enrico Fermi School in Varenna (Cini Castagnoli \\& Provenzale 1997). In the meantime a lot of new work has emerged, but we feel that the ideas presented here are still relevant. Particularly important has been the work of Beer et al.\\ (1998) in demonstrating the persistence of a solar activity cycle throughout the time of the Maunder Minimum. Such behavior emerged from a purely temporal model (Pasquero 1996) as well as from a one-dimensional dynamo model (Tobias 1996) with small turbulent magnetic Prandtl number (so that the viscous diffusion timescale is much longer than that for magnetic diffusion time). Similar results --- also for small turbulent magnetic Prandtl numbers --- have been obtained for two-dimensional models by incorporating quenching in the $\\Lambda$-effect that drives the differential rotation (K\\\"uker et al.\\ 1999). Intermittent behavior in dynamo models has been studied further by Tworkowski et al.\\ (1998) using time-dependent alpha-quenching, and by Charbonneau et al.\\ (2005) using algebraic alpha-quenching. However, intermittency can equally well occur in stochastically forced models (e.g.\\ John et al.\\ 2002).  In the following, we discuss the spatio-temporal variability of similarly forced dynamo models in two dimensions assuming axisymmetry. ",
        "conclusions": ""
    },
    "0801/0801.2967_arXiv.txt": {
        "abstract": "The observed behavior of $\\eta$ Car from 1860 to 1940 has not been considered in most recent accounts, nor has it been explained in any quantitative model.  We have used modern digital processing techniques to examine Harvard objective-prism spectra made from 1892 to 1941.  Relatively high-excitation \\ion{He}{1} $\\lambda$4471 and [\\ion{Fe}{3}] $\\lambda$4658 emission, conspicuous today, were weak and perhaps absent throughout those years.  Feast et al.\\ noted this qualitative fact for other pre-1920 spectra, but we quantify it and extend it to a time only three years before Gaviola's first observations of the high-excitation features.  Evidently the supply of helium-ionizing photons ($\\lambda < 504$ {\\AA}) grew rapidly between 1941 and 1944. {\\it The apparent scarcity of such far-UV radiation before 1944 is difficult to explain in models that employ a hot massive secondary star,} because no feasible dense wind or obscuration by dust would have hidden the photoionization caused by the proposed companion during most of its orbital period. We also discuss the qualitative near-constancy of the spectrum from 1900 to 1940, and $\\eta$ Car's photometric and spectroscopic transition between 1940 and 1953. ",
        "introduction": "Eta Carinae is famous for its many superlatives as the most massive, most luminous star in our region of the Milky Way and for its ``great eruption'' (1837-1858) more than 160 yrs ago, when it ejected 6 M$_{\\odot}$ or more (Davidson \\& Humphreys 1997, Morris et al 1999, Smith et al 2003a) forming its bipolar Homunculus Nebula. $\\eta$ Car is also our closest and best studied example of a ``supernova impostor'', but there are still numerous  unsolved questions about the cause of the great eruption and the star's subsequent and on-going recovery. One of the most outstanding is its continuing instability and what happened when $\\eta$ Car experienced two post-eruption  episodes of spectroscopic and photometric change: the second eruption (1888 - 1895) and again from  about 1941 to 1952 when it  rapidly brightened twice and the first mention of the high excitation emission lines appears in the astronomical literature (Gaviola 1953). The latter event has received little attention, although it may be fundamental to understanding the star's process of recovery.   It is now well known that $\\eta$ Car experiences a spectroscopic/photometric 5.5 year cycle (Whitelock et al. 1994, Damineli 1996) . Modern groundbased spectra normally show high excitation emission lines, notably of He I, [Fe III], [Ne III], etc. sometimes referred to as the excitation maximum or ``high state''. During a spectroscopic event or excitation minimum (low state) however, the high excitation lines weaken and even disappear. These ``events'' apparently last from a few weeks to a few months, and were noticed by several  authors before 1996 (Gaviola 1953, Thackeray 1967, Rodgers \\& Searle 1967, Viotti 1968, Whitelock et al. 1983, Zanella et al. 1984, Bidelman et al. 1993). The 5.5 year spectroscopic cycle has been confirmed beginning with Gaviola's (1953) first recorded evidence for a spectroscopic change in 1948 up to the most recent events in 1998.0 and 2003.5 (Damineli et al. 2000, Davidson et al. 2000, 2005).   Feast, Whitelock and Marang (2001) examined the spectra of $\\eta$ Car obtained at the Radcliffe Observatory in South Africa from 1947 to 1974 and confirmed the cycle as expected.  However, they also reported  that the earliest spectra from 1899 to 1919 never showed the high excitation state. Gaviola (1953) had also noted that the He I emission lines were not present in the spectra described by Moore and Sanford (1913, 1916) and Lunt (1919). The plates examined by Feast et al. (2001)  included two  from 1912-1914  described by Moore and Sanford plus an additional plate from 1914 and two others discussed by Lunt. Feast et al.  also discussed Gill's early spectrum (Gill 1899) and reexamined his measurements, but the plate could not be located. Thus they examined five plates in total which they said were ``of excellent quality'', looking specifically for the He I line at 4471{\\AA} which should have been easily discernible.  {\\it Therefore from approximately 1899 to 1919 the star's spectrum  was never in the high excitation state, and there were  no spectroscopic events as we understand them now.} The 5.5 yr cycle has become an important part of $\\eta$ Car's expected behavior, and has been variously attributed to 1) the presence of a hot secondary star which is eclipsed by $\\eta$ Car (Damineli et al. 2000) blocking the UV radiation for the high excitation lines, 2)  a disturbance in the wind and mass ejection (Zanella et al. 1984, Davidson 1999, Davidson et al.  2000), or 3) a combination of the two (Davidson 2005). The absence of high excitation lines and the spectroscopic cycle in these early spectra naturally raises questions about the origin of the cycle and when and why  the high excitation lines first appeared.   The spectroscopic record from 1920 to the time of Gaviola's spectral series beginning in 1944 was missing in the Feast et al. survey. However, the famous Harvard plate collection contains a series of objective prism plates of the Southern sky covering 1891 to 1951. We examined all of the available plates with spectra of $\\eta$  Car at the Harvard-Smithsonian Center for Astrophysics (CfA). In 2003 the staff at CfA  had just begun a trial program to digitize the Harvard plate collection, and  we were permitted to select several representative spectra for digitization.  In the next section, we describe the Harvard spectra of $\\eta$ Car obtained from 1892 to 1941, and our rectification, calibration and extraction process. In \\S{3} and \\S{4}, respectively, we discuss the earliest spectra from the time of the second eruption, and those from 1902 to 1941. Measurements of the digitized spectra and comparison with recent groundbased spectra show  that  {\\it He I $\\lambda$ 4471 emission was absent or only marginally present.} If present at all, it  was {\\it much weaker} than reported by Gaviola and insignificant compared with current spectra. Similarly, the [Fe III] lines in these spectra are either absent or  marginal,  and there were apparently no spectroscopic events. This result has serious implications for the origin of the He I and the high excitation lines in $\\eta$ Car discussed in the rest of this paper. In \\S{5} we describe the the apparent changes in $\\eta$ Car's wind over more than 50 yrs and the onset of the high excitation emission during the  period of rapid photometric change from $\\sim$ 1941 to 1952. We review the He I emission problem in \\S{6} including  the  possible sources of present day He I emission  in $\\eta$ Car and its absence prior to 1944. Our conclusions are summarized in the last section. ",
        "conclusions": "Our measurement and analysis of the Harvard objective prism spectra of $\\eta$ Car from 1902 to 1941, has shown that He I 4471{\\AA} emission may be marginally present, but  no stronger than about  one-fourth of  its relative intensity  as estimated by Gaviola in spectra from 1944-51.  Our results are equally consistent with no contribution from the He I $\\lambda$4471 line. The [Fe III]  high excitation lines are either absent or very weak in these spectra. The star was thus always in the ``low excitation'' state and there were no apparent spectroscopic events from the time of the second eruption in the 1890's to 1944.  Obscuration of the proposed companion star by either dust or an optically thick wind cannot explain the lack of He I emission prior to 1944. We discussed the possible sources of He I emission in the $\\eta$ Car system and demonstrated  that it is very difficult to explain its absence if the UV flux comes from the proposed hot secondary star. Alternatively, we suggest that the equatorial wind of the primary may be the source of the UV flux due to $\\eta$ Car's post-eruption rotational spin-up and the onset of its present bipolar wind, coincident with its brightening between 1940 and 1952 (Smith et al. 2003b).  This work has also demonstrated that with careful processing and calibration these historical spectra are quite useful especially for objects of astrophysical interest."
    },
    "0801/0801.3293_arXiv.txt": {
        "abstract": "We examine the effect of magnetic fields on star cluster formation by performing simulations following the self-gravitating collapse of a turbulent molecular cloud to form stars in ideal MHD. The collapse of the cloud is computed for global mass-to-flux ratios of $\\infty$, $20$, $10$, $5$ and $3$, that is using both weak and strong magnetic fields. Whilst even at very low strengths the magnetic field is able to significantly influence the star formation process, for magnetic fields with plasma $\\beta < 1$ the results are substantially different to the hydrodynamic case. In these cases we find large-scale magnetically-supported voids imprinted in the cloud structure; anisotropic turbulent motions and column density structure aligned with the magnetic field lines, both of which have recently been observed in the Taurus molecular cloud. We also find strongly suppressed accretion in the magnetised runs, leading to up to a 75\\% reduction in the amount of mass converted into stars over the course of the calculations and a more quiescent mode of star formation. There is also some indication that the relative formation efficiency of brown dwarfs is lower in the strongly magnetised runs due to the reduction in the importance of protostellar ejections. ",
        "introduction": "Understanding how stars like the Sun form is one of the key questions central to our understanding of the Universe we live in. Whilst we have come a long way in this understanding since the pioneering work of Jeans, many of the fundamental questions such as the rate, distribution and efficiency of star formation remain either unknown or are the subject of vigourous debate. One of the key areas of uncertainty is whether or not star formation is a rapid \\citep{mk04,hbb01,elmegreen07} or slow \\citep{sla87,tkm06} process, central to which is the relative importance of magnetic fields to the star formation process.  Whether or not star formation is rapid or slow, the fact remains that molecular clouds are observed to contain magnetic fields of sufficient strengths that they cannot be ignored in any complete theory of how stars form from such clouds \\citep[e.g.][]{crutcher99,bourkeetal01,hc05}. Furthermore molecular clouds are observed to contain supersonic turbulent motions \\citep{larson81}, so the interaction of turbulence and magnetic fields is critical to our understanding of the star formation process. This interaction has been the subject of a number of studies which have shown that magnetic fields are \\emph{not} effective in preventing the rapid dissipation of supersonic turbulence in the absence of continued driving \\citep{sog98,maclowea98,vsetal00}, though the presence of magnetic fields does change the dynamics of the turbulence \\citep{pn02,padoanetal07}. However, to date, there has been only a handful of simulations which have attempted to follow the self-gravitating collapse of a turbulent cloud in the presence of a magnetic field \\citep[e.g.][]{lietal04,ln06,vsetal05,tp07}.  The ability of magnetic fields to provide support against gravitational instability is determined, for an enclosed region of gas threaded by a magnetic field, by the ratio of the mass contained within the region to the magnetic flux passing through the surface, ie. the \\emph{mass-to-flux} ratio, which for a spherical cloud is given by \\begin{equation} \\frac{M}{\\Phi} \\equiv \\frac{M}{4\\pi R^{2} B_{0}}, \\label{eq:masstoflux} \\end{equation} where $M$ is the mass contained within the cloud volume, $\\Phi$ is the magnetic flux threading the cloud surface at radius $R$ assuming a uniform magnetic field $B_{0}$. The critical value of $M/\\Phi$ below which a cloud will be supported against gravitational collapse is given by \\citep[e.g.][]{ms76,mestel99,mk04} \\begin{equation} \\left(\\frac{M}{\\Phi}\\right)_{\\rm crit} = \\frac{2 c_{1}}{3} \\sqrt{\\frac{5}{\\pi G \\mu_{0} }}, \\label{eq:mphicrit} \\end{equation} where $G$ and $\\mu_{0}$ are the gravitational constant and the permeability of free space respectively and $c_{1}$ is a constant determined numerically by \\citet{ms76} to be $c_{1}\\approx 0.53$. Star forming cores with mass-to-flux ratios less than unity are stable against collapse (``subcritical'') and conversely, cores with mass-to-flux ratios greater than unity (termed ``supercritical'') will collapse on the free-fall timescale. Throughout this paper we use the mass-to-flux ratio, defined in terms of the critical value, to quantify the magnetic support of a molecular cloud against collapse and we perform simulations using initially supercritical clouds (though to varying degrees). Giant molecular clouds are generally thought to have magnetically subcritical envelopes but to be supercritical in their inner parts \\citet{cm95,ccw05}, a picture which is largely confirmed by observational results on both large scales \\citep{mckee89,mckeeetal93} and in dense cores \\citep{crutcher99,bourkeetal01}.  \\citet{lietal04} examined the role of magnetic fields in self-gravitating core formation within a turbulent molecular cloud in a periodic box. They found that cores formed within a supercritical cloud were also locally supercritical by at least an order of magnitude, indicating that a globally supercritical magnetic field does \\emph{not} evolve to produce locally magnetically subcritical cores, contrary to some earlier expectations \\citep{ms56}. Furthermore, even with supercritical cores they found a central magnetic field strength in cores $B \\propto \\rho^{1/2}$, similar to observations \\citep{crutcher99} which has often been used as an argument for magnetic support in molecular cloud cores. They also found strong interaction between cores and rotationally supported discs. \\citet{tp07} also performed simulations of self-gravitating collapse in the presence of magnetic fields and found that the observed near-critical cores could form naturally from a globally highly supercritical cloud, and that these cores generally have the same gas-to-magnetic pressure ratio, $\\beta$ as the mean $\\beta$ in the global cloud.  Given that magnetic fields always act to oppose gravitational collapse, it has often been suggested that they may play the dominant role in regulating the star formation efficiency in molecular clouds \\citep[e.g.][]{kt07}, though there are also several other good candidates for doing so, including turbulence-inhibited collapse \\citep[e.g.][]{tp04,lietal04,vsetal05}, feedback from jets and outflows \\citep{nl05,ln06}, the dispersal of initially unbound clouds \\citep{clarketal05} or radiation feedback from the stars themselves. Secondly magnetic fields are often invoked to solve the `angular momentum problem' in star formation via the magnetic braking of star forming cores. Recent simulations by \\citet{pb07} have shown that magnetic fields can have a dramatic effect on circumstellar disc formation and on fragmentation to form binary systems (these results have since been confirmed by \\citealt{ht08}). \\citet{tp07} looked at the effect of the magnetic fields on the star formation efficiency in their simulations, though no clear trend was apparent, in part due to limitations of the numerical model.  Despite the apparent importance of magnetic fields to the star formation problem, it was therefore somewhat surprising that the purely hydrodynamic calculations of \\citet{bbb03} (hereafter BBB03) produced largely the `right answer' in terms of being able to reproduce observed the stellar initial mass function (albeit at low number statistics) as well as several other observational characteristics such as the frequency of binary stars and stellar velocity dispersions. In this and subsequent calculations the initial mass function is built up due to the competition between dynamically interacting protostars in order to accrete from the global cloud \\citep{bb06}. Thus, low mass stars are simply those which have been quickly ejected from multiple systems and thus have only a short accretion history \\citep{bbb02a,bb05}, whereas higher mass stars are those which form and remain at the bottom of deep potential wells and build up their mass through accretion over time \\citep{bonnelletal97}. However subsequent calculations (Bate 2008, in prep) have established that purely hydrodynamic calculations produce an excess of brown dwarfs relative to observations.  Thus it is crucial to extend these types of calculations, with the hope of resolving some of the above issues, by adding the two major pieces of missing physics -- magnetic fields and the effect of radiation transport. This paper presents our first attempt to address the former in large-scale simulations. Whilst in a sense the two are complementary, since we expect the magnetic field to have an effect primarily on larger scales (given the strong physical diffusion of the magnetic field on smaller scales), whilst radiation might be expected to influence the smaller scale dynamics (ie. fragmentation), it is clear that it is imperative to incorporate both pieces of physics into these types of calculations.  The paper presents our first investigation of how magnetic fields change the picture of star cluster formation painted by BBB03. The numerical method is discussed in \\S\\ref{sec:numerics} and the initial conditions for the simulations are discussed in \\S\\ref{sec:initconds}.  Results are presented in \\S\\ref{sec:results} and discussed in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion} We have performed a study of how magnetic fields affect the large scale collapse of turbulent molecular clouds to form star clusters, computing a range of models with mass-to-flux ratios ranging from highly to moderately supercritical (with a corresponding range in the ratio of gas to magnetic pressure, $\\beta$). Whilst even the weakest field runs show differences when compared to the hydrodynamic case (e.g. lower accretion rates, different star formation sequences), strong differences in the gas dynamics were found to be present in the runs where the magnetic pressure dominates the gas pressure (ie. $\\beta < 1$), in the form of magnetically-supported voids, column density striations due to anisotropic turbulent motions and much lower accretion rates from the global cloud. The implication of these results for both observations of star-forming molecular clouds and for our theoretical understanding of star cluster formation are discussed below.  \\subsection{Relevance to observations} Zeeman measurements of magnetic fields in molecular clouds \\citep{crutcher99,bourkeetal01} suggest typical field strengths of $B\\sim 10\\mu$G (and generally $B \\lesssim 30\\mu$G) for regions with $T \\sim 10K$, mass-to-flux ratios which are supercritical by a factor of $\\sim 2-3$, marginally (or perhaps firmly, see \\citealt{padoanetal04}) super-Alfv\\'enic turbulent velocities and $\\beta \\sim 0.03-0.6$ (similar values for $\\beta$ and the ratio of turbulent to Alfv\\'enic velocities are also inferred by \\citealt{ht05} in the wider Cold Neutral Medium). The inference therefore is that the most realistic of our calculations are actually the two strongest field runs ($M/\\Phi = 3$ and $M/\\Phi = 5$). Given that this is the case, we should expect that \\emph{all} of the magnetically-driven features observed in these simulations to also be present in observed molecular clouds.  Recent first results from a large-scale mapping project of the Taurus molecular cloud complex by \\citet{goldsmithetal05} report ``ring, arc and bubble-like features'' and ``striated structure... correlated with the magnetic field direction'' in $^{12}$CO maps. In fact almost the whole of the low density material in Taurus in the \\citet{goldsmithetal05} map appears striated parallel to the global magnetic field threading the cloud as mapped from polarization measurements, as was found to occur in the simulations in the low-density outer regions of the clouds. This is suggestive of a low plasma $\\beta$ in these regions, since we only find striated structure in the two strongest field simulations with $\\beta < 1$ (and most prominently in the $M/\\Phi = 3$ case where $\\beta = 0.3$). This is in broad agreement with the polarisation measurements of \\citet{crutcher99} giving lower limits of $\\beta \\gtrsim 0.06$ for regions within Taurus.  One of the most striking features of the magnetised collapse simulations is the appearance of large-scale magnetic-pressure supported voids produced by the large scale magnetic flux tubes threading the cloud. In these regions the integrated magnetic pressure appears to anti-correlate with the column density of the cloud (Figure~\\ref{fig:magcushion}). Turning again to the observations of Taurus, \\citet{goldsmithetal05} report a ``very interesting feature'' at $4^{h}30^{m} + 25^{\\deg}$ which ``appears as a hole'', where it appears that ``some agent has been responsible for dispersing the molecular gas''. We propose that this may be a magnetic-pressure driven feature. The lack of polarization measurements from this region and the orientation of the ``hole'' suggests that the magnetic field should be aligned parallel to the line of sight and would therefore be best detectable using Faraday rotation measurements. Whilst very few such observations exist, the measurements of \\citet{wr04} give some hint of increased emission in this regions (suggesting a field which is parallel to the line of sight), though somewhat ambiguously. Thus further observations are necessary to confirm this picture.  \\subsection{Relevance to theory} One of the primary effects of magnetic fields in the simulations is that, even at field strengths which do not prevent collapse, the field can have a significant influence on the star formation rate in the cloud. For example overall we find that, after 1.5 cloud free-fall times, only 4\\% of the gas has been converted into stars for a marginally supercritical collapse ($M/\\Phi = 3$) compared with 16\\% in the hydrodynamic case. Similar effects of the magnetic field on the star formation rate have been found in MHD calculations of star formation in the presence of driven turbulence \\citet{vsetal05}. In part this can be attributed to the simple fact that the magnetic field adds an extra source of pressure support to the cloud. Thus we would expect that in the present context it would be possible to similarly decrease the star formation rate by, for example, simply scaling up the turbulent velocity field or by increasing the temperature of the cloud. However we would also expect the change in the cloud geometry to be rather different under these circumstances, since the magnetic field preferentially supports material perpendicular to field lines. The degree to which this pressure support is anisotropic depends on the ratio of gas-to-magnetic pressure, $\\beta$ -- for example a weak field will be much more readily tangled and thus exert a more isotropic pressure, whereas a stronger field will exert a pressure that has a much stronger dependence on the initial field geometry. Furthermore we would not expect to find any of the magnetically driven cloud structures, such as the filaments in the cloud envelope aligned with the global field direction and magnetically supported voids. \\citet{pb07} found that the degree to which the magnetic field can be replaced by an equivalent increase in thermal pressure is strongly dependent on the field geometry.   The net result of the lower star formation rate in the magnetised runs is also a decrease in the importance of violent interactions, leading to fewer ejections and therefore also a trend (though tentative) towards more massive stars being formed (ie. relatively fewer brown dwarfs). Such a trend might have important implications for the variation in the IMF in different environments \\citep[e.g.][]{kroupa01} and as a function of galaxy evolution. In the local environment the IMF appears to have a universal shape (at least within the observational uncertainties), although this may be because there is also remarkable uniformity in the inferred level of magnetic field support in observations of local star formation regions \\citep{crutcher99} (namely that, as discussed above, most star forming cores appear to be marginally supercritical with super-Alfv\\'enic velocity dispersions and $\\beta \\sim 0.1$).  \\subsection{Limitations and future directions} Perhaps the most significant limitation of the calculations presented in this paper is that, at the resolution employed, we were not able to study the effect of magnetic fields on the small scale fragmentation of cores, which may be important with respect to the formation of circumstellar discs \\citep{pb07} and in affecting fragmentation to form binary systems \\citep{machida05b,pb07,ht08,machidaetal07}. A related question which would be interesting to examine is how or whether the global magnetic field affects the frequency and/or strength of jets and outflows in molecular clouds \\citep[e.g.][]{mt04,bp06}.  Whilst increasing the numerical resolution will improve the resolution of the small scale magnetohydrodynamics, there is not strong motivation to do so as the small scale dynamics is also affected by other physical processes which have not been modelled in the present calculations. The most important of these are the effects of radiative transfer (ie. replacing the barytropic equation of state) and physical diffusion in the magnetic field due to non-ideal MHD.  As mentioned in \\S\\ref{sec:eos}, we expect from preliminary simulations which incorporate a self-consistent treatment of radiative transfer (in the flux-limited diffusion approximation) rather than a barytropic equation of state \\citep{wb06} that a full treatment will affect fragmentation where the assumption of spherical symmetry is poor and in the region surrounding a collapsing core where the temperature can increase because of the propagation of radiation from the high temperature central condensation into the lower density gas (a process not captured by a barytropic equation of state where temperature is proportional to density). These issues in particular may be important for fragmentation in discs, which, it should be noted is the prime source of much of the star formation activity in the hydrodynamic and very weak-field ($M/\\Phi = 20$) calculations presented here. It is therefore crucial to quantify the degree to which this fragmentation is physical by performing calculations which explicitly evolve the radiation field.  Secondly, whilst flux-freezing is thought to be a good approximation for molecular cloud dynamics on large scales, non-ideal MHD effects including ambipolar (ion-neutral) diffusion \\citep[e.g.][]{mp81,sla87}, the effect of finite conductivity and the Hall effect are all important at some level on smaller scales \\citep[e.g.][]{wn99}. This in particular applies to the process of core fragmentation (though \\citealt{om06} suggest that ambipolar diffusion is unable to set a characteristic mass scale in a turbulent flow because of the continued propagation of compressive slow MHD waves below the ambipolar diffusion scale). It has also been suggested that ambipolar diffusion rates may be enhanced in turbulent flow \\citep{heitschetal04,ln04} and therefore that ion-neutral diffusion may be important also in the earlier stages of collapse. Thus it is imperative that the calculations should be extended to include non-ideal MHD effects in order to quantify these effects in the present calculations (for example \\citet{hw04a} have already implemented a two-fluid scheme for treating ion-neutral diffusion in an SPH context which could be used in future calculations). Studying physical diffusion processes also requires that the numerical diffusion (e.g. due to the artificial resistivity introduced in order to capture shocks) be reduced to a level below the physical diffusion scale, leading to a much more stringent criterion for fully resolved MHD simulations (of which those presented here are \\emph{not}) compared to purely hydrodynamic runs (where it is sufficient to resolve the Jean's length in order to capture the fragmentation scale)."
    },
    "0801/0801.4685_arXiv.txt": {
        "abstract": "{The nature and origin of the soft X--ray excess in radio--quiet AGN is still an open issue. The interpretation in terms of thermal disc emission has been challanged by the discovery of the constancy of the effective temperature despite the wide range of Black Hole masses of the observed sources. Alternative models are reflection from ionized matter and absorption in a relativistically smeared wind.} { We analyzed XMM--$Newton$ observations of four luminous radio--quiet AGN with the aim of characterising their main properties and in particular the soft excess.} {Different spectral models for the soft excess were tried: thermal disc emission, Comptonization, ionized reflection, relativistically smeared winds.}{Comptonization of thermal emission and the smeared winds provide the best fits, but the other models also provide acceptable fits. All models, however, return parameters very similar from source to source, despite the large differences in luminosities, Black Hole masses and Eddington ratios. Moreover, the smeared wind model require very large smearing velocities. The UV to X--ray fluxes ratios are very different, but do not correlate with any other parameter.}{No fully satisfactory explanation for the soft X--ray excess is found. Better data, like e.g. observations in a broader energy band, are needed to make further progresses. } ",
        "introduction": "While relatively low luminosity AGN are well studied in X-rays, not the same can be said for high luminosity ones, due to their paucity in the local Universe and therefore on their relative low fluxes. With the aim to populate the high L -- high flux portion of the parameter space, we proposed and obtained XMM--$Newton$ observations of four sources, selected from the Grossan HEAO1 LMA catalogue to be bright (F$_{2-10}$ $>$ 10$^{-11}$ erg cm$^{-2}$ s$^{-1}$) and luminous (L$_{2-10}$ $>$ 10$^{44}$ erg s$^{-1}$): H0439-272, Ark 374, Fairall 1116 and PG0052+251 (see Table 1).  HEAO1 instruments scanned the whole sky from 1977 to 1979. The A-1 instrument Large Sky Survey (LASS, Wood et al. 1984) catalogue contains 322 sources brighter than 0.0036 LASS counts/s/cm$^2$ (F$_{2-10}$$>$ 1.82 $\\cdot$ 10$^{-11}$ erg cm$^{-2}$ s$^{-1}$, assuming a power law spectrum with $\\Gamma$ = 1.7) and with $\\mid$b$\\mid$ $>$ 20 (Grossan, 1992). The A-1 instrument consisted of a set of proportional counters sensitive in the 2-20 keV band, which observed the sky through passive collimators with field of view of 1x4 and 0.5x1 degrees. The precision with which the position of each X-ray source was determinated was greatly improved by the simultaneous observation of the MC instrument (Modulation Collimator or A-3). The MC instrument produced a pattern of small, diamond-shaped error regions, only 1x4 arcmin each (90$\\%$ confidence limit). The size of the final LASS/MC error box is therefore of $\\sim$ 0.3 deg$^2$. Grossan and collaborators identified optical counterparts of 287 of the 322 objects in the LASS catalog (86$\\%$). 96 of these objects are identified with AGN: the LASS-MC-AGN (LMA) sample (Grossan, 1992), by correlation with catalogs of known sources and through optical photometry and spectroscopy performed by the MIT group.  Information on the four sources and on the related observations are summarized in Tables~\\ref{log} and ~\\ref{fluxes}. Fluxes are derived from the best fit models (see Sec.3). Unfortunately, the sources were all found at fluxes (and therefore luminosities) lower by at least a factor 2 than in the HEAO-1 observations (not unexpectedly, as they all were just above the flux threshold). Fortunately, they were still bright enough to allow for a detailed spectral analysis.  In this paper, we report on the spectral analysis of these four sources, with particular emphasis on the soft excess.  A soft excess is present in the X-ray spectra of most Seyfert 1s and radio-quiet quasars. First noted by Arnaud et al. (1985), it is an emission below $\\sim$ 1 keV in excess of the extrapolation of the power law component dominating at higher energies.  The origin of the soft excess is still an open issue. In the past, it was often associated with the thermal emission of the accretion disc. The large effective temperatures usually obtained (when compared to what expected in a standard Shakura $\\&$ Sunyaev, 1973, accretion disc), were attributed to a slim accretion disc in which the temperature is raised by photon trapping (Abramowicz et al. 1998, Mineshige et al. 2000), in which case the accretion is super-Eddington, or to Comptonisation of EUV accretion disc photons (Porquet et al. 2004).  However, it has been recently shown that the disc temperature is remarkably constant around 0.1 - 0.2 keV, regardless of the central object luminosity and mass (Gierlinski $\\&$ Done 2004, Crummy et al. 2006). This result is difficult to explain in any model for the soft excess related to disc continuum emission, as in any disc model the temperature is expected to vary with both the Black hole mass and the accretion rate.  On the other hand, the constancy of the temperature may arise naturally if the soft excess is not related to thermal emission at all, but due to absorption/emission processes, the ``temperature'' actually measuring atomic transitions. In fact, there is a great increase in opacity in partially ionized material due to lines and edges of ionized O VII, O VIII and Iron at $\\sim$ 0.7 keV, which can make the apparent soft excess, provided that velocity smearing broadens and blurs the atomic features.  On this line of tought, two alternative scenarios have been proposed: a partially ionized, relativistically blurred reflection from the accretion disc (Crummy et al. 2006) and a velocity smeared, partially ionized absorption (Gierlinski $\\&$ Done 2004), the latter model further developed by Schurch $\\&$ Done (2006) by including the emission associated with the absorbing material and proposing an origin in a `failed wind'.  In the reflection model, the partially ionized material is naturally identified with the accretion disc, but to produce a strong soft excess the intrinsic continuum must be significantly suppressed, perhaps by disc fragmentation or by light--bending effects (Fabian et al. 2002, 2004, 2005; Miniutti $\\&$ Fabian 2004).  In the absorption model, the partially ionized material is supposed to be a wind likely starting from the disc (Schurch $\\&$ Done 2006). If this is the case, the velocity structure  has to be complex, giving rise to a substantial broadening to mask the sharp atomic features, which otherwise should have been observed in high resolution gratings observations. A problem with this interpretation, as noted by the same authors, is that such winds tend to be produced within $\\sim$ 25$^{\\circ}$ of the equatorial plane. Most of type-1 AGN are expected to have lines of sight which would not intercept much of this material, but this is in conflict with the fact that most type-1 AGN show a soft excess.\\ Moreover, the required large velocity smearing implies a mass-loss rate much larger than the accretion rate required to power the observed luminosity. Schurch $\\&$ Done (2006) then proposed a `failed wind' instead, which does not require a large mass-loss rate. In such a case, the central X-ray source is strong enough to overionize the wind removing the acceleration before the material reaches escape velocity, allowing the material to fall back to the disk. ",
        "conclusions": "\\begin{figure} \\epsfig{file=SX.ps,width=6cm,angle=-90} \\caption{A 100 ks Simbol-X simulations of the baseline model spectrum of Fairall 1116 (upper panel), and the data/model ratio when the spectrum is fitted with the smeared absorption model (lower panel).} \\label{sx} \\end{figure}   We have analyzed 4 moderately luminous AGN observed with the XMM--$Newton$ satellite. The main results can be summarized as follows:  a) The hard X-ray emission is characterized by a power law with spectral indices with depends very much on the modeling of the soft excess. If the latter is modelled as a power law, the hard X-ray photon index is quite flat, $\\Gamma$ ranging from 1.25 to 1.56. If, on the other hand, other parameterization are chosen (some of them giving fits of comparable quality), much steeper $\\Gamma$ are found, similar to those usually found in AGN (e.g. Porquet et al. 2004, Piconcelli et al. 2005, Bianchi et al. in prep.), apart when using a smeared absorbing wind, when quite steep values, $\\Gamma >$2, are found.   b) An iron line is also detected, even if in PG 0052+251 only an upper limit is found. The error bars are very large, but the decrease of the best fit values of the EW with both the luminosity and the the Eddington ratio is in agreement with what found by Bianchi et al. (2007) in a much larger sample (the Iwawasa--Taniguchi effect).  c) No satisfactory modeling for the soft excess is found.  A power law, besides resulting in suspiciously flat hard power laws, has no obvious physical meaning, and it will not be discussed further. We just note that the Comptonization model (see below) is basically a power law (and indeed the fits are statistically comparable) with a cutoff (which results in a steeper hard power law).  A multicolor thermal disc emission, apart from providing significantly worse fits, results in very high temperatures, almost constant despite the factor $\\sim$10 differences in Black Hole masses, which should result in a factor $\\sim$2 differences in $kT$.  A Comptonization model provides both good fits and ``normal'' hard power law indices, but both the temperature and the optical depth of the Comptonizing electron are suspiciously the same for all objects, with no obvious reasons for such a constancy.  A (relativistically blurred) ionized reflection model provides not very good fits (which may be attributed to the oversimplifications of the model, like a single ionization zone while a centrally illuminated disc obviously has a strong radial dependence of $\\xi$, e.g. Matt et al. 1993). Moreover, the range of ionization parameters spanned by the four sources is rather small, even if this may be due to the fact that the zone of the disc with that ionization parameters are the ones contributing most to the soft excess (smaller ionizations resulting in a smaller albedo, higher ionizations in a smaller energy dependence of the albedo).  Finally, the relativistically smeared absorbing wind model provides good fits to the data, but the three model parameters (column density, ionization parameter and smearing velocity) are almost constant among the sources. While there may be an explanation for the constancy of the column and the ionization parameter (see discussion in Middleton et al. 2006), the constancy of the smearing velocity and, above all, its high value (about 0.5 $c$) are difficult to explain. Outflowing winds with velocities as high as 0.1-0.2 $c$ have been observed in a few high accretion rate sources (e.g. Pounds et al. 2003). Fixing the smearing velocity to 0.15, however, results in unacceptable fits. It is interesting to remark that Schurch \\& Done (2007b), using a more refined version of the model (not availaible for spectral fitting), found that indeed extremely large, unrealistic velocities are required to explain the soft X-ray excesses in AGN.  In summary, none of the abovementioned model provides a satisfactory description of the data, even if none of them can be completely ruled out. The constancy of the parameters  we found for {\\it all} models is clearly telling us that some characteristic energy is involved. Future, high sensitivity broad band measurements like those provided by e.g. Simbol-X (Ferrando et al. 2006) will hopefully tell us which, if any, of these model is tenable (Matt 2007, Ponti et al. 2007). As an example, in Fig.~\\ref{sx} the Simbol-X simulation of the baseline model for Fairall 1116 is shown, as well as the data/model ratio after fitting with the smeared absorption model.  d) The UV/X ratio is very different from source to source -- more than a factor 30 -- and does not correlate with any other parameter. It is important to note that the soft--to--hard X-rays ratio is instead basically constant, suggesting that the UV and soft X--rays are not physically related -- another clue against a thermal disc origin for the soft X--rays."
    },
    "0801/0801.4961_arXiv.txt": {
        "abstract": "Simultaneous observations by the large number of gamma-ray burst detectors operating in the GLAST era will provide the spectra, lightcurves and locations necessary for studying burst physics and testing the putative relations between intrinsic burst properties. The detectors' energy band and the accumulation timescale of their trigger system affect their sensitivity to hard vs. soft and long vs. short bursts.  Coordination of the Swift and GLAST observing plans consistent with Swift's other science objectives could increase the detection rate of GLAST bursts with redshifts. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.1069_arXiv.txt": {
        "abstract": "We present Very Large Array 21-cm observations of the massive edge-on spiral galaxy NGC 5746.  This galaxy has recently been reported to have a luminous X-ray halo, which has been taken as evidence of residual hot gas as predicted in galaxy formation scenarios.  Such models also predict that some of this gas should undergo thermal instabilities, leading to a population of warm clouds falling onto the disk.  If so, then one might expect to find a vertically extended neutral layer.  We detect a substantial high-latitude component, but conclude that almost all of its mass of $1.2-1.6 \\times 10^9$ \\msun\\ most likely resides in a warp.  Four features far from the plane containing about $10^8$ \\msun\\ are found at velocities distinct from this warp.  These clouds may be associated with the expected infall, although an origin in a disk-halo flow cannot be ruled out, except for one feature which is counter-rotating. The warp itself may be a result of infall according to recent models.  But clearly this galaxy lacks a massive, lagging neutral halo as found in NGC 891. The disk HI is concentrated into two rings of radii 1.5 and 3 arcminutes. Radial inflow is found in the disk, probably due to the bar in this galaxy.  A nearby member of this galaxy group, NGC 5740, is also detected.  It shows a prominent one-sided extension which may be the result of ram pressure stripping. ",
        "introduction": "Gaseous thick disks or halos in spiral galaxies hold promise for answering key questions about galaxy formation and evolution.  These vertically extended layers are multiphase, with detailed studies of the nearest edge-on galaxies revealing halos of neutral hydrogen \\citep[e.g.][]{1994ApJ...429..618I, 1997ApJ...491..140S}, diffuse ionized gas \\citep[DIG; e.g.][]{ 1990ApJ...352L...1R, 1990A&A...232L..15D, 2003A&A...406..505R}, hot X-ray emitting gas \\citep[e.g.][]{1994ApJ...420..570B, 2004ApJ...606..829S, 2006A&A...448...43T}, radio continuum emission \\citep[e.g.][]{ 1999AJ....117.2102I,2001A&A...374...42D} and dust \\citep[e.g.][]{1999AJ....117.2077H,2000A&AS..145...83A, 2006A&A...445..123I}.  Especially for the extended DIG, X-ray, dust and radio continuum components, it has been well established that their prominence is correlated with the star formation activity in the underlying disk \\citep[e.g.][]{1996ApJ...462..712R, 2003A&A...406..493R,1999AJ....117.2077H,2006A&A...457..779T,2006A&A...457..121D}. Large shell-like and vertically oriented filamentary structures are also seen in many of the more actively star-forming edge-ons, in, e.g., DIG \\citep {1990ApJ...352L...1R} and HI \\citep[e.g.][]{2001A&A...377..759L}.  All of this evidence has led to a picture of the origin of these layers in a star-formation-driven disk-halo cycle, which has been theoretically modeled as a general galactic fountain flow \\citep{1976ApJ...205..762S,1980ApJ...236..577B} with later models incorporating the fact that mass and energy input into the halo may efficiently occur through localized structures such as supershells and chimneys \\citep{1989ApJ...345..372N}.  As for the Milky Way, the WHAM survey \\citep{2003ApJS..149..405H} has characterized in detail the so-called Reynolds layer of vertically-extended DIG \\citep{1973ApJ...179..651R}, revealing further instances of possible superbubbles \\citep{2001ApJ...558L.101R}.  Very relevant to this paper are the High Velocity Clouds (HVCs) and Intermediate Velocity Clouds (IVCs).  For the latter, a few are known to be 0.3--4 kpc above the midplane, with metallicities close to solar, and it is quite possible that these clouds originate in a disk-halo flow (\\citealt{2004Ap&SS.289..381W}, and references therein).  The situation for the HVCs is different, however.  Not considering the contribution from the Magellanic Stream \\citep{2003ApJ...586..170P}, some well studied complexes are found to be many kpc from the plane (e.g. 8--10 kpc for part of complex A, $>$4 kpc for complex C; \\citealt {2004Ap&SS.289..381W} and references therein).  Although information is scarce, their metallicities are also well below solar (e.g. \\citealt{2007ApJ...657..271C}), suggesting that such clouds are not part of a disk-halo flow but may be extragalactic clouds that are mixing with metal-enriched gas in the halo as they fall onto the disk.  Such a population of infalling warm clouds is a prediction of recent models of galaxy formation, in which galaxies are still growing by this mechanism.  The idea that galaxies form by gas cooling out of a shock-heated hot halo was presented by \\citet{1978MNRAS.183..341W}.  This hot gas should be thermally unstable and given to fragmentation \\citep{1965ApJ...142..531F, 1985ApJ...298...18F,1990ApJ...363...50M}.  \\citet{2004MNRAS.355..694M} explored the consequences of such an inflow of fragmenting material.  The remaining hot gas has such a low density that it can support itself against infall for a long time.  This may explain the \"over-cooling\" problem in galaxy formation models that leads to overly massive galaxies and the finding that most of the baryons in the universe never ended up in galaxies (\\citealt{2003ApJ...599...38B} and references therein).  Instabilities in cooling inflows have also been studied by \\citet{2006MNRAS.370.1612K} and \\citet{2006ApJ...644L...1S}.  The cloud populations and distributions in these high-resolution models vary, with clouds having an uncertain but probably high ionization fraction.  An alternative explanation to the \"over-cooling\" problem in the most massive galaxies is feedback.  Reheating of the halo gas by core-collapse supernovae or a (possibly recurrent) AGN prevents the gas from cooling and flowing in \\citep{2004MNRAS.347.1093B}.  Heating by Type Ia supernovae is considered by \\citet{2005ASPC..331..329W}.  It is also argued that, except in the most massive halos, infalling gas never heats to the virial temperature but flows in cold \\citep{1977ApJ...215..483B,2003MNRAS.345..349B,2004MNRAS.347.1093B, 2005MNRAS.363....2K,2006MNRAS.368....2D}.   The rotation of gaseous halos may also provide clues as to their origin and the physical processes occurring within them.  In recent years, the manner in which rotation speeds change as a function of height has begun to be characterized in the DIG \\citep{2006ApJ...647.1018H, 2006ApJ...636..181H, 2007ApJ...663..933H, 2007AJ....134.1019O}.  Heald and coworkers have measured the gradient in rotation speed with height ($dV_{rot}/dz$) in three edge-ons which form a decreasing sequence of star forming activity and DIG scale-height: NGC 5775 (--8 km s$^{-1}$ kpc$^{-1}$), NGC 891 (--17 km s$^{-1}$ kpc$^{-1}$), and NGC 4302 (--30 km s$^{-1}$ kpc$^{-1}$).  The gradients, when expressed in terms of km s$^{-1}$ per DIG scale-height, show much less range: --15 to --25 km s$^{-1}$ (scale height)$^{-1}$.  The authors show that simple ballistic models of disk-halo flow \\citep{2002ApJ...578...98C} greatly underpredict the magnitude of these gradients and also predict the wrong trend with scale-height.  Two possible resolutions are: 1) the physics of disk-halo flows are not well described by ballistic models and that hydrodynamical effects such as pressure \\citep{2006A&A...446...61B} or magnetic or viscous forces \\citep{2002ASPC..276..201B} dominate the rotation, or 2) part or all of the gas does not originate in a flow from the high-angular-momentum disk, but rather from infalling low-angular-momentum primordial gas.  For the DIG, the latter explanation must also account for morphological and other connections with ongoing star formation discussed above.  For NGC 891, $dV_{rot}/dz$ has also been measured for the HI and is --15 km s$^{-1}$ kpc$^{-1}$ \\citep{2007AJ....134.1019O}, in good agreement with the DIG value .  The authors conclude from this gradient that at least some of the $1.2 \\times 10^9$ \\msun\\ HI halo must be accreted, low angular momentum material (see also \\citealt{2006MNRAS.366..449F}).  Furthering our understanding of both galactic infall and disk-halo cycling, as well as possible interactions between the two, would be made easier if one could isolate galaxies where one or the other origin is expected to obtain. High star-forming galaxies like NGC 5775 presumably have gaseous halos dominated by disk-halo flows.  The recent discovery of a bright ( $L_x = 7.3 \\pm 3.9 \\times 10^{39}$ erg s$^{-1}$) X-ray halo around the massive, nearby (29.4 Mpc is the commonly adopted distance), low star forming edge-on Sb galaxy NGC 5746 \\citep{2006NewA...11..465P,2006astro.ph.10893R} presents a challenge to disk-halo models as the X-ray luminosity clearly exceeds that expected \\citep{2006A&A...457..779T,2004ApJ...606..829S} for a galaxy with no detected DIG halo (\\citealt{1996ApJ...462..712R}; in an image with an Emission Measure noise level of 3.7 pc cm$^{-6}$) and little star formation. \\citet{1996ApJ...462..712R} roughly characterized the star formation rate per unit disk area of many edge-ons using the far infrared luminosity measured by the {\\it Infrared Astronomical Satellite} (IRAS) divided by the optical disk area: $L_{FIR}$/$D_{25}^2$ (values calculated from the NASA/IPAC Extragalactic Database\\footnote{The NASA/IPAC Extragalactic Database (NED) is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.}).  NGC 5746 is one of the lower of the 16 edge-ons thus characterized, at $4 \\times 10^{39}$ erg s$^{-1}$ kpc$^{-2}$.  Rather, the X-ray luminosity puts NGC 5746 on the expected steep relationship $L_x \\propto V_{rot}^7$ for hot halos in galaxy formation models \\citep{2002MNRAS.335..799T}, and, in such an interpretation, it is only because of the very high rotation speed (measured in this paper to be about 310 km s$^{-1}$) that this residual hot halo can be detected at all.  The metallicity of the hot gas is found to be low, at about 0.04 solar \\citep{2006astro.ph.10893R}, but is uncertain and could be biased toward low values (see \\citealt{2000MNRAS.311..176B}).  But overall, this galaxy may well be an attractive test case where disk-halo cycling is minimized.  There are some caveats to this interpretation, however.  First, NGC 5746 is in a group of 26 cataloged members \\citep{2000ApJ...543..178G}, and there may well have been interactions in the past leading to gas in the group environment.  Second, although there is little ongoing star formation, it may be that Type Ia supernovae or an AGN create a hot wind or provide energy to keep a pre-existing halo heated, as discussed above. However, \\citet{2006astro.ph.10893R} argue that the likely Type Ia SN rate in a galaxy like NGC 5746 would be far too low to explain the X-ray luminosity. They also argue that there is little evidence for at least current nuclear activity in this galaxy, and that a soft, thermal X-ray halo would not be expected for an AGN outflow anyway.  We do note, however, that NGC 5746 is classified as a LINER by \\citet{1999RMxAA..35..187C}, and there is a compact X-ray source \\citep{2006A&A...460...45G} at the center, indicating at least low level activity at present, although no compact radio source has been detected \\citep{2005A&A...435..521N}.  If the metallicity of the hot gas is indeed as low as 0.04, these alternative sources are not very likely.  If this X-ray halo is a residual of galaxy formation and thus an indication of missing baryons, then an obvious question is whether the predicted thermal instabilities are occurring, leading to an infalling warm component.  If it exists, the lack of a DIG halo may simply mean that there are few disk sources capable of significant ionization.  But there may be a neutral halo.  Either a detection or an upper limit will constrain galaxy formation models and the \\lq\\lq missing baryon\\rq\\rq\\ question, and shed light on the origin of the Milky Way's HVCs, more so if individual clouds can be seen.  A detection of a significant neutral halo would also be challenging to explain in terms of a disk-halo cycle origin.  We have therefore observed NGC 5746 in 21-cm emission with the Very Large Array (VLA), as described in \\S 2.  We analyze high-latitude emission in \\S 3, and end with a brief discussion in \\S 4 in terms of the theories of the halo gas discussed above. ",
        "conclusions": "We have observed the massive edge-on spiral NGC 5746 with the VLA, in an attempt to find a vertically extended component of atomic gas as predicted by recent galaxy formation models.  A high-latitude component has been discovered, but almost all of its mass, summing to $1.2-1.6 \\times 10^9$ \\msun\\ -- or about 15\\% of the total HI mass -- is more readily explained as a warp than as a halo.  The warp must be asymmetric in its column density distribution and geometry, but this is not unusual.  If the high-latitude component is a lagging halo, it must feature a large mass in a rather narrow radial range, be centered at a large radius, and have a much larger degree of asymmetry than the warp.  Even then, such a model has difficulty reproducing the narrowness of the high-latitude spectra as well as the warp model.  We therefore conclude that a warp is the more likely explanation.  The warp itself may be a result of infall according to recent models.  It partially resolves into clouds of mass $4 \\times 10^6$ \\msun\\ to $10^8$ \\msun, accounting for about half the mass in the warp.  We have found four high-latitude features at velocities distinct from the warp, totaling about 10$^8$ \\msun.  These could be accreting onto the disk, or they may originate in a disk-halo flow, although there is no other sign of such a flow in this galaxy with little star forming activity.  One cloud is counter-rotating and must have an external origin.  The cloud properties are roughly comparable to those expected in the galaxy formation models of \\citet{2006MNRAS.370.1612K} and \\citet{2006ApJ...644L...1S} where warm clouds form by thermal instabilities in the residual hot halos.  However, these models have not been tailored to galaxies as massive as NGC 5746, and it would be interesting to compare our results with such a model.  Given the calculated pressure in the hot halo, we expect any clouds to be primarily neutral. Alternatively, we cannot rule out a weak disk-halo flow as the origin for these clouds, except for the counter-rotating one.  How neutral gaseous halos relate to other galaxian properties needs to be tested with more observations.  The disk shows the signature of radial inflow at the level of about 10 km s$^{-1}$.  This is most likely due to the bar, which is known to have a strong kinematical signature from optical long-slit spectra.  The group member NGC 5740 is also detected, albeit outside the half-power points of the primary beam.  The most important result here is a broad extension of gas to the NE, which is more likely to be due to ram pressure than tidal stripping, although whether the condition for such stripping is met is unknown."
    },
    "0801/0801.1633_arXiv.txt": {
        "abstract": "{{\\it Aims.} We aim to characterise the morphology and the physical parameters governing the shock physics of the Herbig-Haro object HH99B. We have obtained SINFONI-SPIFFI IFU spectroscopy ({\\it R} $\\sim$ 2000 $-$ 4000) between  1.10 and 2.45 $\\mu$m detecting more than 170 emission lines, to a large extent never observed before in a Herbig-Haro object. Most of them come from ro-vibrational transitions of molecular hydrogen (v$_{up}$\\,$\\le$\\,7, E$_{up}\\la$ 38~000 K) and [\\fe] (E$_{up}\\la$ 30~000 K). In addition, we observed several hydrogen and helium recombination lines, along with fine structure lines of ionic species. All the brightest lines appear resolved in velocity.\\\\ {\\it Methods.} Intensity ratios of ionic lines have been compared with predictions of NLTE models to derive bi-dimensional maps of extinction and electron density, along with estimates of temperature, fractional ionisation and atomic hydrogen post-shock density. H$_2$ line intensities have been interpreted in the framework of Boltzmann diagrams, from which we have derived maps of extinction and temperature of the molecular gas. From the intensity maps of bright lines (i.e. H$_2$ 2.122$\\mu$m and [\\fe] 1.644$\\mu$m) the kinematical properties of the shock(s) at work in the region have been delineated. Finally, from selected [\\fe] lines, constraints on the spontaneous emission coefficients of the 1.257, 1.321 and 1.644 $\\mu$m lines are provided.\\\\ {\\it Results.} Visual extinction variations up to 4 mag emerge, showing that the usual assumption of constant extinction could be critical. The highest A$_V$ is found at the bow-head (A$_V$ $\\sim$ 4~mag) while diminishing along the flanks. The electron density increases from $\\sim$ 3~10$^3$ cm$^{-3}$ in the receding parts of the shock to $\\sim$ 6~10$^3$ cm$^{-3}$ in the apex, where we estimate, from [\\fe] line ratios, a temperature of $\\sim$ 16~000 K. Molecular gas temperature is lower in the bow-flanks (T$\\sim$ 3000 K), then progressively increasing toward the head up to T$\\sim$ 6000 K. In the same zone, we are able to derive from the [\\fe]1.257/[\\pii]1.187 line ratio, the iron gas-phase abundance ($\\sim$ 60\\% of the solar value) along with the hydrogen fractional ionisation (up to 50\\% at the bow-head) and the atomic hydrogen post-shock gas density ($\\sim$ 1 10$^4$ cm$^{-3}$). The kinematical properties derived for the molecular gas substantially confirm those published in Davis et al.(1999), while new information (e.g. v$_{shock}$ $\\sim$ 115 km s$^{-1}$) is provided for the shock component responsible for the ionic emission. We also provide an indirect measure of the H$_2$ breakdown speed (between 70 and 90 km s$^{-1}$) and compute the inclination angle with respect to the line of sight. The map parameters, along with images of the observed line intensities, will be used to put stringent constraints on up-to-date shock models. ",
        "introduction": "Mass loss phenomena in the form of powerful bipolar jets and molecular outflows are often associated with the early evolution of protostars. Together with indirectly regulating the accretion process, they play a crucial r\\^{o}le also in the interaction between the protostar and the natal environment, causing injection of  momentum, kinetic energy and turbulence in the ISM, along with irreversible modifications of its chemical structure and physical conditions. The most violent interaction occurs at the terminal working surface, where the supersonic flow impacts the undisturbed medium and most of the ambient material is entrained by the jet (e.g. Reipurth \\& Bally, 2001). In a schematic model, the interaction occurs via two shocks: an internal working surface (Mach disk) which decelerates the jet gas and a forward shock which accelerates the ambient gas producing a mixture of shock velocities (e.g. Hartigan, 1989). This latter has often a curve-shaped morphology: therefore, only the component of motion along the jet axis is slowed going from the apex of the shock toward the receding parts, where large transverse motions occur. In a classical scenario, for sufficiently fast shocks, the head of the bow is a pure dissociative (J)ump-type shock (Hollenbach \\& McKee, 1989), which changes along the bow flanks, where the impact occurs obliquely, to slower, non-dissociative (C)ontinuos-type shocks (Draine, 1980). In this framework, highly excitated ionic emission should arise at the head of the bow, while molecular emission mainly originates from the cooling regions along the flanks and behind the bow. Recent models, however, predict that mixtures of J- and C-type shocks can occur along the overall structure of the bow, or, in the presence of a sufficiently strong magnetic field, that a J-type shock can evolve into a C-type, remaining embedded at early time (Smith \\& Rosen, 2003; Smith \\& Mac Low, 1997; Flower et al., 2003; Le Bourlot et al., 2002, Lesaffre et al., 2004a,b).  The large number of parameters which regulate the shock physics (e.g. strength and direction of the local magnetic field, ionisation fraction and density gradient between the shocked gas and the local environment) can be constrained only through dedicated observations aimed at probing the physical and kinematical properties of the gas all along the bow structure. In this respect, observations in the near-infrared represent a well suited tool: both H$_2$ ro-vibrational lines, which are the main gas coolants in continuos shock components, and fine structure transitions of abundant atomic species (e.g. {\\ion{Fe}{ii}}, {\\ion{C}{i}}, {\\ion{S}{ii}}), which are expected to emit in dissociative shocks, fall in this wavelength range. Indeed, in recent years multiple observations of jets and bow shocks have been conducted in the near-infrared by means of long-slit spectroscopy (e.g. Eisl\\\"{o}ffel, Smith \\& Davis, 2000; Giannini et al., 2004; Smith, Froebrich \\& Eisl\\\"{o}ffel, 2003; Nisini et al., 2002; O'Connell, Smith \\& Davis, 2004; Nisini et al., 2005). The observed spectra are typically rich in both molecular (H$_2$) and ionic ([FeII]) lines, implying that strong gradients in the local conditions (e.g. temperature, fractional ionisation) do occur and that multiple shock components are simultaneously at work.\\\\ The main limitation of long-slit spectroscopy for studying in detail bow-shock morphologies arises from the poor coverage of the extended bow surface usually obtainable in a reasonable amount of observing time; on the contrary, Integral Field Spectroscopy (IFU) represents a well-tailored tool to overcome this problem, since it allows us to obtain simultaneously 2-D maps at different wavelengths. In this paper we present the spectral images obtained with the IFU facility SINFONI (Eisenhauer et al., 2003, Bonnet et al., 2004) of a prototype bow-shock, namely the Herbig-Haro object HH99B. This is located in the RCrA molecular core at d$\\sim$ 130 pc (Marraco \\& Rydgren, 1981), and it was firstly discovered in the optical by Hartigan \\& Graham (1987), who suggested it is the red-shifted lobe of the outflow powered by the HH100-IR source. More recently, Wilking et al.(1997) have proposed, on the basis of near-infrared images, the infrared source IRS9 and the Herbig Ae star RCrA as other possible exciting source candidates. HH99B was firstly imaged in the near-infrared by Davis et al. (1999, hereafter D99) who identified three different emission zones: one at the head of the bow (subsequently named B0 by M$^c$Coey et al., 2004, hereafter MC04) where the bulk of the emission comes from ionised gas, and two bow-flanks (knots B1 and B3), which emit mainly in H$_2$ lines. A further H$_2$ knot, immediately behind the bow apex, was identified as knot B2. In the framework of both bow- (D99) and planar- (MC04) shock models, some attempt has been made to model the line emission observed in HH99B: the H$_2$ morphology is well fitted by a C-type bow shock, while the physical conditions of the molecular gas (measured  by means of long-slit spectroscopy) have been accounted for by a planar J-type shock with a magnetic precursor. None of these two models, however, is able to reproduce the copious ionic emission, that requires the presence of a further fully dissociative shock component.\\\\ With the present work we aim at putting strong observational constraints on bow-shock models. (whose detailed application will be subject of a separated paper). In particular we intend to: {\\it (i)} morphologically characterise the emission of the different lines; {\\it (i)} derive maps of the main physical parameters which govern the shock physics; {\\it (iii)} study the velocity field(s) along the bow structure.\\\\  Our work is organised as follows: Section 2 describes our observations and the obtained results; the line excitation analysis and the kinematical properties are then presented in Sections 3 and 4. Concluding remarks are given in Section 5, while in Appendix A we describe the procedure we have applied to derive the inclination angle of HH99B. ",
        "conclusions": "We have presented bi-dimensional, deep near-infrared spectral images of the bow shock HH99B. These have allowed us, for the first time, to accurately derive the physical parameters of both the molecular and ionic gas components (summarised in Table\\,\\ref{tab:tab4}, where our results are compared with those derived in previous works), and, at the same time, to characterise the geometry and the kinematical properties of the flow. The main results of this study are the following: \\begin{itemize} \\item[-] More than 170 emission lines have been detected, mainly ro-vibrational H$_2$ and [\\fe] lines, many of them never observed before in an Herbig-Haro object. In addition, transitions of hydrogen and helium recombination and fine structure lines of [\\pii], [\\tiii], and possibly [\\coii] have been observed. \\item[-] A clear bow-shape morphology emerges from the line intensity maps. As shown in Figures\\, \\ref{fig:spettro_j}-\\ref{fig:spettro_k}, [\\fe] and other ionic emission peaks definitely at the bow-head (B0), being strong also in the knot B2 immediately behind. In contrast, H$_2$ emission delineates the bow flanks, peaking in the knots B1 and B3. Notably, the H$_2$ lines with the highest excitation energy (E$_{up}$ $\\ga$ 30\\,000 K) show a different morphology, being strong towards the bow-head. This implies that H$_2$ still survives in this zone, in spite of the significant temperature enhancement there. \\item[-] Extinction maps have been derived from the analysis of both [\\fe] and H$_2$ lines. These give similar results, with A$_V$ ranging between 1 and 4 mag. \\item[-] A detailed electron density map has been obtained in the framework of NLTE approximation for [\\fe] line emission. This remains almost constant in the [\\fe] emission zone, peaking towards the bow-head. From the same emission, we are able to probe a variation of the electron temperature, which falls from $\\sim$ 16\\,000 K at the apex to less than 10\\,000 K in the receding parts of the bow. \\item[-] An iron depletion degree not higher than 30\\% has been inferred at the bow apex, which testifies in favour of a J-type shock as the main excitation mechanism in this part of the bow. In this same zone, we infer a fractional ionisation of $\\sim$ 0.6 and a post-shock density of $\\sim$ 10$^4$ cm$^{-3}$. \\item[-] Analysis of H$_2$ line emission has allowed us to probe the molecular temperature variation. We find that at the bow apex thermalisation has been reached at T$\\sim$ 6\\,000 K, likely due to the presence of a fast, non-dissociative shock. On the contrary, along the flanks different temperature components are simultaneously present. A decrease of about a factor of ten is registered in the H$_2$ column density going from knots B1/B3 toward knot B0. Both these circumstances are not accounted for by models which have attempted to interpret the H$_2$ emission on HH99B on the basis of fewer lines. Therefore, the conclusions of such models should be tested in the light of this new piece of information. \\item[-] From the brigthest [\\fe] and H$_2$ lines we have been able to probe the kinematical properties (e.g. shock velocity) of the shocked gas. In particular, we confirm the result by D99 according to which HH99 is a red-shifted, receding bow. \\item[-] The radial velocity map of [\\fe] emission has been interpreted in the framework of the bow geometry. From this map we have consistently inferred the bow inclination angle and defined a range of 70-90 km s$^{-1}$ for the H$_2$ breakdown velocity. These values are in agreement with the prediction of Le Bourlot et al.(2002) of v$_{dis}$ up to 80 km s$^{-1}$. We propose our method as a valuable tool to derive the jet inclination angle (if larger than 10-20$^{\\circ}$) in cases where proper motion is unknown. \\item[-] The kinematical parameters of the [\\fe] emission estimated in this work do not confirm the model predictions by MC04. In particular, the argument that the [\\fe] lines originate in a pure J-type shock with $v_{shock}$ $\\sim$ 50 km s$^{-1}$ contrasts with our measure of $v_{shock}$ $\\sim$ 110-120 km s$^{-1}$. Thus, the modelling of the [\\fe] emission may have to be revised. \\end{itemize}"
    },
    "0801/0801.3400_arXiv.txt": {
        "abstract": "{Extended Red Emission (ERE) was recently attributed to the photo-luminescence of either doubly ionized Polycyclic Aromatic Hydrocarbons (PAH$^{++}$), or charged PAH dimers ($[$PAH$_{2}$]$^+$).} {We analysed the visible and mid-infrared (mid-IR) dust emission in the North-West and South photo-dissociation regions of the reflection nebula NGC 7023.} {Using a blind signal separation method, we extracted the map of ERE from images obtained with the Hubble Space Telescope, and at the Canada France Hawaii Telescope. We compared the extracted ERE image to the distribution maps of the mid-IR emission of Very Small Grains (VSGs), neutral and ionized PAHs (PAH$^0$ and PAH$^+$) obtained with the Spitzer Space Telescope and the Infrared Space Observatory.} {ERE is dominant in transition regions where VSGs are being photo-evaporated to form free PAH molecules, and is not observed in regions dominated by PAH$^+$. Its carrier makes a minor contribution to the mid-IR emission spectrum.} {These results suggest that the ERE carrier is a transition species formed during the destruction of VSGs. $[$PAH$_{2}$]$^+$ appear as good candidates but PAH$^{++}$ molecules seem to be excluded.} ",
        "introduction": "\\label{int}   Unveiling the composition, structure and charge state of the smallest interstellar dust particles remains one of today's challenges in astrochemistry. Progress in this field requires a detailed analysis of the spectral signatures of these dust populations. Amongst these signatures is the Extended Red Emission (ERE), a broad emission feature ranging from 540 to beyond 900 nm and with a peak wavelength longward of 600 nm to beyond 800 nm \\citep{smi02b}. ERE is likely due to the photo-luminescence of carbonaceous macromolecules or nanograins exposed to UV photons (see Witt et al. 2006 and references therein). The ERE is observed in many environments exposed to UV photons (Photo-dissociation Regions; PDR) found for instance in the diffuse interstellar medium, reflection nebulae and planetary nebulae. It has been observed in galaxies, NGC 3034 \\citep{per95} and NGC 4826 by \\citep{pie02}.  The possible link with the carriers of the Aromatic Infrared Bands (AIBs; also called unidentified infrared bands at 3.3, 6.2, 7.7, 8.6 and 11.3 \\,$\\mu$m) has been discussed by several authors. Both types of  emission features, ERE and AIBs, were found to be relatively cospatial although not matching \\citep{fur90}, pointing to different but related materials for their carriers \\citep{fur92}. Recently, a detailed study of the spatial distribution of the ERE in the northern PDR of NGC\\,7023 has been performed by \\citet{wit06} who concluded that the ERE mechanism is a two-step  process involving the formation of the carrier and then the excitation of the luminescence, and proposed doubly-ionized Polycyclic Aromatic Hydrocarbons (PAHs) as plausible carriers. On the other hand, recent quantum chemistry calculations point to PAH dimers ($[$PAH$_{2}$]$^+$) as the carrier of ERE \\citep{rhe07}.  In this letter, we provide an efficient and non biased way to extract the ERE map in NGC 7023 from the Hubble Space Telescope (HST) data and new Canada France Hawaii Telescope (CFHT) images using a blind signal separation method. We then compare the extracted ERE map to the maps of the different mid-IR emission carriers, namely Very Small Grains (VSGs), neutral and ionized PAHs (PAH$^0$ and PAH$^+$), as extracted by  \\citet{rap05} and \\citet{ber07} from the Infrared Space Observatory (ISO) and Spitzer Space Telescope (\\emph{Spitzer}) observations.  \\begin{figure} \\begin{center} \\includegraphics[width=\\hsize]{9158fig1.ps} \\vspace{-0.0cm} \\caption{On the left: HST images of the NGC 7023 North-West PDR in three SDSS wide-band filters (cf. Witt et al. 2006). On the right: scattered light and ERE images extracted with \\emph{FastICA} from the observations} \\end{center} \\label{ext} \\vspace{-0.5cm} \\end{figure} ",
        "conclusions": "\\label{con} In this letter we have compared in NGC 7023 the spatial distribution of ERE and that of the mid-IR emitters : VSGs,  PAH$^0$,  PAH$^+$. We find strong evidence that the ERE carrier is produced in the region of destruction of VSGs and show that it has a negligible contribution to the mid-IR emission. We show that PAH$^{++}$ are unlikely to be the carrier of ERE and conclude that  $[$PAH$_{2}$]$^+$ are attractive candidates, on the basis of qualitative arguments on the balance between photodissociation and reformation of these clusters. Further investigations are needed, following the strategy presented in this letter, but in other PDRs. Spectro-imagery of such regions in the mid-IR with a high spatial resolution is needed in order to be able to trace the thin frontier between PAHs and VSGs where small clusters are likely to be present. Finally, laboratory spectroscopic studies on $[$PAH$_{2}$]$^+$ are needed to progress in this identification. In particular,  closed-shell cation dimers as proposed by \\citet{rhe07} deserve additional studies."
    },
    "0801/0801.4827_arXiv.txt": {
        "abstract": "We investigate clustering properties of Lyman-break galaxies (LBGs) at $z\\sim 3$ based on deep multi-waveband imaging data from optical to near-infrared wavelengths in the Subaru/\\textit{XMM-Newton} Deep Field. The LBGs are selected by \\uv and \\vz\\ colors in one contiguous area of 561 arcmin$^2$ down to $z'=25.5$. We study the dependence of the clustering strength on rest-frame UV and optical magnitudes, which can be indicators of star formation rate and stellar mass, respectively. The correlation length is found to be a strong function of both UV and optical magnitudes with brighter galaxies being more clustered than faint ones in both cases. Furthermore, the correlation length is dependent on a combination of UV and optical magnitudes in the sense that galaxies bright in optical magnitude have large correlation lengths irrespective of UV magnitude, while galaxies faint in optical magnitude have correlation lengths decreasing with decreasing UV brightness. These results suggest that galaxies with large stellar masses always belong to massive halos in which they can have various star formation rates, while galaxies with small stellar masses reside in less massive halos only if they have low star formation rates. There appears to be an upper limit to the stellar mass and the star formation rate which is determined by the mass of hosting dark halos. ",
        "introduction": "\\label{sec:intro}  Galaxy evolution in dark halos is a central issue of cosmology in the Cold Dark Matter (CDM) universe. A critical key to the problem is to reveal what kind of dark halos host what kind of galaxies, specifically, the relationship between the mass of dark halos and fundamental quantities of galaxies such as star formation rate and stellar mass. The clustering strength of galaxies can be used to infer the mass of hosting dark halos, since in the standard theoretical framework of galaxy evolution, the large-scale distribution of galaxies is determined by the distribution of underlying dark halos, whose spatial clustering depends on their mass, with more massive halos having stronger spatial clustering \\citep{mo1996}.  A number of studies have examined clustering properties of various high-redshift galaxies \\citep{giavalisco1998,giavalisco2001,foucaud2003,daddi2003, ouchi2004a,ouchi2005,adelberger2005a,lee2006,kashikawa2006, hildebrandt2007,quadri2007,ichikawa2007}. Measuring clustering strength requires a large galaxy sample from a wide sky coverage. One considerable case in which the relationship between dark halos and galaxies has been successfully explored by means of clustering analysis is that of Lyman-break galaxies (LBGs). LBGs, which are selected by their continuum features in rest-frame UV spectra (the Lyman-break, the Ly$\\alpha$ forest, and otherwise nearly flat continuum longward of the Ly$\\alpha$) redshifted into optical bandpasses \\citep{guhathakurta1990, steidel1996}, are normal, young star-forming galaxies with modest dust extinction. Requiring only optical imaging in a few bands, the LBG selection method has yielded the largest and well-controlled samples of galaxies in the young universe at $2\\lesssim z\\lesssim 6$ \\citep[e.g.,][]{steidel1999,steidel2003,ouchi2004b,giavalisco2004, dickinson2004,sawicki2006,yoshida2006,iwata2007,bouwens2007}. LBGs are as numerous as the present-day galaxies, suggesting that they play a significant role in the early stage of galaxy evolution. For a review of LBGs including the history of their discovery, see \\citet{giavalisco2002}.  An important piece of evidence which links the mass of hosting dark halos to physical properties of LBGs has been first discovered for the UV luminosity; the clustering strength of LBGs increases strongly with UV luminosity \\citep[e.g.,][]{giavalisco2001,ouchi2004a,adelberger2005a, lee2006,hildebrandt2007}. Since UV luminosity is sensitive to star formation rate, this implies that the star formation activity of LBGs is somehow controlled by the mass of dark halos which host them. To make a next step forward in our understanding of how the evolution of galaxies is related to the mass of dark halos, it would be necessary to examine the clustering strength as a function of other properties of galaxies like stellar mass, age, and dust extinction. In particular, stellar mass is a robust quantity in terms of the star formation history of galaxies compared to star formation rate, which can vary with time. This kind of analysis is, however, not easy because it requires deep and wide-field observational data covering wavelengths simultaneously from optical to near-infrared. Only the dependence of clustering strength on rest-frame optical luminosity has been examined by \\citet{adelberger2005b} for 300 galaxies at $z\\sim 2$ (BX objects) selected by UV wavelength in a similar manner to LBGs.  Recently, clustering properties of near-infrared selected galaxies at high redshift have been examined against various quantities of galaxies \\citep{quadri2007,ichikawa2007}. The selection based on a near-infrared magnitude allows us to compile a nearly stellar mass-selected sample. However, the studies at high redshift on the basis of near-infrared selected galaxies have so far been inevitably limited to shallow \\citep[$K\\lesssim 23.0$;][]{quadri2007} or small-field \\citep[$\\lesssim 25$ arcmin$^2$;][]{ichikawa2007} samples. This is due to the lower sensitivity and smaller area coverage of near-infrared cameras compared to optical ones.  This paper studies the dependence of clustering strength on both UV (apparent $z'$ band) and optical (apparent $K$ band) luminosities for LBGs at $z\\sim 3$ to investigate how the stellar mass and the star formation rate of galaxies depend on the mass of dark halos which host them at high redshift. The redshift $z\\sim 3$ is the best target for such studies, since it is the highest redshift at which ground-based near-infrared imaging can access rest-frame optical wavelengths. To construct a large sample of LBGs at $z\\sim 3$, we performed $U$-band imaging observation with Subaru/Suprime-Cam in the Subaru/\\textit{XMM-Newton} Deep Field (SXDF), which has a set of very deep optical multi-waveband imaging data from the Subaru/\\textit{XMM-Newton} Deep Survey \\citep[SXDS;][]{sekiguchi}, and deep $J$ and $K$ data taken with the wide-field near-infrared camera WFCAM on UKIRT by the UKIDSS Ultra Deep Survey \\citep[UDS;][]{lawrence2007,warren2007}.  The outline of this paper is as follows: In \\S\\ref{sec:data}, a brief account of the observations and the data is presented. A large sample of LBGs at $z\\sim 3$ is constructed in \\S\\ref{sec:lbg}. Simulations to assess the redshift distribution function and the contamination by interlopers of the sample are also described. Clustering analysis is made in \\S\\ref{sec:acf}. A summary and conclusions are given in \\S\\ref{sec:sum}.  Throughout this paper, the photometric system is based on AB magnitude \\citep{oke1983}. The cosmology adopted is a flat universe with $\\mathrm{\\Omega}_m = 0.3$, $\\mathrm{\\Omega}_\\mathrm{\\Lambda} = 0.7$, $\\sigma_8 = 0.9$, baryonic density $\\mathrm{\\Omega}_b = 0.04$, and a Hubble constant of $H_0 = 100\\ h$ km s$^{-1}$ Mpc$^{-1}$ with $h=0.7$. The correlation length is expressed in units of $h^{-1}$ Mpc to facilitate comparison with previous results. ",
        "conclusions": "\\label{sec:sum}  We have investigated clustering properties of Lyman-break galaxies (LBGs) at $z\\sim 3$ based on deep multi-waveband imaging data from optical to near-infrared wavelengths in the Subaru/\\textit{XMM-Newton} Deep Field. The LBGs are selected by \\uv and \\vz\\ colors in one contiguous area of 561 arcmin$^2$ down to $z'=25.5$. The number of LBG candidates detected is 962 in total, among which $708$ are found in the region observed in the $J$ and $K$ bands. We use Monte Carlo simulations to estimate the redshift distribution function and the fraction of contamination by interlopers of the LBG samples. The fraction is found to be at most 6\\% at any magnitude.  We explored the dependence of the clustering strength on rest-frame UV and optical magnitudes, which can be indicators of star formation rate and stellar mass, respectively. The correlation length is found to be a strong function of both UV and optical magnitudes with brighter galaxies being more clustered than faint ones in both cases. It is found that the correlation length also increases with $z'-K$ color, which is correlated with $K$ magnitude in our LBG sample. These results imply that galaxies with larger star formation rates and larger stellar masses are hosted by more massive dark halos.  Furthermore, the correlation length is interestingly dependent on a combination of UV and optical magnitudes in the sense that galaxies bright in optical magnitude have large correlation lengths irrespective of UV magnitude, while galaxies faint in optical magnitude have correlation lengths decreasing with UV magnitude. One implication of this result in terms of the stellar mass assembly of galaxies is that galaxies which have accumulated large stellar masses always belong to massive halos in which they can have various star formation rates, while galaxies with small stellar masses reside in less massive halos only if they have low star formation rates. To put it another way, massive halos can host any galaxies from ones with small stellar masses to ones with large stellar masses, while in less massive halos, only galaxies with small stellar masses reside. There appears to be an upper limit to the stellar mass accumulated which is determined by the mass of hosting dark halos. Moreover, galaxies in massive halos can have various star formation rates whatever stellar masses they have at $z\\sim 3$. On the other hand, galaxies in less massive halos accumulated only small stellar masses and have lower star formation rates at $z\\sim 3$. It is suggested that the mass of dark halos governs the current star formation acitivity as well as the past star formation history at $z\\sim 3$.  \\begin{deluxetable}{cccc}[p] \\tablecolumns{4} \\tablewidth{0pt} \\tablecaption{Clustering Measurements \\label{tbl:clustering}} \\tablehead{ \\colhead{\\hspace{0.3cm} Sample\\hspace{0.3cm} } & \\colhead{\\hspace{0.5cm} $N$\\tablenotemark{\\ast}\\hspace{0.5cm} } & \\colhead{\\hspace{0.3cm} $A_\\omega$\\hspace{0.3cm} } & \\colhead{\\hspace{0.3cm} $r_0$ ($h^{-1}$Mpc)\\hspace{0.3cm} } } \\startdata $23.0<z'\\le 25.5$  & 962 & $1.17^{+0.29}_{-0.29}$ & \\phn$5.5^{+0.8}_{-0.9}$ \\\\ \\\\ $23.0<z'\\le 24.5$ & 288 & $3.44^{+1.00}_{-1.00}$ & $10.7^{+1.9}_{-2.1}$ \\\\ $24.0<z'\\le 25.0$ & 467 & $1.46^{+0.56}_{-0.56}$ & \\phn$6.3^{+1.4}_{-1.6}$ \\\\ $24.5<z'\\le 25.5$ & 674 & $0.76^{+0.36}_{-0.35}$ & \\phn$4.2^{+1.1}_{-1.3}$ \\\\ \\\\ \\phantom{$22.96<$}$K\\le 23.46$ & 171 & $3.85^{+1.28}_{-1.28}$ & $11.5^{+2.3}_{-2.6}$ \\\\ $22.96<K\\le 23.96$             & 225 & $1.80^{+0.92}_{-0.92}$ & \\phn$7.1^{+2.1}_{-2.6}$ \\\\ $23.46<K$\\phantom{$\\le 23.96$} & 537 & $0.96^{+0.39}_{-0.39}$ & \\phn$4.8^{+1.1}_{-1.3}$ \\\\ \\\\ \\zk$\\le 0.8$ & 365 & $1.14^{+0.52}_{-0.52}$ & \\phn$5.4^{+1.4}_{-1.7}$ \\\\ \\zk$>0.8$    & 343 & $2.50^{+0.78}_{-0.79}$ & \\phn$8.8^{+1.6}_{-1.9}$ \\\\ \\\\ $23.0<z'\\le 24.5$,\\ \\ \\ $K\\le 23.46$ & 105     & $4.63^{+2.14}_{-2.13}$ & $12.9^{+3.5}_{-4.1}$ \\\\ $24.5<z'\\le 25.5$,\\ \\ \\ $K\\le 23.46$ & \\phn 66 & $3.98^{+3.29}_{-3.29}$ & $11.7^{+5.4}_{-7.8}$ \\\\ $23.0<z'\\le 24.5$,\\ \\ \\ $23.46<K$    & \\phn 96 & $4.41^{+2.49}_{-2.49}$ & $12.5^{+4.0}_{-5.1}$ \\\\ $24.5<z'\\le 25.5$,\\ \\ \\ $23.46<K$    & 441     & $0.91^{+0.51}_{-0.52}$ & \\phn$4.7^{+1.5}_{-1.9}$ \\\\ \\enddata \\tablenotetext{\\ast\\ }{Number of LBGs contained in the sample.} \\end{deluxetable}"
    },
    "0801/0801.2589_arXiv.txt": {
        "abstract": "Analysis of ten years of high-precision timing data on the millisecond pulsar \\psr\\ has resulted in a model-independent kinematic distance based on an apparent orbital period derivative, \\pbdot, determined at the $1.5\\%$ level of precision ($D_{\\rm k} = 157.0\\pm 2.4$\\,pc), making it one of the most accurate stellar distance estimates published to date.  The discrepancy between this measurement and a previously published parallax distance estimate is attributed to errors in the DE200 Solar System ephemerides. The precise measurement of \\pbdot\\ allows a limit on the variation of Newton's gravitational constant, $|\\dot{G}/G| \\leq 23 \\times 10^{-12}$\\,yr$^{-1}$.  We also constrain any anomalous acceleration along the line of sight to the pulsar to $|a_{\\odot}/c| \\leq 1.5\\times 10^{-18}$\\,s$^{-1}$ at $95\\%$ confidence, and derive a pulsar mass, $m_{\\rm psr} = 1.76 \\pm 0.20\\,M_{\\odot}$, one of the highest estimates so far obtained. ",
        "introduction": "In \\citeyear{jlh+93}, \\citeauthor{jlh+93} reported the discovery of \\psr, the nearest and brightest millisecond pulsar known. Within a year, the white dwarf companion and pulsar wind bow shock were observed \\citep{bbb93} and pulsed X-rays were detected \\citep{bt93a}. The proper motion and an initial estimate of the parallax were later presented along with evidence for secular change in the inclination angle of the orbit due to proper motion \\citep{sbm+97}. Using high time resolution instrumentation, the three-dimensional orbital geometry of the binary system was determined, enabling a new test of general relativity \\citep[GR;][]{vbb+01}. Most recently, multi-frequency observations were used to compute the dispersion measure structure function \\citep{yhc+07}, quantifying the turbulent character of the interstellar medium towards this pulsar.  The high proper motion and proximity of \\psr\\ led to the prediction \\citep{bb96} that a distance measurement independent of parallax would be available within a decade, when the orbital period derivative (\\pbdot) would be determined to high accuracy. Even if the predicted precision of $\\sim$1\\% would not be achieved, such a measurement would be significant given the strong dependence of most methods of distance determination on relatively poorly constrained models and the typically large errors on parallax measurements. Even for nearby stars, both the Hubble Space Telescope and the Hipparcos satellite give typical distance errors of $3\\%$ \\citep{val07} and so far only two distances beyond $100$\\,pc have been determined at $\\sim1\\%$ uncertainty \\citep{tlmr07}. This kinematic distance is one of the few model-independent methods that does not rely upon the motion of the Earth around the Sun.  As demonstrated by \\citet{dt91}, \\pbdot\\ can also be used to constrain the variation of Newton's gravitational constant. The best such limit from pulsar timing to date \\citep[$|\\dot{G}/G| = (4\\pm5)\\times 10^{-12}$\\,yr$^{-1}$ from PSR B1913+16]{tay93} is compromised due to the poorly constrained equation of state for the neutron star companion \\citep{nor90}. The slightly weaker but more reliable limit of $|\\dot{G}/G| = (-9\\pm 18) \\times 10^{-12}$\\,yr$^{-1}$ \\citep[from PSR B1855+09, which has a white dwarf companion]{ktr94} should therefore be considered instead. A more stringent limit can be obtained from the \\citet{sns+05} timing of PSR J1713+0747: $|\\dot{G}/G| = (2.5\\pm 3.3) \\times 10^{-12}$\\,$yr^{-1}$ (at 95\\% certainty). This limit is, however, based upon the formal errors of \\pbdot, $P_{\\rm b}$ and parallax, which are easily underestimated by standard methodologies, as we shall demonstrate later. Because of this we believe the \\citet{sns+05} limit is probably underestimated, but still of relevance.  However, none of these limits are as strong as that put by lunar laser ranging \\citep[LLR;][]{wtb04}: $\\dot{G}/G = (4\\pm 9) \\times 10^{-13}$\\,yr$^{-1}$. Besides limiting alternative theories of gravity, bounds on $\\dot{G}$ can also be used to constrain variations of the Astronomical Unit ($AU$). Current planetary radar experiments \\citep{kb04} have measured a significant linear increase of $d AU/dt =0.15\\pm 0.04$\\,m yr$^{-1}$, which may imply $\\dot{G}/G = (-10 \\pm 3)\\times 10^{-13}$\\,yr$^{-1}$, just beyond the sensitivity of the limits listed above.  As mentioned before, the equation of state for dense neutron star matter is very poorly constrained.  Pulsar mass determinations can probe the range of permissible pulsar masses and thereby limit possible equations of state \\citep{lp07}.  Presently, only the pulsars NGC 6440B, Terzan 5\\,I and Terzan 5\\,J have predicted masses higher than the typical value of $1.4\\,M_{\\odot}$\\citep{frb+07,rhs+05}; however, as discussed in more detail in Section \\ref{Mass}, such predictions do not represent objective mass estimates.  The remainder of this paper is structured as follows: Section \\ref{Obs} describes the observations, data analysis and general timing solution for \\psr. Section \\ref{Dist} describes how the measurement of \\pbdot\\ leads to a new and highly precise distance. In Section \\ref{GandA}, this measurement is combined with the parallax distance to derive limits on $\\dot{G}$ and the Solar System acceleration. Section \\ref{Mass} presents the newly revised pulsar mass and our conclusions are summarised in Section \\ref{Conc}. ",
        "conclusions": "\\label{Conc}  We have presented results from the highest-precision long-term timing campaign to date. With an overall residual rms of 199\\,ns, the 10\\,years of timing data on \\psr\\ have provided a precise measurement of the orbital period derivative, \\pbdot, leading to the first accurate kinematic distance to a millisecond pulsar: $D_{\\rm k} = 157.0\\pm 2.4$\\,pc. Application of this method to other pulsars in the future can be expected to improve distance estimates to other binary pulsar systems \\citep{bb96}.  Another analysis based on the \\pbdot\\ measurement places a limit on the temporal variation of Newton's gravitational constant. We find a bound comparable to the best so far derived from pulsar timing: $\\dot{G}/G = (-5\\pm18)\\times 10^{-12}\\,{\\rm yr}^{-1}$.  An ongoing VLBI campaign on this pulsar is expected to improve this limit, enabling an independent confirmation of the LLR limit.  Previous estimates of the mass of \\psr\\ have been revised upwards to $m_{\\rm psr} = 1.76 \\pm 0.20\\, M_{\\odot}$, which now makes it one of the few pulsars with such a heavy mass measurement.  A new generation of backend instruments, dedicated observing campaigns, and data prewhitening techniques currently under development should decrease the error in this measurement enough to significantly rule out various equations of state for dense nuclear matter."
    },
    "0801/0801.0290_arXiv.txt": {
        "abstract": "We constrain the thermal evolution of the universe with a decaying cosmological term by using the method of the analysis for the Wilkinson Microwave Anisotropy Probe (WMAP) observation data.  The cosmological term is assumed  to be a function of the scale factor that increases toward the early universe, and the radiation energy density is lower compared to that in the model with the standard cosmological constant ($\\Lambda$CDM). The decrease in the radiation density affects the thermal history of the universe; e.g. the photon decoupling occurs at higher-$z$ compared  to the case of the standard $\\Lambda$CDM model. As a consequence, a decaying cosmological term affects the cosmic microwave background (CMB)  anisotropy. Thanks to the Markov-Chain Monte Carlo method, we compare the angular power spectrum in the decaying $\\Lambda$CDM model with the CMB data, and we get severe constraints on parameters of the model. ",
        "introduction": "Recent astronomical observations such as high redshift type Ia supernovae (SNIa) \\cite{Perlmutter1998, Riess2007}, cosmic microwave background (CMB) anisotropy \\cite{WMAP2006}, suggest the existence of dark energy strongly. Although many researchers have investigated dark energy, its nature is still unknown. In the proposed models or the equation of state of dark energy \\cite{Copeland2006}, the cosmological constant cannot be excluded \\cite{Riess2007}.  If we assume that the dark energy is equivalent to a cosmological constant $\\Lambda$, there rises again so called a \\textit{cosmological constant problem}\\cite{Weinberg1989}: the present value of $\\Lambda$ is extraordinarily small compared with an inferred vacuum energy during the Planck time. To solve this problem, it is natural to consider that $\\Lambda$ decreases from a large value at the early epoch to the present value. Many functional forms of $\\Lambda$ have been suggested : for instance decaying-$\\Lambda$ has been introduced as a function of a scalar field in the Brans-Dicke gravitational theory \\cite{Nakamura2005}. The evolution of the universe under various models of $\\Lambda$ which are included in the energy-momentum tensor has been investigated analytically \\cite{Overduin}.  In addition, interacting $\\Lambda$ with other kinds of energy has been also discussed. The vacuum energy of $\\Lambda$ coupled with baryon could be ruled out, because baryon-antibaryon created through vacuum decay causes pair-annihilation. The produced high energy gamma ray flux is contradicted with the observations \\cite{Freese1986}. On the other hand, the vacuum energy decayed into the photon could affect the cosmological evolution significantly. Assuming that the ratio of the vacuum energy to the radiation is constant at the radiation dominated era $(z>10^5_{})$, Freese et al. \\cite{Freese1986} investigated the effects on primordial nucleosynthesis and obtained a limit of vacuum to photon energy ratio, which is less than 0.07. Furthermore, it is pointed out that observational constraints from the CMB intensity put the limit on the ratio of the vacuum to the radiation energy to be $\\sim 10^{-3}_{}$  \\cite{Overduin1993}.  From the thermodynamical point of view, the temperature-redshift relation is modified by including adiabatic photon creation due to vacuum decay \\cite{Lima1995, Lima1996}. A phenomenological decaying-$\\Lambda$ has been found to affect the cosmological evolution after the recombination \\cite{Kimura, Kamikawa, Puy}. In models having $\\Lambda$ terms as a function of the scale factor,  the radiation and matter temperatures would be significantly lower compared to the standard cold dark matter model with a constant $\\Lambda$ (S$\\Lambda$CDM) \\cite{Kimura}: the molecular formation is occurred at earlier epoch by $\\Delta z<10^3$ \\cite{Kamikawa}. Furthermore, it is shown that in some parameter regions, the radiation temperature could become higher compared with the S$\\Lambda$CDM model, which is found to be consistent with the observational result of $z<4$ \\cite{Puy}.  Related to  the recent observations, the first star formation that occurred at the end of the {\\it dark age} has been investigated progressively \\cite{Bromm2003}. In decaying $\\Lambda$ models, using a cooling diagram, the first star formation was estimated to occur at an earlier epoch by $\\Delta z\\sim 20$ with its mass $\\sim10^6_{}M^{}_{\\odot}$ \\cite{Kamikawa}. In the meantime, from an observational approach, the CMB polarization observed by the WMAP satellite predicts  via measured reionization redshift with use of the S$\\Lambda$CDM model that a first object was formed around $z =10$ \\cite{WMAP2006, Page06}. Since the CMB anisotropies give severe constraints on parameter regions concerning the cosmological evolution, we can also estimate the era of the first star formation in the decaying-$\\Lambda$ cosmology using the CMB anisotropy.  In the present paper, we constrain the parameter regions that determine the thermal history of the universe with a $\\Lambda$ decaying into the photon (hereafter we call it D$\\Lambda$CDM ). In Sec. \\ref{sec:model}, we describe briefly the thermal evolution and clarify the effects on photon decoupling in the D$\\Lambda$CDM model. In Sec. \\ref{sec:cmba} we examine the consistency between D$\\Lambda$CDM and the CMB anisotropy data of Winlinson Microwave Anisotropy Probe (WMAP) using the Markov-Chain Monte Carlo (MCMC) method. Summary and discussion are given in Sec. \\ref{sec:rad}. ",
        "conclusions": ""
    },
    "0801/0801.2281_arXiv.txt": {
        "abstract": "Physical parameters of galaxies (as luminosity, stellar mass, age) are often derived by means of the model templates which best fit their spectro-photometric data. We have performed a quantitative test aimed at exploring the ability of this procedure in recovering the physical parameters of early-type galaxies at $1<z<2$. A wide range of simulated SEDs, reproducing those of early-type galaxies at $1<z<2$ with assigned age and mass, are used to build mock photometric catalogs with wavelength coverage and photometric uncertainties similar to those of two topical surveys (i.e. VVDS and GOODS). The best fitting analysis of the simulated photometric data allows to study the differences among the recovered parameters and the input ones. Results indicate that the stellar masses measured by means of optical bands are affected by larger uncertainties with respect to those obtained from near-IR bands, and they frequently underestimate the real values. The M/L ratio in the V band results strongly underestimated, even when derived from the recently proposed recipe based on rest-frame optical colours (e.g. (B-V)). ",
        "introduction": "As a first step aimed at exploring the ability of different mass estimators in recovering the stellar mass content of early-type galaxies at $1<z<2$, we built a set of mock photometric catalogs.  We adopted the BC03 \\cite{ref:1} code, assuming Salpeter IMF and solar metallicity, to generate the SEDs of a wide family of models reproducing the observations of early-type galaxies. The selected SF histories are described by an exponentially declining SFR with time scale $\\tau=0.6$ Gyr, and by the superimposition of $\\tau=0.1$ Gyr models at different times simulating secondary bursts involving between 5\\% and 20\\% of the total final mass. The relevant range of redshift is described distributing the model templates at $z=1.0,1.5,2.0$. The chosen combinations among ages, redshift and times of the secondary burst constitute 57 template spectra, among which 11 with no secondary star forming episodes and 16 with recent (i.e.  $<1$ Gyr) starburst, half of which forms 5\\% of the total final mass while the contribution of the other half is 20\\%. Finally, each generated SED is proposed with three possible values of dust extinction A$_{V}=0.0,0.2,0.5$, assuming the reddening law of Calzetti et al. \\cite{ref:2}. ",
        "conclusions": ""
    },
    "0801/0801.2886_arXiv.txt": {
        "abstract": "{Accretion disks around young stars evolve in time with time scales of few million years.  We present here a study of the accretion properties of a sample of 35 stars in the $\\sim 3$ million year old star-forming region $\\sigma$~Ori. Of these, 31 are objects with evidence of disks, based on their IR excess emission. We use near-IR hydrogen recombination lines (Pa$\\gamma$) to measure their mass accretion rate. We find that the accretion rates are significantly lower in $\\sigma$~Ori than in younger regions, such as $\\rho$~Oph, consistently with viscous disk evolution. The He~I 1.083 $\\mu$m line is detected (either in absorption or in emission) in 72\\% of the stars with disks, providing evidence of accretion-powered activity also in very low accretors, where other accretion indicators dissapear.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.2909_arXiv.txt": {
        "abstract": "% The halo of NGC 891 has been the subject of studies for more than a decade. One of its most striking features is the large asymmetry in H$\\alpha$  emission. We have taken a quantitative look at this asymmetry at different wavelengths for the first time. We propose that NGC 891 is intrinsically almost symmetric, as seen in {\\em Spitzer} observations, and the large asymmetry in H$\\alpha$ emission is mostly due to dust attenuation. We quantify the additional optical depth needed to cause the observed H$\\alpha$ asymmetry. A comparison of large strips on the North East side of the galaxy with strips covering the same area in the South West we can quantify and analyze the asymmetry in the different wavelengths. From the 24 $\\mu$m emission we find that the intrinsic asymmetry in star-formation in NGC 891 is small i.e., approximately 30\\%. The additional asymmetry in H$\\alpha$ is modeled as additional symmetric dust attenuation which extends up to $\\sim$ 40'' (1.9 kpc) above the plane of the galaxy with a mid-plane value of  $\\tau$=0.8 and a scale height of 0.5 kpc. This observational technique offers the possibility  to quantify the effects of vertical ISM disk stability as an explanation for dust lanes in massive galaxies \\citep{Dalcanton04}. ",
        "introduction": "Extinction by dust manifests itself most conspicuously in the dark dust lanes seen in optical images of edge-on galaxies. Occasionally dark structures can be observed extending out of the plane \\citep{Howk97, Howk99a, Howk99b, Thompson04, Howk05}. The study of these dust-rich structures is hampered by the need to have enough stellar light backlighting the dark structures. Previous observations of the scale-height in nearby, mostly edge-on, galaxies appeared to find small scale-heights for dust \\citep[e.g,][]{Xilouris99,Holwerda05b,Bianchi07} but recent observations indicate a much higher scale-height for the dust \\citep[e.g.,][]{Seth05}, one that is similar to the stellar scale-height.  In many edge-on galaxies an H$\\alpha$ halo can be observed. \\cite{Heald06} and \\cite{Kamphuis07a} reported on the kinematics of this H$\\alpha$ halo for NGC 891. Its H$\\alpha$ distribution is extremely asymmetric: the NE side halo is much more extended and brighter (Figure \\ref{f:map}). Is this due to an intrinsic asymmetry in the H$\\alpha$ emission or an extinction effect?  \\begin{figure}[h] \\centering \\includegraphics[width=\\textwidth]{./Holwerda_p41_f1.eps}\\\\% picture filename \\caption{\\label{f:map} The H$\\alpha$ emission \\citep[false color from][]{Rand94} with 24 $\\mu$m contours overlaid, below are the integrated vertical profiles in H$\\alpha$(blue) and 24 $\\mu$m (red dashed line), normalized to the top in the NE side. Both types of emission are powered by ongoing star-formation. The 24 $\\mu$m is very symmetrical while the H$\\alpha$ is not.  Figure from \\cite{Kamphuis07}.} \\end{figure}  In \\cite{Kamphuis07}, we compared two star-formation driven emissions (H$\\alpha$ and {Spitzer} 24 $\\mu$m). The 24 $\\mu$m is much more symmetric (Figure \\ref{f:map}). Assuming both types of emission originate from star-formation regions in the spiral arms of NGC 891 and that the NE arm is closest to us (Figure \\ref{f:mod}), the difference in symmetry between the two types of emission can be attributed to extinction in the inner disk extending out of the plane of the disk ($\\tau_{add}$). The inferred scale-height of the out-of-the plane dust (Figure \\ref{f:prof}) is similar to the heights to which \\cite{Howk97} report dark structures in NGC 891. Still this height is much less than the height of dust emission in the halo as reported by M. Burgdorf \\& M. Ashby, ({\\em this volume}).  \\begin{figure}[h] \\centering \\includegraphics[width=0.9\\textwidth]{./Holwerda_p41_f2.eps}\\\\% picture filename \\caption{\\label{f:mod} A sketch of our model of NGC891. The emission originates predominantly from HII regions (blue) and much of the dust is in the spiral arms (thick black S-shape). The H$\\alpha$ emission is more absorbed in the SW case ($\\tau_{add}$).  Figure from \\cite{Kamphuis07}. } \\end{figure}  \\cite{Dalcanton04} link the appearance of dust lanes to the vertical stability of spiral disks. In more massive disks, gravitation wins over the star-formation driven turbulence and the dust-rich ISM collapses into the characteristic dust lane, notably in the inner scale-length of the disk. In contrast, low-mass galaxies display a more flocculant appearance of dust structures. The additional extinction out-of-the plane of NGC 891, a massive galaxy seems to be in contradiction to this. However, NGC 891 is undergoing significant star-formation ($log( ~ L(FIR)~) ~ = ~ 10.42 ~ L_\\odot $) and much of its ISM is in the dustlane.   \\begin{figure}[h] \\centering \\includegraphics[width=\\textwidth]{./Holwerda_p41_f3.eps}% \\caption{\\label{f:prof} The vertical extent in arseconds of the additional absorption ($\\tau_{add}$), assuming perfect NE/SW symmetry (left) and taking the 24 $\\mu$m symmetry as the intrinsic one (right). The blue solid line is from the \\cite{Rand94} H$\\alpha$ data and the dashed red line based on the \\cite{Kamphuis07a} H$\\alpha$ data. Figure from \\cite{Kamphuis07}.} \\end{figure} ",
        "conclusions": ""
    },
    "0801/0801.0417_arXiv.txt": {
        "abstract": "{}{This paper reports molecular observations towards the Bright Eastern Knot (BEK) in the SNR Puppis A, a feature where radio and X-ray studies suggest that the shock front is interacting with a dense molecular clump.}{We performed high-resolution millimetric observations towards the BEK of Puppis A using the SEST telescope in the \\12 J=1--0 and 2--1 lines (beams of 45\\s~and 23\\s~respectively). More extended, lower angular resolution \\12 J=1--0 observations taken from NANTEN archival data were also analyzed to obtain a complete picture.}{ In the velocity range near 16 \\k, the Puppis A systemic velocity, our study revealed two important properties: (i) no dense molecular gas is detected immediately adjacent to the eastern border of the BEK and (ii) the molecular clump detected very close to the radiocontinuum maximum is probably located in the foreground along the line of sight and has not yet been reached by the SNR shock front. We propose two possible scenarios to explain the absence of molecular emission eastwards of the BEK border of Puppis A. Either the shock front has completely engulfed and destroyed a molecular clump or the shock front is interacting with part of a larger cloud and we do not detect CO emission immediately beyond it because the molecules have been dissociated by photodissociation and by reactions with photoionized material due to the radiative precursor.}{}   \\titlerunning{High resolution CO observations towards the Puppis A BEK} \\authorrunning{S. Paron et al.}   \\keywords {ISM: molecules --- ISM: clouds --- ISM: supernova remnants --- ISM: individual objects: Puppis A} ",
        "introduction": "Shock waves generated by supernova remnants (SNRs) can accelerate, compress, heat, fragment or even destroy surrounding interstellar clouds. Strong shock-cloud interactions can enhance or reduce abundances of different molecular species with respect to quiescent cloud conditions. Observations of molecular gas associated with SNRs provide information essential to understand the physics and chemistry involved in these processes.  Puppis A is a Galactic SNR that has been extensively studied in the whole electromagnetic spectrum. In radio continuum, Puppis A appears as an asymmetric clumpy shell with the brightest section along the eastern border (Figure \\ref{observations}a) (\\citealt{gabi06} and references therein) presenting a good correlation with soft X-ray emission \\citep{petre82,aschen88}. The X-ray emission includes both extended features and compact knots, the most conspicuous of which is the bright eastern knot (BEK; \\citealt{petre82}) that coincides with an indentation in the shock front of the SNR as seen in radio. Such morphology suggests an interaction between the SNR and a dense interstellar clump. \\citet{hwang05} presented ACIS {\\it Chandra} X-ray images and spectral data of the region around the BEK. They conclude that a cloud-shock interaction in a relatively late phase of evolution is taking place near the BEK, while closer to the forward shock in the BEK region, the SNR has recently interacted with a more dense and extended obstacle.  CO studies of the interstellar medium surrounding Puppis A performed by \\citet{dubner88} with angular resolution of 8\\m.7, revealed the existence of a chain of molecular clouds concentrated along the E and NE periphery of the remnant. Studies in infrared wavelengths of Puppis A also suggest the existence of interstellar clouds along the eastern border of the SNR \\citep{arendt91}. From the atomic and molecular studies \\citep{dubner88,reyno95}, a systemic velocity of $v_{LSR} \\sim +16$ \\ks and a kinematical distance of $2.2 \\pm 0.3$ kpc were derived for Puppis A, a distance later confirmed by \\citet{reyno03} based on interferometric, high resolution HI data in the direction of the associated radio quiet neutron star RX J0822-4300 located within the remnant.  In this work, we present high resolution SEST observations performed in the \\12 J=1--0 and J=2--1 lines towards the BEK region, complemented with lower resolution molecular mapping of a larger region including the BEK based on archival data of the NANTEN telescope. The posibility of interaction between the SNR shock front and the surrounding gas is locally investigated in an attempt to understand the origin of the enhancement observed both in X-ray and radio continuum emissions. ",
        "conclusions": "The present observations enabled us to map in detail the molecular material towards the brightest X-ray emission region (BEK) in Puppis A. Our \\12 J=1--0 and J=2--1 high resolution observations show components at $v \\sim 3$ \\ks and $v \\sim 16$ \\k. The $v \\sim 3$ \\ks component is most probably related to the Gum Nebula.  In the velocity range near 16 \\k, our study revealed two important properties: (i) no dense molecular gas is detected immediately adjacent to the eastern border of the BEK of Puppis A (towards the SNR indentation) and (ii) the only detected cloud apparently associated with the BEK is probably located in the foreground along the line of sight and has not been reached by the SN shock yet.  To explain the absence of molecular emission eastwards of Puppis A in the BEK region, two scenarios were proposed. Either the shock front has completely engulfed and destroyed a molecular clump and the X-ray emission is due to the evaporation of the molecular gas, or the indentation is the consequence of an interaction with part of a larger cloud and we do not detect CO emission immediately beyond it because the molecules have been dissociated by photodissociation and by reactions with photoionized material due to the radiative precursor, as optical emission suggests. Far-infrared observations of the dust continuum and the fine structure of [CII] and [OI] may help to test this suggestion.  Concerning the molecular clump detected at $v \\sim 16$ \\k, we estimated for it a linear size of about 0.3 pc, a molecular mass of $\\sim 9.5$ M$_{\\odot}$, and a volume density of $\\sim 1.7 \\times 10^{4}$ cm$^{-3}$. The low value obtained in the ratio between the integrated lines towards this clump, $R_{2-1/1-0} \\sim 0.6$, suggests that we are observing quiescent gas located in front of Puppis A along the line of sight."
    },
    "0801/0801.3414_arXiv.txt": {
        "abstract": "We present an axisymmetric, equilibrium model for late-type galaxies which consists of an exponential disk, a Sersic bulge, and a cuspy dark halo.  The model is specified by a phase space distribution function which, in turn, depends on the integrals of motion. Bayesian statistics and the Markov Chain Monte Carlo method are used to tailor the model to satisfy observational data and theoretical constraints.  By way of example, we construct a chain of $10^5$ models for the Milky Way designed to fit a wide range of photometric and kinematic observations.  From this chain, we calculate the probability distribution function of important Galactic parameters such as the Sersic index of the bulge, the disk scale length, and the disk, bulge, and halo masses.  We also calculate the probability distribution function of the local dark matter velocity dispersion and density, two quantities of paramount significance for terrestrial dark matter detection experiments.  Though the Milky Way models in our chain all satisfy the prescribed observational constraints, they vary considerably in key structural parameters and therefore respond differently to non-axisymmetric perturbations.  We simulate the evolution of twenty-five models which have different Toomre $Q$ and Goldreich-Tremaine $X$ parameters.  Virtually all of these models form a bar, though some, more quickly than others.  The bar pattern speeds are $\\sim 40 - 50\\,{\\rm km\\,s^{-1}\\,kpc^{-1}}$ at the time when they form and then decrease, presumably due to coupling of the bar with the halo. Since the Galactic bar has a pattern speed $\\sim 50\\,{\\rm km\\,s^{-1} \\,kpc^{-1}}$ we conclude that it must have formed recently. ",
        "introduction": "Dynamical galactic models serve a variety of purposes.  They may be used to interpret the structural and kinematical observables of galaxies -- surface brightness profiles, rotation curves and velocity dispersion profiles -- in terms of intrinsic three-dimensional density and velocity distributions.  Dynamical models also provide a starting point for controlled simulations of complicated processes such as the formation of bars and spiral structure.  In short, galactic modeling provides the essential interface between observations and detailed theories of galaxy formation.   In this paper, we introduce a new dynamical model for late-type galaxies which comprises a disk, bulge, and dark halo.  The model is derived from equilibrium solutions to the collisionless Boltzmann and Poisson equations.  It extremely flexible and may be tailored to satisfy observational data and theoretical constraints.  We use Bayesian statistics and the Markov Chain Monte Carlo (MCMC) method to implement these constraints and to determine the probability distribution function (PDF) of the model in the full multi-dimensional parameter space.  Our model builds upon earlier work by \\citet{kui95} and \\citet{wid05}. The original Kuijken \\& Dubinski model consists of an exponential disk, a King-model bulge, and a lowered Evans-model halo and has the attractive feature that the phase space distribution function (DF) is built from analytic functions of the integrals of motion.  No additional assumptions about the velocity-space distribution are made. By contrast, the widely-used approach described in \\citet{her93} (see also \\citet{springel99}) assumes that the local velocity distributions of the halo and bulge particles are Gaussian with dispersions estimated from the Jeans equations.  This approach leads to models which are slighly out of equilibrium.  When used as initial conditions in N-body experiments, they readjust to a different state from the one proposed (see, for example, \\citet{kaz04}).  There are two main disadvantages of the Kuijken and Dubinski models. First, the bulge and halo have constant density (or weakly cuspy) centers whereas actual bulges and dark halos may have central density cusps.  Second, the structure of the bulge and halo are determined {\\it implicitly} by the model parameters.  (By construction, the disk's structural parameters, namely its radial and vertical scale lengths, its mass, and its truncation radius, are determined explicitly by the model parameters.)  \\citet{wid05} built a galactic model with $r^{-1}$-density cusps for both the bulge and halo but again, with a DF that determines the structure of the bulge and halo implicitly.  For our new model, the closed-form DF is abandoned in favor of a numerical DF which is designed to yield, to a very good approximation, user-specified density profiles for the bulge and halo.  That is, the density profiles of the bulge and halo are now {\\it explicit} functions of the model parameters.  The present version of the model allows for a Sersic bulge and a halo with $\\rho\\propto r^{-\\gamma}$ as $r\\to 0$ where $\\gamma$ is between $0$ and $2$.  Our model is specified in terms of fifteen or so parameters.  How are these parameters selected?  One approach is to choose models at random and identify the ones that satisfy certain general constraints (e.g., the Tully-Fisher and size-luminosity relations).  The result would be a catalog of disk-bulge-halo systems which could be used to study kinematical and dynamical trends such as the circular speed-central velocity dispersion ($V_c-\\sigma_0$) relation (see, for example, \\citet{cou07}).  A catalog of this type could also be used for N-body studies of mergers or interactions between galaxies in a cosmological environment.  A second approach, and the one pursued here, is to build models for specific galaxies.  Following \\citet{kui95} and \\citet{wid05}, we use the Milky Way as our illustrative example.  Modeling the Milky Way is a time-honoured endeavor; notable examples include \\citet{inn73}, \\citet{clu77}, \\citet{bs80}, \\citet{cal81}, \\citet{kui91}, \\citet{roh88}, \\citet{mal95}; \\citet{koc96}, \\citet{eva99}, and \\citet{kly02}.  Our construction of dynamical Milky Way models is in the spirit of the mass model survey by \\citet{deh98}.  Their models comprise a multi-component disk, a bulge, and a halo and are characterized by ten parameters.  Twenty-two examples are presented, each the result of a maximum likelihood analysis in which some parameters are held fixed while others are allowed to vary.  The Dehnen \\& Binney likelihood function is constructed by comparing model predictions with six sets of observational data: the inner and outer Galactic rotation curves, the Oort constants, the mass at large radii, the local vertical force, and the line-of-sight velocity dispersion in Baade's window. The advantage of our models is that they not only describe the potential-density pair for the Galaxy but also, the underlying DF.  We therefore have the ability to examine the stability of our Galactic models using N-body simulations, an issue that is often ignored (but see \\citet{sel85} and \\citet{fux97}).  For the most part, we adopt Dehnen \\& Binney's choice of observational data though we include more complete observations of the line-of-sight dispersion in the bulge region as well as photometric data from the COBE satellite.  We also present what we believe to be a more balanced treatment of the likelihood function.  Most significantly, we bring to the problem the powerful tools of Bayesian statistics and MCMC.  These tools allow us to map out PDFs of both input parameters and derived quantities.  Though our model represents an axisymmetric, equilibrium system, it is susceptible to non-axisymmetric instabilities and therefore provides a natural starting point for numerical studies of galactic dynamics.  An N-body realization of the model can be easily generated from the DF and then used as the initial conditions for a numerical simulation.  The Milky Way models in our MCMC series all satisfy the observational constraints but vary considerably in their structural properties.  We simulate a selection of twenty-five models which span a wide range in Toomre $Q$ \\citep{toomre64} and Goldreich-Tremaine $X$ \\citep{gt78, gt79} parameters and find that a bar develops in virtually all of the cases.  The onset of the bar instability can occur immediately or after several Gyr, depending on the model.  We present the model in Section 2, review the observational constraints in Section 3, and provide a summary of the essentials of Bayesian statistics and the MCMC method in Section 4.  We discuss some preliminaries including our choice of prior probabilities, in Section 5.  We present the results of our MCMC analysis in Section 6 and the results of our bar formation simulations in Section 7.  In Section 8 we summarize our main conclusions and speculate on how we might improve upon and extend the models and MCMC analysis. ",
        "conclusions": "We have introduced a dynamical model for late-type galaxies that incorporates our current understanding of disk-bulge-halo systems.  In particular, the bulge has a Sersic surface density profile and the halo has a central density cusp.  We have carried out an MCMC analysis of dynamical models for the Milky Way using a variety of kinematic and photometric constraints.  The results are presented in the form of PDFs for both input parameters and derived quantities.  The MCMC analysis provides a picture of the distribution of models in parameter space that is more complete than can be obtained by other approaches.  Avoided is the awkward procedure of fixing a subset of parameters while allowing the remaining parameters to vary in some minimization scheme.  Instead, a sequence of models is generated which contains all of the desired information. Marginalization over a subset of parameters is accomplished by simply projecting the model distribution onto the appropriate parameter subspace.  Our analysis suggests that the Milky Way has a pseudo-bulge with a Sersic index of $1.3\\pm 0.3$.  Our results for the masses of the disk, bulge, and halo are consistent with those of \\citet{deh98} but call into question choices for these quantities in some popular models from the literature.  For example, the disk and bulge masses in \\citet{johnston99} are entirely inconsistent with our results.  The inferred bulge mass for the standard model used by the MACHO collaboration \\citep{alcock00} is inconsistent with our findings by a factor of 2.5.  On the other hand, the standard values used by terrestrial dark matter detection experiments for the local dark matter density and velocity dispersion are consistent with our results.  A weak point of our analysis is the inability to tightly constrain the halo mass at large radii.  Planetary nebulae, globular clusters, and satellite galaxies may be used as tracers of the Galactic potential. Dynamical models for the tracer populations are required to properly model kinematic data.  In principle, it is straightforward to construct such models but there are subtle issues.  Previous studies showed that velocity anisotropy in a tracer population can affect interpretation of kinematic data.  Velocity anisotropy requires a DF that depends on at least one integral of motion in addition to the energy.  The standard practice is to use the total angular momentum, $L$, but since our models include a disk, $L$ is not conserved.  (Previous analyses side-stepped this issue by using a spherically symmetric Galactic potential.)  Our MCMC analysis provides an ideal starting point for studies of disk stability and bar formation in that we have some $10^5$ models, each of which can serve as initial conditions for a numerical experiment. We have performed a suite of simulations which focuses on the susceptibility of the disk to bar formation as a function of the stability parameters $Q$ and $X$.  Many of the models provide a good match to the inferred properties of the Galactic bar with the proviso that the bar has formed recently."
    },
    "0801/0801.3622_arXiv.txt": {
        "abstract": "Internal gravity waves (IGWs) are naturally produced by convection in stellar envelopes, and they could be an important mechanism for transporting angular momentum in the radiative interiors of stars. Prior work has established that they could operate over a short enough time scale to explain the internal solar rotation as a function of depth.  We demonstrate that the natural action of IGWs is to produce large scale oscillations in the solar rotation as a function of depth, which is in marked contrast to the nearly uniform rotation in the outer radiative envelope of the Sun.  An additional angular momentum transport mechanism is therefore required, and neither molecular nor shear-induced turbulent viscosity is sufficient to smooth out the profile. Magnetic processes, such as the Tayler-Spruit dynamo, could flatten the rotation profile.  We therefore conclude that IGWs must operate in conjunction with magnetic angular momentum transport processes if they operate at all. Furthermore, both classes of mechanisms must be inhibited to some degree by mean molecular weight gradients in order to explain the recent evidence for a rapidly rotating embedded core in the Sun. ",
        "introduction": "\\label{sec:intro}  During their pre-main sequence contraction, young solar-type stars are spun up to rotational velocities of the order of 100\\,\\kms. However, during their subsequent main sequence (MS) evolution the surface rotation slows down as a result of angular momentum loss through magnetized stellar winds (\\citealt{k88,mp08}). If a convective envelope and a radiative core of a solar-type MS star rotated independently of one another then the surface spindown would lead to a strong differential rotation beneath the envelope. In contradiction with this, helioseismic data (e.g., \\citealt{cea03}) reveal that the solar radiative core rotates as a solid body and almost synchronously with the convective envelope at least down to the radius $r\\approx 0.2\\,R_\\odot$. Besides, the spindown of young cluster stars (\\citealt{sh87,kea95,bea97,kea97,b03}) requires that the internal differential rotation only persists for timescales of the order of 20 Myr to 100 Myr. These data indicate that, in radiative interiors of solar-type MS stars, there is an efficient mechanism of angular momentum redistribution that couples their core and envelope rotation. Unfortunately, its physical nature remains elusive in spite of many attempts to understand it made in the last years. Matters are further complicated by the requirement that an appropriate model of angular momentum transport in the solar-type MS stars should also reproduce the intricate variations (depletion) of the surface Li abundance as functions of age and effective temperature observed in the same stars (\\citealt{sr05}).  A breakthrough in solving this complex problem has recently been announced by \\cite{tch05} and \\cite{cht05}. They have shown how internal gravity waves (hereafter, IGWs) generated by the envelope convection can extract angular momentum from rapidly rotating radiative cores of the solar-type MS stars on short enough timescales to explain both the spindown of young cluster stars and the quasi-solid-body rotation of the Sun. Furthermore, they have found that the quick flattening of the internal rotation profile by IGWs reduces element diffusion coefficients associated with hydrodynamic instabilities induced by differential rotation to values consistent with those constrained by the Li data.  Pursuing the goal of uncovering intrinsic causes of canonical extra mixing in low-mass red giant branch stars (\\citealt{dvb03}) and trying to understand the origin of fast rotation of red horizontal branch stars (\\citealt{sp00}), we have been undertaking a critical review of different transport mechanisms in stellar radiative zones. When applying our test model of angular momentum redistribution by IGWs to a model of the present-day Sun, we have found that IGWs strongly disturb the solar internal rotation making it disagree with the helioseismic data. The disturbance can be eliminated only if there is another transport mechanism competing with IGWs. Its efficiency should exceed that of the rotational shear mixing used by \\cite{tch05} by more than three orders of magnitude. Alternatively, a magnetic transport mechanism could be a competitor for IGWs. But in that case the mechanism preventing IGWs from disturbing the solid-body rotation of the Sun could itself be responsible for both shaping the internal rotation of the solar-type MS stars and assisting in the depletion of their surface Li abundance. This paper presents a discussion of computational results that support our conclusions. ",
        "conclusions": "Our numerical solutions of the angular momentum transport equation (\\ref{eq:ampdf}, or \\ref{eq:ampdf2}) have demonstrated that neither the molecular viscosity nor the shear-induced turbulent viscosity can reduce the large scale oscillations of angular velocity in the solar outer radiative core caused by selective damping of IGWs, provided that the net energy flux of these waves is strong enough to shape the Sun's solid body rotation. If waves were the sole mechanism, we would therefore expect to see large deviations from rigid rotation. The amplitudes of these oscillations are found to be too large to agree with the helioseismic data. Our result holds even when the molecular and shear-induced viscosity are multiplied by large factors. We have proved that only a viscosity exceeding the ohmic diffusivity by a factor of $\\ga$\\,10 can smooth out the IGW-induced oscillations of the rotation profile. This may be an indirect indication that some magnetic processes are at work here. To be more precise, our finding actually satisfies a necessary condition for such processes to work because the effective magnetic diffusivity $\\eta_{\\rm e}$ associated with them must exceed the ohmic diffusivity $\\eta_{\\rm mag}$. Otherwise magnetic fields generated by them will decay through the ohmic resistivity too quickly. For example, we have found that magnetic torques are strong enough to successfully compete with the action of IGWs only if the original prescription for the Taylor-Spruit dynamo, in which $\\eta_{\\rm e}\\gg\\eta_{\\rm mag}$, is used. However, this particular case turns out to be irrelevant to our problem because the Tayler-Spruit mechanism can shape the solar solid body rotation alone (\\citealt{eea05}), without being assisted by IGWs.  The helioseismic data suggest that either there is an efficient angular momentum transport mechanism in addition to IGWs that smooths out the SLOs produced by the waves or the spectral energy distribution of IGWs is different (lower) from those used by us. Although the latter assumption leads us outside the scope of our formulated problem we will comment on it. It is possible that strong toroidal magnetic fields in the solar tachocline filter out the IGWs with the doppler-shifted frequency $\\sigma$ above the Alfv\\'{e}n frequency (\\citealt{kea99,kmg03}). For the minimum frequency $0.6\\,\\mu$Hz used in our computations of the large scale SLOs in the solar outer radiative core the magnetic field strength required to prevent the inward wave propagation is about $(3\\times 10^5)/l$ Gauss (\\citealt{kea99}). For $l=1$, this corresponds to quite a strong field. If it is present in the tachocline then the waves with high spherical degrees will be trapped there. However, we have only considered IGWs with $\\omega \\geq 0.6\\,\\mu$Hz and low l values. Most of them are likely to propagate below the tachocline. In order to find out if an enhanced viscosity in the tachocline can hinder the propagation of low-degree high-frequency waves into the solar radiative core we have done test computations in which the shear-induced viscosity was increased by a factor of $10^3$. Their results plotted with black curves in Fig.~\\ref{fig:f7} show that this does not help to solve the problem.  Another, more radical possibility is that the form of the IGW spectra employed by us is completely wrong. For instance, the IGW spectrum estimated in the 2D hydrodynamic simulations by \\cite{rg05} has a flat energy distribution which goes three orders of magnitude below the peak luminosity in our Fig.~\\ref{fig:f1}. Apparently, if we applied that spectrum in our computations then even the molecular viscosity could easily smooth out the SLOs produced by such IGWs. However, it is evident as well that IGWs with this energy distribution could not produce the uniform rotation of the Sun by its present age (multiply the ages of the rotation profiles in Fig.~1 from \\citealt{cht05} by a thousand).  To summarize, we do not see what microscopic or pure hydrodynamic processes could smooth out the large scale SLOs induced by IGWs in the solar outer radiative core. Therefore, we agree with the conclusion made by \\cite{gmci98} about ``the inevitability of a magnetic field in the Sun's radiative interior''. Indeed, if IGWs are as strong as described by our employed spectra then this magnetic field is required to trigger magnetic processes that will counteract the disturbing action of IGWs on the solar rotation profile. On the contrary, if IGWs are weak then we are in need of such magnetic processes to extract an excess angular momentum from the solar interior on the early MS.  During the preparation of this work, results of a preliminary analysis of the GOLF data on solar g-mode oscillations have been published by \\cite{gea07} suggesting that the solar core at $r\\la 0.15\\,R_\\odot$ may rotate three to five times as fast as the rest of the radiative core. If these results prove to be correct, they will seem to rule out the angular momentum redistribution in the Sun by IGWs, as proposed by \\cite{cht05}, because in that case it would have been the Sun's inner core to be spun down first. If these results are confirmed, it will mean that the strong $\\mu$-gradient in the Sun's central region has prevented any angular momentum transport mechanism from operating there. There would also be strong implications for magnetic angular momentum transport, ruling out prescriptions (such as \\citealt{s02}) that predict a weak sensivity to $\\mu$ gradients."
    },
    "0801/0801.3078_arXiv.txt": {
        "abstract": "A study of cicumnuclear star-forming regions (CNSFRs) in several early type spirals has been made in order to investigate their main properties: stellar and gas kinematics, dynamical masses, ionising stellar masses, chemical abundances and other properties of the ionised gas. Both high resolution (R$ \\sim $20000) and moderate resolution (R $ \\sim $ 5000) have been used.  In some cases these regions, about 100 to 150\\,pc in size, are seen to be composed of several individual star clusters with sizes between 1.5 and 4.9\\,pc estimated from Hubble Space Telescope (HST) images.  Stellar and gas velocity dispersions are found to differ by about 20 to 30\\,km/s with the H$\\beta$ emission lines being narrower than both the stellar lines and the [OIII]\\,$\\lambda$\\,5007\\,\\AA lines. The twice ionized oxygen, on the other hand, shows velocity dispersions comparable to those shown by stars. We have applied the virial theorem to estimate dynamical masses of the clusters, assuming that systems are gravitationally bounded and spherically symmetric, and using previously measured sizes. The measured values of the stellar velocity dispersions yield dynamical masses of the order of 10$^7$ to 10$^8$ solar masses for the whole CNSFRs.    We obtain oxygen abundances which are comparable to those found in high metallicity disc HII regions from direct measurements of electron temperatures and consistent with solar values within the errors. The region with the highest oxygen abundance is R3+R4 in NGC~3504, 12+log(O/H) = 8.85, about 1.5 solar if the solar oxygen abundance is set at the value derived by \\citet{AGS05}, 12+log(O/H)$_{\\odot}$ = 8.66$\\pm$0.05. The derived N/O ratios are in average larger than those found in high metallicity disc HII regions and they do not seem to follow the trend of N/O vs O/H which marks the secondary behaviour of nitrogen. On the other hand, the S/O ratios span a very narrow range between 0.6 and 0.8 of the solar value. As compared to high metallicity disc HII regions, CNSFR show values of the O$_{23}$ and the N2 parameters whose distributions are shifted to lower and higher values respectively, hence, even though their derived oxygen and sulphur abundances are similar, higher values would in principle be obtained for the CNSFR  if pure empirical methods were used to estimate abundances. CNSFR also show lower ionisation parameters than their disc counterparts, as derived from the [SII]/[SIII]. Their ionisation structure also seems to be different with CNSFR showing radiation field properties more similar to HII galaxies than to disc high metallicity HII regions. ",
        "introduction": "\\label{introduction}  The inner ($\\sim$ 1Kpc) parts of some spiral galaxies show star formation complexes frequently arranged in an annular pattern around their nuclei. These complexes are sometimes called \"hot spots\" and we will refer to these as circumnuclear starforming regions (CNSFRs). Their sizes go from a few tens to a few hundreds of pc \\citep[e.g.][]{DATTSA00} and they seem to be made of an ensamble of HII regions ionised by luminous compact stellar clusters whose sizes, as measured from high spatial resolution HST images, are seen to be of only a few pc.  The luminosities of CNSFRs are rather large with absolute visual magnitudes (M$_v$) between -12 and -17 and H$\\alpha$ luminosities which are comparable to those shown by 30 Dor, the largest HII region in the LMC, and overlap with those shown by HII galaxies \\citep[][and references therein]{MTM88,DATTSA00,HD06}. In the ultraviolet (UV), massive stars dominate the observed circumnuclear emission even in the presence of an active nucleus \\cite{GD98,CGML02}.  In many cases, CNSFR show emission line spectra similar to those of disc HII regions. However, they show a higher continuum from background stellar populations as expected from their circumnuclear location.  The analysis of these spectra gives us the oportunity to measure the gas abundances close to galactic nuclei which,  in the case of early type spirals, are expected to be amongst the highest metallicity regions.  The importance of an accurate determination of the abundances of high metallicity HII regions cannot be overestimated since they constitute most of the HII regions in early spiral galaxies (Sa to Sbc) and the inner regions of most late  type ones (Sc to Sd) \\cite{D89} without  which our description of the metallicity distribution in galaxies cannot be  complete.  In particular, the effects of the choice of different calibrations  on the derivation of abundance gradients can be very important since any abundance  profile fit will be strongly biased towards data points at the ends of the  distribution. It should be kept in mind that abundance gradients are widely  used to constrain chemical evolution models, histories of star formation over galactic discs or galaxy formation scenarios. Also, the question of how high is the highest oxygen abundance in the gaseous phase of galaxies is still standing and extrapolation of known radial abundance gradients would point to CNSFR as the most probable sites for these high metallicities.  Accurate measures of elemental abundances of high metallicity regions are also crucial to obtain reliable calibrations of empirical abundance estimators, widely used but poorly constrained, whose choice can severely bias results obtained for quantities of the highest relevance for the study of galactic evolution like the luminosity-metallicity (L-Z) relation for galaxies. ",
        "conclusions": "\\label{Kinematics} Gas velocity dispersions were measured  by performing Gaussian fits to the H$\\beta$\\,$\\lambda$\\,4861\\,\\AA\\ and [OIII] $\\lambda$ 5007 \\AA\\ lines on the high dispersion spectra (Figure \\ref{spectra-kin}). Stellar velocity dispersions dispersions were measured using the  CaT lines at $\\lambda\\lambda$\\,8494, 8542, 8662\\,\\AA using the cross-correlation technique  described in detail by \\citet{TD79}. Late type giant and supergiant stars that have strong CaT absorption lines were used as stellar velocity templates.  \\begin{figure}[tb] \\hspace{0.2cm} \\includegraphics[width=.40\\textwidth,height=.31\\textwidth,angle=0]{1dn_R6-s2-b-rest.eps}\\\\ \\vspace{0.2cm} \\includegraphics[width=.40\\textwidth,height=.31\\textwidth,angle=0]{1dn_R6-s2-r-rest.eps} \\caption{Blue (upper panel) and red (lower panel) rest frame normalised spectra of R6. Notice the absence of conspicuous emission lines in the red spectral range for this region.} \\label{spectra-kin} \\end{figure}  For the 5 CNSFR observed in NGC~3351, stellar velocity dispersions are found to be between 39 and 67\\,km\\,s$^{-1}$, about 20\\,km\\,s$^{-1}$ larger than those measured for the gas, if a single Gaussian fit is used. However, the best fits obtained involved two different components for the gas: a ``broad component\" with a velocity dispersion similar to that measured for the stars, and a ``narrow component\" with a dispersion lower than the stellar one by about 30\\,km\\,s$^{-1}$. These two components were found both in the hydrogen recombination lines (Balmer and Paschen) and in the [OIII] $\\lambda$ 5007 \\AA\\ line. The narrow component is dominant in the H recombination lines, while the broad component dominates the [OIII] one.  Figure \\ref{velocities} shows this effect.   \\begin{figure}[tb] \\plotone{ngauss-Hb-R2.eps} \\vspace{0.2cm} \\plotone{ngauss-OIII-R2.eps} \\caption{Sections of the normalised spectrum of region R2 in NGC~3351. Left:  H$\\beta$ and right  [O{\\sc iii}]\\,$\\lambda$\\,5007\\,\\AA\\ . The dashed-dotted line is the broad component, the dotted line is the narrow component and the dashed line is the sum of both} % \\label{velocities} \\end{figure}   CNSFR are seen to consist of several individual star clusters, although some of them seem to have an only knot, at the HST resolution. The derived masses for the individual clusters as derived using the sizes measured on the HST images are between 1.8 and 8.7\\,$\\times$\\,10$^6$\\,M$_\\odot$. These values are between 5.5 and 26 times the mass derived for the SSC A in NGC\\,1569 by \\citet{HF96} and larger than other kinematically derived SSC masses. Values for the dynamical masses of the CNSFRs are in the range between 4.9\\,$\\times$\\,10$^6$ and 4.3\\,$\\times$\\,10$^7$\\,M$_\\odot$. Masses derived from the H$\\beta$ velocity dispersion under the assumption of a single component for the gas would have been underestimated by factors between approximately 2 to 4.  The masses of the ionising stellar clusters of the CNSFRs have been derived from their H$\\alpha$ luminosities under the assumption that the regions are ionisation bound and without taking into account any photon absorption by dust. Their values are between 8-0\\,$\\times$\\,10$^5$ and 2.5\\,$\\times$\\,10$^6$\\,M$_\\odot$.  Therefore, the ratio of the ionising stellar population to the total dynamical mass is between 0.02 and 0.16.  The SSC in the observed CNSFRs seem to contain composite stellar populations. Although the youngest one dominates the UV light and is responsible for the gas ionisation, it constitutes only about 10\\,\\% of the total. This can explain the low equivalent widths of emission lines measured in these regions.  This may well apply to the case of other SSC and therefore conclusions drawn from fits of SSP (single stellar population) models should be taken with caution \\citep[e.g.][]{MGG03}. Also the composite nature of the CNSFRs  means that star formation in the rings is a process that has taken place over time periods much longer than those implied by the properties of the ionised gas."
    },
    "0801/0801.3849_arXiv.txt": {
        "abstract": "Recent work has suggested that the fraction of obscured AGN declines with increasing luminosity, but it has been difficult to quantify this trend. Here, we attempt to measure this fraction as a function of luminosity by studying the ratio of mid-infrared to intrinsic nuclear bolometric luminosity in unobscured AGN.  Because the mid-infrared is created by dust reprocessing of shorter wavelength nuclear light, this ratio is a diagnostic of $f_{\\rm obsc}$, the fraction of solid angle around the nucleus covered by obscuring matter. In order to eliminate possible redshift-dependences while also achieving a large dynamic range in luminosity, we have collected archival 24 micron MIPS photometry from objects with $z$$\\sim$1 in the Sloan Digital Sky Survey (SDSS), the Great Observatories Origins Deep Survey (GOODS) and the Cosmic Evolution Survey (COSMOS). To measure the bolometric luminosity for each object, we used archival optical data supplemented by GALEX data. We find that the mean ratio of 24~$\\mu$m to bolometric luminosity decreases by a factor of $\\sim$3 in the $L_{\\rm bol}$=10$^{44}$-3$\\times$10$^{47}$~ergs~s$^{-1}$ range, but there is also a large scatter at constant $L_{\\rm bol}$. Using radiation transfer solutions for model geometries, we show how the IR/bolometric ratio relates to $f_{\\rm obsc}$ and compare these values with those obtained obtained from samples of X-ray selected AGN.  Although we find approximate agreement, our method indicates somewhat higher values of $f_{\\rm obsc}$, particularly in the middle range of luminosities, suggesting that there may be a significant number of heavily obscured AGN missed by X-ray surveys. ",
        "introduction": "Following the initial recognition by \\citet{antonucci85} that the nucleus of the prototypical type~2 Seyfert galaxy NGC~1068 must be surrounded by a toroidal belt of gas and dust, a tremendous amount of evidence has accumulated in support of the idea that similar obscuring tori surround the nuclei of other type~2 Seyfert galaxies (as reviewed, for example, in \\citealp{antonucci93,krolik99}).  Spectropolarimetry (the method originally employed by Antonucci and Miller) applied to both type~2 Seyfert galaxies and radio galaxies \\citep{diserego94} often reveals in polarized light broad optical/UV emission lines and a strong non-stellar continuum that are essentially invisible in the total flux spectrum.  In many examples of both these types of AGN, conical regions of bright line emission from highly-ionized elements point at the otherwise obscured nucleus \\citep{ferruit00,schmitt03}.  Similarly, both type~2 Seyfert galaxies and radio galaxies often exhibit X-ray continua that either show evidence for very large column densities of material along the line of sight or are extremely weak, likely indicating that the obscuration is Compton thick \\citep{risaliti99,treister04}.  In two cases, NGC~1068 and the Circinus galaxy, the obscuring torus can be seen directly in interferometric infrared imaging \\citep{jaffe04,tristram07}.  The principal effect of this toroidal obscuration is that observers trying to see the nucleus along lines of sight that pass through it are prevented from seeing anything but dust-reprocessed infrared continuum and (perhaps) hard X-rays.  Consequently, most of the classic signatures of AGN---broad optical/UV emission lines, strong optical/UV non-stellar continuum, strong X-ray continuum---are obliterated when the object is seen from such a direction.  Only when the line of sight passes through the central hole of the torus can all these features be seen, and the object is perceived as a type~1 AGN.  In the nearby Universe, where we can see large numbers of low luminosity AGN, large statistical samples have been assembled in order to measure the luminosity functions of both type~1 and type~2 AGN. Ideally compared at matched bolometric luminosity, but more often in terms of their luminosity in a single band or feature, the ratio of their numbers can be immediately interpreted as the ratio of unobscured to obscured solid angle.  For example, \\citet{hao05} and \\citet{simpson05}, using samples of optically-selected AGN from the Sloan Digital Sky Survey, found that the ratio of type~1 objects to type~2 objects increases with increasing luminosity, from $\\simeq 2/3$ at an [OIII]~5007 luminosity of $10^6L_{\\odot}$ to $\\simeq 2$ for line luminosities $\\sim 3 \\times 10^7L_{\\odot}$. Using samples of X-ray selected AGN, \\citet{steffen03} and \\citet{barger05} argued for a similar shift in the ratio of unobscured to obscured, suggesting that the unobscured variety come to dominate the total population for 2--8~keV luminosity $\\gtrsim 10^{44}$~erg~s$^{-1}$. Similar results were obtained by \\citet{ueda03}, \\citet{lafranca05} and others, also using X-ray selected AGN samples. In fact, hints of such a trend were already seen in the much smaller 3CR radio sample \\cite{lawrence91}.  However, there remain significant uncertainties in the measured population ratio at high luminosity.  For example, Compton-thick AGN are completely missing from X-ray-selected samples. Other sample construction methods (e.g., the very large sample of type~2 quasars selected from the SDSS based primarily on [OIII] emission by \\citealp{zakamska03}) can potentially yield quantitative estimates of comparative luminosity functions, but require substantial work in order to turn catalogs into space densities.  Although much effort has been made in the last few years to construct AGN samples from mid-infrared observations with the {\\it Spitzer Space Telescope}, the selection effects remain too poorly understood and the sample size is sometimes too small to permit extraction of a trend in the type ratio as a function of luminosity \\citep{stern05,alonso-herrero06,martinez06,lacy07}.  Here we attempt a different approach to the problem of measuring the unobscured/obscured ratio: using the ratio of reradiated mid-infrared continuum to bolometric luminosity in type~1 AGN as a surrogate (cf. the suggestions in \\citealp{lawrence91,barger05}).  Because the obscuration transforms essentially all the nuclear radiation incident upon it into mid-infrared continuum via dust reradiation, and detailed radiation transfer calculations suggest that most of that continuum is radiated toward observers in the type~1 direction no matter what the detailed arrangement of dust may be \\citep{pier92,granato94,efstathiou95,nenkova02,vanbemmel03,dullemond05}, this ratio should give a good indicator of the ratio of unobscured/obscured solid angle.  The relative ease of identifying type~1 AGN across a wide dynamic range in luminosity also gives this method a number of advantages relative to those dependent on actually counting type~2 objects. \\citet{maiolino07} have recently made an effort to implement this program, but the details of their method differ in significant ways from ours; we will comment on specific contrasts as they arise.  This paper is structured as follows: In section \\S2 we present the selection criteria and basic observational properties of the sources used in this work. In section \\S3 we study the correlations found in this sample, while the significance and interpretation of these trends are discussed in \\S 4. Our conclusions are presented in \\S 5. When required, we assume a $\\Lambda$CDM cosmology with $h_0$=0.71, $\\Omega_m$=0.3 and $\\Omega_\\Lambda$=0.7, in agreement with the most recent cosmological observations \\citep{spergel07}. ",
        "conclusions": "We presented in this paper an alternative way to estimate the fraction of obscured AGN and its possible dependence on luminosity, by computing the relative fraction of IR emission in unobscured AGN spanning a large of luminosities. Using a sample of 206 optically-selected high luminosity AGN from the SDSS and 24 X-ray selected lower luminosity AGN from the GOODS and COSMOS surveys with 0.8$\\leq$$z$$\\leq$1.2, nearly all having Spitzer detections at 24~$\\mu$m, we found a decrease of a factor of $\\sim$3 in the relative IR emission at rest-frame 12~$\\mu$m from $L_{\\rm bol}$=10$^{44}$~ergs~s$^{-1}$ to 10$^{47.5}$~ergs~s$^{-1}$. In addition, we found a significant scatter of a factor of 2--3 at a given bolometric luminosity. Some of this scatter can likely be explained by a dependence of the 11--13~$\\mu$m fraction $f_{12}$ on viewing angle; some may also be due to an intrinsic scatter in $f_{\\rm obsc}$ at fixed bolometric luminosity due to a dependence on other parameters such as black hole mass, magnetic forces, etc.  Using IR re-emission models in order to convert $L_{\\rm MIPS}/L_{\\rm bol}$ into a fraction of obscured AGN, we found that $f_{\\rm obsc}$ changes from $\\sim 90\\%$ at $L_{\\rm bol} = 10^{44}$~ergs~s$^{-1}$ to $\\sim 30$--$40\\%$ at $L_{\\rm bol} \\simeq 10^{47}$~ergs~s$^{-1}$. The derived dependence of the obscured fraction on luminosity is in good agreement with the direct, but possibly biased (by the omission of Compton-thick objects) observation of this fraction in X-ray surveys.  We caution, however, that there may be systematic error in the precise values of $f_{\\rm obsc}$ inferred from $L_{\\rm MIPS}/L_{\\rm bol}$ because our estimate of $\\langle f_{12}\\rangle$ is likely to depend somewhat on the details of the density distribution within the torus.  Nonetheless, it is interesting that in general the values of $f_{\\rm obsc}$ we infer are larger than what is seen in X-ray surveys.  This contrast hints that there may be a significant population of heavily obscured, even Compton thick, AGN that are missed in X-ray observations but included in our samples."
    },
    "0801/0801.4372_arXiv.txt": {
        "abstract": "This paper is the first in a series that aims to understand the long-term evolution of neutron star magnetic fields. We model the stellar matter as an electrically neutral and lightly-ionized plasma composed of three moving particle species: neutrons, protons, and electrons; these species can be converted into each other by weak interactions (beta decays), suffer binary collisions, and be affected by each other's macroscopic electromagnetic fields. Since the evolution of the magnetic field occurs over thousands of years or more, compared to dynamical timescales (sound and Alfv\\'en) of milliseconds to seconds, we use a slow-motion approximation in which we neglect the inertial terms in the equations of motion for the particles. This approximation leads to three nonlinear partial-differential equations describing the evolution of the magnetic field, as well as the movement of two fluids: the charged particles (protons and electrons) and the neutrons. These equations are first rather than second order in time (involving the velocities of the three species but not their accelerations). In this paper, we restrict ourselves to a one-dimensional geometry in which the magnetic field points in one Cartesian direction, but varies only along an orthogonal direction. We study the evolution of the system in three different ways: (i) estimating timescales directly from the equations, guided by physical intuition; (ii) a normal-mode analysis in the limit of a nearly uniform system; and (iii) a finite-difference numerical integration of the full set of nonlinear partial-differential equations. We find good agreement between our analytical normal-mode solutions and the numerical simulations. We show that the magnetic field and the particles evolve through successive quasi-equilibrium states, on timescales that can be understood by physical arguments. Depending on parameter values, the magnetic field can evolve by ohmic diffusion or by ambipolar diffusion, the latter being limited either by interparticle collisions or by relaxation to chemical quasi-equilibrium through beta decays. The numerical simulations are further validated by verifying that they satisfy the known conservation laws in highly nonlinear situations. ",
        "introduction": "\\label{intro}  Observations of the surface magnetic fields of neutron stars have shown a correlation between the age of the star and the magnetic-field strength. For instance, young radio pulsars and high-mass X-ray binaries have surface magnetic fields of the order of $10^{12}$ G, while millisecond pulsars and low-mass X-ray binaries, which are older objects, have magnetic fields $\\sim 10^{8}$ G. These observations suggest that the magnetic field is decaying, although this might be a by-product of accretion (\\cite{B-95}; \\cite{CZB-01}; \\cite{C-02}; \\cite{PM-04}). On the other hand, it is thought that the spontaneous decay of the ultrastrong $10^{14-15}$ G magnetic field in \\emph{magnetars} is the main source of their X-ray luminosity since these objects appear to radiate substantially more power than that available from their rotational energy loss (\\cite{DT-92}; \\cite{TD-96}; \\cite{ACT-04}). Since these objects appear to be isolated, the field decay must be attributed to processes intrinsic to the neutron stars.  Ref.\\cite{GR-92}; (hereafter GR-92) identified three long-term mechanisms that can promote the spontaneous decay of the magnetic field in the interior of neutron stars: \\begin{enumerate} \\item  \\emph{Ambipolar diffusion}, i.e. the drift of the magnetic field and the charged particles relative to the neutrons. Ref. \\cite{TD-96}; (hereafter TD-96), using magnetar parameters, estimated the timescale of the magnetic field decay by means of ambipolar diffusion, finding that was consistent with the timescale of $10^{3-5}$ years observed for these objects. \\item \\emph{Hall drift}, a non-dissipative advection of the magnetic field by the associated electrical current (\\cite{J-88}; \\cite{RG-02}; \\cite{GR-02}; \\cite{HR-02}, \\cite{HR-04}; \\cite{RKG-04}; \\cite{CAZ-04}; \\cite{R-05}, \\cite{R-07}; \\cite{Rei-07}). \\item \\emph{Ohmic diffusion}, a dissipative process caused by electrical resistivity (see also Baym et al. \\cite{Bay-69}). \\end{enumerate} All of these processes occur over timescales of thousands of years or more, compared to typical dynamical timescales (sound and Alfv\\'en) of miliseconds to seconds. The work of GR-92 was analytical, and therefore useful to identify general processes and relevant timescales, but it did not address the action of the identified processes in their full nonlinear development and their interactions with each other. The full evolution of the magnetic field can only be addressed by numerical simulations.  Regarding the geometrical configuration of the magnetic field, three-dimensional magnetohydrodynamics (MHD) simulations showed that in a stable stratified, non-rotating star, the initial field evolves on a short, Alfv\\'en-like timescale to a large-scale equilibrium configuration that can be axisymmetric (\\cite{BS-04},~\\cite{BS-06}; \\cite{BN-06}) or non-axisymmetric (\\cite{B-08}), depending on the radial dependence of the initial magnetic-field strength. These simulations, being focused on the evolution of the magnetic field over short dynamical timescales, were carried out within the framework of the MHD theory, which considers the stellar matter to be a single fluid. The evolution of these configurations has not yet been studied over longer timescales, on which relative motions of different particle species are present and the processes studied by GR-92 become important. The description of these long-term processes requires a multi-fluid theory.  In this paper we extend the model of GR-92 by allowing both neutron and charged-particle movements, rather than considering the neutrons as a fixed background (see Sect.~\\ref{phys}). For simplicity, we restrict ourselves to a plane-parallel system where the magnetic field points in one Cartesian direction but varies only along an orthogonal direction. Although this one-dimensional model is unrealistic, it allows us to study the basic evolutionary processes in an analytical and numerical fashion with no serious complications before considering a realistic stellar geometry in more dimensions. GR-92 considered all the species in the star as non-interacting fermions; in contrast, we use equations of state including interactions among protons and neutrons. Since in neutron stars the magnetic pressure is much less than the degeneracy pressure of the particles, we treat the particle densities in terms of small perturbations around a non-magnetized background in full equilibrium. Thus, we obtain a system of equations that is nonlinear with respect to the magnetic field but linear with respect to the particle densities (see Sect.~\\ref{1dmodel}).  We will show that the magnetic field and particle species evolve through successive quasi-equilibrium states. We describe the basic physical processes of this evolution and estimate analytically the characteristic timescales required to reach these quasi-equilibrium states. We find that these timescales depend on the values of physical parameters in the system, such as the magnetic-field intensity, the collision rate between the charged particles and the neutrons, the weak interaction rate, and the ohmic resistivity (Sect.~\\ref{carsca}). To obtain an analytic solution of the system of partial differential equations, we linearize the equations with respect to the magnetic field and find normal-mode solutions (Sect.~\\ref{normmodes}) and obtain analytical approximations for the decay times in Appendix \\ref{ApA}. We develop numerical simulations (whose basic algorithm is described in Appendix \\ref{ApB}) in the linear regime and compare with the normal-mode results (Sect.~\\ref{finline}). In Sect.~\\ref{finnoli}, we carry out numerical simulations in the nonlinear regime and verify the conservation laws of our model. Finally in Sect.~\\ref{conc}, we summarize the main results of this paper and provide our conclusions. ",
        "conclusions": "\\label{conc}  We have studied the long-term evolution of the magnetic field and the densities of neutrons and charged particles (protons and electrons in the interior of a neutron star, using a multi-fluid model with a simplified geometry in which the magnetic field points in one Cartesian direction and varies along an orthogonal direction. We found a set of three non-linear partial differential equations of first order in time describing this evolution, and we analyzed them in three different ways: (i) estimating evolutionary timescales directly from the equations, guided by physical intuition; (ii) a normal-mode analysis of the equations in the limit of a nearly uniform system; and (iii) a finite-difference numerical integration of the equations. In this section, we summarize the main results of each one of these approaches and present the main conclusions of this work.  \\subsection{Evolutionary timescales} \\label{conc1}  The three partial differential equations of our model constitute a dynamical system of three independent variables (magnetic-field, charged-particle, and neutron density perturbations). We identified three characteristic evolutionary timescales on which the system approaches successive quasi-equilibrium states. In the following, we summarize the basic physical processes governing the evolution and the estimates of each timescale.  \\subsubsection{Timescale to achieve magnetohydrostatic quasi-equilibrium} \\label{conc11} During the early stages of a neutron star's life, the Lorentz force moves the bulk stellar fluid, inducing perturbations in its pressure. In this short time, all particles move together as a single fluid with the same bulk velocity. There is also not enough time for weak interactions between particles (beta decays) to operate, so the composition is frozen. The system evolves until it reaches a magnetohydrodystatic quasi-equilibrium state, in which the Lorentz and the fluid forces are close to balancing. Neutron stars with no rotation or convection reach this quasi-equilibrium state in a short time, not much longer than the Alfv\\'en time, which for typical magnetar core parameters scales as  \\begin{equation} \\label{tafr} t_{Alfven}=5.7~\\times 10^{-2}~ R_{6}~B_{15}^{-1}~s, \\end{equation} where $R_{6}\\equiv R/(10^6~cm)$ denotes the radius of the star in units of $10^6~cm$ and $B_{15}\\equiv B_z/(10^{15}~G)$ the magnetic field in units of $10^{15}~G.$ On timescales far longer than the Alfv\\'en time, relative movements between the different species of particles in the star as well as weak interactions between them become relevant (GR-92). Therefore, a description of the system during these stages requires a multi-fluid theory. In this paper, we develop a model to study the decay of the magnetic field induced by these long-term mechanisms. The Alfv\\'en time is far shorter than the timescales of these long-term processes (see Sects.~\\ref{conc12} and \\ref{conc13}), which are of the order of years or much more. A numerical code simulating the evolution during the Alfv\\'en timescale would require a time step many orders of magnitude shorter than that required to simulate the long-term evolution in a computational time not prohibitively long. In our model, we overcome this difficulty by replacing the short-term dynamics by a ficticious friction force term acting on the neutrons (the most abundant species). Its strength is controlled by an artificial parameter $\\alpha$, chosen so that the timescale to reach  magnetohydrostatic quasi-equilibrium is long enough for the numerical code to be able to deal with it (and therefore much longer than the Alfv\\'en time), but shorter than the timescales of the long-term processes that we are interested in modeling.  \\subsubsection{Timescale for the evolution of the particle densities} \\label{conc12}  On timescales far longer than required to reach magnetohydrostatic quasi-equilibrium, the neutrons and charged particles move relative to each other with different velocities and are affected by collisional drag. Also, weak interactions (beta decays) convert particles from one species into another, erasing departures from chemical quasi-equilibrium among neutrons, protons, and electrons caused by the induced particle density perturbations. Both processes contribute to the evolution of the particle densities. In estimating the timescale for this evolution, it is useful to distinguish the regimes in which each one of these processes is dominant. If the collisions between the charged particles and the neutrons are rare and the weak interactions are slow, the charged particles move easily relative to the neutrons (ambipolar diffusion) allowing the particles to reach a diffusive quasi-equilibrium state (in which gravitational, electromagnetic, and pressure forces are closely balanced) but staying avoid chemical quasi-equilibrium. The diffusive quasi-equilibrium is reached in a timescale controlled by the collision rate between neutrons and charged particles given by  \\begin{equation} t_{drag} \\sim~4.5~\\times 10^{-1}~ L_5^2~ T_8^2 ~ yr, \\end{equation} where $L_5\\equiv L/(10^5~cm)$ and $T_8\\equiv T/(10^8 K).$  In the opposite regime, when the collisions between charged particles and neutrons are frequent, and the weak interaction rate is high, the relative diffusion of particles will be impeded by the collisional drag, but the system reaches the chemical quasi-equilibrium in a timescale controlled by weak interactions of magnitude (assuming modified Urca reactions)  \\begin{equation} t_{weak} \\sim 4.3 ~ \\times 10^5~ T_8^{-6} ~yr. \\end{equation}  \\subsubsection{Timescale for magnetic field evolution} \\label{conc13} The Ohmic dissipation promotes the decay of the magnetic field in a timescale  \\begin{equation} t_{ohmic}~\\sim 1.4~\\times~ 10^{11}~ L_5^{2}~ T_8^{-2}~yr. \\end{equation}  This extremely long time implies that Ohmic dissipation is not very effective in producing magnetic-field decay in neutron stars. A more rapid decay of the magnetic field can occur through ambipolar diffusion, i.e., a joint drift of the magnetic field and the charged particles relative to the neutrons, in which magnetic energy becomes dissipated by collisions.  This diffusion process induces local density perturbations that deviate from chemical quasi-equilibrium. Weak interactions tend to erase the chemical imbalance by converting particles of one species into the other. During these conversions, neutrinos, and antineutrinos remove part of the magnetic energy. In estimating the timescale of magnetic field decay induced by these processes, it is useful to distinguish between the two regimes we discussed above.  If ambipolar diffusion occurs far more rapidly than the weak interaction process, i.e.,  $t_{drag}\\ll t_{weak},$ diffusive quasi-equilibrium is quickly reached, but chemical quasi-equilibrium is not. Thus, the rate at which chemical quasi-equilibrium is restored by weak interactions determines the characteristic timescale over which the magnetic field evolves,  \\begin{equation} \\label{tbb1} t_B~\\sim~ 1.7\\times~ 10^{9}~ B_{15}^{-2}~T_8^{-6}~yr. \\end{equation} This timescale is the corresponding one to reach chemical quasi-equilibrium $t_{weak}$ but amplified by a factor of the order of the ratio of the charged fluid pressure to the magnetic pressure. This is expected since the Lorentz force drives the ambipolar diffusion that prevents weak interactions from restoring chemical quasi-equilibrium.  In the opposite regime, collisions are frequent and weak interactions are fast, i.e. $t_{weak}\\ll t_{drag},$ so the system reaches chemical quasi-equilibrium long before diffusive quasi-equilibrium. Therefore, the evolution of the magnetic field is limited by the collisions that control the ambipolar diffusion, on a timescale of  \\begin{equation} \\label{tbb2} t_B~\\sim~ 1.8\\times~ 10^{3}~ B_{15}^{-2}~ L_{5}^2~ T_8^{2}~yr. \\end{equation} Since the magnetic field sustains the diffusive imbalance and drives the diffusion, the timescale given in Eq.~(\\ref{tbb2}) is similar to that taken by particles to achieve diffusive quasi-equilibrium but, as above, amplified by a factor that depends on the magnetic field strength.  In the limit of non-interacting protons and neutrons studied by GR-92, Eqs.~(\\ref{tbb1}) and (\\ref{tbb2}) correspond to the estimate of the timescale for magnetic field evolution due to the ambipolar diffusion process of Eq.~(35) in GR-92. We note that GR-92 considered the neutrons to be a static species in diffusive quasi-equilibrium. From this correspondence, we confirm that the motion of the neutrons, which was included in our model, plays no important role in the one-dimensional evolution.  \\subsection{Normal-mode solutions} \\label{conc2}  In the spirit of finding  an analytical solution to the set of non-linear partial differential equations that describes the evolution of the system, we completed a linear perturbation analysis of the equations in the limit of a nearly uniform system and tried a normal mode solution (Sect.~\\ref{normmodes}). We found three exponentially-decaying modes that could be interpreted as the episodes of successive relaxation to the quasi-equilibrium states identified in the previous physical analysis (see  Sects.~\\ref{carsca} and ~\\ref{conc1}). We found analytical and numerical solutions for the decay times of these modes (see Appendix \\ref{ApA}) that agree with the estimated formulae for the evolutionary timescales of the general non-linear system (see  Sects.~\\ref{carsca}, \\ref{conc1} and Figs.~\\ref{t1},\\ref{t2}, and \\ref{t3}).  \\subsection{Numerical integration of the partial differential equations} \\label{conc3}  We constructed a finite-difference numerical code to solve the full system of non-linear partial differential equations. This was tested against the normal-mode solutions (Fig.~ \\ref{linen}) and by verifying the known conservation laws of magnetic flux and baryon number in highly non-linear situations (Fig.~ \\ref{gauss}). Applying this numerical code, we propose to consider, in the future, conditions in which the background is non-homogeneous, as well as non-linear situations.  \\subsection{Conclusions} \\label{summa}  We have established a general multifluid formalism in which it is possible to study the magnetic field decay processes in neutron stars identified by GR-92, which are likely to represent the basis of the magnetar phenomenon. As a first step, we have focused on a simplified geometry in which the magnetic field points in one Cartesian direction and varies in an orthogonal direction to it. We estimated the timescales of the relevant processes, and using numerical simulations we followed the temporal evolution of some simple magnetic-field configurations. The present work is far from exhausting the possibilities of this formalism, which should be applied to more realistic situations, including neutron stars with spherical symmetry, non-uniform background density and three-dimensional magnetic fields with components that depend on two or three spatial coordinates. Further steps might consider convective motions and additional species of particles as well as the effects of superfluidity and superconductivity."
    },
    "0801/0801.1761_arXiv.txt": {
        "abstract": "We report here a study of the long term properties of Quasi Periodic Oscillations (QPO) in an unusual accreting X-ray pulsar, 4U 1626--67. This is a unique accretion powered X-ray pulsar in which we have found the QPOs to be present during all sufficiently long X-ray observations with a wide range of X-ray observatories. In the present spin-down era of this source, the QPO central frequency is found to be decreasing. In the earlier spin-up era of this source, there are only two reports of QPO detections, in 1983 with EXOSAT and 1988 with GINGA with an increasing trend. The QPO frequency evolution in 4U 1626--67 during the last 22 years changed from a positive to a negative trend, somewhat coincident with the torque reversal in this source. In the accretion powered X-ray pulsars, the QPO frequency is directly related to the inner radius of the accretion disk, as per Keplerian Frequency Model (KFM) and Beat Frequency Model (BFM). A gradual depletion of accretion disk is reported earlier from the X-ray spectral, flux and pulse profile measurements. The present QPO frequency evolution study shows that X-ray flux and mass accretion rate may not change by the same factor, hence the simple KFM and BFM are not able to explain the QPO evolution in this source. This is the only X-ray pulsar to show persistent QPOs and is also the first accreting X-ray pulsar in which the QPO history is reported for a long time scale relating it with the long term evolution of the accretion disk. ",
        "introduction": "The X-ray source 4U 1626--67 was discovered with the Uhuru satellite (Giacconi et al. 1972) in 2-6 keV band. Pulsations, with a period of 7.68 s were first discovered by Rappaport et al. (1977) with SAS-3 observations and has been extensively monitored since then, especially with the BATSE detectors onboard CGRO (Chakrabarty et al 1997; Bildsten et al 1997). Optical counterpart of the pulsar was identified as KZ TrA, a faint blue star (V $\\approx$ 18.5) with little or no reddening (McClintock et al. 1977; Bradt and McClintock 1983). Optical pulsations with 2\\% amplitude were detected at the same frequency as the X-ray pulsations (Ilovaisky, Motch, \\& Chevalier 1978) and are understood to be due to reprocessing of the pulsed X-ray flux by the accretion disk (Chester 1979). A faint optical counterpart and the observed optical pulsed fraction requires the companion star to be of very small mass (McClintock et al. 1977, 1980). The X-ray light curve does not show any orbital modulation or eclipse. However, from the reprocessed pulsed optical emission and a close sideband in the power-spectrum of optical light curve, an orbital period of 42 minutes was inferred (Middleditch et al. 1981). Therefore, it falls under the category of ultra compact binaries (P$_{orb}$ $<$ 80 minutes), which have hydrogen-depleted secondaries to reach such short periods (Paczynski \\& Sienkiewicz 1981; Nelson et al. 1986).  Despite extensive searches, the orbital motion of this binary has never been detected in the X-ray pulse timing studies (Rappaport et al, 1977; Levine et al. 1988; Jain et al. 2007). A very low mass secondary, in a nearly face on orbit can possibly account for the lack of pulse arrival time delay. Recently Jain et al (2007) have also proposed this source to be a candidate for a neutron star with a supernova fall back accretion disk. From the extensive timing and spectral observations both in optical and X-ray bands, it has not yet been possible to establish the presence of a binary companion, and the upper limit of the companion mass has been determined to be very low. However, the presence of an accretion disk in 4U 1626--67 is beyond any doubt. Optical spectral and timing studies confirm that most of the optical emission is strongly dominated by the accretion disk (Grindlay 1978; McClintock et al. 1980). The X-ray spectrum also shows bright hydrogen-like and helium-like oxygen and neon emission lines with red and blue shifted components, a certain sign for accretion disk origin (Schulz et al. 2001, Krauss et al. 2007). Another direct evidence of an accretion disk in 4U 1626--67 is found from the detection of quasi-periodic oscillations, at a frequency of 40 mHz, from Ginga observations (Shinoda et al. 1990) and subsequently at a higher frequency of about 48 mHz from Beppo-SAX, ASCA, RXTE and XMM-Newton (Owens et al. 1997; Angelini et al. 1995; Kommers et al. 1998, Krauss et al. 2007). The QPOs have also been detected in reprocessed optical emission from both ground based and HST observations (Chakrabarty et al. 1998, 2001).  For more than a decade since its discovery, 4U 1626--67 was found to be spinning up with a characteristic timescale P/\\.{P}  $\\approx$ 5000 yr. It was found to be spinning down at about the same rate by BATSE onboard CGRO in the beginning of 1991 (Chakrabarty et al. 1997). Even though the torque reversal was abrupt, the decrease in bolometric X-ray flux has been gradual and continuous over the past $\\approx$ 30 yr (Chakrabarty et al. 1997, Krauss et al 2007). Recently, from a set of Chandra monitoring observations Krauss et al (2007) have established that the bolometric X-ray flux and various emission line fluxes have decreased continuously over the last few years, indicating a gradual depletion of the accretion disk. The X-ray flux and mass accretion rate are directly related and these are likely to be related to the mass and extent of the material in the accretion disk. Therefore, the observed gradual decrease in X-ray flux indicates a depletion of material in the accretion disk of the pulsar. Another signature of this is seen by Krauss et al. (2007) as a change in the pulse profile of the pulsar as compared to the earlier observations.  In the present work, we have investigated the QPO frequency evolution of 4U 1626--67 over a long period and discuss the relation of the change in QPO frequency with the a possible recession of the inner accretion disk. ",
        "conclusions": "\\begin{itemize} \\item We have detected very persistent quasi-periodic oscillations in the unique accretion powered X-ray pulsar 4U 1626--67. \\item Using data from several observatories, we have detected a gradual evolution of the oscillation frequency over a period of 22 years. \\item The frequency evolution indicates a possible recession of the accretion disk of the pulsar during the present spin-down era. \\end{itemize}"
    },
    "0801/0801.0371_arXiv.txt": {
        "abstract": "If ultra-high-energy cosmic rays (UHECRs) originate from extragalactic sources, understanding the propagation of charged particles through the magnetized large scale structure (LSS) of the universe is crucial in the search for the astrophysical accelerators. Based on a novel model of the turbulence dynamo, we estimate the intergalactic magnetic fields (IGMFs) in cosmological simulations of the formation of the LSS. Under the premise that the sources of UHECRs are strongly associated with the LSS, we consider a model in which protons with $E \\geq 10^{19}$ eV are injected by sources that represent active galactic nuclei located inside clusters of galaxies. With the model IGMFs, we then follow the trajectories of the protons, while taking into account the energy losses due to interactions with the cosmic background radiation. For observers located inside groups of galaxies like ours, about 70\\% and 35\\% of UHECR events above 60 EeV arrive within $\\sim 15^{\\circ}$ and $\\sim 5^{\\circ}$, respectively, of the source position with time delays of less than $\\sim 10^7$ yr. This implies that the arrival direction of super-GZK protons might exhibit a correlation with the distribution of cosmological sources on the sky. In this model, nearby sources (within $10 - 20$ Mpc) should contribute significantly to the particle flux above $\\sim10^{20}$ eV. ",
        "introduction": "Over the past several decades, significant progress has been made on both theoretical and observational fronts in understanding the nature and origin of ultrahigh energy cosmic rays (UHECRs), those with $E \\ga 1$ EeV ($ = 10^{18}$ eV) \\citep[for recent reviews, see][]{nagano00,bgg06}. Yet, the acceleration mechanism and the corresponding astrophysical ``accelerators\" of these energetic particles are still largely unknown. Observational data from several experiments such as the High Resolution Fly's Eye (HiRes) indicate that the mass composition of UHECRs becomes lighter at higher energies \\citep{abbasi05}. However, composition analyses that include high energy interactions are often model-dependent and inconclusive \\citep{watson06}. According to a recent report from the Pierre Auger Observatory, the mass composition is likely mixed, possibly becoming heavier above 30 EeV \\citep{unger07}. The overall distribution of UHECR arrival directions is considered to be consistent with isotropy \\citep{bo03}. Exceptions to this general isotropy include the small-scale clusterings of doublets and triplets found in data collected by the Akeno Giant Air Shower Array (AGASA) experiment \\citep{takeda99,uchihori00}. In addition, a possible correlation of AGASA events with BL Lacertae objects has been suggested \\citep{tinyakov01}. But these claims have been confirmed neither by HiRes \\citep{abbasi06} nor by Auger \\citep{harari07,armen07}. However, the Auger Collaboration recently reported that the arrival directions of UHECRs above $60$EeV in their data show a correlation with the position of active galactic nuclei (AGNs) lying within 75 Mpc \\citep{augerScience}.  Since protons with $E \\ga 1$ EeV cannot be confined within the Galactic plane,  UHECRs likely originate from extragalactic sources. In particular, the overall isotropy of arrival directions suggests that there may be a large number of sources distributed over cosmological distances \\citep{nagano00,bo03}. During their propagation through intergalactic space, such protons will lose energy by means of pion and pair production processes while interacting with the cosmic background radiation \\citep{g66,zk66}. The flux of ultra-high-energy (UHE) protons from cosmological sources is thus expected to be strongly attenuated, resulting in a significant suppression in the observed spectrum above the GZK threshold energy, $E_{\\rm GZK}\\approx 40$ EeV \\citep{bgg06}. Although the AGASA data show no indication of the GZK suppression \\citep{nagano00}, both the Yakutsk Extensive Air Shower array and HiRes have reported a suppression of flux above $E_{\\rm GZK}$, contradicting the AGASA finding \\citep{pravdin99,zech04}. Indeed the same suppression was seen in a recent Auger measurement \\citep{facal07}, which seems to have ended the controversy over the presence of the GZK cutoff. The so-called dip-calibrated UHECR spectra from different experiments compiled by \\citet{bgg06}, appear to be in good agreement with each other and to be consistent with GZK suppression. However, it has yet to be understood whether this suppression is actually due to the GZK cutoff or due to the maximum acceleration energy, $E_{\\rm max}$, of astrophysical accelerators. If UHECRs are protons, the GZK energy loss should operate at acceleration sites as well, leading to an $E_{\\rm max}$ close to $E_{\\rm GZK}$ \\citep[see, \\eg][]{krb97}.  Most UHE protons observed above the GZK energy must come from within the so-called GZK sphere of radius $R_{\\rm GZK} \\sim 100 $ Mpc, although the proton interaction length at $E_{\\rm GZK}$ is $l_{\\rm 40EeV} \\approx 1$ Gpc, corresponding to $z\\sim 0.2$  \\citep{berezinsky88}. However, finding cosmological sources inside the GZK sphere from the arrival directions of UHECRs is not straightforward, since their paths are deflected by intergalactic magnetic fields (IGMFs) \\citep[e.g.][]{sigletal03,dolagetal04}. Sigl and collaborators have extensively studied the propagation of UHECRs in a structured and magnetized universe, adopting a numerical model for the IGMFs \\citep{sigletal03,sigletal04,armen05}. In this model, the IGMFs are generated by means of the Biermann ``battery'' mechanism at shocks and then evolved passively in a cosmological hydrodynamic simulation \\citep{kulsrude97,rkb98}. The strength of the resulting fields is rescaled to match the simulated field strength in Coma-like clusters to the observed strength, which is on the order of microgauss. UHECRs with $E\\ge10$EeV are then injected at cosmological sources. These particles propagate through the magnetized large scale structure (LSS) of the universe, and arrive at a mock observer with deflection angle $\\theta$, the angle between the arrival direction and the source's location on the sky. Sigl \\etal\\ found that the deflection due to IGMFs is significant, with $\\theta \\ga 20^{\\circ}$ above 100 EeV. On the other hand, Dolag \\etal\\ adopted an IGMF model from a ``constrained'' cosmological simulation employing a magnetohydrodynamic (MHD) version of a smoothed particle hydrodynamics (SPH) code. They found the deflection angle of protons with 100 EeV to be less than $1^{\\circ}$, contradicting the estimate by Sigl \\etal\\ \\citep{dolagetal04,dolagetal05}.  This controversy over the predicted deflection angles demonstrates the importance of modeling the IGMFs in identifying the astrophysical sources and studying the origin of UHECRs. In order to reexamine this issue, here we adopt IGMFs, based on a novel models of the turbulence dynamo \\citet{ryuetal08} In this model, the strength of the magnetic fields is estimated from the local vorticity and turbulent kinetic energy in cosmological structure formation simulations. For the field direction, the passive fields from these simulations are used.  The maximum energy of nuclei of charge $Z$ that can be confined and accelerated by astrophysical sources is given by \\begin{equation} E_{\\rm max} \\approx Z \\left({V \\over c}\\right) \\left({B \\over \\mu{\\rm G}}\\right) \\left({L \\over {\\rm kpc}}\\right) 10^3 {\\rm EeV}, \\end{equation} where $V$, $B$, and $L$ are the characteristic flow speed, magnetic field strength, and linear size of the accelerator, respectively \\citep{hillas84}. There are a few viable candidates that can produce the required $E_{\\rm max} \\sim 100$ EeV: jets from AGNs \\citep[\\eg][]{biermann87}, gamma-ray bursts \\citep[\\eg][]{waxman95}, and cosmological shocks \\citep{krj96,krb97}. In this study, we consider AGNs inside galaxy clusters as the sources of UHECRs. Thus, the source position is in effect correlated with the LSS of the universe. Protons with $E\\ge 10$EeV are injected at the sources and travel through the simulated magnetized space until they lose energy down to 10 EeV, meanwhile visiting mock observers placed inside groups of galaxies. In this study we focus mainly on the deflection angle, the time delay relative to rectilinear flight, and the energy spectrum of the UHE protons.  In the next section, we describe our model IGMFs, cosmic-ray sources, and observers, and the simulations of the propagation of UHE protons in intergalactic space. The results are presented in \\S 3. Finally, we conclude in \\S 4. ",
        "conclusions": "In the search for the astrophysical sources of UHECRs, it is important to understand how the propagation of these charged particles is affected by intergalactic magnetic fields in the large-scale structure of the universe. On the other hand, the information imprinted on the distribution of the UHECR arrival directions may help us to understand the nature of the IGMFs and their roles in the formation and evolution of the LSS and constituent galaxies. Considering the limitations of current observational techniques in measuring the IGMFs in very low density regions such as filaments and voids, it is crucial to construct a physically motivated model to estimate the IGMFs in the LSS.  In this study, we adopted a new model based on the turbulence dynamo \\citep{ryuetal08} to predict the strength of the IGMFs. The magnetic field energy is estimated from the local vorticity and turbulent kinetic energy of flow motions in cosmological simulations of the LSS formation in a concordance $\\Lambda$CDM universe. For the direction of the IGMFs, the topology of the passive magnetic fields followed in the cosmological simulations is used. This approach provides an IGMF model that is independent of the initial seed fields and does not require any renormalization to yield the observed field strength in the intracluster medium. We predict highly structured IGMFs with characteristic field strengths on the order of $10^{-6}$ G in clusters of galaxies and $10^{-8}$ G in filaments. The fields should be much weaker in sheets and voids.  Protons with 10 EeV $\\le E_{\\rm inj}\\le 10^3$ EeV are injected at the locations of luminous X-ray clusters with $kT>1$ keV. These sources may represent a population of AGNs residing inside host clusters. This X-ray temperature criterion naturally places the sources at strongly magnetized regions with $ B_s \\sim 0.1 \\mu$G with a comoving density of $2-3 \\times 10^{-5}(\\Mpc)^{-3}$. Then the propagation of the UHE protons is followed through the structured IGMFs, including the energy losses due to interactions with the cosmic background radiation The UHE protons are recorded at the positions of mock observers located in groups of galaxies, with halo temperatures in the range $0.05 {\\rm keV} < kT < 0.5 {\\rm keV}$.  Below the GZK energy, the UHE protons come from sources as distant as $\\sim 1$ Gpc. They are significantly scattered by the IGMFs, resulting in a wide range of deflection angles, up to 180$^{\\circ}$, and have time delays ranging from $10^7$ to $10^{9.5}$ yr. On the other hand, the protons above 60 EeV come mostly from sources within $\\sim 100$ Mpc. About 70\\% of them avoid strong deflection and arrive at the observers within $\\sim 15^{\\circ}$ of their source position on the sky with a time delay of less than $\\sim 10^7$ yr. About 35\\% arrive within $\\sim 5^{\\circ}$. This implies that there may exist a correlation between the arrival direction of super-GZK cosmic rays and the sky position of the corresponding AGNs. We thus conclude that in the present scenario, UHECR astronomy may be possible at $E>$ 60 EeV. Our prediction seems to be consistent with a recent report by the Auger Collaboration \\citep{augerScience} in which the arrival directions of cosmic rays above 60 EeV in their data were found to be correlated with the sky position of AGNs within 75 Mpc.  For any cosmological sources, we expect to see GZK suppression in the energy spectrum of UHECRs if the injection spectrum has a power-law distribution and extends well beyond the GZK energy. In this case, nearby sources, within $10 - 20$ Mpc, are expected to make a significant contribution to the flux above $\\sim 100$ EeV.  Finally, as recently reported by Auger \\citep{unger07}, some UHECRs might be heavy nuclei. In the cosmological-shock model, for example, protons can be accelerated up to a few times 10 EeV, so heavy nuclei should dominate the particle flux above that energy \\citep{krb97,inoue07}. In a future study, we will consider the propagation of heavy nuclei from cosmological sources in our model IGMFs, taking into account photo disintegration, photo pair production, and photo pion production processes. We expect that the propagation of UHE iron nuclei ($Z=26$), at least, will be in the diffusive transport regime as a results of their much smaller gyroradius ($r_g\\propto E/Z$) \\citep{armen05}. For intermediate-mass nuclei such as He, C, N, and O, detailed propagation simulations including secondary particles produced by photo disintegrations are necessary in order to determine whether astronomy with UHE nuclei is possible or not."
    },
    "0801/0801.0147_arXiv.txt": {
        "abstract": "% There is now unequivocal evidence that the jets in FR I radio galaxies are initially relativistic, decelerating flows. On the assumption that they are axisymmetric and intrinsically symmetrical (a good approximation close to the nucleus), we can make models of their geometry, velocity, emissivity and field structure whose parameters can be determined by fitting to deep VLA observations. Mass entrainment -- either from stellar mass loss within the jet volume or via a boundary layer at the jet surface -- is the most likely cause for deceleration. This idea is quantitatively consistent with the velocity field and geometry inferred from kinematic modelling and the external gas density and pressure profiles derived from X-ray observations.  The jets must initially be very light, perhaps with an electron-positron composition. ",
        "introduction": "The morphological division of extragalactic radio sources into two classes introduced by \\cite{FR74} has proved to be remarkably robust. Fanaroff \\& Riley had already noted that the edge-brightened (FR\\,II) sources tend to have higher radio luminosities than the edge-darkened (FR\\,I) sources, but there is also a dependence on stellar luminosity of the host galaxy which makes the division even cleaner \\citep{Ledlow}.  That said, the FR\\,I class is not homogeneous: examples of the range of structures are shown in Fig.~\\ref{fig:fris}. Although almost all FR\\,I sources show jets on small scales, some appear to be confined to the nuclear regions (e.g.\\ 3C\\,84; Fig.~\\ref{fig:fris}a) while others persist almost to the outer edges of the source structure (e.g.\\ 3C\\,296; Fig.~\\ref{fig:fris}f).  It has long been surmised (e.g.\\ \\citealt{Simon78,Fanti82}) that the characteristic structures of FR\\,I sources result from deceleration and disruption of their jets, perhaps triggered by interaction with the surrounding intergalactic medium. Turning this perception from a vague idea into a quantitative description of source dynamics proved to be a slow process, but significant progress has recently been made using a combination of deep radio imaging, sophisticated jet modelling and an understanding of the external environments of the sources from X-ray observations.  These results bear on five of \\citet{Blandford}'s tasks: map jet velocity fields; understand the changing composition; measure jet pressures; deduce jet confinement mechanisms and infer jet powers and thrusts.   \\begin{figure}[!ht] \\plotone{fris.eps} \\caption{Examples of the radio structures of FR\\,I sources from the 3CRR catalogue \\citep{3CRRAtlas}.\\label{fig:fris}} \\end{figure} ",
        "conclusions": ""
    },
    "0801/0801.2004_arXiv.txt": {
        "abstract": "% The basic idea of the kd-tree algorithm is to recursively partition a point set P by hyperplanes, and to store the obtained partitioning in a binary tree. Due to its immense popularity, many applications in astronomy have been implemented. The algorithm can been used to solve a near neighbor problem for cross-identification of huge catalogs and realize the classification of astronomical objects. Since kd-tree can speed up query and partition spaces, some approaches based on it have been applied for photometric redshift measurement. We give the case studies of kd-tree in astronomy to show its importance and performance. ",
        "introduction": "K-dimensional tree (kd-tree), as a computer science term, is a space-partitioning data structure for organizing points in a $k$-dimensional space (Bentley, 1975). Technically, the letter $k$ refers to the numbers of dimensions. A 3-dimensional kd-tree can be called as $3d$-tree. A kd-tree organizes datapoints in such a way that once built, whenever a query arrives requesting a list all points in a neighborhood, the query can be answered quickly without needing to scan every single point. Each tree node represents a subvolume of the parameter space, with the root node containing the entire $k$ dimensional volume spanned by the data. Non-leaf nodes have two children, obtained by splitting the widest dimension of the parent's bounding box, the left child owning those data points that are strictly less than the splitting value in the splitting dimension, and the right child owning the remainder of the parent's data points. A kd-tree is usually constructed top-down, beginning with the full set of points and then splitting in the center of the widest dimension. This produces two child nodes, each with a distinct set of points. Repeat this procedure recursively; a kd-tree can be constructed. kd-trees have been successfully applied in astronomy for some problems. Hsieh et~al. (2005) used the kd-tree algorithm to divide up their sample to improve the redshift accuracy of galaxies. The kd-tree method is used on the 5 flux-space indexing in the SDSS Science Archive to partition the bulk data (Kunszt et~al. 2000). In the following, we give three kinds of case studies with kd-trees. ",
        "conclusions": "From the case studies, it is concluded that kd-trees have a wide application in astronomy due to its own characteristics. Kd-trees are useful data structure for many applications, such as searches involving a multi-dimensional search key. Therefore many researches related to optional and fast search may be searched with the kd-tree approach. With the increase of astronomical data in volume, kd-tree may play a more important role."
    },
    "0801/0801.1082.txt": {
        "abstract": "We present high- and medium-resolution spectroscopic observations of the cataclysmic variable BF~Eri during its low and bright states. The orbital period of this system was found to be 0.270881(3) days. The secondary star is clearly visible in the spectra through absorption lines of the neutral metals \\MgI, \\FeI\\ and \\CaI. Its spectral type was found to be K3$\\pm$0.5. A radial velocity study of the secondary yielded a semi-amplitude of $K_2=182.5\\pm0.9$ \\kms. The radial velocity semiamplitude of the white dwarf was found to be $K_1=74\\pm3$ \\kms\\ from the motion of the wings of the \\Halpha\\ and \\Hbeta\\ emission lines. From these parameters we have obtained that the secondary in BF~Eri is an evolved star with a mass of 0.50--0.59 $M_\\odot$, whose size is about 30 per cent larger than a zero-age main-sequence single-star of the same mass. We also show that BF~Eri contains a massive white dwarf ($M_1\\geq1.23 M_\\odot$), allowing us to consider the system as a SN Ia progenitor. BF~Eri also shows a high $\\gamma$-velocity ($\\gamma\\ =-94$ \\kms) and substantial proper motion. With our estimation of the distance to the system ($d\\approx700\\pm200$ pc), this corresponds to a space velocity of $\\sim$350 \\kms\\ with respect to the dynamical local standard of rest. The cumulative effect of repeated nova eruptions with asymmetric envelope ejection might explain the high space velocity of the system. We analyze the outburst behaviour of BF~Eri and question the current classification of the system as a dwarf nova. We propose that BF~Eri might be an old nova exhibiting ``stunted'' outbursts. ",
        "introduction": "Cataclysmic Variables (CVs) are close interacting binaries that contain a white dwarf (WD) accreting material from a companion, usually a late main-sequence star. CVs are very active photometrically, exhibiting variability on time scales from seconds to centuries (see review by \\citealt{Warner}). Dwarf novae (DNs) are an important subset of CVs. Observable features that distinguish DN from other CVs such as nova-like variables (NLs) are the recurrent outbursts of 2--6 mag that occur on timescales of days to years. In this paper we concentrate on BF Eridani, a little-studied dwarf nova candidate.  BF Eri was discovered and first classified as a slowly varying variable star by \\citet{HanleyShapley}. During the following half a century no observations of the star were reported. As a consequence, the specific nature of the star remained unknown. The identification of the Einstein survey source 1ES 0437-046 with BF~Eri \\citep{Elvis} and optical spectroscopy by \\citet{Schachter} eventually led to its correct identification as a CV. Later, the ROSAT identification of the Einstein source confirmed the classification of BF~Eri as a CV of an unknown type \\citep{Chisholm}.  \\begin{table*} \\label{ObsTab} \\begin{center} %\\begin{flushleft} \\caption{Log of observations of BF~Eri} \\begin{tabular}{ccclcccll} \\hline\\hline Set     &  Date        & HJD Start  & Instrument &$\\Delta\\lambda$$^a$& $\\lambda$~range & Exp.Time & Number   & Duration\\\\ &              &  2450000+  &            &  (\\AA)            & (\\AA)           &  (sec)   & of exps. &  (hours)\\\\ \\hline 2005-N1 & 2005-Oct-30  &  3673.982  & B\\&Ch$^b$& 6.0          & 3680--6770      &  300     & 16       & 1.44   \\\\ 2005-N2 & 2005-Oct-31  &  3674.821  & B\\&Ch   &  6.0          & 6025--8000      &  300     & 40       & 3.50   \\\\ 2005-N3 & 2005-Nov-01  &  3675.821  & B\\&Ch   &  2.1          & 6145--7225      &  300     & 58       & 5.25   \\\\ \\hline 2006-Nov \\\\ 2006-N1 & 2006-Nov-22  &  4061.884  & Echelle & 0.234         & 3915--7105      & 1200     & 9        & 3.13   \\\\ 2006-N2 & 2006-Nov-23  &  4062.728  & Echelle & 0.234         & 3915--7105      & 1200     & 17       & 6.81   \\\\ 2006-N3 & 2006-Nov-24  &  4063.710  & Echelle & 0.234         & 3915--7105      & 1200     & 19       & 7.08   \\\\ 2006-N4 & 2006-Nov-25  &  4064.684  & Echelle & 0.234         & 3915--7105      & 1200     & 3        & 0.70   \\\\ &              &  4064.879  & Echelle & 0.234         & 3915--7105      & 1200     & 3        & 0.70   \\\\ 2006-N5 & 2006-Nov-26  &  4065.698  & Echelle & 0.234         & 3915--7105      & 1200     & 20       & 7.26   \\\\ 2006-N6 & 2006-Nov-27  &  4066.689  & Echelle & 0.234         & 3915--7105      & 1200     & 14       & 4.78   \\\\ \\hline 2006-Dec \\\\ & 2006-Dec-13  &  4082.840  & Echelle & 0.234         & 3915--7105      & 1200     & 2        & 0.35   \\\\ & 2006-Dec-14  &  4083.850  & Echelle & 0.234         & 3915--7105      & 1200     & 2        & 0.35   \\\\ & 2006-Dec-15  &  4084.808  & Echelle & 0.234         & 3915--7105      & 1200     & 2        & 0.36   \\\\ \\hline \\end{tabular} %\\end{flushleft} %\\end{center} \\begin{tabular}{l} %\\begin{flushleft} $^a$ -- $\\Delta\\lambda$ is the FWHM spectral resolution \\\\ $^b$ -- B\\&Ch - Boller \\& Chivens spectrograph  \\\\ \\end{tabular} \\end{center} %\\end{flushleft} \\end{table*}   The long term light curve was first analysed by \\citet{Watanabe} who detected two outbursts and concluded that BF~Eri is a DN with a recurrence period of 40--50 days and the observed magnitude range between 13.2 and fainter than 14.7 mag. \\citet{Kato99} and \\citet{KatoUemura} confirmed this classification but mentioned the low amplitude of its outbursts and the difference in the outburst activity between seasons as the amplitudes of outbursts were remarkably different -- 1.0 mag and 1.6 mag. \\citet{Kato99} was unable to detect any orbital variability.  Prior to the beginning of this work there was no other study of the system in the literature\\footnote{ %  When this work was almost ready, a new paper on BF~Eri has appeared. \\citet{BF-Eri} have recently presented optical medium-resolution spectroscopy of BF~Eri, obtained primarily to determine the orbital period. Below we discuss their results.}. This motivated us to perform time-resolved spectroscopy of BF~Eri in order to study its properties in more detail. In this paper we present and discuss the first high- and medium-resolution spectroscopic observations of BF~Eri during its low and bright states in 2005 and 2006. In Section~\\ref{ObsSec} we  describe our  observations and  data reduction. The data  analysis and the results are presented  in Section~\\ref{DatAnSec}. Here we find the orbital period, measure the radial velocities of the WD and the secondary star, and determine the spectral type of the secondary and its rotational velocity. In Section~\\ref{DopMapSec} we discuss the results of Doppler tomography of the emission lines from BF~Eri. The system parameters are obtained in Section~\\ref{SysParSec}, while a discussion and a summary are given in Sections~\\ref{DiscSec} and \\ref{SumSec}, respectively.  %************************  Average Spectra  ************************************ \\begin{figure*} \\centering \\includegraphics[width=17cm]{spectrum2.eps} \\caption{The average spectra of BF~Eri from the \\textit{2005} set, uncorrected for orbital motion. The lower spectrum is combined from the \\textit{2005-N1} and \\textit{2005-N2} sets while the inset spectrum is an average of all spectra from the \\textit{2005-N3} set.} \\label{aver_spec} \\end{figure*} %************************  Average Spectra  ************************************ ",
        "conclusions": "\\label{DiscSec}  \\subsection{System parameters} \\label{SysParDiscSect}  We would like to make some final remarks on the system parameters of BF~Eri, derived in Section~\\ref{SysParSec}. These parameters were determined using a traditional spectroscopic solution based on the derived radial velocity semi-amplitudes. In order to calculate the masses of the binary we have assumed that the semi-amplitudes of the measured radial velocities represent the true orbital motion of the stars. In CVs, however, this may not be an accurate assumption, as the emission lines arising from the disc may suffer several asymmetric distortions. Likewise, absorption line profiles may be subjected to irradiation or hot-spot contaminations in the surface of the secondary, distorting therefore the true value of $K_2$.  Taking these potential sources of errors into consideration we, however, believe we could avoid them. Indeed, though those problems were severe in the early determinations of the radial velocities, nevertheless modern high resolution spectroscopy has greatly improved the methods in detecting asymmetric distortions to provide more reliable radial velocity values \\citep{Warner}. In our study of BF~Eri, the value of $K_2$ is beyond question. Scattering of its values derived with the use of the different data, is very small. Doppler tomography does not show any emission from the secondary star that can distort the absorption line profiles from the latter, neither in quiescence nor in the outburst. Furthermore, a non-uniform absorption distribution across the surface of the secondary would result in a non-sinusoidal radial velocity curve that is not observed in BF~Eri (Fig.~\\ref{FigRadVel}). All the obtained values of $K_1$ are also consistent with each other. Though the detected asymmetric emission structure of BF~Eri (particularly, the spot in the lower-left of the Doppler maps) may potentially influence the velocity determination, we believe we could avoid this as that strong emission spot is situated well inside of the chosen Gaussian separation. The correct phasing of the radial velocity curve also strengthens our confidence.  Alternatively, one may also calculate the mass ratio of BF~Eri using the observed rotational velocity of the secondary star, independently of $K_1$. Unfortunately, we were unable to derive a consistent value of the rotational velocity, getting instead a broad range of \\vsini\\ values. This gives us the $q$ values in the range of 0.33--0.62 that is consistent with the former solution, but also very uncertain, so we decided to reject it.  Thus, we have shown that BF~Eri contains a massive WD and an evolved secondary. This places this system among a small group of CVs with a high-mass WD (such as RU~Peg and U~Sco). Such systems are specially interesting since they are considered as possible Type Ia supernova progenitors (\\citealt{SNreview}, and references therein). The identity of the progenitor has become important for cosmology, since Type Ia supernovae are used as standard candles and their peak luminosity depends on the nature of the progenitor. The exploding star is now thought to be a carbon--oxygen WD that accretes mass from a binary companion until it approaches the Chandrasekhar limit, ignites carbon under electron-degenerate conditions, undergoes a thermonuclear instability, and disrupts completely. Further studies of the systems with massive WDs similar to BF~Eri are needed to understand if they evolve to SNe Ia.  \\subsection{$\\gamma$-velocity, proper motion and distance}  A decade ago \\citet{GamVel1} drew attention to the distribution of space velocities, as a means of probing the age profile of the CV population. \\citet{KolbStehle} derived the theoretically expected distribution of $\\gamma$-velocities and the dispersion of $\\gamma$ as a function of orbital period. In particular, they showed that the CVs having periods longer than the upper limit of the period gap ($P_{orb}\\geq$3 h) should be younger ($\\leq$1.5 Gyr), and therefore have a smaller line-of-sight velocity dispersion according to the empirical age-velocity dispersion relation (predicted value $\\sim$15 \\kms). Conversely, those CVs with orbital periods shorter than the lower limit of the period gap, should be older ($\\geq$ 3--4 Gyr) and show a larger velocity dispersion (predicted as $\\sim$30 \\kms). Paying a special attention to determine accurate absolute systemic velocities, \\citet{North2002} found its values for four long-period DNs. They obtained an average $\\gamma$-velocity of $\\sim$5 \\kms\\ with a dispersion of $\\sim$8 \\kms\\, indicating that the above postulate is indeed true.  In our study of BF~Eri, we have mostly followed the method of \\citet{North2002} and found the \\textit{absolute} $\\gamma$-velocity, after correction for the solar motion, to be $-94$ \\kms, referred to the dynamical local standard of rest (LSR). This value is very far away from any expectations of the theory (as set out by \\citealt{KolbStehle}). In principle, such a kick velocity could be caused by an event in which a significant fraction of the total mass of the system is violently, asymmetrically ejected. In the case of CVs, the progenitors indeed eject most of their mass during the common-envelope period, but it is supposed that this process does not alter the space velocity of the system. \\citet{KolbStehle} note, however, that this velocity may be affected by the cumulative effect of repeated nova eruptions with asymmetric envelope ejection.  It is worthy to mention that BF~Eri shows not only the high $\\gamma$-velocity but also has a substantial proper motion: $(\\mu_x, \\mu_y) = (+34, -97)$ \\masyr \\citep{Klemola2000,Hanson2004}. In order to convert this value to the space velocity we need to know the distance to the system, which we can evaluate here by oblique methods only.  One of these estimations is based on the \\textit{statistical} period-luminosity-colours relation of \\citet{Ak}. We have taken the $JHK$ magnitudes of BF~Eri from the 2MASS (Two Micron All Sky Survey) observations \\citep{2MASS2,2MASS1} and found the distance to the system, from the dependence of the absolute magnitude $M_J$ on the orbital period $P_{orb}$ and colours $(J - H)$ and $(H - K_s)$, to be $700\\pm200$ pc.  Alternatively, the distances can be estimated based on a \\textit{semi-empirical} donor sequence for CVs of \\citet{Knigge}. Unfortunately, Knigge's donor sequence is intended to describe the unique evolution track followed by \\textit{unevolved} secondaries, and he limited it to $P_{orb} \\la 6.2$ h. BF~Eri is a system with a longer period and likely an evolved secondary. However, we have slightly extrapolated Knigge's sequence and found from it and the $JHK$ magnitudes the lower limit of the distance $d$ to be $\\sim500$ pc. Applying the \\textit{empirical} offsets to the donor absolute magnitudes (see, however, comments by Knigge on this) we have obtained $d$ to be $\\sim900$ pc.  Finaly, the distances can also be derived from the absolute magnitude at outburst $M_V$(max) versus the orbital period relation of \\citet{Harrison}, if one intends to consider the flares of BF~Eri as DN outbursts (see corresponding discussion in Section~\\ref{OutBehSec}). Analysis of the AAVSO database of BF~Eri's observations reveals 8 events brighter than 13 mag with an averaged value of $12.5 \\pm 0.2$ mag. This yields the distance to be $650\\pm60$ pc.  Fortunately, all these independent estimations of the distance are consistent with each other, leading us to a final value of $700\\pm200$ pc. This corresponds to a transverse velocity $v_T=340\\pm100$ \\kms\\ or a space velocity of $\\sim$350 \\kms\\ with respect to the LSR, that is again unusually high. Based on kinematics, \\citet{BF-Eri} surmised that BF~Eri may belong to the halo population. Our Galactic components of the space velocity of the system $(U, V, W) = (-220, -280, 4)$ \\kms are slightly different from those obtained by \\citeauthor{BF-Eri} but also support a system motion being essentially in the Galactic plane and lagging far behind the rotation of the Galactic disc. In this connection the referee has pointed out that the high space velocity may be simply a sign of great age and not a result of any asymmetric mass ejections, that would be consistent with the secondary star having started to evolve prior to the time the system came into contact. Indeed, though \\citet{KolbStehle} do mention asymmetric nova explosions as a possible mechanism for increasing the dispersion in the $\\gamma$-velocity distribution, but they give no numbers, and it is not yet clear whether it is really plausible that the effect could produce a $\\gamma$-velocity as large as the one we found.  However, if such a high space velocity of BF~Eri was indeed caused by repeated nova eruptions then it might be worth searching for its past (recurrent) nova outbursts in the astronomical archive. Although the system is known only since 1940 and during the following 50 years there were no observations, some collections of photographic plates cover over a century of data. It is straightforward to look at plates taken over the years to see if past eruptions have been recorded.  %************************  AAVSO1  ************************************ \\begin{figure*} \\centering \\includegraphics[height=6.0cm]{aavso.eps} \\includegraphics[height=6.0cm]{aavsohist.eps} \\caption{\\textbf{Left:} The AAVSO light curve of BF~Eri. Blue stars represent the data labelled in the AAVSO database as V-magnitude, black filled dots - as Visual-magnitude, open dots - as inaccurate Visual-magnitude, green triangles - as the ``Fainter Than'' Visual-data. \\textbf{Right:} Histogram of the reliable AAVSO data for BF~Eri.} \\label{fig:aavso1} \\end{figure*} %************************  Chi2  ************************************   \\subsection{The outburst behaviour and the nature of BF~Eri} \\label{OutBehSec}  The current classification of BF~Eri as a DN is based on observations by Watanabe and Kato \\citep{Watanabe, Kato99, KatoUemura} who detected ``outbursts'' in the light curve of the system. We note that all those outbursts were of low amplitude (\\textit{1--1.5 mag}) forcing Kato to classify BF~Eri as ``a low-amplitude dwarf nova''. However from the formal side, DNs are an subset of CVs, which undergo recurrent outbursts of \\textit{2--8 mag}. The observed ``outbursts'', including the one observed during the \\textit{2006-Nov} set, are of much lower amplitude, which leads us to the question: are those events the real DN outbursts rather than some kind of flares?  It is now widely accepted that the reason for a DN outburst is a thermal instability in the accretion disc, which switches the disc from a low-viscosity to a high-viscosity regime \\citep{Smak, Osaki, Lasota}. From the observational side this is (usually) accompanied by substantial spectral changes occurring continuously during all stages of the outburst. More than half of DNs exhibit the transition from emission-line spectrum at quiescence to absorption-line spectrum at maximum. For the rest if any emission can be seen it is often in the form of emission lines buried in absorption troughs \\citep{Warner, MRM}. Quite often the spectral line profiles undergo dramatic changes as new emission sources arise in the system such as the inner hemisphere of the secondary star. This emission is caused by increased irradiation from the accretion regions during the outburst and can be seen in hydrogen and neutral helium lines through Doppler tomography \\citep{Neustroev}. Many systems show also strengthening of the \\HeII\\ and \\CIII/\\NIII\\ line emissions during an outburst.  None of these outburst signs was observed in BF~Eri. Thus the observed flares can hardly be associated with a DN outburst. This in turn raises a question on the classification of the system. Actually, BF~Eri does not also show many other usual observational signs of DNs. For example, in spite of the high spectral resolution used, all the emission lines show not the double-peaked profiles but single-peaked ones. From the above-estimated system parameters we can expect the double peak separation to be more than 500 \\kms, that is more than enough to be resolved. However, though the profiles are highly variable, they do not show any signs of the double-peaked accretion disc structure. In consequence, a ring-like emission distribution, characteristic of a Keplerian accretion disc about the WD, is also absent from the Doppler maps. On the other hand, the bulk of the Balmer emission is located to the lower left quadrant of the tomograms. Both the emission line profiles and the appearance of the Doppler maps resemble those in nova-like CVs which in turn resemble those typical for old novae at minimum light.  By definition, NLs should not display any DN outbursts. Their almost steady brightness is thought to be due to their mass transfer rate $\\dot{M}$ exceeding the upper stability limit $\\dot{M}_{crit}$. These stars are thus thermally and tidally stable. However, \\citet{Honeycutt98} reported that a significant fraction of NLs displays so-called ``stunted'' outbursts. \\citet{Honeycutt01} found many similarities between them and outbursts in ordinary DNs but the former are of much smaller amplitude (0.4--1 mag). We suppose thus that BF~Eri might be a NL system exhibiting ``stunted'' outbursts.  The AAVSO database has almost fourteen hundred BF~Eri observations obtained largely after 1999 \\citep{AAVSO}. We have analyzed the light curve compiled from the `Visual magnitude' and `V magnitude' data and found that the outburst activity of BF~Eri is quite complex (Fig.~\\ref{fig:aavso1}, left). Histogram of the AAVSO data reveals at least three types of flashes (Fig.~\\ref{fig:aavso1}, right). Besides the above-mentioned low-amplitude flashes with the peak magnitude $\\sim$13.0 mag and $\\sim$13.5 mag, there were also 8 events with the amplitude more than 1.6 mag but always less than 2.3 mag, with $V_{max}\\approx12$. Formally, the brightest flashes lie on the bottom edge of the accepted range of the outburst amplitudes in ordinary DNs and might be such outbursts. The nature of the low-amplitude flashes is still unclear; further investigation is needed.  Finally, we have obtained the periodogram, using the AAVSO data (Fig.~\\ref{fig:aavso2}). It shows significant excess power in the interval 60--90 d, reflecting a reliable stability of the flashes.  %************************  AAVSO1  ************************************ \\begin{figure} \\centering \\includegraphics{aavsopower.eps} \\caption{The Lomb-Scargle power spectrum of BF~Eri, derived from the AAVSO data.} \\label{fig:aavso2} \\end{figure} %************************  Chi2  ************************************"
    },
    "0801/0801.3234_arXiv.txt": {
        "abstract": "We present a method to include stereoscopic information about the three dimensional structure of flux tubes into the reconstruction of the coronal magnetic field. Due to the low plasma beta in the corona we can assume a force free magnetic field, with the current density parallel to the magnetic field lines. Here we use linear force free fields for simplicity. The method uses the line of sight magnetic field on the photosphere as observational input. The value of $\\alpha$ is determined iteratively by comparing the reconstructed magnetic field with the observed structures. The final configuration is the optimal linear force solution constrained by both the photospheric magnetogram and the observed plasma structures. As an example we apply our method to SOHO MDI/EIT data of an active region. In the future it is planned to apply the method to analyse data from the SECCHI instrument aboard the STEREO mission. ",
        "introduction": "Due to the low average plasma $\\beta$ the structure of the corona is determined by the coronal magnetic field. Knowledge of the structure of the coronal magnetic field is therefore of prime importance to understand the physical processes in the solar corona.  At the present time there is no general method available which allows the direct and accurate measurement of the magnetic field at an arbitrary point in the corona, although some progress has been made using radio observations above active regions (e.g. \\opencite{golub97}). We therefore have to extrapolate the coronal magnetic field from measurements taken at photospheric or chromospheric level.  If we want to do this we have to make assumptions about the current density in the corona. The low average plasma $\\beta$ allows us to assume that to lowest order the magnetic field is force-free, i.e. the current density is aligned with the magnetic field. With a few exceptions (e.g. \\opencite{zhao93}; \\opencite{zhao94}; \\opencite{petrie00}; \\opencite{zhao00}; \\opencite{rudenko01b}), most of the extrapolation and reconstruction methods proposed so far are based on this assumption including the use of potential fields (${\\bf j}= {\\bf 0}$) (e.g.\\ \\opencite{schmidt64}; \\opencite{semel67}; \\opencite{schatten69}; \\opencite{sakurai82}; \\opencite{rudenko01a}), linear force-free fields (e.g.\\ \\opencite{nakagawa72}; \\opencite{chiu77}; \\opencite{seehafer78}; \\opencite{semel88}; \\opencite{gary89}; \\opencite{lothian95}) and nonlinear force-free fields (e.g.\\ \\opencite{sakurai81}; \\opencite{wu85}; \\opencite{roumeliotis96}; \\opencite{amari97}; \\opencite{mcclymont97}; \\opencite{wheatland00}; \\opencite{yan00}).  Ideally, the information contained in a (perfect) vector magnetogram together with the force-free condition would be sufficient to calculate the coronal magnetic field. However, despite the increasing availability of vector magnetogram data, we are still far from this ideal situation due to both the quality of the data and to the difficulty of the calculation. It is therefore still easier to use line-of-sight magnetograms as input for potential or force-free extrapolation.  Potential fields are completely determined by fixing the line-of-sight component of the magnetic field \\cite{semel67}, but do not necessarily give good fits to observed emission structures. In the case of linear force-free fields, the normal (e.g.\\ \\opencite{chiu77}) or line-of-sight component \\cite{semel88} is not sufficient to determine the field uniquely and one has the freedom to choose a value for the linear force-free parameter $\\alpha$. Even though some methods have been suggested to determine $\\alpha$ by using the ambiguity of the full linear force-free solution to fit vector magnetogram data \\cite{amari97, wheatland99}, the usual method is to try to choose the value for $\\alpha$ in such a way that a subset of the field lines matches the observed emission pattern as good as possible (e.g.\\ \\opencite{pevtsov95}).  So far, our information about the emitting plasma structures has been largely limited to two-dimensional projections of intrinsically three\\--dimen\\-sio\\-nal objects. However, within the next few years the STEREO mission will hopefully give us the possibility to get three\\--di\\-men\\-si\\-onal information about the coronal plasma structures, in addition to magnetogram data which are routinely taken by ground- or space-based instruments.  In the present paper we want explore the possibility to determine a value of $\\alpha$ for a linear force-free field by comparing the reconstructed magnetic field to the observed three\\--di\\-men\\-sio\\-nal structure of coronal plasma loops. This is done by defining one or several three-dimensional curves in space representing the spatial structure of the observed loops, and a mathematical measure of the deviation of the reconstructed magnetic field lines from these three\\--di\\-men\\-sio\\-nal space curves. The value of $\\alpha$ is then determined by minimising the deviation from the observed loops. We emphasise that this is to be considered as a first step only and that a generalisation of the method to non-linear force-free fields is planned. As the STEREO mission is yet to be launched, the method is tested using the loop shapes deducted by \\inlinecite{aschwanden99} from applying dynamic stereoscopy to the solar active region NOAA 7986 using SOHO/MDI and SOHO/EIT data taken on 29, 30 and 31 August 1996.  The outline of the paper is as follows. In section \\ref{program} we describe the basic algorithm of the reconstruction method. A brief description of the method used to calculate the linear force-free fields is given in section \\ref{constalpha}. The code is then applied to the data of \\inlinecite{aschwanden99} in section \\ref{applications}. Conclusions and an outlook to future research is given in section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  In this paper we undertook a first step towards including three\\--di\\-men\\-sio\\-nal information from stereoscopic observations into a reconstruction of coronal magnetic fields from photospheric magnetic field measurements. Due to the low plasma $\\beta$ in the solar corona force-free magnetic fields can be used and in the present paper, we restricted our analysis to linear force-free fields for simplicity. The extrapolation method uses the line-of-sight photospheric magnetic field as a boundary condition.  As the coronal plasma has a very high conductivity, the magnetic field is frozen into the plasma. Therefore one can assume that the coronal plasma structures also outline the magnetic field. This is the fundamental assumption of our method. The idea of our reconstruction method is to take both the photospheric magnetogram and the three\\--di\\-men\\-sio\\-nal stereoscopic information into account to derive the magnetic field configuration in the solar corona. The magnetogram provides information regarding the strength and the distribution of the magnetic fields, whereas the three\\--di\\-men\\-sio\\-nal loop shapes restrict the current density. The method works iteratively in that we first calculate a magnetic field configuration from the line of sight magnetogram for a given value of $\\alpha$ without considering the stereoscopic information. To test the configuration we compare a magnetic loop of this magnetic field with the stereoscopic observations. This can be done in various ways of which we have presented two (foot-point method and loop distance method). The value of $\\alpha$ is then systematically optimised by minimising the difference between the observed and the reconstructed loop shape.  We have applied this method to loops deduced by \\inlinecite{aschwanden99} using the method of dynamic stereoscopy for the active region NOAA 7986 observed with SOHO/MDI and EIT on 1996 August 30. For the future we are planing to apply our method to data taken by the SECCHI instrument aboard the STEREO mission.  The results obtained for the \\inlinecite{aschwanden99} data are promising, but indicate the need for improving the method further by the use of non-linear force-free fields. Work along these lines is in progress and will be presented in a forthcoming publication."
    },
    "0801/0801.4552_arXiv.txt": {
        "abstract": "The Cygnus X-ray Emission Spectroscopic Survey (CyXESS) sounding rocket payload was launched from White Sands Missile Range on 2006 November 20 and obtained a high resolution spectrum of the Cygnus Loop supernova remnant in the soft X-rays.  The novel X-ray spectrograph incorporated a wire-grid collimator feeding an array of gratings in the extreme off-plane mount which ultimately dispersed the spectrum onto Gaseous Electron Multiplier (GEM) detectors.  This instrument recorded 65 seconds of usable data between 43-49.5 \\AA\\ in two prominent features.  The first feature near 45 \\AA\\ is dominated by the He-like triplet of \\ion{O}{7} in second order with contributions from \\ion{Mg}{10} and \\ion{Si}{9}-\\ion{Si}{12} in first order, while the second feature near 47.5 \\AA\\ is first order \\ion{S}{9} and \\ion{S}{10}.  Fits to the spectra give an equilibrium plasma at $\\log(T)=6.2$ ($kT_e=0.14$ keV) and near cosmic abundances.  This is consistent with previous observations, which demonstrated that the soft x-ray emission from the Cygnus Loop is dominated by interactions between the initial blast wave with the walls of a precursor formed cavity surrounding the Cygnus Loop and that this interaction can be described using equilibrium conditions. ",
        "introduction": "Supernovae greatly influence the dynamics within the interstellar medium (ISM).  Their ubiquitous nature and the size of their remnants allow them to influence multiple phases of the ISM while influencing the evolution and structure of the galaxy.  They are responsible for chemical enrichment of the ISM, play a major role in energy input, and are dominant sources for hot ionized gas, which fills much of the galaxy \\citep{McKeeOstr77}.  One of the most studied SNRs is the Cygnus Loop.  It is a nearby remnant at 540 pc \\citep{blair05}, quite large, filling $\\sim3^{\\circ}\\times3^{\\circ}$ \\citep{L97}, and is relatively unabsorbed ($E(B-V)=0.08$; \\cite{Fesen}).  The blast wave is currently encountering the surrounding ISM leading to bright emission at all wavelengths.  Therefore, the Cygnus Loop serves as an excellent test bed for SNR evolution theory.  A comprehensive and detailed imaging study of the remnant's X-ray and optical emission was performed by \\citet{L97,L98}.  One of the major findings of these studies is that the Cygnus Loop morphology is not indicative of the typical theoretical picture of a blast wave propagating through a uniform medium \\citep{sedov,spitzer}, but instead is an interaction of the blast wave with an inhomogeneous medium.  The precursor formed cavity is nearly spherical and surrounded by high density clumps and lower density gas.  This causes most emission to originate from a limb-brightened shell.  The optical emission occurs when the shock is decelerated rapidly in high density clumps.  The X-ray emission occurs in these clumps as the blast wave interacts with the cloud and hardens toward the interior as a reflection shock propagates inward reschocking the already shocked gas.  X-rays also originate as expected in the lower density gas that follows the theoretical model more closely.  The Cygnus Loop has been spectroscopically observed in the soft X-ray to some extent.  \\citet{vedder} obtained a high resolution (E/$\\Delta$E$\\sim50$) spectrum at energies below 1 keV using the Focal Plane Crystal Spectrometer (FPCS) on the Einstein Observatory, which observed a $3' \\times 30'$ bright northern section midway between the center and the limb.  They detected \\ion{O}{8} Lyman-$\\alpha$ at 653 eV, the resonance line in the He-like triplet of \\ion{O}{7} at 574 eV and a blend of \\ion{Ne}{9} lines between 905-920 eV.  They do not detect the forbidden or intercombination lines of the \\ion{O}{7} triplet.  Even though the detected fluxes are consistent with an equilibrium plasma with solar abundances at $T\\sim3\\times10^6$ K, the authors argue that the plasma is not yet in collisional ionization equilibrium (CIE) due to the lack of forbidden \\ion{O}{7} emission.  ROSAT Position Sensitive Proportional Counter (PSPC) observations by \\citet{L99} and \\citet{M01} show results consistent with the previous imaging studies.  A \\citet{L99} spectrum of the decelerated cloud shock was best fit by a Raymond-Smith equilibrium plasma at cosmic abundances \\citep{RaySmith} with $kT_e=0.06(+0.04, -0.02)$ keV and absorbing column $N_H=7\\times10^{20}(\\pm4\\times10^{20})$ cm$^{-2}$.  A spectrum extracted from a region interior to the cloud shock had a similar low temperature component, but also included a high temperature contribution, $kT_{e1}=0.25(+0.14, -0.06)$ keV, from the reflected shock, which hardens the interior emission.  The low temperature corresponds to a shock velocity of $v_s=230$ km s$^{-1}$, which has slowed considerably from the initial blast wave velocity of $v_s\\sim500$ km s$^{-1}$ \\citep{L98} resulting in an equilibrium plasma.  \\citet{M01} also find that equilibrium plasmas with cosmic abundances fit the low temperature components of their extracted spectra ($kT_e=0.043$ to $0.067$ keV) and that these contributions are greatest toward the limb.  However, they use ASCA data and a nonequilibrium model with depleted abundances to fit the interior harder spectra ($kT_e=0.27$ to $0.34$ keV).  The Chandra X-ray Observatory Advanced CCD Imaging Spectrometer (ACIS) also observed the Cygnus Loop \\citep{L02,leahy}.  \\citet{L02} used the Xspec \\citep{Xspec} spectral fitting software package and more specifically the MEKAL equilibrium model \\citep{Mewe85,Mewe86,kaastra,arnaud85,arnaud92,liedahl95} to fit the extracted spectra.  Again, the softer emission originates near the limb ($kT_e\\sim0.03$ keV) while the reflected shock hardens the interior spectrum ($kT_e=0.12$ to $kT_e=0.18$ keV).  The best spectral fits all required a depletion in oxygen.  Attempts were made at nonequilibrium models but they did not improve the fit statistics.  Furthermore, the fit values for the ionization parameter ($n_et$, where $n_e$ is electron density and $t$ is the elapsed time since the gas was initially shocked) were all $>10^{12}$ cm$^{-3}$ s where values of $\\gtrsim3\\times10^{11}$ cm$^{-3}$ s signify equilibrium.  The authors do not rule out the possibility of nonequilibrium but state that due to low spectral and spatial resolution the low temperature equilibrium plasma is indistinguishable from a higher temperature nonequilibrium plasma.  Using 21 different extraction regions \\citet{leahy} finds the same result; the spectra are best fit by an equilibrium MEKAL model with variable elemental abundances, and that there is no evidence for nonequilibrium.  The author also states that the abundances are considerably depleted and vary on small spatial scales suggesting that the region is geometrically complex with multiple clouds and even more shocks.  The final spectrum of note \\citep{M07} was taken using the Suzaku Observatory X-ray Imaging Spectrometer (XIS) CCD camera and has the highest spectral resolution other than \\citet{vedder}.   The authors clearly detected \\ion{O}{7} at $562\\pm10$ eV, \\ion{O}{8} at $653\\pm10$ eV, \\ion{C}{6} at $357\\pm10$ eV and \\ion{N}{6} at $425\\pm10$ eV, but the 1/4 keV band emission is unresolved.  Band ratio maps using strong emission features at different radii show that the outermost emission is dominated by the 1/4 keV emission and therefore a low temperature component.  Furthermore, the \\ion{O}{8}/\\ion{O}{7} ratio increases inward showing that the ionization state is higher in the interior.  The best fit models were two component nonequilibrium plasmas with variable abundances (\\textit{vnei} in Xspec).  The lower temperature components of the fits varied from $kT_e=0.10-0.15$ keV and the high temperature component ranged from $kT_e=0.18-0.34$ keV.  The temperature of both components increased towards the interior.  Column densities ranged from $N_H=3-6\\times10^{20}$ cm$^{-2}$ and abundances were heavily depleted.  The Cygnus Loop is clearly complex both spatially and spectrally.  However, spectral resolution is lagging spatial resolution.  Chandra can image fine structures indicative of different physical regions, but the nearly broadband spectra can only show trends.  Determining plasma diagnostics from model fits to these low spectral quality data is uncertain.  Higher spectral resolution is required not only to constrain the parameters of a model, but to test the assumed validity of the model.  Currently there is no efficient, well developed technology that permits high resolution x-ray spectroscopy from large solid angle sources, making spectra of diffuse x-ray sources rare.  In order to address this issue a soft X-ray spectrometer was designed for a rocket payload, the Cygnus X-ray Emission Spectroscopic Survey (CyXESS).  Scientifically, the payload was designed to observe emission in the 1/4 keV bandpass.  This is the least understood range of astrophysical soft X-ray energies.  A diffuse high resolution spectrum has never been achieved even though there is a large amount of flux in this band.  The soft emission from the Cygnus Loop is contained within a narrow shell and is dominated by the early interactions of the initial blast wave with the surrounding cavity.  Therefore, a global spectrum of the remnant will be dominated by the physics of this interaction. ",
        "conclusions": "\\label{disc} The results of the data analysis reveal a departure from cosmic abundances; S is enriched while Si and N are depleted.  The enrichment of S can be explained by confusion due to multiple temperature components.  The single temperature fit at $kT_e=0.14$ keV is intermediate in comparison to \\citet{L99} and \\citet{M01}, but consistent with the low temperature component from \\citet{M07}.  Therefore, there may be some contamination from a higher temperature component in the CyXESS data.  In this passband, increasing the temperature increases the contribution from second order oxygen while decreasing the contributions of other lines.  Therefore, a high temperature component will only add flux to the line complex at 44 \\AA, thus requiring additional flux in the form of a S enhancement in the 47.5 \\AA\\ complex in order to maintain the spectral shape defined by the fit.  Since there are only 2 features to fit, it is impossible to accurately discern the relative contribution from each plasma.  More spectral resolution is required so that individual lines can be used to define the components.  As for the second issue, even though some previous observations argue for an equilibrium plasma, especially in the case of the softest X-rays which occur in cloud shocks, nonequilibrium conditions could still be important.  If nonequilibrium is considered, then the ionization state of the gas should be decreased.  In nonequilibrium the electron temperature is higher than what is expected from the ion state of the gas and the gas is underionized.  Therefore, favoring lower ion species in the MEKAL model will test nonequilibrium conditions.  Also, within the lower ion species the ratio of higher energy line flux to lower energy line flux should be increased relative to equilibrium since the electrons have more energy to excite the ions to higher levels than typical in equilibrium.  In the case of Si, the lines that contribute the problematic long wavelength flux occur around 49-50 \\AA\\ and are \\ion{Si}{9}, \\ion{Si}{10} and \\ion{Si}{11}.  Decreasing the influence of these lines will definitely have a desired effect, but will also require an increase of major \\ion{Si}{7} and \\ion{Si}{8} lines at 52 \\AA, thus reintroducing the problem.  This occurs for other elements as well.  Favoring lower ion species will not explain the depletions, and furthermore, there is no evidence for nonequilibrium conditions contributing to these data.  Another explanation could be depletion into dust, especially in the case of Si depletion.  If dust is a major constituent of the shocked ISM then refractory elements could be contained within the dust and depleted from the gas phase.  However, the dust present in the cloud must be able to survive not only the initial blast wave, but also a reflection shock and sublimation over time in the hot gas (not to mention precursor winds, cloud-cloud interactions, cosmic rays and photodesorption \\citep{DraineSalp1,DraineSalp2}).  Given the conditions present in the Cygnus Loop, \\citet{DraineSalp1} show that thermal sputtering rates for dust grains in $\\log(T)=6.2$ gas are $\\sim$0.001$\\mu$m every 1000 years.  Therefore, small grains (size $\\lesssim$100\\AA) will be quickly destroyed with the mass fraction returning to the gas.  In addition, \\citet{DraineSalp2} have modelled dust sputtering as a function of blast wave velocity.  Their results show that a blast wave with shock velocity $>300$ km s$^{-1}$ will sputter nearly all graphite, silicate, and iron dust grains up to a size of 0.1 $\\mu$m in a cloud with density $n_H=10-100$ cm$^{-3}$.  Calculations for the initial blast wave velocity for the Cygnus Loop vary from 330 km s$^{-1}$ \\citep{L02} to 400 km s$^{-1}$ \\citep{Ku}.  This also suggests that if a significant fraction of refractory elements are depleted into dust in the dense clouds consituting the Cygnus Loop cavity walls, then the size distribution of dust particles favors large grains.  Infrared (IR) emission from the Cygnus Loop has been observed and \\citet{arendt} show that it can be explained by dust emission.  Using \\textit{IRAS} observations at 12 $\\mu$m, 25 $\\mu$m, 60 $\\mu$m, and 100 $\\mu$m, they find that there are two infrared (IR) components, one that correlates well with X-ray emission and another that correlates with the optical emission.  A lack of observed emission in the 12 $\\mu$m and 25 $\\mu$m bands suggests an underabundance of small grains.  This is supported by their models of the X-ray/IR correlated gas, which favor a minimum grain size of $\\sim$150 \\AA\\, consistent with our estimates.  However, the models used to fit the broadband observations assume all emission is due to thermal dust emission.  The authors state that line emission could theoretically contribute a significant fraction to the observed IR emission.  Therefore, constraining the dust fraction is impossible without higher resolution IR spectra, which should become available with upcoming Spitzer observations.  Depletion into dust cannot be ruled out and is supported by our data, especially in the case of silicate grains.  We see no evidence for graphite grains because there are no important C lines in the wavelength range of our spectrum.  The N and Si depletions necessary for our fit are due to a line complex of second order \\ion{N}{7} and first order \\ion{Si}{10} and \\ion{Si}{11}.  Therefore, given this confusion only an upper limit to the Si abundance can be applied suggesting a large depletion of at least 56\\% of the Si into grains."
    },
    "0801/0801.4764_arXiv.txt": {
        "abstract": "We use a set of high-resolution $N$-body simulations of binary galaxy mergers to show that the morphologies of the tidal features  that are seen around a large fraction of nearby, massive ellipticals  in the field, cannot be reproduced by equal-mass dissipationless mergers; rather, they are well explained by the accretion of disk-dominated galaxies. In particular, the arm- and looplike morphologies of the observed tidal debris can only be  produced by the kinematically cold material of the disk components of the accreted galaxies. The tidal features that arise from such ``cold-accretion'' events onto a massive elliptical are visible for significantly longer timescales than the features produced by elliptical-elliptical mergers (about 1-2 Gyr vs. a few hundred million years). Mass ratios of the order of $1:10$ between the accreting elliptical and the accreted disk are sufficient to match the brightness of the observed debris. Furthermore,  stellar population synthesis models and simple order-of-magnitude calculations  indicate that the colors of the tidal features generated in such minor cold-accretion events are relatively red, in agreement with the observations.  The minor cold-accretion events  that explain the presence, brightness, and structural and color properties of the tidal debris cause only a modest mass and luminosity increase in the accreting massive elliptical.  These results, coupled with the relative statistical frequencies of disk- and bulge-dominated galaxies in the field, suggest that massive ellipticals assemble most of their mass well before their tidal debris forms through the accretion of relatively little, kinematically cold material rather than in very recent, dissipationless major mergers. ",
        "introduction": "Bright elliptical galaxies are the most massive stellar systems in the Universe and host a large fraction of the stellar mass that has been produced since the Big Bang~\\citep{1998ApJ...503..518F,2006ARAaA..44..141R}. The boxy isophotes and slow internal rotation of the most massive ellipticals argue for dissipationless mergers playing a role in the last stage of their assembly history~\\citep{2006ApJ...636L..81N,2005MNRAS.359.1379K}. Semianalytic models indicate that major mergers that lead to massive ellipticals are typically not disk-disk mergers but rather elliptical-elliptical or mixed mergers~\\citep{2003ApJ...597L.117K}. Another scenario, as suggested by cosmological SPH simulations~\\citep{2007ApJ...658..710N}, is the  early build-up of ellipticals at high redshift with only minor growth in mass due to small accretion events below $z<1$. The red, stellar tidal arms and loops in the outskirts of many nearby bright ellipticals are widely believed to be fast-dissolving transient beacons of \\emph{recent} dissipationless mergers between smaller, equal-mass ellipticals~\\citep{2005AJ....130.2647V}.\\footnote{The merger rates derived in~\\cite{2005AJ....130.2647V} are calculated from a close galaxy pair analysis and hence should not suffer from a misidentification of the origin of the tidal features.} Such mergers are often called ``dry'' to highlight the absence of gaseous material (as opposed to ``wet'', or gas-rich, mergers). Numerical simulations of elliptical-elliptical mergers indicate that tidal debris disappears within a few hundred million years~\\citep{1995AaA...297...37C,2005AJ....130.2647V,2006ApJ...640..241B}. This would imply a substantial increase in the massive elliptical galaxy population in the last few billion years~\\citep{2004ApJ...608..752B}. This picture faces, however, several observational challenges~\\citep{2006AaA...453L..29C,2007ApJS..172..494S}.  Other studies~\\citep{2006ApJ...648..969K} had hinted at the gradual disruption of a tiny satellite galaxy by a massive elliptical as a source for red tidal features. Yet, these disruption events yield tidal debris too faint to match the typical surface brightnesses of the observed tidal features~\\citep{2005AJ....130.2647V}. Early numerical experiments~\\citep{1988ApJ...331..682H} showed that sharp tidal arms and loops could arise from the tidal disruption of dynamically cold stellar disks. These experiments however adopted a spherical model for the elliptical galaxy and lacked the proper resolution to robustly establish the morphology and lifetime of the tidal streams. Moreover, those early studies were carried out before the current cosmological framework, the cold dark matter model with a cosmological constant, was established.  We use a set of 13 high-resolution $N$-body simulations of galaxy-galaxy mergers between a massive elliptical galaxy and an elliptical or disk galaxy companion to investigate the origin of tidal features observed in bright ellipticals. The companion mass ranges between 10\\% and 100\\% of the mass of the main elliptical (Table~\\ref{tab:modParam}). The numerical galaxy models are constructed to properly match the observed structural properties of both elliptical~\\citep{1992ApJ...399..462B} and disk galaxies~\\citep{1998MNRAS.295..319M}; in particular, the main massive elliptical has a stellar mass of $2.7\\times{}10^{11}$ $M_\\odot$, an effective radius of 9 kpc, and a stellar mass-to-light ratio of 4 in the $R$-band. These parameters ensure that the model galaxy lies on the observed fundamental plane of ellipticals~\\citep{1987ApJ...313...59D,1987ApJ...313...42D}. The stellar component of the galaxies is embedded in a cold dark matter halo with properties derived from cosmological simulations~\\citep{1997ApJ...490..493N}.  The initial orbits in the mergers are set to be either parabolic, as expected in low-density environments, or eccentric, with an apocenter-to-pericenter ratio of 6:1, as is typically seen in cosmological simulations in denser environments, i.e., in galaxy groups and in the cores of galaxy clusters~\\citep{1998MNRAS.300..146G}.  \\begin{table*} \\begin{center} \\caption{Parameters of the Disk Galaxy Models\\label{tab:modParam}} \\begin{tabular}{lcccccccccccccc} \\tableline \\tableline & & & & & & & & & & & & & &\\\\[-0.20cm] & $r_{200}$ & & & & & & & & & & & $h_\\textrm{DM}$ & $h_{D}$ & $h_{B}$ \\\\ Description & (kpc) & $c$ & $\\lambda$ & $m_D$ & $R_d$ & $z_0/R_d$ & $m_B$ & $r_b/R_d$ & $N_\\textrm{DM}$ & $N_{D}$ & $N_{B}$ & (kpc) & (kpc) & (kpc) \\\\ & & & & & & & & & & & & & &\\\\[-0.20cm] \\tableline & & & & & & & & & & & & & &\\\\[-0.12cm] 1:4 disk/bulge & 206.3  & 12 & 0.031 & 0.04  & 2.59 & 0.1 & 0.008 & 0.2 & $1\\times{}10^6$ & $1\\times{}10^5$ & $1\\times{}10^5$ & 0.13 & 0.13 & 0.13 \\\\ 1:4 disk   & 273        & 8  & 0.0398 & 0.03 & 5.64 & 0.15 & 0   & --- & $2.8\\times{}10^6$ & $2\\times{}10^5$ & 0              & 0.3  & 0.2 & --- \\\\ 1:10 disk/bulge & 171.4 & 12 & 0.031 & 0.04  & 1.84 & 0.1 & 0.008 & 0.2 & $1\\times{}10^6$ & $1\\times{}10^5$ & $1\\times{}10^5$ & 0.13 & 0.13 & 0.13 \\\\ 1:10 disk  & 200        & 8  & 0.0398 & 0.03 & 4.14 & 0.15 & 0   & --- & $2.8\\times{}10^6$ & $2\\times{}10^5$ & 0              & 0.3  & 0.2 & --- \\\\ & & & & & & & & & & & & & &\\\\[-0.20cm]\\tableline \\end{tabular} \\tablecomments{The first column describes the model type and its (stellar) mass ratio with respect to the mass of the massive elliptical used in the merger set up. The other columns denote the virial radius, the concentration, the halo spin, the disk mass fraction, the disk scale length, the disk height in units of disk scale length, the bulge mass fraction, the bulge scale length in units of disk scale length, the number of dark, disk, and bulge particles, and their gravitational softenings. We assume that disk material and dark matter have the same specific angular momentum, i.e. $j_D=m_D$. The Hubble parameter is $H=70$ km s${^-1}$ Mpc$^{-1}$.} \\end{center} \\end{table*}  This paper is organized as follows: In section \\ref{sect:initialCond} we describe the set-up of the simulated mergers. In section \\ref{sect:analysis} we define the tidal parameter and the merger time, which are of relevance in the further discussion. In section \\ref{sect:tidalStrength} we discuss the strength and lifetime of the tidal features, their morphology in section \\ref{sect:tidalMorph}, and their expected $B-R$ color evolution in section \\ref{sect:tidalColor}. In section \\ref{sect:fraction} we estimate the expected relative frequency of elliptical-elliptical and elliptical-disk mergers based on statistical studies of galaxies in the field. We discuss our results and conclude in section \\ref{sect:discussion}. ",
        "conclusions": "\\label{sect:discussion}  We have found major differences between the elliptical-elliptical and elliptical-disk mergers in both morphology of tidal tails and their timescales of observability. Only mergers that involve a dynamically cold disk produce strong tidal features with loops and arms resembling those observed around bright ellipticals~\\citep{2005AJ....130.2647V}; these features last much longer (1-2 Gyr) than the smooth features produced in hot-hot mergers.  Dynamical friction and tidal stripping determine the specifics of the merger evolution. For a given satellite galaxy type, either disk or elliptical, the tidal tails generated by less massive satellites extend to larger distances from the center of the accreting elliptical, as the smaller the satellite, the longer it needs to sink to the bottom of the gravitational potential. However, the smaller the satellite mass, the weaker the tidal debris, since less mass is available to produce the tidal features. We find that satellite masses of order one fourth of the mass of the parent elliptical exhibit the most long-lived tidal features.  Since mergers with disk-dominated, relatively low mass companions produce the strongest and longest-lasting tidal streams, such mergers are expected to be the typical events that are revealed in observations of real galaxies. In first approximation, the probability that the main elliptical merges with a companion of a given type will be proportional to the number density of galaxies of that type. The ratio of observable tidal features that are expected from mergers involving a disk galaxy, relative to mergers that involve another elliptical, can thus be approximated by the product of the ratio between the respective tidal lifetimes and the ratio of the respective number densities. Using published luminosity functions for early-type and disk galaxies~\\citep{2006MNRAS.368..414D} and typical mass-to-light ratios~\\citep{2001ApJ...550..212B}, we estimate that, at the mass scale of the host galaxy and  for a 1:4 merger, $>60\\%$ of the observed tidal features stem from an encounter with a disk-dominated system; the fraction increases up to $>80\\%$ for 1:10 mergers. We thus expect that the merger fraction that is estimated from observed tidal features around bright ellipticals reflects primarily the fraction of spheroids that underwent (relatively minor) accretion events of dynamically cold, disk galaxies over the past several billion years, rather than very recent major mergers between elliptical galaxies.  \\begin{figure*}[htbp] \\begin{center} \\begin{tabular}{cc} \\includegraphics[width=80mm, bb=124 242 466 533]{f10a} & \\includegraphics[width=80mm, bb=124 242 466 533]{f10b} \\\\ \\end{tabular} \\caption{(Relative) number densities of disk and early-type galaxies. ($a$) Stellar mass functions of early-type (red dashed line) and disk (blue solid line) galaxies in the local universe~\\citep{2006MNRAS.368..414D}. The vertical, black dot-dashed line corresponds to galaxies of $2.7\\times{}10^{11}$ $M_\\odot$. The inset shows the ratio of the two stellar mass functions and indicates its value for a 4 times (magenta dashed line) and 10 times (green dot-dashed) less massive satellite galaxy. The ratio lies in the range of unity for the considered masses and increases with decreasing stellar mass. ($b$) Relative number density $f_n(R)$ of disk vs. early-type galaxies for a given mass ratio $R$ after integrating over all masses of the elliptical host galaxy (black solid line). Mass ratios of 1:4 and 1:10 are indicated by the green dot-dashed and the magenta dashed lines. The respective values of the average relative number density are 2.0 and 3.0, respectively. Under the assumption that the merger probabilities are proportional to the number densities, a typical 1:4 or 1:10 merger with an elliptical galaxy will involve much more likely a disk than another early-type system.} \\label{fig:numberDens} \\end{center} \\end{figure*}  Our simulations show that the observable tidal debris mostly originates from regions of the disk corresponding to about 2-5 exponential scale lengths. Using observed photometric profiles of spiral~\\citep{1994AaAS..106..451D} and elliptical~\\citep{1990AJ....100.1091P} galaxies and stellar synthesis models~\\citep{2003MNRAS.344.1000B}, we find that the $B-R$ color of the tidal features is rather similar to the color of the main accreting elliptical at several half-light radii from its center, where the tidal tails become visible. We estimate that both the average $B-R$ color difference between tidal debris and main elliptical's stars, and its standard deviation, are only of order $\\sim{}$ 0.3 mag. The tidal tails produced by the accretion of disk galaxies are thus ``red'' in $B-R$ color.  Disk satellites responsible for the observed tidal features would likely host a gas component, which is not included in our $N$-body simulations.  Such a gas component might cause a burst of star formation in the early phases of the interaction~\\citep{1991ApJ...370L..65B}. Simulations show that tidally triggered stellar bars can develop soon after the first pericenter passage and drive substantial gas inflow within the central $\\sim$1 kpc of the satellite galaxy~\\citep{1996ApJ...471..115B,2005ApJ...623L..67K}, where the gas is quickly exhausted in a burst of star formation. Star formation is also expected in the tidal tails at these early stages of the merger. Ram pressure stripping in the hot halo of  the massive elliptical will however remove the reservoir of diffuse gas and thus quench star formation at later times~\\citep{2008ApJ...672L.103K}. Ram pressure becomes increasingly efficient as tidal shocks decrease the potential well of the satellite at each subsequent pericenter and is expected to remove most if not all the gas that has not plunged into its central region~\\citep{2007Natur.445..738M}. This is in agreement with the absence of a  young stellar component in -- and the red colors of -- the observed tidal features. We predict, however, that very faint and diffuse gaseous tidal tails and plumes should be detectable around elliptical galaxies with red stellar tidal signatures.  After the merger the newly formed stars in a disk-dominated progenitor might affect the overall $B-R$ color of the remnant. Making a sensible guess for the parameters of the central starburst, an $e$-folding time of 100 Myr and an initial star formation rate of 20 $M_\\odot$yr$^{-1}$, we find  that within 500 Myr the $B-R$ color of the remnant will closely resemble the original $B-R$ color of the elliptical merger progenitor to less than 0.1 mag (Fig.~\\ref{fig:RedSelection}). In addition, a nonnegligible fraction of nearby elliptical galaxies are found to host a few, up to a few tens, percent of stellar mass in young stars~\\citep{2000AJ....120..165T}.  The existence of red tidal features around the most massive elliptical galaxies can thus be well explained with minor accretion events of disk-dominated satellite galaxies. Consequently, our work supports an early ($z>1$) mass assembly epoch for the most massive galaxies in the Universe with a later accretion of smaller mass units that do not contribute significantly to the total mass. Future observations will have to clarify the relative importance of dry and wet mergers at such early epochs in the formation of the bright elliptical galaxy population."
    },
    "0801/0801.3658_arXiv.txt": {
        "abstract": "High frequency quasi-periodic oscillations (HFQPOs) have been detected in microquasars and neutron star systems. The resonance model suggested by Klu\\'{z}niak \\& Abramowicz (2000,2001) explains twin QPOs as two weakly coupled nonlinear resonant epicyclic modes in the accretion disk. Although this model successfully explains many features of the observed QPOs, it still faces difficulties and shortcomings. Here we summarize the aspects of the theory that remain a puzzle and we briefly discuss likely developments. ",
        "introduction": "\\label{sec:intro}  Many galactic black hole and neutron star sources in low-mass X-ray binaries show quasi-periodic variability in their observed X-ray fluxes. Some of the quasi-periodic oscillations (QPOs) are of the order of few hundred Hz (hence the name high frequency) and often come in pairs of twin peaks in the Fourier power spectra \\citep[see][for a review]{van00,rem06}. High frequency QPOs draw much interest because their frequencies lie in the range of orbital frequencies few Schwarzschild radii outside the central source. Furthermore for different sources they scale with $1/M$, where $M$ is the mass of the central compact object \\citep[][]{mcc06}. These two characteristics make QPOs attractive tools to possibly test General Relativity in strong fields. Many models have been proposed in order to understand HFQPOs: some of them involve orbital motions \\citep[e.g.,][]{klu90,ste98,lam03,sch04}, others consider accretion disk oscillations \\citep[e.g.,][]{wag01,kat01,rez03,kat07,tas07}. None is definitive. Klu\\'{z}niak \\& Abramowicz (2000,2001) proposed that certain properties of HFQPOs in neutron stars could be explained by resonant motions of accreting fluid in strong gravity, and predicted that in black hole systems two HFQPOs in a frequency ratio of 1:2, 1:3, n:m,  could be detected. \\citet{abr01} observed that the newly discovered HFQPOs in GRO J 1655-05 \\citep{str01} were in a $3:2$ ratio, and proposed a more specific model of resonance. This idea guided successive studies \\citep[e.g.,][]{abr03,reb04,hor04,klu04} aimed at clarifying and expanding the theory. Wlodek Klu\\'{z}niak talked about the accomplishments of the model  (see the contribution in the same Proceedings), here we focus on its limits.   \\begin{figure*} \\plotone{bursa.eps} \\caption{Lower ($\\nu_{down}$) versus upper ($\\nu_{upp}$) frequencies in observed twin peak HFQPOs. The four microquasars lie on the lower left part of the plot, on the line with slope $3:2$ . Neutron star systems (all the others) cross the $3:2$ line, but at different times of observation may be in a different position and ratio (courtesy of Michal Bursa). } \\label{fig:bursa} \\end{figure*} ",
        "conclusions": "\\label{sec:concl}  The nonlinear resonance model for high frequency QPOs encounters many internal difficulties. This may mark the end of the model or the beginning of a new era.\\\\ In our opinion the presence of twin peaks in (almost) commensurate ratio and the detection of subharmonics still strongly support the idea that weak nonlinear resonance is a fundamental ingredient in the physics of HFQPOs.\\\\ It is like smelling chocolate in a birthday cake: the smell is the unmistakable signature of the presence of chocolate. However the other (mathematical and physical) ingredients, their relative importance and a draft of the recipe (mechanisms) are needed to know  which cake it is."
    },
    "0801/0801.3870.txt": {
        "abstract": "We calculate the advances in near-infrared astronomy made possible through the use of fibre Bragg gratings to selectively remove hydroxyl emission lines from the night sky spectrum.  Fibre Bragg gratings should remove OH lines at high resolution ($R=10,000$), with high suppression (30dB)  whilst maintaining high throughput ($\\approx90$ per cent) between the lines.  Devices currently under construction should remove 150 lines in each of the $J$ and $H$ bands, effectively making the night sky surface brightness $\\approx 4$ magnitudes fainter.  This background reduction is greater than the improvement adapative optics makes over natural seeing;  photonic OH suppression is at least as important as adaptive optics for the future of cosmology.  We present a model of the NIR sky spectrum, and show that the interline continuum is very faint ($\\approx 80$ \\bright\\ on the ecliptic plane).  We show that OH suppression by high dispersion, i.e.\\ `resolving out' the skylines, cannot obtain the required level of sensitivity to reach the interline continuum due to scattering of light.  The OH lines must be suppressed prior to dispersion.  We have simulated observations employing fibre Bragg gratings of first light objects, high redshift galaxies and cool, low-mass stars.  The simulations are of complete end-to-end systems from object to detector.  The results demonstrate that fibre Bragg grating OH suppression will significantly advance our knowledge in many areas of astrophysics, and in particular will enable rest-frame ultra-violet observations of the Universe at the time of first light and reionisation. ",
        "introduction": "The future of cosmology is dependent in part on the ability to achieve deep observations at near-infrared  (NIR) wavelengths.  As we strive to observe the early Universe, our observations must be performed at increasingly longer wavelengths, as the diagnostic optical spectroscopic features become more redshifted.  At very high redshifts, observations in the optical become futile due to the redshifted Lyman limit.  Figure~\\ref{fig:observedlambda} shows the observed wavelengths of several useful spectroscopic lines as a function of redshift, illustrating the requirement for NIR observations.  \\begin{figure} \\centering \\includegraphics[scale=0.3,angle=270]{observedlambda.eps} \\caption{The observed wavelength of several useful spectral lines as a function of redshift.  The shaded areas show the wavelength coverage of standard zJH filters.} \\label{fig:observedlambda} \\end{figure}  The need for NIR observations is well documented.  Over the last few years there have been several major reviews of the current status of our knowledge of the Universe, and in particular of gaps in our understanding, serving as a starting point for strategic plans for future observations (e.g.\\ the Australian Decadal Plan for Astronomy 2006-2015\\footnote{http://www.atnf.csiro.au/nca/DecadalPlan\\_web.pdf}, the UK PPARC roadmap\\footnote{http://www.so.stfc.ac.uk/roadmap/rmQuestionsHome.aspx}, the US National Research Council's `Astronomy and Astrophysics in the New Millennium'  \\citealt{aanm01} and `Connecting Quarks to the Cosmos' \\citealt{cqwc03}).  These reviews provide similar lists of the `big questions' for the future of astronomy, showing a consensus of opinion at an international level.  Some of the key issues they identify are: \\emph{ What is the nature of dark energy and dark matter? How and when do the first stars form? How and when did the Universe reionise? How and when do galaxies assemble, and how do they evolve? What is the relationship between super-massive black-holes and galaxies? How do stars and planetary systems form? How common are planetary systems capable of supporting life? }   The success of answering the questions above is reliant on NIR observations.  The first half of the list depends upon the ability to observe the high redshift Universe, whilst the problems concerning planetary systems will also benefit from dust-penetrating NIR observations, especially since planetary sized bodies emit most of their radiation at NIR wavelengths.    Observations of high redshift galaxies carried out in the NIR offer significant advantages over optical observations: extinction due to local dust is smaller in the NIR; $k$-corrections are smaller, and therefore less critical; adaptive optics is easier to achieve in the NIR.  However, there are also significant disadvantages  in NIR astronomy, the chief concern being the high background level.  Table~\\ref{tab:skybright} lists typical surface-brightnesses for a good observing site, as used in the ESO ELT exposure time calculator\\footnote{http://www.eso.org/observing/etc/bin/ut3/conica/script/eltsimu}.  \\begin{table} \\begin{center} \\caption{Typical sky-brightnesses at a good observing site.} \\label{tab:skybright} \\begin{tabular}{ccc} & \\multicolumn{2}{c}{Surface brightness/ mag arcsec$^{-2}$} \\\\ Band & Natural & No OH emission \\\\ \\hline U  &     21.5& -- \\\\ B    &   22.4& -- \\\\ V      & 21.7& -- \\\\ R     &  20.8& 21.7 \\\\ I      & 19.9 &  21.7 \\\\ J     &  16.0 &  22.1\\\\ H     &  14.0&  22.2\\\\ \\end{tabular} \\end{center} \\end{table}  At wavelengths longer than the $H$ band, thermal emission from the telescope and surroundings dominate the background emission.  Therefore the detector and camera must be designed to operate at cryogenic temperatures -- a significant technical challenge.  At shorter wavelengths the thermal background is less problematic but instead there is the supplementary frustration of very bright hydroxyl emission (\\citealt{mei50}; \\citealt{duf51}), resulting from the vibrational decay of an excited OH-radical, formed from the combination of a hydrogen atom with an ozone molecule.  Furthermore the intensity of the OH emission is highly variable (see section~\\ref{sec:ohvar} below), making accurate sky subtraction very difficult (e.g.\\ \\citealt{dav07}).  The intensity and variability of the NIR background have traditionally restricted the depth of NIR observations compared to optical observations.  However, two attributes of the NIR background spectrum suggest that in principle it should be possible to obtain very deep observations in the NIR.  Firstly the OH spectrum comprises nearly all ($\\approx 99.9$ per cent from our model, see \\S~\\ref{sec:sky}) of the $J$ and $H$ band background.  The interline continuum, due mainly to zodiacal scattered light, is expected to be very faint, e.g.\\ the models of \\citet{kel98} give a flux density of only $\\approx 77$ \\bright\\ on the ecliptic plane.  Secondly, the OH lines are intrinsically very narrow (\\citealt{tur83}), meaning that large sections of the NIR sky spectrum will be very dark if the OH lines can be efficiently removed, thus enabling deep observations.  There are several methods which can, in principle, exploit the dark interline continuum for deep NIR observations, including narrow-band filters which allow observations of a very short bandpass between lines (e.g.\\ DAzLE; \\citealt{hor04}), or space-based telescopes (e.g.\\  JWST; \\citealt{gar06}).  Here we discuss ground-based OH suppression, by which we mean the selective removal (by means of low transmission, reflection etc.) of very narrow regions of the spectrum corresponding to the OH lines, whilst maintaining high throughput in the interline regions.  There have been several methods proposed to achieve OH suppression, including spectroscopic masking at high dispersion (e.g.\\  \\citealt{mai00}), rugate filters (\\citealt{off98}) and multi-notch holographic filters (\\citealt{blai04}).  We discuss these methods in section~\\ref{sec:old}; each offers potential improvements over natural conditions, but in practice there are problems in achieving \\emph{efficient} OH suppression by any of these means.  By efficient we mean achieving high suppression of the lines whilst maintaining high interline throughput, for a large number of lines ($\\gsim 100$).   In this paper we make the case for using fibre Bragg gratings (FBGs, see \\citealt{bland04}) to achieve OH suppression (\\S~\\ref{sec:fbgs}); FBGs  do not suffer from any of the drawbacks that have prevented successful application of the methods mentioned above.  There is now a real case for believing that highly efficient OH suppression is close to being demonstrated.  Before discussing methods of OH suppression, we first review the background sources contributing to the NIR in section~\\ref{sec:sky}.  We then address the scattering properties of dispersive optics (\\S~\\ref{sec:scatter}), the proper understanding of which is essential to achieve efficient OH suppression.   The study of scattering shows that the OH lines must be suppressed prior to dispersion; the combination of the extremely bright OH lines and the strong scattering wings, that are an unavoidable consequence of any diffraction system, make attempts to remove or `resolve out' the OH skylines after dispersion futile.  We then make the case for OH suppression, i.e.\\ selectively removing the OH lines prior to dispersion (\\S~\\ref{sec:ohsupp}), including a review of previously proposed methods (\\S~\\ref{sec:old}), and their limitations.   We then show that fibre Bragg gratings offer the prospect of cleanly and efficiently removing OH lines photonically  (\\S~\\ref{sec:fbgs}).   We calculate the required level of suppression and spectral resolution to achieve scientifically useful observations.  We summarise our case in section~\\ref{sec:summary}, and conclude with some remarks on the future in section~\\ref{sec:real}. ",
        "conclusions": "We have demonstrated the possibility to achieve extremely deep near-infrared observations using fibre Bragg gratings.  Hitherto, attempts to obtain deep NIR observations have been frustrated by the extremely bright and variable OH emission from the atmosphere.  We have shown that in order to efficiently remove OH lines from astronomical spectra the OH lines must be suppressed prior to the light entering the camera.  Scattering from dispersive optics results in Lorentzian line profiles, the wings of which are much brighter than the true interline continuum, and in most cases are brighter than the sources being observed.  The variability of the OH lines makes it extremely difficult to subtract cleanly.  FBGs can suppress OH lines at high resolution, $R=10,000$, high suppression (30dB) and with high throughput out of the suppression regions ($\\approx 90$ per cent).  Suppressing 150 lines in the $J$ or $H$ band, with a maximum suppression of 30dB, decreases the sky background by a factor of $\\approx 45$, or 4 magnitudes.  Achieving this will have the potential to extend NIR observation to objects orders of magnitude fainter than currently possible. The resulting observations will  have a great impact on many areas of astronomy, and we highlight a few now.  \\subsubsection{First light and reionisation}  FBGs offer the opportunity to study the rest-frame UV light from objects at the time of reionisation.  Figure~\\ref{fig:qsosim} shows a simulated spectrum of a $z=11$ QSO with a clearly visible \\citet{gun65} absorption trough.  Spectroscopy with FBGs would therefore be ideally suited to identifying first light sources, performing follow-up spectroscopy from surveys of \\lya\\ emission (e.g.\\ \\citealt{tan05}), continuum drop-out due to the Lyman-break (e.g.\\ \\citealt{sta04}) and ultra-narrow band surveys (e.g.\\ \\citealt{hor04}).  A spectrum of a complete Gunn-Peterson trough would confirm the redshift of the source, which would allow estimates of the luminosity function and number-magnitude relations of the first sources of light in the Universe.  If the sources are bright enough, then it may be possible to also measure  the line profile of \\lya\\, which is expected to be distinctively asymmetrical due to resonant scattering bluewards of the line-centre.  The intrinsic width of the line itself will depend upon the type of source; quasars will have broad \\lya\\ lines, FWHM$\\approx 1500$km s$^{-1}$, whereas galaxies will have narrower lines, FWHM$\\approx 100$km s$^{-1}$.  The difference in line widths leads to the possibility of distinguishing the type of sources responsible for ionisination.  For the purpose of confirming the presence of a redshifted \\lya\\ line it is necessary to resolve the lines sufficiently to measure the asymmetry, requiring resolutions of $R\\gsim 300000/100 =  3000$.  The combination of FBGs and  ELTs will make it possible to thoroughly explore the ionisation history of the IGM using NIR spectroscopy of first light objects.  The detection of complete Gunn-Peterson absorption troughs in the spectra of high redshift quasars (\\citealt{bec01}) show that the IGM was partly neutral at $z\\approx 6.5$.  However, a neutral fraction of only $x_{\\rm HI}\\approx 10^{-4}$ is all it takes for complete absorption due to the high oscillator strength of the \\lya\\ transition.  Thus to probe further back in time to uncover the details of the ionisation of the IGM, we must look to other means.  \\citet{oh02} suggest that metal absorption lines may provide a useful means of measuring the ionisation state of the IGM.  O{\\sc i} and Si{\\sc ii} have ionisation potentials close to that of neutral hydrogen, and thus measuring the neutral fraction of either of these metals could be used to infer the neutral fraction of hydrogen.  Furthermore there are useful lines just longer than \\lya\\ so a single measured spectrum may incorporate \\lya\\ 1215.67\\AA, O{\\sc i} 1302.17\\AA\\ and Si{\\sc ii} 1260.42\\AA.  These metals have previously been observed in the spectra of high redshift quasars (\\citealt{pet01}) verifying their usability for this purpose.  Since it is unlikely that metals will be smoothly distributed through the IGM, the signature will be in the form of a forest of lines rather than a Gunn-Peterson trough.  A measurement of the equivalent width of the lines, or even simply the number of lines will provide precious information on the ionisation state.  The detection or non-detection of a radiation damping wing has the potential to divulge important information on the ionisation state of the IGM.  If a radiation damping wing is found it will provide a very accurate measure of the neutral fraction of hydrogen (\\citealt{mir98}).  If the ionising sources reside in significant cosmic Str\\\"{o}mgren spheres then a damping wing will be impossible to measure (\\citealt{cen00}; \\citealt{mad00}).  However, the line profile may still be identified as \\lya\\ due to its asymmetrical shape, and thus the lack of a damping wing can place constraints on the ionising fluxes of the sources.  If the sources are galaxies then it may also place constraints on the clustering of the objects, since it is unlikely that a single primordial galaxy would produce a significant ionised region on its own.  If spectra can be obtained with a sufficiently high signal-to-noise, then the tail of the damping wing may provide detailed information on the size of the Str\\\"{o}mgren spheres.   Figure~\\ref{fig:qso30m} shows the same simulation as described in section~\\ref{sec:qsosim}, except the telescope was assumed to have a diameter of 30m.  Metal absorption lines are now visible in the spectrum.%, such that the ionisation and enrichment of the IGM may be measured as described above.   \\begin{figure*} \\centering \\includegraphics[scale=0.7]{qso30mzoom.eps} \\caption{A simulated spectrum of an $H=24.6$ Vega mag, $z=11$ QSO as observed by the system described in \\S~\\ref{sec:sys} with a 30m telesscope.  The exposure time was 70hr, and the spectral resolution was $R=1000$.  The top panel shows a system with FBGs and the bottom panel shows an identical system without FBGs.  The red lines indicate the true object spectrum.} \\label{fig:qso30m} \\end{figure*}   \\subsubsection{Galaxy evolution}  Deep NIR imaging and spectroscopy with FBGs will allow galaxy evolution to be studied over most of the history of the Universe.  Figure~\\ref{fig:galsim} shows a $H=23$ Vega mag galaxy at $z=3$.  Absorption lines are clearly visible.  It would thus be possible to measure galaxy masses, luminosities and star-formation rates from emission lines, and to derive the evolution of the star-formation history, the luminosity function, the \\citet{fab76} relation and if high spatial resolution imaging is available, the fundamental plane.  \\subsubsection{Cool stars}  NIR infrared observations are essential for identifying cool stars such as L, T and Y dwarfs.  Measuring the number density of such objects as a function of mass is a critical test of our understanding of both star and planet formation processes.  NIR spectroscopy is essential for fully characterising their temperature, surface gravity and atmospheric conditions.  The use of FBGs would allow observations of objects orders of magnitude fainter than currently possible (see Figure~\\ref{fig:t5sim}), thus extending the survey volumes and numbers of such objects, and placing strong constraints on the low mass end of the initial mass function."
    },
    "0801/0801.2433_arXiv.txt": {
        "abstract": "We extend the covariant, parametrized post-Friedmann (PPF) treatment of cosmic acceleration from modified gravity to an arbitrary admixture of matter, radiation, relativistic components and spatial curvature.  This generalization facilitates the adaptation of Einstein-Boltzmann codes for solving CMB and matter perturbations in the linear regime.   We use such a code to study the effect of metric evolution on the CMB through the integrated Sachs-Wolfe effect.   We discuss the ability of modified gravity to alter the low multipole spectrum, including lowering the power in the quadrupole. From a principal component description of the primary metric ratio parameter, we obtain general constraints from WMAP on modified gravity models of the acceleration. ",
        "introduction": "\\label{sec:introduction}  In the absence of theoretically compelling models for the acceleration of the expansion, it is useful to have a parametrized description of possible deviations from the standard cosmological constant model.   To explore modified gravity explanations of the acceleration we need a parametrized post-Friedmann (PPF) description of gravity that parallels the parametrized post-Newtonian description of solar system tests.  Many attempts to parametrize deviations phenomenologically (e.g.~\\cite{IshUpaSpe06,KnoSonTys06,WanHuiMayHai07,Son06,HutLin07}) or across a limited range of scales (e.g.~\\cite{CalCooMel07,Amendola:2007rr,ZhaLigBeaDod07,Jain:2007yk,AmiWagBla07}) exist in the literature.  Hu \\& Sawicki \\cite{HuSaw07b} recently introduced a PPF approach that spans all linear scales and includes an ansatz for non-linear phenomenology. The linear framework is based on enforcing the requirements of a metric theory, with small deviations from the Friedmann-Robertson-Walker metric, and covariant conservation of the stress energy tensor. This approach provides an excellent description of at least two specific models of modified gravity, the self-accelerating branch of the braneworld acceleration model (DGP \\cite{DvaGabPor00}) and the modified action $f(R)$ model \\cite{Caretal03,NojOdi03,Capozziello:2003tk}.  In this Paper, we extend the PPF approach to include relativistic matter components and spatial curvature.  This generalization yields a self-consistent means of evolving metric perturbations from initial conditions in the radiation dominated era through to the present. We also phrase the PPF description in a manner that is useful for adapting an Einstein-Boltzmann code (e.g.~\\cite{SelZal96,Lewetal00}) for modified gravity. Furthermore,  specific modified gravity models already require such components at low redshift, e.g. the self-accelerating branch of DGP requires spatial curvature to fit distance measures \\cite{SonSawHu06} and the normal branch requires dark energy to explain acceleration.   We use this generalization to study the impact of evolution in the metric on the cosmic microwave background (CMB) through the integrated Sachs-Wolfe (ISW) effect.  In the DGP and $f(R)$ models of cosmic acceleration it is well known that the ISW effect provides one of the most powerful and robust tests available  \\cite{LueScoSta04,SonSawHu06,SonPeiHu07,Schmidt:2007vj,CalCooMel07,Pogosian:2007sw}.  Since it arises from fluctuations near the horizon scale during the acceleration epoch, it  requires a fully covariant description such as the one developed here.    It also therefore tests gravity on the largest scales observable.  We begin in \\S \\ref{sec:formalism} with the multicomponent generalization of the PPF description.  This derivation is given in the comoving gauge but we describe in the Appendix the corresponding relations in the synchronous gauge. In \\S \\ref{sec:ISW} we explore the phenomenology of the ISW effect in the CMB and examine current constraints.   We discuss these results in \\S \\ref{sec:discussion}. ",
        "conclusions": "We have generalized the parametrized post-Friedmann (PPF) description of cosmic acceleration from modified gravity to include multiple relativistic matter species and spatial curvature.  This generalization facilitates the adaptation of Einstein-Boltzmann codes for modified gravity.  We have adapted a comoving gauge code to study constraints on  deviations from general relativity on the largest scales through the CMB.  The integrated Sachs-Wolfe (ISW) effect probes the evolution of the metric $\\Phi_{-}$ on scales near the horizon during the acceleration epoch.    Under the PPF parametrization, modifications to gravity can change the amplitude and shape of the ISW contributions to the low multipole CMB temperature power spectrum even if the expansion history is indistinguishable from that of a cosmological constant.   For example in some regions of parameter space, the ISW effect can be nearly eliminated at the quadrupole bringing the predicted ensemble average quadrupole nearer to its observed value on our sky.  We use the WMAP data to constrain such modifications and find that deviations in the metric ratio parameter $g$ are constrained at the level of $g_{\\rm eff}=-0.12 \\pm 0.27$  at $z \\sim 0.7$ and $k \\sim 10^{-3}$ Mpc$^{-1}$. Correspondingly, specific models such as the DGP  and some $f(R)$ models are constrained by the data at levels consistent with the findings of previous work (e.g. \\cite{SonSawHu06,SonPeiHu07}).  By phrasing these and other constraints in the model-independent PPF description, we gain a more general understanding of what aspects of general relativity are tested by the observations.  The modified gravity framework studied here should also enable future studies of other cosmological observables on large scales that are beyond the scope of this work such as the lensing of the CMB and faint galaxies as well as CMB-galaxy correlations.           \\appendix"
    },
    "0801/0801.2119_arXiv.txt": {
        "abstract": "The use of specific tracers of the dense molecular gas phase can help to explore the feedback of activity on the interstellar medium (ISM) in galaxies. This information is a key to any quantitative assessment of the efficiency of the star formation process in galaxies. We present the results of a survey devoted to probe the feedback of activity through the study of the excitation and chemistry of the dense molecular gas in a sample of local universe starbursts and active galactic nuclei (AGNs). Our sample includes also 17 luminous and ultraluminous infrared galaxies (LIRGs and ULIRGs). From the analysis of the LIRGs/ULIRGs subsample, published in Graci\\'a-Carpio \\etal\\ (\\cite{gr07}), we find the first clear observational evidence that the star formation efficiency of the dense gas, measured by the L$_{FIR}$/L$_{HCN}$  ratio, is significantly higher in LIRGs and ULIRGs than in normal galaxies. Mounting evidence of overabundant HCN in active environments would even reinforce the reported trend, pointing to a significant turn upward in the Kennicutt-Schmidt law around L$_{FIR}$=10$^{11}$L$_{\\odot}$. This result has major implications for the use of HCN as a tracer of the dense gas in local and high-redshift luminous infrared galaxies. ",
        "introduction": "Millimeter telescopes can provide a detailed picture of the distribution and kinematics of molecular gas in galaxy disks through extensive CO line mapping. However, the use of tracers more specific to the dense molecular gas phase, the one directly involved in the fueling of star formation and AGN episodes, is required to probe the feedback of activity on the ISM. Provided that high spatial resolution and high sensitivity requirements are met, the observation of complex molecular species can help to track down galaxy evolution. Extragalactic chemistry can be used to challenge current chemical models of molecular gas, as active galaxies can drive chemical complexity on much larger scales compared to our Galaxy. Furthermore, this type of studies are a key to constrain conversion factors, which are required to derive gas masses from line luminosities. This is mostly relevant, as the excitation and chemistry of some of the more routinely used dense gas tracers (e.g., HCN lines) can be heavily affected in active environments.  Different processes can shape the evolution of molecular gas along the typical evolutionary track of a starburst: large-scale shocks, strong UV-fields, cosmic-rays, and eventually X-rays, in the presence of an AGN. These actors are expected to be at play at different stages but also at different locations in the disk and in the disk-halo interface of galaxies showing activity. The use of interferometers makes possible to unveil a strong chemical differentiation in the molecular gas disks of active galaxies. The example of M82, a prototypical starburst, is paradigmatic in this respect. Virtually all of the large-scale SiO emission detected in the PdBI map of M~82, published by Garc\\'{\\i}a-Burillo \\etal\\ (\\cite{gb01}), traces the disk-halo interface of the galaxy where episodes of mass injection are building up the gaseous halo. In contrast, widespread emission of the formyl radical, HCO, mapped in M~82 with the PdBI, reveals the propagation of photo-dissociation region (PDR) chemistry inside the disk of this starburst (Garc\\'{\\i}a-Burillo \\etal\\ \\cite{gb02}). The scenario of a giant PDR in the disk of M~82 has received further support from the high abundances measured for molecular tracers specific to PDR such as CN, HOC$^+$, CO$^+$ as well as a set of small hydrocarbon chains (Fuente \\etal\\ \\cite{fu05,fu06}). ",
        "conclusions": ""
    },
    "0801/0801.0436_arXiv.txt": {
        "abstract": "In order to explain the slow rotation observed in a large fraction of accreting pre-main-sequence stars (CTTSs), we explore the role of stellar winds in torquing down the stars.  For this mechanism to be effective, the stellar winds need to have relatively high outflow rates, and thus would likely be powered by the accretion process itself.  Here, we use numerical magnetohydrodynamical simulations to compute detailed 2-dimensional (axisymmetric) stellar wind solutions, in order to determine the spin down torque on the star.  We discuss wind driving mechanisms and then adopt a Parker-like (thermal pressure driven) wind, modified by rotation, magnetic fields, and enhanced mass loss rate (relative to the sun).  We explore a range of parameters relevant for CTTSs, including variations in the stellar mass, radius, spin rate, surface magnetic field strength, the mass loss rate, and wind acceleration rate.  We also consider both dipole and quadrupole magnetic field geometries.  Our simulations indicate that the stellar wind torque is of sufficient magnitude to be important for spinning down a ``typical'' CTTS, for a mass loss rate of $\\sim 10^{-9} M_\\odot$ yr$^{-1}$.  The winds are wide-angle, self-collimated flows, as expected of magnetic rotator winds with moderately fast rotation.  The cases with quadrupolar field produce a much weaker torque than for a dipole with the same surface field strength, demonstrating that magnetic geometry plays a fundamental role in determining the torque.  Cases with varying wind acceleration rate show much smaller variations in the torque suggesting that the details of the wind driving are less important. We use our computed results to fit a semi-analytic formula for the effective Alfv\\'en radius in the wind, as well as the torque.  This allows for considerable predictive power, and is an improvement over existing approximations. ",
        "introduction": "\\label{sec_intro}    For more than half a century, the spin rates and the angular momentum evolution of stars have been topics of vigorous study.  We know that stellar winds are responsible for the spinning down of late-type (later than F2) main sequence stars \\citep[][]{parker58, schatzman62, kraft67, skumanich72, soderblom83, kawaler88, macgregorbrenner91, barnessofia96, bouvier3ea97}.  There is still progress to be made on main sequence star spins \\citep{barnes03}, but perhaps the largest open questions remain at the pre-main-sequence phase, which determines the ``initial conditions'' for the spin histories of stars.   By the time intermediate/low mass ($\\la 2 M_\\odot$) pre-main-sequence stars become optically visible (T Tauri stars; TTSs), they already have ages around $10^{5}$ -- $10^{6}$ yrs.  A large fraction of TTSs (called classical TTSs; CTTSs) are observed to actively accrete material from a disk at a rate within a wide range of $\\sim 10^{-8} M_\\odot$ yr$^{-1}$ \\citep[e.g.,][]{johnskrullgafford02}.  At this rate, the angular momentum accreted from the orbiting disk should spin up the stars to a substantial fraction of breakup speed in a short amount of time (comparable to their ages).  The fact that the stars are also still contracting \\citep[e.g.,][]{rebullea02}, and that they presumably were accreting at much higher rates before they became optically visible, further adds to expectation of fast rotation.  Large data sets for the spins of TTSs in star formation regions and clusters of different ages \\citep[see][for a compilation]{rebull3ea04} show that approximately half of the stars are rotating rapidly and do seem to spin up as expected as they approach zero-age main sequence \\citep{vogelkuhi81, bouvier3ea97, rebull3ea04, herbstea07}.  However, the surprise is that the other $\\sim$half of TTSs exhibit much slower rotation rates ($\\sim 10$\\% of breakup speed) at all ages.  Recent studies have shown a correlation between slow rotation and the presence of an a accretion disk \\citep[see especially][]{ciezabaliber07}, though this idea has been controversial in the past \\citep[e.g.,][]{stassunea99, herbstea00, stassunea01, herbstea02}.  This is still an open issue, but it is clear than an efficient angular momentum loss or regulation mechanism is operating for the slow rotators.  Although alternative ideas have been proposed since \\citep[][see \\citealp{mattpudritz07coolstars} for a history]{konigl91, shuea94}, \\citet{hartmannstauffer89} offered the first potential explanation for the slow rotators, namely that massive stellar winds may be responsible for carrying off substantial angular momentum \\citep[see also][]{toutpringle92}.  In \\citet[][hereafter Paper I]{mattpudritz05l}, we extended this idea to consider the effects of the magnetic interaction between the star and disk, and we used a 1-dimensional scaling from the solar wind angular momentum loss to estimate the torque for TTSs.  The scaling suggested that, for an observationally constrained dipole magnetic field strength of 200 G \\citep[e.g.,][]{johnskrullea99, bouvierea07, johnskrull07iau, johnskrull07, smirnovea03, yang3ea07}, it might indeed be possible for the stellar wind to extract enough angular momentum to explain the slow rotators.  For stellar winds to balance the accreted angular momentum, the wind outflow rate needs to be a substantial fraction of the accretion rate.  In Paper I, we suggested that this is possible, if a fraction of the energy liberated by the accretion process actually powers the stellar wind.  The pre-main-sequence phase is, in fact, marked by powerful outflows \\citep{reipurthbally01}.  In the most powerful sources, due to the large linear momenta of the outflows \\citep{koniglpudritz00}, the X-ray luminosities \\citep{decampli81}, and possible detection of rotation \\citep{bacciottiea02, andersonea03, coffeyea04, ferriera3ea06, coffeyea07}, it appears that most of the flow arises from the accretion disk, rather than the star.  It is not clear what fraction of the total outflow may actually originate from the star, and thus how powerful are the stellar winds compared to main-sequence phase winds or compared to their accretion rates.  There is some observational evidence for powerful stellar winds from CTTSs, as distinguished from inner disk winds.  In particular, \\citet{edwardsea03, edwardsea06} observed the He I 10830 \\AA\\ line in 39 CTTSs and saw several cases with a broad, deep, blue-shifted absorption, indicating outflow velocities of typically a few hundred, and up to $\\sim 400$ km s$^{-1}$.  They concluded that this feature is best interpreted as arising in an optically thick stellar wind \\citep[see also][]{dupreeea05}.  They also suggested the winds may be accretion-powered, since the wind signatures are most prevalent in the stars with highest accretion rates and absent in non-accreting systems.  Subsequent modeling of the He I 10830 \\AA\\ line by \\citet{kwan3ea07} indicates that approximately half of these CTTSs show evidence for a powerful stellar wind.  Furthermore, \\citet{kurosawa3ea06} modeled the H$\\alpha$ emission line in these systems and suggested that a stellar wind component could most naturally explain the profiles observed in $\\sim 7$\\% of the stars in a sample compiled by \\citet{reipurth3ea96}.   There already exists some theoretical work on stellar winds, specifically from pre-main-sequence stars, with a focus on the wind driving mechanism \\citep{decampli81, hartmann3ea82, hartmannea90} or the collimation of the winds \\citep{fendt3ea95, fendtcamenzind96}. These do not discuss the expected angular momentum outflow rates, however.  The works that do calculate stellar wind torques for pre-main-sequence stars \\citep[][Paper I]{hartmannmacgregor82, mestel84, hartmannstauffer89, toutpringle92, paatzcamenzind96} are either based on a 1-dimensional formulation and/or have made a priori simplifying assumptions regarding the stellar magnetic field structure, wind flow speed, and latitudinal dependence of the wind. Calculating the stellar wind torque reliably is a complex, multi-dimensional problem, and more work is needed to develop the stellar wind theory further.  The primary goal of this paper therefore, is to take the next major step in developing the accretion-powered stellar wind picture by rigourously computing the steady-state solutions of winds from spinning magnetized stars.  We carry out a parameter study to provide a range of possible solutions that are expected to charaterize accretion-powered stellar winds.  Where possible, we compare our results to analytic magnetohydrodynamic (MHD) stellar wind theory.  In a companion paper, we will use these solutions to compare the stellar wind torques and wind driving power with the torque and energy deposition expected to arise from the interaction of the star with its accretion disk.  In the following section (\\S \\ref{sec_theory}), we give a brief introduction to basic stellar wind theory.  This provides the motivation for using a numerical approach and sets the stage for comparing our numerical results with the analytic theory.  Section \\ref{sub_driving} contains a discussion of our adopted wind driving mechanism.  We describe our numerical method for obtaining solutions in section \\ref{sec_method}, and present the results in section \\ref{sec_solutions}.  Section \\ref{sub_rasim} contains a semi-analytic formulation for the torque and a comparison to previous theory. ",
        "conclusions": "\\label{sec_discussion}    Using 2D (axisymmetric) MHD simulations, we computed steady-state, stellar wind solutions for a parameter range appropriate for T Tauri stars.  We carried out a parameter study including variations of the stellar mass, radius, surface magnetic field strength, and rotation rate, as well as mass loss rate, wind acceleration rate, and two different magnetic geometries (dipole and quadrupole).  Our solutions enabled us to determine the angular momentum carried in the wind, and its dependence on many of the parameters of the system.  Our main conclusions can be summarized as follows:    \\begin{enumerate}  \\item{For fiducial parameters, the torque is of the same order ($\\sim 10^{36}$ erg) as estimated in Paper I.  Therefore, if the stellar winds of TTSs have similar parameters to those considered here, they should have a significant influence on the stellar spin.}  \\item{The stellar winds are in the regime of moderately fast magnetic rotator winds.  They produce jets, as well as a wide-angle flow \\citep[see, e.g.,][]{mattbalick04}, which should interact with, and be modified by, surrounding material \\citep[not included in our simulations; e.g.,][]{gardinerfrank03, shangea06}.}  \\item{The cases with quadrupole fields resulted in a torque that is much weaker than cases with a dipole field of the same surface field strength.  Specifically, we find that a 200 G dipole field exerts the same stellar wind torque upon a star as a 1--2 kG quadrupole.  This illustrates the very strong effect of magnetic geometry on the stellar wind torque.}  \\item{We ran cases where the mass loss rate and other parameters were fixed, but the thermal wind driving parameters were varied.  For large variations in the wind acceleration, the torque changed by less than a factor of 2.  This suggests that the details of the velocity profile are not of fundamental importance, and our solutions should be a reasonable approximation for winds with other wind driving mechanisms. However, it will still be important for future work to compare our torque results to stellar wind solutions that use alternative driving mechanisms.}  \\item{Our determination of the torque allowed us to calculate the Alfv\\'en radius (via eq.\\ \\ref{eqn_tw}), which is a fundamental quantity in MHD wind theory.  We compared our numerical solutions to previous analytic work and obtained a semi-analytic formulation for $r_{\\rm A} / R_* \\propto [B_*^2 R_*^2 / (\\dot M_{\\rm w} v_{\\rm esc})]^m$, with $m \\approx 0.22$ (eq.\\ \\ref{eqn_rasim}), that well-describes many of our simulations with dipole fields.  This formulation appears to be an improvement over existing work, and the exponent $m$ is significantly smaller than usually assumed.}  \\end{enumerate}    We will continue to develop the theory of accretion-powered stellar winds in forthcoming work.  In a companion paper (the third in our series), we compare the stellar wind torques computed here to the torques expected to arise from the interaction between the star and an accretion disk.  We find spin-equilibrium (net zero torque) solutions and test the suggestion of Paper~I.  In a later paper, we will use the stellar wind solutions of this work to compute emission properties of TTS coronal winds."
    },
    "0801/0801.3329_arXiv.txt": {
        "abstract": "We report the detections of the anti-correlated soft and hard X-rays, and the  time lags of $\\sim$  hecto-second from the neutron star low-mass X-ray binary  Cyg X-2, a well-known Z-type luminous source. Both the anti-correlation and the positive correlation were detected during the low-intensity states, while  only the latter showed up during high-intensity states. Comparing with the lower part of normal branch and flaring branch, more observations located on  the horizontal and the upper normal branches are accompanied with the anti-correlation, implying   the occurrence of the anti-correlation   under circumstance of a low mass accretion rate. So far  the anti-correlated hard lag of thousand-second timescale are only reported from the Galactic black hole candidates  in their hard states. Here we provide the first evidence that a similar feature can also establish in a neutron-star system like Cyg X-2. Finally, the possible origins of the observed time lags are discussed under the current LMXB models. ",
        "introduction": "The compact X-ray binary system hosts  either a neutron star (NS) or a black hole (BH) candidate. Both systems show the similarity in their dependences on the X-ray luminosity and spectral shape ( van der Klis 1994a, 1994b, 2000; Belloni et al. 2002). The spectra of X-ray binaries are generally dominated by soft component, probably from the accretion disk and/or the surface of neutron star, and hard component from thermal Comptonization of soft photons in a hot plasma close to the compact object (Nowak et al. 1999b; Church 2001; Di Salvo et al. 2006b). Low-Mass X-ray Binary (LMXB)  contains a companion  with the mass $\\lesssim$ 1 $M_{\\odot}$.  NS LMXBs are divided into Z-type  and Atoll-type according to the track shape in the color-color diagrams (CCDs), where the  different position on the track correlates with  the different timing behavior (Hasinger \\& van der Klis 1989; Hasinger 1990).   Cyg X-2 is a  persistent LMXB, where  the type-I X-ray bursts are detected, suggesting its compact star is a NS with a low magnetic field (e.g. Kahn \\& Grindlay 1984), and its mass is measured  as 1.78 $\\pm$ 0.23 $M_{\\odot}$ (Orosz \\& Kuulkers 1999). The binary system has an evolved, late-type companion V1341 Cyg, with an orbital period of $\\sim$ 9.8 days (Cowley et al. 1979; Casares et al. 1998). Cyg X-2 is classified as a Z source according to its track on the CCD (Hasinger \\& van der Klis 1989; Hasinger 1990; van der Klis 2000), where the so-called horizontal branch (HB) is located at the top, and the normal branch (NB) and the flaring branch (FB) at the bottom. Along with the evolution of the source, it  moves continuously on the Z-track, with each position related to the different mass accretion rate. The accretion rate reaches a minimum at the left end of the HB and  a maximum at the right end of the FB (Hasinger \\& van der Klis 1989; Hasinger 1990). In addition, the quasi-periodic oscillations (QPOs) of Cyg X-2 are connected with the position on CCD as well:   the frequency varies between $\\sim$ 15 and $\\sim$ 60 Hz on the HB and the upper part of the NB,  and the low frequency  about 5-7 Hz on the lower part of the NB (so-called NBOs) (Piraino et al. 2002; van der Klis 2000, 2006).  Although Z sources are the well-known  LMXBs, the origin and evolution of the observed hard tails ($\\gtrsim$ 30 keV) are still remaining unclear. The hard power-law tails have been detected in the Z sources of e.g. Sco X-1 (Manchanda 2006), GX 5-1 (Asai et al. 1994), GX 17+2 (Di Salvo et al. 2000),  GX 349+2 (Di Salvo et al. 2001), Cyg X-2 (Frontera et al. 1998), GX 340+0 (Lavagetto et al. 2004), and Cir X-1 (Iaria et al. 2001; Ding et al. 2003). Except for Sco X-1 whose hard tail does not show any clear correlation with the position of CCD as detected by the Rossi X-Ray Timing Explorer ({\\it RXTE}) (D'Amico et al. 2001; Manchanda 2006), and could be correlated with the periods of radio flaring (Strickman \\& Barret 2000). for most of the others, the hard component appears to become weaker at the higher accretion rates. However, for Sco X-1, Di Salvo et al. (2006a) find that the flux of the power-law component slightly decreases when the source moves in the CCD from HB to NB, and becomes not significantly detected in FB with {\\it INTEGRAL} and {\\it RXTE}. The jets  are also observed in the low magnetic field NS systems (Fender \\& Kuulkers 2001). For Cyg X-2, the hard tail is detected only in the HB spectra, which is consistent  with the fact that the radio flux is observed to have a maximum in HB (Hasinger et al. 1990; Di Salvo et al. 2002). Therefore, both the hard tail and radio emission could be produced by a jet (Penninx et al. 1988; Hasinger et al. 1990; Di Salvo et al. 2002). Such an accretion-ejection phenomena (Das \\& Chakrabarti 1999; Merier 2001) for the low magnetic NS systems is well addressed in the model of advection-dominated accretion flow by Narayan \\& Yi (1994).      Recently, an anti-correlated time lag ($\\lesssim 2000 s$) of the hard X-rays (20-50 keV) with respect to the soft X-rays (2-7 keV) is found in the hard state of some black hole X-ray binaries (BHXBs) (Choudhury \\& Rao 2004; Choudhury et al. 2005; Sriram et al. 2007). The timescale of the delay could be ascribed to the viscous timescale of matter flow in the optically thick accretion disk (Choudhury \\& Rao 2004). In addition, the outflow as a jet appears in the low-hard state of all BH sources and the truncated accretion disk also appears in low-hard state, which implies that the accretion disk emission is connected with the jet emission (Fender 2001; Choudhury et al. 2005). The purpose of this paper is to investigate whether the anti-correlation between the soft  (2-5 keV)  and hard (16-30 keV) X-rays exists as well in the NS LMXB system of Cyg X-2. The paper is  organized as the follows. In section 2, the observations and data analysis of  {\\it RXTE} are introduced. In section 3, the results obtained on the  QPO and  the cross-correlation between the soft and hard X-rays are presented. Finally, in section 4 are  the discussions and  summary. ",
        "conclusions": "We have analyzed all the  data available from {\\it RXTE} observations on the NS Cyg X-2, and detected the anti-correlation between the soft (2-5 keV) and hard X-rays (16-30 keV). We find that, during the low-intensity state of Cyg X-2, there exists different observations showing separately either the anti-correlation or the positive correlation. These  observations are located on the NB and FB in the CCD. But for the majority of the observations, the fraction of the observations with anti-correlation in HB and upper part of NB are larger than that in low part of NB and FB. During the high-intensity state only positive correlation is detected. Specially, during the FB of high-intensity state, the intensities decrease, but the colors remain the same.   The pivoting of the wide-band X-ray spectrum is also detected in all the observation with anti-correlation.   The mass accretion rate of individual Z source increases along HB, NB and FB (van der Klis 2000; 2006). We discuss only the result in Figure 6 where most observations are included. The anti-correlation between the soft and hard X-rays is mostly found on the HB and upper parts of NB, which suggest that the anti-correlation could be associated with the low mass accretion rate. In the low-hard state of BHXBs, an anti-correlation of the soft (2-7 keV) and hard X-rays (20-50 keV) is detected, which suggests the existence of a truncated accretion disk (Choudhury et al. 2005). At low accretion rates, the disk is truncated far away from the compact object and the X-ray spectrum is dominated by a thermal Compton spectrum. With the increase of the accretion rate, the location of the disk truncation moves in, resulting in the Comptonizing cloud to decrease as well as the amount of seed thermal photons to increase, which cools the Comptonizing cloud more efficiently, giving rise to the anti-correlation of hard and soft X-rays and the pivoting of the wide-band X-ray spectrum (Choudhury et al. 2005).   The timescale of this anti-correlated hard X-ray time lag is about  thousand seconds in BHXBs, which is explained to be attributed to the viscous timescale of the matter flow in the radiation pressure dominated optically thick accretion disk (Choudhury \\& Rao 2004), i.e., the influence by the matter flow propagates from the outer disk (soft photon region) to the innermost disk (hard photon region). Our results show that, for NS LMXB Cyg X-2, the anti-correlated hard X-ray time lag is about several hundred seconds. It is considered that the hard X-ray emission could come from the Comptonizing of the soft seed photons for X-ray binary system. For BHXBs, the soft seed photons are produced from the accretion disk. However, for NS systems, the soft seed photons are produced from the surface of NS and/or the accretion disk, and the Comptonizing cloud might be a hot corona, a hot flared inner disk, or even the boundary layer between the NS and accretion disk if it is sufficiently thin and hot (also see Popham \\& Sunyaev 2001). If the soft seed photon comes from the NS surface, the hard photons come from the inverse Compton scattering of soft photons by hot electrons, and a hard lag of this Compton process is expected with timescale of about  $<$ 1s (Hasinger 1987; van der Klis et al. 1987; Nowak et al. 1999a). If the soft X-rays come from the cool and dense blob, with the blob spiraling inward, it emits the soft X-ray radiation that is Comptonized in an hot and dense corona, which results in Comptonized X-ray spectrum to be harder with time, causing the hard time lags of timescale $\\sim$ 1 s (B{\\\"o}ttcher \\& Liang 1999). Therefore, the above-mentioned models should be ruled out as the possible mechanisms for the several hundred-second time lags, and our results suggest that the hecto-second  time lag should be produced by other mechanisms.    The timescale of the anti-correlated hard X-ray time lag is comparative to  the viscous timescale of the compact systems, therefore, Choudhury and Rao (2004) ascribe the time lag of BH system to the viscous timescale of the matter flow in the radiation pressure dominated optically thick accretion disk. Assuming that the soft seed photons of Cyg X-2 come from the accretion disk, the viscous time $t_{vis}$ in the disk can be calculated by $$t_{vis} = 30 \\alpha^{-1} M^{-1/2} R^{7/2} \\dot{M}^{-2} s$$ where $\\alpha$ is the viscosity parameter in units of 0.01, M is the mass of the compact object in solar mass, R is the radial location in the accretion disk in units of 10$^7$ cm, $\\dot{M}$ is the mass accretion rate in units of 10$^{18}$ g/s (see Fender \\& Belloni 2004). For Cyg X-2, taking $\\alpha$ = 1, M = 1.78 and $\\dot{M}$ = 1.3 ($\\sim$ 2$\\times$$10^{-8}$ M$_\\odot$ $yr^{-1}$, see Smale 1998), we get R $\\sim$ 1.8 for a viscous timescale of $\\sim$ 100 s. Thus, the observed delay will infer  the location of the truncation of $\\sim$ 18 NS radius. Choudhury \\& Rao (2004) obtain the location of the disk truncation of $\\sim$ 25 Schwarzschild radius for the BHXB Cyg X-3. Table 1 shows the locations of the truncation of every ObsID with anti-correlated hard X-rays time lag, by assuming  $\\alpha$ = 1 and $\\dot{M}$ = 1.3. The results of Table 1 show the locations of the truncation are not evolutionary  along the branches. However, unlikely the BH binaries, for the Z sources, the X-ray intensity may not be proportional to $\\dot{M}$, or $\\dot{M}$ can not be estimated accurately. In addition, the viscosity parameter $\\alpha$ could be different for the different position of Z-track. Therefore, we can not obtain the locations of the truncation accurately, maybe, different  from the results of Table 1, the locations of the truncations are evolutionary  along Z-track.    The above model can explain the anti-correlated hard X-ray lag, however, it can not explain the anti-correlated soft X-ray lag. The anti-correlated soft X-ray lag ($< 200 s$) is also detected in Cyg X-2, whose timescale is comparative to  the viscous timescale. If a fluctuation produced from the innermost accretion disk could propagate to the outer accretion disk and modulates the soft X-ray emission regions, the soft lag is about to be explained. It is possible to produce   such a fluctuation, since the model by Li et al. (2007) provides a probability for it, which  ascribes the accretion disk to be thermally unstable and to exhibit the limit-cycle behavior (see also e.g. Kato et al. 1998). In the cycle process, the expansion wave in the innermost disk can be formed and moves outward with time, which perturbs the outer material, modulating the soft X-ray emission, and resulting in the soft X-ray lag, and the timescale is several hundred seconds.     The anti-correlation in BH binary systems implies that the geometric structure of the disk-jet could be connected with a truncated disk (Choudhury \\& Rao 2004; Choudhury et al. 2005). The truncated accretion disk could occur in the low-hard state where the outflow in the form of a jet is present. In NS LMXB Cyg X-2, a hard power-low tail is significantly detected only in the HB spectra (Di Salvo et al. 2002). The strong non-thermal radio flares are observed in the HB and upper NB, and the source is radio quiet in the lower NB and FB (Hasinger et al. 1990). Therefore, the hard power-law component could be related with the radio flares, and the hard power-law component is related to the high-velocity electrons (probably from a jet) of producing the radio emission (Di Salvo et al. 2002). Our results show that in Cyg X-2 the anti-correlations mostly present in HB and upper NB, at where the hard power-law component and radio flares also occur."
    },
    "0801/0801.4023_arXiv.txt": {
        "abstract": "An evaporating black hole in the presence of an extra spatial dimension would undergo an explosive phase of evaporation. We show that such an event, involving a primordial black hole, can produce a detectable, distinguishable electromagnetic pulse, signaling the existence of an extra dimension of size $L\\sim10^{-18}-10^{-20}$~m. We derive a generic relationship between the Lorentz factor of a pulse-producing ``fireball\" and the TeV energy scale. For an ordinary toroidally compactified extra dimension, transient radio-pulse searches probe the electroweak energy scale ($\\sim$0.1~TeV), enabling comparison with the Large Hadron Collider. ",
        "introduction": "A new generation of radio telescopes will search for transient pulses from the universe \\cite{LWA,MWA,LOFAR,ETA1,ETA2}. Such searches, using pre-existing data, have recently found surprising pulses of galactic and extragalactic origin \\cite{McLaughlin,Lorimer,Vachaspati}. While the results will be of obvious astrophysical importance, they could also answer basic questions in physics which are difficult to address. In particular, as we will discuss here, searches for transient pulses from exploding primordial black holes (PBHs) can yield evidence of the existence of an extra spatial dimension, and explore electroweak-scale physics.\\footnote{The existence of primordial black holes is an open question. However, there exist models of the early universe which produce large numbers of primordial black holes and are consistent with all current observational data (see, for example \\cite{kaw}).} The potential impact could be timely and cut across many areas of investigation. For example, the Large Hadron Collider (LHC) is poised to investigate electroweak-scale physics, and may also yield evidence of the existence of extra spatial dimensions. Also, intensive work on the unification of quantum mechanics and gravitation has yielded insightful theoretical advances, often requiring extra spatial dimensions \\cite{GSW}, yet there is little experimental observation which gives feedback on this proposed phenomenon. Furthermore, mapping of the anisotropies in the cosmic microwave background radiation has enabled ``precision cosmology,'' yet searches for PBHs, which would explore smaller scale primordial irregularities (a source of PBHs), would be valuable \\cite{Carr}. Searches for transient pulses from exploding primordial black holes can provide information impacting all of these areas of investigation, which at first glance appear unrelated, but are intimately connected.  The defining relation governing the Hawking evaporation of a black hole \\cite{Hawking} is \\be T = \\frac{\\hbar c^3}{8\\pi G k} \\frac{1}{M}, \\label{tem} \\ee for mass $M$ and temperature $T$. The power emitted by the black hole is \\be P \\propto \\frac{\\alpha(T)}{M^2}, \\label{pow} \\ee where $\\alpha(T)$ is the number of particle modes available. Equations~(\\ref{tem}) and (\\ref{pow}), along with an increase in the number of particle modes available at high temperature, leads to the possibility of an explosive outburst as the black hole evaporates its remaining mass in an emission of radiation and particles.\\footnote{The behavior of the evaporation process, as the Planck mass is reached, is not certain \\cite{Bee3}. However, the details of this late stage of evaporation will not alter the analysis presented here.} PBHs of sufficiently low mass would be reaching this late stage now \\cite{Carr}. Searches for these explosive outbursts have traditionally focused on $\\gamma$-ray detection \\cite{Halzen}. However, Rees noted that exploding primordial black holes could provide an observable coherent radio pulse that would be easier to detect \\cite{Rees}.  Rees \\cite{Rees} and Blandford \\cite{Blandford} describe the production of a coherent electromagnetic pulse by an explosive event in which the entire mass of the black hole is emitted. If significant numbers of electron-positron pairs are produced in the event, the relativistically expanding shell of these particles (a ``fireball'' of Lorentz factor $\\gamma_f$) acts as a perfect conductor, reflecting and boosting the virtual photons of the interstellar magnetic field. An electromagnetic pulse results only for $\\gamma_f \\sim 10^5$ to $10^7$, for typical interstellar magnetic flux densities and free electron densities. Below $\\gamma_f \\sim 10^5$ the energy emitted by the PBH goes primarily into sweeping up the ambient interstellar plasma, and not into an electromagnetic pulse; above $\\gamma_f \\sim 10^7$ the number of electron-positron pairs is insufficient to carry the fireball surface current necessary to expel the interstellar magnetic flux density. The energy of the electron-positron pairs is \\begin{equation} \\label{gen1} kT \\approx \\frac{\\gamma_f}{10^5} \\ 0.1 \\ {\\rm TeV}. \\end{equation} Thus the energy associated with $\\gamma_f\\sim10^5$ corresponds roughly to the electroweak scale. ",
        "conclusions": ""
    },
    "0801/0801.4679_arXiv.txt": {
        "abstract": "{A bright X-ray transient was seen during an XMM-Newton observation in the direction of the Small Magellanic Cloud (SMC) in October 2006.} {The EPIC data allow us to accurately locate the source and to investigate its temporal and spectral behaviour.} {X-ray spectra covering 0.2$-$10 keV and pulse profiles in different energy bands were extracted from the EPIC data.} {The detection of 6.85 s pulsations in the EPIC-PN data unambiguously identifies the transient with \\xtej, discovered as 6.85 s pulsar by RXTE. The X-ray light curve during the XMM-Newton observation shows flaring activity of the source with intensity changes by a factor of two within 10 minutes. Modelling of pulse-phase averaged spectra with a simple absorbed power-law indicates systematic residuals which can be accounted for by a second emission component. For models implying blackbody emission, thermal plasma emission or emission from the accretion disk (disk-blackbody), the latter yields physically sensible parameters. The photon index of the power-law of $\\sim$0.4 indicates a relatively hard spectrum. The 0.2$-$10 keV luminosity was 2\\ergs{37} with a contribution of $\\sim$3\\% from the disk-blackbody component. A likely origin for the excess emission is reprocessing of hard X-rays from the neutron star by optically thick material near the inner edge of an accretion disk. From a timing analysis we determine the pulse period to 6.85401(1) s indicating an average spin-down of $\\sim$0.0017 s per year since the discovery of \\xtej\\ in May 2003.} {The X-ray properties and the identification with a Be star confirm \\xtej\\ as Be/X-ray binary transient in the SMC.} ",
        "introduction": "The transient pulsar \\xtej\\ with a period of 6.8482$\\pm$0.0007 s was discovered during RXTE observations of the Small Magellanic Cloud (SMC). The pulsations were detected on 2003, Apr. 29, May 7, May 15 and May 19 but not on Apr. 24 and May 28 \\citep{2003ATel..163....1C}, suggesting an X-ray outburst which lasted about three to four weeks. This behaviour is characteristic for Be/X-ray binaries undergoing a type II outburst which is caused by the ejection of matter from the Be star and enhanced accretion onto a compact object, in most cases a neutron star \\citep{1998A&A...338..505N}. Shorter outbursts (type I) with a duration of a few days are often separated by the orbital period of the binary system. They generally occur close to the time of periastron passage of the neutron star when the neutron star approaches the circumstellar disc of the Be star \\citep[see e.g.][]{2001A&A...377..161O}.  Be/X-ray binaries form the major class of High Mass X-ray Binaries (HMXBs). In the remaining systems - the supergiant HMXBs - the compact object accretes matter from the fast stellar wind of an early type O or B supergiant star. In the SMC more than 60 Be-HMXBs are known, while only one HMXB is established as supergiant system (SMC\\,X-1). Recent reviews of the optical and X-ray properties of these systems can be found in \\citet{2005MNRAS.356..502C} and \\citet{2004A&A...414..667H}, respectively.  The X-ray transient \\xtej\\ is one of several Be/X-ray binary candidates in the SMC which were discovered through their pulsations with RXTE. This non-imaging instrument is able to monitor the activity of pulsars in the SMC \\citep{2005ApJS..161...96L} but provides only a coarse determination of their sky coordinates. The latter hampers the optical identification and final confirmation of the HMXB nature of the pulsars. The pulsar \\xtej\\ was detected in X-ray outburst during an \\xmm\\ observation of RX\\,J0103.8$-$7254, a candidate super-soft X-ray source in the direction of the SMC discovered by ROSAT. The precise localization in the EPIC images allowed us to identify the optical counterpart of \\xtej\\ which shows optical brightness and colours consistent with a Be star \\citep[see][]{2007ATel.1095....1H}. Analysis of MACHO and OGLE data of the counterpart reveals long-term variations in the optical light curve \\citep{2007ATel.1181....1S} on time scales of 620 to 660 days which are not periodic, suggesting that they are related to the Be phenomenon and not to the binary orbit \\citep{2007arXiv0711.4010M}.  Here we present a temporal and spectral analysis of the EPIC data of \\xtej\\ obtained in October 2006. ",
        "conclusions": "After its discovery with RXTE in April 2003, the 6.85 s pulsar \\xtej\\ was detected in X-rays for the first time with imaging instruments which allowed us to better localize the source and identify the optical counterpart. During the \\xmm\\ observation in October 2006 the source was seen in outburst with a luminosity of $\\sim$2\\ergs{37} in the 0.2$-$10.0 keV band. During the observation the source strongly varied in intensity with changes by a factor of two within ten minutes. Pulse profiles obtained by folding the data from the EPIC-PN detector in different energy bands with 30 phase bins (corresponding to 228.5 ms per phase bin) are highly structured. The fast intensity variations indicate that the bulk of the emission originates close to the neutron star surface and suggests a complicated emission geometry with narrow beams contributing to the X-ray emission at various neutron star spin phases and with energy dependent amplitudes.  Deviations from a power-law model in the phase averaged EPIC spectra suggest an additional emission component mainly contributing at energies below $\\sim$1.5~keV. This component is harder than the soft component seen in other Be-HMXBs in the SMC, which can be modeled by thermal plasma emission with kT$\\sim$0.15 keV \\citep[see e.g. RX\\,J0103.6$-$7201;][]{2005A&A...438..211H}. On the other hand the excess emission in the spectra of \\xtej\\ is softer than the hot blackbody component seen in \\xmm\\ spectra of two persistent Be-HMXBs in the Milky Way \\citep{2006A&A...455..283L,2007A&A...474..137L} which is interpreted as emission from hot polar caps of the neutron star. When modelled as blackbody emission, the large inferred emission area (with radius $\\sim$30 km) for \\xtej\\ is incompatible with an origin on the neutron star surface.  In their work on the origin of the soft excess in X-ray pulsars, \\citet{2004ApJ...614..881H} conclude that for luminous X-ray pulsars ($\\gsim$\\oergs{38}) the soft excess can only be explained by reprocessing of hard X-rays from the neutron star by optically thick material. At intermediate luminosity ($\\sim$\\oergs{37}) also other processes such as emission from photo-ionized or collisionally heated gas can contribute. The spectra of \\xtej\\ do not show significant line emission suggesting that in this source the excess emission is mainly caused by reprocessing in optically thick material, most likely located near the inner edge of an accretion disk. The relatively low contribution ($\\sim$3\\%) to the total luminosity and the estimate for the inner disk radius of $\\sim$19~km inferred from the disk-blackbody model are consistent with such a model.  The luminosity of the soft component seen from RX\\,J0103.6$-$7201 is strongly correlated with the total source luminosity over at least a factor of ten variation in source intensity. At a maximum of $\\sim$6.4\\ergs{36} RX\\,J0103.6$-$7201 is still a factor of $\\sim$3 below the luminosity of \\xtej\\ during its October 2006 outburst. This is consistent with the conclusions of \\citet{2004ApJ...614..881H}, that for lower luminosities emission from diffuse, optically thin gas dominates. Which of the emission components becomes visible in the X-ray spectra of Be-HMXBs certainly depends not only on source luminosity but also on geometric effects."
    },
    "0801/0801.4486_arXiv.txt": {
        "abstract": "We have studied the non-resonant streaming instability of charged energetic particles moving through a background plasma, discovered by Bell \\cite{bell04}. We confirm his numerical results regarding a significant magnetic field amplification in the system. A detailed physical picture of the instability development and of the magnetic field  evolution is given. ",
        "introduction": "The diffusive shock acceleration process (Krymsky \\cite{krymsky77}; Axford et al. \\cite{axford77}; Bell \\cite{bell78}; Blandford and Ostriker \\cite{blandford78}) is considered as the principal mechanism for the production of the galactic cosmic rays in supernova remnants (SNRs). A great strength of the random magnetic fields that provide the scattering of energetic particles both upstream and downstream of a supernova shock is necessary for efficient acceleration. It was originally suggested that this may be the result of a gyroresonant streaming instability that develops due to the presence of a diffusive streaming of accelerated particles (Bell \\cite{bell78}, Blandford \\& Ostriker \\cite{blandford78}).  The corresponding magnetohydrodynamic (MHD) waves have wavelengths of the order of the gyroradii of the energetic particles. This kind of resonant cyclotron instability had been suggested earlier to regulate the propagation of the galactic cosmic rays (Lerche \\cite{lerche67}, Wentzel \\cite{wentzel74}). If the generated random fields are amplified up to wave amplitudes that correspond to the strength of the mean magnetic field, the scattering mean free path of resonant energetic particles will decrease to a value comparable with their gyroradius. This regime of diffusion is called Bohm diffusion and is often used to estimate the maximum energy of the accelerated particles.  However, even this rather optimistic regime of diffusion may ensure the acceleration of cosmic ray protons only up to energies of about $10^{14}$ eV (Lagage \\& Cesarsky \\cite{lagage83}, Berezhko \\cite{ber96}) if the magnetic field in the remnant is comparable with the interstellar magnetic field, which is typically less than $10$ $\\mu $G. Significant magnetic field amplification is necessary for acceleration up to higher energies. Amplification might be possible since quasilinear theory of the resonant streaming instability allows it (see e.g. McKenzie \\& V\\\"olk \\cite{mckenzie82}). Since this perturbation theory breaks down in the case of resonant interaction of particles with high amplitude MHD waves, Lucek and Bell \\cite{lucek00} performed MHD simulations combined with calculations of energetic particle trajectories and found that the magnetic field can be amplified to a level where the random field exceeds the initial mean field. The necessity to perform detailed trajectory calculations leads to strong numerical limitations of such simulations.  Different qualitative treatments of magnetic field amplification at supernova shocks were suggested on an analytical level (see e.g. Bell \\& Lucek \\cite{bell01}, Ptuskin \\& Zirakashvili \\cite{ptuskin03}, Vladimirov et al. \\cite{vladimirov06}, Amato \\& Blasi \\cite{amato06}). Using the observed synchrotron emission, field amplification was included on a phenomenological level in numerical solutions of the coupled gas dynamic and particle acceleration equations (e.g. Berezhko et al. \\cite{berezhko02}, Berezhko \\& V\\\"olk \\cite{berezhko04a, berezhko04b}, V\\\"olk et al. \\cite{voelk07}). It became clear that such amplification can lead to particle accelaration to knee energies at $\\approx 3\\cdot 10^{15}$ eV and  -- according to speculative extrapolations -- possibly even beyond.  On the basis of the dispersion relation for collisionless MHD waves derived by Achterberg \\cite{achterberg83}, Bell \\cite{bell04} found a non-resonant streaming instability that had been overlooked before. He argued that the diffusive streaming of particles accelerated at supernova shocks may be so strong that it modifies the dispersion relation of MHD waves in an essential way. Then a non-oscillatory purely growing MHD mode appears at scales smaller than the gyroradius of the particles which excite the instability.  The growth rate of this mode can be larger than the growth rate of the resonant mode. Since the particle trajectories are only weakly deflected by small-scale magnetic inhomogeneities, one can avoid complicated trajectory calculations. Only the mean cosmic-ray flux needs to be specified. Bell \\cite {bell04} performed corresponding MHD simulations and showed that the magnetic field can indeed be significantly amplified.  In the following we investigate this non-resonant instability in considerably more detail. We have performed MHD simulations of the instability with very good  numerical resolution and have found a simplified analytical description for the magnetic field amplification. Our results may be applied to different astrophysical objects, where energetic particles exist.  The present work deals with the instability in the presence of a non-resonant energetic particle population streaming through the thermal plasma. And it simulates the nonlinear instability development.  In a companion paper (Zirakashvili \\& Ptuskin \\cite{zirakashvili08}, Paper II) this modeling will be combined with an analytical treatment of the diffusive acceleration of particles in a plane, steady shock with its intrinsically non-uniform energetic particle distribution in the precursor.  The paper is organized as follows: The basic equations are given in the next section. The scattering of cosmic ray particles by small-scale random magnetic fields is considered in Sect.3. The cosmic ray electric current is calculated in Sect. 4.  The non-resonant instability and the MHD simulations are described in Sect.5 and 6. The summary is given in the last Section. ",
        "conclusions": "We have  modeled the non-resonant instability produced by a flux of charged energetic particles that is driven through a scattering thermal plasma. Using a significantly better numerical resolution we basically confirm the results obtained earlier by Bell \\cite{bell04}. The magnetic field may be amplified significantly. The unstable magnetic spirals collide with each other and shocks of moderate strength are formed  as the instability develops (see previous Section). These shocks lead to significant gas heating. Since free expansion of the magnetic spirals after collision is impossible, the field grows only linearly in time at later epochs (see Eq. (B3)). If the system has enough time to evolve, the magnetic field growth will be stopped when the gyroradius $pc/qB$ of the energetic particles in the amplified field $B$ will drop down to the scale of the amplified field $k^{-1}\\sim cB/4\\pi j_\\mathrm{d}$. This determines the value of the saturated magnetic field (Bell \\cite{bell04}, Pelletier et al. \\cite{pelletier06}):  \\begin{equation} \\frac {B^2}{4\\pi }\\sim \\frac {u_\\mathrm{cr}}{v}\\epsilon _\\mathrm{cr}. \\end{equation}  We should note that the small-scale approximation considered in Sect. 3 becomes invalid when the magnetic field reaches this saturation value.  Since the instability is driven by the Lorentz force, the corresponding MHD turbulence has specific properties. It has non-zero magnetic helicity and a non-zero mean electric field ${\\bf E}_0$ parallel to the mean magnetic field.  We expect that the formation of multiple shocks in three-dimensional MHD turbulence will also occur for other instabilities, in particular for the resonant streaming instability, driven by cosmic rays.  The scattering of energetic particles by the small-scale magnetic inhomogeneities can be described using the Dolginov-Toptygin approximation (Dolginov \\& Toptygin \\cite{dolginov67}). The appearance of a mean second order electric field ${\\bf E}_0=-c^{-1}\\left< \\bf \\delta u\\times \\delta B\\right> $, which is oppositely directed to the electric current of the energetic particles, modifies the cosmic ray transport equation (see Appendix A).  This non-resonant streaming instability may be important in any astrophysical site where a strong electric current of energetic particles exists and where the initial magnetic strength is small enough (see condition (18)). Supernova remnants, starburst galaxies, galaxy cluster accretion shocks and AGN jets are  possible candidates for an application of this instability.  The  results  obtained will be used in a companion paper by Zirakashvili \\& Ptuskin \\cite{zirakashvili08} (Paper II) for a  model of diffusive shock acceleration in young SNRs in the presence of the non-resonant streaming instability."
    },
    "0801/0801.0166_arXiv.txt": {
        "abstract": "We present chemical abundance measurements from medium resolution observations of 8 sub-damped Lyman-$\\alpha$ absorber and 2 strong Lyman-limit systems at z $\\la$ 1.5 observed with the MIKE spectrograph on the 6.5m Magellan II Clay telescope. These observations were taken as part of an ongoing project to determine abundances in \\za $\\la$ 1.5 quasar absorption line systems (QSOALS) focusing on sub-DLA systems. These observations increase the sample of Zn measurements in \\za $\\la$ 1.5 sub-DLAs by $\\sim$50$\\%$. Lines of Mg I, Mg II, Al II, Al III, Ca II, Mn II, Fe II, and Zn II were detected and column densities were determined. Zn II, a relatively undepleted element and tracer of the gas phase metallicity is detected in two of these systems, with [Zn/H]$=-$0.05$\\pm$0.12 and [Zn/H]$>$+0.86. The latter system is however a weak system with \\nhI$<$18.8, and therefore may need significant ionisation corrections to the abundances. Fe II lines were detected in all systems, with an average Fe abundance of $<$[Fe/H]$>$$=-$0.68, higher than typical Fe abundances for DLA systems at these redshifts. This high mean [Fe/H] could be due to less depletion of Fe onto dust grains, or to higher abundances in these systems. We also discuss the relative abundances in these absorbers. The systems with high metallicity  show high ratios of [Mn/Fe] and [Zn/Fe], as seen previously in another sub-DLA. These higher values of [Mn/Fe] could be a result of heavy depletion of Fe onto grains, unmixed gas, or an intrinsically non-solar abundance pattern. Based on Cloudy modeling, we do not expect ionisation effects to cause this phenomenon. ",
        "introduction": "There are still many open questions concerning the processes of galaxy formation and evolution. Measurements of the abundances of heavy elements in galaxies give important information about the ongoing processes of star formation and death and the overall chemical enrichment of these galaxies. Studying galaxies at higher redshifts through emission is often difficult and time consuming, and biases towards more luminous galaxies are often present. Quasar absorption line systems (QSOALS) provide a means of studying the interstellar medium (ISM) of high redshift galaxies independent of the  morphology of galaxies (that is, not based on the selection of galaxies of certain morphology). In addition, absorption lines in QSO spectra allow us to study the diffuse intergalactic medium using lines of O VI (e.g., \\citealt{Sim02}), or using X-Ray absorption lines of more highly ionized species \\citep{Fang02}.  Quasar absorption line systems with strong Lyman-$\\alpha$ lines are often divided into two classes: Damped Lyman-$\\alpha$ (DLAs, log N$_{\\rm H \\ I}$ $\\ge$ 20.3) and sub-Damped Lyman-$\\alpha$ (sub-DLA 19 $\\la$ log N$_{\\rm H \\ I}$ $<$ 20.3, \\citealt{Per01}) are expected to contain a major fraction of the neutral gas in the Universe, while the majority of the baryons are thought to lie in the highly ionized and diffuse Lyman-$\\alpha$ forest clouds with log N$_{\\rm H \\ I}$ $\\la$ 14 in intergalactic space \\citep{Petit93}. The DLA and sub-DLA systems are generally believed to be associated directly with galaxies, at all redshifts in which they are seen. Observations of the galaxies in emission responsible for the absorption line systems have been met with mixed results. In several instances, a galaxy at the same redshift in emission could not be found near the line of sight of the QSO (see for example \\citealt{Kul00, Kul01, Chen03, Rao03}.)  DLAs have been the preferred systems for chemical abundance studies thanks to their high gas content. However, most DLAs are found to be metal poor, typically far below the solar level (e.g., \\citet{Kul05} and references therein). Based on Fe II abundance measurements, \\citet{Per03a} suggested stronger metallicity evolution in sub-DLA systems than DLA systems. Specifically, \\citet{Per03a} showed that the \\nhI-weighted mean Fe metallicity in sub-DLAs increased from $\\sim$ 1/100 Z$_{\\sun}$ at z$\\sim$4.5 to $\\sim$ 1/3 Z$_{\\sun}$ at z$\\sim$0.5. This has been validated with Zn, the more reliable metallicity inidicator by \\citet{Kul07}.  To date, few observations have been made of $z < 1.5$ sub-DLAs due to the lack of spectrographs with enough sensitivity at short wavelengths and the the paucity of known sub-DLAs in this redshift range. Redshifts $z < 1.5$ span ~70$\\%$ of the age of the Universe (assuming a concordance cosmology of $\\Omega_{m}=0.3, \\Omega_{\\Lambda}=0.7$). This redshift range is clearly important for understanding the nature of sub-DLA systems and galactic chemical evolution as well.  Recently \\citet{Per03a, Per06a, York06, Kul07, Pro06, Kh07, Mei07} have suggested that sub-DLAs may be more metal rich than DLAs, and may contribute significantly to the cosmic metal budget. Several sub-DLA systems have actually been seen with super-solar abundances \\citep{Kh04, Per06a, Pro06, Mei07}. Super-solar abundances have also been reported in strong Lyman-$\\alpha$ forest lines with log \\nhI$<$16, although ionisation correction factors are important and not easily constrained \\citep{Char03, Ding03, Mas05, Sch07}.  As the DLA systems dominate the \\nhI weighted mean metallicity, previous studies have been dominated by DLAs. At least an equally large sample of sub-DLAs is needed to determine the overall metallicity evolution of QSO absorbers. The H I column density distribution for QSO absorbers shows that systems with 19 $\\la$ log N$_{\\rm H \\ I} \\la$ 20.3 are $\\sim$8 times more numerous than the higher column density classical damped Ly$\\alpha$ systems, making sub-DLAs more readily available probes of neutral gas in the distant Universe.  Many elements have been detected in QSO absorber systems, including C, N, O, Mg, Si, S, Ca, Ti, Cr, Mn, Fe, Ni, and Zn. Zn is often the preferred tracer of the gas-phase metallicity as it is relatively undepleted in the Galactic ISM, especially when the fraction of H in molecular form is low, as is the case in most DLAs. Zn also tracks the Fe abundance in Galactic stars (e.g., \\citet{Niss04}), and the lines of Zn II $\\lambda$$\\lambda$ 2026,2062 are relatively weak and typically unsaturated. These lines can be covered with ground based spectroscopy over a wide range of redshifts, 0.65 $\\la$ z $\\la$ 3.5, which covers a large portion ($\\sim 45\\%$) of the history of the Universe. Abundances of refractory elements such as Cr and Fe relative to Zn also give a measure of the amount of dust extinction \\citep{York06}.  The number of sub-DLAs studied to date is still small, with the largest samples coming from \\citet{Per03a} and \\citet{Mei07}. In this paper we add 10 new systems to our previous medium-resolution measurements of sub-DLAs taken with the Magellan Inamori Kyocera Echelle (MIKE) spectrograph  on the 6.5m Clay telescope at Las Campanas observatory. As the amount of published data on abundances of low redshift sub-DLAs is still small, we present these observations here with a more detailed analysis of the abundances (both absolute and relative) and kinematics in a future paper with our complete sample. The structure of this paper is as follows. In $\\S$ 2, we discuss details of our observations and data reduction techniques. In $\\S$ 3, column density determinations are discussed. $\\S$ 4 gives information on the individual objects, and in $\\S$ 5 we discuss the abundances of these absorbers. We also provide an appendix at the end of this paper, showing plots of the UV spectra with the fits to the Lyman-$\\alpha$ lines. ",
        "conclusions": "The abundances for the observed systems are given in table 14. We have used the total column densities (i.e. the sum of the column densities in the individual components of a system that were determined via profile fitting method) along with the total \\nhI as given in table 1, to determine the abundances of these systems. It would be interesting to determine the abundances in the individual components, as they may possibly vary from component to component. This however is impossible at the present time due to the lack of high resolution UV spectrographs that would be necessary to determine \\nhI in individual components. We have ignored any ionisation corrections while determining these abundances, and have assumed the first ions to be the dominant ionisation species of the elements for which these abundances have been determined, namely Zn, Fe, Mn, Cr, and Si. We discuss the issue of ionisation further below. Solar systems abundances from \\citet{Lodd03} are also given in table 14. The abundances for Ca II are likely only lower limits, as the ionisation potential of 11.868 eV for Ca$^{+}$ is lower than ionisation potential of H$^{0}$ (13.6 eV).  Relative abundances of various elements are also given in table 14, with the column densities determined from the profile fitting analysis. We give the metallicities of these systems, along with the [Zn/Fe] ratio, which is often used as an indicator of dust depletion. We also provide the ratios of [Si/Fe], [Ca/Fe], [Cr/Fe], and [Mn/Fe]. Finally, we provide ratios of the column densities of the adjacent ions Al III/Al II and Mg II/Mg I as well as Mg II/Al III and Fe II/Al III, any of which may provide information about the ionisation in these systems. By matching the observed ratios of the ions mentioned above to those generated through grids of photoionisation models, one can constrain the levels of ionisation, and any ionisation correction factors for the observed abundances.  Based on grids of Cloudy models, is has previously been shown that there appears to be little need for ionisation corrections in most sub-DLA systems \\citep{Des03, Mei07}. We have therefore not applied any ionisation corrections to to the abundances in table 14. We do note however, that the \\za=0.8301 system  with \\nhI=18.95$\\pm$0.18 in Q1054-0020 and the \\za=0.9281 system with \\nhI$<$18.8 in Q1436-0051 both show possible signs of high levels of ionization, with Fe II/Al III=0.74 and Fe II/Al III=0.80 respectively. Substantial ionisation corrections could be necessary. Cloudy models for the \\za=0.9281 system in Q1436-0051 indicate a positive ionization correction factor for Zn in this system of $\\sim$+0.2 to +0.3 dex for $-6 <$log U$<-3$ (indicating that [Zn II/H I] underestimates the true abundance by this amount). We caution that these results may be inaccurate due to uncertainties in the atomic data for Zn.  The ionisation correction factors for Si, Mn, and Fe are all negative over this range of ionisation parameters with values of $\\sim-$0.3 to $-$0.5 dex for Mn and Fe, and values ranging from $\\sim$0.0 to $\\sim-$1.25 for Si. Based on these models, it appears that [Fe/Zn] is suppressed as one goes to lower and lower \\nhI values, the true depletion is actually much higher than it appears based on Zn II and Fe II. Similar values were derived for the \\za=0.8301 system  with \\nhI=18.95$\\pm$0.18 in Q1054-0020 based on Cloudy models. As no adjacent ions that could constrain the ionisation parameter are available for these systems, we leave the abundances unaltered in table 14. We prefer to leave the corrections to a later time when more constraints are available (i.e. when we have UV spectra with coverage of Fe III, S III, Si IV, etc and better atomic data become available).  Variations by component of the relative abundances of elements can also provide information about the local star formation history and possibly ionization. \\citet{Pro02} noticed a remarkably small ($\\sim$0.3 dex) component-by-component dispersion in [Zn/Fe] over a wide range of metallicities. As we only have Zn detections for two of the systems here, we instead look at [Mn/Fe], which is likely less affected by dust depletion (Mn and Fe have similar condensation temperature) and more affected by nucleosynthetic enrichment. Unfortunately, other ions such as Mg II and Al II are too saturated to determine accurate column densities in most of the individual components. The \\za=1.4244 system in Q1037+0028 and the \\za=0.8301 system in Q1054-0020 show similarly low dispersions of [Mn/Fe], but interestingly the high metallicity \\za=0.7377 and \\za=0.9281 systems of Q1436-0051 show higher spreads of [Mn/Fe] of $\\sim$0.6 and $\\sim$0.5 dex respectively. The metal rich sub-DLA observed in \\citet{Per06a} also showed a large dispersion in [Mn/Fe]. The \\za=0.7377 system in Q1436-0051 also shows a 0.6 dex range of [Zn/Fe] in the components. One possible explanation for these larger dispersions could be enrichment by recent supernovae explosions in which the gas has not had time to sufficiently mix.   Zn II was detected in only 2 of the systems, and upper limits on the Zn abundance were placed on the rest of the systems. The Cr II $\\lambda$ 2056 line was detected only in the \\za=1.0929 system in Q1455-0045. Due to the weakness of the Lyman-$\\alpha$  line in the \\za=0.9281 system in Q1436-0051, only an upper limit could be placed on the neutral hydrogen column density, and therefore lower limits on the abundances in this system. This system shows an unusual abundance pattern, with high metallicity ([Zn/H]$>$+0.86, [Fe/H]$>-$0.07, and [Si/H]$>$+0.68) and high depletion ([Zn/Fe]=+0.94). This system also showed an overabundance of Mn, with [Mn/Fe]=+0.22. As Mn is an odd element, and Fe is even, [Mn/Fe] might be below solar in material whose source is dominated by the products of Type II supernovae, but likely not higher if differential dust depletion effects are negligible. In the other case of a Zn detection, the system at \\za=0.7377 in Q1436-0051, Mn is also slightly more abundant relative to Fe, with [Mn/Fe]=+0.04. \\citet{Led02} found a large dispersion in [Mn/Fe] based on a sample of DLA absorbers and concluded that this large range of values not seen in other abundance ratios such as [Cr/Fe] or [Si/Fe] must be due to a combination of both depletion and nucleosynthetic effects.  A similar over-abundance of Mn relative to Fe was also seen by \\citet{Per06a} in another sub-DLA at \\za=0.716 with [Zn/H]=+0.61 and [Mn/Fe]=+0.14. This could perhaps be due to heavy depletion of Fe compared to Mn onto dust grains, unmixed gas that has recently been expelled through stellar winds and supernovae, or an intrinsically non-solar abundance pattern. Cloudy models for this system indicate that the corrections for [Mn/Fe] due to ionisation are small ($\\la$ 0.10 dex for $-6<$log U$<-1$). In the other 3 cases where we have detections of Mn and Fe, but no Zn, Mn is slightly under-abundant relative to Fe, with $<$[Mn/Fe]$>$=$-$0.31. The systems in our sample have an average Fe abundance of $<$[Fe/H]$>$=$-$0.68, which is higher than for typical DLA systems ($<$[Fe/H]$>_{\\rm DLA}\\sim-$1.5). This may be due to less dust depletion or truly higher abundances in these systems, or to undue influence of ionisation effects in the low \\nhI systems.  To determine an average [Zn/H] abundance of the absorbers in this sample, the spectra were shifted to rest wavelengths and averaged. Any obvious non-associated features in the region around the Zn II $\\lambda$ 2026 line in the individual spectra were replaced by random noise with S/N similar to the uncontaminated regions. We show the average spectrum around Zn II $\\lambda$ 2026 \\ang in Figure 11. A feature was seen in the region with \\w=8$\\pm$2 m\\ang and log N$_{\\rm Zn \\ II}$=11.63. The feature at $\\sim$+70 \\kms may be due to absorption from Mg I $\\lambda$ 2026.477, and was not included when determining \\w. The systems in this sample have an average H I column density of log $<$N$_{\\rm H \\ I}$$>$=19.80, giving this sample an average Zn abundance of $<$[Zn/H]$>=-$0.80 based on this composite spectrum. To increase the S/N, we also binned the averaged spectrum to a dispersion of 0.5 \\ang/pixel ( R$\\sim$4000 ). After binning, a feature was detected at a $\\ga$2$\\sigma$ significance level with \\w=11m\\ang. If the feature is assumed to be a true detection, the Zn II column density is log N$_{\\rm Zn \\ II}$=11.80, with a corresponding mean abundance of $<$[Zn/H]$>$=$-$0.63 for these absorbers at 0.7 $\\la$ \\za$ \\la$ 1.5 based on the binned spectrum. In comparison, DLA absorbers have a mean Zn abundance of $<$[Zn/H]$>$$\\sim-$1.0 in this redshift range \\citep{Kul05, Kul07}.  The \\nhI-weighted mean metallicity given by: \\begin{equation}\\rm{[\\langle X/H \\rangle]=log \\langle (X/H) \\rangle - log(X/H)_{\\sun}}\\end{equation} where \\begin{equation}\\rm{ \\langle(X/H)\\rangle=\\frac{\\sum_{i=1}^nN(X)_{i}}{\\sum_{i=1}^nN(H)_{i}}}\\end{equation} has often been used as a quantitative way of estimating the amount of metal enrichment in the Universe in a given epoch \\citep{Kul05, Kul07} For the absorption systems in this sample, the  \\nhI-weighted mean metallicity based on Fe is [$<$Fe/H$>$]=$-$0.76. Using Zn which is less affected by dust depletion, we determine the \\nhI-weighted mean metallicity for these systems at 0.7 $\\la$ \\za$ \\la$ 1.5. to be [$<$Zn/H$>$]=$-$0.40 assuming the upper limits to be detections, and [$<$Zn/H$>$]=$-$0.70 assuming the limits to be zero.   \\setcounter{figure}{10} \\begin{figure} \\includegraphics[height=\\linewidth,  angle=90]{fig11.eps} \\caption{The  unbinned average spectrum of the 10 absorbers in this sample near Zn II $\\lambda$ 2026 \\ang. The horizontal red line indicates the continuum level, and the horizontal dashed green lines show the best fit continuum level increased or decreased by $\\pm$40$\\%$ of the RMS variance in the noise. The vertical dashed lines show the integration limits. The feature seen was detected at $\\ga4\\sigma$ with \\w=8$\\pm$2 m\\ang. } \\end{figure}     In this paper we have presented medium-resolution spectra of 10 sub-DLAs at \\za$\\la$1.5. Although to date the DLA systems have been the preferred tracer of metallicity at high redshift, every absorber with solar or higher metallicity that has been detected has been a sub-DLA \\citep{Pet00, Kh04, Per06a, Pro06, Mei07}. With this latest sample of sub-DLAs, we have found two more systems with high metallicity, [Zn/H]$=-$0.05 and [Zn/H]$>$+0.86 Two other systems in this sample that did not have a Zn II detection, did nonetheless have [Fe/H]$>-0.30$.  These observations suggest that the largely ignored sub-DLAs may provide important clues into the chemical evolution of the Universe. Clearly, not all sub-DLAs have solar or higher metallicity, nor would one expect them to if they trace typical galaxies which have a range of metallicities. Sub-DLAs do appear, however, to have a higher mean metallicity  and show faster evolution than do DLAs \\citep{Kul07}. We continue to find systems with low \\nhI and high metallicities, but higher S/N is needed to determine if the systems with upper limits of [Zn/H]$\\sim$-0.3 are indeed at that level, or are much less, comparable to DLA abundances. With our sample sizes of sub-DLAs now increasing, we can begin to address some of the questions relating to the chemical evolution of metals. Future UV spectra would also allow measurements of key lines of S III and Fe III which could yield important information about ionisation. In forthcoming work, and with an extended sample, we will examine differences in the properties of sub-DLAs and DLAs, including abundances and kinematical structure, as well as the redshift evolution of these properties in sub-DLAs.     \\begin{sidewaystable*} \\begin{minipage}{220mm} \\textbf{Table 14.} Relative abundances and abundance ratios for the observed sub-DLAs. The first line below the heading gives the solar reference abundance for that column. \\begin{tabular}{lccccccccccccc} \\hline \\hline QSO\t\t&\t\\za\t\t&\tLog \\nhI\t&\t[Zn/H]\t \t\t&\t[Fe/H]\t\t\t&\t[Zn/Fe]\t\t\t&\t[Si/Fe]\t\t&\t[Ca/Fe]\t\t\t&\t[Cr/Fe]\t\t&\t[Mn/Fe]\t\t\t&Al III/Al II$^{a}$ \t&\tMg II/Mg I$^{a}$&\tMg II/Al III$^{a}$\t&\tFe II/Al III$^{a}$\\\\ \\hline [X/Y]$_{\\sun}$\t&\t\t\t&\t\t\t&\t$-$7.37\t\t\t&\t$-$4.53\t\t\t&\t$-$2.85\t\t\t&\t+0.07\t\t&\t$-$1.13\t\t\t&\t$-$1.82\t\t&\t$-$1.97\t\t\t&\t\t\t&\t\t\t&\t\t\t\t&\t\t\t\\\\ \\hline Q1037+0028\t& \t1.4244\t \t& \t20.04$\\pm$0.12 \t&\t$<-0.63$\t\t& \t$-$0.67$\\pm$0.12 \t&\t$<+0.04$\t\t& \t0.18$\\pm$0.04\t&\t \t-\t \t&\t$<-0.76$\t&\t $-$0.34$\\pm$0.05 \t&\t$<-0.96$\t&\t -\t \t&\t$>$2.29\t\t \t&\t1.65$\\pm$0.03\t\\\\ Q1054-0020A\t& \t0.8301\t \t& \t18.95$\\pm$0.18 \t&\t$<+0.18$\t\t& \t$-$0.09$\\pm$0.18 \t&\t$<+0.28$\t\t&\t \t-\t&\t \t-\t \t&\t$<-0.04$\t&\t $-$0.09$\\pm$0.07 \t&\t \t- \t&\t$>$2.28\t \t&\t$>$1.24\t\t \t&\t 0.74$\\pm$0.06 \t\\\\ Q1054-0020B\t& \t0.9514\t \t& \t19.28$\\pm$0.02 \t&\t$<-0.21$\t\t& \t$-$1.09$\\pm$0.02 \t&\t$<+0.88$\t\t&\t$<+0.52$\t&\t \t-\t \t&\t$<+0.36$\t&\t$<+0.23$\t\t&\t \t- \t&\t$>$1.68\t \t&\t$>$1.61\t\t \t&\t$>$1.55\t\t\\\\ Q1215-0034\t& \t1.5543\t \t& \t19.56$\\pm$0.02 \t&\t$<-0.56$\t\t& \t$-$0.64$\\pm$0.02 \t&\t$<+0.08$\t\t&\t \t-\t&\t \t-\t \t&\t$<-0.20$\t&\t$<-0.79$\t\t&\t \t- \t&\t-\t \t&\t-\t\t \t&\t-\t\t\\\\ Q1220-0040\t& \t0.9746\t \t& \t20.20$\\pm$0.07\t&\t$<-1.14$\t\t& \t$-$1.33$\\pm$0.07 \t&\t$<+0.19$\t\t&\t \t-\t&\t \t-\t \t&\t$<-0.14$\t&\t$<-0.42$        \t&\t \t- \t&\t$>$2.69\t \t&\t$>$2.19\t\t \t&\t 1.72$\\pm$0.06\t\\\\ Q1228+1018\t& \t0.9376\t \t& \t19.41$\\pm$0.02 \t&\t $<-0.37$\t\t& \t$-$0.30$\\pm$0.02\t&\t$<-0.07$\t\t&\t \t-\t&\t \t-\t \t&\t$<-0.55$\t&\t$<-0.47$\t \t&\t \t- \t&\t$>$2.29\t \t&\t-\t\t \t&\t-\t\t\\\\ Q1330-2056\t& \t0.8526\t \t& \t19.40$\\pm$0.02 \t&\t$<-0.07$\t\t& \t$-$1.07$\\pm$0.02 \t&\t$<+1.00$\t\t&\t \t-\t&\t \t- \t\t&\t$<+0.20$\t&\t$<-0.22$\t\t&\t \t- \t&\t$>$1.88\t \t&\t$>$1.62\t\t \t&\t 1.21$\\pm$0.05\t\\\\ Q1436-0051A\t& \t0.7377\t \t& \t20.08$\\pm$0.11 \t&\t$-$0.05$\\pm$0.12 \t& \t$-$0.61$\\pm$0.11 \t& \t+0.58$\\pm$0.04\t \t&\t \t-\t&\t $-$0.98$\\pm$0.03 \t&\t$<-0.41$\t& \t+0.04$\\pm$0.03 \t\t&\t \t- \t&\t-\t \t&\t-\t\t \t&\t- \t\t\\\\ Q1436-0051B\t& \t0.9281\t \t& \t$<$18.8\t\t&\t$>$+0.86\t \t& \t$>-$0.07\t \t& \t+0.94$\\pm$0.06 \t\t& \t+0.75$\\pm$0.05 \t&\t $-$0.79$\\pm$0.04 \t&\t$<+0.20$\t& \t+0.22$\\pm$0.05 \t\t&\t \t- \t&\t$>$2.63\t \t&\t$>$1.72\t\t \t&\t0.80$\\pm$0.03   \\\\ Q1455-0045\t& \t1.0929\t \t& \t20.08$\\pm$0.06\t&\t$<-0.80$\t\t& \t$-$0.98$\\pm$0.06 \t&\t$<+0.18$\t\t& \t+0.00$\\pm$0.10 \t&\t \t-\t \t&    $-$0.35$\\pm$0.13\t& \t-0.51$\\pm$0.15 \t\t&\t$<-$0.18 \t&\t$>$2.39\t \t&\t$>$1.27\t\t \t&\t0.92$\\pm$0.02\t\\\\ \\hline \\end{tabular} $^{a}$Ratios of column densities \\end{minipage}  \\setcounter{table}{14} \\end{sidewaystable*}"
    },
    "0801/0801.2944.txt": {
        "abstract": "\\noindent In these lecture notes, we present the equations presently used in stellar interior models in order to compute the effects of axial rotation. We discuss the hypotheses made. We suggest that the effects of rotation might play a key role at low metallicity. ",
        "introduction": "Axial rotation modifies the hydrostatic equilibrium configuration by adding a centrifugal acceleration term in the hydrostatic equation, induces many instabilities driving the transport of angular momentum and of chemical species in radiative zones and changes the mass loss rates. In the present lecture notes we shall consider the case of models without magnetic fields.  \\subsection{Hydrostatic effects of rotation}  \\subsubsection{The equations of stellar structure.}  In a rotating star, the equations of stellar structure need to be modified  \\cite{KipTho70}. The usual spherical coordinates must be replaced by new coordinates characterizing the equipotentials. The classical method applies when  the effective gravity can be derived from a potential $\\Psi = \\Phi - \\frac{1}{2} \\Omega^2 r^2 \\sin^2 \\theta$, i.e. when the  problem is conservative. There, $\\Phi$ is the gravitational potential, $\\Omega$ the angular velocity, $r$ the radius at the colatitude $\\theta$. If  the rotation law is shellular ({\\it i.e.} such that $\\Omega$ is constant on isobaric surfaces see below), the problem is non--conservative. Most existing models of rotating stars apply, rather inconsistently, the classical scheme by \\cite{KipTho70}. However, as shown by \\cite{MM97}, the equations of stellar structure can still be written consistently, in term of a coordinate referring to the mass inside the isobaric surfaces\\footnote{For shellular rotation, the shape of the isobaric surfaces are given by the same expression as the one giving the shape of the equipotentials in conservative cases provided some changes of variables are performed.} . Thus, the problem of the stellar structure of a differentially rotating star in a shellular rotation state can be kept one--dimensional.  \\subsubsection{The Roche model.}  In all the derivations, we shall use the Roche model, {\\it i.e.} we approximate the gravitational potential by $GM_{\\overline{r}}/\\overline{r}$ where $M_{\\overline{r}}$ is the mass inside the isobaric surface with a mean radius $\\overline{r}$. The radius $\\overline{r}$ which labels each isobaric surface is defined by $\\overline{r}=\\left( V_{\\overline{r}}/(4/3 \\pi)\\right )^{1/3}$ where $V_{\\overline{r}}$ is the volume (deformed by rotation) inside the isobaric surface considered. Apart from the case of extreme rotational velocities, the parameter $\\overline{r}$ is close to  the average radius of an isobar, which is the radius at $P_2 (\\cos \\vartheta) =0$, namely for $\\vartheta = 54.7$ degrees.  In the frame of the Roche model, the shape of a meridian at the surface of the star (which is an isobaric surface) is given by couples of $R$ and $\\theta$ values satisfying the following equation: \\begin{equation} {GM \\over R}+{1 \\over 2}\\Omega^2 R^2 \\sin^2 \\theta={GM \\over R_{\\rm p}}, \\label{eq1} \\end{equation} where $R$ is the radius at colatitude $\\theta$, $\\Omega$ the angular velocity, $M$ the mass inside the surface and $R_{\\rm p}$, the polar radius. Thus the shape of the surface (as well as of any isobaric surface inside the star) is determined by three parameters $M$, $\\Omega$ and $R_{\\rm p}$. The first two $M$ and $\\Omega$ are independent variables. The third one is derived from the first two and the equations of stellar structure. Setting $x=\\left({GM \\over \\Omega^2}  \\right)^{-1/3} R$, one can write Eq.~\\ref{eq1} (see \\cite{KipTho70}) \\begin{equation} {1 \\over x}+{1 \\over 2}x^2 \\sin^2 \\theta={1 \\over x_{\\rm p}}. \\label{eq1b} \\end{equation} With this change of variable, the shape of an equipotential is uniquely determined by only one parameter $x_{\\rm p}$.  Setting \\begin{equation} f={R_{\\rm e} \\over R_{\\rm p}}, \\label{eq2} \\end{equation} where $R_{\\rm e}$ is the equatorial radius, one easily obtains from Eq.~\\ref{eq1} that \\begin{equation} R_{\\rm p}=\\left({GM \\over \\Omega^2} \\right)^{1/3} \\left({2 (f-1)\\over f^3} \\right)^{1/3}=\\left({GM \\over \\Omega^2} \\right)^{1/3} x_{\\rm p}. \\label{eq3} \\end{equation} The above equation relates the inverse of the oblateness $f$ to $R_{\\rm p}$.  \\subsubsection{The von Zeipel theorem and its consequences.}  The von Zeipel theorem \\cite{vZ24}  expresses that the radiative flux $\\vec{F}$ at some colatitude $\\vartheta$ in a rotating star is proportional to the local effective  gravity $\\vec{g_\\mathrm{eff}}$. \\cite{Ma99} has generalized this theorem to the case of shellular rotation and the  expression of the flux $\\vec{F}$ for a star with angular velocity $\\Omega$ on the isobaric stellar surface  is  \\begin{equation} \\vec{F}  =  - \\frac{L(P)}{4 \\pi GM_{\\star}} \\vec{g_{\\rm{eff}}} [1 + \\zeta(\\vartheta)]  \\quad {\\mathrm{with}} \\quad \\label{Ma23} \\end{equation} %eqn 2.3  \\begin{equation} M_{\\star} = M \\left( 1 - \\frac{\\Omega^2} {2 \\pi G \\rho_{\\rm{m}}}  \\right) \\quad {\\mathrm{and}} \\quad \\label{Ma24} \\end{equation}%eqn 2.4  \\begin{equation} \\zeta(\\vartheta) = \\left[\\left(1 - \\frac{\\chi_T}{\\delta}\\right) \\Theta + \\frac{H_T}{\\delta} \\frac{d\\Theta}{dr}\\right] P_{2}(\\cos \\vartheta). \\label{Ma25} \\end{equation}%eqn 2.5  \\noindent There, $\\rho_{\\rm{m}}$ is the internal average density, $\\chi = 4acT^3/(3 \\kappa \\rho)$ and $\\chi_{T}$ is the partial derivative with respect to T.  The quantity $\\Theta$ is defined by $\\Theta = \\frac{\\tilde{\\rho}}{\\bar{\\rho}}$, i.e. the ratio of the horizontal density fluctuation to the average density on the isobar \\cite{Z92}. One has the thermodynamic coefficients $\\delta = - (\\partial \\ln\\rho / \\partial \\ln T)_{P, \\mu}$, $H_{T}$ is the temperature scale height. The term $\\zeta(\\vartheta)$, which expresses the deviations of the von Zeipel theorem due to the baroclinicity of the star, is generally very small, (cf. \\cite{Ma99}).   Let us emphasize that the flux is proportional to $\\vec{g_\\mathrm{eff}}$ and not to $\\vec{g_\\mathrm{tot}}$. This results from the fact that the equation of hydrostatic equilibrium is $\\frac{\\vec{\\nabla} P}{\\rho} = - \\vec{g_\\mathrm{eff}}$. The effect of radiation pressure is already counted in the expression of $P$, which is the total pressure. We may call $M_{\\star}$ the effective mass, i.e. the mass reduced by  the centrifugal force. This is the complete form of the von Zeipel theorem in a differentially rotating star with shellular rotation, whether or not one is close to the Eddington limit.  This theorem has numerous consequences. Some of them are discussed below. \\vskip 2mm \\noindent{\\underline{The position of the star in the HR diagram}} \\vskip 2mm A fast rotating star has stronger radiative fluxes at the pole than at the equator. Therefore the position of such a star in the HR diagram will depend of the angle between the line of sight and the rotational axis (inclination angle). If for instance that angle is 90 degrees, a great part of the light will come from the equatorial belt characterized by lower radiative fluxes and cooler effective temperatures, while when the star is observed pole-on most of the light will come from the hot polar region characterized by stronger fluxes and higher effective temperatures. Thus the perceived luminosity and effective temperature (and also effective gravity) of a star depend on the inclination angle. This has to be kept in mind when comparisons are made with observed quantities. Computations of the effect of the inclination angle on the emergent luminosity, colors and spectrum have been performed by \\cite{MP70}. The effect of the inclination angle on the determination of the effective gravity is discussed in \\cite{HG06}. In general, a theoretical evolutionary track is given in term of total luminosity and of an average effective temperature defined by $T_{\\rm eff}^4= L/(\\sigma S(\\Omega))$, where $\\sigma$ is Stefan's constant and $S(\\Omega)$ the total actual stellar surface. The total luminosity (corresponding to the integrated flux over the surface) does not depend on the angle of view, but cannot be directly compared to the ``observed'' luminosity deduced from the apparent luminosity coming from the hemisphere directed toward us. Let us note however that for surface velocities inferior to about 70\\% of the critical velocity these effects remain quite modest. As a numerical example, the ratio ($T_{\\rm eff}$(pole)-$T_{\\rm eff}$(equator))/$T_{\\rm eff}$(equator) becomes superior to 10\\% only for $\\omega > 0.7$. At break-up, the effective temperature of the polar region is about a factor two higher than that of the equatorial one. \\vskip 2mm \\noindent{\\underline{The Eddington luminosity}} \\vskip 2mm Let us express  the total gravity at some colatitude $\\vartheta$, taking into account the radiative acceleration (cf. \\cite{Ma99})  \\begin{equation} \\vec{g_\\mathrm{rad}} = \\frac{1}{\\rho} \\vec{\\nabla} P_\\mathrm{rad} = \\frac{\\kappa(\\vartheta)\\vec{F}}{c} \\; , \\end{equation}%eqn 2.6  \\noindent thus one has  \\begin{eqnarray} \\vec{g_\\mathrm{tot}}    =\\vec{g_\\mathrm{eff}}+\\vec{g_\\mathrm{rad}}= \\vec{g_\\mathrm{eff}}+\\frac{\\kappa(\\vartheta)\\vec{F}}{c} % \\vec{g_\\mathrm{eff}} \\left[1 -  \\frac{\\kappa(\\vartheta)L(P) %[1 + \\zeta (\\vartheta)] } %{4 \\pi cGM (1 - \\frac{\\Omega^2} {2 \\pi G \\rho_{\\mathrm{m}} } )}\\right]. \\end{eqnarray}%eqn 2.7     \\noindent The rotation effects appear both in $\\vec{g_\\mathrm{eff}}$ and in $\\vec{F}$. We may also consider the local limiting flux. The condition $\\vec{g_\\mathrm{tot}}= \\vec{0}$ allows us to define a limiting flux,  \\begin{equation} \\vec{F_{\\mathrm{lim}}}(\\vartheta) = - \\frac{c}{\\kappa(\\vartheta)} \\vec{g_{\\mathrm{eff}}}(\\vartheta) \\; . \\end{equation}%eqn 2.8   \\noindent From that we may define the ratio  $\\Gamma_{\\Omega}(\\vartheta)$ of the  actual flux (see Eq.~\\ref{Ma23}) $F(\\vartheta)$ to the limiting local flux in a rotating star,  \\begin{eqnarray} \\Gamma_{\\Omega}(\\vartheta) = \\frac{\\vec{F}(\\vartheta)}{\\vec{F_{\\mathrm{lim}}}(\\vartheta)}= \\frac{ \\kappa (\\vartheta) \\; L(P)[1+\\zeta (\\vartheta)]}{4 \\pi cGM \\left( 1 - \\frac{\\Omega^2}{2 \\pi G \\rho_{\\rm{m}}}  \\right) }\\; . \\label{Ma29} \\end{eqnarray}%eqn 2.9   \\noindent As a matter of fact, $\\Gamma_{\\Omega}(\\vartheta)$ is the local Eddington ratio and \\begin{eqnarray} L_{\\rm Edd} = \\frac{4 \\pi cGM \\left( 1 - \\frac{\\Omega^2}{2 \\pi G \\rho_{\\rm{m}}}  \\right) }{\\kappa (\\vartheta) \\; [1+\\zeta (\\vartheta)] } \\label{Lmax} \\end{eqnarray}%eqn 2.9 is the local Eddington luminosity. For a certain angular velocity $\\Omega$ on the isobaric surface, the maximum permitted luminosity of a star is reduced by rotation, with respect to the usual Eddington limit. In the above relation, $\\kappa(\\vartheta)$ is the largest value of the opacity on the surface of the rotating star. For O--type stars with photospheric opacities dominated by electron scattering, the opacity $\\kappa$ is the same everywhere on the star.  For zero rotation, the usual expressions are found: $\\Gamma_{\\Omega}(\\vartheta) = \\Gamma=\\frac{\\kappa L}{4 \\pi c GM}$ and $L_{\\rm Edd} =\\frac{4 \\pi cGM }{\\kappa}$ . \\vskip 2mm \\noindent{\\underline{Critical limits.}} \\vskip 2mm Critical limits correspond to values of respectively the luminosity and/or the velocity which impose that the total gravity become equal to zero at least at some peculiar places at the surface. We may identify different limits \\cite{mm6}:  \\vspace{2mm} \\noindent -- We speak of the Eddington or $\\Gamma$--limit, when rotation effects can be neglected  and $\\vec{g_\\mathrm{rad}} + \\vec{g_\\mathrm{grav}} = \\vec{0}$\\footnote{$\\vec{g_\\mathrm{grav}}$ is the gravitational acceleration.}, which implies that  \\begin{equation} \\Gamma = \\frac{\\kappa L}{4 \\pi c GM}  \\; \\rightarrow  \\; 1. \\end{equation}%eqn 1.2  \\noindent In that case $L = L_{\\rm Edd}=4 \\pi c GM/\\kappa$. The opacity $\\kappa$ considered here is the total opacity.  \\vspace{2mm} \\noindent --The critical velocity or $\\Omega$--limit is reached for a star with an angular velocity $\\Omega$ at the surface, when the effective gravity  $\\vec{g_\\mathrm{eff}} = \\vec{g_\\mathrm{grav}} + \\vec{g_\\mathrm{rot}} = \\vec{0}$ and in addition when radiation pressure effects can be neglected.  \\vspace{2mm} \\noindent --\\textit{The $\\Omega \\Gamma$--limit is  reached when the total gravity $\\vec{g_\\mathrm{tot}} = \\vec{0}$, with significant effects of both rotation and radiation.} This is the  general case. It should lead to the two above cases in their respective limits.  Using relation \\ref{Ma29}, we may write the expression for the total gravity  as  \\begin{equation} \\vec{g_\\mathrm{tot}} = \\vec{g_\\mathrm{eff}} \\left[ 1 - \\Gamma_{\\Omega}(\\vartheta) \\right] \\; . \\end{equation}%eqn 2.10   \\noindent This  shows that the expression for the total acceleration in a rotating star is similar to the usual one, except  that $\\Gamma$ is replaced by the local value  $\\Gamma_{\\Omega}(\\vartheta)$. Indeed, contrarily to expressions such as $\\vec{g_\\mathrm{tot}} = \\vec{g_\\mathrm{eff}} \\left( 1 - \\Gamma \\right)$ often found in literature, we see that the appropriate Eddington factor given by Eq.~\\ref{Ma29} also depends on the angular velocity $\\Omega$ on the isobaric surface.  From Eq.~\\ref{Ma29}, we note that over the surface of a rotating star, which has a varying gravity and T$_\\mathrm{eff}$,  $\\Gamma_{\\Omega}(\\vartheta)$ is the highest at the latitude where $\\kappa(\\vartheta)$ is the largest, (if we neglect the effects of $\\zeta (\\vartheta)$, which is justified in general). If the opacity increases with decreasing T as in  hot stars, the opacity is the highest at the equator and there the limit $\\Gamma_{\\Omega}(\\vartheta) = 1 $ may by reached first. Thus, it is to be stressed that if the limit  $\\Gamma_{\\Omega}(\\vartheta) = 1 $ happens to be met at the equator, it is not because $\\vec{g_\\mathrm{eff}}$ is the lowest there, but because the opacity is the highest ! The reason for no direct dependence on $\\vec{g_\\mathrm{eff}}$ is because both terms $\\vec{g_\\mathrm{eff}}$ cancel each other in the expression~\\ref{Ma29} of the flux ratio.  The critical velocity is reached when somewhere on the star one has $\\vec{g_\\mathrm{tot}}=\\vec{0}$, i.e.  \\begin{equation} \\vec{g_\\mathrm{eff}} \\;\\left[1 - \\Gamma_{\\Omega}(\\vartheta) \\right] = \\vec{0} . \\label{Ma313} \\end{equation}%eqn 3.13  \\noindent This equation has two roots. The first one $v_\\mathrm{crit,1}$ is given by the usual condition $\\vec{g_\\mathrm{eff}}= \\vec{0}$, which implies the equality $\\Omega^2 R^{3}_\\mathrm{eb}/(GM) =1$ at the equator. This corresponds to an equatorial critical velocity  \\begin{equation} v_\\mathrm{crit, 1} = \\Omega \\; R_\\mathrm{eb} = \\left( \\frac{2}{3} \\frac{GM}{R_\\mathrm{pb}} \\right)^{\\frac{1}{2}} \\; . \\label{Ma314} \\end{equation}%eqn 3.14  \\noindent $R_\\mathrm{eb}$  and $R_\\mathrm{pb}$ are respectively the equatorial  and polar radius at  the critical velocity. We notice that the critical velocity $v_\\mathrm{crit, 1}$ is independent on the Eddington factor. To this extent, this is in agreement with \\cite{Gla98}. The basic physical reason for this independence is quite clear: the radiative flux tends toward zero when the effective gravity is zero, thus there is no effect of the radiative acceleration when this occurs.  Equation~\\ref{Ma313}  has a second root, which is given by the condition $\\Gamma_{\\Omega}(\\vartheta)$ = 1. If we call the  Eddington ratio $\\Gamma_\\mathrm{max}$ the maximum value  of $\\kappa(\\vartheta) L(P)/(4 \\pi c GM)$ over the surface (in general at equator), \\cite{mm6} has shown that for  $\\Gamma_\\mathrm{max} < 0.639$, no value of $\\Omega$ can lead to $\\Gamma_{\\Omega}(\\vartheta)$ = 1. Physically this means that when the star is sufficiently far from the Eddington limit, the effects of rotation on the radiative equilibrium are not sufficient for lowering the Eddington luminosity such that it may have an impact on the value of the critical velocity. In that case, Eq.~\\ref{Ma313} has only one root given by the classical expression.  For a given large enough $\\Gamma_\\mathrm{max}$ (i.e. larger than 0.639), a second root is obtained given by  \\begin{eqnarray} v_\\mathrm{crit, 2}^2 = \\frac{9}{4} \\;v_\\mathrm{crit, 1}^2 \\; \\frac{1-\\Gamma}{V^{\\prime}(\\omega)} \\; \\frac{ R^{2}_\\mathrm{e}(\\omega)}{R^2_{\\mathrm{pb}}} \\end{eqnarray}%eqn 3.19  The quantity $\\omega$ is the fraction $\\Omega/\\Omega_\\mathrm{c}$ of the angular velocity at break--up. The quantity $V^{\\prime}(\\omega)$ is the ratio of the actual volume of a star with rotation $\\omega$ to the volume of a sphere  of radius   $R_\\mathrm{pb}$. $V^{\\prime}(\\omega)$ is  obtained by the integration of the solutions of the surface equation for a given value of the parameter $\\omega$.  For $\\Gamma_\\mathrm{max} > 0.639$, this second root is inferior to the first one. Thus it is encountered first and is therefore the expression of the critical velocity that has to be used. \\vskip 2mm \\noindent{\\underline{The mass loss rates}} \\vskip 2mm Due to the von-Zeipel theorem, the radiative flux, which is the driving force for the stellar winds of massive hot stars, varies as a function of the colatitude. This effect, when accounted for in the theory of the line driven wind theory, lead to an enhancement of the quantities of mass lost and to wind anisotropies. We shall describe in more details these effects in the Section 1.3 below. \\vskip 2mm \\noindent{\\underline{Three remarks}} \\vskip 2mm 1) The classical critical angular velocity or the $\\Omega$-limit (to distinguish it from the $\\Omega\\Gamma$-limit as defined by \\cite{mm6}) in the frame of the Roche model is given by \\begin{equation} \\Omega_{\\rm crit}=\\left({2 \\over 3}\\right)^{3 \\over 2}\\left({GM \\over R^3_{\\rm pb}}\\right)^{1 \\over 2}, \\label{eq4} \\end{equation} where $R_{\\rm pb}$ is the polar radius when the surface rotates with the critical velocity. The critical velocity is given by \\begin{equation} \\upsilon_{\\rm crit}=\\left({2 \\over 3}{GM \\over R_{\\rm pb}}\\right)^{1 \\over 2}. \\label{eq4b} \\end{equation} Replacing $\\Omega$ in Eq.~\\ref{eq3} by $(\\Omega/\\Omega_{\\rm crit}) \\cdot \\Omega_{\\rm crit}$ and using Eq.~\\ref{eq4}, one obtains a relation between $R_{\\rm p}$, $f$ and $R_{\\rm pb}$, \\begin{equation} R_{\\rm p}={3 \\over 2} R_{\\rm pb} \\left({\\Omega_{\\rm crit} \\over \\Omega}\\right)^{2/3} \\left({2 (f-1)\\over f^3} \\right)^{1/3}, \\label{eq5} \\end{equation} Thus one has \\begin{equation} {\\Omega \\over \\Omega_{\\rm crit}}=\\left({3 \\over 2}\\right)^{3/2} \\left({R_{\\rm pb} \\over R_{\\rm p}}\\right)^{3/2} \\left({2 (f-1)\\over f^3} \\right)^{1/2}. \\label{eq6} \\end{equation} With a good approximation (see below) one has that $R_{\\rm pb}/ R_{\\rm p} \\simeq 1$ and therefore \\begin{equation} {\\Omega \\over \\Omega_{\\rm crit}} \\simeq \\left({3 \\over 2}\\right)^{3/2} \\left({2 (f-1)\\over f^3} \\right)^{1/2}. \\label{eq7} \\end{equation} This equation is quite useful since it allows the determination of $\\Omega_{\\rm crit}$ from quantities obtained with a model computed for $\\Omega$. This is not the case of Eq.~\\ref{eq5} which involves $R_{\\rm pb}$ whose knowledge can only be obtained by computing models at the critical limit. In general Eq.~\\ref{eq7} gives a very good approximation of $\\Omega_{\\rm crit}$ (see \\cite{EMM} for a discussion of this point).  Setting $\\upsilon$ the velocity at the equator, one has that \\begin{equation} {\\upsilon \\over \\upsilon_{\\rm crit}}={\\Omega R_{\\rm e}\\over\\Omega_{\\rm c}R_{\\rm eb}} ={\\Omega \\over \\Omega_{\\rm crit}}{R_{\\rm e}\\over R_{\\rm p}}{R_{\\rm p}\\over R_{\\rm pb}}{R_{\\rm pb}\\over R_{\\rm eb}}, \\label{eq7bis} \\end{equation} where $R_{\\rm eb}$ is the equatorial radius when the surface rotates with the critical velocity. Using Eq.~\\ref{eq5} above, and the fact that in the Roche model $R_{\\rm pb}/ R_{\\rm eb}=2/3$, one obtains \\begin{equation} {\\upsilon \\over \\upsilon_{\\rm crit}}= \\left({\\Omega \\over \\Omega_{\\rm crit}}2(f-1)\\right)^{1/3} \\label{eq8} \\end{equation} The relations between $\\upsilon/\\upsilon_{\\rm crit}$ and $\\Omega/\\Omega_{\\rm crit}$ obtained in the frame of the Roche model (see Eq.~\\ref{eq8}) for the 1 and 60 M$_\\odot$ stellar models at $Z=0.02$ are shown in Fig.~\\ref{ocvc}. In case we suppose that the polar radius remains constant $R_{\\rm pb}/R_{\\rm p}=1$, then  Eq.~\\ref{eq7} can be used and one obtains a unique relation between $\\Omega/\\Omega_{\\rm crit}$ and $\\upsilon/\\upsilon_{\\rm crit}$, independent of the mass, metallicity and evolutionary stage considered. One sees that the values of $\\upsilon/\\upsilon_{\\rm crit}$ is smaller than that of $\\Omega/\\Omega_{\\rm crit}$ by at most $\\sim$25\\%. At the two extremes the ratios are of course equal.  \\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{ocvc.eps}} \\caption{Relation between $\\upsilon/\\upsilon_{\\rm crit}$ and $\\Omega/\\Omega_{\\rm crit}$ obtained in the frame of the Roche model. The continuous line is obtained assuming $R_{\\rm pb}/R_{\\rm p}=1$ (see text and Eqs.~\\ref{eq7} and \\ref{eq8}). The dot-dashed lines show the relations for the Z=0.02 models with 1 and 60 M$_{\\odot}$ using Eq.~\\ref{eq6}. The dotted line is the line of slope 1.} \\label{ocvc} \\end{figure}  An interesting quantity is the ratio of the centrifugal acceleration, $a_{\\rm cen}$, to the gravity, $g_{\\rm e}$, at the equator \\begin{equation} {a_{\\rm cen} \\over g_{\\rm e}}={\\Omega^2 R^3_{\\rm e}\\over G M}= \\left({\\Omega \\over \\Omega_{\\rm crit}}\\right)^2 \\left({2 \\over 3}\\right)^3 f^3 \\left({R_{\\rm p}\\over R_{\\rm pb}}\\right)^3, \\label{eq9} \\end{equation} where we have used Eq.~\\ref{eq4} and divided/multiplied by $R_{\\rm pb}^3$. Replacing $R_{\\rm pb}/R_{\\rm p}$ by its expression deduced from Eq.~\\ref{eq5}, we obtain \\begin{equation} {a_{\\rm cen} \\over g_{\\rm e}}=2(f-1). \\label{eq10} \\end{equation} We can check that at the critical limit, when $f=3/2$, then $a_{\\rm cen}=g_{\\rm e}$. \\vskip 2mm 2) It is interesting to note that the fact that the effective temperature varies as a function of the colatitude on an isobaric surface does not necessarily imply that the temperature varies as a function of the colatitude on an isobaric surface. For instance in the conservative case, the temperature is constant on equipotentials which are also isobaric surfaces. This simply illustrates the difference between the effective temperature whose definition is related to the radiative flux ($F=\\sigma T_{\\rm eff}^4$) and hence to the {\\it temperature gradient} and the temperature itself. Interestingly, one has that in a conservative case, $\\Gamma_\\Omega(\\theta)$ is constant on isobaric surfaces (neglecting the term $\\zeta(\\vartheta)$). This means that when the $\\Omega\\Gamma$-limit is reached, it is reached over the whole stellar surface at the same time. This is in contrast with the $\\Omega$-limit which is reached first at the equator. \\vskip 2mm 3) Recently we have examined the effects of rotation on the thermal gradient and on the Solberg--Hoiland term by analytical developments and by numerical models \\cite{MCM08}. Writing the criterion for convection in rotating envelopes, we show that the effects of rotation on the thermal gradient are much larger and of  opposite sign to   the effect of the Solberg--Hoiland criterion. On the whole, rotation favors convection in stellar  envelopes at the  equator  and to a smaller extent at the poles. In a rotating 20 M$_{\\odot}$ star at 94\\% of the critical angular velocity, there are two convective envelopes, the biggest one has a thickness of 13.2\\% of the equatorial radius. The convective  layers are  shown in  Fig. \\ref{P94}.  They are more extended than without rotation. In the non-rotating model, the corresponding convective zone has a thickness of only 4.6\\% of the radius. The occurrence of outer convection in massive stars has many consequences (see \\cite{MCM08}).  \\begin{figure}[t] \\includegraphics[angle=00,width=8.8cm]{P020z20S7ZCsansSH.eps} \\caption{2--D representation of the  external convective zones (in red) and of the convective core (blue)  in a model of 20 M$_{\\odot}$ with $X=0.70$ and $Z=0.020$ at the end of MS evolution with fast rotation ( $\\Omega/\\Omega_{\\mathrm{crit}}=0.94$). The axis are in units of cm. Figure taken from \\cite{MCM08}.} \\label{P94} \\end{figure}  \\vskip 2mm \\subsection{Transport mechanisms of angular momentum and of chemical species}  In a solid body rotation state, the first instability to set up is a thermal instability called the meridional circulation (see below). It consists in large meridional currents which transport angular momentum either inside-out or conversely from the outer regions toward the inner ones. Such meridional currents rapidly build ut gradients of the angular momentum both in the ``horizontal direction'' ({\\it i.e.} along isobaric surface) and in the vertical one. Along isobaric surface, any gradient of $\\Omega$ triggers a strong horizontal turbulence. Indeed in that direction the instability can develop without having to overcome any stable density gradient. As a consequence any gradient of $\\Omega$ along isobaric surfaces is rapidly erased and the star settles into a ``shellular'' rotation state \\cite{Z92}. This means that $\\Omega$ can be considered as nearly constant along isobars. In the vertical direction, where in a radiative zone, a stable density gradient counteracts any instability, the gradients of $\\Omega$ are eroded on much longer timescales  (see below). The equations below describe the interactions of meridional currents and of shear instabilities in a state of shellular rotation \\cite{Z92}.  %Let us suppose that stars begin their life on the ZAMS in a state of solid body rotation. This %is probably not what happens in reality but this initial state is rapidly erased and %forgotten . Indeed, it has been shown (\\cite{Den99}; \\cite{}) that, %under the actions of both meridional currents and of shear instabilities (see below), the profile of $\\Omega$ evolves towards an equilibrium configuration. %The timescale for the %adjustment of the rotation law, i.e.  $t_{\\rm circ} \\simeq R/U$, %is short with respect to %the MS lifetime for any significant rotation (1 to 2\\% of the Main-Sequence lifetime).  Later on during the bulk %of MS evolution, %the profiles $\\Omega(r)$ will evolve more slowly.    \\subsubsection{Meridional circulation.}  Meridional circulation is an essential mixing mechanism in rotating stars and there is a considerable literature on the subject (see ref. in \\cite{Tass90}). The velocity of the meridional circulation in the case of shellular rotation was derived by \\cite{Z92}. The velocity of meridional circulation is derived from the equation of energy conservation \\cite{Mes53} \\begin{eqnarray} \\rho T\\left[{\\partial S \\over \\partial t}+({\\bf e}_r \\dot r +{\\bf U})\\cdot{\\bf \\nabla}S\\right]={\\rm div}(\\chi{\\bf \\nabla}T)+\\rho\\epsilon-{\\rm div}{\\bf F}_h \\label{eqn9} \\end{eqnarray} where $S$ is the entropy per unit mass, $\\chi$ the thermal conductivity, $\\epsilon$ the rate of nuclear energy per unit mass and ${\\bf F}_h$ the flux of thermal energy due to horizontal turbulence. All the quantities are expanded linearly around their average on a level surface or isobar, using Legendre Polynomials $P_2(\\cos\\theta)$. For instance $$T(P,\\theta)=\\bar T(P)+\\tilde{T}P_2(\\cos\\theta).$$ Then  Eq.~\\ref{eqn9} is linearized and an expression for $U_2$ can be deduced \\cite{Z92}. Using the same method \\cite{MZ98} revised the expression for $U_2$ to account for expansion and contraction in non--stationary models. They also studied the effects of the $\\mu$--gradients (mean molecular weight gradients), of the horizontal turbulence and considered a general equation of state. They obtained \\begin{eqnarray} U_2(r)={P \\over \\bar \\rho \\bar g C_P\\bar T[\\nabla_{\\rm ad}-\\nabla+(\\varphi/\\delta)\\nabla_\\mu ]} \\times\\left[{L\\over M_*}(E_\\Omega+E_\\mu)+{C_P \\over \\delta}{\\partial \\Theta \\over \\partial t}\\right], \\label{Umer} \\end{eqnarray} where $M_{\\star}=M \\left( 1 - \\frac{\\Omega^2}{2 \\pi G \\rho_{\\rm{m}}}  \\right)$ is the reduced mass and the other symbols have the same meaning as in \\cite{Z92} and \\cite{MZ98} \\footnote{$U_2$ is the same as $U$ in Eq.~\\ref{urt}}. The driving term in the square brackets in the second member is $E_{\\Omega}$.  It behaves mainly like $E_{\\Omega}  \\simeq  \\frac{8}{3} \\left[ 1 - \\frac{{\\Omega^2}} {2\\pi G\\overline{\\rho}}\\right] \\left( \\frac{\\Omega^2r^3}{GM}\\right)$ The term $\\overline{\\rho}$ means the average on the considered equipotential. The term with the minus sign in the square bracket is the Gratton--\\\"{O}pik term, which becomes important in the outer layers when the local density is small. This term  produces negative values of $U_2(r)$ (noted $U(r)$ from now), meaning that the circulation is going down along the polar axis and up in the equatorial plane. This makes an outward transport of angular momentum, while a positive $U(r)$ gives an inward transport. At lower $Z$, the Gratton--\\\"{O}pik term is negligible, which contributes to make larger $\\Omega$--gradients in lower $Z$ stars.  Recently \\cite{Mat04} rederived the system of partial differential equations, which govern the transport of angular momentum, heat and chemical elements. They expand the departure from spherical symmetry to higher order and include explicitly the differential rotation in latitude, to first order. Boundary conditions for the surface and at the frontiers between radiative and convective zones are also explicitly given in this paper. %The effects of the vertical %$\\mu$--gradient  $\\nabla_{\\mu}$  and of the horizontal %turbulence  on meridional circulation are very important %and   they were taken into account by \\cite{MZ98}. %The expression of the radial term of the vertical component %$U(r)$ of the meridional circulation %by \\cite{MZ98} is   %\\begin{eqnarray} %U(r) =  \\frac{P}{{\\rho} {g} C_{\\!P} {T} %\\, [\\nabla_{\\rm ad} - \\nabla + % (\\varphi/\\delta) \\nabla_{\\mu}] } %\\left\\{  \\frac{L}{M_\\star} (E_\\Omega + E_\\mu) \\right\\} \\; . %\\label{Umer} %\\end{eqnarray}  %\\noindent %$P$ is the pressure, $C_P$ the specific heat, %$E_{\\Omega}$ and $E_{\\mu}$ are  terms depending on the $\\Omega$-- %and $\\mu$--distributions respectively, up to the third order derivatives %and on various thermodynamic quantities (see details in \\cite{MZ98}). %The term $\\nabla_{\\mu}$ may be %one or two orders of magnitude larger than $\\nabla_{\\rm ad}-\\nabla$ %in some layers, %so that $U (r)$ may be reduced by the same ratio. This is one of the %important differences introduced by the work \\cite{MZ98}.  \\subsubsection{Shellular rotation.}  The differential rotation which results from the evolution and transport of the angular momentum makes the stellar interior highly turbulent. As explained above, the turbulence is  very anisotropic, with a much  stronger geostrophic--like transport in the horizontal direction than in the vertical one \\cite{Z92}, where stabilisation is favoured by the stable density gradient. This strong horizontal transport is characterized by a large diffusion coefficient $D_{\\rm{h}}$. Various expressions have been proposed: \\begin{itemize} \\item The usual expression for the coefficient $\\nu_{\\mathrm{h}}$ of viscosity  due to horizontal turbulence and for the coefficient $D_{\\mathrm{h}}$ of horizontal diffusion, which is of the same order, is, according to \\cite{Z92}, \\begin{equation} D_{\\mathrm{h}} \\simeq \\nu_{\\mathrm{h}} = \\frac{1}{c_{\\mathrm{h}}} r \\;|2V(r) - \\alpha U(r)| \\,\\, , \\label{Zahn92} \\end{equation}%Eq. 1 where $r$ is the appropriately defined eulerian coordinate of the isobar \\cite{MM97}. $V(r)$ is defined by $u_\\theta(r,\\theta)=V(r) {{\\rm d}P_2(\\cos\\theta) \\over {\\rm d}r}$ where $u_\\theta$ is  the horizontal component of the velocity of the meridional currents\\footnote{$V(r)$ can be obtained from $U(r)$ see Eq.~2.10 in \\cite{Z92}}, $\\alpha = \\frac{1}{2} \\frac{d \\ln r^{2} \\Omega} {d \\ln r}$ and $c_{\\mathrm{h}}$ is a constant of  order of unity or smaller. This equation was derived assuming  that the differential rotation on an isobaric surface  is small \\cite{MZ98}. \\item \\cite{Ma99} has derived an expression for the coefficient $D_{\\mathrm{h}}$ of diffusion by horizontal turbulence in rotating stars. He has obtained \\begin{eqnarray} D_{\\mathrm{h}} \\propto \\; r \\; \\left(r \\overline{\\Omega}(r) \\; V \\; \\left[ 2 V - \\alpha U \\right]\\right)^\\frac{1}{3} \\;. \\label{nuh} \\end{eqnarray}%Eq. 19 This expression  can be written  in the usual form  $\\nu_{\\mathrm{h}}= \\frac{1}{3} \\; l \\cdot v$ for a viscosity, where the appropriate velocity $v$ is a geometric mean of  3 relevant velocities: a velocity $(2 V - \\alpha U)$ as in Eq~\\ref{Zahn92} by \\cite{Z92}, the horizontal component $V$ of the meridional circulation, the average local rotational velocity $ r \\overline{\\Omega}(r)$. This rotational velocity is usually much larger than either $U(r)$ or $V(r)$, typically by 6 to 8 orders of a magnitude in an upper Main Sequence star rotating with the average velocity. \\item From torque measurements in the classical Couette-Tayler experiment \\cite{RZ99}, \\cite{Mati04} have found the following expression \\begin{eqnarray} \\nu_{\\mathrm{h}} = \\left({\\beta \\over 10}\\right)^{1/2}  \\; \\left(r^2 \\overline{\\Omega}(r) \\; \\left[r \\left|2 V - \\alpha U \\right|\\right]\\right)^\\frac{1}{2} \\; , \\label{mat} \\end{eqnarray}%Eq. 19 with $\\beta\\approx 1.5 \\times 10^{-5}$ \\cite{RZ99}. \\end{itemize} The horizontal turbulent coupling  favours an essentially constant angular velocity $\\Omega$   on the isobars. This rotation law, constant on shells,  applies to fast as well as to slow rotators. As an approximation, it is often represented by a law of the form $\\Omega = \\Omega (r)$ (\\cite{Z92}; see also \\cite{ES76}). Let us note here that the exact value of the diffusion coefficient $D_h$ is not well known. Indeed the values of the numerical factors intervening in the various expressions shown above may vary to some extent. Since the expression of $D_h$ intervenes in the formulas for $U_r$, for $D_{\\rm shear}$ and $D_{\\rm eff}$ (see below), these uncertainties have some impact on the amplitudes of the transport mechanisms.  \\subsubsection{Shear turbulence and mixing.} In a radiative zone, shear due to differential rotation is likely to be a most efficient mixing process. Indeed shear instability grows on a dynamical timescale that is of the order of the rotation period \\cite{Z92}. The usual criterion for shear instability is the Richardson criterion, which compares the balance between the restoring force of the density gradient and the excess energy present in the differentially rotating layers,  \\begin{equation} Ri = \\frac{N^{2}_{\\mathrm{ad}}}{(0.8836\\ \\Omega\\frac {d\\ln\\Omega}{d\\ln r})^2} < \\frac {1}{4}, \\end{equation}   \\noindent where we have taken the average over an isobar, $r$ is the radius and $N_{\\rm{ad}}$ the Brunt-V\\\"ais\\\"al\\\"a frequency given by  \\begin{eqnarray} N^2_{\\rm{ad}} = \\frac{g \\delta}{H_{P}} \\left[ \\frac{\\varphi} {\\delta} \\nabla_{\\mu} + \\nabla_{ad} - \\nabla_{\\rm{rad}} \\right]. \\end{eqnarray}   \\noindent When thermal dissipation is significant, the restoring force of buoyancy is reduced and the instability occurs more easily. Its timescale is however longer, being the thermal timescale. This case is  referred to as ``secular shear instability''.  The criterion for low Peclet numbers $Pe$ (i.e. of large thermal dissipation, see below) has been considered by \\cite{Za74}, while the cases of general Peclet numbers  $Pe$ have been considered by \\cite{Mae95}, \\cite{MM96}, who give  \\begin{equation} Ri = \\frac{g \\delta}{(0.8836\\ \\Omega\\frac {d\\ln\\Omega}{d\\ln r})^{2} H_{P}} \\left[ \\frac{\\Gamma}{\\Gamma +1} (\\nabla_{ad} -\\nabla) + \\frac{\\varphi}{\\delta} \\nabla_{\\mu} \\right] < \\frac{1}{4} \\end{equation}    \\noindent The quantity $\\Gamma =  Pe/6 $, where the Peclet number $Pe$ is the ratio of the thermal cooling time to the dynamical time, i.e.  $Pe = \\frac{v \\ell}{K}$ where $v$ and $\\ell$ are the characteristic velocity and  length scales, and $K = (4acT^3)/ (3 C_P \\kappa \\rho^2 )$ is the thermal diffusivity. A discussion of shear--driven turbulence by \\cite{Ca98} suggests that the limiting   $Ri$ number may be larger than $\\frac{1}{4}$.  To account for shear transport and diffusion, we need a diffusion coefficient. Amazingly, a great variety of coefficients  $D_{\\mathrm{shear}} = \\frac{1}{3} v \\ell$ have been derived and applied (see a more extended discussion in \\cite{MMV}):  \\begin{enumerate} \\item \\cite{Z92} defines the diffusion coefficient cor\\-responding to the eddies which have the largest $Pe$ number so that the Richardson criterion is just marginally satisfied. However, the effects of the vertical $\\mu$--gradient are not accounted for and the expression only applies to low Peclet numbers. The same has been done by \\cite{MM96}, who considers also the effect of the vertical $\\mu$--gradient, the case of general Peclet numbers and, in addition they account for the coupling due to the fact that the shear also modifies the local thermal gradient. This coefficient has been used by \\cite{MM97} and by \\cite{Den99}. The comparisons  of model results and observations of surface abundances have led many authors to conclude that the $\\mu$--gradients  appear to inhibit  the shear mixing too much with respect to what is required by the observations (\\cite{Cha95a},\\cite{MM97},\\cite{He00}).  \\item Instead of using a gradient $\\nabla_{\\mu}$ in the criterion for shear mixing, \\cite{Cha95a} and \\cite{He00} write  $f_{\\mu} \\nabla_{\\mu}$ with a factor $f_{\\mu} = 0.05$ or even  smaller. This procedure is not satisfactory since it only accounts for a small fraction of the existing $\\mu$--gradients in stars. The problem is that the models depend at least as much (if not more) on $f_{\\mu}$ than on rotation, {\\it i.e.} a change of $f_{\\mu}$ in the allowed range (between 0 and 1) produces as important effects as a change of the initial rotational velocity. This situation has led to two other more physical approaches discussed below. Also \\cite{He00} introduces another factor $f_c$ to adjust the ratio of the transport of the angular momentum and of the chemical elements like  \\cite{Pin89}.  \\item Around the convective core in the region where the $\\mu$--gradient inhibits mixing, there is anyway some turbulence due to both the  horizontal turbulence and to the  semiconvective instability, which is  generally present in massive stars. This situation has led  to the hypothesis  \\cite{Mae97} that the  excess energy in the shear, or a fraction $\\alpha$ of it of the order of unity, is degraded by turbulence on the local thermal timescale. This progressively changes the entropy gradient and consequently the $\\mu$--gradient. This hypothesis leads to  a diffusion coefficient $D_{\\rm shear}$ given by \\begin{eqnarray} D_{\\rm shear} = 4 \\frac{K}{N^{2}_{\\rm{ad}}} \\left[ \\frac{1}{4} \\alpha \\left(0.8836\\  \\Omega \\frac{d\\ln\\Omega}{d\\ln r} \\right)^2 - (\\nabla^{\\prime} -\\nabla) \\right]. \\label{Mae97} \\end{eqnarray} The term $\\nabla^{\\prime} -\\nabla$ in Eq.~\\ref{Mae97} expresses either the stabilizing effect of the thermal gradients in radiative zones or its destabilizing effect in semiconvective zones (if any). When the shear is negligible, $D_{\\rm shear}$ tends toward the diffusion coefficient for semiconvection by \\cite{La83} in semiconvective zones. When the thermal losses are large ($\\nabla^{\\prime} =\\nabla$), it tends toward the value \\begin{equation} D_{\\rm shear} = \\alpha (K/N^2_{\\rm{ad}}) \\left(0.8836\\  \\Omega \\frac{d\\ln\\Omega}{d\\ln r}\\right)^2 , \\label{orig} \\end{equation} given by \\cite{Z92}. Eq.~\\ref{Mae97} is completed by the three following equations expressing the thermal effects \\cite{Mae97} \\begin{eqnarray} D_{\\rm shear} = 2 K \\Gamma \\;\\;\\;\\;\\;\\;\\;\\; \\nabla=\\frac{\\nabla_{rad}+ (\\frac{6 \\Gamma^2}{1+\\Gamma}) \\nabla_{ad}}{1+(\\frac{6 \\Gamma^2}{1+\\Gamma})}, \\label{Mae972} \\end{eqnarray} \\begin{eqnarray} \\nabla^{\\prime} -\\nabla = \\frac{\\Gamma}{\\Gamma +1} (\\nabla_ {\\mathrm{ad}} - \\nabla). \\label{Mae973} \\end{eqnarray} The system of 4 equations given by Eqs.~\\ref{Mae97}, \\ref{Mae972} and \\ref{Mae973} form a coupled system  with 4 unknown quantities $D_{\\rm shear}$, $\\Gamma$, $\\nabla$ and $\\nabla^{\\prime}$. The system is of the third degree in $\\Gamma$. When it is solved   numerically, we find  that as a matter of fact the thermal losses in the shears are rather large in massive stars  and thus that the Peclet number $Pe$ is very small (of the order of 10$^{-3}$  to 10$^{-4}$). For very low Peclet number $Pe =6 \\Gamma$, the differences $(\\nabla^{\\prime} -\\nabla)$ are also very small as shown by Eq.~\\ref{Mae973}. Thus, we conclude that Eq.~\\ref{Mae97} is essentially equivalent, at least in massive stars, to the original Eq.~\\ref{orig} above, as given by \\cite{Z92}. We may suspect that this is not necessarily true in low and intermediate mass stars since there the $Pe$ number may be larger.  \\item \\cite{TZ97} found that the diffusion coefficient for the shears is modified by the horizontal turbulence. The change can be an increase or a decrease of the diffusion coefficient depending on the various parameters, as discussed below. Thus, we have  \\begin{eqnarray} D =  \\frac{ (K + D_{\\mathrm{h}})} {\\left[\\frac{\\varphi}{\\delta} \\nabla_{\\mu}(1+\\frac{K}{D_{\\mathrm{h}}})+ (\\nabla_{\\mathrm{ad}} -\\nabla_{\\mathrm{rad}}) \\right] }\\; \\times \\label{dsh} \\\\[2mm] \\nonumber \\frac{H_{\\mathrm{p}}}{g \\delta} \\; \\left [ \\alpha\\left( 0.8836\\Omega{d\\ln \\Omega \\over d\\ln r} \\right)^2 -4 (\\nabla^{\\prime}  -\\nabla) \\right] \\end{eqnarray}%Eqn 4  \\noindent where $D_{\\mathrm{h}}$ is the coefficient of  horizontal diffusion (cf. \\cite{Z92}). We ignore here the thermal coupling effects discussed by Maeder (\\cite{Mae97}) because they were found to be relatively small and  they increase the numerical complexity. Interestingly, we see that in regions where $\\nabla_{\\mu} \\simeq 0$, Eq.~\\ref{dsh}  leads us to replace $K$ by  $(K+D_{\\mathrm{h}})$ in the usual expression (cf. \\cite{TZ97}), i.e. it reinforces slightly the diffusion in regions which are close to chemical homogeneity. On the contrary, in regions where $\\nabla_{\\mu}$ dominates with respect to $(\\nabla_{\\mathrm{ad}} -\\nabla_{\\mathrm{rad}})$, the transport is proportional to $D_{\\mathrm{h}}$ rather than to $K$, which is quite logical since the diffusion is then determined by  $D_{\\mathrm{h}}$ rather than by thermal effects. The above result shows the importance of the treatment for the meridional circulation, since in turn it determines the size of  $D_{\\mathrm{h}}$ and to some extent the diffusion by shears. \\end{enumerate} Of course, the Reynolds condition $D_{\\rm shear} \\geq \\frac{1}{3} \\nu Re_c$ must be satisfied in order that  the medium is turbulent. The quantity $\\nu$ is the  total viscosity (radiative + mole\\-cular) and $Re_c$  the critical Reynolds number estimated to be around 10 (cf. \\cite{Den99}; \\cite{Z92}). The numerical results indicate that the conditions for the occurrence of turbulence are satisfied.   \\subsubsection{Transport of the angular momentum}  Let us express the rate of change of the angular momentum, ${{\\rm d}{\\mathcal L}\\over{\\rm d}t}$, of the element of mass in the volume ABCD represented in Fig.~\\ref{schema}: $${{\\rm d}{\\mathcal L}\\over{\\rm d}t}={\\bf M},$$ where ${\\bf M}$ is the momentum of the forces acting on the volume element. We assume that angular momentum is transported only through advection (by a velocity field {\\bf U}) and through turbulent diffusion, which may be different in the radial (vertical) and tangential (horizontal) direction. The component of the angular momentum aligned with the rotational axis is equal to \\footnote{The components perpendicular to the rotational axis cancel each other when the integration is performed over $\\varphi$.} $$\\underbrace{\\rho r^2 \\sin\\theta {\\rm d}\\theta {\\rm d}\\varphi {\\rm d} r}_{\\rm Mass\\ of\\ ABCD}\\ \\ \\underbrace{r\\sin\\theta \\Omega}_{\\rm velocity}\\underbrace{r\\sin\\theta ,}_{\\rm distance\\ to\\  axis}$$ where $ \\Omega=\\dot\\varphi$. Since the mass of the volume element ABCD does not change, the rate of change of the angular momentum can be written \\begin{eqnarray} \\rho r^2 \\sin\\theta {\\rm d}\\theta {\\rm d}\\varphi {\\rm d} r {{\\rm d}\\over {\\rm d}t} (r^2 \\sin^2\\theta \\Omega)_{M_r}. \\label{eqn1} \\end{eqnarray} Due to shear, forces apply on the surfaces of the volume element. The force on the surface AB is equal to $$\\underbrace{\\eta_v}_{\\rm vertical\\  viscosity}\\ \\ \\underbrace{r \\sin\\theta {\\partial\\Omega\\over\\partial r}}_{\\rm vertical\\ shear} \\underbrace{r^2\\sin\\theta{\\rm d}\\theta{\\rm d}\\varphi}_{\\rm surface\\  AB}.$$ The component of the momentum of this force along the rotational axis is $$\\underbrace{\\eta_v r^3\\sin^2\\theta {\\partial\\Omega\\over\\partial r}{\\rm d}\\theta{\\rm d}\\varphi}_{\\rm force\\  on\\  AB} \\ \\underbrace{r\\sin\\theta .}_{\\rm distance\\  to\\  axis}$$ The component along the rotational axis of the resultant momentum of the forces acting on AB and CD is equal to \\begin{eqnarray} {\\partial \\over \\partial r}(\\eta_v r^4 \\sin^3\\theta {\\rm d}\\theta{\\rm d}{\\varphi}{\\partial\\Omega\\over\\partial r}){\\rm d}r. \\label{eqn2} \\end{eqnarray} The force on the surface AC due to the tangential shear is equal to $$\\eta_h \\underbrace{r\\sin\\theta {\\partial \\Omega\\over r\\partial\\theta}}_{\\rm tangential\\  shear} \\underbrace{r\\sin\\theta{\\rm d}\\varphi{\\rm d}r }_{\\rm surface\\  AC},$$ where $\\eta_h$ is the horizontal viscosity. The component along the rotational axis of the resultant momentum of the forces acting on AC and BD is equal to \\begin{eqnarray} {\\partial \\over r \\partial \\theta}(\\eta_h r^2\\sin^3\\theta{\\rm d}r{\\rm d}\\varphi {\\partial \\Omega \\over \\partial \\theta})r{\\rm d}\\theta. \\label{eqn3} \\end{eqnarray} Using Eqs.~\\ref{eqn1}, \\ref{eqn2} and \\ref{eqn3}, simplifying by ${\\rm d}r{\\rm d}\\theta{\\rm d}\\varphi$, one obtains the equation for the transport of the angular momentum \\begin{eqnarray} \\rho r^2 \\sin\\theta {{\\rm d}\\over {\\rm d}t} (r^2 \\sin^2\\theta \\Omega)_{M_r}= {\\partial \\over \\partial r}(\\eta_v r^4 \\sin^3\\theta {\\partial\\Omega\\over\\partial r})+ {\\partial \\over \\partial \\theta}(\\eta_h r^2\\sin^3\\theta {\\partial \\Omega \\over \\partial \\theta}). \\label{eqn4} \\end{eqnarray} Setting $\\eta_v=\\rho D_v$ and $\\eta_h=\\rho D_h$ and dividing the left and right member by $r^2\\sin\\theta$, one obtains \\begin{eqnarray} \\rho{{\\rm d}\\over {\\rm d}t} (r^2 \\sin^2\\theta \\Omega)_{M_r}={\\sin^2\\theta \\over r^2} {\\partial \\over \\partial r}(\\rho D_v r^4  {\\partial\\Omega\\over\\partial r})+ {1\\over \\sin\\theta} {\\partial \\over \\partial \\theta}(\\rho D_h \\sin^3\\theta {\\partial \\Omega \\over \\partial \\theta}). \\label{eqn43} \\end{eqnarray} Now, the left--handside term can be written $$\\rho{{\\rm d}\\over {\\rm d}t} (r^2 \\sin^2\\theta \\Omega)_{M_r}={{\\rm d}\\over {\\rm d}t} (\\rho r^2 \\sin^2\\theta \\Omega)_{M_r}-r^2 \\sin^2\\theta \\Omega{{\\rm d}\\rho\\over {\\rm d}t}|_{M_r}.$$ Using the relation between the Lagrangian and Eulerian derivatives, one has \\begin{eqnarray} \\rho{{\\rm d}\\over {\\rm d}t} (r^2 \\sin^2\\theta \\Omega)_{M_r}=\\nonumber \\end{eqnarray} \\begin{eqnarray} {\\partial\\over \\partial t} (\\rho r^2 \\sin^2\\theta \\Omega)_{r} +{\\bf U}\\cdot{\\bf \\nabla}(\\rho r^2 \\sin^2\\theta \\Omega)-r^2 \\sin^2\\theta \\Omega{{\\rm d}\\rho\\over {\\rm d}t}|_{M_r}. \\label{eqn42} \\end{eqnarray}  \\begin{figure}[t] \\begin{center} \\resizebox{5cm}{!}{\\includegraphics[angle=0]{Meynet_G.f1.eps}} \\caption{The momentum of the viscosity forces acting on the element ABCD is derived in the text and the general form of the equation describing the change with time of the angular momentum of this element is deduced. The star rotates around the vertical axis with the angular velocity $\\Omega$; $r$ and $\\theta$ are the radial and colatitude coordinates of point A. } \\label{schema} \\end{center} \\end{figure}% fig. 1 macro wr/NeHCHE   \\noindent Using $${{\\rm d} \\rho \\over {\\rm d}t}|_{M_r}={\\partial \\rho \\over \\partial t}|_{r}+{\\bf U} \\cdot {\\bf \\nabla} \\rho,$$ and the continuity equation $${\\partial \\rho \\over \\partial t}|_{r}=-{\\rm div}(\\rho {\\bf U}),$$ one obtains ${\\rm d}\\rho/{\\rm d}t|_{M_r}+\\rho {\\rm div}{\\bf U}=0$, which incorporated in Eq.~\\ref{eqn42} gives $$\\rho{{\\rm d}\\over {\\rm d}t} (r^2 \\sin^2\\theta \\Omega)_{M_r} ={\\partial \\over \\partial t} (\\rho r^2 \\sin^2\\theta \\Omega)_{r}+{\\bf \\nabla}({\\bf U}\\rho r^2 \\sin^2\\theta \\Omega).$$ Developing the divergence in spherical coordinates and using Eq.~\\ref{eqn43}, one finally obtains the equation describing the transport of the angular momentum (\\cite{MZ98}; \\cite{Mat04}) \\begin{eqnarray} {\\partial\\over \\partial t} (\\rho r^2 \\sin^2\\theta \\Omega)_{r}+{1 \\over r^2}{\\partial \\over \\partial r}(\\rho r^4\\sin^2\\theta w_r\\Omega)+{1 \\over r\\sin\\theta} {\\partial \\over \\partial \\theta}(\\rho r^2\\sin^3\\theta w_{\\theta} \\Omega)= \\nonumber \\\\ {\\sin^2\\theta \\over r^2} {\\partial \\over \\partial r}(\\rho D_v r^4  {\\partial\\Omega\\over\\partial r})+ {1\\over \\sin\\theta} {\\partial \\over \\partial \\theta}(\\rho D_h \\sin^3\\theta {\\partial \\Omega \\over \\partial \\theta}), \\label{eqn5} \\end{eqnarray} where $w_r=U_r+\\dot r$ is the sum of the radial component of the meridional circulation velocity and the velocity of expansion/contraction, and $w_\\theta=U_\\theta$, where $U_\\theta$ is the horizontal component of the meridional circulation velocity. Assuming, as in \\cite{Z92}, that the rotation depends little on latitude due to strong horizontal diffusion, we write $$\\Omega(r,\\theta)=\\bar\\Omega(r)+\\hat \\Omega(r,\\theta),$$ with $\\hat\\Omega \\ll \\bar\\Omega$. The horizontal average $\\bar \\Omega$ is defined as being the angular velocity of a shell rotating like a solid body and having the same angular momentum as the considered actual shell. Thus $$\\bar \\Omega={\\int \\Omega \\sin^3 \\theta{\\rm d} \\theta \\over \\int \\sin^3 \\theta {\\rm d}\\theta}.$$ Any vector field whose Laplacian is nul can be decomposed in spherical harmonics. Thus, the meridional circulation velocity can be written \\cite{Mat04} $${\\bf U}=\\underbrace{\\sum_{l > 0} U_l (r) P_l (\\cos\\theta)}_{u_r} {\\bf e}_r+ \\underbrace{ \\sum_{l > 0} V_l(r) {{\\rm d} P_l (\\cos\\theta) \\over {\\rm d} \\theta}}_{u_\\theta}{\\bf e}_\\theta,$$ where ${\\bf e_r}$ and ${\\bf e_\\theta}$ are unit vectors along the radial and colatitude directions respectively. Multiplying Eq.~\\ref{eqn5} by $\\sin\\theta{\\rm d}\\theta$ and integrating it over $\\theta$ from 0 to $\\pi$, one obtains \\cite{MZ98} \\begin{eqnarray} {\\partial \\over \\partial t}(\\rho r^2 \\bar \\Omega)_r={1 \\over 5 r^2}{\\partial \\over \\partial r}(\\rho r^4 \\bar \\Omega [U_2(r)-5\\dot r]) +{1 \\over r^2}{\\partial \\over \\partial r}\\left(\\rho D_v r^4 {\\partial \\bar \\Omega \\over \\partial r} \\right). \\label{eqn6} \\end{eqnarray} It is interesting to note that only the $l=2$ component of the circulation is able to advect a net amount of angular momentum. As explained in \\cite{Spie92} the higher order components do not contribute to the vertical transport of angular momentum. Note also that the change in radius $\\dot r$ of the given mass shell is included in Eq.~\\ref{eqn6}, which is the Eulerian formulation of the angular momentum transport equation. In its Lagrangian formulation, the variable $r$ is linked to $M_r$ through ${\\rm d}M_r=4\\pi r^2 \\rho {\\rm d}r$, and the equation for the transport of the angular momentum can be written \\begin{eqnarray} \\rho{\\partial \\over \\partial t}(r^2 \\bar \\Omega)_{M_r}={1 \\over 5 r^2}{\\partial \\over \\partial r}(\\rho r^4 \\bar \\Omega U_2(r)) +{1 \\over r^2}{\\partial \\over \\partial r}\\left(\\rho D_v r^4 {\\partial \\bar \\Omega \\over \\partial r} \\right). \\label{eqn7} \\end{eqnarray} The characteristic time associated to the transport of $\\Omega$ by the circulation is \\cite{Z92} \\begin{eqnarray} t_\\Omega\\approx t_{KH} \\left({\\Omega^2 R \\over g_s}\\right)^{-1}, \\label{eqn8} \\end{eqnarray} where $g_s$ is the gravity at the surface and $t_{KH}$ the Kelvin--Helmholtz timescale, which is the characteristic timescale for the change of $r$ in hydrostatic models. From Eq.~\\ref{eqn8}, one sees that $t_\\Omega$ is a few times $t_{KH}$, which itself is much shorter that the Main Sequence lifetime.  %In order to resolve Eq.~\\ref{eqn7}, one needs expressions for  $U_2 (r)$, $D_v$, and $D_h$. %In the following we indicate the general lines of the reasoning and give %the references where more complete mathematical derivations can be found.   %In the above equation $D_v$ is the shear diffusion coefficient whose expression is taken as in \\cite{TZ97}. %The usual estimate of $D_{\\mathrm{h}} = %\\frac{1}{c_{\\mathrm{h}}} r \\;|2V(r) - \\alpha U(r)| $ was given by %\\cite{Z92}. Recent studies suggest that this coefficient %is at least an order of magnitude larger (\\cite{Ma03}; %\\cite{Mati04}).  For shellular rotation, the equation of transport of angular momentum in the vertical direction is in lagrangian coordinates (cf. \\cite{Z92}; \\cite{MZ98})  \\begin{eqnarray} \\lefteqn{\\rho \\frac{d}{d t} \\left( r^2 \\Omega\\right)_{M_r} = } \\nonumber \\\\[2mm] && \\frac{1}{5 r^2}  \\frac{\\partial}{\\partial r} \\left(\\rho r^4 \\Omega U(r) \\right) + \\frac{1}{r^2} \\frac{\\partial}{\\partial r} \\left(\\rho D r^4 \\frac{\\partial \\Omega}{\\partial r} \\right) . \\label{full} \\end{eqnarray}  \\noindent $\\Omega(r)$ is the mean angular velocity at level $r$. The vertical component $u(r,\\theta)$ of the velocity of the meridional circulation at a distance $r$ to the center and at a colatitude $\\theta$ can be written  \\begin{eqnarray} u(r,\\theta)=U(r)P_2(\\cos \\theta), \\label{urt} \\end{eqnarray}  \\noindent where $P_2(\\cos \\theta)$ is the second Legendre polynomial. Only the radial term $U(r)$ appears in Eq.~\\ref{full}. The quantity $D$ is  the  total diffusion coefficient representing the various instabilities considered  and which transport the angular momentum, namely  convection, semiconvection and  shear turbulence. As a matter of fact, a very large diffusion coefficient as in convective regions implies a rotation law which is not far from solid body rotation. In this work, we take $D = D_{\\rm{shear}}$ in radiative zones, since as extra--convective mixing we  consider shear mixing and meridional circulation.  In case the outward transport of the angular momentum by the shear is compensated by an inward transport due to the meridional circulation, we obtain the local conservation of the angular momentum. We call this solution the \\emph{stationary solution}. In this case, $U(r)$ is given by (cf. \\cite{Z92})  \\begin{equation} U(r)= - \\frac{5 D}{\\Omega} \\frac{\\partial \\Omega}{\\partial r} \\; . \\label{stati} \\end{equation}  \\noindent The full solution of Eq.~\\ref{full} taking into account $U(r)$ and $D$ gives the \\emph{non--stationary solution} of the problem. In this case, $\\Omega (r)$ evolves as a result of the various transport processes, according to their appropriate timescales, and in turn differential rotation influences the various above processes. This  produces a feedback and,  thus, a self--consistent solution for the evolution of $\\Omega (r)$ has to be found.  Fig.~\\ref{Ur} shows the evolution of $U(r)$  in a model of a 20 M$_{\\odot}$  star with $Z$ = 0.004 and an initial rotation velocity $v_{\\mathrm{ini}}$ = 300 km s$^{-1}$ \\cite{MMVII}. $U(r)$ is initially positive in the  interior, but progressively  the fraction of the star where  $U(r)$ is negative is growing. This is due to the Gratton--\\\"{O}pik term in Eq. (2), which favors a negative $U(r)$ in the outer layers, when the density decreases. This negative velocity causes an outward transport of the angular momentum, as well as the shears\\footnote{When U is negative, the meridional currents turn anticlockwise, {\\it i.e.} go inwards along directions parallel to the rotational axis and go outwards in directions parallel to the equatorial plane.}.  \\begin{figure}[tb] \\resizebox{\\hsize}{!}{\\includegraphics{MS1279f2.eps}} \\caption{Evolution of $U(r)$ the radial term of the vertical component of the velocity of meridional circulation in a 20 M$_{\\odot}$  star with $Z$ = 0.004 and an initial rotation velocity $v_{\\mathrm{ini}}$ = 300 km s$^{-1}$. $X_c$ is the hydrogen mass fraction at the center. Figure taken from \\cite{MMVII}} \\label{Ur} \\end{figure}%Fig. 2   The transport of angular momentum by circulation has often been treated as a diffusion process (\\cite{ES76}; \\cite{Pin89}; \\cite{He00}). From Eq.~\\ref{full}, we see that the term with $U$ (advection) is functio\\-nally not the same as the term with $D$ (diffusion). Physically advection and diffusion are quite different: diffusion brings a quantity  from where there is a lot to other places where there is little. This is not necessarily the case for advection. A circulation with a positive value of $U(r)$, i.e.\\ rising along the polar axis and descending at the equator, is as a matter of fact making an inward transport of angular momentum. Thus, we see that when this process is treated as a diffusion, like a function of $\\frac{\\partial \\Omega}{\\partial r}$, even the sign of the effect may be wrong.  The expression of $U(r)$ given above (Eq.~\\ref{Umer}) involves derivatives up to the third order, thus Eq.~\\ref{full} is of the fourth order, which makes the system very difficult to solve numerically. In practice, we have applied a Henyey scheme to make the calculations. Eq.~\\ref{full} also implies four boundary conditions. At the stellar surface, we take (cf. \\cite{Ta97})  \\begin{eqnarray} \\frac{\\partial \\Omega}{\\partial r} = 0  \\;\\;\\;\\; \\mathrm{and}  \\;\\;\\; \\; U(r) =0 \\label{BE} \\end{eqnarray}  \\noindent and at the edge of the core we have   \\begin{eqnarray} \\frac{\\partial \\Omega}{\\partial r} = 0  \\; \\; \\; \\; \\mathrm{and}  \\; \\; \\; \\;  \\Omega(r) = \\Omega_{\\mathrm{core}.} \\end{eqnarray}   \\noindent We  assume that the mass lost by stellar winds is just embarking its own angular momentum. This means that we ignore any possible magnetic coupling, as it occurs in low mass stars. It is interesting to mention here, that in case of no viscous, nor magnetic coupling at the stellar surface, {\\it i.e.} with the boundary conditions \\ref{BE}, the integration of Eq.~\\ref{full} gives for an external shell of mass $\\Delta M$ \\cite{Ma99}  \\begin{eqnarray} \\Delta M {d \\over dt} (\\Omega r^2)=-{4\\pi \\over 5} \\rho r^4 \\Omega U(r). \\end{eqnarray}  \\noindent This equation is valid provided the stellar winds are spherically symmetric. When the surface velocity approches the critical velocity, it is likely that there are  anisotropies of the mass loss rates (polar ejection or formation of an equatorial ring) and thus the surface condition should be modified according to the prescriptions of  \\cite{Ma99}.  \\subsubsection{Mixing and transport of the chemical elements.}  A diffusion--advection equation like Eq.~\\ref{full} should normally be used to express the transport of chemical elements. However, if the horizontal  component of the turbulent diffusion $D_{\\rm{h}}$ is large, the vertical advection of the elements  can be treated as a simple diffusion \\cite{Cha92} with a diffusion coefficient $D_{\\rm eff}$. As emphasized by \\cite{Cha92} , this does not apply to the transport of the angular momentum. $D_{\\rm eff}$ is given by  \\begin{equation} D_{\\rm eff} = \\frac{\\mid rU(r) \\mid^2}{30 D_h} \\; , \\label{deff} \\end{equation}  \\noindent where $D_{\\rm{h}}$ is the coefficient of horizontal turbulence. Eq.~\\ref{deff} expresses that the vertical advection of chemical elements is severely inhibited by the strong horizontal turbulence characterized by $D_{\\rm{h}}$. Thus, the change of the mass fraction $X_i$ of the chemical species $i$ is simply  \\begin{eqnarray} \\left( \\frac{dX_i}{dt} \\right)_{M_r} = \\left(\\frac{\\partial  }{\\partial M_r} \\right)_t \\left[ (4\\pi r^2 \\rho)^2 D_{\\rm mix} \\left( \\frac{\\partial X_i} {\\partial M_r}\\right)_t \\right] +  \\left(\\frac{d X_i}{dt} \\right)_{\\rm nucl} . \\label{fullxi} \\end{eqnarray}     \\noindent The second term on the right accounts for composition changes due to nuclear reactions. The coefficient $D_{\\rm mix}$ is the sum $D_{\\rm mix} = D_{\\rm shear}+D_{\\rm eff}$ and $D_{\\rm eff}$ is given by Eq.~\\ref{deff}. The characteristic time for the mixing of chemical elements is therefore $t_{\\rm mix} \\simeq \\frac{R^2}{D_{\\rm mix}} $ and is not given by  $t_{\\rm circ} \\simeq \\frac{R}{U}$, as has been generally considered \\cite{Sch58}. This makes the mixing of the chemical elements much slower, since $D_{\\rm eff}$ is very much reduced. In this context, we recall that several authors have  reduced  by large factors, up to 30 or 100, the coefficient for the transport of the chemical elements, with respect to the transport of the angular momentum, in order to better fit the observed surface compositions (cf. \\cite{He00}). This reduction of the diffusion of the chemical elements is no longer necessary with the more appropriate expression of $D_{\\rm eff}$ given here.  Surface enrichments due to rotation are illustrated in Fig~\\ref{ncp}. The tracks are plotted in the plane ${\\rm (N/C)/(N/C)_{ini}}$ versus $P$ where $P$ is the rotational period in hours. During the evolution the surface is progressively enriched in CNO burning products, {\\it i.e.} is enriched in nitrogen and depleted in carbon. At the same time, the rotational period  increases.  \\begin{figure} \\includegraphics[height=0.755\\textheight]{ncplow.eps} \\caption{Evolutionary tracks in the plane surface N/C ratio, normalized to its initial value, versus the rotational period in hours for different initial mass stars, various initial velocities and for the metallicities $Z$=0.02 and 0.002. Positions of some periods in days are indicated at the bottom of the figure. The dotted tracks never reach the critical limit during the MS phase. The short dashed tracks reach the critical limit during the MS phase. The dividing line between the shaded and non-shaded areas corresponds to the entrance into the phase when the star is at the critical limit during the MS phase. If Be stars are stars rotating at or very near the critical limit, present models would predict that they would lie in the vicinity of this dividing line or above it. Note the different vertical scales used when comparing similar masses at different metallicities. Figure taken from \\cite{EMM}.} \\label{ncp} \\end{figure}  When the effects of the shear and of the meridional circulation compensate each other for the transport of the angular momentum (\\emph{stationary solution}), the value of $U$ entering the expression for $D_{\\rm eff}$ is given by Eq.~\\ref{stati}.  \\subsection{Rotation and mass loss}  We can classify the effects of rotation on mass loss in three categories.  \\begin{enumerate} \\item The structural effects of rotation. \\item The changes brought by rotation on the radiation driven stellar winds. \\item The mass loss induced by rotation at the critical limit. \\end{enumerate}  Let us now consider in turn these various processes.  \\subsubsection{Structural effects of rotation on mass loss.}  Rotation, by changing the chemical structure of the star, modifies its evolution. For instance, moderate rotation at metallicities of the Small Magellanic Cloud (SMC) favors redward evolution in the Hertzsprung-Russel diagram. This behavior can account for the high number of red supergiants observed in the SMC \\cite{MMVII}, an observational fact which is not at all reproduced by non-rotating stellar models.  Now it is well known that the mass loss rates are greater when the star evolves into the red part of the HR diagram, thus in this case, rotation modifies the mass loss indirectly, by changing the evolutionary tracks. The $\\upsilon_{\\rm ini}=0$, 200, 300 and 400 km s$^{-1}$ models lose respectively 0.14, 1.40, 1.71 and 1.93 M$_\\odot$ during the core He-burning phase (see Table~1 in \\cite{MMVII}). The enhancement of the mass lost reflects the longer lifetimes of the red supergiant phase when velocity increases. Note that these numbers were obtained assuming that the same scaling law between mass loss and metallicity as in the MS phase applies during the red supergiant phase. If, during this phase, mass loss comes from continuum-opacity driven wind then the mass-loss rate will not depend on metallicity (see the review by \\cite{vL06}). In that case, the redward evolution favored by rotation would have a greater impact on mass loss than that shown by the computations shown above.   Of course, such a trend cannot continue forever. For instance, at very high rotation, the star will have a homogeneous evolution and will never become a red supergiant \\cite{M87}. In this case, the mass loss will be reduced, although this effect will be somewhat  compensated by other processes: first by the fact that the Main-Sequence lifetime will last longer, second, by the fact that the star will enter the Wolf-Rayet phase (a phase with high mass loss rates) at an earlier stage of its evolution, and third by the fact that the star may encounter the $\\Omega$-limit.      \\subsubsection{Radiation driven stellar winds with rotation.}  The effects of rotation on the radiation driven stellar winds result from the changes brought by rotation to the stellar surface. They induce changes of the morphologies of the stellar winds and increase their intensities. \\vskip 2mm \\noindent\\underline{Stellar wind anisotropies} \\vskip 2mm Naively we would first guess that a rotating star would lose mass preferentially from the equator, where the effective gravity (gravity decreased by the effect of the centrifugal force) is lower. This is probably true when the star reaches the $\\Omega$-limit (i.e. when the equatorial surface velocity is such that the centrifugal acceleration exactly compensates the gravity), but this is not correct when the star is not at the critical limit. Indeed as recalled above, a rotating star has a non uniform surface brightness, and the polar regions are those which have the most powerful radiative flux. Thus one expects in case the opacity does not vary at the surface, that the star will lose mass preferentially along the rotational axis. This is correct for hot stars, for which the dominant source of opacity is electron scattering. In that case the opacity only depends on the mass fraction of hydrogen and does not depends on other physical quantities such as temperature. In that way, rotation induces anisotropies of the winds   (\\cite{MD01};\\cite{DO02}). This is illustrated in the left panel of Fig.~\\ref{ani}. Wind anisotropies have consequences for the angular momentum that a star retains in its interior. Indeed, when mass is lost preferentially along the polar axis, little angular momentum is lost. This process allows loss of mass without too much loss of angular momentum a process which might be important in the context of the evolutionary scenarios leading to Gamma Ray Bursts. Indeed in the framework of the collapsar scenario (\\cite{W93}), one has to accommodate two contradictory requirements: on one side, the progenitor needs to lose mass in order to have its H and He-rich envelope removed at the time of its explosion, and on the other hand it must have retained sufficient angular momentum in its central region to give birth to a fast rotating black-hole.  \\vskip 2mm \\noindent\\underline{Intensities of the stellar winds} \\vskip 2mm The quantity of mass lost through radiatively driven stellar winds is enhanced by rotation. This enhancement can occur through two channels: by reducing the effective gravity at the surface of the star, by increasing the opacity of the outer layers through surface metallicity enhancements due to rotational mixing.  \\begin{itemize}  \\item{\\it reduction of the effective gravity: } The ratio of the mass loss rate of a star with a surface angular velocity $\\Omega$ to that of a non-rotating star, of the same initial mass, metallicity and lying at the same position in the HR diagram is given by \\cite{mm6}  \\begin{equation} \\frac{\\dot{M} (\\Omega)} {\\dot{M} (0)} \\simeq \\frac{\\left( 1  -\\Gamma\\right) ^{\\frac{1}{\\alpha} - 1}} {\\left[ 1 - \\frac{4}{9} (\\frac{v}{v_{\\mathrm{crit, 1}}})^2-\\Gamma \\right] ^{\\frac{1}{\\alpha} - 1}} \\; , \\end{equation} \\noindent where $\\Gamma$ is the electron scattering opacity for a non--rotating star with the same mass and luminosity, $\\alpha$ is a force multiplier \\cite{La95}. The enhancement factor remains modest for stars with luminosity sufficiently far away from the Eddington limit \\cite{mm6}. Typically, $\\frac{\\dot{M} (\\Omega)} {\\dot{M} (0)} \\simeq 1.5$ for main-sequence B--stars. In that case, when the surface velocity approaches the critical limit, the effective gravity decreases and the radiative flux also decreases. Thus the matter becomes less bound when, at the same time, the radiative forces become also weaker. When the stellar luminosity approaches the Eddington limit, the mass loss increases can be much greater, reaching orders of magnitude. This comes from the fact that rotation lowers the maximum luminosity or the Eddington luminosity of a star. Thus it may happen that for a velocity still far from the classical critical limit, the rotationally decreased maximum luminosity becomes equal to the actual luminosity of the star. In that case, strong mass loss ensues and the star is said to have reached the $\\Omega\\Gamma$ limit \\cite{mm6}.   \\item {\\it Effects due to rotational mixing: } During the core helium burning phase, at low metallicity, the surface may be strongly enriched in both H-burning and He-burning products, {\\it i.e.} mainly in nitrogen, carbon and oxygen. Nitrogen is produced by transformation of the carbon and oxygen produced in the He-burning core and which have diffused by rotational mixing in the H-burning shell \\cite{MMVIII}. Part of the carbon and oxygen produced in the He-core also diffuses up to the surface. Thus at the surface, one obtains very high value of the CNO elements. For instance a 60 M$_\\odot$ with Z=$10^{-8}$ and $\\upsilon_{\\rm ini}=800$ km s$^{-1}$ has, at the end of its evolution, a CNO content at the surface equivalent to 1 million times its initial metallicity! In case the usual scaling laws linking the surface metallicity to the mass loss rates is applied, such a the star would lose due to this process more than half of its initial mass.  \\end{itemize}  \\subsubsection{Mass loss induced by rotation}  As recalled above, during the Main-Sequence phase the core contracts and the envelope expands. In case of local conservation of the angular momentum, the core would thus spin faster and faster while the envelope would slow down. In that case, it can be easily shown that the surface velocity would evolve away from the critical velocity (see e.g. \\cite{Vl06}). In models with shellular rotation however an important coupling between the core and the envelope is established through the action of the meridional currents. As a net result, angular momentum is brought from the inner regions to the outer ones. Thus, would the star lose no mass by radiation driven stellar winds (as is the case at low Z), one expects that the surface velocity would increase with time and would approach the critical limit. In contrast, when radiation driven stellar winds are important, the timescale for removing mass and angular momentum at the surface is shorter than the timescale for accelerating the outer layers by the above process and the surface velocity decreases as a function of time. It evolves away from the critical limit. Thus, an interesting situation occurs: when the star loses little mass by radiation driven stellar winds, it has more chance to lose mass by reaching the critical limit. On the other hand, when the star loses mass at a high rate by radiation driven mass loss then it has no chance to reach the critical limit and thus to undergo a mechanical mass loss. This is illustrated in the right panel of Fig.~\\ref{ani}.  \\begin{figure} \\resizebox{5cm}{!}{\\includegraphics{Fig1.eps}} \\hfill %\\hspace{1cm} \\includegraphics[width=2.5in,height=2.5in]{V60Z.eps} \\caption{{\\it Left panel}: Iso-mass loss distribution for a 120 M$_\\odot$ star with Log L/L$_\\odot$=6.0 and T$_{\\rm eff}$ = 30000 K rotating at a fraction 0.8 of critical velocity (figure from \\cite{MD01}). {\\it Right panel}: Evolution of the surface velocities for a 60 M$_ {\\odot}$ star with 3 different initial metallicities. } \\label{ani} \\end{figure}   \\subsubsection{Discussion.}  At this point it is interesting to discuss three aspects of the various effects described above. First what are the main uncertainties affecting them?  Second, what are their relative importance? And finally what are their consequences for the interstellar medium enrichment? \\vskip 2mm \\noindent\\underline{Uncertainties} \\vskip 2mm In addition to the usual uncertainties affecting the radiation driven mass loss rates, the above processes poses three additional problems: \\begin{enumerate}  \\item {\\it What does happen when the CNO content of the surface increases by six orders of magnitude as was obtained in the 60 M$_\\odot$ model described above?} Can we apply the usual scaling law between Z and the mass losses? This is what we have done in our models, but of course this should be studied in more details by stellar winds models. For instance, for WR stars, \\cite{V05} have shown that at $Z=Z_\\odot/30$, 60\\% of the driving is due to CNO elements and only 10\\% to Fe. Here the high CNO surface enhancements result from rotational mixing which enrich the radiative outer region of the star in these elements, but also from the fact that the star evolves to the red part of the HR diagram, making an outer convective zone to appear. This convective zone plays an essential role in dredging up the CNO elements at the surface. Thus what is needed here is the effects on the stellar winds of CNO enhancements in a somewhat red part of the HR diagram (typical effective temperatures of the order of Log T$_{\\rm eff}\\sim$3.8).  \\item {\\it Do stars can reach the critical limit?} For instance, \\cite{BO70} obtain that during pre-main sequence evolution of rapidly rotating massive stars, ``equatorial mass loss'' or ``rotational mass ejection'' never occur (see also \\cite{BO73}). In these models the condition of zero effective gravity is never reached.  However, these authors studied pre-main sequence evolution and made different hypotheses on the transport mechanisms than in the present work. Since they were interested in the radiative contraction phase, they correctly supposed that ``the various instabilities and currents which transport angular momentum have characteristic times much longer than the radiative-contraction time''. This is no longer the case for the Main-Sequence phase. In our models, we consistently accounted for the transport of the angular momentum by the meridional currents and the shear instabilities. A detailed account of the transport mechanisms shows that they are never able to prevent the star from reaching the critical velocity. Another difference between the approach in the work of \\cite{BO70} and ours is that \\cite{BO70} consider another distribution of the angular velocity than in our models. They supposed constant $\\Omega$ on cylindrical surface, while here we adopted, as imposed by the theory of \\cite{Z92}, a ``shellular rotation law''. They resolved the Poisson equation for the gravitational potential, while here we adopted the Roche model. Let us note that the Roche approximation appears justified in the present case, since only the outer layers, containing little mass, are approaching the critical limit. The majority of the stellar mass has a rotation rate much below the critical limit and is thus not strongly deformed by rotation. Thus these differences probably explain why in our models we reach situations where the effective gravity becomes zero.  \\item {\\it What does happen when the surface velocity reaches the critical limit?} Let us first note that when the surface reaches the critical velocity, the energy which is still needed to make equatorial matter to escape from the potential well of the star is still important. This is because the gravity of the system continues of course to be effective all along the path from the surface to the infinity and needs to be overcome. If one estimates the escape velocity from the usual equation energy for a piece of material of mass $m$ at the equator of a body of mass $M$, radius $R$ and rotating at the critical velocity, \\begin{equation} {1 \\over 2}m \\upsilon_{\\rm crit}^2+{1 \\over 2}m \\upsilon_{\\rm esc}^2-{GMm \\over R}=0, \\end{equation} one obtains, using $\\upsilon_{\\rm crit}^2={GM/R}$  that the escape velocity is simply reduced by a factor $1/\\sqrt{2}=0.71$ with respect to the escape velocity from a non-rotating body \\footnote{We suppose here that the vector $\\upsilon_{\\rm esc}$ is normal to the direction of the vector $\\upsilon_{\\rm crit}$.}. Thus the reduction is rather limited and one can wonder if matter will be really lost. A way to overcome this difficulty is to consider the fact that, at the critical limit, the matter will be launched into a keplerian orbit around the star. Thus, probably, when the star reaches the critical limit an equatorial disk is formed like for instance around Be stars. Here we suppose that this disk will eventually dissipate by radiative effects and thus that the material will be lost by the star.  Practically, in the present models, we remove the supercritical layers. This removal of material allows the outer layers to become again subcritical at least until secular evolution will bring again the surface near the critical limit (see \\cite{MEM06} for more details in this process). Secular evolution during the Main-Sequence phase triggers two counteracting effects: on one side, the stellar surface expands. Local conservation of the angular momentum makes the surface to slow down and the surface velocity to evolve away from the critical limit. On the other hand, meridional circulation continuously brings angular momentum to the surface and accelerates the outer layers. This last effect in general overcomes the first one and the star rapidly reach again the critical limit. How much mass is lost by this process? As seen above, the two above processes will maintain the star near the critical limit for most of the time. In the models, we adopt the mass loss rate required to maintain the star at about 95-98\\% of the critical limit. Such a mass loss rate is imposed  as long as the secular evolution brings back the star near the critical limit. In general, during the Main-Sequence phase, once the critical limit is reached, the star remains near this limit for the rest of the Main-Sequence phase. At the end of the Main-Sequence phase, evolution speeds up and the local conservation of the angular momentum overcomes the effects due to meridional currents, the star evolves away from the critical limit and the imposed ``critical'' mass loss is turned off.   %One can wonder also, as was indicated by REF, if a convective zone could not %appear in the equatorial region when the star is near the critical limit. If true, the presence of such a %convective may have important impact on the mass loss triggered by this process \\end{enumerate} ",
        "conclusions": ""
    },
    "0801/0801.2814_arXiv.txt": {
        "abstract": "We use high-resolution {\\sl Hubble Space Telescope} imaging observations of the young ($\\sim 15-25$ Myr-old) star cluster NGC 1818 in the Large Magellanic Cloud to derive an estimate for the binary fraction of F stars ($1.3 < M_\\star/M_\\odot < 1.6$). This study provides the strongest constraints yet on the binary fraction in a young star cluster in a low-metallicity environment ($\\mbox{[Fe/H]} \\sim -0.4$ dex). Employing artificial-star tests, we develop a simple method that can efficiently measure the probabilities of stellar blends and superpositions from the observed stellar catalog. We create synthetic color-magnitude diagrams matching the fundamental parameters of NGC 1818, with different binary fractions and mass-ratio distributions. We find that this method is sensitive to binaries with mass ratios, $q \\ga 0.4$. For binaries with F-star primaries and mass ratios $q > 0.4$, the binary fraction is $\\sim 0.35$. This suggests a total binary fraction for F stars of 0.55 to unity, depending on assumptions about the form of the mass-ratio distribution at low $q$. ",
        "introduction": "The majority of stars are thought to form in binary or multiple\\footnote{For brevity, by `binaries' we generally also mean `multiples', since many stars are found in triple and higher-order multiple systems. We will distinguish the two only when it is important to do so. Note that the analysis in this paper is of {\\em binary} systems.} systems (Goodwin \\& Kroupa 2005; Duch\\^ene et al. 2007; Goodwin et al. 2007), and the initial binary properties of stars place important constraints on star formation and the origin of the stellar initial mass function (IMF; Goodwin et al. 2007, 2008). The majority of stars with masses greater than approximately $0.5 M_\\odot$ are also thought to form in star clusters (Lada \\& Lada 2003), and the binary content of a star cluster plays an important role in both its observational properties and its dynamical evolution (e.g., Kroupa et al. 1999). In addition, many exotic objects observed in star clusters, such as blue stragglers, cataclysmic variables, and X-ray sources, are believed to be related to binary systems. Almost all studies of binarity have been limited to nearby, solar-metallicity populations. However, it might be expected that metallicity (e.g., through its effects on cooling and, hence, on the opacity limit for fragmentation) will play a role in the fragmentation of cores to produce binary systems (Bate 2005; Goodwin et al. 2007).  In general, the most direct way in which to study binary fractions is by examining whether a given star is part of a binary system, on an individual basis. Over the past two decades, the binary fractions of field stars in the solar neighborhood have been studied carefully in this conventional fashion (e.g., Abt 1983 for B stars; Duquennoy \\& Mayor 1991 for G dwarfs; Fischer \\& Marcy 1992 for M dwarfs; Kouwenhoven et al. 2006 for A and B stars; see also Goodwin et al. 2007 for a review). Nearby clusters and associations have also been examined in detail (see Duch\\^ene 1999 and Duch\\^ene et al. 2007 for reviews and comparisons). However, the binary fractions in more distant, massive clusters have not yet been studied thoroughly, because these environments are too crowded and their distances too great, so that their member stars are too faint to be examined individually for binarity. Fortunately, there is an alternative approach, i.e., by means of an artificial-star-test technique, which allows us to estimate the binary fractions in crowded environments. By studying the morphology of their color-magnitude diagrams (CMDs), Romani \\& Weinberg (1991) determined the observed binary fractions in M92 and M30 at $\\la 9$ and $\\la 4$\\%, respectively, for the full populations of cluster stars down to $m_V \\sim 23$ and $\\sim 22$ mag, respectively, based on two-dimensional maximum-likelihood analysis. Rubenstein \\& Bailyn (1997, hereafter RB97) investigated the binary fraction of main-sequence stars with $15.8 < V < 28.4$ mag in the $\\sim 13.5$ Gyr-old (e.g., Pasquini et al. 2007) post-core-collapse Galactic globular cluster NGC~6752. They found a binary fraction of 15--38\\% inside the inner core radius, falling to $\\la 16$\\% at larger radii, with a power-law mass-ratio distribution. For other old globular clusters, Bellazzini et al. (2002) estimated the binary fraction in NGC 288 for stars with $20 < V < 23$ mag (corresponding to masses of $0.54 \\la M_\\star/M_\\odot \\la 0.77$) at 8--38\\% inside the cluster's half-mass radius (and at $< 0.10$ in the outer regions, most likely close to zero), regardless of the actual mass-ratio distribution. Zhao \\& Bailyn (2005) claimed 6--22\\% of main-sequence binaries ($19.2 \\le m_{\\rm F555W} \\le 21.2$ mag) for M3, within the cluster's core radius, and between 1 and 3\\% for stars between 1 and 2 core radii. By applying similar techniques to the post-core-collapse Galactic globular cluster NGC 6397, Cool \\& Bolton (2002) derived a binary fraction of 3\\% for main-sequence stars with primary masses between 0.45 and $0.80 M_\\odot$, for a range of mass ratios. Based on an extrapolation to all mass ratios, they estimated the total main-sequence binary fraction in the cluster core at 5 to 7\\%. Davis et al. (2008), using the method pioneered by Romani \\& Weinberg (1991) in combination with numerical simulations by Hurley et al. (2007), concluded that the outer regions of this cluster (at 1.3--2.8 half-mass radii) exhibit a binary fraction ($1.2 \\pm 0.4$\\%) close to the primordial fraction of $\\sim$1\\% predicted by the simulations, while the inner region has a substantially higher fraction, $5.1 \\pm 1.0$\\%. However, {\\it all clusters thus far studied in this way are old stellar systems} (cf. table 1 in Davis et al. 2008) in which dynamical evolution is expected to have significantly altered the initial binary population.  In this paper, we use accurate photometric observations of the young, low-metallicity star cluster NGC 1818 in the Large Magellanic Cloud, taken with the Wide Field and Planetary Camera-2 (WFPC2) on board the {\\sl Hubble Space Telescope (HST)}, to study its binary fraction. The photometric data and the cluster CMD are discussed in \\S 2. Our newly developed method to correct for stellar blends and superpositions, based on the artificial-star-test technique, is presented in \\S 3. The fitting of the binary fraction is discussed in \\S 4, and \\S 5 contains a further discussion and our conclusions. ",
        "conclusions": "The CMD of NGC 1818, obtained from {\\sl HST} photometry, shows a clearly asymmetric broadening of the main sequence, which implies that this cluster contains a large fraction of binary systems. Using the artificial-star-test method, we estimate that the binary fraction in the mass range from 1.3 to $1.6 M_\\odot$ is $f_{\\rm b} \\sim 0.35$ for systems with an approximately flat mass-ratio distribution for $q>0.4$. This is consistent with a {\\em total} binary fraction of F stars of 0.55 to unity. Elson et al. (1998) found the fraction of roughly equal-mass ($q \\sim 1$) systems in NGC 1818 to be 30--40\\% in the core and 15--25\\% outside the core, which is consistent with our result.  Observationally determined binary fractions are rather scant in general. Binary fractions are a clear function of stellar type, age, and environment. We note that our result is quite close to the fraction of binary dwarfs in the field of the same spectral type, which is a little smaller than the fraction in Sco OB2 (Kouwenhoven 2006) and much higher than in Galactic globular clusters (e.g., RB97; Bellazzini et al. 2002). The most recently published relevant results include the determination by Sana et al. (2008) of the binary fraction of O-type stars in NGC 6231 (at least 63\\%), while Reipurth et al. (2007) reported a visual-binary fraction of late-type stars in the Orion Nebula Cluster of only 8.8\\%.  At 15--25~Myr old, NGC 1818 is several crossing times old, and the binary population would be expected to have been modified by dynamical interactions (see Goodwin et al. 2007; and references therein). In particular, soft (i.e., wide) binaries would be expected to have been destroyed by this age. Therefore, the high binary fraction found for F stars suggests that these binaries are relatively `hard' and able to survive dynamical encounters.\\footnote{We argue that the binaries in NGC 1818 must be `hard,' since the cluster is dynamically old and soft binaries are destroyed in roughly a crossing time (e.g., Parker et al. 2009), so soft binaries cannot be a major component of the binary population (some may form by capture but will be quickly destroyed).}  The relatively flat mass-ratio distribution in NGC~1818 compared to similar-mass stars in the loose association Sco OB2 (Kouwenhoven et al. 2007; $\\sim q^{-0.4}$) may be evidence for a difference in the initial populations (see also Sana et al. 2008 for an alternative mass-ratio distribution biased toward unity). However, it is more likely to be a product of the different dynamical evolution of the two populations. The larger number of encounters suffered by the binary population in NGC~1818 would be expected to disrupt less-bound (i.e., wide and/or low-$q$) systems, and to form more equal-mass systems, leading to a mass-ratio distribution more biased to high $q$.  We conclude that the binary fraction of F stars in the young, low-metallicity LMC cluster NGC~1818 is high and consistent with the field and lower-density clusters. This suggests that, at least among intermediate-mass stars, metallicity down to $\\mbox{[Fe/H]} \\sim -0.4$ dex does not suppress fragmentation and binary formation, and the binarity of these stars is at least as high as at solar metallicity."
    },
    "0801/0801.0811_arXiv.txt": {
        "abstract": "The 6dF Galaxy Survey provides a very large sample of galaxies with reliable measurements of Lick line indices and velocity dispersions. This sample can be used to explore the correlations between mass and stellar population parameters such as age, metallicity and [$\\alpha$/Fe]. Preliminary results from such an analysis are presented here, and show that age and metallicity are significantly anti-correlated for both passive and star-forming galaxies. Passive galaxies have strong correlations between mass and metallicity and between age and $\\alpha$-element over-abundance, which combine to produce a downsizing relation between age and mass. For old passive galaxies, the different trends of $M/L$ with mass and luminosity in different passbands result from the differential effect of the mass--metallicity relation on the luminosities in each passband. Future work with this sample will examine the Fundamental Plane of bulge-dominated galaxies and the influence of environment on relations between stellar population parameters and mass. ",
        "introduction": "The 6dF Galaxy Survey (6dFGS; www.aao.gov.au/local/www/6df; Jones et~al.\\ 2004, 2005) is a redshift and peculiar velocity survey of galaxies in the local universe. The observations for the survey were obtained during 2001--2006 using the UK Schmidt Telescope and the 6dF spectrograph (Watson et~al.\\ 1998). The 6dFGS covers 92\\% of the southern sky with $|b|>10^\\circ$. Its primary sample is drawn from the 2MASS Extended Source Catalog (XSC; Jarrett et~al.\\ 2000) and consists of galaxies with $K_{\\rm tot}<12.65$; for this sample the redshift completeness is 88\\%. It also includes secondary samples complete down to $H<12.95$, $J<13.75$ (from the 2MASS XSC) and $r_F<15.6$, $b_J<16.75$ (from the SuperCosmos Sky Survey; Hambly et~al.\\ 2001). The 6dFGS peculiar velocity survey uses the Fundamental Plane relation to derive distances and velocities for about 15,000 bright early-type galaxies. The 6dFGS database comprises 137k spectra and 124k galaxy redshifts, plus photometry and images. The final data release will be made public in August 2007 (Jones et~al.\\ in prep.; see www-wfau.roe.ac.uk/6dfgs).  As well as its intended purpose as a survey of the structure and motions in the local universe, the 6dFGS also provides a benchmark sample for studying the properties of the low-redshift galaxy population. Here we present some preliminary results from an analysis of the stellar populations in a sample of about 6000 galaxies with high-quality spectra from the 6dFGS Second Data Release (DR2; Jones et~al.\\ 2005). Velocity dispersions ($\\sigma$) have been measured for these galaxies using the cross-correlation technique of Tonry \\& Davis (1979). Comparisons with other high-quality samples show good agreement and imply the 6dFGS dispersions have a median error of 10.9\\% (Campbell et~al., in prep.). ",
        "conclusions": "The 6dFGS provides a large sample of galaxies with spectra of sufficiently high quality and spectral range to allow reliable measurements of both line indices and velocity dispersions. This sample can be used to explore the correlations between mass, age, metallicity and [$\\alpha$/Fe]. Some preliminary results from such an analysis are presented here; a more detailed study will appear in Proctor et~al.\\ (2007, in prep).  We find that age and metallicity are significantly anti-correlated (younger galaxies are more metal-rich) for both passive and star-forming samples. Passive galaxies show the well-known strong correlation between mass and metallicity (more massive galaxies are more metal-rich), but also a strong correlation between age and $\\alpha$-element over-abundance (older galaxies have higher over-abundances). These two effects combine to produce a downsizing relation between age and mass, such that the age of the {\\it youngest} galaxies decreases with decreasing mass. For old ($>$10\\,Gyr) passive galaxies, the different trends of $M/L$ with mass and luminosity in different passbands result from the differential effect of the mass--metallicity relation on the luminosities in each passband.  Future work will examine the Fundamental Plane of bulge-dominated galaxies and the influence of environment on relations between stellar population parameters and mass."
    },
    "0801/0801.3215_arXiv.txt": {
        "abstract": "Knowledge of the structure of the coronal magnetic field is important for our understanding of many solar activity phenomena, e.g. flares and CMEs. However, the direct measurement of coronal magnetic fields is not possible with present methods, and therefore the coronal field has to be extrapolated from photospheric measurements. Due to the low plasma beta the coronal magnetic field can usually be assumed to be approximately force free, with electric currents flowing along the magnetic field lines. There are both observational and theoretical reasons which suggest that at least prior to an eruption the coronal magnetic field is in a nonlinear force free state. Unfortunately the computation of nonlinear force free fields is way more difficult than potential or linear force free fields and analytic solutions are not generally available. We discuss several methods which have been proposed to compute nonlinear force free fields and focus particularly on an optimization method which has been suggested recently. We compare the numerical performance of a newly developed numerical code based on the optimization method with the performance of another code based on an MHD relaxation method if both codes are applied to the reconstruction of a semi-analytic nonlinear force-free solution. The optimization method has also been tested for cases where we add random noise to the perfect boundary conditions of the analytic solution, in this way mimicking the more realistic case where the boundary conditions are given by vector magnetogram data. We find that the convergence properties of the optimization method are affected by adding noise to the boundary data and we discuss possibilities to overcome this difficulty. ",
        "introduction": "\\label{sec:intro}  The magnetic field is the most important quantity for understanding the plasma structure of the solar corona and the activity phenomena it shows. Unfortunately a systematic direct measurement of the coronal magnetic field is not feasible with the observational methods presently available. Therefore the extrapolation of the magnetic field into the corona from measurements of the magnetic field at the photospheric level is extremely important.  It is generally assumed that the magnetic pressure in the corona is much higher than the plasma pressure (small plasma $\\beta$) and that therefore the magnetic field is nearly force-free \\citep[for a critical view of this assumption see][]{gary01}. The extrapolation methods based on this assumption include potential field extrapolation \\citep[e.g.,][]{schmidt64,semel67}, linear force-free field extrapolation \\citep[e.g.,][]{chiu:hilton77,seehafer78,seehafer82,semel88} and nonlinear force-free field extrapolation \\citep[e.g.,][]{amari:etal97}. Methods for non-force-free field extrapolation have also been developed \\citep{petrie:neukirch00} but are not used routinely.  Whereas potential and linear force-free fields can be used as a first step to model the general structure of magnetic fields in the solar corona, the use of non-linear force free fields is essential to understand eruptive phenomena as there are both observational and theoretical reasons which suggest that the pre-eruptive magnetic fields are non-linear force-free fields (see below). The calculation of nonlinear force free fields is complicated by the intrinsic nonlinearity of the underlying mathematical problem. Another problem which becomes especially important when nonlinear force free field are used for magnetic field extrapolation is the correct formulation of the problem with respect to boundary values. The available magnetograph observation provide either the line-of-sight magnetic field ($B_{los}$), which is sufficient for potential and linear force free fields, or all three components of the photospheric magnetic field. Only the latter information is sufficient to determine non-linear force free fields completely. There are, however, tremendous difficulties to be overcome both in extracting the necessary information about the magnetic field components on the boundary from the data in an unambiguous way and in the use of this information in the computation of the corresponding magnetic field. Even though a lot of research has been devoted to these problems in the past so far none of the proposed methods has been found to be outstanding in combining simplicity, usability and robustness.  The purpose of the present contribution is to assess some of the methods presently available for calculating non-linear force-free fields, with a view to implementing them into a general extrapolation code which also takes stereoscopic information into account in the extrapolation process. A first version of such a method based on linear force free extrapolation has recently been proposed \\citep{twtn02}. As stated in \\citet{twtn02} and bearing in mind the remarks above, it is would be very desirable to generalize the extrapolation method to nonlinear force free fields.  As a first step in this direction we investigate the robustness of the optimization method recently proposed by \\citet{wheatland00}, and compare its performance with another method based on MHD relaxation. The paper is organized as follows. In Sect. \\ref{sec:basic} we discuss the basic equations and background theory. Section \\ref{sec:methods} describes various methods which have been proposed for nonlinear force free field extrapolation. Out of these we focus on the optimisation method and, for the case of analytically given boundary data, on the MHD relaxation method. These methods are tested using an analytical solution in Sect. \\ref{sec:tests}. The optimisation method is also subjected to tests with more general boundary conditions where we add noise to the given analytical boundary data. The paper closes with the conclusions in Sect. \\ref{sec:end}. ",
        "conclusions": "\\label{sec:end}  In the present paper we have presented an assessment of the properties of an optimization method which has recently been proposed for the calculation of nonlinear force-free fields from given boundary data \\citep{wheatland00}. One part of the assessment was a comparison the performance of the optimization method with the performance of an MHD relaxation method. For both methods new parallelized codes have been developed. Both methods have been applied to finding a known semi-analytic solution \\citep{lowandlou} from a given non-equilibrium initial condition. Both methods converge to the exact solution, but the optimization method has higher accuracy. The MHD relaxation method is only applied to this case since in its present numerical implementation, the boundary conditions could give rise to inconsistencies.  To simulate a more realistic situation which is closer to working with observational data from vector magnetograms we added noise to the boundary conditions. We have only applied the optimization method to this type of problem and found that already a relatively small noise level of 10 \\% can affect the convergence of the method quite considerably. More work is needed to see how this difficulty can be overcome. We also intend to generalize the relaxation method in such way that is able of coping with more realistic boundary data.  The ultimate aim for the future is to apply these methods to data from vector magnetograms. Unfortunately vector magnetograms are less accurate than line-of-sight magnetograms (e.g. from MDI on SOHO). Therefore, in the light of our results in the cases where noise was added to the boundary conditions it will be interesting to see whether and how quickly the methods converge. We also plan to use stereoscopic information as a further constraint in the reconstruction process. This  will become especially important for the analysis of data from the STEREO mission. A method which is based on linear force free fields has recently been proposed by \\citet{twtn02}, but linear force free fields seem to be too restrictive to describe coronal phenomena appropriately. Therefore, the natural next step will be the use of nonlinear force free fields in such a code.       \\balance %"
    },
    "0801/0801.3679_arXiv.txt": {
        "abstract": "The study of young star cluster (YSC) systems, preferentially in starburst and merging galaxies, has seen great interest in the recent past, as it provides important input to models of star formation. However, even some basic properties (like the luminosity function [LF]) of YSC systems are still under debate. Here we study the photometric properties of the YSC system in the nearest major merger system, the Antennae galaxies. We find evidence for the existence of a statistically significant turnover in the LF. ",
        "introduction": "\\label{s:intro}  Star clusters (SCs) form nearly instantaneously through the collapse of giant molecular gas clouds. Hence, all stars within a SC are approximately coeval, share the same chemical composition, and therefore represent a simple stellar population. A small number of parameters, in particular their initial chemical composition and initial stellar mass function, are enough to describe their colour and luminosity evolution on the basis of a given set of stellar evolutionary tracks or isochrones (e.g. \\citealt{1999ApJS..123....3L,2003A&A...401.1063A,2003MNRAS.344.1000B}). Therefore, observed spectrophotometric properties of SCs are relatively easy and straightforward to interpret.  SC formation is a major mode of all star formation, and possibly even the dominant mode in strong starbursts triggered in gas-rich galaxy mergers (e.g., \\citealt{1995Natur.375..742M,2003NewA....8..155D}). In addition, as SCs inherit and conserve the chemical abundances at the place and time of their birth up to old ages, they are excellent tracers of their parent galaxy's properties in terms of star formation and chemical enrichment history.  One of the most basic and commonly used diagnostics to explore the properties of entire SC {\\sl systems} is their LF. While for old globular cluster systems the Gaussian shape of their LFs (and indeed mass functions) is well established (see e.g. \\citealt{1998gcs..book.....A,1991ARA&A..29..543H}), the situation for young SCs is still under discussion. While mainly LFs consistent with a power-law, resembling the MF of nearby molecular cloud cores, are quoted for YSC systems ranging from Galactic open cluster to YSCs in major mergers (e.g. \\citealt{1984AJ.....89.1822V,2003AJ....126.1836H,1998AJ....116.2206S,1999AJ....118.1551W}, see \\citealt{2003MNRAS.343.1285D} for a recent compilation), some studies find deviations from a power-law or direct evidence for Gaussian distributions (\\citealt{2006MNRAS.366..295D,1999A&A...342L..25F,2004ApJ...613L.121G}).  To contribute to our understanding of this astrophysical question, which ties directly to the fundamental physical conditions of star formation in general, we performed an analysis of the young ($\\sim 0 - 100$ Myr) SC system formed during the, still ongoing, merging process in the nearest major merger of two giant gas-rich spiral galaxies, the Antennae galaxies, i.e. NGC 4038/39. ",
        "conclusions": ""
    },
    "0801/0801.4573_arXiv.txt": {
        "abstract": "The structure of the solar corona is dominated by the magnetic field  because the magnetic pressure is about four orders of magnitude higher than the plasma pressure. Due to the high conductivity the emitting coronal plasma (visible e.g. in SOHO/EIT) outlines the magnetic field lines. The gradient of the emitting plasma structures is significantly lower parallel to the magnetic field lines than in the perpendicular direction. Consequently information regarding the coronal magnetic field can be used for the interpretation of coronal plasma structures. We extrapolate the coronal magnetic field from photospheric magnetic field measurements into the corona. The extrapolation method depends on assumptions regarding coronal currents, e.g. potential fields (current free) or force-free fields (current parallel to magnetic field).  As a next step we project the reconstructed 3D magnetic field lines on an EIT-image and compare with the emitting plasma structures. Coronal loops are identified as closed magnetic field lines with a high emissivity in EIT and a small gradient of the emissivity along the magnetic field. ",
        "introduction": "\\label{sec1} The automatic identification of features in solar images provide an increasingly helpful tool for the analysis of solar physics data. Previous methods of solar feature identification used, e.g, contrast, intensity and contiguity criteria in Ca II K images to identify structures like sunspots and plage \\cite{1998ApJ...496..998W,2001SoPh..202...53P}. Here we are mainly interested in identifying closed loop structures in active regions. Coronal plasma loops are optical thin and coronal images, e.g. in X-rays or EUV radiation, provide only a 2D projection of the 3D structures. The line-of-sight integrated character of these images complicates their interpretation. The corresponding problem that loop structures visible in images are often a superposition of individual loops occurs particularly in high resolution TRACE images \\cite{1999SoPh..187..261S, 2004ApJ...615..512S}, although occasionally individual loops are clearly visible over a large fraction of their length \\cite{2000ApJ...541.1059A}. To get insights regarding the structure of the coronal plasma it is important to consider the distribution of the plasma $\\beta$ with height. The plasma $\\beta$ varies from $\\beta > 1$ in the photosphere over $\\beta \\ll 1$ in the mid-corona to $\\beta > 1$ in the upper corona (See \\citeauthor{2001SoPh..203...71G}(\\citeyear{2001SoPh..203...71G}) for details). For the interpretation of coronal loops it is helpful that they are located in the low $\\beta$ corona and consequently their structure is dominated by the coronal magnetic field. In the following we will therefore discuss how information about the magnetic field can be used for the analysis of coronal images, e.g., from SOHO/EIT, SOHO/SUMER, or in future from the two STEREO spacecrafts.   The strongest emissivity in  EUV-images originates from closed magnetic loops in active regions. Due to the high conductivity the coronal plasma becomes trapped on close magnetic field lines and consequently the emitting plasma outlines the magnetic field \\cite{1975SoPh...44...83P,1977SoPh...52..397S,1989SSRv...51...11S}. This feature has previously been used (in combination with line of sight photospheric magnetograms) to compute free parameters (e.g., the linear force-free parameter $\\alpha$) in coronal magnetic field models. \\citeauthor{2004ApJ...615..512S}(\\citeyear{2004ApJ...615..512S}) used global potential fields and scaling laws to fill the coronal magnetic field lines with plasma and computed so called synthetic EIT images from their magnetic field model, which they compared with real EIT observations. While potential fields give a rough impression regarding the global structure of the solar corona (e.g. the location of coronal holes and active regions, \\citeauthor{2002JGRA.107lSSH13N} \\citeyear{2002JGRA.107lSSH13N}) they often fail to describe magnetic fields in active regions accurately. Here we use different magnetic field models (potential, force-free fields and coronal magnetic fields reconstructed directly from measurements) in combination with coronal plasma images to identify coronal structures, where we concentrate on closed loops in an active region.  While the space in the corona is basically filled with magnetic field, only some fraction of all field lines shows a strong emissivity. Our aim is to identify out of the many (theoretical infinite) number of field lines the ones which are accompanied by a strong plasma radiation. \\citeauthor{1999ApJ...515..842A} (\\citeyear{1999ApJ...515..842A}) picked out individual loops from two EIT images and applied a geometrical fit (assumption of circular loops) and dynamic stereoscopy to draw conclusions regarding the 3D structure of the loops. \\citeauthor{2002SoPh..208..233W} (\\citeyear{2002SoPh..208..233W}) used these circular 3D loops and a MDI magnetogram to compute the optimal linear force-free field consistent with these observations. The method was extended by \\citeauthor{2003SoPh..218...29C} (\\citeyear{2003SoPh..218...29C}) to work with 2D coronal images, in which an individual loop  is visible in coronal images along it's entire length (including the approximate location of both footpoints). An alternative approach was taken by \\citeauthor{2004A&A...428..629M} (\\citeyear{2004A&A...428..629M}), who identified several loops or (large) parts of loops in EIT-images (by manual tracing). The purpose of these papers was not feature identification, but to use coronal images and line of sight magnetograms to compute the coronal magnetic field within a linear force-free model. A manual selection of the loops was necessary.  The aim of this work is to automize these procedures and allow more general magnetic field models. We use linear and non-linear force-free magnetic field models and also give an example with measured magnetic fields. In the special case of linear force-free fields our method automatically calculates the optimal linear force-free parameter $\\alpha$.  In this sense the method is a further development for computing linear force-free fields from line-of sight magnetograms and coronal images. We outline the paper as follows: In section \\ref{sec2} we outline the loop identification method and in section \\ref{sec3} we describe how our method finds the optimal values for magnetic field models with free parameters. Section \\ref{sec4} contains an example where we tested our algorithm with different magnetic field models. Finally we draw conclusions and give an outlook for future work in section \\ref{sec5}. ",
        "conclusions": "\\label{sec5} Within this work we have developed a method for the automatic identification of coronal loops. The method uses a suitable coronal magnetic field model, e.g. from a force-free extrapolation, and compares the projection of 3D magnetic field lines with structures on images, e.g. from SOHO/EIT. As a result coronal structures are represented by 1D curves. A further use of the method is planned in particular for the analysis of data from the two STEREO spacecrafts. The reconstruction of curves is attractive for the STEREO mission for two reasons. Firstly structures on the solar surface are dominated by thin flux tubes which often can well be approximated by idealized curves. Secondly, the stereoscopic reconstruction problem for 1D objects is well posed. Formally, a projected curve in an image can be projected back along the view direction giving rise to a 3D solution surface on which the 3D curve must lie. From two images, we obtain two such surfaces and their intersection must be the solution curve. This is different for a point-like object. Its projection yields two coordinate parameters in two images while only 3 spatial coordinates are required. Formally, the triangulation of point-like objects is therefore over-determined. Conversely, if we want to solve for optically thin, curved surfaces we have an underdetermined problem because the bright limb visible in the two images might project back to different parts of the surface. It is therefore useful to identify 1D objects (loops) out of 2D images for a stereoscopic reconstruction by e.g. tie-point-like methods \\cite{2004AGUFMSH21B0422H}. Let us remark that the structure identification method by magnetic field lines provides already a 3D structure of the loop which might be further refined by stereoscopic methods. In particular, the advent of the STEREO mission has produced an increased interest of the solar physics community in stereoscopic (\\citeauthor{1999ApJ...515..842A},\\citeyear{1999ApJ...515..842A},\\citeyear{2000ApJ...531.1129A} and tomographic techniques \\cite{1999JGR...104.9727Z,2000ApJ...530.1026F,2003SoPh..214..287W} used in other branches of physics, image and information sciences. However, unlike many of the ``experiments'' conducted in the latter disciplines, the STEREO mission does not provide sufficient control over many of the parameters which determine the conditions for an optimal 3D reconstruction. Tomographic methods will in particular be useful in situations where individual loops cannot be identified due to the superposition of many loops due to the line-of-sight integrated character of 2D images."
    },
    "0801/0801.3165_arXiv.txt": {
        "abstract": "{ We argue that X-ray flaring variability observed in the transient X-ray pulsar A0535+26 is due to low-mode magnetospheric instability. This instability develops at the onset of accretion, in the thin boundary layer between the accretion disk and neutron star magnetosphere. As a result, the matter collected in the boundary layer can rapidly fall onto the NS surface close to the magnetic poles, but not exactly along the field lines by which the stationary accretion proceeds. This explains the shift in cyclotron line energy measured using RXTE data in a pre-outburst spike, with respect to the line energy observed during the main outburst. Furthermore, the instability can account for the difference in pulse profiles, and their energy evolution that is different in the pre-outburst flare and main outburst. ",
        "introduction": "The transient X-ray binary pulsar A0535+26 was discovered by the \\textit{Ariel V} satellite during a giant outburst (\\cite{coe:1975, rosenberg:1975}). The neutron star (NS) rotates with a period $P\\simeq 103$ s and is in an eccentric ($e=0.47$) orbit about a massive O9.7IIIe star  with an orbital period of $P_{orb}\\sim 110$ d (\\cite{coe:2006}). The X-ray activity of this transient source is complicated: several giant outbursts were observed in 1975, 1980, 1983, 1989, 1994 and 2005, with weaker outbursts sometimes observed at successive periastron passages. This type of X-ray activity is common for Be/X-ray binary systems, which represent the largest subclass of massive X-ray binaries. The giant (type II, according to \\cite{stella:1986}) outbursts can take place at any orbital phase and most probably are triggered by the enhanced activity of the optical star (\\cite{coe:2006}). More regular, weaker type I outbursts are associated with enhanced interaction of the neutron star with a circumstellar disk surrounding  the Be-companion (\\cite{okazaki:2001}). A0535+26 did not show outbursting activity between 1994 and 2005. The last giant outburst of A0535+26 occurred after a prolonged quiescence period in May-June 2005 (\\cite{tueller:2005}). In September 2005, the subsequent normal type I outburst was extensively studied by RXTE and INTEGRAL (\\cite{caballero:2007}, 2008). These observations revealed somewhat unexpected features, which we summarize below.  \\textit{Feature 1}. A short (fraction of a day) pre-outburst X-ray spike, with a luminosity comparable to the outburst maximum, is observed during the rise of the outburst.  \\textit{Feature 2}. The NS pulse period at the onset of the outburst, remained constant within the margins of error, but rapidly decreased after the periastron passage.  \\textit{Feature 3}. RXTE data shows that the energy of the cyclotron resonance scattering feature (CRSF) during the spike at 90\\% confidence is $52.0^{+1.6}_{-1.4}$ keV (\\cite{caballero:2008}), which is notably higher than the value $46.1\\pm 0.5$~keV measured during the main outburst at the same level of X-ray luminosity. The latter value is consistent with INTEGRAL measurements (\\cite{caballero:2007}).  \\textit{Feature 4}. The CRSF energy in the main outburst (within the margins of error) is independent of the X-ray luminosity.  \\textit{Feature 5}. The X-ray pulse profiles during the spike are different from the pulse profiles observed during the main outburst. Pulse profiles during the outburst show an abrupt change at energies above the cyclotron resonance, while those in the pre-outburst spike do not.  The additional inspection of the \\textit{Swift} BAT (15-50 keV) light curve of the source (Fig. 1) reveals that the pre-outburst spike observed by RXTE is only one of a collection of flares, with the characteristic time up to a few times $10^4$ s, appearing at the rising part of the outburst. The flux evolution then proceeds more smoothly  during the spin-up of the NS (\\textit{Feature 6}), which begins close to periastron passage. A similar flaring activity was observed during the onset of the preceding giant outburst, at the same flux level, in April 2005, and at the next periastron passage in December 2005.  \\begin{figure*} \\resizebox{\\hsize}{!}{\\includegraphics{f1_9277.ps}} \\caption{The Swift/BAT light curve of A0535+26 during the normal outburst in August/September 2005. The smooth solid line shows the pulse period development as measured by RXTE (from Caballero et al. 2007b). The dashed line shows the time of periastron passage.} \\label{batlc} \\end{figure*}  In this Letter we argue that the flaring activity and pre-outburst spikes are most likely due to magnetospheric instability. This would take place at the onset of a smooth increase in the mass-accretion rate through the accretion disk surrounding the NS, close to periastron passage. This instability would cause plasma that had accumulated in the thin boundary layer, between the accretion disk and NS magnetosphere, to fall onto the NS surface as large blobs. These blobs would be channeled to regions close to the NS magnetic poles by different magnetic field lines than those guiding the accretion flow smoothly increasing from the disk. ",
        "conclusions": "The magnetospheric instability model proposed for the pre-outburst flares in A0535+26 is generic and can be applied to other transients. The prerequisite, however, is that the source must be on the verge of low-mode instability, which depends on the NS magnetic field, the spin period, the accretion rate and possibly other parameters (e.g. the misalignment of the magnetic dipole and/or NS spin axis relative to the orbital angular momentum). It is possible that the transition from propeller to accretion stage can occur without a strong KS instability being observed. The model predicts changes in pulse profiles and in the cyclotron line energy during the short flares, as observed in A0535+26 (\\cite{caballero:2008}), which can be compared with observations of other sources.  Are there other possible explanations for the observed pre-outburst X-ray spikes?  The SPH modeling of an accretion disk surrounding the NS (Hayasaki \\& Okazaki, 2006) reproduces normal outbursts at successive periastron passages. In some instances the modeling shows a single peak preceding outburst maximum. The accretion disk is formed from the Roche lobe overflow of the coplanar circumstellar disk surrounding the Be companion star. The important feature of these SPH simulations is the transient, one-armed spiral structure of the accretion disk, induced by a phase-dependent mass accretion. Such accretion disks can be responsible for mass transfer enhancement close to periastron. The model by Hayasaki \\& Okazaki, however, ignores the disk-magnetosphere interaction critical to transient accretion onto magnetized NS, and can not reproduce the observed flaring activity.   Spruit \\& Taam (1993) also found a viscous, disk-magnetosphere instability that was associated with the magnetospheric boundary around the corotation radius. This instability results in a cyclic enhancement of the mass accretion rate on the viscous time scale at the magnetospheric boundary. While this timescale can be as short as $10^4$ s for the corotation radius $10^9$ cm, the chaotic behaviour of flares observed in A0535+26 and, most importantly, other features (different pulse profiles and CRSF energy, and absence of the flaring during the NS spin-up phase) imply that this model can not provide the correct explanation.  Our model for short flares observed at the onset of outbursts in the transient X-ray pulsar A0535+26 is based on magnetospheric instability. It succesfully explains all features of the outburst observed during  August-September 2005 by RXTE and INTEGRAL, and makes clear predictions for future observations. The present-day accuracy of cyclotron line measurements in accreting neutron stars, precise timing analysis and evolution of X-ray pulse shapes with luminosity at different energies, have became the working tools to probe the non-stationary accretion onto magnetized neutron stars."
    },
    "0801/0801.4567_arXiv.txt": {
        "abstract": "{Isolated cooling neutron stars with thermal X-ray emission remain rarely detected objects despite many searches investigating the ROSAT data.} {We simulate the population of close-by young cooling neutron stars to explain the current observational results. Given the inhomogeneity of the neutron star distribution on the sky it is particularly interesting to identify promising sky regions with respect to on-going and future searches.} {Applying a population synthesis model the inhomogeneity of the progenitor distribution and the inhomogeneity of the X-ray absorbing interstellar medium are considered for the first time. The total number of observable neutron stars is derived with respect to ROSAT count rates. In addition, we present sky maps of neutron star locations and discuss age and distance distributions of the simulated neutron stars. Implications for future searches are discussed.} {With our advanced  model we can successfully explain the observed log~N~--~log~S distribution of close-by neutron stars. Cooling neutron stars will be most abundant in the directions of rich OB associations. New candidates are expected to be identified behind the Gould Belt, in particular in the  Cygnus-Cepheus region. They are expected to be on average  younger and then hotter than the known population of  isolated cooling neutron stars. In addition, we propose to use data on runaway stars to search for more radio-quiet cooling neutron stars. } {} ",
        "introduction": "\\label{intro}  More than 10 years after the discovery of its brightest member RX J1856-3754 (Walter et al. 1996), the small group of seven radio-quiet isolated neutron stars (NSs) detected by the ROSAT satellite  have gained an important place in the rich zoo of compact objects. Together with Geminga and several close-by young radio pulsars, these objects form the local population of cooling NSs. Studies of this group of sources already provided a wealth of information on NSs physics (see e.g. \\citealt{Haberl2007,Page2007,Zane2007} for recent reviews).  Since 2001 the number of known close-by radio-quiet NSs has not been growing despite all attempts to identify (mainly in the ROSAT data) new candidates. Partly this is due to the fact that all these searches are blind, i.e. not necessarily targeted to promising sky regions. To advance the identification of new near-by cooling NSs it is necessary to perform a realistic modeling of this population. The seven ROSAT radio-quiet NSs form the main part of the known local population of cooling NSs. For these seven sources we use the term {\\it Magnificent Seven}. There is still no general term to address the whole population to which these sources belong (terms such as XDINS -- X-ray dim isolated NSs, DINS -- dim isolated NSs, RINS -- ROSAT (or radio quiet) isolated NSs,  are used sometimes by different authors). In this paper we use the term ICoNSs (Isolated Cooling NSs). In our opinion, this variant better reflects the nature of these sources than, for example, XDINS, as many of them are not dim in X-rays. Finally, we call {\\it coolers} all NSs, whose residual  thermal emission can be detected. Besides the radio-quiet NSs, such sources can be also known as normal radio pulsars, or they may demonstrate some kind of radio emission or even $\\gamma$-ray emission distinct from classical pulsars. In this study we do not care if an object shows some activity in addition to thermal surface emission. The only crucial point is the detection of X-ray emission related to the cooling of initially hot compact objects.  The population synthesis (PS) approach is used here to investigate the population of close-by young cooling NSs. There are two main versions of this method \\citep{Fritze-v.Alvensleben00}. One, which sometimes is also called {\\it empirical population synthesis}, is a kind of top-down approach, when objects, for which only integral characteristics are available, are studied. Numerical models are applied to reproduce observed properties of such sources by using different sets of subpopulations which constitute an object class. For example, the stellar content of a galaxy can be derived by modeling its integral spectra.  Here, we focus on the second -- bottom-up -- approach, often called {\\it evolutionary synthesis}. In this case the evolution of a population of objects is modelled from their birth based on defined initial distributions of parameters and some equations which describe evolution of these distributions in time. The main reasons to use this second technique are the following. At first, if initial parameters or/and evolutionary laws for some kind of sources are not well known, then this kind of PS  can be applied to test hypotheses about these uncertain properties by comparison between modeled and observed populations. Determination of the initial parameters of radio pulsars can be a good example here \\citep{faucher2006}. Next, even if the initial distributions and their evolution are known, but at the moment just a small part of the population -- a tip of an iceberg -- is observed, then PS calculations can be used to predict properties of the unobserved faint part of the population, and to plan a search strategy to identify new members. A more detailed review of the PS technique can be found, for example, in \\citet{pp2005}.  In this paper an upgraded population synthesis model is discussed for the population of close-by ($< 3$~kpc) isolated NSs which can be observed via their thermal emission in soft X-rays. Previously, our models were applied to confirm the link between the Magnificent Seven and the Gould Belt \\citep{p03} and to test theories of thermal evolution of NSs \\citep{pgtb04}. The major interest of the present study is to understand how to find more objects of this type. After  the description of the PS model ingredients in Sec.~\\ref{model}, we present our results -- $\\log$~N~-~$\\log$~S-curves, age and distance diagrams as well as sky maps of expected NS locations -- in Sec.~\\ref{results}, and discuss possible applications in Sec.~\\ref{search}. An outlook concerning the search for ICoNSs is given in Sec.~\\ref{outlook} Finally, we give our conclusions in Sec.~\\ref{concl}. ",
        "conclusions": "\\label{concl}  In this paper we presented our new more advanced model for population synthesis of close-by cooling NSs. The two, sligthly different, mass distributions we consider result in nearly the same observable number of NSs. Detailed treatment of the initial spatial distribution of NS progenitors and a detailed ISM structure up to 3 kpc allow us to discuss the strategy to look for new ICoNS. Our main results in this respect are the following: new candidates are expected to be identified behind the Gould Belt, in directions to rich OB associations, in particular in the  Cygnus-Cepheus region; new candidates, on average, are expected to be hotter than the known population of ICoNS. Besides the usual approach (looking for soft X-ray sources), the search in 'empty' $\\gamma$-ray error boxes or among run-away OB stars may yield new X-ray thermally emitting  NS candidates."
    },
    "0801/0801.1668_arXiv.txt": {
        "abstract": "We propose that the flat decay phase in the first 10$^2$--10$^4$ seconds of the X--ray light curve of Gamma Ray Bursts can be interpreted as prolonged activity of the central engine, producing shells of decreasing bulk Lorentz factors $\\Gamma$. The internal dissipation of these late shells produces a continuous and smooth emission, usually dominant in X--rays and sometimes in the optical. When $\\Gamma$ of the late shells is larger than $1/\\theta_j$, where $\\theta_j$ is the jet opening angle, we see only a portion of the emitting surface. Eventually, $\\Gamma$ becomes smaller than $1/\\theta_j$, and the entire emitting surface is visible. When $\\Gamma=1/\\theta_j$ there is a break in the light curve, and the plateau ends. During the plateau phase, we see the sum of the ``late--prompt'' emission (due to late internal dissipation), and the ``real afterglow'' emission (due to external shocks). A variety of different optical and X--ray light curves is possible, explaining why the X--ray and the optical light curves often do not track each other, and why they often do not have simultaneous breaks. ",
        "introduction": "The so called ``Steep--Flat--Steep\" behavior \\cite{gt05, nousek05} of the early (up to $\\sim$a day) X--ray afterglow was unpredicted before we could observe it with {\\it Swift}. It has been interpreted in several ways (for reviews, see e.g. \\cite{zhang07}) none of which seems conclusive. The spectral slope does not change across the temporal break from the shallow to the normal decay phase, ruling out a changing spectral break as a viable explanation. An hydrodynamical or geometrical nature of the break is instead preferred. Furthermore, the X--ray and optical lightcurves often do not track one another (e.g. \\cite{pana06, pana07}) suggesting a possible different origin.   To solve these difficulties Uhm \\& Beloborodov \\cite{uhm07} and Genet, Daigne \\& Mochkovitch \\cite{genet07} suggested that the X--ray plateau emission is not due to the forward, but to the reverse shock running into ejecta of relatively small (and decreasing) Lorentz factors. This however requires an appropriate $\\Gamma$--distribution of the ejecta, and also the suppression of the X--ray flux produced by the forward shock.  We (\\cite{gg07}) instead suggested that the plateau phase of the X--ray emission (and sometimes even of the optical) is due to a prolonged activity of the central engine (see also \\cite{lp07}), responsible for a ``late--prompt'' phase: after the early ``standard\" prompt the central engine continues to produce for a long time (i.e. days) shells of progressively lower power and bulk Lorentz factor. The dissipation process during this and the early phases occur at similar radii (namely close to the transparency radius). The reason for the shallow decay phase, and for the break ending it, is that the $\\Gamma$--factors of the late shells are motonically decreasing, allowing to see an increasing portion of the emitting surface, until all of it is visible. Then the break occurs when $\\Gamma=1/\\theta_j$.  \\begin{figure} \\includegraphics[height=0.5\\textheight]{ghisellini_f1.eps} \\caption{ Cartoon of the proposed model, and schematic illustration of the different components contributing to the X--ray and optical light curves, as labelled. Scales are arbitrary. The early prompt phase is erratic, with shells of varying $\\Gamma$ and power. Then the central engine produces shells of progressively less power and bulk Lorentz factors, producing a smoother light curve. Since the average $\\Gamma$--factor is decreasing, the observer sees an increasing portion of the emitting area, until all of it becomes visible when $\\Gamma \\sim 1/\\theta_j$. When this occurs there is a break in the light curve, associated with the ending of the shallow phase. The case illustrated here is only one (likely the most common) possible case, when the X--ray flux is dominated by late prompt emission (solid line, the dotted line corresponds to an extrapolation at very late times), while the optical flux is dominated by the real afterglow (dashed). Adapted from \\cite{gg07}. } \\end{figure} ",
        "conclusions": ""
    },
    "0801/0801.1342_arXiv.txt": {
        "abstract": "I investigate the discrepancy between the evolution and pulsation masses for Cepheid variables. A number of recent works have proposed that non-canonical mass-loss can account for the mass discrepancy. This mass-loss would be such that a 5$M_{\\odot}$ star loses approximately 20\\% of its mass by arriving at the Cepheid instability strip; a 14$M_{\\odot}$ star, none. Such findings would pose a serious challenge to our understanding of mass-loss. I revisit these results in light of the Padova stellar evolutionary models and find evolutionary masses are ($17\\pm5$)\\% greater than pulsation masses for Cepheids between $5<M/M_{\\odot}<14$. I find that mild internal mixing in the main-sequence progenitor of the Cepheid are able to account for this mass discrepancy. ",
        "introduction": "$\\delta$ Cepheid variables are an essential step in our determination of extragalactic distances. Apart from their use as distance indicators, the regularity of Cepheid pulsation provides a well-defined set of observational parameters with which to probe the course of stellar evolution for intermediate mass stars.  Over the preceeding decades much effort has been devoted to reconciling mass determinations for Cepheids from the various methods at our disposal (see \\citet{Cox80} for a review). The longest standing of these, the Cepheid pulsation mass discrepancy was first revealed by \\cite{Stobie69} who showed that pulsation masses for bump Cepheids\\footnote{Bump Cepheids are distinguished by a secondary local maximum of the light curve seen in Cepheids with periods in the range of 6 to 16 days \\citep{Bono00b,Hertzsprung26}.} were significantly lower than those predicted by stellar evolutionary models. \\cite{Andreasen88} showed that by artifically enhancing the opacity over the temperature range $1.5\\times10^{5}K<T<8\\times10^{5}K$ by a factor of 2.5 it was possible to remove the discrepancy. More detailed modelling of opacities by the OPAL \\citep{Rogers92} and the Opacity Project \\citep{Seaton94} confirmed an increase in this temperature range due to metal opacity.  The implementation of new opacities largely resolved the bump Cepheid mass discrepancy \\citep{Moskalik92}. However, despite convergence, a number of subsequent studies have shown that the discrepancy remains significant and requires explanation. Studies of Galactic \\citep{Natale07,Caputo05}, LMC \\citep{Wood97,Keller02,Keller06,Bono02} and SMC \\citep{Keller06} Cepheids have shown that the masses determined via pulsation modelling are $\\sim$15-20\\% less massive than those expected from evolutionary models. Dynamical masses for Cepheids are difficult to obtain given the low spatial density of Cepheids. The works of \\citet{Benedict07,Evans07, Evans06, Evans98} present dynamical masses for six Cepheids. Albeit with large associated uncertainties, these results confirm the conclusions drawn for pulsation modelling; that the evolutionary masses appear $\\sim15$\\% larger.  From an evolutionary perspective, Cepheids are understood to be post-red giant stars crossing the instability strip on so-called blue loops following the initiation of core-He burning. To ascribe an evolutionary mass one takes the Cepheid's luminosity and, using a mass-luminosity (M-L) relation that is derived from evolutionary models, derives the mass of the Cepheid. The M-L relation can be modified substantially by the treatment of internal mixing and mass-loss. Both processes feature complex hydrodynamical and radiative mechanisms for which we, at present, only possess empirical approximations.  The treatment of internal mixing modifies the size of the helium core established during the star's main-sequence (MS) evolution. Overshoot at the edge of the convective core of the Cepheid progenitor mixes additional hydrogen into the core and hence increases the helium core mass. As a consequence the post-MS evolution occurs at a higher luminosity. Mass loss, in an ad-hoc manner at least, offers a mechanism to modify the M-L relation by directly reducing the mass of a Cepheid.  The properties of pulsation, on the other hand, are dependent on the structure of the atmosphere of the Cepheid. \\citet{Keller06} show that the morphology of a bump Cepheid light curve is highly sensitive to the mass, luminosity, effective temperature and metallicity. Hence modelling of the light curve can be used to determine a pulsation Cepheid mass that is entirely independent of stellar evolution calculations.  In the work of \\cite{Bono02} and \\cite{Caputo05} it is proposed that mass-loss can account for the mass discrepancy between pulsation and evolutionary masses. \\citet{Caputo05} also conclude that models that incorporate additional internal mixing in the vicinity of the convective core are not able to explain the mass discrepancy. In this paper we revisit these conclusions and present a scenario for resolution of the mass discrepancy. ",
        "conclusions": "In this paper I revisit the conclusions of \\citet{Caputo05} regarding the cause of the observed discrepancy between the evolution and pulsation masses for Cepheids. \\citet{Caputo05} find that Cepheids of 5$M_{\\odot}$ have 20\\% less mass, as determined from pulsation analysis, than expected from evolutionary calculations, while the evolution and pulsation masses for Cepheids of 14$M_{\\odot}$ agree. In order to explain this finding Caputo et al.\\ propose a scenario of non-standard mass-loss:- one that sees increased mass-loss at lower masses and drops to negligible mass-loss at $M \\ge 14M_{\\odot}$. This scheme of mass-loss would be counter to the observationally grounded evidence that accumulated mass-loss by the epoch of core-He burning increases with increasing stellar mass. Furthermore, Caputo et al.\\ claim that convective core overshoot is unable to provide a solution as it can not account for the trend of mass discrepancy with Cepheid mass.  The findings of Caputo et al.\\ are based on a Cepheid mass-luminosity relationship that proves erroneous when extrapolated to higher masses. In this paper I show that including a full description of the mass-luminosity relation results in a mass discrepancy in which evolutionary masses are ($17\\pm5$)\\% greater than pulsation masses. The trend of mass discrepancy with Cepheid mass is removed.  I propose that additional internal mixing, as parametrised in the convective core overshooting paradigm, is the primary mechanism giving rise to the mass discrepancy. The level of convective core overshoot so derived is $\\Lambda_{c} = 0.67\\pm0.17$, a value which agrees with previous determinations from a range of techniques."
    },
    "0801/0801.1803_arXiv.txt": {
        "abstract": "We present an updated all-particle energy spectrum of primary cosmic rays in a wide range from $10^{14}$ eV to $10^{17}$ eV using $5.5 \\times 10^{7}$ events collected in the period from 2000 November through 2004 October by the Tibet-III air-shower array located at 4300 m above sea level (atmospheric depth of 606 g/cm$^{2}$). The size spectrum exhibits a sharp knee at a corresponding primary energy around 4 PeV. This work uses increased statistics and new simulation calculations for the analysis. We performed extensive Monte Carlo calculations and  discuss the model dependences involved in the final result assuming interaction models of QGSJET01c and SIBYLL2.1 and primary composition models of heavy dominant (HD) and proton dominant (PD) ones. Pure proton and pure iron primary models are also examined as extreme cases. The detector simulation was also made to improve the accuracy of determining the size of the air showers and the energy of the primary particle. We confirmed that the all-particle energy spectra obtained under various plausible model parameters are not significantly different from each other as expected from the characteristics of the experiment at the high altitude, where the air showers of the primary energy around the knee reaches near maximum development and their features are dominated by electromagnetic components leading to the weak dependence on the interaction model or the primary mass. This is the highest-statistical and the best systematics-controlled measurement covering the widest energy range around the knee energy region. ",
        "introduction": "\\label{intro}  Although almost 100 years have passed over since the discovery of cosmic rays, their source and acceleration mechanism are still not fully understood. The energy spectrum and chemical composition of cosmic rays can be key information for probing their origin, acceleration mechanism and propagation mechanism. The cosmic-ray spectrum has been observed by many ground-based experiments to resemble two power laws, having a form dj/dE $\\propto$ $E^{-\\gamma}$, with $\\gamma$ = 2.7 below the energy around 4 $\\times$ $10^{15}$ eV, and then steepening to $\\gamma$ = 3.1 above this energy \\cite{Hoerandel2}. The change of the power index at this energy is called the spectral ``knee''.  Although the existence of the knee has been well established experimentally, there are still controversial arguments on its origin. Proposals for its origin range from astrophysical scenarios like the change of acceleration mechanisms \\citep{Berezhko,Stanev,Kobayakawa,Volk} at the sources of cosmic rays (supernova remnants, pulsars, etc.), the single source assumption \\cite{Erlykin}, or effects due to the propagation \\citep{Ptuskin,Candia} inside the galaxy (diffusion, drift, escape from the Galaxy) to particle-physics models like the interaction with relic neutrinos \\cite{Wigmans} during transport or new processes in the atmosphere \\cite{Nikolsky} during air-shower development. Common to all models is the prediction of a change of the chemical composition over the knee region. Direct measurements of primary cosmic rays on board balloons or satellites are the best ways for the study of the chemical composition, however, the energy region covered by them with sufficient statistics are limited to $10^{14}$ eV. The energy spectrum and chemical composition of primary cosmic rays around the knee, therefore, has to be studied with ground-based air-shower experiments using surface array and/or detectors for Cerenkov light.  Many reports have so far been made on the energy spectrum as well as the chemical composition of primary cosmic rays. Although the global features of the all-particle spectrum agree well when we take into account the systematic errors of about 20\\% involved in the energy scale \\cite{Hoerandel2}, there are still serious disagreements in the chemical composition depending on experimental methods, for example, the knee composition obtained by the Tibet and KASCADE experiments can be summarized as follows. We have already reported the energy spectrum of protons and helium in the energy range from 200 TeV to 10,000 TeV \\citep{Amenomori5,Amenomori2} from air-shower core observation, suggesting a steep power index of approximately -3.1. This indicates that the power index of light component is changed from approximately -2.7 as measured by direct observations to -3.1 around a few hundred TeV. Hence, the light component should become less abundant at the knee and the main component responsible to the structure of the knee must be heavier than helium. Furthermore, the spectral shape of light component seems to keep the power law instead of the exponential cutoff. On the contrary, the KASCADE using electron-muon size analysis \\cite{Antoni} claims that the knee in the all-particle spectrum is due to the steepening of the spectra of light elements with exponential type cutoff.   The accurate measurement of the all-particle energy spectrum around the knee is essential to establish the chemical composition at this energy range.  There is no precise measurement of the chemical composition around the knee region yet and it is in fact impossible to discriminate individual elements clearly by indirect observations. Therefore, most of the works made so far simply discussed the average mass $<\\ln A>$. Another approach is the unfolding of the all-particle spectrum using shower characteristics like electron-muon ratio, depth of the shower maximum and so on. In these methods, the detailed information of the all-particle spectrum plays an important role in determining the chemical composition. It is also expected that the specific features of each component like cutoff energy or source characteristics should be reflected in the shape of the all-particle spectrum as discussed in the single source model \\cite{Erlykin}. The important features of the all-particle spectrum are the absolute intensity, the position of the knee, the difference of the power index before and after the knee and the sharpness in the size spectrum, which are deeply connected with the acceleration mechanism and the source of cosmic rays.  The merit of the air-shower experiment in Tibet is that the atmospheric depth of the experimental site (4300 m a.s.l., 606 g/{cm}$^{2}$) is close to the maximum development of the air showers with energies around the knee almost independent of the masses of primary cosmic rays as demonstrated in Fig.~\\ref{fig:1} for vertically incident cosmic rays. It should be also noted that the number of shower particles is dominated by electromagnetic component with minor contribution of muons, whose interaction model dependence is known to be rather large among current interaction models leading to a large systematic error in the experiments carried out at the sea level because of the large contribution of muons. In other words, the air shower observation at high altitude is sensitive to the most forward region of the hadronic interactions in the center of momentum system (CMS) where high energy secondaries are produced, and the electromagnetic component as a decay product of neutral pions dominate the number of shower particles, while it is insensitive to the central region of the CMS where large number of muons as decay product of charged pions are produced. The differences among current interaction models are mainly related to the central region as seen in the problem of the electron-muon correlation. Hence, the air-shower experiment in Tibet can determine the primary cosmic-ray energy much less dependently upon the chemical composition and the interaction model than experiments at the sea level.  \\begin{figure} \\epsscale{.80} \\plotone{./f1.eps} \\caption{Average transition curves of air-shower size induced by protons and iron nuclei for a vertical incidence. } \\label{fig:1} \\end{figure}   We have already reported the first result on the all-particle spectrum around the knee region based on data from 2000 November to 2001 October observed by the Tibet-III air-shower array \\cite{Amenomori8}. In this paper, we present an updated all-particle energy spectrum using data set collected in the period from 2000 November through 2004 October. The updates are due to (1) increased statistics by approximately 2.6 times, (2) use of new simulation codes and (3) improvement of the lateral structure function used for the size estimation of air showers. The previous result was obtained using almost the same analysis as used in Tibet I \\cite{Amenomori1} except for the parameters which depend on the detector configuration. In the present paper, the simulation code Cosmos is replaced by Corsika with interaction models QGSJET01c and SIBYLL2.1, which is now widely used in many analyses by other works and enables the comparison of this work with others easier. The third update on the structure function is made to cover wider energy range than before (see section \\ref{ModNKG}).  Thus, we obtained the all-particle  energy spectrum of cosmic rays in a wide range over 3 decades between $10^{14}$ eV  and $10^{17}$ eV and the updated result is compared with previous ones. ",
        "conclusions": "\\subsection{The model dependence of the size spectrum}  \\begin{figure} \\epsscale{.80} \\plotone{./f15.eps} \\caption{The model dependence of the size spectrum of nearly vertical air showers (sec($\\theta$) $\\leq$ 1.1).} \\label{fig:15} \\end{figure}  As a first step in checking the model dependence, the difference of the size spectra derived by the lateral fitting is examined using the structure functions based on five models (QGSJET interaction model with four primary composition models and SIBYLL+HD model). Shown in Fig.~\\ref{fig:15} is the model dependence of the air-shower size spectrum of nearly vertical air showers. It is seen that the model dependence of the air-shower size is small (less than 5\\% in absolute intensity on the primary composition model, and also less than 5\\% on the hadronic interaction models).   \\subsection{All-particle energy spectrum} \\subsubsection{Determination of the primary energy} \\begin{figure} \\epsscale{.80} \\plotone{./f16.eps} \\caption{Scatter plots of the primary energy $E_0$ and the estimated shower size $N_e$ based on the QGSJET+HD model. } \\label{fig:16} \\end{figure}   In Fig.~\\ref{fig:16} we show the scatter plots of the primary energy $E_0$, and the estimated shower size $N_e$ based on the QGSJET+HD model. \\begin{figure} \\epsscale{.80} \\plotone{./f17.eps} \\caption{The correlations between the estimated shower size $N_e$ and the primary energy $E_0$ for given models. }  \\label{fig:17} \\end{figure}  \\begin{figure} \\epsscale{.80} \\plotone{./f18.eps} \\caption{The differential energy spectra of all particles obtained by the present work using 5 models and they are compared with direct observations. JACEE \\cite{Asakimori},RUNJOB \\cite{Apanasenko}, Grigorov \\cite{Grigorov} } \\label{fig:18} \\end{figure}   Fig.~\\ref{fig:17} shows the correlations between the estimated shower size $N_e$ and the primary energy $E_0$ for QGSJET + 4 primary models and SIBYLL+HD model. In this figure, the results of pure composition models are also shown just to help the understanding of the composition dependence in the energy determination. Using these pure composition model will result in remarkably different primary energy spectrum as shown later. The correlation between the estimated shower size $N_e$ and the primary energy $E_0$ for $\\sec\\theta$ $\\leq$ 1.1 can be well fitted using a following conversion function for each interaction models and composition models.  \\begin{equation} \\label{conver} E_0 = a \\times {(\\frac{N_e}{1000})}^{b} \\mbox{\\hspace{2cm}  [TeV]}, \\end{equation} \\noindent where the numerical values of $a$ and $b$ in eq.(\\ref{conver}) are summarized in Table ~\\ref{tab:01a}. These approximations are valid for nearly vertical air showers (sec($\\theta$) $\\leq$ 1.1 ) at Yangbajing altitude.    As seen in Fig.~\\ref{fig:17}, the difference of the conversion factor between two mixed composition models of HD and PD is not significant and almost disappears above a few times 1000 TeV in spite of the large difference of the fractional abundances of chemical components. From the comparison between QGSJET+HD and SIBYLL+HD, one can see that the dependence on the interaction model is very small.    It is also noted that the energy resolutions in different interaction models and two mixed composition models are very close to each other. The differences are estimated as 36\\% and 17\\% at energies around 200 TeV and 2000 TeV, respectively, in the case of QGSJET+HD model. Corresponding values for SIBYLL+HD model are 38\\% and 19\\% and that for QGSJET+PD model are 39\\% and 19\\%. It is commonly understood that there is an energy dependent overestimation problem of the flux due to the steep power index of the cosmic ray energy spectrum when the error of the estimated energy increases with decreasing primary energy. It may be, however, worthwhile to note here that this effect is already included in our method of the energy determination using the conversion function, because the primary energy at a given shower size is determined including the contribution of all possible primary energies which leads to a smaller value than the case of without fluctuation reflecting the larger population of low energy primaries than high energies. To avoid methodical systematic error, the reproducibility of the primary flux was carefully examined using MC events and no significant deviation was found between MC input spectrum and the reconstructed spectrum.  \\begin{figure} \\epsscale{.80} \\plotone{./f19.eps} \\caption{The differential energy spectra of all particles obtained by the present work using mixed composition models compared with other experiments. JACEE \\cite{Asakimori},RUNJOB \\cite{Apanasenko}, Grigorov \\cite{Grigorov}, KASCADE \\cite{Antoni}, CASA-MIA \\cite{Glasmacher}, AKENO(1984) \\cite{Nagano1}, Tibet-III(ICRC2003) \\cite{Amenomori8}. } \\label{fig:19} \\end{figure}   \\subsubsection{Energy spectrum of all-particles and the knee parameters}   \\begin{figure} \\epsscale{.80} \\plotone{./f20.eps} \\caption{The differential energy spectra of all particles obtained by the present work above the knee compared with other experiments at the highest energy range. KASCADE \\cite{Antoni}, AKENO(1984) \\cite{Nagano1}, AKENO(1992) \\cite{Nagano2}, AGASA(2003) \\cite{AGASA}, AUGER(ICRC2007)(SD vertical) \\citep{Yamamoto,Roth} } \\label{fig:20} \\end{figure}   The all-particle energy spectrum of primary cosmic rays in a wide range over 3 decades between 1 $\\times$ $10^{14}$ eV  and 1 $\\times$ $10^{17}$ eV is shown in Fig.~\\ref{fig:18} for 5 models. In this figure, it is important to note that the position of the knee is clearly seen at the energy around 4 PeV irrespective of the model used.   The model dependence of the conversion function shown in Fig.~\\ref{fig:17} may lead to different shapes of the all-particle energy spectrum. The model dependence on the primary composition can be checked by comparing the results of QGSJET+HD and QGSJET+PD. Although the composition is fairly different between the HD and PD model at energies above $10^{15}$ eV, the difference of the absolute intensity is 20\\% at most between the two models and decreases with increasing primary energy. As mentioned above, it may be helpful to show the composition dependence by the pure component assumption. The difference between pure proton and pure iron primary model becomes large at the lower energy region of the range shown in this figure, where the difference of the intensities between two models exceeds a factor of 3 at 10$^{14}$ eV, although these results do not belong to the argument for the reality at least in the energy range lower than a few times 10$^{14}$ eV. It should be noted that the mixed composition models give more realistic extension of the direct observations and the composition model dependence almost disappears above 10$^{16}$ eV. This is the remarkable characteristics of the Tibet experiment.   The interaction model dependence is seen by comparing the results of QGSJET+HD and SIBYLL+HD. This comparison shows that the shapes of the spectrum from both interaction models are almost the same and the difference in the absolute intensity is within 10\\%. In Fig.~\\ref{fig:19}, we show the results using the mixed composition model in comparison with other works including our previous work presented in ICRC2003. Our results seem to lead to the data of direct observation smoothly in the low energy side. Fig.~\\ref{fig:20} shows the higher energy part above the knee to compare with the surface array experiments done at the highest energy region. It is very interesting to see how the smooth extension of our spectrum above $10^{16}$ eV is connected with the data at the highest energy region.  The intensities of all-particle energy spectra measured by the Tibet-III array are also posted in Table~\\ref{tab:03} and the summary of the measured all-particle energy spectrum and knee parameters are listed in Table~\\ref{tab:02}, where $\\gamma_1$ is the best fitted-index for the energy range 100 TeV $<$ $E_0$ $<$ 1 PeV, and $\\gamma_2$ is for the energy above 4 PeV.   \\begin{table}[t] \\begin{center} \\caption{The summary of the knee parameters. The symbol $\\gamma_1$ is the best fitted-index for the energy range 100 TeV $<$ $E_0$ $<$ 1 PeV, and $\\gamma_2$ is for the energy above 4 PeV. } \\begin{tabular}{l|c|c} \\tableline\\tableline Model & knee position & Index of \\\\ & (PeV) & spectrum \\\\ \\tableline QGSJET& (4.4 $\\pm$ 0.1) & $\\gamma_1$= -2.81 $\\pm$ 0.01 \\\\ +Iron&  & $\\gamma_2$= -3.21 $\\pm$ 0.01 \\\\ \\tableline QGSJET& (4.0 $\\pm$ 0.1) & $\\gamma_1$= -2.67 $\\pm$ 0.01 \\\\ +HD&  & $\\gamma_2$= -3.10 $\\pm$ 0.01 \\\\ \\tableline QGSJET& (3.8 $\\pm$ 0.1) & $\\gamma_1$= -2.65 $\\pm$ 0.01 \\\\ +PD &  & $\\gamma_2$= -3.08 $\\pm$ 0.01 \\\\ \\tableline SIBYLL & (4.0 $\\pm$ 0.1) & $\\gamma_1$= -2.67 $\\pm$ 0.01  \\\\ +HD &  & $\\gamma_2$= -3.12 $\\pm$ 0.01  \\\\ \\tableline QGSJET& (3.4 $\\pm$ 0.1) & $\\gamma_1$= -2.60 $\\pm$ 0.01 \\\\ +Proton &  & $\\gamma_2$= -3.03 $\\pm$ 0.01 \\\\ \\tableline \\tableline \\end{tabular} \\label{tab:02} \\end{center} \\end{table}"
    },
    "0801/0801.1032_arXiv.txt": {
        "abstract": "Modern hydrodynamical simulations offer nowadays a powerful means to trace the evolution of the X--ray properties of the intra--cluster medium (ICM) during the cosmological history of the hierarchical build up of galaxy clusters. In this paper we review the current status of these simulations and how their predictions fare in reproducing the most recent X--ray observations of clusters. After briefly discussing the shortcomings of the self--similar model, based on assuming that gravity only drives the evolution of the ICM, we discuss how the processes of gas cooling and non--gravitational heating are expected to bring model predictions into better agreement with observational data. We then present results from the hydrodynamical simulations, performed by different groups, and how they compare with observational data. As terms of comparison, we use X--ray scaling relations between mass, luminosity, temperature and pressure, as well as the profiles of temperature and entropy. The results of this comparison can be summarised as follows: {\\em (a)} simulations, which include gas cooling, star formation and supernova feedback, are generally successful in reproducing the X--ray properties of the ICM outside the core regions; {\\em (b)} simulations generally fail in reproducing the observed ``cool core'' structure, in that they have serious difficulties in regulating overcooling, thereby producing steep negative central temperature profiles. This discrepancy calls for the need of introducing other physical processes, such as energy feedback from active galactic nuclei, which should compensate the radiative losses of the gas with high density, low entropy and short cooling time, which is observed to reside in the innermost regions of galaxy clusters. ",
        "introduction": "\\label{Introduction} Clusters of galaxies form from the collapse of exceptionally high density perturbations having typical size of $\\sim 10$ Mpc in a comoving frame. As such, they mark the transition between two distinct regimes in the study of the formation of cosmic structures. The evolution of structures involving larger scales is mainly driven by the action of gravitational instability of the dark matter (DM) density perturbations and, as such, it retains the memory of the initial conditions. On the other hand, galaxy--sized structures, which form from initial fluctuations on scales of $\\sim 1$ Mpc, evolve under the combined action of gravity and of complex gas--dynamical and astrophysical processes. On such scales, gas cooling, star formation and the subsequent release of energy and metal feedback from supernovae (SN) and active galactic nuclei (AGN) have a deep impact on the observational properties of the diffuse gas and of the galaxy population.  In this sense, clusters of galaxies can be used as invaluable cosmological tools and astrophysical laboratories (see \\citealt{2002ARA&A..40..539R} and  \\citealt{2005RvMP...77..207V} for reviews). These two aspects are clearly interconnected with each other. From the one hand, the evolution of the population of galaxy clusters and their overall baryonic content provide in principle powerful constraints on cosmological parameters. On the other hand, for such constraints to be robust, one has to understand in detail the physical properties of the intra--cluster medium (ICM) and its interaction with the galaxy population.  The simplest model to predict the properties of the ICM and their evolution has been proposed by \\citet{1986MNRAS.222..323K}. This model is based on the assumption that the evolution of the thermodynamical properties of the ICM is determined only by gravity, with gas heated to the virial temperature of the hosting DM halos by accretion shocks (\\citealt{bykov2008a} - Chapter 7, this volume). Since gravitational interaction does not introduce any preferred scale, this model has been called ``self--similar''\\footnote{Strictly speaking, self--similarity also requires that no characteristic scales are present in the underlying cosmological model. This means that the Universe must obey the Einstein--de-Sitter expansion law and that the shape of the power spectrum of density perturbations is a featureless power law. In any case, the violation of self--similarity introduced by the standard cosmological model is negligible with respect to that related to the non--gravitational effects acting on the gas.}. As we shall discuss, this model provides precise predictions on the shape and evolution of scaling relations between X--ray luminosity, entropy, total and gas mass, which have been tested against numerical hydrodynamical simulations (e.g., \\citealt{1998ApJ...503..569E,1998ApJ...495...80B}). These predictions have been recognised for several years to be at variance with a number of observations. In particular, the observed relation between X--ray luminosity and temperature (e.g. \\citealt{1998ApJ...504...27M,1999MNRAS.305..631A,2004MNRAS.350.1511O}) is steeper and the measured level of gas entropy higher than expected (e.g., \\citealt{2003MNRAS.343..331P,2005A&A...429..791P}), especially for poor clusters and groups. This led to the concept that more complex physical processes, related to the heating from astrophysical sources of energy feedback, and radiative cooling of the gas in the central cluster regions, play a key role in determining the properties of the diffuse hot baryons.  Although semi--analytical approaches (e.g., \\citealt{2001ApJ...546...63T,2005RvMP...77..207V}, and references therein) offer invaluable guidelines to this study, it is only with hydrodynamical simulations that one can capture the full complexity of the problem, so as to study in detail the existing interplay between cosmological evolution and the astrophysical processes.  In the last years, ever improving code efficiency and supercomputing capabilities have opened the possibility to perform simulations over fairly large dynamical ranges, thus allowing to resolve scales of a few kiloparsecs (kpc), which are relevant for the formation of single galaxies, while capturing the global cosmological environment on scales of tens or hundreds of Megaparsecs (Mpc), which are relevant for the evolution of galaxy clusters. Starting from first attempts, in which only simplified heating schemes were studied (e.g., \\citealt{1995MNRAS.275..720N,2001ApJ...555..597B,2002MNRAS.336..409B}), a number of groups have studied the effect of introducing also  cooling (e.g. \\citealt{1993ApJ...412..455K,2000ApJ...536..623L,2001ApJ...552L..27M,2002ApJ...579...23D,2003MNRAS.342.1025T}), of more realistic sources of energy feedback (e.g.,  \\citealt{2004MNRAS.348.1078B,2007MNRAS.377..317K,2007astro.ph..3661N,2007arXiv0705.2238S}), of thermal conduction (e.g., \\citealt{2004ApJ...606L..97D}), and of non--thermal pressure support from magnetic fields (e.g.,  \\citealt{2001A&A...369...36D}) and cosmic rays (e.g.,  \\citealt{2007MNRAS.378..385P}).  In this paper, we will review the recent advancement performed in this field of computational cosmology and critically discuss the comparison between simulation predictions and observations, by restricting the discussion to the thermal effects. As such, this paper complements the reviews by \\citealt{borgani2008} - Chapter 18, this volume, which reviews the study of the ICM chemical enrichment and by \\citealt{dolag2008b} - Chapter 15, this volume, which reviews the study of the non--thermal properties of the ICM from simulations. We refer to the reviews by \\citealt{dolag2008a} - Chapter 12, this volume for a description of the techniques of numerical simulations and by \\citealt{kaastra2008} - Chapter 9, this volume for an overview of the observed thermal properties of the ICM.  The scheme of the presentation is as follows. In Sect. \\ref{models} we  briefly discuss the self--similar model of the ICM and how the action of non--gravitational heating and cooling are expected to alter the predictions of this model. Sect.~\\ref{scalings} and \\ref{profiles}  overview the results obtained on the comparison between observed and simulated scaling relations and profiles of X--ray observable quantities, respectively. In Sect.~\\ref{summary} we  summarise and critically discuss the results presented. ",
        "conclusions": "\\label{summary} In this paper we reviewed the current status of the comparison between cosmological hydrodynamical simulations of the X--ray properties of the intra--cluster medium (ICM) and observations. We first presented the basic predictions of the self--similar model, based on the assumption that only gravitational processes drive the evolution of the ICM. We then showed how a number of observational facts are at variance with these predictions, with poor clusters and groups characterised by a relatively lower density and higher entropy of the gas. This calls for the need of introducing some extra physical processes, such as non--gravitational heating and radiative cooling, which are able to break the self--similarity between objects of different size. The results based on simulations, which include these effects, can be summarised as follows.\\\\ {\\bf (1)} The observed scaling relation between X--ray luminosity and temperature can be reproduced by simulations including cooling only, but at the expense of producing a too large fraction of cooled gas. Introducing ad--hoc schemes of entropy injection at high redshift (the so--called pre--heating) can eventually produce the correct $L_X$--$T$ relation. However, no simulations have been presented so far, in which a good agreement with observations is achieved by using a feedback scheme in which the energy release and thermalisation from SN or AGN is self--consistently computed by the simulated star formation or accretion onto a cosmologically evolving population of black holes.\\\\ {\\bf (2)} Simulations which include cooling, star formation and SN feedback, correctly reproduce the observed mass--temperature and ``pressure''--temperature relations. Quite interestingly, this agreement is achieved once the masses of simulated clusters are estimated by applying the same mass estimators, based on the assumption of hydrostatic equilibrium, which are applied to the analysis of real clusters. The violation of hydrostatic equilibrium, related to the non--thermal pressure support from subsonic gas motions, would otherwise lead to a $\\sim 20$ per cent overestimate of the normalisation of the above scaling relations for simulated clusters.\\\\ {\\bf (3)} Both non--radiative and radiative runs naturally predict the observed negative gradients of the temperature profiles outside the cluster core regions, $r\\gtrsim 0.2 r_{200}$. This suggests that the ICM thermal structure in these regions is indeed dominated by the action of gravitational processes, such as heating from shocks associated to supersonic gas accretion.\\\\ {\\bf (4)} Introducing cooling has the effect of steepening the temperature profiles in the innermost regions, as a consequence of the adiabatic compression of inflowing gas, caused by the lack of pressure support. Therefore, gas in the cores of simulated clusters generally lies in a high--temperature phase, which is very well separated from the cold ($\\sim 10^4$ K) phase. This is at variance with the observed ``cool core'' structure of real clusters, which show instead the presence of a fair amount of gas, down to a limiting temperature of $\\sim 1/3$ of the virial temperature, which formally has a rather short cooling time. The presence of this gas makes the observed temperature profiles of relaxed clusters to peak at $r\\simeq 0.2r_{200}$, thereby decreasing by about a factor two in the innermost sampled regions. The discrepancy between the simulated and the observed ``cool core'' calls for the need of introducing in cluster simulations a suitable feedback mechanism which is able to compensate the radiative losses, keeps pressurised gas in the central regions, and suppresses the mass deposition rate.\\\\ {\\bf (5)} Simulations are generally rather successful in reproducing the observed slope of the entropy profiles outside the core regions, $S\\propto r^\\alpha$ with $\\alpha \\simeq 1$. This slope is also close to that predicted by models based on gravitational heating from spherical accretion, thus lending further support to the picture that gravity drives the evolution of the ICM outside the core regions. Quite intriguingly, however, the normalisation of the profiles out to the largest sampled radii, $\\lesssim r_{500}$, is observed to have a milder temperature dependence than expected from the self--similar model, $S\\propto T^\\alpha$ with $\\alpha\\simeq 0.65$ instead of $\\alpha\\simeq 1$. This entropy excess in poor systems can be reproduced in simulations only by adding a diffuse pre--heating mechanism, which smoothes the pattern of gas accretion, thereby leading to an amplification of entropy generation from shocks in poorer systems. However, this mechanism is able to provide the correct explanation only for clusters with temperature $T\\lesssim 3$ keV.  Generating entropy amplification in hotter systems would require an exceedingly large amount of pre--heating, which would generate too shallow entropy profiles.  In general, the above results demonstrate the capability of cosmological hydrodynamical simulations to predict the correct thermodynamical properties of the ICM outside the central regions of galaxy clusters. On the other hand, even simulations based on an efficient SN feedback fail in preventing overcooling and providing the correct description of the thermodynamical ICM properties in the central regions. The ``cool core'' failure of simulations based on stellar feedback is also witnessed by the exceedingly massive and blue central galaxies that are generally produced (see also \\citealt{borgani2008} - Chapter 18, this volume).  The fact that cool cores are observed for systems spanning at least two orders of magnitude in mass, from $\\sim 10^{13}M_\\odot$ to $\\sim 10^{15}M_\\odot$, suggests that only one self--regulated mechanism should be mainly responsible for this, rather than the combination of several mechanisms possibly acting over different time--scales. Although AGN are generally thought to naturally provide such a mechanism, a detailed comparison between observational data and an extended set of cluster simulations, including this kind of feedback is still lacking. Furthermore, understanding in detail the role of AGN in determining the ICM properties requires addressing in detail two major issues. Firstly, the accretion process onto the central black hole involves scales of the order of the parsec, while the X--ray observable effects involve scales of $\\sim 10$--100 kpc. This requires understanding the cross--talk between two ranges of scales, which differ by at least four orders of magnitude. Secondly, the total energy budget available from AGN is orders of magnitude larger than that required to regulate cooling flows. Therefore, one needs ultimately to understand the channels for the thermalisation of the released energy and how this naturally takes place without resorting to any ad--hoc tuning. These problems definitely need to be addressed by simulations of the next generation, which will be aimed at understanding in detail the evolution of the cosmic baryons and the observational X--ray properties of galaxy clusters."
    },
    "0801/0801.1204_arXiv.txt": {
        "abstract": "In the present work, it is shown that a chameleon scalar field having a nonminimal coupling with dark matter can give rise to a smooth transition from a decelerated to an accelerated phase of expansion for the universe. It is surprising to note that the coupling with the chameleon scalar field hardly affects the evolution of the dark matter sector, which still redshifts as $a^{-3}$. ",
        "introduction": "In the absence of a clear verdict in favour of any `dark energy' candidate, which drives the alleged late time acceleration of the universe \\cite{1}, every possibilities are being thoroughly investigated so as to find out what can really negotiate this latest challenge that theoretical physics is exposed to. The accelerated expansion is counter-intuitive, as the dominating interaction in the dynamics of the universe is gravity which is attractive. So the `dark energy' sector has to produce an effective anti-gravity effect. However, the actual problem is even more stringent, both observations and theoretical requirements demand that the acceleration must have set in quite late in the evolution of the universe \\cite{2}. The models are also required to be consistent with other observational data like those on the cosmic microwave background radiation. There are already some excellent reviews which summarize the problem and the efforts towards finding a consistent solution \\cite{3}. Amongst the most talked about models, the quintessence models enjoy a fair deal of popularity. These models contain a scalar field endowed with a potential, so that the contribution to the pressure sector, $p_{\\phi}=\\frac{1}{2}{\\dot{\\phi}}^2 - V(\\phi)$, can evolve to attain an adequately large negative value and drive the observed accelerated expansion. Given a particular temporal behaviour of the scale factor of the universe, it is always possible to find the requisite potential \\cite{4}. Naturally there is a long list of potentials that can do the trick \\cite{5}. None of them however can boast of a proper theoretical support from field theory. \\par In these scalar field models, the cold dark matter and dark energy are normally allowed to evolve independently. However, there are attempts to include an interaction amongst them so that one grows at the expense of the other ( see for example \\cite{6} ). Nonminimally coupled scalar fields, where there is a coupling between the scalar field and geometry, also posed themselves as possible candidates for explaining the present acceleration. For example, Brans - Dicke theory cropped up as a possible arena \\cite{7}. \\par A modification of the Brans - Dicke theory, where an interaction between the scalar field and the dark matter could be allowed, showed that the matter dominated era can have a transition from a decelerated to an accelerated expansion without the requirement of any dark energy \\cite{8}. On the other hand, a dark energy interacting with Brans - Dicke scalar field can give rise to a late time acceleration for a very wide range of potentials \\cite{9}. \\par A completely different approach has also been explored, where, unlike Brans - Dicke theory, the scalar field is allowed to have a ``non-minimal'' coupling with the dark matter sector, ( but not with the Ricci scalar ) through an interference term in the action \\cite{10}. This kind of a scalar field is dubbed as the ``chameleon field''. This field also proved quite useful in modelling an accelerated expansion of the universe. Many interesting cosmological possibilities with this chameleon field have been recently pointed out by Das, Corasaniti and Khoury \\cite{11}. A recent review indicates the other gravitational implications of such a field \\cite{12}. \\par In the present work, it has been shown that a simple chameleon field can indeed make room for a decelerated expansion in the early matter dominated era and allow for an accelerated expansion at the present epoch. This transition can take place quite smoothly at a recent past. \\par An intriguing feature of this field, as will be shown later, is that in spite of the interaction with the chameleon field the dark matter density falls off as $\\frac{1}{a^3}$, where $a$ is the scale factor of the universe, exactly similar to the behaviour expected where the dark matter sector does not interact with dark energy and satisfies its conservation equation by itself. So the apprehension that there could be discrepancy between a chameleon model and the expected $(1+z)^3$ dependence in the clustering of the cold dark matter \\cite{11} is ruled out. ",
        "conclusions": "The smooth transition from a decelerated to an accelerated phase of expansion driven by the chameleon field makes the latter worthy of attention. The model presented here is restricted, but it clearly shows that the nonminimal coupling of the field with the CDM sector governs the onset of the acceleration, but it hardly affects the $a^{-3}$ behaviour of the CDM. This field thus warrants more attention. If one could find a coupling $f(\\phi)$ so that the chameleon field has an oscillatory behaviour ( with a small amplitude ) at the beginning of the matter dominated epoch, but grows later to dictate the dynamics of the universe, it would solve many a fine tuning problems."
    },
    "0801/0801.2388_arXiv.txt": {
        "abstract": "We study the nature of faint, red-selected galaxies at $z \\sim 2-3$ using the \\textit{Hubble} Ultra Deep Field (HUDF) and \\spitzer\\ Infrared Array Camera (IRAC) photometry.  Given the magnitude limit of the \\hst\\ data, we detect candidate galaxies to $H_\\mathrm{AB}<26$~mag, probing lower-luminosity (lower mass) galaxies at these redshifts.  We identify 32 galaxies satisfying the $(\\nicj-\\nich)_\\mathrm{AB} > 1.0$~mag color selection, 16 of which have unblended \\mone\\ and \\mtwo\\ photometry from \\spitzer.  Using this multiwavelength dataset we derive photometric redshifts, masses, and stellar population parameters for these objects.  We find that the selected objects span a diverse range of properties over a large range of redshifts, $1 \\lsim z \\lsim 3.5$.  A substantial fraction (11/32) of the $(\\nicj - \\nich)_\\mathrm{AB} > 1$~mag population appear to be lower-redshift ($z \\lsim 2.5$), heavily obscured dusty galaxies or edge-on spiral galaxies, while others (12/32) appear to be galaxies at $2 \\lsim z \\lsim 3.5$ whose light at rest-frame optical wavelengths is dominated by evolved stellar populations.  We argue that longer wavelength data ($\\gsim 1$~\\micron, rest-frame) are essential for interpreting the properties of the stellar populations in red-selected galaxies at these redshifts.  Interestingly, by including \\spitzer\\ data many candidates for galaxies dominated by evolved stellar populations are rejected, and for only a subset of the sample (6/16) do the data favor this interpretation.  These objects have a surface density of $\\sim1$~arcmin$^{-2}$.  We place an upper limit on the space density of candidate massive evolved galaxies with $2.5 < z < 3.5$ and $\\nich(\\mathrm{AB}) \\leq 26$~mag of $n = 6.6^{+2.0}_{-3.0} \\times 10^{-4}$ Mpc$^{-3}$ with a corresponding upper limit on the stellar mass density of $\\rho^\\ast = 5.6^{+4.4}_{-2.8}\\times 10^{7}$ $M_\\odot$ Mpc$^{-3}$.  The $z > 2.5$ objects that are dominated by evolved stellar populations have a space density at most one-third that of $z\\sim 0$ red, early-type galaxies.  Therefore, at least two-thirds of present-day early-type galaxies assemble or evolve into their current configuration at redshifts below 2.5.  We find a dearth of candidates for low-mass ($\\lsim 2\\times 10^{10}$~$\\msun)$ galaxies at $1.5 < z \\lesssim 3$ that are dominated by passively evolving stellar populations even though the data should be sensitive to them; thus, at these redshifts, galaxies whose light is dominated by evolved stellar populations are restricted to only those galaxies that have assembled high stellar mass. ",
        "introduction": "In the hierarchical picture of galaxy evolution, galaxies assemble within dark matter haloes that merge over time \\citep[\\eg,][]{freedman01,sp03}.  This model has been highly successful at reproducing galaxy clustering properties based on well-understood gravitational physics \\citep[\\eg,][]{weinberg04}. However, it is difficult to test how galaxies assemble their stellar content within the context of these models because of uncertain physics coupling the dark matter and baryonic matter in these haloes.  Observationally, the cosmic star-formation-rate (SFR) density rises by roughly an order of magnitude from $z\\sim 0$ to 1 and remains roughly constant from $z\\sim 1$-6 \\citep[and references therein]{hopkins04}. At the same time, roughly half of the stellar-mass density in galaxies formed between $z\\sim $3 and 1, corresponding to this peak in the SFR density \\citep[\\eg,][]{dickinson03,rudnick03,glazebrook04}.  While theoretical models broadly reproduce this evolution in the global galaxy population \\citep[\\eg,][]{somerville01,hernquist03}, they have difficulty in reproducing star-formation trends in individual galaxies, particularly in accounting for the formation and evolution of specific SFRs of the most massive objects \\citep{baugh03,somerville04,delucia06,croton06}.  In particular, it is difficult for models to produce galaxies (at any redshift) that are mostly devoid of current star-formation and evolve only passively \\citep[\\eg,][]{croton06,dave06}.  Quantifying the number of passively evolving galaxies --- and as a function of redshift and stellar mass --- is helpful for constraining the processes that form stars within galaxies.  Searches for high-redshift galaxies that are dominated by older, passively evolving stellar populations are challenging.  The light from these stars peaks at rest-frame optical and near-IR wavelengths ($\\simeq 0.4-2$~\\micron), which at $z\\gsim 1$ are shifted into the observed-frame near-IR.  Several surveys have used deep, near-IR observations to search for galaxies up to $z\\sim 1- 2$ whose light is dominated by older stellar populations \\citep[see \\eg,][]{mccarthy04}. Such objects should have very red $R-K$ colors, which span the 4000\\AA/Balmer--break at these redshifts; several searches have focused on so-called extremely red objects (EROs), with selected objects typically satisfying $(R-K)_\\mathrm{Vega} \\gsim 5$~mag.  As a class, EROs include both passively evolving early-type galaxies up to $z\\sim 1$ \\citep[\\eg,][]{cimatti02} and dust-reddened starbursts \\citep[\\eg,][]{smail02}.  However, with only optical and near-IR photometry, it was not possible to uniquely interpret the majority of these objects \\citep[\\eg,][]{moustakas04,wilson07}.  Recent observations at longer wavelengths from \\spitzer\\ indicate that more than half of the ERO population have strong emission in the thermal IR \\citep{wilson04,yanl04b}, suggesting that such objects constitute a significant portion of the red-selected galaxy population.  Recently, observers have used deep near-IR data to study red-selected galaxies at higher redshifts.  \\citet{franx03} used deep VLT/ISAAC $JH\\ks$ observations to identify a population of ``distant red galaxies'' (DRGs) with $(J_s - \\ks)_\\mathrm{Vega} > 2.3$~mag, which should select galaxies that have a strong Balmer/4000~\\AA\\ break between the $J_s$ and \\ks\\ bands at $z\\sim 2$-3.5 \\citep[see also][]{saracco01}.  Like EROs, the DRG color selection is also sensitive to high redshift ($z\\gsim 1$) galaxies dominated by dust-enshrouded starbursts also have red observed $J-K$ colors \\citep[\\eg,][]{smail02}.  Indeed, subsequent study of DRGs has shown that the majority are massive galaxies that are actively forming stars at $z\\sim 1.5$-3 \\citep{vd03,fs04,rubin04,knudsen05,reddy05,papovich06,webb06}, and only a small subset appear to be completely devoid of star formation and passively evolving \\citep{labbe05,kriek06}.  However, nearly all surveys of red-selected objects have been restricted to using near-IR images from ground-based telescopes, which in practice limits these searches to $K_\\mathrm{AB}\\lsim 24$~mag \\citep[\\eg,][]{labbe03}.  For passively evolving stellar populations at $z\\sim 3$, this magnitude limit acts as a limit in \\textit{stellar mass} of $\\gsim 10^{11}$~\\msol.  In contrast, the $z\\sim 2-3.5$ UV-selected, star-forming galaxies (so-called Lyman-break galaxies, LBGs) that dominate the SFR density \\citep{reddy05} have typical stellar masses $\\sim 10^{10}$~\\msol\\ (for ``\\lstar'' LBGs; Papovich et al.\\ 2001).  The inferred ages for LBGs are generally less than 1~Gyr \\citep{shapley01,shapley05} --- significantly less than the age of the Universe at these epochs.  Thus, it is unknown if there exists a population of ``faded'' LBGs, identifiable as red, passively evolving galaxies with stellar masses $\\sim 10^{10}$~\\msol.  Deriving constraints on the density of such objects would improve our understanding of the evolution of these galaxies and on how galaxies assemble their stellar content at these early epochs.  However, detecting lower-mass, red galaxies at high redshift requires near-IR surveys with higher sensitivity than what is practically available from the ground.  Here, we use observations in the \\textit{Hubble} Ultra Deep Field \\citep[HUDF; Beckwith et al.\\ 2006 and][]{thompson05}.  This field has the deepest optical \\hst/ACS images to date, which combined with extremely deep, near-IR \\hst/NICMOS, and Spitzer (3.6, 4.5~\\micron) datasets form part of the Great Observatories Origins Deep Survey (GOODS).  The depth of the images in this field provides a means to explore high-redshift galaxies to a relatively low limiting stellar-mass.  Here, we focus on a sample of galaxies selected with red $\\nicj - \\nich$ colors in an effort to identify candidates for passively evolving galaxies at $z\\gsim 2$.  Similarly, \\citet{brammer07} use a $J-H$ color selection to identify red galaxies at $z > 2$ using the deep FIRES near-IR survey imaging.  The \\hst/NICMOS data for the HUDF allow us to probe objects with lower stellar masses than has been previously possible.  For example, \\citet{dickinson00} used deep NICMOS data in the \\textit{Hubble} Deep Field \\citep[HDF;][]{williams96} to study an unusually red ($\\nicj - \\nich \\gsim 2$~mag) source, possibly a passively evolving $z\\sim 3$ galaxy with $\\mcal > 3\\times 10^{10}$~\\msol.  Similarly, Chen \\& Marzke (2004) used the HUDF \\hst\\ images to identify red galaxies ($\\acsi-\\nich > 2$) with photometric redshifts $z_\\mathrm{ph} > 2.8$ (see also Yan et al.\\ 2004).  In this work we use the HUDF data to study the number density and mass density of objects at $2.5 < z < 3.5$ with colors consistent with older passively evolving stellar populations.  To study the HUDF $\\nicj - \\nich > 1$ sources, when possible we include GOODS IRAC 3.6-4.5~\\micron\\ and MIPS 24~\\micron\\ data, which allows us to better constrain the nature of the stellar populations in these objects (as we discuss below).  Throughout this paper we assume the following cosmological parameters: $\\Omega_M =$ 0.3, $\\Omega_{\\Lambda} =$ 0.7 and H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$ \\citep[consistent with the latest \\textit{WMAP} results, e.g.,][]{sp03}.  We use \\acsb\\acsv\\acsi\\acsz\\nicj\\nich\\ to denote magnitudes derived in the the \\hst\\ ACS and NICMOS bandpasses F435W, F606W, F775W, F850LP, F110W and F160W respectively.  We denote magnitudes derived from the IRAC channel 1 and channel 2 data as $[3.6\\micron]$ and $[4.5\\micron]$.  For reference, the effective wavelengths of these bandpasses are $0.43$, 0.60, 0.77, 0.91, 1.1, 1.6, 3.6 and 4.5~\\micron. All magnitudes henceforth are on the AB system, $m(\\mathrm{AB}) = 23.9 - 2.5\\log(f_\\nu / 1\\;\\ujy)$. ",
        "conclusions": "\\label{section:discussion}  We have investigated the physical properties of $32$ galaxies selected with $(\\nicj - \\nich) > 1.0$ colors in the HUDF using deep, broadband imaging from \\hst\\ and \\spitzer.  Exactly half of these objects (16/32) have unblended IRAC $[3.6\\micron]$ and $[4.5\\micron]$ photometry.  Here, we focus on the subset of galaxies with IRAC photometry because the data span a longer wavelength baseline, allowing us to separate more cleanly the objects into those whose colors are consistent with dust--enshrouded starbursts and those dominated by old stars.  Thus, where we derive number densities and mass densities, we include a factor of 2 correction for incompleteness due to blending in the IRAC image.  Furthermore, including both the \\spitzer/IRAC data allows us to better constrain the properties of the galaxies' stellar populations (see further discussion in, e.g., Labb\\'e et al.\\ 2005; Papovich et al.\\ 2006).  \\subsection{The Properties of Galaxies with Evolved Stellar Populations at $z\\gsim 2.5$}  Based on the full SED fitting, we split the subsample of objects with IRAC photometry into two broad categories: those objects whose broad-band SEDs are well represented by our two component model with a substantial component of evolved stars combined with a young star-forming component, and those objects whose interpretation changes substantially by including the IRAC photometry compared to the model fits to only the ACS and NICMOS data.  For 12/16 (75\\%) of the objects, including the IRAC data give roughly consistent best-fit models as to the ACS and NICMOS data only (NIC IDs 88, 219, 223, 269, 319, 625, 706, 921, 989, 1172, 1211, and 1237).  Because we are primarily interested in placing constraints on the number density of galaxies dominated by passively evolving stellar populations, we focus here on the subset of galaxies whose model fits have $z \\geq 2.5$ and $f_\\mathrm{old} \\geq 0.9$.  That is, the models favor a scenario where the vast majority of the galaxy's stellar mass resides in an older, passively evolving component.  There are 6 objects satisfying these conditions (NIC IDs 219, 223, 269, 625, 989, and 1211).  For three of these objects (NIC IDs 223, 625, 989), we have higher-confidence redshifts due to the presence of a Lyman ``break'' between in the $\\acsb -\\acsv$ colors.  This occurs for galaxies with relatively blue rest-frame UV colors, and a characteristic ``break'' in the colors as neutral hydrogen attenuation from the IGM absorbs the intrinsic flux shortward of 1216~\\AA\\ (e.g., Madau 1995).  For the ACS filters, absorption owing to IGM neutral hydrogen affects the \\acsb-band at $z \\gsim 3$, and the presence of this break drives the estimate of the redshifts for these galaxies ($z_\\mathrm{phot} \\simeq 3.5$, 3.2, and 3.1, respectively).  Interestingly, at $z\\sim 3-3.5$ the Balmer/4000~\\AA\\ is moving through the \\nich-band, diminishing the $J-H$ color. Therefore the fact that these galaxies have $\\nicj - \\nich \\geq 1$~mag implies they need a relatively large fraction of their stellar mass in older populations, and in all cases we find best fits with $f_\\mathrm{old} > 0.93$.  Chen \\& Marzke (2004) used a NICMOS-selected catalog in the HUDF combined with photometric redshifts to select six candidates for massive, red ($\\acsi - \\nich \\geq 2$) early-type galaxies at $z > 3$. Interestingly, for only 2 of their candidates where we have IRAC data does our analysis reach the conclusion that old stellar populations dominate the observed colors (NIC IDs 223, 625).  We find for three of their candidates (NIC IDs 219, 1078, 1237) that when the IRAC data are included the SED fits favor solutions where most of the light originates from dusty, young stellar populations.  For the remaining candidate of Chen \\& Marzke, we do not have IRAC data but our SED fit to the existing ACS + NICMOS photometry does favor a solution where the galaxy is dominated by older stellar populations.  For three of our six objects with IRAC photometry, $z\\geq 2.7$, and $f_\\mathrm{old} \\geq 0.9$ (NIC IDs 269, 989, and 1211), Chen \\& Marzke (2004) do not include these in their sample of evolved galaxies.  We derive a photometric redshift, $z=2.7$, for NIC ID 269, slightly below the $z > 2.8$ selection of Chen \\& Marzke, which may be the reason they excluded it.  Chen \\& Marzke interpret object NIC ID 1211 to be dominated by a dusty starburst.  However, the $\\nich$ to \\mone\\ and \\mtwo\\ colors imply that while its rest-frame UV light stems from dust-obscured young stars, most of the mass ($>$88\\% at 68\\% confidence) resides in evolved stellar populations (see figure~8; however, we note that the upper limit on the X-ray detection of this galaxy may imply that it hosts an AGN, see \\S~3.2).  NIC ID 989 appears to have some ongoing star formation (with relatively blue $\\acsb$ to $\\acsz$ colors), but older stars appear to dominate the rest-frame optical and near-IR light (figure~7).  We suspect this object was excluded from the selection of Chen \\& Marzke because of its relatively blue ACS colors.  Regardless, we find that data at wavelengths greater than rest-frame 1~\\micron\\ are required for interpreting the properties of the stellar populations of red galaxies.  Parenthetically, we note that there are two objects in the HUDF with $\\nicj - \\nich > 2$~mag (NIC IDs 1078, 1211) and one object has $\\nicj - \\nich = 1.95$~mag (NIC ID 1237).  All three have IRAC photometry. These objects have colors comparable to the ``unusual'' red, $\\nicj - \\nich$-selected object in the HDF-N reported by Dickinson et al.\\ (2000), which has $\\nicj - \\nich \\simeq 2.3$~mag.  Based on our fits to the ACS, NICMOS, and IRAC photometry of the HUDF objects, we conclude that two (NIC IDs 1078, 1237) are dominated by dusty starbursts, while only one (NIC ID 1211) is dominated by a substantial population of old stars.  Dickinson et al.\\ (2000) are unable to distinguish between these two possibilities (or, for that matter, the possibility that their object is a $z\\approx 12$ Lyman-break-type galaxy).  This is partly because the HDF-N object is undetected in the \\hst/WFPC2 data, nor did it have longer wavelength data from \\spitzer/IRAC.  Two of the three objects here have \\hst/ACS photometry $\\gsim 1$~mag fainter than the detection limit of the WFPC2 data in the HDF-N.  Therefore, we conclude that one requires \\textit{very} deep \\hst\\ optical imaging and \\spitzer/IRAC imaging to constrain the nature of these very red, $\\nicj - \\nich \\gsim 2$~mag objects.  None of our six candidates for galaxies with substantial population of evolved stellar stars is entirely devoid of recent star formation (i.e., all objects have $f_\\mathrm{old} < 1.0$).  Indeed, in the entire HUDF there are \\textit{no} candidates for galaxies consistent with a purely passively-evolving stellar population formed at much higher redshift (and in the NICMOS HDF-N, there is only one possible candidate, see discussion above; Dickinson et al.\\ 2000).  Therefore, such objects are extremely rare if they exist at this redshift.  We conclude that at $z\\gsim 2.5$ even galaxies with a substantial amount of evolved stellar populations still maintain some ongoing star-formation.  Papovich et al.\\ (2005) argue based on galaxies' morphologies that at $z\\gsim 2$, recurrent star-formation dominates the rest-frame UV and blue light, which will maintain some level of homogeneity in the galaxies' internal colors.  Our interpretation here is consistent with that scenario. However, we note that it appears that while the young stellar populations dominate the UV and blue light, in a small fraction of galaxies a substantial population of older stars dominates at rest-frame wavelengths $\\gsim 1$~\\micron.  \\begin{figure*}[t] \\begin{center} \\scalebox{0.48}{{\\includegraphics{fig9a.eps}}{\\includegraphics{fig9b.eps}}} \\caption{$\\nich$ magnitudes vs.\\ photometric redshifts for the subset of $(\\nicj-\\nich)_\\mathrm{AB} > 1.0$~mag objects with best-fit $f_\\mathrm{old} > 0.9$, for objects without IRAC photometry (left panel) and with IRAC photometry (right panel). The dashed line shows the $\\nich_\\mathrm{AB} < 26$ magnitude cut.  In the right-hand panel, cyan stars indicate those sources among the IRAC--detected subsample with $f_\\mathrm{old} > 0.9$ which have 24~\\micron\\ detections.  Thick--circles denote objects with $f_\\mathrm{old} > 0.9$ and $z_\\mathrm{phot} > 2.5$} \\label{fig:co} \\end{center} \\end{figure*}  \\begin{figure*}[t] \\begin{center} \\scalebox{0.48}{{\\includegraphics{fig10a.eps}}{\\includegraphics{fig10b.eps}}} \\caption{Best-fit model masses and photometric redshifts for all HUDF objects with $(\\nicj-\\nich)_\\mathrm{AB} > 1.0$~mag (left panel) and the subset of these with unblended IRAC photometry (right panel).  The dashed line shows the limiting stellar mass for galaxies with a stellar population formed in a burst at $z = 7$ with observed magnitude $\\nich = 26.0$.  In both panels the open circles indicate objects dominated by old stellar populations ($f_\\mathrm{old} > 0.9$).  In the right-hand panel, cyan stars indicate those sources among the IRAC--detected subsample with 24~\\micron\\ detections.  Thick--circles denote those objects with $f_\\mathrm{old} > 0.9$ and $z_\\mathrm{phot} > 2.5$, which are the discussion of \\S~4.3.} \\label{fig:co} \\end{center} \\end{figure*}  \\subsection{Galaxies with Evolved Stellar Populations at $z > 1.5$}  With the $\\nicj - \\nich > 1$~mag selection, we identify massive galaxies with substantial stellar mass and with SEDs dominated by the light from passively evolving stellar populations.  Figure~9 shows the measured $\\nich$ magnitudes as a function of redshift for galaxies with $f_\\mathrm{old} > 0.9$, for objects with IRAC photometry (right panel) and without IRAC photometry (left panel).  Although the data should be sensitive to such objects down to our selection limit, $H_{160} < 26$ mag, only one such object has $25 < H_{160} < 26$ mag --- the rest have $H_{160} < 25$ mag.  Figure~10 shows the derived stellar masses as a function of galaxy redshift for all galaxies in the $\\nicj - \\nich > 1$~mag sample.  The figure shows the stellar masses derived for the whole sample using only the ACS through NICMOS data (left panel), and it shows the stellar masses derived including the IRAC data for the IRAC--detected subsample (right panel).  Objects with more than 90\\% of their stellar mass in evolved stellar populations are circled in the figure.  There exist objects whose light (and stellar mass) is dominated by passively evolving stellar populations with total stellar masses up to and exceeding $>$$10^{11}$~\\msol\\ over the entire redshift range.  The existence of these galaxies is not entirely unexpected, given that such objects likely dominate the stellar mass density even at $z\\sim 2$ (e.g., Rudnick et al.\\ 2006; van Dokkum et al.\\ 2006).  However, there is an apparent dearth of objects whose light is dominated by passively evolving stellar populations with lower stellar masses.  For example, none of the 10 IRAC--detected galaxies with $z_\\mathrm{phot} > 1.6$ and with more than 90\\% of their stellar masses evolved stellar populations are fainter than $\\nich > 25$~mag (see figure~9), which is one magnitude above our detection limit.  (We note there is one object with $z_\\mathrm{phot} \\simeq 1.5$ and $H = 25.5$~mag, and a derived total stellar mass $4\\times 10^{9}$~\\msol, see figure~9.)  Therefore, at these high redshifts, the candidates for passively evolving galaxies are already massive.  In terms of total stellar mass, the right-hand panel of figure~10 shows that all the IRAC--detected objects with $f_\\mathrm{old} > 0.9$ and $z > 1.5$ have stellar masses greater than $\\sim$$2 \\times 10^{10}$~\\msol.  If we instead consider the results for galaxies from their stellar masses derived without the IRAC data, there is only a single candidate as a low-mass galaxy whose mass is dominated by passively evolved stellar populations (and we consider the stellar masses derived in the old component without IRAC data to be less robust).  The data should be sensitive to objects dominated by old stellar populations with stellar masses with $\\sim$$3\\times 10^{9}$ to $10^{10}$~\\msol (see figure 4 and figure 10).  We note that this is not a bias related to the sources with IRAC detections.  The IRAC data from GOODS (5.8 \\micron\\ flux density of 0.11~$\\mu$Jy, $5\\sigma$) would detect a passively evolving stellar population formed at z=7 and M $> 10^{9.5}$~$\\msun$ for all redshifts considered here.  Any plausible mass function for galaxies at $z\\sim 2$ would predict a greater number of galaxies with lower stellar masses than at high stellar mass.  For example, using the Fontana et al.\\ (2006) empirical galaxy mass function, we would expect roughly double the number of galaxies with $0.5 - 2.0 \\times 10^{10}$~\\msol\\ compared to those with $>$$2 \\times 10^{10}$~\\msol.  This strongly contrasts with the findings here for the red galaxies dominated by passively evolving stellar populations.  Therefore, the data suggest that galaxies with more than 90\\% of their stellar mass in evolved stellar populations have a minimum mass threshold of roughly several times $10^{10}$~\\msol\\ at $z\\sim 2$.  The lack of such objects with lower masses may be related to the processes governing stellar mass build--up in galaxies at these redshifts.   \\subsection{Constraints on the Densities of Galaxies with Evolved Stellar Populations at $2.5 < z < 3.5$}  With the full multiwavelength dataset, we have separated our sample into those galaxies that are dominated by dust-enshrouded starbursts and those galaxies whose rest-frame optical and near-IR emission is consistent with being dominated by a substantial population of evolved stars (see, e.g., section 4.2, above).  Of the 16 $\\nicj - \\nich > 1$~mag galaxies with robust IRAC photometry, 6 have best-fit models favoring a substantial population of evolved stellar populations, which is 37.5\\% of the sample.  Extending this to the full sample of 32 galaxies with $\\nicj - \\nich > 1$~mag, we estimate there should be 12 such objects in the area of the HUDF.  Using these numbers, we place constraints on the number and mass density for these types of objects.  If some of our 6 candidates for evolved stellar populations turned out to be dusty starbursts, our numbers would be reduced.  Due to the lower mass to light ratios of younger models, the fitted masses would also be reduced.  Therefore the densities derived here should be regarded as upper limits.  The HUDF data reach very deep limits in stellar mass for galaxies dominated by evolving stellar populations.  For a galaxy whose stellar mass formed \\textit{in situ} at $z_f=7$, the $\\nich = 26.0$~mag detection limit corresponds to a limit in stellar mass ranging from $6\\times 10^9$~\\msol\\ at $z=2$ to $1\\times 10^{11}$~\\msol\\ at $z=4$ (see figures~4 and 10).  Although the $(\\nicj - \\nich) > 1$~mag color selection should be, in principle, sensitive to old stellar populations at $z \\gsim 2$, in practice we find candidates for galaxies dominated by passively evolving stellar populations $f_\\mathrm{old} > 0.94$ at $z\\geq 2.7$.  Likewise, at $z > 3.5$, the Balmer/4000\\AA-breaks move beyond the F160W bandpass, making our color selection less efficient.  Therefore, to set an upper limit on the number and mass densities of galaxies dominated by passively evolving stellar populations, we limit our redshift range to $2.5 < z < 3.5$. We note that requiring an IRAC detection does not limit our stellar mass selection; the limiting stellar mass for the IRAC GOODS data sensitivity limit is always less than that for the NICMOS limit at these redshifts (see above).  The volume in the area of the HUDF is $1.9 \\times 10^{4}$~Mpc$^3$ between $2.5 < z < 3.5$.  Restricting our analysis to only those galaxies with $2.5 < z < 3.5$ we derive a number density of 3.1$\\times$$10^{-4}$~Mpc$^{-3}$, or 6.3$\\times$$10^{-4}$~Mpc$^{-3}$ including our 50\\% incompleteness.  However, these six objects are not visible over the whole redshift range.  In the number density calculation, we compute the effective volume, $V_\\mathrm{eff}$ for each galaxy from $z=2.5$ to $z_\\mathrm{eff}$, where $z_\\mathrm{eff}$ is the lesser of $z=3.5$ and the maximum redshift to which the galaxy could be detected.  The number density derived by summing $\\Sigma_i 1/V_\\mathrm{eff}^{(i)}$ over all $i$ galaxies gives $3.3^{+1.0}_{-1.5} \\times 10^{-4}$~Mpc$^{-3}$, and after accounting for incompleteness, $n = 6.6^{+2.0}_{-3.0} \\times 10^{-4}$~Mpc$^{-3}$.  Assuming these objects will evolve to become present-day, passively evolving red galaxies, then this number density is significantly lower than what we would expect using local mass functions.  Using the $z\\sim 0$ mass function parameters listed in table 4 of \\citet{bell03}, we expect the number density of red, early-type galaxies to be $n(z=0) = 3.5\\times 10^{-3}$~Mpc$^{-3}$ in the volume spanned between $2.5 < z < 3.5$ in the HUDF (and again taking the redshift-dependent mass limit from figure~\\ref{fig:masslimit})\\footnote{Note that \\citep{bell03} assumed a formation redshift, $z_f = 4$, in their mass functions.  This is only marginally lower than $z_f=7$ used in this analysis.  However, the change in lookback time from $z\\approx 0$ to 4 relative to that from $z\\approx 0$ to 7 corresponds to a change in the mass-to-light ratio of $M/L_B = 5.8$ to 6.2.  Applying our higher formation redshift to Bell et al.\\ would increase their stellar masses by $\\approx 5$\\%, which is negligible compared to other uncertainties.}.  This is larger than the observed number density (including incompleteness) by roughly a factor of 5.  Therefore, while we would have expected 65 objects in the HUDF (i.e., 32 after our $50\\%$ completeness) we have found only 6.  This difference is substantially larger than the Poisson noise ($\\pm \\sqrt{6}/6$, i.e., 40\\%).  It may be that these objects suffer strong field-to-field clustering (possibly expected for red galaxies, e.g., Daddi et al.\\ 2003; Quadri et al. 2007).  We estimate the number error from density variations in the volume between $2.5 < z < 3.5$ in the HUDF as $0.3\\times \\sigma_8$.  While we have no measurement for $\\sigma_8$ for these galaxies, based on theoretical models we expect $\\sigma_8 \\sim 1.0$ \\citep{adelberger05}, implying that the error from density variations is 30\\%.  Summing these errors in quadrature brings our total error estimate to $\\approx$50\\% for the expected number of these galaxies in the HUDF.  Even so, the difference between the 32 expected galaxies and the 6 we find is substantially larger than the combined errors from counting statistics and cosmic variance. Therefore, by number, and recalling that these candidates only define upper limits, the objects in the HUDF which are dominated by evolved stellar populations make up less than one-third of present-day red, early-type galaxies.  Thus, at least two-thirds of present-day red, early-type galaxies must assemble and/or evolve into their current configuration at $z \\lsim 2.5$.  Much of the stellar mass density in galaxies $2.5 < z < 3.5$ could reside in red-selected galaxies, whose light is dominated by older, passively evolving stellar populations. For the $\\nicj - \\nich > 1$~mag galaxies in the HUDF, we derive an upper limit on the stellar mass density of $\\log (\\rho^\\ast / [M_\\odot \\mathrm{Mpc}^{-3}] ) = 7.4\\pm{0.3}$, including the effect of our redshift-evolving mass limit, where our errors are derived from a Monte--Carlo simulation of the data.  When we include our 50\\% incompleteness, our total stellar mass density becomes $\\log (\\rho^\\ast / [M_\\odot \\mathrm{Mpc}^{-3}] ) = 7.7\\pm{0.3}$.  Poisson sampling of the 6 objects gives a factor of two error and therefore is the dominant source of uncertainty in the mass density estimate.  We consider this measurement to be an upper limit on the mass density of distant evolved galaxies for two reasons: a) it is possible that some of our evolved galaxy candidates are in fact dusty starbursts, and b) because our two-component model fit incorporates an a priori upper limit on the mass--to--light ratio in the form of a maximally old SSP component (see \\S 3).  \\citet{rudnick03} derive a stellar mass density at $z = 2.80^{+0.40}_{-0.39}$ of log $\\rho_* = 7.5^{+0.1}_{-0.1}$ and at $z = 2.01^{+0.40}_{-0.41}$ of log $\\rho_* = 7.5^{+0.1}_{-0.1}$.  Because our mass density provides an upper limit on the mass limit, we conclude that our derived mass density is consistent those of Rudnick et al.  In this case, it may be that passively stellar populations dominate the stellar mass density.  This is consistent with the analysis of Dickinson et al.\\ (2003) who used two--component star--formation history models to set an upper limit on the stellar mass density from passively evolving stellar populations hidden beneath the glare of younger stars in galaxies at $z < 3$.  Dickinson et al.\\ find that at $z\\sim 2.5-3$ an upper limit on the stellar mass density of $\\log (\\rho^\\ast / [\\msol \\mathrm{Mpc}^{-3}]) = 7.89^{+0.20}_{-0.15}$.  If this reflects reality, then based on our analysis, almost all of this mass density could reside in faint, red--selected galaxies at these redshifts.  \\citet{vd06} found that at $z\\sim 2$ red-selected galaxies with masses greater than $10^{11}$~\\msol\\ dominate the mass density.  Our analysis of the HUDF galaxies suggests that at lower stellar masses red-selected galaxies still contribute significantly to the total mass budget.  Lastly, we stress that with only rest-frame UV and optical data, it is not possible to discriminate between those galaxies dominated by dusty starbursts and older stellar populations (e.g., Moustakas et al.\\ 2004; Smail et al.\\ 2002).  Although the IRAC data does not resolve this degeneracy for all galaxies, we show that a sub-set of objects do have consistent best-fit results and we therefore use these objects in our analysis.  We wish to thank our colleagues at Steward Observatory for stimulating conversations.  In particular we would like to thank George Rieke for useful comments, Xiaohui Fan for generously providing models of the IGM attenuation, and Karl Gordon for his suggestions and help with the dust models used in this work.  AMS would also like to thank Richard Cool, Aleks Diamond-Stanic, Brandon Kelly, Juna Kollmeier, Andy Marble, John Moustakas, Jane Rigby and Yong Shi for helpful discussions.  Support for CJP was provided by NASA through the Spitzer Space Telescope Fellowship Program, through a contract issued by the Jet Propulsion Laboratory (JPL), California Institute of Technology under a contract with NASA.  DJE is also supported by an Alfred P.\\ Sloan Research Fellowship."
    },
    "0801/0801.0491_arXiv.txt": {
        "abstract": "We consider cosmological consequences of a heavy axino, decaying to the neutralino in R-parity conserving models. The importance and influence of the axino decay on the resultant abundance of neutralino dark matter depends on the lifetime and the energy density of axino. For a high reheating temperature after inflation, copiously produced axinos dominate the energy density of the universe and its decay produces a large amount of entropy. As a bonus, we obtain that the upper bound on the reheating temperature after inflation  via gravitino decay can be moderated, because the entropy production by the axino decay more or less dilutes the gravitinos. ",
        "introduction": "Neutralino, if it is the lightest supersymmetric particle (LSP) in R-parity conserving models, is a natural candidate for dark matter. Because of the TeV scale sparticle interactions, the thermal history of neutralinos allows the neutralino dark matter possibility. But, imposing a solution of the strong CP problem, the thermal history involves contributions from the additional sector.  The strong CP problem is naturally solved by introducing a very light axion $a$. Most probably, it appears when the Peccei-Quinn (PQ) symmetry is broken at a scale of $f_a$. Below the PQ scale, the effective axion interaction with gluons is \\dis{ {\\cal L} =\\frac{g_s^2}{32\\pi^2 f_a}a F \\tilde{F}, } where $g_s$ is the strong coupling constant~\\cite{Review}. The PQ scale is constrained by the astrophysical and cosmological considerations in the narrow window $10^{10} \\gev \\lesssim f_a \\lesssim 10^{12} \\gev$~\\cite{fabound}.   TeV scale supersymmetry (SUSY) suggests axino $\\tilde a$, the superpartner of axion, around the electroweak scale in the gravity mediation scenario. Here, we consider the effects of {\\it heavy} axinos in cosmology. The axino cosmology depends crucially on the axino decoupling temperature \\cite{Rajagopal:1990yx}, \\dis{ \\tadec =10^{11}\\gev \\left(\\frac{f_a}{10^{12}\\gev} \\right)^2\\left(\\frac{0.1}{\\alpha_s} \\right)^3,\\label{T_dec} } where $\\alpha_{s}=g_s^2/4\\pi$.%  The axion supermultiplet includes axion, saxion (the scalar partner) and axino. Both saxion and axino masses are split from the almost vanishing axion mass if SUSY is broken. The precise value of the axino mass depends on the model, specified by the SUSY breaking sector and the mediation sector to the axion supermultiplet~\\cite{susyreview}. In principle, the axion supermultiplet is independent from the observable sector in which case we may take the axino mass as a free parameter of order from keV to a value much larger than the gravitino mass \\cite{axinomass,ChunLukas}. Light axinos can be a dark matter (DM) candidate, which has been studied extensively \\cite{axinoDM,Roszkowski:2006kw,Asaka:2000ew}. Heavy axinos, however, cannot be the LSP and can decay to the LSP plus light particles. This heavy axino decay to neutralino was considered in the literature \\cite{ChunLukas} where the neutralino relic density was not considered seriously. Some considered the axino as the next LSP decaying to the gravitino LSP in the gauge mediated SUSY breaking scenario \\cite{Chun:1993vz}. Recently, supersymmetric axion models were studied with an emphasis on saxion \\cite{Kim:1992eu}, where the heavy axino possibility was also considered briefly \\cite{Kawasaki:2007mk}.  In this paper, we present a more or less complete cosmological analysis of a {\\it heavy} axino with mass in the TeV region so that it is heavier than the LSP neutralino. Compared with the saxion study, the heavy axino study probes SUSY directly because the axino carries the odd R-charge.   If kinematically allowed, a heavy axino decays predominantly to a gluino and a gluon and the gluino subsequently decays, finally producing the neutralino LSP. However, if axino is lighter than gluino, this axino to gluino decay is forbidden and the axino predominantly decays to a neutralino and a photon. The neutralinos from axino decay can annihilate in the cosmos if the neutralino number density, $n$, is large enough so that $n \\langle \\sigma_{ann} v\\rangle > H$ where $\\sigma_{ann}$ is the annihilation cross section. In this case, the neutralino abundance is modified, which is obtained by solving the Boltzmann equation. For this to be compatible with the observed DM density, we find the required thermally-averaged cross section of neutralino, $ \\langle \\sigma_{ann} v\\rangle $, to be around $10^{-8}\\gev^{-2}$. The Higgsino-like neutralino can give this kind of large cross section.  The thermally produced ${\\cal O}(100 \\gev)$ gravitinos after inflation induce severe problems on the light element abundances, which restricts the reheating temperature to $\\treh< 10^{6-7} \\gev$. We find that the entropy production from the heavy axino decay can dilute the primordial gravitinos. For axino mass smaller than gluino mass and for a low PQ scale, $f_a\\sim 10^{10}\\gev$, we find as shown below that there is no gravitino problem for the reheating temperature up to the axino decoupling temperature $\\tadec$.  A heavy axino leads to different physical consequences depending on its mass being greater or smaller than the gluino mass $m_\\gluino$. If axino is heavier than gluino, it decays dominantly to a gluino plus a gluon.  If axino is lighter than gluino, it decays dominantly to a $b$-ino-like neutralino and a photon and, if kinematically allowed, to a neutralino and a $Z$-boson. Thus, the axino decay width is given by,  \\noindent (i) $m_{\\tilde a}>m_\\gluino$ \\dis{ \\Gamma(\\axino \\rightarrow \\gluino + g)=\\frac{8\\alpha_{s}^2}{128\\pi^3} \\frac{m_\\axino^3}{f_a^2} \\left(1-\\frac{m_\\gluino^2}{m_\\axino^2} \\right)^{3}, \\label{Gamma_axino} } \\noindent (ii) $m_\\chi<m_{\\tilde a}<m_\\gluino$: \\begin{equation} \\Gamma(\\axino \\rightarrow \\chi_i + \\gamma)= \\frac{\\alpha_{em}^2C^2_{a\\chi_i\\gamma}}{128\\pi^3} \\frac{m_\\axino^3}{f_a^2} \\left(1-\\frac{m_{\\chi_i}^2}{m_\\axino^2} \\right)^{3}, \\end{equation} with $C_{a\\chi_i\\gamma}=(C_{a Y Y}/\\cos\\theta_W)Z_{\\chi_i B}$~\\cite{axinoDM}, where $Z_{\\chi_i B}$ is the $b$-ino fraction of the $i$-th neutralino and  $\\theta_W$ is a Weinberg mixing angle. Hereafter we will use $C_{a Y Y}=Z_{\\chi_i B}=1$ for simplicity. For the case (ii), as one can see, the lifetime can be easily longer than $0.1$ second. If the lifetime is shorter than about $0.1$ second, the axino decay does not harm the standard big bang nucleosynthesis (BBN). Note that for Eq.~(\\ref{Gamma_axino}) the axino lifetime is about $3.3\\times 10^{-7} \\textrm{s} \\left({\\alpha_{s}}/{0.1} \\right)^{-2} \\left({f_a}/{10^{11} \\gev} \\right)^2 \\left({m_\\axino}/{1 \\tev} \\right)^{-3}$.  Right after the axino decay, the temperature of radiation $\\td= { \\sqrt{\\Gamma_\\axino M_P}}/{\\left(\\pi^2g_*/90\\right)^{1/4}} $ is given by \\dis{ \\td = 1.4\\ {\\rm GeV} \\left(\\frac{70}{g_*} \\right)^{1/4} \\left(\\frac{3 \\times 10^{-7} {\\rm sec}} {\\tau_\\axino}\\right)^{1/2}, \\label{TD} } where $M_P$ is the reduced Planck mass. For the case that axinos dominate the energy density of the universe before they decay, we can regard $\\td$ as the second reheating temperature due to the axino decay. This temperature is marginal when we discuss its effect on neutralino DM as we discuss below. This is because, for the axino lifetime of $10^{-7}$ second, this temperature is comparable to the typical neutralino freeze-out temperature $T_{fr} \\approx m_{\\chi}/25$ \\cite{KolbTurner}. If $\\td > T_{fr}$ or equivalently $\\tau_{\\tilde{a}}<{\\cal O}(10^{-7})$ second, the axino decay has no effect on the neutralino DM abundance, because the neutralino produced by the axino decay also could reach the thermal equilibrium. From Eq.~(\\ref{Gamma_axino}), we can see that this would be the case for a heavier axino or for a lower PQ scale.  The axino abundance depends on the thermal history of the early universe after inflation. Another relevant temperature we introduce is the reheating temperature after inflation $\\treh$. So, the temperatures we introduce are \\dis{ &\\tadec={\\rm axino\\ decoupling\\ temperature}\\\\ &\\treh={\\rm reheating\\ temperature\\ after\\ inflation}\\\\ &T_{fr}={\\rm neutralino\\ freeze-out\\ temperature}\\\\ &\\tarad={\\rm axino-radiation\\ equality\\ temperature}\\\\ &\\td = {\\rm radiation\\ temperature\\ right\\ after\\ \\tilde {\\it a}\\ decay}. \\label{Temps} }  For $\\treh>\\tadec$, axinos were in the thermal equilibrium and the axino number density is comparable to the photon number density. Below the axino decoupling temperature, the axino number in the comoving volume is conserved if there is no new physics below $\\tadec$; then the axino abundance is given by \\dis{ Y_\\axino=\\left.\\frac{n_\\axino}{s}\\right|_{\\tadec} =\\frac{135\\zeta(3)}{8\\pi^4} \\frac{g_\\axino}{g_*(\\tadec)},\\quad \\textrm{for}\\quad \\treh>\\tadec, } where $\\zeta(3)\\simeq1.202$, $g_\\axino$ is the degrees of freedom of axino and $s=\\frac{2\\pi^2}{45}g_{*s}T^3$ is the entropy density.  For $\\treh<\\tadec$, axinos could not be in thermal equilibrium after inflation. Nevertheless, they are regenerated by thermal scattering and by decays of gluinos, squarks and neutralinos in the thermal plasma. For a high enough reheating temperature, $T\\gg 10^{4}\\gev$, the hard thermal loop approximation is good enough and the axino abundance can be approximated as \\cite{Brandenburg:2004du} \\dis{ Y_{\\tilde{a} } = 2.0\\times 10^{-7}g_s^6 \\ln \\left( \\frac{1.108}{g_s}\\right)\\left(\\frac{10^{11} \\gev}{f_a} \\right)^2\\left( \\frac{T_R}{10^4\\gev}\\right). } We can see that the axino abundance might be large in terms of $Y_{\\tilde{a}}$ going up to ${\\cal O}(0.002)$, strongly depending on $\\treh$ and $f_a$. Hence, if many axinos are produced after inflation, axinos might dominate the energy density of the universe before they decay.  Let us first examine the condition for axino domination. For $T < m_{\\tilde{a}}$, axino becomes nonrelativistic. When axinos are nonrelativistic, the energy densities of radiation and axinos  are given by \\begin{eqnarray} \\rho_R = \\frac{\\pi^2 g_*}{30} T^4 ,\\quad \\left.\\rho_{\\tilde{a}} \\right|_{T<m_{\\tilde{a}}}= m_{\\tilde{a}} Y_{\\tilde{a}} s  \\label{axino:Tma}, \\end{eqnarray} respectively. From Eq.~(\\ref{axino:Tma}), we find the axino-radiation equality temperature, $\\tarad$, \\begin{equation} \\tarad\\equiv \\frac{4}{3} m_{\\tilde{a}} Y_{\\tilde{a}}> \\td. \\label{DominationCondition} \\end{equation} If this equality occurs before axino decays as shown as the inequality, then axino can dominate the universe. On the other hand, for $ \\frac{4}{3} m_{\\tilde{a}} Y_{\\tilde{a}} < \\td$, axinos never dominate the energy density of the universe before they decay. In this regard, we find the lowest reheating temperature, $\\treh^{min}$, above which axinos can dominate the universe before they decay, by solving the equality $ \\frac{4}{3} m_{\\tilde{a}} Y_{\\tilde{a}}(\\treh^{min}) = \\td. $  In case axinos dominate the universe, the entropy production by axino decay dilutes the previously existing number densities. The ratio of the entropy per comoving volume before and after the axino decay is \\cite{KolbTurner} \\dis{ r\\equiv\\frac{S_{f}}{S_{0}}\\simeq \\frac{4 m_{\\tilde{a}} Y_{\\tilde{a}}}{3 \\td} , \\label{EntropyProduction} } for $ \\tarad > \\td $. Thus, from Eqs.~(\\ref{TD}) and (\\ref{EntropyProduction}), the entropy ratio can be \\dis{ r\\equiv\\frac{S_{f}}{S_{0}} \\simeq 1.3 \\left(\\frac{m_\\axino}{1\\tev}\\right) \\left(\\frac{2\\gev}{\\td}\\right) \\left(\\frac{Y_\\axino}{0.002}\\right) \\label{entropyratio}. } This is shown in Figs.~\\ref{fig_fa10} and~\\ref{fig_fa12} as magenta lines. Above the solid line denoted as $r=1$, axinos can dominate the universe and can produce additional entropy.  Similarly, in case that an axino can decay only into a neutralino and a photon [Case (ii)], $T_{\\rm D}$ looks to be able to become as small as of several MeV and $S_{f}/S_{0}$ might be as large as ${\\cal O}(10^2-10^3)$. ",
        "conclusions": "We have discussed the heavy axino possibility so that the heavy axino decay leads to a reasonable DM density of the LSP neutralino and enough radiation to dilute the gravitinos. For this scenario to be realized, we need a large annihilation cross section of neutralinos $\\langle\\sigma v \\rangle$. Such a neutralino is Higgsino-like, in which case the cross section with nuclei for the direct DM search can be as large as $\\sigma_{SI} \\simeq 10^{-(7-8)}$ pb, which can be probed in future DM search experiments.  \\vskip 0.5cm"
    },
    "0801/0801.2177_arXiv.txt": {
        "abstract": "We show that space-based telescopes such as the proposed Terrestrial Planet Finder Coronagraph will be able to detect the light scattered by the interstellar grains along lines of sight passing near stars in our Galaxy.  The relative flux of the scattered light within one arcsecond of a star at 100~pc in a uniform interstellar medium of 0.1 H~atoms cm$^{-3}$ is about $10^{-7}$.  The halo increases in strength with the distance to the star and is unlikely to limit the coronagraphic detection of planets around the nearest stars.  Grains passing within 100~AU of Sun-like stars are deflected by radiation, gravity and magnetic forces, leading to features in the scattered light that can potentially reveal the strength of the stellar wind, the orientation of the stellar magnetic field and the relative motion between the star and the surrounding interstellar medium. ",
        "introduction": "}  Scattered starlight is a valuable tool for measuring the properties of interstellar dust.  In optical reflection nebulae \\citep{wp92,c95,b02}, dark clouds \\citep{fs76,wo90} and the diffuse Galactic light \\citep{hb91,mh95,wf97} the scattering grains typically lie many arcseconds from the illuminating sources.  The dust can be nearer to or further from us than the source and the detected photons have been deflected through a wide range of angles.  Further information about interstellar grains comes from stellar X-rays scattered by foreground dust within a few arcseconds of the star \\citep{m95,sd98,ws01,dt03}.  The measurements sample the whole column of material between us and the star because the grains are strongly forward scattering at X-ray energies.  Here we show that high-contrast imaging with a space-based telescope such as the proposed Terrestrial Planet Finder Coronagraph enables a new kind of measurement for probing dust at projected separations similar to the X-ray measurements.  The technique samples a range of scattering angles because at visible wavelengths the scattering is only mildly biased in the forward direction.  As a consequence, most of the light comes from foreground grains lying close to the star. The scattered light halo carries information on the interstellar dust passing through the cavity opened in the interstellar gas by the stellar wind.  Evidence for the existence of interstellar grains within our own heliosphere was obtained by the Ulysses spacecraft when 55 dust impacts as measured by the Cosmic Dust Experiment were identified as interstellar by their speed, mass and arrival direction \\citep{gg94}. The overall contribution of interstellar dust to the zodiacal cloud --- the debris disk of the solar system --- is unknown, although the fraction is presumably greater in the outer solar system.  Collisional debris from the asteroid belt dominates inside 3~AU \\citep{gd01}. However the interstellar grains, charged by photoionization and strongly influenced by the solar gravity, radiation pressure and magnetic field, are able to penetrate deep into the inner solar system.  A uniform incoming spatial distribution of interstellar grains becomes strongly clumped as a function of particle size and phase of the solar cycle, although the contribution of this component to the all-sky thermal emission as viewed from 1~AU is negligible due to the local dominance of asteroidal and cometary material \\citep{gd96}.  Starlight scattered by interstellar dust also is a source of noise in coronagraphic planet searches.  We calculate the expected brightness and distribution of the scattered light and show that the halo is unlikely to affect the direct detection of Earth-like planets around the nearest stars.  We compute the trajectories of interstellar grains under the stellar gravity, radiation and magnetic forces and show that the stellar wind produces observable signatures in the scattered light.  For a large number of stars at distances greater than 100~pc, high-contrast imaging can potentially yield detailed information about the stellar wind and the adjacent interstellar medium.  The remaining five sections of the paper cover the radiative transfer method (\\S2), the results for uniformly-distributed dust, obtained analytically assuming isotropic scattering (\\S3) and numerically including the anisotropy (\\S4), the results for dust passing through a model stellar wind (\\S5), and a summary and conclusions (\\S6). ",
        "conclusions": "}  We have shown that stars appear surrounded by halos of light scattered by interstellar dust.  The halos are bright enough for detection with the proposed Terrestrial Planet Finder Coronagraph around stars at distances of 100~pc and greater.  The halo reveals the distribution of dust in the vicinity of the star, providing opportunities to probe the interstellar medium and measure the properties of the stellar wind. Particles of different sizes have different distributions, with larger particles approaching the star most closely \\citep{gd96}.  This size-sorting can potentially be used to measure the distribution of grain sizes at selected locations in the interstellar medium and to resolve whether particles larger than one micron are present in significant numbers, as suggested by dust impact detector data from the Ulysses and Galileo spacecraft \\citep{fd99,wd01}.  Since the grains are deflected by stellar wind magnetic forces, the features in the scattered light halo can be used to estimate the wind speed and the relative motion between the star and the surrounding interstellar medium.  The results will complement the wind mass loss rates estimated using the Lyman-$\\alpha$ absorption by the gas collected near stellar wind bow shocks \\citep{wm05}.  The scattering halos will make it possible in addition to determine the wind magnetic field strength and orientation (figure~\\ref{fig:halo}) and, by observations spanning a decade or so, the variation over the stellar activity cycle \\citep{l00}.  Imaging the halos will yield fresh information about the winds of stars with a range of ages and could lead to a better understanding of stellar spin histories \\citep{bf97} and the long-term evolution of stellar activity."
    },
    "0801/0801.4492_arXiv.txt": {
        "abstract": "{{\\chandra} and {\\xmm} surveys show that the fraction  of obscured AGN decreases rapidly with increasing luminosity. Although this is usually explained by assuming that  the covering factor of the central engine is much smaller at luminous QSOs, the exact origin of this effect remains unknown.  We perform toy simulations to test whether photo-ionisation of the obscuring screen in the presence of a strong radiation field  can reproduce this effect. In particular, we create X-ray spectral simulations using  a warm absorber model assuming a range of input column densities and ionisation parameters. We fit instead the simulated spectra with a simple cold absorption power-law model that is the standard practice in X-ray surveys. We find that the fraction of absorbed AGN should fall with luminosity as $L^{-0.16\\pm0.03}$ in rough agreement with the observations. Furthermore, this apparent decrease in the obscuring material is consistent with the dependence of the FeK$\\alpha$ narrow-line equivalent width on luminosity, ie. the X-ray Baldwin effect. { ",
        "introduction": "The deepest {\\chandra} X-ray surveys find a high surface density of  AGN (Alexander et al. 2003, Giacconi et al. 2002). At the faintest fluxes probed, $\\rm f_{X}(2-10 keV)\\approx 10^{-16}$ \\funits, the large majority of these sources ($\\sim$ 80\\%) are obscured, having a hydrogen column density of $\\rm N_H>10^{22}$ \\cunits (e.g. Akylas et al. 2006; Tozzi et al. 2006). The high fraction of obscured AGN witnessed by {\\chandra} is well in agreement with previous X-ray background synthesis models (e.g. Comastri et al. 1995, 2001). However, the situation is drastically different at the brighter fluxes probed by the {\\xmm} surveys. These show a marked deficit of X-ray obscured AGN (Piconcelli et al. 2002;  Perola et al. 2004; Georgantopoulos et al. 2004;  Caccianiga et al. 2004). The fraction of the obscured AGN drops to $\\sim30\\%$ at fluxes of $\\rm f_{X}(2-10~keV)\\approx 10^{-14}$ \\funits. In flux limited samples, a small increase in the obscured AGN fraction with decreasing observed flux is expected, as the absorbed sources present lower flux. However the observed increase is much steeper than that predicted (eg. Piconcelli et al. 2002).  This discrepancy has been resolved with recent {\\chandra} and {\\xmm} observations (e.g. Akylas et al. 2006; La Franca et al. 2005; Ueda et al. 2003) which demonstrate that the X-ray obscured AGN fraction is lower at higher luminosities. This  can also explain the deficit of type-2 narrow-line QSOs at high redshift (Steffen et al. 2003). Recent X-ray background synthesis models that take into account the dependency of the obscured AGN fraction on luminosity (Gilli, Comastri \\& Hasinger 2007) are very successful in predicting the observed flux distribution.  The important question is what is the physical mechanism that can produce the decrease in the fraction of obscured AGN. It is possible that in the presence of a strong radiation field the fraction covered by the torus may be reduced. K\\\"{o}ningl \\& Kartje (1994) for example proposed a disc-driven hydromagnetic wind model of the torus where for high bolometric luminosities the radiation pressure is expected to flatten the torus. A similar model, the receding torus model,  has been invoked by Lawrence (1991) and Simpson  (2005), to explain the fraction of radio AGN optically classified as type-2. In the presence of a strong radiation field the dust is heated and sublimates when it reaches a temperature of about 1500-2000K. At high luminosities the sublimation radius increases and thus effectively the torus opening angle increases, resulting in a higher number of type-1 AGN.  A similar effect is expected to take place in X-ray wavelengths. The X-ray absorbing gas is situated close to the nucleus as inferred from X-ray variability studies of type-2 AGN (Risaliti, Elvis \\& Nicastro 2002). Thus it is reasonable to assume that a strong radiation field will also ionise an appreciable fraction of the surrounding obscuring screen, resulting in the decrease in the ``effective'' column density. In this paper, we perform toy X-ray spectral simulations to estimate the strength of this effect. We also discuss any possible connection with the X-ray Baldwin effect. ",
        "conclusions": "We employ X-ray spectral simulations to determine the effect of photo-ionisation on the obscuring torus in AGN. In particular,  our goal is to explore whether photo-ionisation can explain the steep decrease in the fraction of obscured AGN with increasing luminosity as witnessed in X-ray surveys. Our conclusions can be summarised as follows:  \\begin{itemize} \\item{Photo-ionisation can reproduce the observed decrease in the absorbed AGN fraction as a function of luminosity. We find that the fraction of absorbed AGN should fall with luminosity as $L^{-0.16\\pm0.03}$ in rough agreement with the observations.}  \\item{The observed decrease in the obscuring material covering factor is roughly consistent  with the dependence of the Fe K$_{\\alpha}$ line equivalent width on luminosity, the X-ray Baldwin effect.}  \\item{A simple  prediction of the photo-ionised model is that a high fraction of  luminous QSOs should present warm absorption features  in their X-ray spectra. Future, high effective area mission, such as {\\it XEUS} will be able to quantitatively assess this scenario.}  \\end{itemize}"
    },
    "0801/0801.1081_arXiv.txt": {
        "abstract": "The colour-magnitude diagrams of resolved single stellar populations, such as open and globular clusters, have provided the best natural laboratories to test stellar evolution theory. Whilst a variety of techniques have been used to infer the basic properties of these simple populations, systematic uncertainties arise from the purely geometrical degeneracy produced by the similar shape of isochrones of different ages and metallicities. Here we present an objective and robust statistical technique which lifts this degeneracy to a great extent through the use of a key observable: the number of stars along the isochrone. Through extensive Monte Carlo simulations we show that, for instance, we can infer the four main parameters (age, metallicity, distance and reddening) in an objective way, along with robust confidence intervals and their full covariance matrix. We show that systematic uncertainties due to field contamination, unresolved binaries, initial or present-day stellar mass function are either negligible or well under control. This technique provides, for the first time, a proper way to infer with unprecedented accuracy the fundamental properties of simple stellar populations, in an easy-to-implement algorithm. ",
        "introduction": "The theory of stellar evolution has been tested extensively mainly through well-studied binary stars ({\\sl e.g.}, Popper 1980, Lastennet \\& Valls-Gabaud 2002, Ribas 2006) or simple, coeval stellar populations such as open and globular clusters ({\\sl e.g.} Vandenberg et al. 1996, Renzini \\& Fusi Pecci 1988). It has been extraordinarily successful at explaining most stellar properties across the different evolutionary stages, by incorporating increasingly complex physics in both stellar interiors and atmospheres (see Salaris \\& Cassisi 2006 for an updated and extensive review). In contrast, and perhaps somewhat paradoxically, the determination of the fundamental parameters of single stellar populations, such as distance, age, metallicity, reddening, convective parameters, etc, has seldom been the subject of a rigorous mathematical determination. More often than not, isochrones are ``fitted'' by eye, and very rough estimates of errors are given, if given at all. The reason for this state of affairs is perhaps the many degeneracies between these parameters, which makes it seemingly impossible to find a unique solution to the set of fundamental parameters.  Alternative, semi-empirical methods have been developed instead to infer {\\sl relative} quantities, rather than absolute ones. For instance, the so-called horizontal method relies on the difference in colour between the turn-off (TO) point and the base of the red giant branch (RGB), and is independent of the assumptions made on the convective theory. The vertical method, on the other hand, relies on the difference in apparent magnitude between horizontal branch (HB) stars and turn-off stars. Both methods rely heavily on ``semi-empirical calibrations'' provided by stellar models, and are subject to many uncertainties ({\\sl e.g.}, How to define properly the TO point when photometric errors blur this region of the colour-magnitude diagram ? What is the vertical offset HB--TO when the HB is not horizontal ?) which nevertheless have been tackled with some success ({\\sl e.g.} Rosenberg et al. 1999, Salaris \\& Weiss 2002).  Here we want to point out that many of the degeneracies are actually mere artifacts, and that a proper and rigorous mathematical technique can be formulated to measure these parameters in a robust and objective way. For example, the well-known age-metallicity degeneracy, which states that a given isochrone can be fitted with a different age provided the metallicity is changed appropriately, is a purely {\\sl geometrical} degeneracy : the {\\sl shape} of the isochrones is the same, yet stellar evolution predicts that stars with a different metal content will evolve at a different {\\sl speed}. Hence, the age-metallicity degeneracy is not a physical one, but simply a result of using the geometrical shape of isochrones as the only discriminant factor. Clearly, a simple statistical estimator that would count the number of stars {\\sl along} the isochrones would give very different results for any pair of age/metallicity values that would give the same geometrical shape to the isochrones. In a more formal way, let us consider a curvilinear coordinate $s$ along an isochrone of given parameters (say, distance, reddening, metallicity, age, alpha elements enhancement, etc). This position depends only on the age $t$ of that isochrone and on the mass $m$ of the star at that precise locus, so that $s=s(m,t)$ only. An offset in age $dt$ is reflected only through changes in mass $m$ and position $s$ along the isochrone of age $t$ since \\beq dt(m,s) \\; = \\; \\left.\\frac{ \\partial t }{\\partial m}\\right|_s \\, dm \\; + \\; \\left.\\frac{ \\partial t }{\\partial s}\\right|_m \\, ds \\qquad . \\eeq For a star in this isochrone the offset is, by definition, zero, hence \\beq \\left.\\frac{ \\partial m }{\\partial s}\\right|_t \\; = \\; - \\; \\left.\\frac{ \\partial m }{\\partial t}\\right|_s \\; \\times \\; \\left( \\left.\\frac{ \\partial s }{\\partial t}\\right|_m \\right)^{-1}   \\quad . \\label{eq2} \\eeq The first term is always finite. The second term is the evolutionary speed, the rate of change for a given mass $m$ of its coordinate along the isochrone when the age changes by some small amount. One can think of $s$ as representing some evolutionary phase, and so this term will be large when the phase is short-lived : a small variation in age yields a very large change in position along the isochrone. Alternatively, for a given age, and since the first term is always finite, a wide variation in position implies a narrow range in mass. This is the case of the red giant branch or the white dwarf cooling sequence\\footnote{excluding the bottom of the cooling sequence, where white dwarfs having a large range of progenitor masses pile up}, for instance. On the other hand, slowly evolving phases such as the main sequence have small evolutionary speeds and wide ranges in mass for a given interval along an isochrone. Clearly, the most important phases to discriminate between alternative ages and metallicities will be the post main-sequence ones, where the range of mass is small (and hence insensitive to the details of the stellar mass function), and at the same time where evolutionary speeds are large. If we consider the mid- to lower main sequence, at fixed metallicity, isochrones of all ages trace the same locus, with only marginal changes in the density distribution of points amongst them. In this sense, one of the parameters, $t$, is to a large extent absent from the main sequence, while phases beyond are always substantially a function of both age and metallicity. The density of stars along an isochrone is therefore \\beq \\frac{d N}{d s} \\; = \\; - \\left(\\frac{dN}{dm}\\right) \\times \\left( \\left. \\frac{\\partial m}{\\partial t}\\right|_s \\right) \\times \\left( \\left. \\frac{\\partial t}{\\partial s}\\right|_m \\right) \\qquad . \\eeq The first term is related to  the initial stellar mass function (IMF), and, as discussed above, the second term is finite while the third is a strong function of the evolutionary speed. If the mass after the turn-off is assumed to be roughly constant, this implies that the {\\sl ratio} in the number of stars in two different evolutionary stages after the turn-off will only depend on the ratio of their evolutionary time scales\\footnote{In the context of stellar population synthesis, this is known as the fuel consumption theorem (Renzini \\& Buzzoni 1983).}.  In terms of observables, this well-know relation has been exploited using luminosity functions as a proxy (as pioneered by Paczy\\'nski 1984), that is, number counts of stars as a function of apparent magnitude. This is not the isochrone coordinate $s$, but rather the projection of the isochrone on the vertical axis of the colour-magnitude diagram. The fact that the sub-giant branch often appears nearly horizontal in CMDs implies that luminosity bins in this regime will not be discriminant. This is unfortunate because this phase is more heavily populated than the red giant branch, and thus the use of luminosity functions is not optimal. Likewise, the projection on the horizontal axis, a colour, is less sensitive to the details along the RGB. Clearly, the optimal strategy is to use the full information provided by the observed density of stars along the isochrone.  To summarize existing methods, one approach has grown out of the fitting of 'geometrical' indicators of the observed CMD to isochrones, from single particular indicators such as turn off point (TO), magnitude or colour difference between two points on the CMD, TO vs. zero-age horizontal branch or tip of the RGB and RGB bump position and combinations thereof. Examples can be found in Renzini \\& Fusi Pecci (1988), Buonanno et al. (1998), Salaris \\& Cassisi (1998), Salaris \\& Weiss (1998), VandenBerg \\& Durrel (1990), Ferraro et al. (1999) and Rosenberg et al. (1999). The accuracy of these approaches has been extended by including several such geometrical indicators simultaneously (e.g. VandenBerg 2000, Meissner \\& Weiss 2006), or a complete geometrical comparison between observed CMD and isochrone (e.g. Straniero \\& Chieffi 1991). The reduction of the CMD to a few numbers, given the varied and complex physics which enters in determining them implies using only a subset of the information available in the CMD, while the unavoidable observational errors present necessarily lead to ambiguities in the identification of the key points to be used, which are hard to quantify formally, and which result in uncertain confidence intervals being assigned to the inferences derived. Neglecting stellar evolutionary effects limits the information content of the isochrones against which the CMDs are compared.  On the other hand, approaches concentrating explicitly on the stellar evolutionary effects have grown out of the already mentioned analysis of luminosity functions presented by Paczy\\'nski (1984), a projection of the CMD onto the vertical axis. Extensions have mostly centered on the easy to measure and age sensitive features of the RGB luminosity function ({\\sl e.g.} Jimenez \\& Padoan, 1998, and Zoccali \\& Piotto 2000). Here again, observational errors together with the binning adopted in generating luminosity functions, imply a limited use of the information available in the observation, while not including the information encoded by the geometry of both the CMD and isochrones limits the accuracy of the inferences.  There are have been a few attempts in the past at using the full information available in a colour-magnitude diagram, as pioneered by Flannery \\& Johnson (1982) and further developed by Wilson \\& Hurrey (2003) and Naylor \\& Jeffries (2006). Unfortunately all these approaches, while mathematically correct, fail to include properly the key discriminating factor discussed above, the stellar evolutionary speed, resulting in wide uncertainties in the inferred parameters. The only other attempt we are aware of, in the context of inferring parameters from a stellar population, was made by J{\\o}rgensen \\& Lindegren (2005), who adapted our formalism developed in a previous paper (Hernandez, Valls-Gabaud \\& Gilmore 1999) to derive stellar ages. The present paper is a further and robust extension of this approach. Paper II in this series (Valls-Gabaud \\& Hernandez 2007) will apply the method to different sets of observations of stellar populations.  In section 2 we present the rigorous probabilistic model, which is tested with synthetic cases in Section 3. In section 4 we explore and quantify the effects of systematic uncertainties, in section 5 we present an application of the method to a real case, NGC~3201.  Finally in section 6 we present our conclusions on this novel approach. ",
        "conclusions": "We have presented a method which treats the inference of structural parameters of single stellar populations in a fully rigorous statistical manner, within a Bayesian approach, to construct a merit function for the comparison of an observed CMD in relation to a point in the 4-D parameter space of age, metallicity, distance modulus and reddening.  Implementing a genetic algorithm simulation we have shown through various examples using synthetic CMDs with parameters as appropriate for current observations of Galactic Single Stellar Populations, that the method accurately recovers the input parameters, and yields reliable confidence intervals, including naturally an analysis of the covariances and correlations between the recovered parameters.  By performing extensive Monte Carlo experiments with systematic effects, which surely apply to real cases and which are explicitly excluded from the formalism presented, we have evaluated the consequences  these effects are expected to have on the results of our method. We found that although the error ranges on the recovered parameters typically increase, no systematic offsets appear, at least within the reasonable ranges tested.  This implies that our formalism, which explicitly takes into account the number density of stars along isochrones, can tentatively test stellar evolution models when using clusters with well-measured abundances, distances and reddenings. This, along with the absolute dating of clusters with precise photometry, will be presented in the second paper of this series (Valls--Gabaud \\& Hernandez 2007)."
    },
    "0801/0801.1048_arXiv.txt": {
        "abstract": "Non-thermal components are key ingredients for understanding clusters of galaxies. In the hierarchical model of structure formation, shocks and large-scale turbulence are unavoidable in the cluster formation processes. Understanding the amplification and evolution of the magnetic field in galaxy clusters is necessary for modelling both the heat transport and the dissipative processes in the hot intra-cluster plasma. The acceleration, transport and interactions of non-thermal energetic particles are essential for modelling the observed emissions. Therefore, the inclusion of the non-thermal components will be mandatory for simulating accurately the global dynamical processes in clusters. In this review, we  summarise the results obtained with the simulations of the formation of galaxy clusters which address the issues of shocks, magnetic field, cosmic ray particles and turbulence. ",
        "introduction": "\\label{Introduction}  Within the intra-cluster medium (ICM), there are several processes which involve non-thermal components, such as the magnetic field and cosmic rays (CRs). At first glance, many of these processes can be studied numerically in simplified configurations where one can learn how the individual processes work in detail. For example, there are investigations which consider magneto-hydro-dynamical (MHD) simulations of cloud-wind interactions \\citep[see][ and references therein]{2000ApJ...543..775G} or simulate the rise of relic radio bubbles \\citep[see][ and references herein]{2005ApJ...624..586J,2005MNRAS.357..242R}. Such work usually focuses on the relevance of CRs, turbulence and the local magnetic fields within and around these bubbles. In this review, we will instead concentrate on simulations of non-thermal phenomena, which aim at understanding the relevance of these phenomena for galaxy cluster properties or at unveiling possible origins of the non-thermal radiation. So far, firm evidence for the presence of non-thermal emission (at radio wavelengths) and for the presence of cosmologically relevant extended magnetic fields has been found only in galaxy clusters; in filaments only weak limits have been found so far. For a more detailed discussion see the reviews on radio observations by \\citet{rephaeli2008} - Chapter 5, this volume, and \\citealt{ferrari2008} - Chapter 6, this volume. ",
        "conclusions": "\\label{sec:discussion}  It seems that, in the last years, a consistent picture of magnetic fields in clusters of galaxies has emerged from both numerical work and observations. Simulations of individual processes like shear flows, shock/bubble interactions or turbulence/merging events predict a super-adiabatic amplification of magnetic fields. It is worth mentioning that this common result is obtained by using a variety of different codes, which are based on different numerical schemes. Further support for such super-adiabatic amplifications comes from analytical estimates of the anisotropic collapse which use the Zel'dovich approximation \\citep{1970A&A.....5...84Z}. When this amplification occurs in fully consistent cosmological simulations, various observational aspects are reproduced (e.g. Fig.~\\ref{fig:RMprof}); moreover, the final strength of the magnetic field reaches a level sufficient to link models that predict magnetic field seed at high redshift with the magnetic fields observed in galaxy clusters today. It is important to note that all the simulations show that this effect increases when the resolution is improved, and therefore all the current values have to be taken as lower limits of the possible amplification.  Despite this general agreement, there are significant differences among the predictions on the structure of the magnetic fields coming from different models of magnetic field seeds. In particular, there are several differences between the up-scaled cosmological battery fields \\citep{2001ApJ...562..233M,2004PhRvD..70d3007S} and the magnetic field predicted by the high-resolution simulations of galaxy clusters that use either AMR \\citep{2005ApJ...631L..21B} or SPH \\citep{1999A&A...348..351D,2002A&A...387..383D,2005JCAP...01..009D}. In the SPH case, it is possible to follow the amplification of the field seeds within the turbulent ICM in more detail.  A good visual impression can be obtained by comparing the regions filled with the high magnetic fields shown in Fig.~\\ref{fig:Sigl_I} and \\ref{fig:Brueggen_I}. It is clear that the high magnetic field regions for the battery fields are predicted to be much more extended; this result leads to a flat profile around the forming structure, whereas for the turbulent amplified magnetic fields the clusters show a magnetic field distribution which is much more peaked. Part of this difference originates from the physical model, as the cosmological shocks are much stronger outside the clusters. Somewhat less clear is how big it is the influence of the different numerical resolutions of these simulations to these discrepancies. Usually, the model parameters for such simulations can only be calibrated using the magnetic field strength and structure within the high-density regions of galaxy clusters. Therefore it is crucial to perform detailed comparisons with all the available observations to validate the simulations. Note that extrapolating the predictions of the simulations into lower density regions, where no strong observational constraints exist, will further amplify the differences among the simulations in the predictions of the magnetic field structure.  Moreover, one has to keep in mind that, depending on the ICM resistivity, the magnetic field could suffer a decay, which is so far neglected in all the simulations. Furthermore, a clear deficiency of the current simulations is that they do not include the creation of a magnetic field due to all the feedback processes happening within the LSS (like radio bubbles inflated by AGN, galactic winds, etc.): this might alter the magnetic field prediction if their contribution turns out to be significant. Also all of the simulations so far neglect radiative losses; if included, this would lead to a significant increase of the density in the central region of clusters and thereby to a further magnetic field amplification in these regions.  Finally, there is an increasing number of arguments suggesting that instabilities and turbulence on very small scales can amplify the magnetic fields to the observed $\\mu$G level over a relatively short timescale. However, it is still unclear how such fields, which are tangled on very small scales, can be further processed to be aligned on large scales (up to hundreds of kpc), as observed.  Concerning radio haloes and relics, the picture is only partially consistent (for a more detailed discussion of primary and secondary models see \\citet{2004JKAS...37..493B} and references therein as well as \\citealt{ferrari2008} - Chapter 6, this volume). Radio relics seem to be most likely related to strong shocks produced by major merger events and therefore produced by direct re-acceleration of CRs, the so-called primary models. Although some of the observed features, like morphology, polarisation and position with respect to the cluster centre, can be reasonably well reproduced, there might be still some puzzles to be solved. On the one hand, direct acceleration of CRs in shocks seems to overestimate the abundance, and maybe the luminosity, of radio relics; on the other hand, simulations which illuminate fossil radio plasma can produce reasonable relics only starting from a small range of parameter settings.  A similar situation arises for modelling the central radio emission of galaxy clusters by secondary models. The total luminosity seems to be reproduced using reasonable assumptions and also the observed steep correlation between cluster temperature/mass and radio power seems to be reproduced quite well. However, these models suffer from two drawbacks. The first one is that, in the framework of such models, every massive cluster produces a powerful radio halo, at odds with observations (Fig.~\\ref{fig:pt}). The other problem of secondary models is that the detailed radio properties are not reproduced. In fact, firstly, in most cases, the profile of radio emission is too steep, so that these models can almost never reproduce the size of the observed radio halos. Secondly, the observed spectral steepening \\citep[e.g.][]{2005A&A...440..867G} cannot be reproduced.  One possible way out, at least for the first concern, would be that magnetic fields are extremely dynamical and transient and therefore {\\it light up} the halo during a merger. However, there is also no indication from observations that clusters which show radio emission contain magnetic fields which are more intense than those in clusters without observable extended diffuse radio emission. On the contrary, the cluster A~2142 has a magnetic field strength similar to the Coma cluster, but the upper limit on its radio emission is at least two orders of magnitude below the value expected from the correlation (Fig.~\\ref{fig:pt}). Note that both clusters are merging systems which are characterized by the presence of two central cD galaxies. This indicates that there should be further processes involved or additional conditions required to produce radio emission. It is worth noting that recent models, based on turbulent acceleration, seem to overcome this problem \\citep[see][ and references therein]{2003ApJ...594..732K,2005MNRAS.363.1173B,2005MNRAS.357.1313C}. \\citet{2005MNRAS.357.1313C} predicted that the probability for a galaxy cluster to show a giant radio halo is an increasing function of the cluster mass; these authors also reproduced the observed probability of $\\approx 30$~\\% for a massive galaxy clusters to show extended radio emission.  Being dominated by large-scale bulk motions, induced by the cosmological structure formation process, different simulation methods reach good agreement in predicting that the ratio of the bulk kinetic energy to thermal energy has an upper limit of 15~\\% in galaxy clusters \\citep[see][]{1999ApJ...525..554F}. However, the exact role of turbulence induced by hydrodynamic instabilities within the complex flow patterns of galaxy clusters is still hard to quantify \\citep[see][]{2005MNRAS.364..753D}. It should also be mentioned that the real viscosity of the ICM is hardly known, and can play a noticeable role within the ICM, as demonstrated in simulations which attempt to include physically motivated viscosity into a cosmological context \\citep[see][]{2006MNRAS.371.1025S}.  The dynamical role of CRs in galaxy clusters is not very well understood yet. Simulations using different codes predict quite different relative pressure contained in CRs within galaxy clusters. Moreover, the simulations indicate that the relative importance of CRs in galaxy clusters also strongly depends on other non-thermal processes, like radiative losses and feedback from star formation \\citep[see][]{2007MNRAS.378..385P}.  In general it has to be pointed out that, although the complexity of the physical processes treated by cosmological simulations have improved dramatically in recent years, this research field is still young. Therefore future work is expected to enlighten such complex processes within the ICM and will help to improve our understanding of the non-thermal components in galaxy clusters. Such work will also be needed to interpret the observational information which is overwhelmingly increasing in quality and size and is expected to become available soon with the next generation of instruments, particularly at radio band, like ALMA, LOFAR, SKA, EVLA and others."
    },
    "0801/0801.2367.txt": {
        "abstract": "\\noindent We present an effective field theory formulation for a class of condensed matter systems with crystalline structures for which some of the discrete symmetries of the underlying crystal survive the long distance limit, up to mesoscopic scales, and argue that this class includes interesting materials, such as $Si$-doped $GaAs$. The surviving symmetries determine a limited set of possible effective interactions, that we analyze in detail for the case of  $Si$-doped $GaAs$ materials. These coincide with the ones proposed in the literature to describe the spin relaxation times for the $Si$-doped $Ga As$ materials, obtained here as a consequence of the choice of effective fields and their symmetries. The resulting low-energy effective theory is described in terms of three (six chiral) one-dimensional Luttinger liquid systems and their corresponding intervalley transitions. We also discuss the Mott transition within the context of the effective theory. ",
        "introduction": "Both electronic and spin dynamics are major research areas in Condensed Matter Physics, specially for materials which induce strong correlating among electrons. In such systems, the interplay between several effective interactions and physical degrees of freedom is believed to lead to interesting effects, such as the colossal magnetoresistance and the high $T_C$ superconductivity. The microscopic structure of these materials is complex: generically, these are compounds grew on some substrate with random impurities, yielding a crystalline structure with some amount of disorder. Moreover, the low density of the free carriers renders often their kinetic energy comparable to the Coulomb energy, pushing the system close to the Hubbard-Mott transition ({\\it i.e.}, metal-insulator transition driven by correlations). In this regime, the approximations leading to the usual band theory results become invalid, and it is assumed that the electronic system is not a Fermi liquid.  The Hubbard-Mott transition has been the subject of much attention in Condensed Matter Physics, often studied in reference to the Hubbard Model. Actually, some authors recognize two types of Mott transitions \\cite{Giamarchi-1}: the so-called $U$-Mott transition, in which at a given filling, a variation of the interaction leads to a metal-insulator transition, and the $\\delta$-Mott transition, for which the interaction is constant and the doping variable, leading to a commensurate to a incommensurate transition. The physics of the transition still remains not well-understood in $d=2$ and $d=3$ spatial dimensions. Moreover, most of the studies used mean-field techniques, in particular, the Dynamical Mean Field Theory \\cite{Georges-Rev}. Some more detailed information of the transition, such as its universality class or critical exponents, still remain to be fully established. Recently, it  has been claimed that the Mott transition is of the gas-liquid type, and belongs to the Ising Universality class.\\cite{Georges-Mott-Universality} \\cite{Kotliar-Rozenberg-orderP}. One may argue, however, that at least for some systems of interest, the separation between $U$-Mott and $\\delta$-Mott transitions is somewhat artificial, since, at low enough density values, changing the number of free particles ({\\it i.e.}, doping) changes simultaneously the interaction and the commensurate-incommensurate character of the $\\delta$\\ Mott transition.  On the other hand, the Mott transition has recently attracted interest in a rather different arena, namely, the non-equilibrium  magnetic properties of spin polarized electrons. In a remarkable experiment, Kikkawa and Awschalom\\cite{Kikkawa1} have found that the spin lifetime in Si-doped $GaAs$ semiconductors can be enlarged by up to three orders of magnitude by the introduction of $n$-doping of the order of $n=10^{-3} cm^{-3}$. Among the early mechanisms that have been used to study the spin decay in semiconductors are the ones proposed by Elliot-Yafet (E-Y)\\cite{Elliot} \\cite{Yafet}, D'yakonov-Perel (D-P)\\cite{D-P} , Bir-Aronov-Pikus \\cite{B.A.P}, the relaxation due to nuclear spins, the Dyaloshinskii- Moriya (D-M) interaction\\cite{Dzyalo}\\cite{Moriya} and others. Recently, some calculations have been repeated for the E-Y mechanism, in the low temperature regime \\cite{Tambo}, and for the D-P mechanism \\cite{Flatte} but they disagree with the experimental results by nearly three orders of magnitude at doping near to $n=10^{16}\\ cm^{-1}$. Moreover, it has been claimed that the enhancement of spin lifetime can be understood within the D-M picture \\cite{kav1}. This claim has been criticized by Gorkov, who has pointed out that the enhanced lifetimes take place at doping levels near to the metal-insulator (Mott) transition \\cite{Gorkov}. In fact, while all the preceding mechanisms and calculations are perturbative, the Mott transition is a strongly correlated effect,thus at first glance, the Hiller-London calculation in ref \\cite{kav1} seems misleading as was pointed out in ref \\cite{Gorkov}. At present, however, we are not aware of any calculations of the Mott transition for systems like $GaAs$. In particular, it is unknown how the crystalline symmetries and the band structure could affect the Mott transition in semiconductor systems. The answer to these questions cannot be given by performing band structure calculations, due to the limitations imposed by the strong correlations.  To summarize: on the one hand, there is not as yet a consensus about the nature of the spin relaxation mechanism in $n$-doped semiconductors: several different models are all based in band structure and dilute gas scattering approximations, which may not be valid for a phenomenon involving strong correlations. On the other hand, several characteristics of the Mott transition are still not well understood, and its eventual relevance to the spin relaxation dynamics it is still an open issue . A first step towards the understanding of both problems would be the formulation of a unified framework for describing systems of electrons including both the band structure information and the strong correlations effects. The main purpose of this paper is to propose such a framework, given by an effective field theory that can be applied for crystalline condensed matter systems of the type discussed above. Effective field theories (EFTs) are a general, systematic and powerful framework to describe the low energy (long distance) dynamics of systems with one or more energy scales. These include and unify some earlier approaches, such as the Landau-Ginzburg theories and the Fermi liquid theories, among others\\cite{Polchinski}. They provide a {\\it systematic tool} to analyze the different interactions and their relative importance through the use of the Renormalization Group (RG) ideas and techniques. Thus, this framework is well suited for the crystalline systems in the regime discussed above, where a naive calculation shows an  energy scale of about $E\\sim {\\hbar }/{100ps}\\sim 6.5\\ 10 ^{-3}\\ mev$ at doping levels of $n\\sim 10^{16}\\ cm^{-3}$, while the typical Fermi energy is of the order of $E_F \\sim 2\\ mev$ at this doping, and the gap energy is of the order of some ev, $E_g\\sim 3\\ ev$. As a test bench for the EFT, we discuss the Mott transition for samples of $Si$-doped $GaAs$, and show that it can be used as a useful tool for discussing some properties of the spin relaxation dynamics processes, although we shall not address those processes directly, given that these are non-equilibrium process and therefore, outside the applicability range of any field theory.  This paper is organized as follows: in the section 2 we write down the general EFT for electrons in crystalline systems. In section 3 we consider the inclusion of low disorder in the EFTs using standard techniques. Section 4 is devoted to the formulation and discussion of the EFT for the case of $Si$-doped $GaAs$ materials, including some perturbative Renormalization Group calculations. In Section 5, we discuss the strong coupling regime of the EFT and its relation to the Mott transition, the scenario for the appearance of a $d$-wave density order parameter and their relevance to the systems considered here. Finally, we give some Conclusions.   %-2------------------------------------ ",
        "conclusions": "In this article we have presented EFTs for electrons in crystalline systems and have studied the effect of quenched impurities (low disorder effects) in them. Based on experimental evidence and consistency and plausibility arguments, we have considered the situation in which impurities form arrays in mesoscopic regions for which some or all of the discrete symmetries of the microscopic lattice are preserved in the long distance limit. We have also discussed EFTs for generic systems with substitutional impurities, although we did not get into details as to how to deal with a general case. We believe that the EFT point of view is a good complement to the one based on microscopic models for studying strongly correlated electron systems in crystals, providing support to them and bringing in a useful bridge between strongly correlated problems and perturbative band theory calculations. We have obtained the EFT for $Si$-doped $GaAs$ systems, obtaining from the symmetries of the effective fields all possible single particle terms in the effective Hamiltonian, which reproduce the known terms of the pseudo-potential calculations of band structure approaches and the so called Dyakonov-Perel Hamiltonian for the spin-orbit coupling. Remarkably, the resulting low-energy effective theory is described in terms of three (six chiral) one-dimensional Luttinger liquid systems and their corresponding intervalley transitions.  \\clearpage   \\vspace {2 cm}    % \\def\\NP{{\\it Nucl. Phys.\\ }}  \\def\\PRL{{\\it Phys. Rev. Lett.\\ }}  \\def\\PL{{\\it Phys. Lett.\\ }}  \\def\\PR{{\\it Phys. Rev.  \\ }}  \\def\\PRB{{\\it Phys. Rev. B  \\ }}  \\def\\CMP{{\\it Comm. Math. Phys.\\ }}  \\def\\IJMP{{\\it Int. J. Mod. Phys.\\ }}  \\def\\MPL{{\\it Mod. Phys. Lett.\\ }}  \\def\\RMP{{\\it Rev. Mod. Phys.\\ }}  \\def\\AP{{\\it Ann. Phys.\\ }} ."
    },
    "0801/0801.0752_arXiv.txt": {
        "abstract": "Chemically peculiar stars define a class of stars that show unusual elemental abundances due to stellar photospheric effects and not due to natal variations. In this paper, we compare the elemental abundance patterns of the ultra metal-poor stars with metallicities [Fe/H] $\\sim -5 $ to those of a subclass of chemically peculiar stars.   These include post-AGB stars, RV Tauri variable stars, and the Lambda Bootis stars, which range in mass, age, binarity, and evolutionary status, yet can have iron abundance determinations as low as [Fe/H] $\\sim -5$. These chemical peculiarities are interpreted as due to the separation of gas and dust beyond the stellar surface, followed by the accretion of dust depleted-gas. Contrary to this, the elemental abundances in the ultra metal-poor stars are thought to represent yields of the most metal-poor supernova and, therefore, observationally constrain the earliest stages of chemical evolution in the Universe. Detailed chemical abundances are now available for HE1327-2326 and HE0107-5240, the two extreme ultra metal-poor stars in our Galaxy, and for HE0557-4840, another ultra metal-poor star found by the Hamberg/ESO survey. There are interesting similarities in their abundance ratios to those of the chemically peculiar stars, e.g., the abundance of the elements in their photospheres are related to the condensation temperature of that element. If HE1327-2326 and HE0107-5240 are ultra metal-poor due to the preferential removal of metals by dust grain formation or dilution through the accretion of metal-poor interstellar gas, then their CNO abundances suggest true metallicities of [$X$/H] $\\sim -2$ rather than their present metallicities of [Fe/H] $\\leq -5$, and, thus their status as truly ultra metal-poor stars is called into question. The initial abundance for HE0557-4840 would be [$X$/H] $\\ge -4$. It is important to establish the nature of these stars since they are used as tests of the early chemical evolution of the Galaxy, yet if they are chemically peculiar, then those tests should be focused on stars in the metal-poor tail of the Galactic halo distribution.  Many, but not all, chemically peculiar stars show a mid-infrared excess from the circumstellar dust.   We examine the JHK fluxes for the ultra metal-poor stars but find no excesses. A more important test of the stars' status as chemically peculiar is provided by the elemental abundances of sulphur and/or zinc. These two elements have low condensation temperatures and do not form dust grains easily; furthermore, the chemically peculiar stars universally show sulphur and zinc to be undepleted or nearly so. We show that near-infrared lines of S\\,{\\sc i} offer a promising test of the possibility that HE1327-2326 may be chemically peculiar. Although there are some parallels between the compositions of the ultra metal-poor stars and chemically peculiar stars, a definitive ruling on whether the former are chemically peculiar requires additional information. ",
        "introduction": "A fossil record of the earliest episodes of stellar nucleosynthesis in the Universe and Galaxy should be revealed by the compositions of the most metal-poor Galactic stars (e.g., Tumlinson 2007a, 2007b; Tominaga \\etal 2007; Umeda \\& Nomoto 2003). The lure of this revelation has driven the search to find and analyse such Rosetta stones. A great leap forward  was achieved recently by the discovery of two stars with  iron abundances [Fe/H] $< -5.3$ (Christlieb \\etal 2002; Frebel \\etal 2005), a limit about 1.5 dex below the abundance of the previously known most metal-poor star. A third star with [Fe/H] $\\simeq -4.8$ also beats the previous lower bound (Norris \\etal 2007). Prior to these remarkable discoveries, the most Fe-poor stars known were HR~4049 and HD~52961 with [Fe/H] $\\simeq -4.8$ (e.g., Waelkens \\etal 1991).   However, these and slightly more iron-rich examples were dismissed - correctly - as irrelevant to the issue of early stellar nucleosynthesis because they are `chemically peculiar', i.e., their present surface compositions are far removed from their initial compositions. In particular, their compositions reflect that of gas from which refractory elements have been removed to varying degrees by a process dubbed `dust-gas separation'. The existence of HR~4049 and HD~52961 has led us to reexamine the question of whether the recently discovered ultra metal-poor stars may themselves be chemically peculiar.  Two of the three ultra metal-poor stars in question are HE1327-2326 and HE0107-5240. HE1327-2326 was discovered by Frebel \\etal (2005) and abundance analyses have been described by Aoki \\etal (2006), Frebel \\etal (2006), Collet \\etal (2006), and most recently by Frebel \\& Christlieb (2007). The latter analysis based on the highest SNR spectra yields abundances for 11 elements and upper limits for an additional nine elements: the new iron abundance at [Fe/H] $=-5.9$ is even lower than the previous determination. HE0107-5240 was discovered by Christlieb \\etal (2002) and abundance analyses have been reported by Christlieb \\etal (2004), Bessell \\etal (2004), and Collet \\etal (2006) with the latter suggesting [Fe/H]$\\simeq -5.6$. The third star HE0557-4840 was discovered and analysed by Norris \\etal (2007): [Fe/H] $= -4.8$ places it between the two ultra metal-poor stars (HE1327-2326 and HE0107-5240) and the lower boundary of the metal-poor tail of Galactic stars.  A marked characteristic of these three discoveries is that some abundance ratios are uncharacteristic  of metal-poor stars of higher Fe abundance. Notably, the stars are C-rich for their Fe abundance, i.e., [C/Fe] $= 3.7$ for HE1327-2326, and also [N/Fe] $= 4.1$ and [O/Fe] $= 3.4$ (Frebel \\& Christlieb 2007). This property of unusual abundance ratios is shared in a qualitative sense with HR~4049 and HD~52961.  In the following sections, we briefly review the classes of known chemically peculiar stars affected by dust-gas separation. Then, we discuss if the three stars HE1327-2326,  HE0107-5240, and HE0557-4840 are chemically peculiar rather than true ultra metal-poor stars. In the final sections, we discuss possible tests of the hypothesis that ultra metal-poor stars may be chemically peculiar.     \\section {Separation of gas and dust and chemically peculiar stars}  The chemically peculiar stars in question are those whose atmospheres betray the operation of the dust-gas separation process. In gas of a sufficiently low temperature, dust condenses out and the gas is depleted in those elements that form the dust. The local interstellar gas, for example, displays such depletion (Savage \\& Sembach 1996). The composition of the gas in a dust-gas mixture depends on several primary factors including the initial composition of the gas, the total pressure, and the  history of the gas-dust mixture. If a star were to then accrete gas, preferentially over dust, and the accreted gas were to comprise a major fraction of the stellar photosphere,  a star with striking abundance anomalies results.   This scenario is one version of how a star could develop chemical peculiarities from dust-gas separation (or ``dust-gas winnowing''). Anomalies plausibly attributable to dust-gas separation have now been reported for three kinds of stars: Lambda Bootis stars, post-AGB A-type spectroscopic binaries, and RV Tauri variables.  The Lambda Bootis stars are main sequence stars, ranging from A- to mid F-types, and found at all evolutionary phases from very young (e.g., they are found in young open clusters, Gray \\& Corbally 1998), to the end of the main-sequence (Kamp \\& Paunzen 2002, Andrievsky \\etal 2002). They are expected to have a solar-like composition, but have long been known for their significant metal-deficiencies (Morgan, Keenan, \\& Kellman 1943; Burbidge \\& Burbidge 1956; Baschek \\& Searle 1969).  These stars have effective temperatures of from about 6500 to 9500~K and surface gravities \\logg $\\simeq 4$. Our (Venn \\& Lambert 1990) abundance analysis of three stars, including the eponym, showed that the compositions of the stars could be attributed to accretion of circumstellar gas, without dust. The iron deficiency in the case of $\\lambda$~Boo was [Fe/H] $\\simeq -2$, and slightly less severe for 29~Cyg, with normal abundances of C, N, O and S for both stars, in agreement with the expectations for an atmosphere contaminated with  dust-free gas (see Figure~1). Even Vega, a `standard' A0 star, was shown to be a mild Lambda Bootis star (Venn \\& Lambert 1990, Lemke \\& Venn 1996).   Vega also shows an infrared excess due to a dusty circumstellar disk (Su \\etal 2005; Aumann \\etal 1984).  The shallow convective envelope of early A-type stars is a key factor in the creation and maintenance of the abundance anomalies in the {\\it diffusion/accretion model} (Turcotte \\& Charbonneau 1993). Abundance anomalies can persist ($\\sim$10$^6$ yr), even after dispersal of the circumstellar dust and gas and, hence, removal of the infrared excess contributed by the dust.   Thus, not all Lambda Bootis stars show an infrared excess; Paunzen \\etal (2003) estimate that 23\\% of {\\it bona fide} Lambda Bootis stars show evidence for circumstellar matter. On the other hand, the circumstellar matter may be of interstellar origin (Kamp \\& Paunzen 2002, G\\'{a}sp\\'{a}r \\etal 2007). Movement of a star through the denser parts of the interstellar medium can create a bow shock which heats the interstellar dust causing an infrared excess.  Meanwhile, radiation pressure from the star repels the grains, while gas is accreted onto the stellar surface. With return of the star to passage through low density gas, the infrared excess dissipates and accretion ceases. This alternate theory for the origin of the Lambda Bootis stars implies the chemical anomalies are transient, but could help to explain why the phenomenon is seen in such a wide range of main-sequence stars.  Abundance anomalies should not survive the transition from the main sequence to the giant branch though.  The deep convective envelope of giant stars will surely dilute abundance anomalies beyond recognition. Yet, a metal deficiency of even greater severity can be found among post-AGB stars in spectroscopic binaries (Waelkens \\etal 1991; Van Winckel 2003). These stars, like HR~4049 and HD~52961 discussed above, are supergiants with \\teff in the range of 6000 to 7600~K and surface gravities \\logg $\\simeq 1$. One must suppose that a new episode of dust-gas separation led to these abundance anomalies.   The original quartet of post-AGB binaries (Van Winckel \\etal 1995) comprised HR~4049, HD~44179, HD~52961, and BD~+39~4926.  Their [Fe/H] values range from $-3.0$ to $-4.8$, but they show an abundance pattern reminiscent of interstellar gas, i.e., quasi-solar abundances of C, N, O, S, and Zn, but severe underabundances of, for example, Si, Ca, and Fe (see Figure~1). The pattern shows that it is the former set of elements that define the initial composition of these stars and not the latter set. Gas is thought to be accreted onto the star from a circumbinary disk, while radiation pressure exerted by the star on the dust grains inhibits the accretion of dust by the star and may also promote a separation of dust and gas in the disk.   The dusty disk is betrayed by an infrared excess: HD~44179, also known as the Red Rectangle, is a rather special proto-planetary nebula with a striking infrared excess.  On the other hand, BD~+39~4926 lacks an infrared excess.   Shallow convective envelopes in these extended stars are presumably a key factor in the appearance of their huge abundance anomalies.  Chemical peculiarities of the post-AGB stars are presumably developed from less extreme peculiarities seen in their immediate progenitors, the RV Tauri variables.  Discovery of this third category of star displaying the marks of dust-gas winnowing began with the analysis of the RV Tauri variable IW~Car (Giridhar \\etal 1994). A RV Tauri star is a post-AGB star with an infrared excess (first noted by Gehrz 1972). Subsequent analyses (Giridhar \\etal 2005) showed that the effects of dust-gas separation among RV Tauri stars appear limited to the warmer stars (\\teff $> 4500$ K) and to the stars with intrinsic metallicities [Fe/H] $\\geq -1$.   Affected stars have \\teff $\\simeq 4500 - 6500$~K (hotter stars fall outside the instability strip and appear as non-variable post-AGB stars) and \\logg $\\simeq 0$ to 1. Gonzalez \\& Lambert (1997) also discussed the potential importance of the composition of the photosphere (e.g., C/O ratio), and environment (e.g., field vs globular cluster stars).  The reasons for these effective temperature and metallicity boundaries are not entirely clear, but the cooler stars possess more extensive convective envelopes that dilute the accreted gas. Dust-gas winnowing may be impaired in low metallicity gas where the dust to gas ratio is necessarily lower. The winnowing site is again presumed to be a circumbinary disk; there is increasing evidence that the affected stars are spectroscopic binaries (Van Winckel 2007).  How the gas is captured onto the star from the circumbinary disk in the presence of a stellar wind remains an unsolved problem, as it does for the post-AGB stars in binaries.  A signature of a star that is affected seriously by dust-gas winnowing is a correlation between an element's abundance and the predicted condensation temperature $T_{cond}$ (or the abundance in interstellar gas). Estimates of $T_{cond}$ depend on the  initial composition and pressure in a gas and the assumption that cooling of the gas and condensation of grains occurs under equilibrium conditions. Lodders (2003) provides a comprehensive discussion of $T_{cond}$ estimates for all elements for the solar (O-rich) composition. In the post-AGB stars, and in some Lambda Bootis and RV Tauri stars, the abundance [$X$/H] is smoothly correlated with $T_{cond}$ (see Figure~1).     In some Lambda Bootis and RV Tauri stars, [$X$/H]  shows no obvious trend with the  $T_{cond}$ (e.g., EQ~Cas, Giridhar \\etal 2005). Rao \\& Reddy (2005) brought order to chaos by noting instead that the [$X$/H] for EQ~Cas were correlated with the ionization potential of the neutral atom (`the first ionization potential' or FIP). This raises the intriguing possibility that an alternative or additional process is operating in this one star. The FIP effect is a well known  phenomena in the solar corona thought to reflect the greater ease with which ions (low FIP elements) rather than neutral atoms (high FIP elements) are fed from the cool chromosphere into the corona. Perhaps, EQ~Cas is a star where there is a selective feeding of the stellar wind. Thus, in single stars, the {\\sl stellar wind} may control the abundance anomalies but where the variable is in a  binary the {\\sl circumbinary disk} may exert control.  Another possibility for affecting the [$X$/H] vs $T_{cond}$ correlation (at high T$_{cond}$) is a competition between accretion of metal free gas and chemical separation due to, e.g., diffusion, gravitational settling, radiative acceleration, or rotational mixing above the convective zone.   The diffusion/accretion model for Lambda Bootis stars by Turcotte \\& Charbonneau (1993) suggests timescales for chemical separation can vary by element, thus complicating the abundance pattern established earlier by accretion. The conclusion must be that the dust-gas separation mechanism does not have the same effects on the elemental abundances in the various stars where it appears to operate. Either the mechanism itself, and/or the re-accretion phase of the dust-depleted gas, and/or additional processes that affect chemical separation impact the observed chemical abundance pattern.  In summary, stars with abundance anomalies attributable to dust-gas separation are seen in several parts of the HR-diagram. Associated properties of affected stars, such as age, mass, binarity, or evolutionary status, do not appear to be uniform across the cases. Details in the observed abundance patterns seem to vary between the cases, particularly for elements with high T$_{cond}$, possibly due to the mechanism itself or additional processes. An accompanying infrared excess may be a warning bell, but it is not a required signature of dust-gas separation. However, to date,  the inferred intrinsic metallicities of affected stars are one-tenth solar or greater. ",
        "conclusions": "In this paper, we have outlined the existing evicence for and against the three ultra metal-poor stars being affected by dust-gas separation. The evidence for dust-gas separation having occured at the stellar photosphere is primarily the abundance pattern, which resembles that of known chemically peculiar stars globally where dust-gas separation has been supported by other observations.   The other observations usually include one or more of the following for at least one object within a sample; infrared excesses (usually beyond the K band at 2.2 microns) or other evidence of circumstellar material, evidence of binarity, or the observation that sulphur or zinc do not have the same depletions as the rest of the metals.   The latter is a critical test because non-depleted abundances of these two elements is more consistent with dust formation, otherwise requiring random and puzzling variations in nucleosynthesic yields.    Most of these observational tests do not yet exist for the ultra metal-poor stars, or the results are currently inconclusive regarding dust-gas separation.  Searchs for radial velocity variations are ongoing in the ultra metal-poor stars to determine if they are in binary systems.  Currently, no variations are found, which suggests these stars are not binaries, but this evidence cannot be conclusive because of orbit inclinations and/or potentially small amplitude variations.     Sulphur and zinc upper-limits exist for all three ultra metal-poor stars, however these upper-limits are presently too high to distinguish between nucleosynthetic origins or dust-gas separation. K band photometry exists for all three ultra metal-poor stars, and none of them show an infrared excess, however many RV Tau, post-AGB, and Lambda Boo chemically peculiar stars only show infrared excesses beyond the K band, if at all (many have no detectable infrared excesses). In addition, known chemically peculiar stars suggest that neither dust not binarity are necessary conditions; the Lambda Bootis stars  exist as single stars, often without dust; the peculiar post-AGB are binaries but one lacks dust; the affected RV Tauri stars have an infrared excess and, although binarity may be a necessary condition, it is as yet unknown observationally that all are binaries. Nevertheless, these observational are possible for the ultra metal-poor stars, e.g.  Spitzer observations at mid-IR wavelengths, and examination of the S\\,{\\sc I} lines in near-IR spectroscopy, but have yet to be performed.  One remaining concern is that two of the ultra metal-poor stars do not occupy similar atmospheric parameter ranges as any of the known chemically peculiar stars. HE0107-5240 and HE0557-4840 have temperatures like the RV~Tauri variables, yet higher gravities and lower intrinsic metallicities (if CNO are the appropriate proxies). Physically, these stars are expected to have atmospheres  with deep convective envelopes, and thus, accreted gas from a stellar wind, a circumstellar shell, or a circumbinary disk is should be diluted beyond detection. Existence of these envelopes makes it difficult to understand how dust-gas separation effects can be created and sustained in these giants. Indeed, the stars must have received several tenths of a solar mass of separated gas in order that the anomalies be detectable in these giants. Of course, this is a concern for the RV Tauri and post-AGB stars as well. These difficulties are  less severe for HE1327-2326, which is significantly hotter and is expected to have a more shallow convective envelope.   HE1327-2326 occupies (nearly) the same stellar parameter range as the Lambda Bootis stars.  Fortunately, detection and analysis of the S\\,{\\sc i} lines in this star offers a critical test for any chemical peculiarities due to dust-gas separation.   Stellar astronomers are conditioned to think about stars with very peculiar abundance anomalies - real or imagined. Commonsense is often a reliable guide to the boundary between real and imagined. In this regard, if the three stars under consideration here are truly stars of a higher intrinsic metallicity, one might ask why examples have not been seen in well studied samples of metal-poor stars. The globular clusters spring to mind. Str\\\"{o}mgren photometry should be able to pick out those few stars with a reduced metallicity from the rest of the stars showing the monometallicity that is a mark of a globular cluster. Of course, if these stars are due to the random effects of passing through dense interstellar clouds, and the effects on the abundances are short lived, then one does not expect to find similar stars in globular clusters, and from objective prism surveys such stars would not stand out as peculiar if the resultant metallicity is greater than $\\sim -3$.  It is only because ultra metal-poor stars are so important as tests of early chemical evolution in the Galaxy that these stars were picked out and studied in detail from high resolution spectroscopy.   It is important that the suspicion be laid to rest that they are not truly ultra metal-poor but chemically peculiar stars  of a more common, if low,  metallicity. Then, the focus may be placed exclusively on finding an explanation in terms of stellar nucleosynthesis and the chemical evolution of the young Galaxy (e.g., Iwamoto \\etal 2005)."
    },
    "0801/0801.2656_arXiv.txt": {
        "abstract": "Two low-dimensional magnetohydrodynamic models containing three velocity and three magnetic modes are described. One of them (nonhelical model) has zero kinetic and current helicity, while the other model (helical) has nonzero kinetic and current helicity. The velocity modes are forced in both these models. These low-dimensional models exhibit a dynamo transition at a critical forcing amplitude that depends on the Prandtl number.   In the nonhelical model, dynamo exists only for magnetic Prandtl number beyond 1, while the helical model exhibits dynamo for all magnetic Prandtl number.   Although the model is far from reproducing all the possible features of dynamo mechanisms, its simplicity allows a very detailed study and the observed dynamo transition is shown to bear similarities with recent numerical and experimental results. ",
        "introduction": "The understanding of magnetic field generation, usually referred to as the dynamo effect, in planets, stars, galaxies, and other astrophysical objects remains one of the major challenges in turbulence research. There are many observational results from the studies of the Sun, the Earth, and the galaxies~\\cite{Moff:book,Krau:book,Bran:PR}. Dynamo has also been observed recently in laboratory experiments~\\cite{Fauv:VKS_PRL06,Fauv:dynamo_reverse_EPL06} that have made the whole field very exciting. Numerical simulations~\\cite{Sche:dynamo_critPm_PRL04,Sche:dynamo_lowPm_PRL07,Pont:dynamo_lowPr_PRL05,Mini:DynamoPoP06,Mini:LowPr} also give access to many useful insights into the physics of dynamo. However, the complete understanding of the dynamo mechanisms has not yet emerged.  The two most important nondimensional parameters for the dynamo studies are the Reynolds number $Re=UL/\\nu$ and the magnetic Prandtl number $P_m=\\nu/\\eta$, where $U$ and $L$ are the large-scale velocity and the large length-scale of the system respectively, and $\\nu$ and $\\eta$ are the kinematic viscosity and the magnetic diffusivity of the fluid. Another nondimensional parameter used in this field is the magnetic Reynolds number $Re_m$, defined as $UL/\\eta$. Clearly $Re_m=Re\\ P_m$, hence only two among the above three parameters are independent. Note that galaxies, clusters, and the interstellar medium have large $P_m$, while stars, planets, and liquid sodium and mercury (fluids used in laboratory experiments) have small $P_m$ \\cite{Moff:book,Krau:book}.  In a typical simulation, the magnetofluid is forced and the dynamo transition is considered to be observed when a nonzero magnetic field is sustained in the steady-state laminar solution or in the statistically stationary turbulent solution, depending on the regime. Typically, dynamos occur for forcing amplitudes beyond a critical value which also defines the critical Reynolds number $Re^{c}$ and the critical magnetic Reynolds number $Re_m^{c}$. One of the objectives of both the numerical simulations and the experiments~\\cite{Fauv:VKS_PRL06,Fauv:dynamo_reverse_EPL06} is the determination of this critical magnetic Reynolds number $Re_m^{c}$. It has been found that $Re_m^{c}$ depends on both the type of forcing and the Prandtl number (or Reynolds number), yet the range of $Re_m^{c}$ observed in the numerical simulations is from 10 to 500 for a wide range of $P_m$ (from $5\\times10^{-3}$ to 2500).  In recent magnetohydrodynamics (MHD) simulations, Schekochihin et al.~\\cite{Sche:dynamo_critPm_PRL04,Sche:DynamoApJ04} applied nonhelical forcings and observed that the dynamo is active for a magnetic Prandtl number larger than a critical Prandtl number $P_m^{c}$ that is around 1. For fluids with small Prandtl number $P_m$ ($P_m<1)$, numerical simulations~\\cite{Sche:dynamo_lowPm_PRL07,Pont:dynamo_lowPr_PRL05,Mini:DynamoPoP06} indicate that the dynamo can be produced using forcings having local helicity (the net helicity of the force could still be zero). The range of the critical magnetic Reynolds number in most of the simulations~\\cite{Sche:dynamo_lowPm_PRL07,Pont:dynamo_lowPr_PRL05,Mini:DynamoPoP06} is 10 to 500. Note that in the Von-Karman-Sodium (VKS) experiment, the critical magnetic Reynolds number is around 30.  There are many attempts to understand the above observations. For large Prandtl number, the resistive length scale is smaller than the viscous scale. For this regime, Schekochihin et al.~\\cite{Sche:dynamo_critPm_PRL04,Sche:DynamoApJ04} suggested that the growth rate of the magnetic field is higher in the small scales because stretching is faster at these scales. This kind of magnetic field excitation is referred to as \\emph{small-scale turbulent dynamo}. For low $P_m$ Stepanov and Plunian \\cite{Step:DynamoGrowth} argue for similar growth mechanism.  Their results are based on shell model calculations.  The numerical results of Iskakov et al.~\\cite{Sche:dynamo_lowPm_PRL07} however are not conclusive in this regard.  The main arguments supporting these explanations are  based on the inertial range (small-scales) properties of turbulence~\\cite{Sche:dynamo_critPm_PRL04,Sche:DynamoApJ04}. In this paper, we present an alternate viewpoint. We show that a low-dimensional dynamical system containing only large scales properties of the fields, three velocity and three magnetic Fourier modes, is able to reproduce some of the above numerical results. For  certain types of  forcing of velocity, a dynamo transition is observed for $Re_m>Re_m^{c}$. These observations indicate that the large-scale eddies may also be responsible for the dynamo excitation, and several important properties can be derived from the dynamics of large-scale modes. These observations are consistent with earlier results of MHD turbulence indicating that the large-scale velocity field provides a significant fraction of the energy (around 40\\%) contained in the large-scale magnetic field \\cite{Dar:flux,Oliv:Simulation,MKV:PR,Mini:I,Mini:II,Olvi:JOT}. ",
        "conclusions": "In this paper, a class of low-dimensional models that exhibit dynamo transition is derived. These models depend on several parameters such as the Prandtl number, the forcing amplitude and a parameter, $h$, that characterizes the ability of the velocity modes to carry kinetic helicity. The fixed points of two simple models corresponding to $h=0$ and $h=1$ are studied in details. The first model is nonhelical (zero kinetic and current helicity) and is compatible with a stationary nonzero magnetic solution only for $P_m>1$. The second model has nonzero kinetic and current helicities, and it has nonzero stationary magnetic solution for all $P_m$. These findings confirm the idea that both kinetic and current helicities may play an important role in the dynamo transition, especially in helical model for $P_m<1$. These two values of $h$ correspond to the only values for which one of the nonlinear coupling in the low dimensional model vanishes. With the specific choice of the forcing proposed in the previous section, these models have then exact and tractable solutions. Arbitrary values of $h$ would lead to much more complex systems.  Obviously, the models presented here are only a small subset of many possible MHD models that could exhibit dynamo. For instance, Rikitake~\\cite{Rikitake:PhilSoc} constructed a dynamo model consisting of two-coupled disks that shows field reversal. Nozi\\`eres \\cite{Nozieres} proposed a model involving one velocity and two magnetic field variables. These two models are phenomenological. Recently,  Donner et al.~\\cite{Donner:Lowdim_physd_2006} proposed a truncation of the MHD equations along the same lines as our model, but derived a much more complex system of equations for 152 modes.  Donner et al.~\\cite{Donner:Lowdim_physd_2006} analyzed the dynamical evolution of these modes for $P_m=1$ and observed steady-state and chaos in their system. The main advantage of our model is that it allows a complete analytical treatment.  The six variables of our model are only representative of the large-scale modes of the system, while a realistic description of a turbulent system exhibiting dynamo should have a large number of modes. Although the small scale variables are often quite important in turbulent flow, in some cases the large scale variables may determine some of the flow properties. The surprising success of the very simple models proposed here in reproducing several feature of the dynamo transition could suggest that we are in such a situation. Schekochihin et al.~\\cite{Sche:dynamo_critPm_PRL04,Sche:DynamoApJ04}, Stepanov and Plunian \\cite{Step:DynamoGrowth}, and Iskakov et al.~\\cite{Sche:dynamo_lowPm_PRL07} have highlighted the role played by inertial range eddies that are absent in our low-dimensional models.  Yet in the absence of a definite theory of dynamo, it is interesting to show that low-dimensional models that focus on the dynamics of the large scale flows may also be successful."
    },
    "0801/0801.1463_arXiv.txt": {
        "abstract": "We predict how protoplanetary disks around low-mass young stars would appear in molecular lines observed with the ALMA interferometer. Our goal is to identify those molecules and transitions that can be used to probe and distinguish between chemical and physical disk structure and to define necessary requirements for ALMA observations. Disk models with and without vertical temperature gradient as well as with uniform abundances and those from a chemical network are considered. As an example, we show the channel maps of HCO$^+$(4-3) synthesized with a non-LTE line radiative transfer code and used as an input to the GILDAS ALMA simulator to produce noise-added realistic images. The channel maps reveal complex asymmetric patterns even for the model with uniform abundances and no vertical thermal gradient. We find that a spatial resolution of $0.2-0.5\\arcsec$ and 0.5--10~hours of integration time will be needed to disentangle large-scale temperature gradients and the chemical stratification in disks in lines of abundant molecules. ",
        "introduction": "The cradles of planet formation -- protoplanetary disks -- are ubiquitous around young stars, at least of low and intermediate masses ($\\la 8\\,M_\\sun$). It is imperative to study disk structure and evolution in detail if one aims to understand initial conditions for planet formation. These objects are best studied at (sub-)millimeter wavelengths in thermal dust continuum and various molecular lines tracing their physical and chemical structure. However, being small ($\\sim 200-1\\,000$~AU) and of relatively low mass ($\\sim 0.01{\\rm M}_\\sun$), such disks are hard to study observationally as high spatial resolution and sensitivity are required. To date only several nearby disks around T Tauri and Herbig Ae/Be stars have been spatially resolved in a handful of molecular lines \\citep{Bea07}. These observational data coupled to advanced modeling allowed us to constrain basic disk parameters such as size, mass, temperature, density distribution, kinematics, and ionization structure \\citep[e.g.,][]{Sea05,Qi_ea06,Pietu_etal2007,Dutrey_ea07}. An interesting result concerns the vertical temperature gradient that seems to be present in some disks (e.g., \\object{DM Tau}, \\object{AB Aur}) while absent in others (e.g., \\object{LkCa 15}), see \\citet{Pietu_etal2005,Pietu_etal2007}.  The situation will change dramatically when the Atacama Large Millimeter Array (ALMA), equipped with 50 12-m antennas, will come into operation in 2012. ALMA will be capable of imaging protoplanetary disks at spatial resolutions up to $\\sim 0.04-0.005\\arcsec$ in a frequency range between 100 and 950~GHz. This will allow direct detection and characterization of disk instabilities, resulting in clumpy structures (vorticities, ``spiral arms'', etc.) as well as inner gaps and holes induced by forming giant planets \\citep{Wolf_ea02a,Wolf_ea02b,WDA_05,Nea06}. The large-scale chemical structure of protoplanetary disks will become accessible in rotational lines of many abundant molecules \\citep{ARea07,Kea07}. As demonstrated by \\citet{Pea07}, excitation of molecular lines in disks with strong gradients of physical conditions and chemical structure can be a complicated process and may be hard to interpret.  The main aim of this work is to unravel the potential of ALMA to distinguish between various temperature and chemical effects in protoplanetary disks, with an emphasis on discerning spatial resolution and integration time needed for this goal. Such observations will allow us to fully explore the validity of state-of-the-art disk chemical models on a global scale \\emph{for the first time}, and advance our understanding of the disk physics enormously. In contrast to previous studies of \\citet{Nea06} and \\citet{Kea07}, we take into account both the non-LTE excitation and realistic chemical structure. ",
        "conclusions": "We showed that the ALMA interferometer will allow us to distinguish the effects of temperature gradients and chemical stratification in disks through molecular line observations, in particular for highly-inclined objects. We found that moderately extended array configurations (with baselines of $\\la 1$~km) and 0.5--10 hours of integration time will be necessary to pursue such observations."
    },
    "0801/0801.1180_arXiv.txt": {
        "abstract": "%  $\\Lambda$CDM numerical simulations predict that the ``missing baryons\" reside in a Warm-Hot gas phase in the over-dense cosmic filaments. However, there are now several theoretical and observational arguments that support the fact that galactic disks may be more massive than usually thought, containing a substantial fraction of the ``missing baryons\". Hereafter, we present new N-body simulations of galactic disks, where the gas content has been multiplied by a factor 5. The stability of the disk is ensured by assuming that the ISM is composed out of two partially coupled phases, a warm phase, corresponding the observed CO and HI gas and a cold collisionless phase corresponding to the unseen baryons. ",
        "introduction": "While the $\\Lambda$CDM scenario predicts that the dark matter in galactic disks is distributed in a dynamically hot heavy halo, there are now several observational facts showing that a fraction of this matter may be baryonic, lying in the outer disk of galaxies. Long ago, \\citet{bosma78} \\citep[see also][]{hoekstra01} observed a correlation between the HI and dark matter in spiral galaxies, indicating a physical link between these two components. This HI-DM correlation is confirmed by the baryonic Tully-Fischer relation \\citep{pfenniger05}. Dark matter in the disk of galaxies is supported by dynamical arguments based on the asymmetries of galaxies, either in the plane of the disks or transversal to it. Large spiral structure present in the very extended HI disk of NGC\\,2915 is naturally explained by a quasi self-gravitating disk \\citep{bureau99,masset03}. Self-gravitating disks are also vertically unstable and give a natural explanation to the old problematic of warps \\citep{revaz04}. \\citet{thilker05} observed UV-bright sources in the extreme outer disk of M83, revealing unexpected star formation, indicating that molecular gas, must be present in abundance in the outer disk of galaxies. Other arguments supporting the idea that dark matter could reside in the galactic disk in form of cold molecular gas are discussed in \\citet{pfenniger94a}. ",
        "conclusions": "New numerical simulations show that the hypothesis where galaxies are assumed to have an additional dark baryonic component is not in contradiction with galaxy properties. Multiplying the HI disk by a factor of 5 let the disk stable if the ISM is composed out of two dynamically decoupled phases. Such galaxy gives a natural solution to different puzzling problems like the the presence of spiral structures in the outer HI disks and the missing mass problem in collisional debris of galaxies.             \\vspace{-0.25cm}"
    },
    "0801/0801.1525_arXiv.txt": {
        "abstract": "We report a dynamical measurement of the mass of the brown dwarf GJ~802B using aperture-masking interferometry and astrometry. In addition, we report the discovery that GJ~802A is itself a close spectroscopic non-eclipsing binary with a 19 hour period. We find the mass of GJ~802~B to be $0.063\\pm0.005$M$_\\sun$. GJ~802 has kinematics inconsistent with a young star and more consistent with the thick disk population, implying a system age of $\\sim$10\\,GYr. However, model evolutionary tracks for GJ~802B predict system ages of $\\sim$2\\,GYr, suggesting that brown dwarf evolutionary models may be underestimating luminosity for old brown dwarfs. ",
        "introduction": "The boundary between stars and brown dwarfs is defined as the mass at which the luminosity of old objects is just dominated by hydrogen burning. This boundary is predicted to occur at a mass between 0.07-0.072\\,$M_\\sun$ at [Fe/H]=0 and $\\sim$0.092\\,$M_\\sun$ at [Fe/H]=-3 \\citep{Chabrier00,Burrows01}.  However, this boundary, and the theoretical relationships that predict effective temperature and luminosity as a function of mass and age, are largely untested by observations. The observations that are required to test these models are dynamical mass measurements of binary and multiple-star systems, combined with accurate photometry and distance determinations, preferably at known age. Many stars between 0.1 and 0.2\\,$M_\\sun$ now have accurate dynamical masses \\citep{Segransan00}, but objects with masses between the hydrogen burning limit and 0.1\\,$M_\\sun$ so far do not have accurate ($\\la10$\\%) mass and luminosity determinations.  Dependence on theoretical models has led to controversy about the mass of brown dwarfs, particularly when objects are plausibly near the canonical planetary-mass boundary of $\\sim$13\\,M$_{\\rm jup}$ \\citep[e.g.][]{Luhman07}. In the few cases where the dynamical mass of a brown dwarf has been measured \\citep{ZapateroOsorio04,Stassun06}, precision is either inadequate to truly constrain models, or the model fits have to make unusual assumptions such as non-coeval systems \\citep{Stassun07}.  GJ~802 is a M5 field dwarf system at $\\sim$16\\,pc that was discovered to have a brown dwarf component through astrometry \\citep{Pravdo05}. The subsequent detection of the companion, GJ~802B, with aperture-masking interferometry \\citep{Lloyd06} made this an ideal target for dynamical mass determination, due to the system's $\\sim$3 year period. The high contrast ($\\sim$100:1) and small separation ($\\la100$\\,mas) of GJ~802AB make this system too difficult for direct imaging observations, but is in an ideal range for interferometric detection.  This paper reports aperture-masking interferometry detections of GJ~802B at six epochs in three colors, enough to make good measurements of the orbit, and the discovery that GJ~802A is itself a close spectroscopic M dwarf binary. Section~\\ref{sectObservations} describes the astrometric, interferometric, spectroscopic and photometric observations that go into the orbit determinations. Section~\\ref{sectDiscussion} describes the orbit of the close binary GJ~802Aab and Section~\\ref{sectDynamMass} describes the constraint on the dynamical mass of GJ~802B from the orbit of GJ~802AB. In Section~\\ref{sectModelComp} we compare our results to theoretical predictions and Section~\\ref{sectConclusion} has a summary and discussion. ",
        "conclusions": "\\label{sectConclusion}  GJ~802 is a triple system with component masses of $\\sim$0.14, 0.14 and 0.063\\,$M_\\sun$. The inner pair of equal-mass stars, GJ~802Aab, has an orbital period of 0.795 days and an inclination between 74 and 83 degrees. The outer component, GJ~802B, has a $\\sim$3 year orbit and a mass of $0.063 \\pm 0.005$\\,M$_\\sun$. The inclination of the GJ~802B orbit, $\\sim$83 degrees, is consistent with co-planarity. The most promising way to decrease the mass uncertainty would be to obtain an accurate radial velocity curve of the 3 year orbit, with an expected amplitude of $\\sim$3\\,km/s.  The brown dwarf ``desert'' is generally defined as a lack of low-mass companions around solar-type stars, but a similar lack of very unequal mass-ratio companions has been found around very low-mass stars \\citep{Close03}. However, as a triple system, GJ~802 no longer fits in this category. In fact, there are a large number of brown dwarfs known in triple or higher order multiple systems. This entirely consistent with dynamical star-formation simulations \\citep{Delgado04}, but may also relate to the details of the fragmentation process as a cloud with relatively high angular momentum collapses. Measuring the mutual inclination (i.e. degree of co-planarity) of this and other low-mass triples with long-baseline interferometry could help to determine the mechanism responsible for producing triples like GJ~802.  The kinematics of GJ~802B place it in the thick disk, at an age of $\\sim$10\\,GYr, while the model-derived age for GJ~802B is $\\sim$2\\,GYr. Although this discrepancy is only significant at the 93\\% confidence level so far (1.5$\\sigma$), we will list several possibilities for reconciling the discrepancy: 1) GJ~802Aab could actually have a total mass of $>$0.30\\,$M_\\sun$, consistent with the astrometry, but be underluminous by $\\sim30$\\% in K-band compared to field stars. GJ~802B would then be placed right on the substellar boundary at $\\sim$0.07\\,$M_\\sun$, a naively unlikely position given the very small sample of astrometric STEPS binaries from which GJ~802 was taken \\citep{Pravdo05}. 2) GJ~802Aab could have an apparent total mass of $>$0.30\\,$M_\\sun$ because of an additional low-mass component in a $\\sim30$\\,day, just stable orbit. This could also increase our dynamical mass of GJ~802B to $\\sim$0.07\\,$M_\\sun$. 3) GJ~802 could actually be $<$2\\,GYr old but have experienced a unique dynamical past that gave it a high space velocity without tearing the wide binary apart. Or, 4) models for old brown dwarfs are systematically under-predicting luminosities. These model errors could relate to, e.g. the effects of magnetic fields that can hinder heat flow in brown dwarfs \\citep{Chabrier07}. Given that GJ~802B is currently the only $\\sim$Gyr or older brown dwarf with an accurate dynamical mass, we find this fourth possibility most likely: a hypothesis that will be testable within the next few years as more field brown dwarfs have accurate mass determinations."
    },
    "0801/0801.2076_arXiv.txt": {
        "abstract": "The observations of a nearby low-luminosity gamma-ray burst (GRB) 060218 associated with supernova SN 2006aj may imply an interesting astronomical picture where a supernova shock breakout locates behind a relativistic GRB jet. Based on this picture, we study neutrino emission for early afterglows of GRB 060218-like GRBs, where neutrinos are expected to be produced from photopion interactions in a GRB blast wave that propagates into a dense wind. Relativistic protons for the interactions are accelerated by an external shock, while target photons are basically provided by the incoming thermal emission from the shock breakout and its inverse-Compton scattered component. Because of a high estimated event rate of low-luminosity GRBs, we would have more opportunities to detect afterglow neutrinos from a single nearby GRB event of this type by IceCube. Such a possible detection could provide evidence for the picture described above. ",
        "introduction": "Gamma-ray burst (GRB) 060218 associated with SN 2006aj discovered by Swift (Campana et al. 2006) provides a new example of low-luminosity GRBs (LL-GRBs), as its isotropic equivalent energy ($\\sim6\\times10^{49}$ erg) is 100 to 1000 times less but its duration ($T_{90}=2100\\pm100$ s) is much longer than those of conventional high-luminosity GRBs. More interestingly, besides an usual non-thermal component in its early X-ray spectrum, a surprising thermal component was observed by the Swift XRT during both burst and afterglow phases. Fitting with a blackbody spectrum, the temperature of this thermal component was inferred to be $kT\\sim0.17$ keV during the first 3 ks. When $t>$10 ks, however, the peak energy of the blackbody decreased and then passed through the Swift UVOT energy range at $\\sim$100 ks (Campana et al. 2006; Blustin 2007).  To explain the prompt emission, in principle, a model based on the internal dissipation of relativistic ejecta may be valid in the case of GRB 060218. The relativistic GRB ejecta, which interacts with a dense wind surrounding the progenitor, has also been required to understand a power-law decaying afterglow in X-ray and radio bands about $\\sim$10 ks after the burst, although some complications beyond the standard afterglow model are involved (Soderberg et al. 2006; Fan et al. 2006; Waxman et al. 2007). Furthermore, as suggested by Wang \\& M\\'esz\\'aros (2006), the soft X-ray thermal emission could arise from a shock breakout, namely, a hot cocoon that breaks out from the supernova ejecta and the stellar wind. In detail, a more rapid part of a jet moving in the envelope and the dense wind is accelerated to a highly-relativistic velocity to produce the GRB, while a slower part of the jet together with the outermost parts of the envelope becomes a mildly relativistic cocoon (M\\'esz\\'aros \\& Rees 2001; Ramirez-Ruiz et al. 2002; Zhang et al. 2003), which locates behind the GRB blast wave.  Although the GRB blast wave runs in front of the shock breakout, the thermal emission from the latter outshines the former persistently until the emission of the breakout is switched off. Thus, the emission properties of the GRB blast wave (consisting of external shock-accelerated electrons and protons) should be influenced by the incoming thermal photons significantly during both burst and early afterglow phases. On one hand, the cooling of the relativistic electrons, which upscatter the thermal photons, could be dominated by inverse-Compton (IC) radiation rather than synchrotron radiation (Wang \\& M\\'esz\\'aros 2006). On the other hand, inferred from the observations, the intensity of the thermal emission could be comparable in the same band to the one of the prompt emission due to internal dissipations and much larger than the one of the afterglow emission due to an external shock. Therefore, the thermal photons as target photons for photopion interactions could also play an important role in the energy loss of the relativistic protons and thus influence or even dominate neutrino emission of the GRB blast wave.  It has been widely studied that conventional GRBs in the standard internal-external shock model emit high energy neutrinos during the burst, early afterglow, and X-ray flare phases (Waxman \\& Bahcall 1997, 2000; Dai \\& Lu 2001; Dermer 2002; Dermer \\& Atoyan 2003; Asano 2005; Murase \\& Nagataki 2006a, 2006b; Murase 2007; Gupta \\& Zhang 2007). In contrast to the conventional GRBs, LL-GRBs may have a much higher event rate (several hundred to thousand events $\\rm Gpc^{-3}yr^{-1}$), which is inferred from the fact that two typical nearby LL-GRBs, i.e., GRB 060218 and GRB 980425, have been observed within a relatively short period of time (Cobb et al. 2006; Soderberg 2006; Liang et al. 2007). The high event rate implies that the contribution from LL-GRBs to the diffuse neutrino background may be important and that we have more opportunities to detect neutrinos from a very nearby single LL-GRB event. Therefore, Murase et al. (2006) and Gupta \\& Zhang (2006) recently studied the neutrino emission properties of LL-GRBs during their burst phase using the internal shock model, in which the target photons for photopion interactions are mainly provided by internal shock-driven non-thermal emission. In this paper, however, we will focus on the early afterglow neutrino emission. During this phase, relativistic protons are accelerated by an external shock (rather than internal shocks) and target photons are dominated by the incoming thermal emission and even its IC scattered component.  This paper is organized as follows. In section 2 we briefly describe the dynamics of a GRB blast wave propagating into a surrounding dense wind. In section 3, we give the photon distribution in the blast wave by considering the incoming thermal emission and its IC scattered component, but the weak synchrotron radiation of the electrons is ignored. In section 4, neutrino spectra are derived formally with an energy loss timescale of protons due to photopion interactions. In section 5, we calculate the timescale and then the neutrino spectra using the target photon spectra obtained in section 3 and an experiential fitting formula of the cross section of photopion interactions. In addition, the peak of a neutrino spectrum is also estimated analytically by using $\\Delta-$approximation. Finally, a summary is given in section 6. ",
        "conclusions": "The surprising soft X-ray thermal emission during both burst and afterglow phases of GRB 060218/SN 2006aj was proposed to be due to the breakout from a strong stellar wind of a radiation-dominated shock. This shock breakout was further thought to locate behind a relativistic GRB jet, which is required by understanding the burst emission and the power-law decaying afterglow emission. Wang \\& M\\'esz\\'aros (2006) suggested that a sub-GeV or GeV emission produced by IC scattering of the thermal photons by the relativistic electrons in the GRB blast wave could give evidence for this astronomical picture. In this paper, we studied another possible implication, namely, afterglow neutrino emission. The neutrinos are produced by photopion interactions of relativistic protons, which could be accelerated by a relativistic external shock. The target photons in the interactions are contributed by the incoming thermal emission and its upscattered component. By considering the high event rate of LL-GRBs, we argue optimistically that the afterglow neutrinos from very nearby (several tens of Mpc) LL-GRBs may be detected by IceCube in the following decades. We believe the detection of these expected afterglow neutrinos is helpful to uncover the nature of GRB 060218-like GRBs."
    },
    "0801/0801.1769_arXiv.txt": {
        "abstract": "}[1]{{ \\footnotesize \\noindent {\\bf Abstract} #1 \\\\}} \\renewcommand{\\author}[1]{\\subsubsection*{\\it#1}} \\newcommand{\\address}[1]{\\subsubsection*{\\it#1}}    \\begin{document}  \\chapter*{Stars, gas and dust in elliptical galaxies\\label{pipino1}} \\author{Antonio Pipino$^{a,b}$\\footnote{email: axp@astro.ox.ac.uk},  Francesca Matteucci$^{b,d}$ and Thomas H. Puzia$^e$} \\address{$^a$Astrophysics, Oxford University, Denys Wilkinson Building, Keble Road, Oxford OX1 3RH, UK\\\\ $^b$Dipartimento di Astronomia, Universit\\`a di Trieste, Via G.B. Tiepolo 11, 34100 Trieste, Italy\\\\ $^d$ INAF, Osservatorio Astronomico di Trieste, Via G.B. Tiepolo 11, 34100 Trieste, Italy\\\\ $^e$ Herzberg Institute of Astrophysics, 5071 West Saanich Road, Victoria, BC V9E 2E7, Canada, }    \\abstract{I will present recent theoretical results on the formation and the high redshift assembly of spheroids. These findings have been obtained by utilising different and complementary techniques: chemodynamical models offer great insight in the radial abundance gradients in the stars; while state semi-analytic codes implementing a detailed treatment of the chemical evolution allow an exploration of the role of the galactic mass in shaping many observed relations. The results will be shown by following the path represented by the evolution of the mass-metallicity relation in stars, gas and dust. I will show how, under a few sensible assumptions, it is possible to reproduce a large number of observables ranging from the Xrays to the Infrared. By comparing model predictions with observations, we derive a picture of galaxy formation in which the higher is the mass of the galaxy, the shorter are the infall and the star formation timescales. Therefore, the stellar component of the most massive and luminous galaxies might attain a metallicity $Z \\ge Z_{\\odot}$ in only 0.5 Gyr. Each galaxy is created outside-in, i.e. the outermost regions accrete gas, form stars and develop a galactic wind very quickly, compared to the central core in which the star formation can last up to $\\sim 1.3$ Gyr.  } ",
        "introduction": "Any model of galaxy evolution presented so far had to overcome the strong challenge represented by the observational fact that elliptical galaxies show a remarkable uniformity in their photometric and chemical properties, one of the strongest constraints being the mass-metallicity relation (e.g. Carollo et al. 1993, Davies et al. 1993). The first proposed scenario of elliptical formation was the so-called monolithic collapse scenario (e.g. Larson, 1974). In this framework, ellipticals are assumed to have formed at high redshift as a result of a rapid collapse of a gas cloud. This gas is then rapidly converted into stars by means of a very strong burst, followed by a galactic wind powered by the energy injected into the interstellar medium (ISM) by supernovae (SNe) and stellar winds. The wind carries out the residual gas from the galaxies, thus inhibiting further star formation. In this way the mass-metallicity (i.e. more massive a galaxies have the higher metal content in stars and gas) relation could be easily explained in terms of metallicity sequences, namely the more massive objects develop the wind later (due to their deeper potential wells) and, thus, have more time to enrich their stellar generations. This scenario has been recently reviewed and modified by Pipino \\& Matteucci (2004, PM04) in order to take into account that ellipticals show an increasing Mg/Fe abundance ratio in the stars as a function of galactic mass (Faber et al. 1992). Due to the different nucleosynthesis leading to the production of Mg (by type II SNe, on short timescales) and Fe (by type Ia SNe, on longer timescale), the Mg/Fe-mass relation implies that the more massive objects should have formed faster than the less massive ones (see Matteucci, this book, and references therein). Pipino \\& Matteucci (2004) implemented an infall term in the chemical evolution equation in order to simulate the creation of galaxies and found that most of the photo-chemical observables, including the Mg/Fe-mass relation can be reproduced in a scenario in which the more massive galaxies formed faster and with a much more efficient star formation process with respect to the low mass objects.  PM04 suggested that a single galaxy should form outside-in, namely the outermost regions form earlier and faster with respect to the central parts.  A natural consequence of this model and of the time-delay between the production of Fe and that of Mg is that the mean [Mg/Fe] abundance ratio in the stars should increase with radius. Pipino et al. (2006, PMC06) compared PM04 best model results with the very recent observations for the galaxy NGC 4697 (Mendez et al. 2005), and found them in excellent agreement.  Metallicity gradients, in fact, are characteristic of the stellar populations inside elliptical galaxies. Evidences come from the increase of line-strength indices (e.g. Carollo et al., 1993; Davies et al., 1993; Trager et al., 2000) and the reddening of the colours (e.g. Peletier et al. 1990) towards the centre of the galaxies.  The study of such gradients provide insights into the mechanism of galaxy formation, particularly on the duration of the chemical enrichment process at each radius.  Metallicity indices, in fact, contain information on the chemical composition and the age of the simple stellar populations (SSPs) inhabiting a given galactic zone. ",
        "conclusions": "A detailed study of the chemical properties of the CSPs inhabiting elliptical galaxies as well as the change of their properties as a function of both time and radius, allow us to gather a wealth of information. Our main conclusions are:  \\begin{itemize} \\item Both observed and predicted stellar metallicity distribution for ellipticals show a sharp truncation at high metallicities that, in the light of our models, might be interpreted as the first direct evidence for the occurrence of the galactic wind in spheroids. \\item The stellar component of the most massive and luminous galaxies might attain a metallicity $Z \\ge Z_{\\odot}$ in only 0.5 Gyr. \\item PM04's best model prediction of increasing [$\\rm <\\alpha/Fe>$] ratio with radius is in very good agreement with the observed gradient in [$\\alpha$/Fe] of NGC 4697. This strongly suggests an outside-in galaxy formation scenario for elliptical galaxies that show strong gradients (see also Pipino et al. 2007b and \\ref{pipino2}, this book). \\item By comparing the radial trend of [$\\rm <Z/H>$] with the \\emph{observed} one, we notice a discrepancy which is due to the fact that a CSP behaves in a different way with respect to a SSP. In particular the predicted gradient of [$\\rm <Z/H>$] is flatter than the observed one at large radii. Therefore, this should be taken into account when estimates for the metallicity of a galaxy are derived from the simple comparison between the observed line-strength index and the predictions for a SSP, a method currently adopted in the literature. \\item Abundance ratios such as [Mg/Fe] are less affected by the discrepancy between the SSPs and a CSP, since their distribution functions are narrower and more symmetric. Therefore, we stress the importance of such a ratio as the most robust tool to estimate the duration of the galaxy formation process. \\item We show that the observed multi-modality in the globular cluster metallicity distributions can be, at least partly, ascribed to the radial variation in the underlying stellar populations in giant elliptical galaxies. In particular, the observed globular cluster systems are consistent with a linear combination of the globular cluster sub-populations inhabiting different galactocentric radii projected on the sky.  \\item A new prediction of our models, which is in astonishing agreement with the spectroscopic observations, is the presence of a super-solar metallicity mode that seems to emerge in the most massive elliptical galaxies. In smaller objects, instead, this mode disappears quickly with decreasing stellar mass of the host galaxy.    \\item Since in our scenario the GCs properties sample the galactic CSPs, we predict an increase of the mean metallicity and mean [$\\alpha/Fe$] of cluster systems with the host galaxy mass, which closely follows the scaling relations for ellipticals. Moreover, we expect that a major fraction of the GCs (i.e. those born inside the galaxy) follows an age-metalliticity relationship, in the sense that the older ones are also more $\\alpha$-enhanced and more metal-poor. \\end{itemize}   A.P. thanks the Organizers for having provided financial support for attending the conference."
    },
    "0801/0801.1127_arXiv.txt": {
        "abstract": "We use a suite of cosmological N-body simulations to study the properties of substructure halos (subhalos) in galaxy-sized cold dark matter halos. We extend prior work on the subject by considering the whole population of subhalos {\\it physically associated} with the main system. These are defined as subhalos that have at some time in the past been within the virial radius of the halo's main progenitor and that have survived as self-bound entities to $z=0$.  We find that this population extends beyond {\\it three times} the virial radius, and contains objects on extreme orbits, including a few with velocities approaching the nominal escape speed from the system. We trace the origin of these unorthodox orbits to the tidal dissociation of bound groups of subhalos, which results in the ejection of some subhalos along tidal streams. Ejected subhalos are primarily low-mass systems, leading to mass-dependent biases in their spatial distribution and kinematics: the lower the subhalo mass at accretion time, the less centrally concentrated and kinematically hotter their descendant population. The bias is strongest amongst the most massive subhalos, but disappears at the low-mass end: below a certain mass, subhalos behave just like test particles in the potential of the main halo. Overall, our findings imply that subhalos identified within the virial radius represent a rather incomplete census of the substructure physically related to a halo: only about {\\it one half} of all associated subhalos are found today within the virial radius of a halo, and many relatively isolated halos may have actually been ejected in the past from more massive systems. These results may explain the age dependence of the clustering of low-mass halos reported recently by Gao et al, and has further implications for (i) the interpretation of the structural parameters and assembly histories of halos neighboring massive systems; (ii) the existence of low-mass dynamical outliers, such as Leo I and And XII in the Local Group; and (iii) the presence of evidence for evolutionary effects, such as tidal truncation or ram-pressure stripping, well outside the traditional virial boundary of a galaxy system. ",
        "introduction": "\\label{sec:intro}  In the current paradigm of structure formation, the concordance $\\Lambda$CDM scenario, the dark matter halos that host galaxy systems are assembled hierarchically, through the merger and accretion of smaller subunits. One relic of this process is the presence of {\\it substructure}, which consists of the self-bound cores of accreted subsystems that have so far escaped full disruption in the tidal field of the main halo (Klypin et al 1999; Moore et al 1999).  Although substructure halos (referred to hereafter as ``subhalos'', for short) typically make up only a small fraction ($5$ to $10\\%$) of the total mass of the system, they chart the innermost regions of accreted subsystems, and are thus appealing tracers of the location and kinematics of the {\\it galaxies} that subhalos may have hosted. Substructure is thus a valuable tool for studying galaxies embedded in the potential of a much larger system, such as satellite galaxies orbiting a primary, or individual galaxies orbiting within a group or cluster of galaxies.  This realization has prompted a number of studies over the past few years, both analytical and numerical, aimed at characterizing the main properties of subhalos, such as their mass function, spatial distribution, and kinematics (e.g. Ghigna et al 1998, 1999; Moore et al 1999; Taylor \\& Babul 2005a,b; Benson 2005; Gao et al 2004; Diemand et al 2007a,b).  Consensus has been slowly but steadily emerging on these issues. For example, the {\\it mass function} of subhalos has been found to be rather steep, $dN/dM \\propto M_{\\rm sub}^{-1.9}$ or steeper, implying that the subhalo population is dominated in number by low-mass systems but that most of the substructure mass resides with the few most massive subhalos (Springel et al 2001, Helmi, White \\& Springel 2002, Gao et al 2004). Confirmation of this comes from the fact that the total fraction of mass in subhalos is rather low (typically below 10\\%) even in the highest resolution simulations published so far (although see Diemand et al 2007a,b for a differing view).  Subhalos have also been found to be {\\it spatially biased} relative to the smooth dark matter halo where they are embedded, avoiding in general the innermost regions. Furthermore, the number density profile of the subhalo population also differs markedly from that of galaxies in clusters, and possibly from the radial distribution of luminous satellites around the Milky Way (Kravstov et al 2004; Willman et al 2004; Madau et al 2008). This precludes identifying directly the population of ``surviving'' subhalos with galaxies in clusters, and highlights the need for either more sophisticated numerical modeling techniques, or for pairing up the N-body results with semi-analytic modeling in order to trace more faithfully the galaxy population (Springel et al 2001, De Lucia et al 2004, Gao et al 2004, Croton et al 2006).   One intriguing result of all these studies has been the remarkably weak dependence of the properties of substructure on subhalo mass. Gao et al (2004) and Diemand, Moore \\& Stadel (2004), for example, find that the radial distribution of subhalos is largely independent of their self-bound mass. This is surprising given the strong mass dependence expected for the processes that dictate the evolution of subhalos within the main halo, such as dynamical friction and tidal stripping. Although efficient mixing within the potential of the main halo is a possibility, an alternative explanation has been advanced by Kravtsov, Gnedin \\& Klypin (2004).  These authors argue that the {\\it present-day} mass of a subhalo may be a poor indicator of the {\\it original} mass of the system, which may have been substantially larger at the time of accretion, and used this idea to motivate how the faintest dwarf companions of the Milky Way were able to build up a sizable stellar mass through several episodes of star formation despite their shallow present-day potential wells. The same idea was also adopted by Libeskind et al (2007) as a possible reason for the peculiar spatial alignment of satellites around the Milky Way (Lynden-Bell 1976, 1982; Kunkel \\& Demers 1976; Kroupa, Thies \\& Boily 2005).  We revisit here these issues with the aid of a suite of high-resolution N-body simulations of galaxy-sized halos. We extend prior work by carefully tracking the orbits of surviving subhalos back in time. This allows us to select a complete set of subhalos {\\it physically associated} with the main halo, rather than only the ones that happen to be within the virial radius at a particular time. As we discuss below, a large fraction of the associated subhalo population are on unorthodox orbits that take them well beyond the virial radius, a result with important implications for studies of satellite galaxies and of halos clustered around massive systems.  The plan of this paper is as follows. We introduce briefly the numerical simulations in \\S~\\ref{sec:numexp} and describe our subhalo detection algorithm and tracking method in \\S~\\ref{sec:anal}. Our main results are presented in \\S~\\ref{sec:res}: we begin by exploring the subhalo spatial distribution and kinematics, as well as their dependence on mass, and discuss the consequences of our findings for the subhalo mass function.  We end with a brief summary and discussion of possible implications and future work in \\S~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have used a suite of cosmological N-body simulations to study the orbital properties of substructure halos (subhalos) in galaxy-sized cold dark matter halos. We extend prior work on the subject by considering the whole population of {\\it associated} subhalos, defined as those that (i) survive as self-bound entities to $z=0$, and (ii) have at some time in the past been within the virial radius of the halo's main progenitor. Our main findings may be summarized as follows. \\begin{itemize}  \\item The population of {\\it associated} subhalos extends well beyond three times the virial radius, $r_{200}$, and contains a number of objects on extreme orbits, including a few with velocities approaching the nominal escape speed from the system. These are typically the low-mass members of accreted groups which are propelled onto high energy orbits during the tidal dissociation of the group in the potential of the main halo.  \\item The net result of this effect is to push low-mass subhalos to the periphery of the system, creating a well-defined mass-dependent bias in the spatial distribution of associated subhalos. For example, only about $\\sim 29\\%$ of subhalos which, at accretion time, had peak circular velocities of order $3\\%$ of the present-day virial velocity ($V_{\\rm max}^{\\rm acc} \\sim 0.03 \\, V_{200}$), are found today within $r_{200}$.  This fraction climbs to $\\sim 61\\%$ and to $\\sim 78\\%$ for subhalos with $V_{\\rm max}^{\\rm acc}\\sim 0.1\\, V_{200}$ and $\\sim 0.3 \\, V_{200}$, respectively.  \\item The strength of the bias is much weaker when expressed in terms of the subhalo {\\it present-day} mass, due to the increased effect of dynamical friction and tidal stripping on the most massive subsystems.  \\item The spatial distribution, kinematics, and velocity anisotropy of the subhalo population are distinct from the properties of the dark matter. Subhalos are less centrally concentrated, have a mild velocity bias, and are, near the center, on more tangential orbits than the dark matter.  \\end{itemize}  The unorthodox orbits of substructure halos that result from the complex history of accretion in hierarchical formation scenarios have a number of interesting implications for theoretical and observational studies of substructure and of the general halo population.  One implication is that subhalos identified within the virial radius represent a rather incomplete census of the substructures physically related to (and affected by) a massive halo. This affects, for example, the interpretation of galaxy properties in the periphery of galaxy clusters, and confirms earlier suggestions that evolutionary effects normally associated with passage through the innermost regions of a massive halo, such as tidal truncation or ram-pressure stripping, should be detectable well outside the traditional virial boundaries of a group or cluster (Balogh, Navarro \\& Norris 2000).  Furthermore, associated subhalos pushed well outside the virial radius of their main halo might be erroneously identified as separate, isolated structures in studies that do not follow in detail the orbital trajectories of each system. This effect would be most prevalent at low masses, and it is likely to have a significant effect on the internal properties of halos in the vicinity of massive systems. We expect, for example, halos in the periphery of groups/clusters to show evidence of truncation and stripping, such as higher concentrations and/or sharp cutoffs in their outer mass profiles.  The same effect may also introduce a substantial environmental dependence in the formation-time dependence of halo clustering reported in recent studies (Gao et al 2005; Zhu et al 2006; Jing et al 2007; see also Diemand et al 2007b). In particular, at fixed mass, early collapsing halos might be more clustered because they are physically associated with a more massive system from which they were expelled.  A proposal along these lines has recently been advanced by Wang, Mo \\& Jing (2007) (see also Hahn et al 2008), who argue that such environmental effects might be fully responsible for the age-dependence of halo clustering. Our physical interpretation, however, differs in detail from theirs. Whereas Wang et al argue for the suppression of mass accretion onto ``old'' halos by ``heating by large-scale tidal fields'' as responsible for their enhanced clustering, our results suggest that the real culprit is the orbital energy gain associated with the tidal dissociation of bound groups of subhalos, which allows ``old'' low-mass halos to evade merging and to survive in the vicinity of massive systems until the present.  A further implication of our results concern the spatial bias of the most massive substructures discussed in S.~\\ref{ssec:subhrdist}. If, for example, luminous substructures in the Local Group trace the most massive associated subhalos at the time of accretion, they may actually be significantly more concentrated and kinematically biased relative to the dark matter, a result that ought to be taken into account when using satellite dynamics to place constraints on the mass of the halos of the Milky Way and M31.  Finally, as already pointed out by Sales et al (2007a,b), gravitational interactions during accretion may also be responsible for the presence of dynamical outliers in the Local Group, such as Leo I and And XII. Further work is needed to assess whether the exceptional orbits of such systems could indeed have originated in the tidal dissociation of groups recently accreted into the Local Group. Since the latest proper motion studies of the Magellanic Clouds seem to suggest that the Clouds are on their first pericentric passage (Kallivayalil et al 2006; Piatek et al 2007), this is a possibility to consider seriously when trying to puzzle out the significance of the motion of the satellites of the Local Group.  \\vskip 2.5cm We thank Simon White and Vincent Eke for useful discussions, and an anonymous referee for a constructive report. ADL would like to thank Jorge Pe\\~{n}arrubia and Scott Chapman for many useful discussion which have improved this work. The simulations reported here were run on the Llaima Cluster at the University of Victoria, and on the McKenzie Cluster at the Canadian Institute for Theoretical Astrophysics. This work has been supported by various grants to JFN and a post-graduate scholarship to ADL from Canada's NSERC. AH gratefully acknowledges financial support from NOVA and NWO.   \\begin{table*} \\center \\caption{Properties of simulated halos used in this study.} \\begin{tabular}{l c c c c c  c c c c c c c c } \\hline \\hline Halo & $\\epsilon_{\\rm G}$& $M_{200}$ & $M_{\\rm DM}^{\\rm assoc}$ &$r_{200}$ & $V_{\\rm max}$ & $r_{\\rm max}$ & $N_{200}$ &$M_{\\rm sub}^{\\rm assoc}$ & $N_{\\rm sub}$ & $N_{\\rm sub}$ & $N_{\\rm sub}$ & $N_{\\rm sub}$ & $N_{\\rm sub}$ \\\\ &  [kpc/h]    & [$M_{\\odot}$/h]&[$M_{\\odot}$/h] & [kpc/h] & [km/s] & [kpc/h] &          & [$M_{\\odot}$/h] &     [assoc]    &($r<r_{200}$)  & ($r<r_{100}$) & ($r<r_{50}$)   & $r_{\\rm apo}>2.5r_{\\rm ta}$\\\\ \\hline 9-1-53     &   0.250     &   9.17$\\times10^{11}$& 12.0$\\times10^{11}$&  158.0 &  184.8 & 36.0 &   3.24e6       & 0.89$\\times10^{11}$& 904       &    513   &  742 & 1017 &    3     \\\\ 9-12-46    &   0.181     &   6.44$\\times10^{11}$& 9.79$\\times10^{11}$& 140.4  &  159.9 & 34.9 &   4.82e6       & 0.47$\\times10^{11}$ & 2020      &    865   &  1314& 1828 &     14    \\\\ 9-13-74    &   0.220     &   8.76$\\times10^{11}$& 12.4$\\times10^{11}$& 155.6   &  175.9 & 32.0 &   4.18e6       & 1.04$\\times10^{11}$&  1683      &    831   &  1232& 1645 &   15    \\\\ 9-14-39    &   0.186     &   12.6$\\times10^{11}$& 17.9$\\times10^{11}$& 175.7  &  203.5 & 37.8 &   3.30e6       & 1.27$\\times10^{11}$ & 1416      &    594   &  1050& 1581 &   4     \\\\ 9-14-56    &   0.275     &   8.57$\\times10^{11}$& 12.2$\\times10^{11}$& 154.4   &  178.6 & 33.7 &   2.61e6       &1.00$\\times10^{11}$ &  1160      &    469   &  646 & 930  &   12    \\\\ \\hline \\end{tabular}\\label{tab:numexp} \\end{table*} \\begin{table*} \\center \\caption{Dynamical properties of subhalo population} \\begin{tabular}{c| c c c c c c c c}\\hline Sample                                   &$N_{\\rm sub}$      &  $n_{0}$    & $r_{-2}$ &$r_{\\rm h}$&$\\alpha$&$\\langle \\beta \\rangle$&$\\langle \\sigma_r \\rangle$ \\\\ &[5 sims] & [arb.units] &$[r_{200}]$&$[r_{200}]$&        &                       & $      [V_{200}]$          \\\\\\hline\\hline Assoc. dark matter                       &         &   5422      &  0.l12   &  0.558    & 0.159  &  0.270                &        0.734              \\\\ &         &             &          &           &        &                                                  \\\\ Assoc. subhalos                          & 7183    & 4691.0      &   1.08   &  1.09    & 0.842  &         0.053         &  0.712                    \\\\ &         &             &          &           &        &                                                  \\\\ ($X=V_{\\rm max}^{\\rm acc}/V_{200}$)  &             &        &          &               &           &            &                                    \\\\ ${\\rm X}> 0.170$                         &  150    & 42848.5      & 0.70     & 0.70      &  0.352 &         0.059          &  0.699                \\\\ $0.063 < X < 0.17$                       & 2110    & 3709.2       & 0.80     & 0.87     &   0.754 &         0.091         &  0.712                \\\\ $0.040 < X < 0.063 $                     & 2110    & 875.9        & 1.09     & 1.10     &  0.991 &         0.018         &  0.699                 \\\\ $0.025 < X < 0.040 $                     &  750    & 222.0        & 1.26     & 1.39     &   1.001 &         0.034         &  0.762                \\\\ &         &              &          &           &         &                                             \\\\ ($Y=\\log_{10}M_{\\rm sub}/M_{200}$)  &         &        &          &          &           &                                                        \\\\ $Y> -3.3  $                              &  150    & 28.3         & 1.12     &  1.14     &  1.536 &         0.207         &  0.769                 \\\\ $-4.6  < Y < -3.3$                       & 1918    & 1181.3       & 0.99     &  1.06     &  0.886 &         -0.001        &  0.723               \\\\ $-5.1  < Y < -4.6$                       & 1918    & 1275.5       & 1.11     &  1.08     &  0.897 &         0.021         &  0.735               \\\\ $-5.42 < Y < -5.1$                       &  750    & 341.6        & 1.09     &  1.13     &  0.939 &         -0.028        &  0.708             \\\\ \\end{tabular}\\label{tab:nrho} \\end{table*}"
    },
    "0801/0801.4534_arXiv.txt": {
        "abstract": "{This paper summarizes highlights of the OG1 session of the 30th International Cosmic Ray Conference, held in Merida (Yucatan, Mexico). The subsessions (OG1.1, OG1.2, OG1.3, OG1.4 and OG1.5) summarized here were mainly devoted to direct measurements, acceleration and propagation of cosmic rays. }  \\begin{document} ",
        "introduction": "The origin of cosmic rays is a problem that has been haunting scientists for almost one century now. Solving this scientific problem means putting together numerous pieces of a complex puzzle, in which the acceleration processes, the inner dynamics of the sources, the propagation, the chemical composition all fit together to provide a satisfactory and self-consistent global picture. We are not there yet, though an increasingly larger number of pieces are finding their place in the puzzle. The International Cosmic Ray Conference is a privileged place to observe the puzzle taking shape.  This rapporteur paper is organized as follows: in \\S \\ref{sec:bird} I present a subjective bird's eye view of how things stand in the field, in order to make it easier for the reader to put in the right context the extremely wide range of results presented at the Conference. In \\S \\ref{sec:measure} I summarize some important observational results about spectra and chemical composition of cosmic rays (including the electronic component) that have been presented. In \\S \\ref{sec:accel} I illustrate some recent developments in the understanding of the acceleration of cosmic rays. In \\S \\ref{sec:prop} I discuss the problem of relating the observed cosmic rays to the sources through propagation in the Galaxy. The transition from the galactic component to the extragalactic cosmic ray component is briefly discussed in \\S \\ref{sec:trans} together with some issues related to the propagation of ultra high energy cosmic rays in the Galaxy and in the intergalactic medium. I conclude in \\S \\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  The collection of high quality data is the main drive for the field of cosmic ray research. In the very low energy region (below $\\sim 10$ GeV) the accurate measurements of the cosmic ray flux by BESS has, among other things, allowed to achieve an excellent understanding of the effects of solar modulation. At energies below the knee in the all-particle spectrum, several experiments (CREAM, ATIC, Tracer) are providing us with large statistics of data and correspondingly excellent quality spectra of nuclei spanning the all periodic table. The domain of ultra-heavy nuclei is being covered by the TIGER experiment. These direct measurements are gradually approaching the energy range around the knee thereby providing us with a test of consistency (or inconsistency) of data collected indirectly by ground experiments using atmospheric showers induced by the cosmic ray interactions. One of the goals of the research in this field in the years to come should be (and in fact it is already) to match the results of direct and indirect measurements of cosmic ray fluxes and chemical composition. Thinking of the possibility of ultra-long duration balloon flights to reach closer to or across the knee appears as a worthy effort.  The achievement of accurate measurements of cosmic ray spectra is also opening the way to the search for weak signatures of rare phenomena, such as propagation of antimatter from nearby regions of the universe and annihilation of dark matter in the Galaxy. The latter has profound implications on the spectra of positrons and antiprotons especially, besides creating features in gamma rays and at other wavelengths. The launch of the PAMELA satellite in 2006 has been a milestone in this direction, and hopefully next ICRC will have plenty of presentations of exciting new results.  The picture that emerges from direct cosmic ray measurements can be summarized as follows. There is a satisfactory agreement on the all-particle spectrum as measured by different experiments. The spectra of the chemical species do not show statistically significant differences in the slopes. The abundances of Ga and Ge do not seem immediately compatible with acceleration scenarios based on pure first ionization potential or volatility. Different determinations of the B/C ratio are consistent with each other up to the highest measured energy ($\\sim 1$ TeV), but the error bars are still large enough to leave open the possibility of a flattening in the slope of the B/C ratio as a function of energy. A long duration flight, for instance of Tracer, is likely to settle this issue. We recall that from the theoretical point of view a naive extrapolation of the observed B/C ratio to the knee region would lead to exceedingly large anisotropy.  A few results of measurements carried out with ground experiments, such as HESS and the Tibet hybrid detector were presented in the OG session. The HESS collaboration presented the positive detection of the direct Cherenkov radiation, and used it to infer the spectrum of iron nuclei (more correctly of nuclei with charge $Z>24$) with energy up to $\\sim 100$ TeV. The measured spectrum appears to be in good agreement with other results in the same energy region.  The Tibet Collaboration discussed the results of a measurement of the proton and helium primary spectra as obtained with the Tibet hybrid experiment. The discrepancy between these results and those of the KASCADE experiment are clear, especially for the helium spectrum. As stressed above, one way of facing these problems is to extend the direct measurements with high statistics of events to the knee region.  The calibration of the ground arrays by an overlap with direct measurements is a crucial goal to pursue, not only to understand the origin of the knee but also to describe correctly the transition from galactic to extragalactic cosmic rays. The two problems are tightly related to each other.  There have been numerous presentations at the Conference on the theory of acceleration of cosmic rays. From the phenomenological point of view it is clear that several independent pieces support the possibility that the bulk of galactic cosmic rays are accelerated in SNRs: 1) X-ray observations have shown that effective magnetic field amplification at SNR shocks takes place; 2) multifrequency observations of SNR RXJ1713 and Vela junior are best fit if the observed gamma ray emission is of hadronic origin; 3) the spectra of radio and gamma rays show some evidence, though probably not conclusive, of curvature (concavity), a phenomenon which is expected in non linear theories of particle acceleration at shock waves, when the dynamical reaction of the particles is not negligible, namely when acceleration is very efficient.  From the theoretical point of view, the origin of the magnetic field amplification is being widely investigated. Two mechanisms might be at work at the same time: streaming instability induced by the efficiently accelerated cosmic rays and turbulent amplification in a shocked medium with density inhomogeneities. Other instabilities, such as firehose instability, might also be at work. There have been several contributed papers on the growth rate of the non resonant streaming instability in the context of a kinetic approach, and as simulated with the help of PIC simulations. The latter do not seem to confirm the large growth rates predicted by quasi-linear theory, but additional work is needed to make sure that the initial conditions are set in a physically meaningful way.  Efficient magnetic field amplification, when it is due to streaming instability, also implies efficient particle acceleration, which in turn induces a concavity in the spectra of both the accelerated particles and the radiations produced by them. Efficient magnetic field amplification is also required in order to explain cosmic ray energies at least in the knee region for protons and correspondingly higher for nuclei with higher charge.  All these ingredients are put together in multifrequency investigations of SNRs, where evidence for magnetic field amplification comes from X-ray astronomy. Convincing evidence was presented that such multifrequency spectra, from radio to gamma rays, can be explained in the framework of particle acceleration at modified shocks. It was also claimed that the propagated spectra of different chemical components, when summed together provide a good fit to the observed all-particle spectrum of cosmic rays.  Unfortunately propagation in the Galaxy is all but well understood, mainly because of our ignorance of both the regular and turbulent components of the galactic magnetic field: the competition between parallel and perpendicular diffusion, the presence of a wind, the topology of the magnetic field in the spiral arms and in between the arms all affect the escape times and their dependence on energy. A hint of what might be going on could come from future extensions of the primary to secondary ratios to higher energies, possibly approaching the knee region."
    },
    "0801/0801.1582_arXiv.txt": {
        "abstract": "The existence of supermassive black holes lurking in the centers of galaxies and of stellar binary systems containing a black hole with a few solar masses has been established beyond reasonable doubt. The idea that black holes of intermediate masses ($\\sim 1000$ \\msun) may exist in globular star clusters has gained credence over recent years but no conclusive evidence has been established yet. An attractive feature of this hypothesis is the potential to not only disrupt solar-type stars but also compact white dwarf stars. In close encounters the white dwarfs can be sufficiently compressed to thermonuclearly explode. The detection of an underluminous thermonuclear explosion accompanied by a soft, transient X-ray signal would be compelling evidence for the presence of intermediate mass black holes in stellar clusters. In this paper we focus on the numerical techniques used to simulate the entire disruption process from the initial parabolic orbit, over the  nuclear energy release during tidal compression, the subsequent ejection of freshly synthesized material and the formation process of an accretion disk around the black hole. ",
        "introduction": "\\label{intro} The existence of two classes of black holes is well-established: supermassive black holes with masses beyond $10^6$ M$_\\odot$ are thought to lurk in the centers of most galaxies\\cite{richstone98} and black holes with just a few solar masses have been identified as the unseen components in X-ray binary systems\\cite{lewin06}. Black holes with $M_{\\rm bh}\\sim 1000$ M$_\\odot$, so-called ``intermediate mass black holes'', represent a plausible, but so far still unconfirmed ``missing link'' between these two well-established classes of black holes. In recent years, the stellar dynamics in the centers of some globular clusters has been interpreted as being the result the gravitational interaction with an intermediate mass black hole\\cite{gebhardt02,gerssen02,gerssen03,gebhardt05}. Further circumstantial evidence comes from  ultraluminous, compact X-ray sources in young star clusters\\cite{zezas02,pooley06} and from n-body simulations\\cite{portegies04} that indicate that runaway collisions in dense young star clusters can lead to rapidly growing black holes. None of these arguments is conclusive by itself, but their different nature suggests that the possibility of intermediate mass black holes must be taken seriously.\\\\ The tidal disruption of white dwarfs offers the unique possibility to explore the presence of an intermediate mass black hole. The corresponding disruption processes have been explored with various approximations in earlier studies \\cite{luminet89b,wilson04,dearborn05}, our simulations for the first time explore the full evolution from the initial parabolic orbit over the disruption process to the subsequent build-up of an accretion disk. We have performed a large set of calculations to identify observational signatures that can corroborate or, alternatively, rule out the existence of intermediate mass black holes. In this paper we focus on the numerical techniques that have been employed in this disruption study, a detailed discussion of the astrophysical implications will be given elsewhere\\cite{rosswog07f,rosswog07d}. ",
        "conclusions": ""
    },
    "0801/0801.3778_arXiv.txt": {
        "abstract": "I discuss two of the possible sources of biases in the determination of the IMF: binning and the existence of unresolved components. The first source is important for clusters with a small number of stars detected in a given mass bin while the second one is relevant for all clusters located beyond the immediate solar neighborhood. For both cases I will present results of numerical simulations and I will discuss strategies to correct for their effects. I also present a brief description of a third unrelated bias source. ",
        "introduction": "%  Most initial mass functions (IMFs) of simple stellar populations are calculated by positioning the observational data in a color-magnitude (CMD) or in a Hertzsprung-Russell (HRD) diagram. From stellar models one extracts the isochrone appropriate for the age and metallicity and compares it with the data to obtain the evolutionary masses for the individual stars. Stars are then grouped into bins with equal width in $\\Delta(\\log m)$ and a function (e.g. $dn/dm = Am^\\gamma$) is fitted to the binned data using $\\chi^2$ minimization.  The above procedure is subject to different biases that can influence the derived parameters. In this contribution, I discuss two of them and present my plans to analyze a third. The first bias is the one induced by using bins of equal width when fitting power-laws to binned data. Such an effect is more general than its application to the calculation of mass functions and was discussed in Paper I \\citep{MaizUbed05}. The second bias is caused by the existence of unresolved multiple systems, either physical or chance alignments, and was discussed in Paper II \\citep{Maiz08a}. The third bias is related to the effect of random uncertainties in the mass determinations on the computed IMF and will be analyzed in a future paper. ",
        "conclusions": ""
    },
    "0801/0801.2915_arXiv.txt": {
        "abstract": "We present a new method to derive the age of young ($<$0.7 Gyr) L dwarfs based on their near-infrared photometry, colours, and distances. The method is based on samples of L dwarfs belonging to the Upper Sco association (5 Myr), the Alpha Per (85 Myr) and Pleiades (125 Myr) clusters, and the Ursa Major (400 Myr) and Hyades (625 Myr) moving groups. We apply our method to a number of interesting objects in the literature, including a known L dwarf binary, L dwarf companions, and spectroscopic members of the young $\\sigma$ Orionis cluster. ",
        "introduction": "Low-mass stars and brown dwarfs undergo a significant change in luminosity over the first hundred million years (hereafter Myr) of their life \\citep{baraffe98,burrows01}. As brown dwarfs fade inexorably with time, estimating their age is important to infer fundamental parameters such as the mass. A large number of age diagnostics and methods exist, including the main-sequence turn-off \\citep{mermilliod81}, the lithium depletion boundary \\citep*{rebolo92}, rotation and activity \\citep{randich01}, kinematics \\citep{reid02,dahn02}, and belonging to a moving group \\citep{montes01}. However uncertainties remain both on the observational and theoretical sides: for example ages derived from the main-sequence turn-off method tend to be half those inferred from the lithium depletion boundary \\citep{jeffries01}.  Large-scale sky surveys have uncovered a large population of ultracool dwarfs (defined as dwarfs with spectral types later than M7) \\citep[e.g.][]{delfosse99,kirkpatrick00,geballe02}. These discoveries required a new class of objects cooler than M dwarfs to be defined, the L dwarfs. L dwarfs were proposed by \\citet{kirkpatrick99} and \\citet{martin99} to describe the disappearance of titanium oxide and vanadium oxide and the onset of metal alkali at optical and infrared wavelengths. Old field L dwarfs represent a mixture of very-low-mass stars and brown dwarfs \\citep{kirkpatrick00}. They have typical effective temperatures between 2200 and 1400 Kelvins \\citep{basri00,leggett00} and exhibit red optical and near-infrared colours \\citep{knapp04}. There are currently about 500 L dwarfs\\footnote[1]{See http://spider.ipac.caltech.edu/staff/davy/ARCHIVE/, a webpage dedicated to L and T dwarfs maintained by C.\\ Gelino, D.\\ Kirkpatrick, and A.\\ Burgasser} known, mainly uncovered by large-scale sky surveys \\citep[e.g.][]{kirkpatrick00,chiu06} and proper motion studies \\citep[e.g.][]{cruz03}, including 76 with parallaxes \\citep{perryman97,dahn02,vrba04}.  In this paper we present a scheme to infer the age of L dwarfs using their near-infrared photometry, assuming that they are single objects with a known distance. We have chosen five samples of young, intermediate-age, and old late-type dwarfs in clusters and moving groups to derive an age-colour relationship independent of theoretical models. In the first part of this paper we describe the samples used in our study (Sect.\\ \\ref{age:sample}), including field L dwarf members of the Hyades and Ursa Major moving groups, Pleiades proper motion members, \\APer{} photometric candidates, and Upper Sco spectroscopic members. In the second part we discuss the method to determine the age of L dwarfs (Sect.\\ \\ref{age:determination}). Finally, we apply our method to several L dwarfs with known distance published to date (Sect.\\ \\ref{age:application}).  \\begin{table} \\begin{center} \\caption{Summary of the four L dwarf members of the Hyades moving group. Magnitudes from 2MASS and infrared spectral types are listed for each object along with its distance derived from parallax measurement.} \\label{tab_age:HyadesMG_memb} \\begin{tabular}{@{\\hspace{0mm}}l c c c c c@{\\hspace{0mm}}} \\hline 2MASS J\\ldots{}  &   $J$  &   $H$  & $K_{s}$ & d$_{p}$       & SpT   \\\\ \\hline 00361617$+$1821104  & 12.466 & 11.588 & 11.058 & 8.8 $\\pm$0.1 & L4    \\\\ 00325937$+$1410371  & 16.830 & 15.648 & 14.946 & 33.2$\\pm$6.9 & L8    \\\\ 01075242$+$0041563  & 15.824 & 14.512 & 13.709 & 15.6$\\pm$1.2 & L5.5  \\\\ 08251968$+$2115521  & 15.100 & 13.792 & 13.028 & 10.7$\\pm$0.1 & L6    \\\\ 12171110$-$0311131  & 15.860 & 15.748 & 15.887 & 11.0$\\pm$0.3 & T7.5  \\\\ \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "\\label{Pleiades:conclusions}  We have presented a new method to infer the age of young (age $\\leq$0.7 Gyr) L dwarfs from their infrared photometry alone, assuming that their distance is known and they are single objects of solar like metallicity. \\citet{eggen96} finds that for a sample of nearby lower main sequence stars 53 percent have ages of less than 1 Gyr and we would expect a similar result for L dwarfs. Thus the technique will be useful for many L dwarfs.  Our method can only be as accurate as the calibrating clusters used and it is possible that differences in metallicity between the clusters could have a significant effect on the determination of the age. It would be very helpful to have some theoretical guidance on how the metallicity affects the isochrones which would involve running the structure and atmospheric models for different metallicities. At the very least the method is capable of saying whether a given dwarf is older or younger than the calibrating clusters with an accuracy of $\\pm$0.2 in $\\log$(age). However, relative ages can be measured more precisely. With a dynamic range $\\sim$2.5 magnitudes in absolute magnitude, it has the potential, with further observations and theoretical input, to become a very accurate and powerful method. Hence, we anticipate that this method will evolve and improve with the discovery of a larger number of L dwarfs with a wider range of ages as well as with updated theoretical models."
    },
    "0801/0801.2124_arXiv.txt": {
        "abstract": "The understanding of the formation, life, and death of Population III stars, as well as the impact that these objects had on later generations of structure formation, is one of the foremost issues in modern cosmological research and has been an active area of research during the past several years. We summarize the results presented at ``First Stars III,'' a conference sponsored by Los Alamos National Laboratory, the Kavli Institute for Particle Astrophysics and Cosmology, and the Joint Institute for Nuclear Astrophysics.  This conference, the third in a series, took place in July 2007 at the La Fonda Hotel in Santa Fe, New Mexico, U.S.A. ",
        "introduction": "In July 2007, more than 130 international researchers met at the La Fonda Hotel in Santa Fe, New Mexico, U.S.A. to discuss the formation, life, and death of zero-metallicity (Population III) and very low metallicity stars, as well as the impact that these objects had on later generations of stars and on structure formation.  This field has made significant theoretical and observational advances since First Stars I, which was a MPA/ESO/MPE Joint conference held in 1999 in Garching b.\\ M\\\"unchen, Germany, and First Stars II, which was held at The Pennsylvania State University in State College, Pennsylvania, U.S.A., in 2003.  Though major advances have been made, the understanding of Population III and very low metallicity stars is still in its infancy, and many important questions remain unanswered.  Donald Rumsfeld is known for many things, but his contribution to the philosophy of science is less well recognized. In an oft-cited speech he delivered on Feb 12, 2002, he said:  \\begin{quote} There are known knowns. These are things we know that we know. There are known unknowns. That is to say, there are things that we know we don't know. But there are also unknown unknowns. There are things we don't know we don't know. \\end{quote}  \\noindent He applied this to intelligence data, but one can describe our task as scientists is to uncover ``unknown unknowns\" by pure thought or by serendipitous observation or experiment, thereby converting them to ``known unknowns;\" and to then use systematic observation, experiment and theory to convert these to \"known knowns.\"  The study of the first stars is a new field. There are almost no ``known knowns,\" except for the background cosmology and an increasing set of data on abundances in very metal poor stars. There are many ``known unknowns:\" What is the nature of dark matter? What is the strength of the magnetic field? What are the abundance and size distribution of dust? What is the rate of mixing of metals into primordial gas? Until recently, the effect of dark matter annihilation on the first stars was an ``unknown unknown,\" but this has become a ``known unknown\" through the work of Freese and her collaborators.  We conjecture that the number of ``unknown unknowns\"---i.e., the discovery potential---in a field is proportional to the ratio of the number of ``known unknowns\" to ``known knowns:\" \\begin{equation} UU \\propto KU/KK. \\end{equation} If so, the study of the First Stars is an excellent field for young people---as reflected in the youth of the audience and the organizers (although not in the summarizer)!  In the following sections, we shall summarize the results presented at this conference, and raise some significant issues that have yet to be explored. In the interests of conserving space, we discuss only the results presented during First Stars III, and direct interested readers to individual contributions for more in-depth information and citations to refereed papers.  In addition, we have made a concerted effort to refer only  to contributions in this proceedings, and do so by the last name of one of the authors (typically the first) of each contribution. ",
        "conclusions": ""
    },
    "0801/0801.0910_arXiv.txt": {
        "abstract": "We present here near-infrared spectroscopy in the H and K bands of a selection of nearly 80 stars that belong to various AGB types, namely S type, M type and SR type. This sample also includes 16 Post-AGB (PAGB) stars. {\\bf From these spectra, we seek correlations between the equivalent widths of some important spectral signatures and the infrared colors that are indicative of mass loss. Repeated spectroscopic observations were made on some PAGB stars to look for spectral variations. We also analyse archival SPITZER mid-infrared spectra on a few PAGB stars to identify spectral features due to PAH molecules providing confirmation of the advanced stage of their evolution. Further, we model the SEDs of the stars (compiled from archival data) and compare circumstellar dust parameters and mass loss rates in different types.}  {\\bf Our near-infrared spectra show that in the case of M and S type stars, the equivalent widths of the CO(3-0) band are moderately correlated with infrared colors, suggesting a possible relationship with mass loss processes. A few PAGB stars revealed short term variability in their spectra, indicating episodic mass loss: the cooler stars showed in CO first overtone bands and the hotter ones showed in HI Brackett lines. Our spectra on IRAS 19399+2312 suggest that it is a transition object. From the SPITZER spectra, there seems to be a dependence between the spectral type of the PAGB stars and the strength of the PAH features. Modelling of SEDs showed among the M and PAGB stars that the higher the mass loss rates, the higher the [K-12] colour in our sample.} ",
        "introduction": "During their asymptotic giant branch (AGB) stage, intermediate mass stars undergo substantial mass loss triggered by pulsation shocks and radiation pressure (e.g., Willson \\cite{will00}; van Winckel \\cite{vanw03}). Near-infrared spectroscopy is one of the recognized tools to study the mass loss process (e.g., Kleinmann \\& Hall \\cite{klei86}; Volk, Kwok \\& Hrivnak \\cite{volk99}; Lancon \\& Wood \\cite{lanc00}; Bieging, Rieke \\& Rieke \\cite{beig02}; Winters et al. \\cite{wint03}). The near-infrared JHK bands, contain several important diagnostic lines that can be used as signatures to probe the atmospheres of these stars (Hinkle, Hall \\& Ridgway \\cite{hink82}).  One of the most dramatic manifestations of the AGB stage is the high rate of mass loss that is mainly attributed to pulsational levitation of atmosphere followed by expulsion of matter by the radiation pressure (Wood \\cite{wood79}; Bowen \\cite{bowe88}; Gail \\& Sedlmayr \\cite{gail88}; Vassiliadis \\& Wood \\cite{vass93}; Habing \\cite{habi96}; Feast \\cite{feas00}; Winters et al. \\cite{wint00}; Willson \\cite{will00}; Woitke \\cite{woit03}; Garcia-Lario \\cite{garc06}). One of the issues that we would address is the possible signatures of pulsational mass loss in the spectra of AGB stars. We would also like to study the infrared spectral signatures that the intermediate mass stars leave while evolving beyond AGB stage to become Planetary Nebulae (e.g., Kwok \\cite{kwok00})  In this work we report near-infrared (H and K bands) spectroscopy on a selection of about 80 AGB stars of different types, namely M type, S type, Semi-regular variables (SR) and post-AGB (PAGB) candidates (Section 2). In Section 3.1, we try to seek significant differences among these types and correlate spectral signature strengths with pulsation parameters or mass loss indicators (such as the infrared color indices) in case of M and S and SR types.   \\begin{figure*} \\centering \\includegraphics[scale=0.8]{figure1.eps} \\caption{Mt Abu near-infrared spectra of a selection of AGB and PAGB stars. From top to bottom the panels show PAGB, M, S and SR type stars giving two examples per category.} \\label{nirspec}% \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[scale=0.6,angle=-90]{figure2.eps} \\caption{Plots of Equivalent Widths (in \\AA) of the CO(3-0) line (left panels) and CO(2-0) line (right panels) with the IR colors [H-K] for M stars (shown as open circles in top panels) and [K-12] for S type stars (filled triangles in bottom panels). The Spearman rank correlation coefficient ($\\rho$) and its significance (s) are given in the corresponding figures.} \\label{co2030}% \\end{figure*}  Further, we present the SPITZER archival spectra in the mid-infrared region on a few post-AGB stars in order to find out the evolutionary stage of these stars in terms of the spectral features seen in the spectra (Section 3.2). Following this, using the DUSTY code, we model the infrared spectral energy distributions of these sample stars, constructed from archival data in the visible, near and far-infrared regions to investigate basic differences in the photospheric and circumstellar parameters among the various types of AGB stars (Section 3.3). Some interesting aspects on variation of spectral lines in a few PAGB stars are discussed in Section 4. Section 5 lists important conclusions of the present work.  \\begin{center} \\begin{table*} \\caption{2MASS \\& IRAS Colors and Mt Abu Equivalent Widths (in \\AA) of the Spectral Features of the Sample M Stars} \\label{mstar1} \\centering \\begin{tabular}{l c c c c c c c c c c r} \\hline\\hline Star &  Sp. Type  & $\\phi$ & H-K & K-12 & CO(3-0) & CO(4-1) & CO(H)   &    CO(2-0)  &  CO(3-1)& CO(K) & $\\dot{M}$    \\\\ \\hline 02234-0024 &  M4e   & 0.34 & 0.44  &  2.31  & 16.1 &  7.8 & 51.5 &  21.1 &  6.9 &  61.9  &  $2.6(-6)^{\\dagger}$ \\\\ 03082+1436 &  M5.5e & 0.33 & 0.45  &  1.99  & 11.3 &  9.2 & 64.2 &  25.1 & 13.3 &   62.3 &  $9.8(-7)^{\\dagger}$\\\\ 03489-0131 &  M4III & ...  & 0.31  &  1.86  &  4.8 &  5.3 & 61.9 &  18.4 & 20.2 &  45.9  &  1.3(-6)\\\\ 04352+6602 &  Sev   & 0.44 & 0.50  &  1.22  &  ... &  ... &  ... &  31.8 & 29.8 & 101.2  &  $1.4(-7)^{\\dagger}$ \\\\ 05265-0443 &  M7e   & 0.57 & 0.49  &  1.32  & 13.6 & 10.1 & 53.1 &  15.1 &  7.8 &  22.9  &  $4.6(-7)^{\\dagger}$ \\\\ 05495+1547 &  S7.5  & 0.49 & 0.59  &  2.79  & 11.6 &  8.4 & 47.4 &  11.9 & 16.0 &  47.6  &  3.5(-6) \\\\ 06468-1342 &  S     & ...  & 0.52  &  1.03  & 14.5 & 16.6 & 85.0 &  32.9 & 31.2 & 110.6  &  ...\\\\ 06571+5524 &  Se    & 0.59 & 0.64  &  1.67  & 20.0 & 22.1 &105.6 &  18.3 & 25.6 & 52.8   &  $1.7(-6)^{\\dagger}$ \\\\ 07043+2246 &  S6.9e & 0.36 & 0.18  &  1.17  & 11.6 & 13.9 & 77.0 &  34.3 & 28.8 &  98.8  &  $3.1(-7)^{\\dagger}$ \\\\ 07092+0735 &  Se    & 0.30 & 0.71  &  3.02  & 16.4 & 15.4 & 72.1 &  20.5 & 25.8 &  81.0  &  2.9(-6) \\\\ 07197-1451 &  S     & 0.26 & 0.55  &  2.32  & 11.8 & 8.0  & 65.0 &  22.3 &  9.7 &  43.5  &  4.6(-6) \\\\ 07299+0825 &  M7e   & 0.57 & 0.36  &  1.90  & 16.2 & 10.5 & 64.1 &  26.8 & 14.1 &  65.3  &  $8.5(-7)^{\\dagger}$ \\\\ 07462+2351 &  SeV   & 0.96 & 0.52  & -0.90  & 12.3 & 13.8 & 68.8 &  28.1 & 33.9 & 108.0  &  ... \\\\ 07537+3118 &        & 0.09 & 0.53  &  1.40  & 13.7 &  9.9 & 76.5 &  22.8 & 25.0 &  77.5  &  ...\\\\ 07584-2051 &  S2,4e & 0.19 & 0.44  &  1.00  & 12.6 & 11.0 & 78.4 &  21.2 & 17.9 &  39.1  &  ... \\\\ 08188+1726 &  SeV   & ...  & 0.41  &  2.28  & 14.3 & 13.6 & 66.1 &  22.2 & 24.8 &  71.3  &  ... \\\\ 08308-1748 &  S     & ...  & 0.43  &  0.72  & 10.1 & 15.9 & 72.6 &  25.7 & 19.1 &  92.9  &  1.7(-7) \\\\ 08416-3220 &        & ...  & 0.41  &  0.97  & 12.2 & 14.2 & 75.8 &  33.0 & 35.0 & 105.3  &  ...\\\\ 08430-2548 &        & ...  & 0.46  &  1.51  & 12.2 & 13.0 & 68.3 &  19.6 & 16.6 & 80.2   &  ...\\\\ 09411-1820 &  Me    & 0.31 & 0.49  &  1.63  &  5.1 & 7.3  & 33.2 &   8.5 &  5.7 &  22.1  &  $7.9(-7)^{\\dagger}$ \\\\ 12372-2623 &  M4    & 0.92 & 0.40  &  1.39  &  9.9 & 14.8 & 64.2 &  20.2 & 23.8 &  50.8  &  ... \\\\ 17001-3651 &  M     & 0.70 & 0.58  &  2.17  &  9.6 &  8.9 & 57.0 &  11.3 &  8.9 &  35.8  &  $1.9(-6)^{\\dagger}$ \\\\ 17186-2914 &  S     & ...  & 0.70  &  1.76  & 14.5 & 13.2 & 80.7 &  23.9 & 26.1 &  89.2  &  2.9(-6) \\\\ 17490-3502 &  M     & 0.60 & 0.54  &  2.26  & 16.4 & 16.1 & 78.1 &  24.6 & 17.6 &  95.8  &  ...\\\\ 17521-2907 &  Se    & 0.11 & 0.56  &  1.64  & 18.9 & 19.5 & 88.4 &  20.1 & 21.9 &  71.8  &  4.6(-6)\\\\ 18508-1415 &  SeV   & 0.44 & 0.38  &  1.00  &  3.4 &  5.1 & 64.3 &  8.8  &  5.5 &   31.1 &  ...\\\\ 19111+2555 &  S     & 0.57 & 0.80  &  4.36  & 19.1 & 17.6 & 66.8 &  13.3 & 11.2 &  24.5  &  $4.6(-6)^{\\dagger}$ \\\\ 19126-0708 &  S..   & 0.90 & 0.79  &  3.81  & 21.6 & 19.1 &102.0 &  21.5 & 23.5 &  59.4  &  $2.9(-6)^{\\dagger}$ \\\\ 19354+5005 &  S..   & 0.24 & 0.52  &  2.28  & 26.6 & 22.8 &108.0 &  30.2 & 33.8 &  98.2  &  $1.5(-6)^{\\dagger}$ \\\\ 23041+1016 &  M7e   & 0.87 & 0.52  &  2.40  &  7.3 &  9.0 & 50.7 &  12.4 &  6.3 &  18.7  &  $1.5(-6)^{\\dagger}$\\\\ \\hline \\end{tabular} \\end{table*}  \\begin{table*} \\caption{2MASS \\& IRAS Colors and Mt Abu Equivalent Widths (in \\AA) of Spectral Features of the Sample S Stars} \\label{sstar1} \\centering \\begin{tabular}{l c c c c c c c c c r} \\hline\\hline Star &  Sp. Type  & H-K  & K-12  & CO(3-0) & CO(4-1) & CO(H)     &  CO(2-0)  &  CO(3-1)& CO(K) & $\\dot{M}$   \\\\ \\hline 04599+1514 &  M0       &  0.49  &  0.96  & 12.8 & 17.3 & 80.8 &  24.2 & 28.6 & 99.4 &  $1.4(-7)^{\\dagger}$\\\\ 05036+1222 &           &  0.60  &  1.09  & 12.7 & 15.1 & 80.9 &  26.2 & 28.1 & 54.3 &  ...\\\\ 05374+3153 &  M2Iab    &  0.16  &  1.44  & 13.7 & 12.2 & 77.5 &  21.6 & 26.6 & 88.3 &  $7.3(-7)^{\\dagger}$  \\\\ 07247-1137 &  Sv..     &  0.34  &  0.85  &  8.6 & 10.7 & 57.8 &  25.8 & 19.1 & 74.5 &  ... \\\\ 07392+1419 &  M3II-III &  0.09  & -1.44  &  5.3 &  8.4 & 38.9 &  16.7 & 14.7 & 55.3 &  7.3(-7)  \\\\ 09070-2847 &  S        &  0.25  &  0.91  & 10.7 & 12.8 & 74.1 &  22.1 & 37.1 & 59.2 &  ...\\\\ 09152-3023 &           &  0.47  &  1.21  &  4.0 &  9.4 & 44.2 &  18.1 & 14.4 & 39.1 &  ... \\\\ 10538-1033 &  M...     &  0.29  &  0.74  &  8.1 &  4.5 & 44.6 &  22.9 &  9.6 & 64.9 &  ... \\\\ 11046+6838 &  M...     &  0.40  &  0.39  &  3.7 &  9.5 & 41.1 &  15.0 & 10.4 & 54.7 &  ... \\\\ 12219-2802 &  S        &  0.30  &  0.72  &  9.2 &  9.4 & 57.0 &  22.1 & 17.8 & 69.8 &  ...\\\\ 12417+6121 &  SeV      &  0.41  &  0.95  & 16.7 & 18.0 & 82.4 &  28.2 & 29.5 & 92.2 &  $1.0(-7)^{\\dagger}$ \\\\ 13421-0316 &  K5       &  0.22  &  0.57  &  ... &  8.3 & 40.9 &  20.0 & 14.9 & 59.8 &  ...\\\\ 13494-0313 &  M...     &  0.42  &  0.56  &  3.0 &  9.2 & 55.5 &  17.4 & 20.0 & 47.4 &  ...\\\\ 14251-0251 &  S...     &  0.34  &  0.71  &  3.7 &  9.1 & 45.8 &  20.2 & 16.7 & 60.0 &  ...\\\\ 15494-2312 &  S        &  0.38  &  0.62  &  5.1 &  8.1 & 46.1 &  19.8 & 17.2 & 59.3 &  ...\\\\ 16205+5659 &  M...     &  0.18  &  0.78  &  5.4 &  7.1 & 45.3 &  13.0 & 19.1 & 61.6 &  ...\\\\ 16209-2808 &  S        &  0.54  &  0.92  & 10.7 & 13.6 & 73.9 &  21.9 & 27.2 & 79.5 &  1.0(-7)\\\\ \\hline \\end{tabular} \\end{table*}   \\begin{table*} \\caption{2MASS \\& IRAS Colors and Mt Abu Equivalent Widths (in \\AA) of Spectral Features of the Sample SR Stars} \\label{srstar1} \\centering \\begin{tabular}{l c c c c c c c c c c r} \\hline\\hline Star &  Sp. Type  & $\\phi$ & H-K & K-12 & CO(3-0) & CO(4-1) & CO(H)     &  CO(2-0)  &  CO(3-1)& CO(K) & $\\dot{M}$   \\\\ \\hline 04030+2435 &  S4.2V    & ...  & 0.40  &   0.63  &  4.1 &   8.3  &  45.0  &\t 16.2 &   17.9  &   34.1  &  ...\\\\ 04483+2826 &  CII      & ...  & 0.38  &   1.04  &  7.0 &  10.9  &  60.3  &\t 24.5 &   17.6  &   58.3  &  1.4(-7)\\\\ 05213+0615 &  M4       & 0.41 & 0.39  &   1.37  &  9.3 &  11.1  &  64.8  &\t 19.3 &   21.4  &   56.9  &  ...\\\\ 05440+1753 &  S0.V     & 0.33 & 0.66  &   1.89  & 22.7 &  19.5  &  99.5  &\t 33.2 &   30.3  &   88.1  &  1.4(-6)\\\\ 06333-0520 &  M6       & 0.85 & 0.38  &   1.20  & 19.7 &   10.1 &   77.8 &\t 19.1 &    24.8 & 61.5    &  1.7(-7)\\\\ 08272-0609 &  M6e      & ...  & 0.40  &   1.47  & 11.3 &   7.8  &  51.7  &\t 19.1 &   15.1  &    81.4 &  5.6(-7)\\\\ 08372-0924 &  M5II     & ...  & 0.36  &   1.57  &  5.5 &   6.5  &  52.3  &\t 18.5 &   12.4  &   49.0  &  1.7(-7)\\\\ 08555+1102 &  M5III    & ...  & 0.26  &   1.28  &  7.4 &   7.9  &   61.8 &\t 15.5 &    13.9 &   43.7  &  2.4(-7)\\\\ 09076+3110 &  M6IIIaSe & ...  & 0.31  &   1.20  & 13.5 &  11.6  &  65.9  &\t 18.5 &    18.8 &   56.7  &  $3.8(-7)^{\\dagger}$\\\\ 10436-3459 &  S        & ...  & 0.37  &   1.61  & 10.5 &  15.4  &  61.2  &\t 30.3 &   23.9  &   79.0  &  $9.5(-7)^{\\dagger}$\\\\ 15492+4837 &  M6S      & 0.77 & 0.40  &   1.58  & 12.7 &   13.0 &  63.8  &\t19.9  &   24.5  &   44.4  &  $9.5(-7)^{\\dagger}$\\\\ 16334-3107 &  S4.7:.V  & ...  & 0.43  &   1.13  & 15.0 &  22.9  &  93.4  &\t 22.5 &   29.7  &  104.8  &  $1.0(-7)^{\\dagger}$\\\\ 17206-2826 &  M..      & ...  & 0.42  &   0.82  & 10.3 &  11.4  &  68.1  &\t 15.2 &   19.6  &  63.8   &  $1.0(-7)^{\\dagger}$\\\\ 17390+0645 &  SV...    & ...  & 0.36  &   0.92  & 12.1 &   7.6  &  56.7  &\t 29.8 &   14.2  &  70.9   &  ...\\\\ 19497+4327 &           & ...  & 0.60  &   0.82  &  6.7 &  11.5  &  59.2  &\t 25.2 &   18.1  &  73.8   &  ...\\\\ \\hline \\end{tabular} \\end{table*}   \\begin{table*} \\caption{2MASS \\& IRAS Colors and Mt Abu Equivalent Widths (in \\AA) of Spectral Features of the Sample PAGB Stars} \\label{pagb1} \\centering \\begin{tabular}{l c c c c c c c c r} \\hline\\hline Star   &  Sp. Type &  H-K & K-12 & Pa$\\beta$ & Br$\\gamma$ & CO(K) & 21$\\mu$m & $T_{d}$ & $\\dot{M}$\\\\ \\hline Z02229+6208& G9a     & 0.44  & 6.38  & ...   &  9.2  &  8.4  & ...   & 300      & 2.9(-5)\\\\ 04296+3429 & G0Ia    & 0.49  & 7.34  &  ...  &  4.1  &  17.7 & -2.11 & 270      & $4.5(-5)^{\\dagger}$\\\\ 05113+1347 & G8Ia    & 0.25  & 5.98  &  ...  &  6.5  &  12.9 & ...   & 350      & 2.1(-5) \\\\ 06556+1623 & Bpe     & 1.11  & 5.10  &  ...  & -52.4 &  ...  & ...   & 1200/150 & 1.0(-5)\\\\ 07134+1005 & F5Ia    & 0.10  & 6.45  &  ...  &  5.1  &  ...  & -1.73 & 280      & $2.3(-5)^{\\dagger}$\\\\ 07284-0940 & K0Ibpv  & 0.23  & 5.64  &  ...  &  ...  &  10.9 & ...   & 450      & 1.0(-5) \\\\ 07430+1115 & G5Ia    & 0.44  & 6.35  &  ...  &  3.9  &  28.9 & ...   & 320      & 3.3(-5) \\\\ 17423-1755 & Be      & 1.48  & 5.30  &  ...  &  -5.6 & -23.6 & ...   & 125      & $4.9(-6)^{\\dagger}$ \\\\ 18237-0715 & Be      & 0.55  & 2.29  & -25.6 & -14.9 &  ...  & ...   & 70\t    & 5.8(-6) \\\\ 19114+0002 & G5Ia    & 0.27  & 4.84  & ...   & 5.4   & ...   & ...   & 120      & $4.9(-5)^{\\dagger}$\\\\ 19157-0247 & B1III   & 1.19  & 5.76  & ...   &  ...  & -10.8 & ...   &  270     & $1.1(-6)^{\\dagger}$\\\\ 19399+2312 & B1e     & 0.11  & 4.80  & -6.2  & -15.6 &  ...  & ...   & ...      & ...\\\\ 22223+4327 & F9Ia    & 0.24  & 4.44  &  7.5  &  5.8  &  ...  & -1.31 & 125      & $1.0(-5)^{\\dagger}$\\\\ 22272+5435 & G5Ia    & 0.38  & 5.55  &  4.4  &  3.4  &  30.2 & -0.83 & 250      & $1.6(-5)^{\\dagger}$ \\\\ 22327-1731 & A0III   & 0.90  & 4.95  &  3.8  &  ...  &  28.7 & ...   & 800      & 5.0(-6)\\\\ 23304+6147 & G2Ia    & 0.36  & 6.48  &  3.4  &  5.1  &  12.8 & -2.53 & 250      & $3.5(-5)^{\\dagger}$\\\\ \\hline \\end{tabular} \\end{table*} \\end{center} ",
        "conclusions": "\\begin{enumerate}  \\item {\\bf The equivalent widths of CO (3-0) show a trend of correlation with [H-K] in M type; while in S stars the CO (3-0) shows a positive correlation with [K-12].}  \\item {\\bf The Mt Abu H and K band spectra also showed marked differences in the three types.}  \\item {\\bf Archival SPITZER spectra of a few PAGB stars showing PAH emission features confirm their advanced stage of evolution into PPNe. There appears to be a dependence of the strength of the PAH features on the spectral type.}  \\item {\\bf Modeling of SEDs of a number of the programme stars showed marked differences in various generic types. Model mass loss rates for M and PAGB stars show an increasing trend with the color [K-12] in our sample stars.}  \\item {\\bf Our observations showed that several PAGB stars undergo short term spectral variability that is indicative of on-going episodic mass loss. Cooler PAGB stars showed variation in CO first overtone lines. In contrast, the hot PAGB stars showed variation in HI lines. The hot PAGB star IRAS 19399+2312 is identified as a transition object.}  \\end{enumerate}"
    },
    "0801/0801.1707_arXiv.txt": {
        "abstract": "%  We present a new set of multi-million particle SPH simulations of the formation of disk dominated galaxies in a cosmological context. Some of these galaxies are higher resolution versions of the models already described in Governato et al (2007).  To correctly compare simulations with observations we create artificial images of our simulations and from them measure photometric Bulge to Disk (B/D) ratios and disk scale lengths.  We show how feedback and high force and mass resolution are necessary ingredients to form galaxies that have flatter rotation curves, larger I band disk scale lengths and smaller B/D ratios. A new simulated disk galaxy has an I-band disk scale length of 9.2 kpc and a B/D flux ratio of 0.64 (face on, dust reddened). ",
        "introduction": "Here we present results from an ongoing project aimed at simulating a representative sample of galaxies that form in a $\\Lambda$CDM cosmology. Simulations are run with GASOLINE \\citep{wadsley04} and use the ``blastwave feedback'' described in \\cite{stinson06}. Three of the four models presented here use the same initial conditions as in Governato et al. (2007, G07) and thus the galaxy halos have the same merging history and form inside the same large scale structure. However, these new simulations use 8 times more resolution elements and a 50\\% smaller force softening than those presented in G07.  To understand the numerical limitations of our work and to better compare with the existing literature, we vary the number of resolution elements and the feedback implementations.  In one case (named MW1) we run three models with the number of resolution elements varying from a few tens of thousand to a few million particles and force resolution from 1.2kpc to 0.3kpc.  To allow a straightforward comparison between simulations and observations, the galaxies' B/D ratios and disk scale lengths are measured using GALFIT on dust reddened, face-on I and K band artificial images produced using {\\sc SUNRISE} \\citep{jonsson06}. Recent works by different groups and using different hydrodynamical codes have highlighted the role of feedback in creating galaxies with a dominant stellar component supported by rotation \\citep{okamoto05,robertson04}. Significant improvements have been made in reproducing a number of scaling properties of late type galaxies: namely the age - circular velocity, the Tully Fisher, the stellar mass - metallicity relations and the abundance of galaxy satellites \\citep{g07,brooks07,libeskind07}. However, a well known problem of most galaxies formed in cosmological simulations is a massive, centrally concentrated spheroidal component that creates unrealistically declining rotation curves \\citep{eke01,g07}.  The explanations proposed in the literature invoke both physical and numerical arguments: catastrophic angular momentum loss due to dynamical friction suffered by gas rich subhalos \\citep{navarro96}, spurious gravitational torques between the disk and the surrounding halo of dark matter and gas caused by the noisy mass distribution at low resolution \\citep{kaufmann07}, or artificial pressure gradients at the cold/hot gas interface \\citep{okamoto03}. \\begin{figure}[!ht] \\plottwo{mayer1.ps}{mayer2.ps} \\caption{Left: The kinematic disk, bulge and halo stellar mass fraction (blue, red and orange) in the MW1-HR, MW1-MR and MW1-MR-TH models. All runs use the ``blastwave'' feedback, but the one marked ``TH'', where ``thermal'' feedback is instead used. MR is the close equivalent of the MW1 galaxy described in G07. Histograms are normalized to 100\\%. The total stellar masses differ by less than 10\\%.  Right: The circular velocity profile from the total mass distribution (V$_c$ = (M/r)$^{1/2}$) for the same galaxy halo, run with different resolutions and feedback implementations. The circular dot marks the radius that contains 85\\% of the stars. HR, MR and LR runs use a total of 5 million, 600k and 70k DM, gas and star particles within the virial radius.} \\end{figure} \\begin{figure}[!ht] \\plottwo{mayer3.ps}{mayer4.ps} \\caption{Left: The  rotation curves for the higher resolution version equivalent of the models in G07. The dashed line is the curve for MW1 published in the paper, equivalent to the new ``MR'' model. The dotted line is the low resolution MW1-LR run. The vertical lines mark the radius at which V$_{rot}$ is measured: 3.5 I-band R$_d$s. This is a radius were the amount of included mass (DM, stars \\& gas) is close to converging (Fig.1b) and gives more reliable results compared to measuring V$_{rot}$ at 2.2 R$_d$ where numerical effects might still be present.  Right: the Tully Fisher relation for a sample of four high resolution runs vs. data in Giovanelli et al 97) Large filled dots: SSPs from Girardi et al (2000). Smaller filled dots: SSPs from Starburst99. Open square: MW1-LR. } \\end{figure}   \\begin{table*} \\centering \\begin{tabular}{cccl} \\hline\\hline \\noalign{\\smallskip} Run & I-band R$_d$ (kpc) & Kinem. B/D  & u -r \\\\ \\hline \\noalign{\\smallskip} MW1-HR       & 3.5 & 0.83 (0.38) & 1.6   \\\\ MW1-MR       & 3.3 & 1.56 & 1.7  \\\\ MW1-MR-TH     & 2.1 &1.51 & 1.6 \\\\ MW1-LR        & 2.1  &1.35 &2.0  \\\\ \\hline \\end{tabular} \\caption[Summary of the photometric properties of the MW1 runs] {Summary of the photometric properties of the MW1 run obtained using SUNRISE and GALFIT. The value in brackets is the photometric B-band B/D ratio for MW1-HR. MW1 has a total mass of 10$^{12}$ M$\\odot$} \\label{tbl-3} \\end{table*}   SN feedback is a promising solution as it decouples baryons from the dark matter (Zavala et al., 2007).  However, \\cite{naab07} recently presented results from simulations that did not include feedback from SNe: the resulting circular velocity profile of their simulated galaxies changed significantly with increasing resolution, with the peak velocity dropping from 320 km/sec to 220 when the number of particles was increased from 40$^3$ to 200$^3$ (their Fig.1). \\cite{kaufmann07} used simulations of isolated galaxies to show that cold, rotationally supported disks lose a large fraction of thier initial angular momentum if the number of gas particles is below $10^6$.  The literature presents encouraging evidence that both a physically motivated description of feedback and a high number of resolution elements help to form galaxies with at least some of the properties of the real ones.  However, an analysis of a uniform sample of simulations with techniques {\\it directly comparable} to those used by observers (who measure the light rather than the underlying mass distribution) has been lacking, making a comparison between observations and theoretical models more open to interpretation. ",
        "conclusions": ""
    },
    "0801/0801.3969_arXiv.txt": {
        "abstract": "Spectroscopic observations of three lenticular (S0) galaxies (NGC~1167, NGC~4150, and NGC~ 6340) and one SBa galaxy (NGC~2273) have been taken with the 6-m telescope of the Special Astrophysical Observatory of the Russian Academy of Sciences aimed to study the structure and kinematic properties of early-type disk galaxies. The radial profiles of the stellar radial velocities and the velocity dispersion are measured. $N$-body simulations are used to construct dynamical models of galaxies containing a stellar disk, bulge, and halo.  The masses of individual components are estimated for maximum-mass disk models. A comparison of models with estimated rotational velocities and the stellar velocity dispersion suggests that the stellar disks in lenticular galaxies are ``overheated''; i.e., there is a significant excess velocity dispersion over the minimum level required to maintain the stability of the disk. This supports the hypothesis that the stellar disks of S0 galaxies were subject to strong gravitational perturbations. The relative thickness of the stellar disks in the S0 galaxies we consider substantially exceed the typical disk thickness of spiral galaxies. ",
        "introduction": "Early-type disk galaxies are galaxies of morphological types S0\\mbox{--}S0/a (lenticular galaxies) and Sa with properties similar to those of lenticular galaxies. The structure of early-type disk galaxies is similar to that of later-type spirals: they have a massive stellar disk and, in many cases, also a dynamically decoupled stellar circumnuclear disk, a developed bulge, and a dark halo that determines the rotational velocity of their outer regions. They differ from most later-type spirals in their higher (on average) bulge luminosities, the low contrast or even total lack of their spiral arms, the very low surface density of gas ({HI}), and, consequently, their extremely weak star formation.  Explaining the observed features of early-type disk galaxies poses a number of problems. First and foremost, it is unclear whether lenticular galaxies represent a logical extension of the  Sd--Sa morphological sequence of galaxies, which reflects the conditions for their formation and the nature of their ensuing ``quiescent'' evolution, or whether their peculiarities are caused by  their interaction with the  environment (mergers, accretion of small satellites, loss of gas due to the pressure of the external medium).  Indeed, lenticular galaxies include many objects whose structure suggests an appreciable external influence (e.g., dynamically and chemically decoupled circumnuclear  disks, peculiarities of the radial brightness profile, or peculiar structural features, such as polar rings). S0 galaxies in rich clusters appear to form as a result of the direct effect of the intergalactic gas on the interstellar medium, directly or indirectly leading to a reduction in the amount of gas and a ``halt'' of active star formation. Environmental effects evidently play a key role in this case, as is demonstrated by the lower percentage of early galaxies in distant clusters (the Butcher--Oemler effect [1, 2]), and the lower contemporary rates of the current star formation in galaxies located in denser environments~[6]. However, early-type galaxies also include quite a few field galaxies, which may have different histories.  The problem of the gas content in lenticular galaxies is equally interesting. Even when HI is present in detectable amounts, its total mass is at least an order of magnitude lower than would be expected to result from a simple return to the interstellar medium of gas ejected by evolved disk stars~[4]. The scarcity of data on the thickness of the stellar disks in   lenticular galaxies makes it difficult to compare them with other galaxies in terms of the volume gas density or the gas pressure in the disk plane.  The relative mass fraction of  the dark halos in early disk galaxies also remains an open question. This is due, first and foremost, to difficulties in estimating rotation curves at large galactocentric distances based on stellar absorption lines. Fairly extended HI rotation curves have been obtained for a small number of lenticular galaxies. According to Noordermeer~[5] and Noordermeer et al.~[6], the rotational velocities in these galaxies often decrease toward the periphery, but they still imply the presence of fairly massive dark halos.  The large scatter of data in the Tully\\mbox{--}Fisher relation (luminosity--rotation velocity diagram) for lenticular galaxies suggests substantial inhomogeneity of their properties (see, e.g.,~[7] and references therein). It also follows from the analysis of the velocity dispersions of old stars in the galaxy disks, which shows that in some early-type disk galaxies, the stellar velocity dispersion substantially exceeds the minimum level required for the gravitational stability of the disk, whereas, in later-type spiral galaxies, the velocity dispersion of the disk stars usually appears to be close to its threshold value~[8]. However, the relatively low accuracy of the estimated velocity dispersions for disk stars beyond the bright bulges and problems with decomposing the velocity dispersion into $r, \\varphi$, and $z$ components lead us to treat this conclusion as being tentative.  A comparison of the observational data with dynamical models in which the velocity dispersion of the disk stars---both in the plane of the disk and perpendicular to this plane---is close to the critical values for the dynamical stability of the disk provides insight into its dynamical evolution. A comparison of model ($c_{\\ell}$) and observed ($c_{\\textrm{obs}}$) line-of-sight velocity dispersions for old disk stars can reveal one of three possible situations.  (1) $c_{\\textrm{obs}}<c_{\\ell}$. If interpreted in terms of disk stability, it would imply that the mass of the disk adopted in the model is overestimated, and that the disk must be ``less massive'' in order to satisfy the conditions of dynamical stability against perturbations in the plane of the disk and against bending perturbations.  (2) $c_{\\textrm{obs}} = c_{\\ell}$ within the measurement  errors. In this case, it appears that dynamical instabilities during the formation of the bulk of the disk mass have brought the stellar disk to a marginally stable state, where the stellar velocity dispersion is determined by the surface density of the disk, its rotational velocity as a function of galactocentric distance $r$, and its internal structure. This appears to be the most common case among noninteracting spiral galaxies.  (3) $c_{\\textrm{obs}}> c_{\\ell}$; i.e., the observed stellar velocity dispersion exceeds the model values corresponding to marginal disk stability. In this case, there is reason to believe that disk stars have acquired an excess (i.e., above the level required for stability of the disk) energy of random motions during their evolution, so that the disk has become overheated. This can be viewed as evidence that the stellar population of the disk subsystem of the galaxy has been subject to strong gravitational perturbations, for example, as a result of mergers of massive stellar or gaseous satellites, or of close interactions with nearby neighbors. In principle,  the dynamical heating in the inner part of the galaxy can also be related to the disruption of a high-contrast stellar bar.  \\begin{figure*} \\includegraphics{Zasov.eps} \\caption{$V$-band images of the galaxies NGC~1167 (left) and NGC~6340 (right) on logarithmic brightness scales, showing the raw images (top) and the images after the  model surface brightness distribution subtracted (bottom).} \\end{figure*}  Constructing dynamical galaxy models that can yield estimates of the disk-to-halo mass ratio or disk thickness requires estimates of the rotation velocity and stellar velocity dispersion at the largest possible galactocentric distances, preferably along the main  axes of the galaxy, to make it possible to determine the velocity dispersion along both the radial and vertical directions.  In this paper, we describe spectroscopic observations and the results of our construction of dynamical models for four early-type disk galaxies. Table~1 lists the basic parameters of these galaxies we have adopted. Figure~1 shows images of two of the galaxies with low-contrast structure taken with the 6-m telescope of the Special Astrophysical Observatory of the Russian Academy of Sciences. Table~2 gives a log of the observations.  The galaxy luminosities in Table~1 correspond   to the total $B_T$ magnitudes from the   \\mbox{HYPERLEDA} database [11]. The last column gives the radius of the exponential stellar disk $R_{\\textrm{max}}$ which was adopted in galaxy models. This radius corresponds to the isophotal radius $R_{25}=D_{25}/2$, or the radius beyond which the photometric profile steepens. ",
        "conclusions": "\\textbf{Dynamical state of the disk.} Of the four galaxies considered, the observed parameters of only one---the SBa galaxy NGC~2273---are consistent with the hypothesis that the stellar disk was not subject to dynamical heating, with the stellar velocity dispersion corresponding to the minimum level sufficient to maintain a quasi-stationary equilibrium state. In this respect, NGC~2273 resembles many later-type spiral galaxies, whose stellar velocity dispersions are close to their threshold levels throughout a considerable fraction of the disk (see, e.g., [8, 22, 36, 37]).  The disks in the other three (lenticular) galaxies are ``hotter'' even for the maximum allowed mass. This can be viewed as evidence that these galaxies were subject to external gravitational perturbations in the past, most likely as a result of close interactions with nearby galaxies or mergers of fairly large satellites. Numerical simulations (see, e.g., [38]) confirm the efficiency of the latter process. Gravitational perturbations of the gaseous component of the disk can also explain why these galaxies have rapidly exhausted their interstellar-gas reserves due to a burst of star formation, thereby acquiring the properties characteristic of lenticular systems.    \\textbf{Estimates of ${c_r}$, ${c_z}$ and the thickness of the stellar disk.} In a disk that is stable against bending perturbations, perturbations in its plane due, e.g., to stochastic spiral arms, would primarily increase $c_r$, and hence decrease $c_z /c_r$. Merging of satellites and their passage across the disk should result in a more isotropic distribution of the velocity dispersion. In this case, heating occurs mostly via bending perturbations of the disk (see Ardi et~al. [38] for the discussion of this issue).  For all four galaxies, the velocity-dispersion components $c_z$ and $c_r$ and their ratio $c_z/c_r$ calculated in the marginal-disk model (Fig.~8) decrease with galactocentric distance. Note that $c_z$ does not decrease rapidly enough to keep the marginally stable disk thickness constant: it increases with $r$ in all cases. The same is true for the measured stellar velocity dispersions in the galaxies, although the insufficient accuracy of the velocity-dispersion estimates prevented us from reconstructing the detailed behavior of the radial variation of $c_z/c_r$ from the observations. However, we estimated the average $c_z /c_r$ values by analyzing the variation of the velocity dispersion along the major and minor axes for three of the four galaxies observed---NGC~1167, NGC~2273, and NGC~4150. We found this ratio to be 0.6--0.8 in the interval (1--2)$r_d$ for all three galaxies, and to decrease to 0.4--0.5 at (2.5--3)$r_d$ (for NGC~2273 and NGC~4150).   \\begin{figure} \\includegraphics[scale=0.9]{Zasov8.eps} \\caption{Radial distributions of $c_z/c_r$ in numerical simulations for maximum-disk models at the gravitational stability limit. } \\end{figure}    Table~3 lists the average disk half-thicknesses $z_0$ within the maximum galactocentric radius covered by the observations for the galaxies considered. We calculated these half-thicknesses for the inferred  $c_z$ and vertical density profile $\\textrm{cosh}^{-2}(z/z_0)$. In all cases, the disks turned out to be thicker than the disk of the Milky Way. The average $z_0$ estimate for NGC~4150 should not be taken at face value, it reduces to a  half the tabulated value in the inner region of the galaxy [(1--1.5)$r_d$] (see previous Section). The disk of this galaxy is very ``chubby,'' strongly widens with $r$, so that it grades into the halo at a distance of several radial disk scales.  Note that the inferred $z_0$ values are fairly typical for  the thick disks of spiral  galaxies, where they coexist with more massive thin disks of galaxies like those observed in the Hubble Space Telescope Ultra Deep Field in early stages of their formation~[39].  \\textbf{Estimates of the masses of the disks and halos.} Table~3 lists estimates of the masses of these two main components of the disk galaxies and the mass-to-luminosity ratio for the stellar disk found for  the maximum-disk model. It is evident from the estimates obtained that the mass of the dark halo is fairly high, even when calculated in the maximum-disk model: the ratio of the halo mass to the total mass of the stellar components (disk~$+$~bulge) within the photometric radius is equal to 0.5--0.8. Similar or even higher ratios are also observed in later-type galaxies~[29, 40]. Thus, lenticular galaxies do not differ strongly from galaxies of other types in terms of their dark-mass content.  The mass-to-luminosity ratio $M_d/L_B$ for the disks of the four galaxies are several solar units, which, at least for three of these objects (all except NGC~4150), is quite consistent with the ratio expected for the old stellar population. For NGC~4150, $M_d/L_B\\approx 1.2$ in solar units. Although this ratio was inferred for a model with the maximum disk mass, it is a factor of two to three lower than expected for the othergalaxies. It is unclear why the $M_d/L_B$ ratio is so low; such values are rarely found in galaxies dominated by their old stellar population. Either the disk of this galaxy is deficient in low-mass stars (due to a shallow initial mass function), or the adopted model underestimates the disk mass, e.g., due to an overestimation of the inclination $i$. However, the latter possibility is unlikely, since the inclination is fairly large (56$^{\\circ}$), if our initial assumption that the structure of the disk is axisymmetric is correct.  Recall that NGC~4150 ranks well below the other three galaxies in terms of its luminosity, mass (by about an order of magnitude), and size. The stellar disk of this galaxy appears to have a low volume density. The formation history and evolution of the disk in this case may have differed from those for the disks in the other more massive galaxies."
    },
    "0801/0801.0647_arXiv.txt": {
        "abstract": "{Westerlund\\,2 is a young, massive, obscured stellar cluster of our Galaxy. It harbors the most massive star with well determined parameters, WR20a (82+83M$_{\\odot}$), a dozen very early O-type stars (O3--7), and a wealth of PMS stars. Although of importance, this cluster is still not well known.} {The high-energy properties of this cluster, especially those of its early-type stars are examined in details. The variability of the X-ray sources is investigated. } {A monitoring of the field was performed using three Chandra observations. This dataset probes daily as well as monthly to yearly timescales and provides the deepest X-ray view of the cluster to date.} {The two Wolf-Rayet stars WR20a (WN6ha+WN6ha) and WR20b (WN6ha) were analyzed in detail. They are both very luminous and display very hard spectra, but WR20b does not seem to vary. On the contrary, WR20a, a known eclipsing, colliding-wind binary, brightens in the X-ray domain during the eclipses, i.e. when the collision is seen face-on. This can be explained by the properties of the wind-wind collision zone, whose high density leads to a large absorbing column ($2 \\times 10^{24}$\\,cm$^{-2}$).\\\\ All twelve O-type stars previously classified spectroscopically, two eclipsing binaries previously identified and nine newly identified O-type star candidates are detected in the high energy domain; ten of them could be analyzed spectroscopically. Four are overluminous, but the others present typical $L_X/L_{BOL}$ ratios, suggesting that several O-type objects are actually binaries. Variability at the $\\sim2\\sigma$ level was detected for a majority of the sources, of unknown origin for the putatively single objects. \\\\ Faint, soft, diffuse emission pervades the entire field-of-view but no clear structure can be identified, even at the position of a blister proposed to be at the origin of the TeV source HESS J1023$-$575.\\\\ Finally, the X-ray properties of PMS objects were also investigated, in particular the brightest flaring ones. They provided an additional argument in favor of a large distance ($\\sim$8kpc) for the cluster.} {} ",
        "introduction": "\\subsection{The Westerlund\\,2 cluster}  The young (1$-$2Myr) stellar cluster Westerlund\\,2 lies in the \\hii\\ region RCW49. It is an intense star forming region with numerous Pre Main Sequence (PMS) objects, as enlighted by recent Spitzer data \\citep[and references therein]{chu05}. The cluster also contains many early-type stars: at least two Wolf-Rayet stars, WR20a and WR20b \\citep{vdH}, and twelve very hot O-stars with spectral types ranging from O3 to O6.5 \\citep{rau07}. In the IR and radio domains, these stars appear surrounded by two shells, one enclosing WR20a and the core of Westerlund\\,2, the other being around WR20b.  There are still many uncertainties about the properties and content of this cluster. For example, its distance is still a matter of intense debate. To cite only the most recent papers: \\citet{asc07} and \\citet{tsu07} proposed a `low' value of 2--5kpc, based on the X-ray and IR luminosities of PMS stars; while \\citet{rau07} rather argue in favor of $8.0 \\pm 1.4$\\,kpc, based on the spectrophotometric properties of the O-type star population and on the eclipsing characteristics of WR20a. \\citet{Dame} recently discussed the distance of the Westerlund\\,2 cluster on the basis of the CO emission of the giant molecular cloud MC8 (likely associated with the cluster) and the H\\,{\\sc i} 21\\,cm absorptions along the line of sight towards Westerlund\\,2. This author found that the cluster must be located on the far side of the Carina spiral arm and inferred a kinematic distance of $6.0 \\pm 1.0$\\,kpc in coarse agreement with the value of \\citet{rau07}. An indirect argument can be added in favor of a `high' value, using the properties of the diffuse emission detected by \\citet{tow04}. The latter authors found an X-ray luminosity of about 8--$9\\times10^{32}$ erg s$^{-1}$ (for an adopted distance of 2.3 kpc). Considering that the diffuse emission from massive star forming regions typically amounts to 2--$5\\times10^{32}$ erg s$^{-1}$ per early-type star \\citep{tow03,tow04}, we would expect values of 3--$7\\times10^{33}$ erg s$^{-1}$ for a cluster containing at least 14 objects earlier than O7 (including the two Wolf-Rayet systems). This is compatible with the observed emission only if the distance is rather large (4-7kpc rather than 2.3kpc). In this paper, we will thus adopt 8kpc as the distance to Westerlund\\,2.  \\begin{figure*} \\centering \\caption{Three-color images of Westerlund\\,2 in X-rays (a, left: full field; b, right: central 8' in the middle and c, below: core of the cluster). The red, green and blue colors correspond to soft (0.3--1. keV), medium (1.--2. keV) and hard (2.--10. keV) photon energies, respectively. The images were corrected by the exposure maps and adaptively smoothed prior to RGB combination.} \\label{3col} \\end{figure*} \\setcounter{figure}{0} \\begin{figure}[ht] \\centering \\caption{Continued.} \\end{figure}  \\subsection{X-ray properties}  Studying the X-ray emission of young clusters provides crucial information, especially if the cluster is heavily obscured as is the case of Westerlund\\,2. PMS objects, mainly T Tauri stars, are more easily detected in the X-ray and IR domains, and their high-energy flares can be analysed to get their decay times, which are related to the physical size of the flaring region. In the high-mass range, the interacting colliding wind binaries and young magnetic objects are more easily identified and their properties constrained in the X-ray domain. Finally, the large-scale interaction with the surrounding medium is better studied through the analysis of the diffuse X-ray emission.  The X-ray properties of Westerlund\\,2 have been studied in increasing detail with years. The first detection of the cluster was made by the Einstein observatory: \\citet{gol87} reported the presence of one X-ray emitter at the position of RCW49, interpreted as unresolved sources or truly diffuse emission (from a supernova remnant or a wind-blown bubble). Using ROSAT, \\citet{bel94} detected possible diffuse emission and 24 point sources (7 of them being in our ACIS-I field-of-view), among which HD90273, MSP18\\footnote{The MSP numbers come from the numbering scheme devised by \\citet{mof91}.} and WR20a. The high X-ray luminosity of WR20a was suggested as a possible evidence of an on-going colliding wind phenomenon. More recently, \\citet{tow04} and \\citet{tsu07} reported the results of a first Chandra observation. Thanks to its high spatial resolution, this facility is the first to resolve the individual sources in the core of Westerlund\\,2. More than 400 X-ray sources were identified by \\citet{tsu07}, mostly PMS stars and a few early-type objects. Faint diffuse emission is also present; it extends to the south-east of the cluster and presents a soft spectrum plus a very faint hard tail \\citep{tow04}.  To investigate further the high-energy properties of Westerlund\\,2, especially those of its massive star population, we have obtained additional X-ray observations, which combine to provide the deepest X-ray view of the cluster. It enables us to study the variability of the X-ray sources on several timescales (short-term variations within one exposure; monthly to yearly changes between the datasets).  The paper is organized as follows. Section 2 presents the datasets and their reduction. Sections 3 and 4 focus on the O-stars and WR-stars content, respectively. The search for a counterpart to the recently discovered TeV source and the presence of diffuse emission are discussed in section 5, while section 6 concentrates on the flaring sources. Finally, we conclude in section 7. ",
        "conclusions": "In this paper, we have analyzed three Chandra observations aimed at monitoring the massive objects of Westerlund\\,2. In this heavily obscured region, 24 O-type stars (identified or candidates) were detected and analyzed. Most of them display typical values of their $L_X/L_{BOL}$ ratio, though the scatter is rather large (about 60\\%) due to the rather low signal-to-noise ratio of the X-ray spectrum, but also the much poorer knowledge of their bolometric luminosities and multiplicities. The most X-ray luminous objects (MSP18, 165, 167 and 188) are probably binaries and need a specific follow-up. About half of the O-type stars also exhibit variability at the 2$\\sigma$ level or more, whose origin is unknown in these supposedly single objects.  WR20a, the most massive binary, displays a hard and heavily absorbed X-ray spectrum. It brightens by about 40\\% in the X-ray domain during the photometric eclipses. The X-ray emission is believed to come mostly from the wind-wind collision region, which would easily explain the presence of a high temperature in the best-fit models. If the colliding wind region is sufficiently large and opaque, a modulation will appear when the system rotates and it should then be brighter when seen `face-on', as is the case during eclipses. The other Wolf-Rayet star of the field, WR20b, is thought to be single and its flux is accordingly constant. However, the star also exhibits a high temperature and a high X-ray luminosity, and its hardness ratio is even larger than for WR20a. These latter characteristics are more typical of colliding-wind binaries.  The photometry of the other X-ray sources was examined for each observation. Many objects appear variable, and we focused on the eight brightest, flaring ones. Their lightcurve decay time and spectral properties enabled us to infer rather compact flaring regions. The comparison between the optical and X-ray properties of the X-ray sources again suggest a large distance for Westerlund\\,2."
    },
    "0801/0801.2532_arXiv.txt": {
        "abstract": "{We investigate the capability of detecting, with Simbol-X, non-thermal emission during stellar flares, and distinguishing it from hot thermal emission. We find that flare non-thermal emission is detectable when at least $\\sim20$\\,cts are detected with the CZT detector in the $20-80$\\,keV band. Therefore Simbol-X will detect the non-thermal emission from some of the X-ray brightest nearby stars, whether the thermal vs. non-thermal relation, derived for solar flares, holds.} ",
        "introduction": "Stellar flares are phenomena where the magnetic field rapidly releases large amount of energy. It is supposed that: 1) magnetic reconnections generate non-maxwellian population of fast particles; 2) these particles hit and heat the chromospheric plasma, evaporating it, and producing non-thermal emission (hard X-rays and $\\gamma$-rays); 3) the evaporated chromospheric plasma fills coronal structures and generates thermal emission (mostly in the UV and soft X-ray bands).  In solar data a power-law spectrum, compatible to the bremsstrahlung emission of fast particles hitting a thick target, is observed \\citep[i.e.][]{HudsonRyan1995}. This hard power-law emission precedes the soft X-ray and UV thermal emission. Information on stellar emission in the hard X-ray band is very poor, since most X-ray observations are dedicated only to the soft band. Recently \\citet{OstenDrake2007} presented the detection of non-thermal emission in the hard X-rays during a stellar flare.  The study of the flare non-thermal emission allows to study: 1) the depicted scenario for stellar flares, and hence whether and how the thermal emission is caused by the non-thermal particles; 2) the flare energy balance, since only a small amount of released energy is lost by thermal radiation.  We investigate the Simbol-X capability of detecting non-thermal emission associated with stellar flares, and distinguishing it from a hot thermal one. Results about this issue are important also for other astrophysical subjects where non-thermal emission is likely present (galaxy clusters, SNR). ",
        "conclusions": "Inspecting all the simulated flares we find that: non-thermal emission is detected and recognized with Simbol-X observations whether the CZT detector collects more than $\\sim20$ photons in the $20-80$\\,keV band.  For each simulated spectrum we compute the thermal and non-thermal luminosity. We compare the results obtained with those derived from solar flares by \\citet[][see Fig.~5]{IsolaFavata2007}. We find that for some of the considered flares the non-thermal emission can be detected with Simbol-X if thermal vs. non-thermal solar flare relation is valid also at high luminosities."
    },
    "0801/0801.1795_arXiv.txt": {
        "abstract": "{} {We determine the central mass distribution of galaxy cluster \\rxj using strong gravitational lensing.} {High resolution HST/ACS images of the galaxy cluster \\rxj have enabled us to identify several new multiple image candidates in the cluster, including a 5 image system with a central image. The multiple images allow us to construct an accurate 2-dimensional mass map of the central part of the cluster. The modelling of the cluster mass includes the most prominent cluster galaxies modelled as truncated isothermal spheres and a smooth halo component that is described with 2 parametric profiles. The mass reconstruction is done using a Markov chain Monte Carlo method that provides us with a total projected mass density as well as estimates for the parameters of interest and their respective errors.} {Inside the Einstein radius of the cluster ($\\sim$35$''$, or $\\sim$200 kpc, for a source at redshift 1.8) we obtain a total mass of \\mbox{(2.6 $\\pm$ 0.1) $\\times$ 10$^{14}$ M$_{\\odot}$}. The mass profile of the cluster is well fitted by both a Navarro, Frenk and White profile with a moderate concentration of $c$ = 5.3$^{+0.4}_{-0.6}$ and $r_{200}$ = 3.3$^{+0.2}_{-0.1}$ Mpc, or a non-singular isothermal sphere with velocity dispersion $\\sigma$ = 1949 $\\pm$ 40 km/s and a core radius of $r_{c}$ = 20 $\\pm$ 2 $''$. The mass profile is in reasonable agreement with previous mass estimates based on the X-ray emission from the hot intra-cluster gas, however the X-ray mass estimates are systematically lower than what we obtain with gravitational lensing.} {} ",
        "introduction": "\\label{sec:intro}  A precise knowledge of the masses and mass profiles of galaxy clusters is a key to better understand what the mass and energy in the universe is made of. The standard model of cosmology that has emerged over the past 100 years is remarkably successful in reproducing a wealth of observational phenomena. These include the formation and evolution of structure from the tiny temperature variations seen in the cosmic microwave background to galaxies and clusters of galaxies, the element abundances produced in the Big Bang nucleosynthesis and the apparent acceleration of the expansion of the universe as seen from the brightnesses of distant supernovae. The greatest drawback of the standard model is the need to include two dark components in the energy density of the Universe. The first is dark matter which is needed to explain the formation of structure and the dynamics of both galaxies and clusters of galaxies. The second is dark energy which in turn is needed to make the Universe flat and to explain the supernova data. Both of these dark components are still unexplained although candidate particles for dark matter exist.  Clusters of galaxies are of prime interest when trying to constrain the properties of both dark matter and dark energy. The mass budget and therefore the dynamics of clusters are dominated by dark matter, whereas the mass function and power spectrum of clusters is influenced by both dark matter and dark energy. In order to obtain strong constraints on the various cosmological parameters such as $\\Omega_{\\mathrm{m}}$, $\\sigma_{8}$ and the equation of state $w$ of dark energy, accurate mass estimates for many clusters are necessary. Gravitational lensing is a powerful tool for determining cluster masses since the deflection of light is independent of the nature and dynamical state of the matter in clusters and can therefore provide unbiased estimates of the total masses of galaxy clusters.  An important prediction of numerical simulations with cold dark matter is the so called universal mass profile. For large radii (greater than the scale radius) the density $\\rho$ depends on $r$ like $\\rho(r) \\propto r^{-3}$, the exact value of the slope at small radii (inside the scale radius) is unknown but is believed to be somewhere between $-1.5$ and $-1$.  Although there is strong support for the universal mass profile, other mass profiles such as the family of isothermal sphere profiles have not been ruled out.  By studying the mass profiles of galaxy clusters we can potentially constrain the nature of dark matter through comparison with simulations.  As the most luminous X-ray cluster \\citep{schindler:95} known to date \\rxj is obviously of great interest. It has been studied spectroscopically by \\citet{cohen:02} and extensively with X-rays and Sunyaev-Zel'dovich effect \\citep{schindler:97,pointecouteau:99,pointecouteau:01,allen:02,gitti:04,gitti:07}. Gravitational lensing based only on the arcs has also been used to estimate the mass \\citep{sahu:98} as well as weak lensing \\citep{fischer:97,kling:05}. Recently a combination of both the weak and strong lensing has also been used \\citep{bradac:05}. The lack of space-based observations with the Hubble Space Telescope (HST) has so far limited the number of multiple images available for the strong lensing modelling. In this paper we use deep HST images taken with the Advanced Camera for Surveys (ACS) to identify new multiple image candidates that enable us to derive well constrained mass maps for the central regions of \\rxj.  We have obtained spectra of some of the multiple images with the FORS2 spectrograph on the Very Large Telescope in Chile. The spectroscopic redshifts obtained are crucial in fixing the mass scale of the cluster.  The cosmology used throughout this paper is $\\Omega_{\\mathrm{m}}$=0.30, $\\Omega_{\\Lambda}$=0.70 and H$_{0}$=70~km/s/Mpc, unless otherwise stated. With this cosmology an arcsecond at the redshift of the cluster ($z_{\\mathrm{cl}}$=0.451) corresponds to 5.77 kpc. ",
        "conclusions": "\\label{sec:conclusion}  In this paper we present the first detailed strong lensing model for the galaxy cluster \\rxj. The models are based on images taken with the Advanced Camera for Surveys on the Hubble Space Telescope. The high resolution of the image along with the three colours have allowed us to identify several new strongly lensed background images in the cluster. The follow-up spectroscopic observations on the FORS2 at the VLT have provided us with redshifts for the images in one multiple image system. This is crucial for an accurate absolute mass determination of the cluster. Among the new multiple image systems is a 5 image system with a central image that is very important in fixing the mass profile at the core of the cluster.  The mass in the cluster is modelled with small-scale mass components in the cluster galaxies described by truncated isothermal ellipsoids, and a large-scale component in two parametric mass profiles (Navarro, Frenk and White profile and non-singular isothermal ellipsoid). The parameters of the large-scale component are constrained by the location and magnifications of the multiple images. The total mass profile of the cluster is well described by both a Navarro, Frenk and White profile with a moderate concentration of $c$ = 5.3$^{+0.4}_{-0.6}$ and $r_{\\rm{200}}$ = 3.3$^{+0.2}_{-0.1}$ Mpc, or a non-singular isothermal sphere with velocity dispersion $\\sigma_{\\rm{nsis}}$ = 1949 $\\pm$ 40 km/s and a core radius $r_{\\rm{c}}$ = 20.3 $\\pm$ 1.8 $''$. The core radius in the isothermal profile is necessary in order to reproduce the central image and to fit the total mass of the profile at the all radii. A singular isothermal sphere fit to the total mass has too much mass at small radii and too little at larger radii. The total mass of the cluster inside the Einstein radius of the main arc ($\\sim35''$, or $\\sim$200 kpc) is \\mbox{(2.6 $\\pm$ 0.12) $\\times$ 10$^{14}$ M$_{\\odot}$}.  A comparison with the X-ray mass estimates by \\citet{gitti:07} shows that the lensing mass is consistently higher than the X-ray mass, although still within the error bars. The X-ray based velocity dispersion in \\citet{allen:02} on the other hand is in good agreement with the one determined here once the effect of the different core radii is taken into account."
    },
    "0801/0801.4386_arXiv.txt": {
        "abstract": "We compile one of the largest ever samples to probe the X-ray normal galaxy luminosity function and its evolution with cosmic time. In particular, we select 207 galaxies (106 late and 101 early-type systems) from the Chandra Deep Field North and South surveys, the Extended Chandra Deep Field South and the XBOOTES survey. We derive the luminosity function separately for the total (early+late), the early and the late-type samples using both a parametric maximum likelihood method, and a variant of the non-parametric $1/V_m$ method.  Although the statistics is limited, we find that the total (early+late) galaxy sample is consistent with a Pure Luminosity evolution model where the luminosity evolves according to $L(z) \\propto  (1+z)^{2.2}$. The late-type systems appear to drive this trend while the early-type systems show much weaker evidence for evolution. We argue that the X-ray evolution of late-type systems is consistent with that of blue galaxies in the optical. In contrast there is a mismatch between the X-ray evolution of early-type systems and that of red galaxies at optical wavelengths. ",
        "introduction": "The Chandra and XMM missions have led to great progress in our understanding of X-ray emission processes in normal galaxies (see review by Fabbiano, 2005). The X-ray emission in the most luminous early-type systems comes mainly from hot gas, while in the less luminous ones, Low-Mass X-ray binaries (LMXRB) prevail. On the other hand, in late-type systems the X-ray binaries dominate the X-ray emission. In the most actively star-forming ones, it is the High-Mass X-ray binaries (HMXRB) in particular which dominate the X-ray energy budget. In star-forming galaxies, the total galaxy X-ray luminosity is an excellent indicator of star-formation (e.g. Ranalli et al. 2003, Gilfanov et al.  2004).  The study of the galaxy luminosity function at X-ray wavelengths provides the opportunity to study, as an ensemble, the evolution of these systems over cosmological timescales. For example, it is important to investigate whether the evolution of the X-ray emission for star-forming systems is the same as in optical and infrared wavelengths. Ghosh \\& White (2001) suggest that there may be a time-lag because of the onset of the LMXRB, a few Gyr after the star-burst event. In this respect the X-ray luminosity function has important implications for the evolution of X-ray binaries.  One predicament, in the construction of X-ray selected galaxy samples is that, galaxies are intrinsically X-ray faint sources. This implies that we need to probe either very deep in flux or over a wide field-of-view to gather sufficiently large samples. The deep Chandra observations in the area of the Hubble Deep Field North provided the first X-ray selected galaxy sample (Hornschemeier et al. 2003). Norman et al. (2004) derived the X-ray luminosity function finding hints for pure luminosity evolution. The study of Norman et al. (2004) pertains to the total galaxy sample i.e. including both early and late type. At the same time XMM surveys with their large field-of-view,  have helped to constrain the galaxy luminosity function in the local Universe at a median redshift of $z\\sim 0.1$ (Georgakakis et al. 2004, Georgantopoulos et al. 2005, Georgakakis et al. 2006). The luminosity function is well represented by either a Schechter or a log-norm form.  Here, we present the galaxy luminosity function derived from the analysis of  a large Chandra sample of galaxies consisting both of Deep Fields and wide area surveys. $\\rm H_0=72 \\ km~s^{-1} \\ Mpc^{-1}, \\Omega_m=1/3$ and $\\rm \\Omega_\\Lambda=2/3$ are adopted. ",
        "conclusions": "\\subsection{The cosmic evolution} The X-ray luminosity function for galaxies has been derived for the first time by Norman et al. (2004). They used data from the Chandra Deep Fields (CDF-N and CDF-S)  but without discriminating for early and late type systems. Norman et al. (2004) find strong pure luminosity evolution with an evolution index $k=2.7$. Recently, Ptak et al. (2007) constructed again the luminosity function in the same fields. Their more conservative selection criteria result in a smaller sample compared to that of Norman et al. (2004). They derive the non-parametric luminosity function separately in various redshift bins, estimating the evolution index $k$ from the difference in the $L_\\star$. Their late-type sample shows strong Pure Luminosity Evolution, with $k=2.3^{+0.8}_{-0.8}$. In the case of the early-type sample they find some mild evolution $k=1.6^{+1.1}_{-1.0}$, which however is not inconsistent with our results. Georgakakis et al. (2006) derived the evolution in the Chandra Deep Fields using number counts, again separately for early- and late-type samples. They found evolution only for late-type systems ($k=2.7\\pm0.35$) while the early-type ones are consistent with no evolution ($k=0.6^{+0.8}_{-0.6}$). Georgakakis et al. (2007) applying a very careful galaxy selection, aiming to minimize AGN contamination, find an evolution index of $k=2.4$ for the star-forming galaxies in the CDF-N. From the above, it appears then that the results are converging as far as the evolution is concerned: under the assumption of luminosity evolution the late-type systems evolve strongly, while the early-type ones show milder or possibly no evolution. However, clearly better statistics is needed to test whether more complex forms of evolution take place, as for example Luminosity Dependent Density Evolution.   \\subsection{Interpretation and Comparison with the optical}  The optical luminosity function has been derived by Faber et al. (2007) for blue and red galaxies using DEEP2 data. They find that both blue systems evolve vigorously over cosmic time between a redshift of z=1 and today, following pure luminosity evolution with $k\\sim2.7$. The red galaxies evolve in a different manner. Although their luminosity function evolves in luminosity in a similar fashion,its normalization is higher at redshift z=0. Now assuming  that roughly the blue galaxies correspond to our late-type systems and the red to the early-type ones, we can make a comparison between the galaxy evolution probed in X-ray and optical wavelengths. The late-type galaxies have a very similar evolution in both the optical and the X-ray regime. This suggests that the X-ray emission in these galaxies is dominated by HMXRB. Indeed HMXRB evolve on Myr scale and therefore they trace closely the bulk of the optical and IR emission coming from massive stars. In contrast, there is a mismatch in the evolution of early-type systems witnessed in optical and X-ray wavelengths. The absence of vigorous evolution in X-ray wavelengths between z=0.7 and today ($\\sim 6$ Gyr), is much longer than the expected lifetimes of LMXRBs ($0.1-1$Gyr), placing interesting constraints on the formation of these systems."
    },
    "0801/0801.3273_arXiv.txt": {
        "abstract": "We test the hypothesis that the host galaxies of long-duration gamma-ray bursts (GRBs) as well as quasar-selected damped \\lya\\ (DLA) systems are drawn from the population of UV-selected star-forming, high $z$ galaxies (generally referred to as Lyman-break galaxies). Specifically, we compare the metallicity distributions of the GRB and DLA populations against simple models where these galaxies are drawn randomly from the distribution of star-forming galaxies according to their star-formation rate and HI cross-section respectively. We find that it is possible to match both observational distributions assuming very simple and constrained relations between luminosity, metallicity and HI sizes. The simple model can be tested by observing the luminosity distribution of GRB host galaxies and by measuring the luminosity and impact parameters of DLA selected galaxies as a function of metallicity. Our results support the expectation that GRB and DLA samples, in contrast with magnitude limited surveys, provide an almost complete census of $z\\approx3$ star-forming galaxies that are not heavily obscured. ",
        "introduction": "The past 10 years has marked the emergence of extensive observational analysis of high redshift ($z>2$) galaxies. This remarkable and rapid advance was inspired by new technologies in space and ground-base facilities for deep imaging, clever approaches to target selection, and the arrival of 10 m class ground-based telescopes for spectroscopic confirmation. A plethora of classes are now surveyed, each named for the observational technique that selects the galaxies: the \\lya\\ emitters (Hu et al.\\ 1998), the Lyman break galaxies (LBGs, Steidel et al.\\ 2003), sub-mm galaxies (e.g., Chapman et al.\\ 2005), distant red galaxies (van Dokkum et al.\\ 2006), damped \\lya\\ (DLA) systems (Wolfe et al.\\ 2005), extremely red objects (Cimatti et al.\\ 2003), long-duration $\\gamma$-ray burst (GRB) host galaxies (e.g., Fruchter et al.\\ 2006), \\ion{Mg}{2} absorbers, radio galaxies (Miley \\& De Breuck 2008), quasar (QSO) host galaxies, etc. Large, dedicated surveys have identified in some cases thousands of these galaxies providing a direct view into the processes of galaxy formation in the young universe.   Because of the significant differences in the sample selection of high $z$ galaxies, there has been a tendency by observers to treat each population separately and/or contrast the populations. However, the various populations will overlap to some extent and it is important to understand how \\citep[see also][]{Adelberger00,moller,fynbo03,reddy05}.  Of the various galactic populations discovered at $z>2$ to date, only two offer the opportunity to study the interstellar medium at a precision comparable to the Galaxy and its nearest neighbors: the damped \\lya\\ systems intervening quasar sightlines (QSO-DLA) and the host galaxies of GRBs which exhibit bright afterglows (GRB-DLA\\footnote{A small fraction of these GRB afterglow spectra have \\ion{H}{1} column densities in the host galaxy that are below the standard DLA definition $\\mnhi = 2\\times10^{20}\\cm{-2}$ (Jakobsson et al.\\ 2006a; Dessauges-Zavadsky al.\\ 2008, in prep). These GRB-LLS will not be considered here. Similarly, the QSO absorption line systems with \\ion{H}{1} column densities between $10^{17}$ and $2\\times10^{20} \\cm{-2}$ (the empirical threshold for the DLAs) will not be considered here.}). These galaxies are characterized by a bright background source which probes the gas along the sightline to Earth. In the case of QSO-DLA it is the background QSO while for GRB-DLA it is the afterglow of the GRB located within the host galaxy. Hence, the GRB-DLA will not probe the full line-of-sight through the host. The gas in the ISM imprints signatures of the total \\ion{H}{1} column density, the metal content, the ionization state, the velocity fields, and the molecular fraction along the line-of-sight. Because the galaxies are identified in absorption, there is no formal magnitude limit for the associated stellar populations. In this respect, they may trace a large dynamic range in stellar mass, morphology, star formation rate, etc.  The connection between long-duration GRBs and star-forming galaxies has been empirically established.  At large redshift, there is an exclusive coincidence of GRBs with actively star-forming galaxies \\citep[e.g.,][]{hogg99, bloom02,fruchter06}, the majority of which show elevated specific star-formation rates \\citep{chg04}.  At low $z$, there is a direct link between GRBs and massive stars via the detection of spatially and temporally coinciding core-collapse supernovae \\citep[][but see also Fynbo et al. 2006]{hsm+03,smg+03}. Furthermore, GRBs are in some cases found in Wolf-Rayet galaxies \\citep{hfs+06}, which further strengthens the link between GRBs and massive star-formation. The simplest hypothesis, therefore, is that these galaxies uniformly sample high $z$ galaxies according to star-formation rate, i.e.\\ $\\fgrb \\propto {\\rm SFR}(L) \\phi(L)$, where $\\phi(L)$ is the luminosity function.  The link between QSO-DLA and star-forming galaxies is less direct, primarily because the bright background quasar precludes the easy detection of stellar light.  Nevertheless, the presence of heavy metals in all QSO-DLAs (and dust in the majority) indicates at least prior star-formation \\citep{pgw+03}. Furthermore, the observation of \\ion{C}{2}* absorption suggests heating of the ISM by far-UV photons from ongoing star-formation in at least half of the sample \\citep{wpg03,wolfe04}.  Finally, a handful of QSO-DLA have been detected in emission and exhibit properties similar to low luminosity LBGs \\citep{moller}. In contrast to the GRB-DLA, however, the QSO-DLA are selected according to their covering fraction on the sky, i.e.\\ the probability of detection is the convolution of the \\ion{H}{1} cross-section with the luminosity function: $\\fdla \\propto \\sigma_{HI}(L) \\phi(L)$.  While both populations of DLAs may be drawn from the full sample of star-forming galaxies, their distribution functions would only be the same if $\\sigma_{HI}(L) \\propto SFR(L)$ (see also Chen et al.\\ 2000).  In this paper, we will test these ideas by comparing the observed metallicity distributions of QSO-DLA and GRB-DLA with simple predictions based on empirical measurements of star-forming galaxies at $z=3$. Specifically, we combine the luminosity function of UV-selected galaxies (LBGs), a metallicity/luminosity relation, and a simple prescription for the radial distribution of \\ion{H}{1} gas with luminosity to predict the metallicity distributions of the QSO-DLA and GRB-DLA. We then compare these models with current observation to test this general picture. This analysis is similar in spirit to the studies of \\cite{fynbo99} who combined the LBG luminosity function with a Holmberg relation for $\\sigma_{HI}$ to predict the luminosities and impact parameters of QSO-DLA galaxies, \\cite{jakobsson} who compared the luminosity distribution function of GRB host galaxies with the LBG luminosity function, and Chen et al.\\ (2005) and Zwaan et al.\\ (2005) who reconciled the properties of local galaxies with the QSO-DLA cross-section and metal abundances.  The paper is structured in the following way: In Sect.~2 we describe our methodology, in Sect.~3 our results and in Sect.~4 a discussion of the results obtained and the uncertainties behind the analysis.  Throughout the paper we assume a cosmology with $\\Omega_\\Lambda = 0.70, \\Omega_m = 0.30, h=0.7$. ",
        "conclusions": "We find that with a simple model including the luminosity function for LBGs, a Holmberg relation for $\\sigma(L)$, an L-Z relation and a metallicity gradient with a slope $\\gamma$ dependent on luminosity it is possible to reconcile the metallicity distributions of QSO-DLA, GRB-DLA and LBGs.  In this model the faint end of the luminosity function plays a very important role. As seen in the lower panel in Fig.~\\ref{fig:illu2} more than 75\\% of star-formation selected galaxies are fainter than the flux limit for LBGs, R=25.5. For QSO-DLA galaxies the fraction is even higher.  Hence, in this model the GRB and DLA samples, in contrast with magnitude limited surveys, provide an almost complete census of $z\\approx3$ star-forming galaxies that are not heavily obscured."
    },
    "0801/0801.3759_arXiv.txt": {
        "abstract": "{} {We investigate the influence of ram-pressure stripping on the star formation and the mass distribution in simulated spiral galaxies. Special emphasis is put on the question where the newly formed stars are located. The stripping radius from the simulation is compared to analytical estimates.} {Disc galaxies are modelled in combined N-body/hydrodynamic simulations (GADGET-2) with prescriptions for cooling, star formation, stellar feedback, and galactic winds. These model galaxies move through a constant density and temperature gas, which has parameters comparable to the intra-cluster medium (ICM) in the outskirts of a galaxy cluster (T=3 keV $\\approx$ 3.6$\\times$10$^7$ K and $\\rho$=10$^{-28}$ g/cm$^3$). With this numerical setup we analyse the influence of ram-pressure stripping on the star formation rate of the model galaxy.} {We find that the star formation rate is significantly enhanced by the ram-pressure effect (up to a factor of 3). Stars form in the compressed central region of the galaxy as well as in the stripped gas behind the galaxy. Newly formed stars can be found up to hundred kpc behind the disc, forming structures with sizes of roughly 1 kpc in diameter and with masses of up to 10$^7$ M$_{\\odot}$. As they do not possess a dark matter halo due to their formation history, we name them 'stripped baryonic dwarf' galaxies. We also find that the analytical estimate for the stripping radius from a Gunn \\& Gott (1972) criterion is in good agreement with the numerical value from the simulation. Like in former investigations, edge-on systems lose less gas than face-on systems and the resulting spatial distribution of the gas and the newly formed stars is different.} {} ",
        "introduction": "Many issues of galaxy formation and evolution can be understood within the framework of the successful Cold Dark Matter (CDM) theory. But besides the hierarchical assembly of galaxies via mergers, other interaction phenomena are important in the cluster environment. Different types of such interactions are discussed: galactic winds, ram-pressure stripping, galaxy-galaxy interactions and outflows from Active Galactic Nuclei (AGN). Most of these processes can occur, with some principal differences, also in the field environment. Ram-pressure stripping, however, is uniquely linked to the presence of a surrounding medium that acts by its ram pressure. The space between the galaxies in a galaxy cluster is filled with a hot ($\\sim 10^8$ K), thin ($\\sim 10^3$ ions/m$^3$) plasma, the so called intra-cluster medium (ICM). Theoretically, the interaction between galaxies in the cluster and the ICM was already described by Gunn \\& Gott (1972). They argued that galaxies lose material if the force due to ram-pressure stripping exceeds the restoring gravitational force of the galaxy. As a consequence, a stripping radius in ram-pressure affected galaxies can be defined, outside which the inter-stellar medium (ISM) cannot be prevented from being stripped by the galactic gravitational potential. Another analytical estimate was proposed by Mori \\& Burkert (2000) where the thermal ISM pressure is compared with the external ram pressure. Later it was also found observationally that ram pressure acts on several galaxies in the Virgo cluster (e.g. Cayatte et al. 1990; Kenney et al. 2004; Vollmer et al. 2004) and in the Coma cluster (Bravo-Alfaro et al. 2000, 2001). Those galaxies, which are subject to ram-pressure stripping, lose parts of their inter-stellar medium to the ICM. As this gas has been enriched with heavy elements by the stars within the galaxy, the chemical abundance of the surrounding intra-cluster medium will be affected (Schindler et al. 2005; Domainko et al. 2006; Kapferer et al. 2007). These heavy elements can then be found in X-ray observations.  Besides this treatment in cluster-scale simulations, several groups have already studied numerically the effect of ram-pressure stripping on individual galaxies. The stripping radius with respect to acting ram pressure and galaxy properties was investigated (e.g. Abadi et al. 1999, Vollmer et al. 2001, Roediger \\& Hensler 2005) and a reasonable agreement between the analytically estimated stripping radius and the stripping radius from numerical simulations was found. Further galaxy-scale numerical simulations for ram-pressure stripping were presented by Quilis et al. (2000), Mori \\& Burkert (2000), Toniazzo \\& Schindler (2001), Schulz \\& Struck (2001), and Roediger \\& Br\\\"uggen (2006,2007). Very recently J{\\'a}chym et al. (2007) used an N-body/SPH code to investigate the influence of a time-varying ram pressure on spiral galaxies.  In contrast to these papers we focus on the effect of ram-pressure stripping on the star formation rate (SFR) of the galaxy, rather than on the mass loss. Nevertheless, for completeness we also analyse the mass loss and the resulting disc structure. The combined effects of ram-pressure stripping and tidal interactions on simulated cluster spirals is investigated in detail in Kapferer et al. (2007b).  This paper is organized as follows: In Sect. 2 we present the numerical method and the simulation setup used for this work. In Sect. 3 the results are presented which are compared to observations in Sect. 4. We end with a summary of the main conclusions in Sect. 5. ",
        "conclusions": "We have investigated the influence of ram-pressure stripping on the star formation and the resulting mass distribution in simulated spiral galaxies. We used a simplified model, where the ICM is distributed homogeneously and with a constant temperature, in order to isolate the effects caused by ram pressure.  \\begin{itemize} \\item We found that ram-pressure stripping enhances the star-formation rate by up to a factor of 3 over several hundred Myr. In total the mass of newly formed stars is about two times higher than in an isolated galaxy after 500 Myr of ram pressure acting.  \\item The new stars are mainly formed in the central parts of the disc but a significant fraction forms also in the wake of the galaxy. As they are collisionless they do not feel the ram pressure and remain bound to the galaxy.  \\item Sub-structures form in the wake of the galaxy from compressed spiral arms of the original galaxy. These irregular shaped structures have sizes of roughly 1 kpc in diameter and harbour newly formed stars. We term these structures 'stripped baryonic dwarf' galaxies, as they do not possess a dark matter halo.  \\item The stripping radius, which is visible in the model galaxy is in good agreement with an analytical estimate using a Gunn \\& Gott (1972) criterion.  \\item In a model galaxy, which moves edge-on through the ICM, the star formation rate is enhanced as well, however the distribution of the newly formed stars is different. The new stars are mainly formed in the central region of the disc.  \\item Observations of enhanced SF in merging galaxy clusters and of individual galaxies with tails of stripped gas support the results presented in this paper. As a galaxy moves through a real cluster environment, it encounters gradients in density, temperature, and the gravitational potential of the cluster and interactions with other galaxies. We will investigate this complex and degenerate interplay of the various processes in an upcoming paper.  \\end{itemize}"
    },
    "0801/0801.1089_arXiv.txt": {
        "abstract": "{Most multi-planetary systems are characterized by hot-Jupiters close to their central star, moving on eccentric orbits. From a dynamical point of view, compact multi-planetary systems form a specific class of the general {\\it N}-body problem (where $N\\ge3$). Moreover, extrasolar planets are found in prograde orbits about their host star, and often in mean motion resonances (MMR).} {In a first step, we study theoretically a new stabilizing mechanism suitable for compact two-planet systems. This mechanism involves counter-revolving orbits forming a retrograde MMR. In a second step, we investigate the feasibility of planetary systems hosting counter-revolving planets. Dynamical stability, observations, and formation processes of these systems are analyzed and discussed.} {To characterize the dynamical behavior of multi-dimensional planetary systems, we apply our technique of global dynamics analysis based on the MEGNO indicator (Mean Exponential Growth factor of Nearby Orbits) that provides the fine structure of the phase space. In a few cases of possible counter-revolving configurations, we carry out new fits to the observations using the Pikaia genetic algorithm. A statistical study of the stability in the neighborhood of different observed, planetary-systems is completed using a Monte-Carlo method.} {We analyse the observational data for the HD\\thinspace73526 planetary system and find that counter-revolving configurations may be consistent with the observational data. We highlight the fine and characteristic structure of retrograde MMRs. We demonstrate that retrograde resonances open a family of stabilizing mechanisms involving new apsidal precession behaviors.} {Considering two possible formation mechanisms (free-floating planet and Slingshot model), we conclude that counter-revolving configurations are feasible.} ",
        "introduction": "At present, 271 extrasolar planets have been detected around 233 stars (both solar and non-solar type)\\footnote{January, the $\\textrm{11}^\\textrm{th}$, 2008. http://exoplanet.eu/catalog.php}. Among them, there are 25 multiple-planet systems~: 17 two-planet systems (e.g. HD$\\thinspace$82943, 47 UMa, HD$\\thinspace$108874, HD$\\thinspace$128311), 6 three-planet systems (e.g. $\\upsilon$ And, HD$\\thinspace$69830, Gliese 876, Gliese 581), 1 four-planet system (HD$\\thinspace$160691) and more recently 1 five-planet system ($55$ CnC). Observations indicate that Mean Motion Resonances (MMR) frequently occur for planets of multiple-planet systems~: Gliese 876 (e.g. Rivera \\etal 2005), HD$\\thinspace$82943 (e.g. Ji \\etal 2003, Mayor \\etal 2004) and HD$\\thinspace$128311 (Vogt \\etal 2005) are in 2:1 MMR, HD$\\thinspace$202206 is in 5:1 MMR (Correia \\etal 2005), while 47 UMa is close to a 7:3 (Fischer \\etal 2002) or 8:3 commensurability (Fischer \\etal 2003).  This work is devoted to compact multi-planetary systems, characterized by (a) giant Jupiter-like planets found close to their central star, and (b) high eccentricities. These two peculiarities lead to strong gravitational interactions between the planets and may result in an unstable, dynamical behavior. However, we\\linebreak observe many such planetary systems suggesting that they are stable, and raising the question of why they are stable. From a dynamical point of view, compact multi-planetary systems form a specific class of the general N-body problem (with $N \\ge 3$)  whose analytical solutions are not necessarily known. A stability analysis of planetary systems, using numerical methods to explore multi-dimensional parameter space, typically leads to stability maps in which rare islands of stability can be identified amidst large chaotic zones. The underlying mechanisms for these stability zones must be identified.  In 2002, Kiseleva-Eggleton \\etal (2002) showed that the currently-published, orbital parameters place the planetary systems HD 12661, HD 38529, HD 37124, and HD$\\thinspace$160691 in very different situations from the point of view of dynamical distribution. Since this first study of the comparative stability of multi-planetary systems, many studies\\linebreak have been carried out in this direction. The role of the orbital mean motion resonances, in particular with a 2:1 ratio, has been intensively studied by several research groups (for example Hadjidemetriou 2002; Lee \\& Peale 2002, 2003; Bois \\etal 2003; Ji \\etal 2003; Ferraz-Mello \\etal 2005b; Psychoyos \\& Hadjidemetriou 2005; Beaug\\'e \\etal 2006). As a result, it has been discovered that an extrasolar planetary system, even with large planetary masses and eccentricities, can be stable if planetary orbits are close to stable, resonant, and periodic orbits. It has also been established (see e.g. Chiang \\& Murray 2002; Lee \\& Peale 2002; Libert \\& Henrard 2006) that orbits in a large number of compact multi-planet systems, are locked in Apsidal Synchronous Precessions (ASP hereafter), i.e. that the apsidal lines precess, on average, at the same rate\\footnote{For the study of 3-D, full 3-body problems, we introduced the terminology ASP, for expressing that the apsidal lines precess on average in a 3-D space at the same rate: see Bois 2005, Bois \\etal 2005. We note that in the planar case this phenomenon is also called ``apsidal corotation'' (ACR; Beaug\\'e \\etal 2003). In a number of papers, one may also find the incorrect expression ``apsidal secular resonances'' (ASR). ACR and ASP are in general not true secular resonances, as highlighted by Ferraz-Mello \\etal (2005a).}. A solution involving\\linebreak both MMRs and ASP describes well the stability of eccentric, compact multi-planetary systems, but may not however be unique. We note, for example, that other multi-planetary systems have been found to be mainly controlled by secular dynamics (cf. Michtchenko \\etal 2006; Libert \\& Henrard 2006; Ji \\etal 2007). In the present paper, we illustrate theoretically that other mechanisms can in addition provide the stability in multi-planetary systems.  In the case of the HD$\\thinspace$73526 system (2:1 MMR), Tinney \\etal (2006) found stability over 1 Myr. Based on the analytical\\linebreak classification of Hadjidemetriou (2002) established according to a hierarchy of masses and eccentricities, this system could instead be classified as unstable. Hadjidemetriou's classification may however be too general to disprove the stability found by Tinney \\etal (2006). Be that as it may, we use the same data as Tinney \\etal (see Table \\ref{tab_param}) and our numerical method is outlined in the following section. Exploring the stability of the HD$\\thinspace$73526 system in orbital parameter space, we find large chaotic regions. We find that the published data can even be described by a chaotic behavior. We note however that Tinney \\etal (2006) used a different definition of stability.\\footnote{In the paper of Tinney \\etal (2006), the claim of stability is obtained from the dynamical behaviors of the resonant angles related to the 2:1 MMR, rather than by characterizations of quasi-periodicity of the orbital solution. Besides, the notion of stability is only presumed to be acquired by the simple absence of planet ejection. We use instead the usual definition of stability related to quasi-periodicity (see Section \\ref{method}) and suitable for conservative dynamical systems.} Of course, we cannot exclude that the observational data were insufficient to allow a reliable orbital fit or the fit itself was not\\linebreak adequate. On the other hand, it is also possible that the underlying assumption of two prograde orbits is wrong. When placing\\linebreak one of the two planets on a retrograde orbit (which forms a system with counter-revolving planets), the stability region\\linebreak becomes very large. We will show below that this does not imply\\linebreak that the orbital fit is consistent with this stability zone. It implies\\linebreak that, in the neighborhood of the observational point, we can\\linebreak theoretically find stable solutions for counter-revolving configurations. To distinguish between two resonance cases when both planets are in prograde orbits, or when one planet is on a retro\\-grade orbit, we call them prograde and retrograde resonances, respectively.  Presently, all known extrasolar planets in multiple systems are believed to revolve in the same direction about their corresponding central star. Most fitted, orbital elements are derived by assuming prograde orbits. This is expected according to current\\linebreak theories for planetary formation in a circumstellar disk. In order to obtain a planet in retrograde resonance, an additional event is necessary such as violent, dynamical evolution of the planetary system, or a capture of the retrograde planet. In our Solar System, comets and the planetary satellites of Neptune, Saturn and Jupiter are known to have retrograde orbits. It is, therefore,\\linebreak important to investigate the stability of exoplanetary systems with a retrograde planet in particular if the observations do not yield a stable system when assuming all planets on prograde\\linebreak orbits.  In Section \\ref{method}, we present our method of global dynamics analysis. We show that there exists theoretically initial conditions in the vicinity of observational data such that stability is only possible for a counter-revolving configuration. It raises the question of whether such a configuration is consistent with the observational\\linebreak data of a given system (Section \\ref{obs}). In Section \\ref{stat}, we focus on the statistical occurence of stable solutions related to both prograde and retrograde resonances. This statistical approach is applied to three systems in 2:1 MMR and two systems in 5:1 MMR. If such systems harboring counter-revolving planets exist, we must also consider how they form: we discuss this issue in Section \\ref{form}. By analyzing the parameter space in the vicinity of the best-fit of the HD\\thinspace73526 planetary system, we highlight the fine structure of the 2:1 retrograde resonance (Sections \\ref{structure} and \\ref{evidence}), and the nature of associated apsidal precessions (Section \\ref{apsidal}). In addition, we complete an analoguous study for a theoretical system in 5:1 retrograde MMR (Section \\ref{HD16}). ",
        "conclusions": "We have found novel mechanisms giving rise to stability that could be suitable for a class of compact planetary systems. Such mechanisms involve counter-revolving orbits forming a retrograde MMR occuring in a quasi-identical plane. High statistical\\linebreak occurence of stable counter-revolving orbits is found. Our study of retrograde MMRs indicates the large stability domains and the specific behaviors of the precession and resonant angles. We propose that these large stability domains are caused by close approaches much faster and shorter for counter-revolving\\linebreak configurations than for the prograde ones. Scanning the HD\\thinspace73526 planetary system, we find evidence for a new type of apsidal precession (the rocking ASP). We find that the\\linebreak difference between the longitudes of periastron reveals a specific alternation mode at retrograde resonances. We emphasize that the counter-revolving configuration studied for the HD\\thinspace73526 planetary system is consistent with the observational data. Free-floating planets or the Slingshot model might explain the origin of such counter-revolving systems."
    },
    "0801/0801.4453_arXiv.txt": {
        "abstract": "We report on a correlation between virial mass $M$ and spin parameter $\\lambda$ for dark matter halos forming at redshifts $z \\gtrsim 10$. We find that the spin parameter decreases with increasing halo mass. Interestingly, our analysis indicates that halos forming at later times do not exhibit such a strong correlation, in agreement with the findings of previous studies. We briefly discuss the implications of this correlation for galaxy formation at high redshifts and the galaxy population we observe today. ",
        "introduction": "\\label{sec:introduction}  The physical mechanism by which galaxies acquire their angular momentum is an important problem that has been the subject of investigation for nearly sixty years (Hoyle 1949). This reflects the fundamental role played by angular momentum of galactic material in defining the size and shapes of galaxies (e.g. Fall \\& Efstathiou 1981). Yet despite its physical significance, a precise and accurate understanding of the origin of galactic angular momentum remains one of the missing pieces in the galaxy formation puzzle.  A fundamental assumption in current galaxy formation models is that galaxies form when gas cools and condenses within the potential wells of dark matter halos (White \\& Rees 1978). Consequently it is probable that the angular momentum of the galaxy will be linked to the angular momentum of its dark matter halo (e.g. Fall \\& Efstathiou 1980; Mo, Mao \\& White 1998; Zavala, Okamoto \\& Frenk 2007). Within the context of hierarchical structure formation models, the angular momentum growth of a dark matter proto-halo is driven by gravitational tidal torquing during the early stages (i.e. the linear regime) of its assembly. This ``Tidal Torque Theory'' has been explored in detail; it is a well-developed analytic theory (e.g. Peebles 1969, Doroshkevich 1979, White 1984) and its predictions are in good agreement with the results of cosmological $N$-body simulations (e.g. Barnes \\& Efstathiou 1987; Warren~\\ea 1992; Sugerman, Summers \\& Kamionkowski 2000; Porciani, Dekel \\& Hoffman 2002). However, once the proto-halo has passed through maximum expansion and the collapse has become non-linear, tidal torquing no longer provides an adequate description of the evolution of the angular momentum (White 1984), which tends to decrease with time. During this phase it is likely that merger and accretion events play an increasingly important role in determining both the magnitude and direction of the angular momentum of a galaxy (e.g. Bailin \\& Steinmetz 2005). Indeed, a number of studies have argued that mergers and accretion events are the primary determinants of the angular momenta of galaxies at the present day (Gardner 2001; Maller, Dekel~\\& Somerville 2002; Vitvitska~\\ea 2002).\\\\  It is common practice to quantify the angular momentum of a dark matter halo by the dimensionless ``classical'' spin parameter (Peebles 1969), \\begin{equation} \\label{eq:lambda} \\lambda = \\frac{J \\sqrt{|E|}}{GM^{5/2}}, \\end{equation} where $J$ is the magnitude of the angular momentum of material within the virial radius, $M$ is the virial mass, and $E$ is the total energy of the system. It has been shown that halos that have suffered a recent major merger will tend to have a higher spin parameter $\\lambda$ than the average (e.g. Hetznecker \\& Burkert 2006; Power, Knebe \\& Knollmann 2008). Therefore one could argue that within the framework of hierarchical structure formation that higher mass halos should have larger spin parameters \\emph{on average} than less massive systems because they have assembled a larger fraction of their mass (by merging) more recently.  However, if we consider only halos in virial equilibrium, should we expect to see a correlation between halo mass and spin? One might na\\\"ively expect that more massive systems will have had their maximum expansion more recently and so these systems will have been tidally torqued for longer than systems that had their maximum expansion at earlier times. This suggests that spin should \\emph{increase} with timing of maximum expansion and therefore halo mass. However, one finds at best a weak correlation between mass and spin for equilibrium halos at $z$=0 (e.g. Cole~\\& Lacey 1996; Maccio et al. 2007; Bett et al. 2007, hereafter B07), and the correlation is for spin to \\emph{decrease} with increasing halo mass, contrary to our na\\\"ive expectation.  In this paper, we report on a (weak) correlation between spin and mass for equilibrium halos at redshift $z$=10. The trend is for higher-mass halos to have smaller spins, and is qualitatively similar to the one reported by B07 for the halo population at $z$=0. We present the main evidence in support of this correlation in Section~\\ref{sec:correlation} and we consider its implications for galaxy formation in Section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} We have performed a careful investigation of the relation between virial mass and dimensionless spin parameter for dark matter halos forming at high redshifts $z \\gtrsim 10$ in a $\\Lambda$CDM cosmology. The result of our study, which is based on a series of cosmological $N$-body simulations in which box size was varied while keeping particle number fixed, indicates that there is a \\emph{weak} correlation between mass and spin at $z$=10, such that the spin decreases with increasing mass. If there is a correlation at $z$=1, we argue that it is significantly weaker than the one we find at $z$=10; this is in qualitative agreement with the findings of previous studies that focused on lower redshifts (Maccio~\\ea 2007, Shaw~\\ea 2005, Lemson~\\& Kauffmann 1999).\\\\  Interestingly, B07 find a weak correlation between median spin and halo mass at $z$=0 in the Millennium Simulation (Springel et al. 2005), in the same spirit as the one presented here for $z$=10: lower mass halos tend to have higher spins. However, as we show, the correlation between halo mass and spin is weaker at $z$=1 than at $z$=10, whereas the correlation reported in B07 for halos at $z$=0 is much stronger than the one we find at $z$=1. This is not what one would expect, and so it is important to try and understand the source of the difference between our result at $z$=1 and the B07 result at $z$=0. B07 fitted a 3$^{\\rm rd}$-order polynomial to the median spins of halos in the mass range $3\\times 10^{11} \\lesssim M/(h^{-1} M_{\\odot}) \\lesssim 3\\times 10^{14}$ at $z$=0. The form of this polynomial is extremely sensitive to the precise values of the best-fit parameters (Bett, private communication) and it is not straightforward to extrapolate its behaviour outside of the given mass range and redshift. We derive our estimates of the power-law exponents from the spin distribution with respect to halo mass at $z$=1. Our halos lie in the mass range $3\\times 10^{9} \\lesssim M/(h^{-1} M_{\\odot}) \\lesssim 5\\times 10^{12}$. B07 base their median spins upon $\\sim 1.5$ million halos with correspondingly small errors, and note that the weak nature of the trend of spin with mass makes it hard to detect. This suggests to us that the correlation between mass and spin at $z$=10 is remarkably strong rather than the correlation at $z$=1 being too weak!  When studying correlations between halo mass and spin, great care must be taken in defining the halo sample. In particular, we find that mass resolution (i.e. the number of particles with which a halo is resolved) and the degree of virialisation of a halo can have a significant effect on the strength of the correlation (at least at $z$=1, cf. \\Table{tab:Nmin}). This -- at least -- is in good agreement with the findings of B07.\\\\  We note that Power \\& Knebe (2006) demonstrated that the size of simulation box can lead to a suppression of angular momentum in smaller boxes, due to the absence of longer wavelength perturbations in the initial conditions. This will lead to a bias in our estimate of $\\lambda$ (approximately a $\\sim 10\\%$ effect) but we have verified that the spin distributions we obtain from a B20 run truncated on scales larger than the longest wavelength perturbation modelled in the B1 run produces results that are consistent. Indeed, we would expect the correlation to be strengthened if the B1 spins were corrected for box size effects.\\\\  It is interesting to speculate on the consequences of this correlation for galaxy formation at high redshifts and the galaxy population we observe today. In the standard picture of galaxy formation, gas cools on to dark matter halos and is shock heated to the virial temperature of the halo. The angular momentum of the gas and the dark matter should (initially) be similar because they are subject to the same tidal field. As the innermost densest parts of the gaseous halo cool, they will settle into a gaseous disk with a scale length determined by the specific angular momentum of the gas, which we would expect to be related to the angular momentum of the halo (e.g. Zavala, Okamoto \\& Frenk 2007).  If more massive halos at high redshifts show a tendency to have smaller spin parameters, the gas disks will have lower specific angular momenta and therefore will be more centrally concentrated. If star formation rate correlates with surface density, then we might expect the star formation rate to be enhanced in more massive halos. Because massive halos tend to form preferentially in high density, highly clustered environments in which the merger rate also tends to be enhanced, then we might expect star formation to proceed more rapidly and at earlier times in these halos. Might this explain the effect of ``downsizing'' (e.g. Cowie et al. 1996), the successive shifting of star formation from high- to low-mass galaxies with decreasing redshift? We shall pursue this in a more quantitative manner in future work."
    },
    "0801/0801.2569_arXiv.txt": {
        "abstract": "Recent Cherenkov observations of BL Lac objects showed that the TeV flux of PKS 2155--304 changed by a factor 2 in just 3--5 minutes. This fast variability can be accounted for if the emitting region is moving with a bulk Lorentz factor $\\Gamma\\sim 50$ and a similar relativistic Doppler factor. If this $\\Gamma$ is adopted, several models can fit the data, but, irrespective of the chosen model, the jet is matter dominated. The Doppler factor requires viewing angles of the order of $1^\\circ$ or less: if the entire jet is as narrow as this, then we have problems with current unification schemes. This suggests that there are small active regions, inside a larger jet, moving faster than the rest of the plasma, occasionally pointing at us. Coordinated X--ray/TeV variability can discriminate between the different scenarios. ",
        "introduction": "The blazars PKS 2155--304 (Aharonian et al. 2007) and Mkn 501 (Albert et al. 2007) showed variations on $t_{\\rm var}=$3--5 minutes of their TeV flux. This ultrafast variability calls for a revision of our current ideas about the properties of the emitting region. Particularly challenging is the case of PKS 2155--304, because the variations occurred during an overall very active state of the source with an observed TeV luminosity of $\\sim 10^{47}$ erg s$^{-1}$. The usual way to infer the size from the variability timescale is $R<ct_{\\rm var}\\delta$ where $\\delta$ is the Doppler factor. Begelman, Fabian \\& Rees (2008; hereafter B08) pointed out that such small $t_{\\rm var}$ cannot be any longer indicative of the size of the black hole (as instead are in the ``internal shock scenario\" (Sikora et al 1994, Ghisellini 1999, Spada et al. 2001, Guetta et al. 2004). The mass of the central black hole of PKS 2155--304 is uncertain (of the order of $10^9$ $M_\\odot$; Aharonian et al. 2007) but even a (prudently small) mass of $M=10^8$ solar masses corresponds to a light travel time across the Schwarzschild radius of $10^3$ s. To avoid to have a too compact source, with the accompanying problem of $\\gamma$--ray absorption through the $\\gamma$--$\\gamma$ $\\to e^\\pm$ process (see e.g. Dondi \\& Ghisellini 1995), the bulk Lorentz factor of the emitting region (hence $\\delta$) must have a value close to 50. This requirement is not unprecedented, since also the large separation, in frequency, of the two broad peaks characterising the spectral energy distribution (SED) requires $\\delta>30$ in single--zone synchrotron self Compton (SSC) models (see e.g. Konopelko et al. 2003). Smaller $\\delta$ and $\\Gamma$ factors are however possible in alternative models, where two or more emission regions have different $\\Gamma$, as in the decelerating jet model of Georganopoulos \\& Kazanas (2003) and in the ``spine--layer\" model proposed by Ghisellini, Tavecchio \\& Chiaberge (2005; hereafter G05). In both models contiguous emitting regions of the jet are moving with different $\\Gamma$, and one part can see the emission of other relativistically boosted, thus enhancing the radiation energy density and the consequent inverse Compton emission. In this way one can relax the requirement of having $\\Gamma$--factors of the order of 30 or more, requiring a more standard value of around 15. The observed rapid variability, instead, cannot be ``cured\" by invoking a structured jet emission region, and is therefore more demanding.  The implied large $\\Gamma$ makes any external photon source strongly boosted in the comoving frame, favouring the ``external Compton\" (EC) process, in which the energy density of the seed photons produced outside the jet is larger than the energy density of the synchrotron radiation produced locally. This argument leads B08 to suggest that the EC process is the main radiation mechanism. They also argue that, since to produce TeV photons one needs highly energetic electrons (with random Lorentz factors $\\gamma \\sim 10^6/\\Gamma$), the jet could be particle starved (i.e. one needs fewer electrons to produce the radiation we see, if they are at high energies).  Therefore the jet should be magnetically dominated, i.e. the power carried by the jet should be mainly in the form of the Poynting flux associated to magnetic fields.   In this paper we investigate if the two claims: 1) EC favoured with respect to the SSC model, and 2) the jet is magnetically dominated (made by B08 on the basis of rather general arguments), are confirmed by detailed modelling of the SED of PKS 2155--304 with the SSC and the EC models. In particular, we consider two constraints not considered in B08: i) there is a limit on the amount of the seed external luminosity, posed by the observed flux, and ii) the predicted synchrotron flux must reproduce, or be smaller, than the observed optical/X--ray flux. We also present a new scenario, analogous to the spine/layer model, in which a compact and active region is moving within a larger jet, with a bulk Lorentz factor larger then the one of the surrounding plasma. We call this version of the spine/layer model the ``needle/jet'' model. ",
        "conclusions": "We have modelled the SED of PKS 2155--304 during the huge TeV flare of July 28, 2006 with three models, assuming always a bulk Lorentz factor $\\Gamma=50$, a size of the TeV emitting region $R\\sim 3\\times 10^{14}$ cm, and a viewing angle of $1^\\circ$, to account for the observed short variability timescale of the TeV flux.  As mentioned in the Introduction, our aim was to check whether (as concluded by B08) in this case 1) the large $\\Gamma$ favours the dominance of seed photons produced externally to the jet to produce the TeV flux, and 2) the jet (or the portion of the jet) producing the variable TeV emission could be Poynting flux dominated.  The answer to 2) is negative since the Poynting flux carried by the jet is never dominating, even assuming the most favourable conditions. It is true that, in the EC case, the larger the external radiation energy density the larger the allowed magnetic field (and thus $L_{\\rm B}$), but the limit posed by the SED itself to the amount of the luminosity produced externally to the TeV emitting region is also limiting the maximum value of the magnetic field.  The answer to 1) (is the external radiation important?) is more complex. In our fits the SSC model reproduces the data better than the EC model. However, this is due to the fact that we wanted to fit the high energy data and, {\\it at the same time}, maximise the magnetic field. Allowing a {\\it smaller} magnetic field we could obtain a better agreement with the SED. In fact the best fitting model is the needle/jet one, that can be thought as a kind of EC model, where the ``external\" radiation is the radiation produced by the jet itself, but externally to the needle.  In the EC case, the power carried by the jet in emitting electrons is smaller than what carried in radiation.  This apparent paradox (how can the radiation have more power than the electrons producing it?) is solved recalling that in this EC model the cooling time of the electron becomes short, with the need of refreshing the electron distribution with the injection of new particles.  Therefore the flow of relativistic electrons across a given jet cross sectional area underestimates the number (and the energy) of the electrons having contributed to the radiation flux. The derived $L_{\\rm e}$ is therefore a lower limit.  Concerning the ``blazar sequence\", in the form of the relation between $\\gamma_{\\rm peak}$ and the energy density $U^\\prime$ (magnetic plus radiative) seen in the comoving frame, it is well obeyed in the SSC and needle/jet case, not so for the EC model.  Intriguingly, the relation between $\\gamma_{\\rm peak}$ and the total jet power is instead followed by all the three models.   The SSC and EC models here discussed cannot describe the ``quiescent\" or ``persistent\" state of blazars,  (but in any case variable on timescale of the order of $10^4$ s) since the combination of jet narrowness and the required small viewing angle cannot be reconciled with current unification schemes. One possibility is that PKS 2155--304 is unique, and the fast TeV variability will not be a general property of TeV blazars. This is unlikely, for two reasons: the first is that already another TeV blazars showed a very small $t_{\\rm var}$ (i.e. Mkn 501), the other reason is that such a fast variability was observed as soon as the Cherenkov telescopes had the required sensitivity to detect it.  It may be dangerous to infer, from the ultrafast TeV variability, the general properties of a jet that can be discontinuous, either in space or in time, or both. In other words: episodes of very fast variability can be produced by ``needles\" of matter moving faster than average, oriented in different directions but contained in a wider cone, and occasionally pointing at the observer. Or the emitting regions are shells that corresponds to an intermittent activity of the central engine, having on average a bulk Lorentz factor that is quite large but not extreme, and occasionally going much faster (and possibly contained in a narrower cone). In the first case a ``needle\" would move through a region already filled with radiation (produced by the rest of the jet), and this would resemble the needle/jet model here discussed, in the second case it is likely that the emitting region would be rather external--photon--starved, and the SSC model could better describe the corresponding SED. How to distinguish the two cases? A crucial point is the UV--X--ray flux, if observed exactly simultaneously with the TeV band. If there is no sign of simultaneous fast variability, then the ``needle\" model is to be preferred, since the UV--X--rays could be produced by the entire jet, while the TeV flux can come from the needle. If instead both the UV--X--ray and the TeV fluxes vary together and quickly, then the SSC model is favoured."
    },
    "0801/0801.0793_arXiv.txt": {
        "abstract": "Asteroid 3200 Phaethon resembles a comet in some ways, including a highly-eccentric orbit ($e\\sim0.89$) and a strong associated meteor shower (the Geminids). Yet this object has never been observed to exhibit any cometary activity, i.e., gas production. We observed 3200 Phaethon with the Caltech Submillimeter Observatory on two occasions, once while it was near its closest approach to Earth as it neared perihelion, and another while it was further from Earth post-perihelion. Observations of the $J=2\\rightarrow1$ and $J=3\\rightarrow2$ rotational transitions of $^{12}$CO, typically strong lines in comets and indicative of gas production, yielded no detection. Upper limits on the $^{12}$CO production of $1.8\\times10^{28}$ molecules~s$^{-1}$ and $7.6\\times10^{28}$ molecules~s$^{-1}$ for Phaethon were determined on these two occasions. ",
        "introduction": "Asteroid 3200 Phaethon was discovered by the Infrared Astronomical Satellite in 1983 \\citep{grekow83}. Almost immediately thereafter its orbit was noted to be very similar to that of the Geminid meteor shower \\citep{whi83}. This makes it the only asteroid linked to a strong meteor shower.  A possible exception is asteroid 2003~EH1, which is associated with the Quadrantid shower \\citep{jen04,wilryabat04,wiebro05}) but which has not yet been observed near perihelion and thus whether it is asteroidal or cometary in nature remains unclear.  The link between Phaethon and the Geminid meteor shower is also based on the closeness of their orbits, thus one could argue that they are simply aligned by coincidence. However, based on the very small difference between their orbits (see Table~1) and the absence of other objects on orbits nearby, \\citet{wiebro04a} determined that probability of the orbital similarity of Phaethon and the Geminids being due to a chance alignment is less than one in a thousand. Though Phaethon's current activity (if any) is certainly low, intermittent or both, a study of the orbits of observed Geminids has indicated that they are consistent with cometary release no more than 2000 years ago and possibly within the last 600 years \\citep{gus89}. Thus ongoing, perhaps sporadic, gas production by this object remains a possibility.  If 3200~Phaethon is not a comet, this does not mean it cannot be the parent body of the Geminids. A collision or impact might have released the material which formed the Geminid stream. A new rare population of {}``main-belt comets'' has been uncovered in the last several years \\citep{hsijew06}, bodies whose gas and dust production may be partly the result of collisions. However, the sublimation of volatiles from the body's surface also has a role to play, as evidenced by the seasonal cycle of activity that has been seen in one of them (133P/Elst-Pizarro) to date \\cite[]{hsijewlow07}.  We note that Geminid meteors are much more durable than typical cometary meteors: studies of their ablation in Earth's atmosphere have shown that they have mean densities of 2.9~g~cm$^{-3}$ \\citep{bab02}, the highest of any of the streams measured. This may imply that the Geminids are made of relatively strong rocky asteroidal material rather than more porous and fragile cometary material. However, the question is complicated by the small (0.14~AU \\citep{coo73}) perihelion distance of the Geminid stream. Smaller than is typical of comets, this low perihelion distance could result in extreme baking and sintering of the meteoroid particles, producing more durable meteoroids than might otherwise be expected of a cometary meteoroid stream.  Phaethon's association with a strong meteor shower and its highly eccentric orbit ($e\\approx0.89$) make it a candidate cometary object. Yet in the over two decades since its discovery (which include many perihelion passages of its 1.43-year orbit), no cometary activity has ever been observed \\citep{cocbar84,chamcfsch96,hsijew05}.  Phaethon's spectral type is B \\cite[]{gremeadav85}, a primitive type associated with the outer portion of the main asteroid belt, though its albedo (0.11-0.17) is somewhat higher than would be expected of such a type \\cite[]{veekowmat84,gremeadav85}, as was pointed out by \\cite{weibotlev02}.  Visible and near IR spectroscopy of Phaethon support its identification as an asteroid rather than a comet nucleus \\cite[]{liccammot07}.  Phaethon has been also linked with the smaller asteroid 2005~UD, with which it shares a bluish colour, an absence of dust production \\cite[]{jewhsi06} and an orbital similarity \\cite[]{ohtsekkin06}.  Is Phaethon then an asteroid, and if so, how did it produce its meteor stream?  The question of the nature of this unusual object can only be answered by ongoing observations. In particular, observations which might reveal gas production by this body are the most important to undertake, as a positive detection would unequivocally identify this object as a comet, though its muted activity might imply it was extinct or dormant. Given the lack of previous detection of gas from Phaethon, its gas production rate is undoubtedly low, and observations should ideally be conducted near perihelion, or as close to it as can be managed given the interfering presence of the Sun.  In this paper we describe our attempts to measure gas production on Phaethon through the detection of $J=2\\rightarrow1$ and $J=3\\rightarrow2$ rotational transitions of the $^{12}$CO molecule at the Caltech Submillimeter Observatory (CSO), located atop Mauna Kea, Hawaii.  Observations in the mm/sub-mm range suffer the constraints of bad weather even more than traditional optical astronomy. This project arose from an idea to use time on the CSO that was unsuitable for projects requiring excellent atmospheric conditions for other purposes.  The choice of Phaethon as a target, despite its lack of observed activity, was motivated by its meteor stream and the intermittency of activity seen in other comets eg. in the main-belt: ongoing monitoring of this object seems justifiable.  The $J=2\\rightarrow1$ and $J=3\\rightarrow2$ rotational transitions of carbon monoxide (CO) are strong in bright comets, though not perhaps ideal indicators of gas production for Phaethon. They were chosen as being among our best bets given our observational window during mediocre conditions ($\\tau_{225~GHz} \\sim 0.15$). In fact our upper limit on gas production using these lines is comparable to that of a moderate-activity comet, and by observing as close to perihelion as possible, we hoped to catch Phaethon in outburst. Though our attempt was unsuccessful, only by ongoing observations will Phaethon's nature be clarified.  We present the three spectra we obtained from corresponding sets of observations and discuss our results in the next sections. ",
        "conclusions": "The submillimeter rotational transitions of $^{12}$CO have been studied in detail for comets and the strength of the lines can be linked to gas production rates \\citep{croleb83}.  Assuming a lifetime of $^{12}$CO against dissociation by solar UV radiation at 1~AU of $\\tau=1.5\\times10^{6}$~s (\\citet{huecar79} from \\cite{croleb83}) and assuming a gas expansion velocity $v_{{\\rm {\\rm }{exp}}}=800$ m/s (the pre-perihelion value determined for Hyakutake from submillimeter observations by \\citet{bivboccro99}), any $^{12}$CO cloud around 3200 Phaethon should have a diameter $\\sim2.4\\times10^{6}$~km, with an angular size of about 0.9 degree at 1~AU. A comparison of this angular size with the CSO main beam sizes quoted earlier for our observations justifies the efficiency adopted for our observations (i.e., 70\\%).  Assuming fluorescence equilibrium and a Haser distribution \\citep{has57,bercosque98} with $\\gamma_{{\\rm {\\rm }{CO}}}=v_{{\\rm {\\rm }{exp}}}\\tau=1.2\\times10^{6}$~km, and using an emission constant $T_{B}\\Delta v/\\langle N_{{\\rm {\\rm }{CO}}}\\rangle$ of $3.65\\times10^{-16}$~K~km~s$^{-1}$~(molec. cm$^{-2}$)$^{-1}$ and $8.08\\times10^{-16}$~K~km~s$^{-1}$~(molec. cm$^{-2}$)$^{-1}$ for the $J=2\\rightarrow1$ and $J=3\\rightarrow2$ transitions, respectively \\citep{croleb83}, we can derive a corresponding upper limit on $Q_{{\\rm {\\rm {\\rm }{CO}}}}$, the rate of $^{12}$CO production, of approximately $1.8\\times10^{28}$~\\ms for 5 November 2004 (from the $J=3\\rightarrow2$ transition). More precisely, at 1.05~AU (the distance between Earth and 3200 Phaethon during the corresponding observations) the dependence on the measured and assumed parameters can be expressed as  \\begin{equation} Q_{{\\rm {\\rm }{CO}}}\\approx3.60\\times10^{28}\\left(\\frac{v_{{\\rm {\\rm }{exp}}}}{800~\\textrm{{m/s}}}\\right)^{2}\\left(\\frac{T_{B}}{42~\\textrm{{mK}}}\\right)\\textrm{molecules}\\cdot\\textrm{s}^{-1}\\label{eq:Q2-1}\\end{equation}   \\noindent for the $J=2\\rightarrow1$ transition, and  \\begin{equation} Q_{{\\rm {\\rm }{CO}}}\\approx1.77\\times10^{28}\\left(\\frac{v_{{\\rm {\\rm }{exp}}}}{800~\\textrm{{m/s}}}\\right)^{2}\\left(\\frac{T_{B}}{69~\\textrm{{mK}}}\\right)\\textrm{molecules}\\cdot\\textrm{s}^{-1}\\label{eq:Q3-2}\\end{equation}   \\noindent for the $J=3\\rightarrow2$ transition. Using an equation similar to equation (\\ref{eq:Q2-1}) for a distance of 1.904~AU we also establish an upper limit of $Q_{{\\rm {\\rm }{CO}}}\\approx7.6\\times10^{28}$~\\ms for 13 September 2005.  These upper limits are lower than production rates observed for very bright comets at the same heliocentric distance, e.g., C/1995 O1 (Hale-Bopp), $Q_{{\\rm {\\rm }{CO}}}=2.07\\times10^{30}$~\\ms \\citep{dismumdel01}. The limits are comparable to the production rate of C/1996 B2 Hyakutake, $Q_{{\\rm {\\rm }{CO}}}\\approx5\\times10^{28}$~\\ms near 1~AU \\citep{bivboccro99}, though it is a long-period comet and expected to be rich in ices.  Comet 2P/Encke, an old low-activity comet, might provide a better benchmark. Though its carbon monoxide production has not been measured to our knowledge, we can estimate this from its H$_{2}$O production rate, which has measured at 0.8 to $2.6\\times10^{28}$ molecules~s$^{-1}$ at similar heliocentric distances \\citep{maksilsch01}. The $^{12}$CO to H$_{2}$O ratio in comets is variable, ranging from 0.5\\% to 22\\% \\citep{bivboccro99, dismumdel01,mumdeldis01b, mumdeldis01, mummcldis01, disdelmag02,felweabur02,dismumdel03,disandvil07}. Taking the ratio to be 5\\%, we can derive a corresponding expected $^{12}$CO production rate of $0.4\\times10^{27}$ to $1.3\\times10^{27}$~molecules~s$^{-1}$ for comet Encke. Though a nearly-extinct comet could produce $^{12}$CO at even lower rates, this presents us with a target measurement sensitivity of $Q_{\\mathrm{CO}}\\sim10^{27}$~molecules~s$^{-1}$.   In view of these numbers it will be necessary to make deeper observations of 3200 Phaethon in order to detect any $^{12}$CO production or to reduce significantly our derived upper limits.  It follows that we need to increase our sensitivity by at least an order of magnitude. The results discussed here were obtained under relatively poor sky conditions and improving these limits should be feasible with a combination of longer exposure times and better observing conditions (e.g., $\\tau_{225\\,{\\rm {GHz}}}\\simeq0.05$), as is relatively common (20\\% to 25\\% of the time) on Mauna Kea. The sensitivity can also be increased by observing Phaethon when it is closest to Earth, as the detection limit drops approximately linearly with the Earth-comet distance. Observations of this object during its upcoming close approach to Earth ($\\las0.13$~AU in December 2007) would thus be particularly recommended."
    },
    "0801/0801.2275_arXiv.txt": {
        "abstract": "We present the first mid-infrared Period-Luminosity (PL) relations for Large Magellanic Cloud (LMC) Cepheids. Single-epoch observations of 70 Cepheids were extracted from Spitzer IRAC observations at 3.6, 4.5, 5.8 and 8.0~$\\mu$m, serendipitously obtained during the SAGE (Surveying the Agents of a Galaxy's Evolution) imaging survey of the LMC. All four mid-infrared PL relations have nearly identical slopes over the period range 6 - 88 days, with a small scatter of only $\\pm$0.16~mag independent of period for all four of these wavelengths. We emphasize that differential reddening is not contributing significantly to the observed scatter, given the nearly two orders of magnitude reduced sensitivity of the mid-IR to extinction compared to the optical.  Future observations, filling in the light curves for these Cepheids, should noticably reduce the residual scatter.  These attributes alone suggest that mid-infrared PL relations will provide a practical means of significantly improving the accuracy of Cepheid distances to nearby galaxies. ",
        "introduction": "Distance measurements to the Large Magellanic Cloud (LMC) have historically played a critical role in the calibration of the extragalactic distance scale (Feast \\& Walker 1987; Madore \\& Freedman 1991; Udalski {\\it et al.} 1999; Freedman {\\it et al.} 2001; Sandage {\\it et al.} 2006). The LMC Cepheid Period-Luminosity relations have been adopted as the fiducial calibration sample for many recent extragalactic distance measurements and for the determination of the Hubble constant ({\\it e.g.,} Freedman {\\it et al.} 2001, Sandage {\\it et al.} 2006, Riess {\\it et al.} 2005).  A number of different methods have been used to measure the distance to the LMC. As tabulated in Gibson (2000) and Freedman {\\it et al.} (2001), most of the values for the LMC distance modulus based on these different methods fall between 18.1 to 18.7 mag, corresponding to a full range of 42 to 55 kpc, and likely reflecting the effects of systematic errors in various methods. More recent values have tended to cluster around a modulus of 18.5 mag (see Alves 2004; and also an interesting commentary by Schaefer 2007) \\footnote[1]{For a regularly updated compilation of published distances to the Large Magellanic Cloud see the web site maintained by Ian Steer and Barry Madore through NED at http://nedwww.ipac.caltech.edu/level5/NED0D/LMC\\_ref.html.}.  In an effort to reduce the systematic errors in the LMC distance, Persson {\\it et al.} (2004) obtained near-infrared J, H, and K$_s$ photometric measurements of 92 LMC Cepheids. These stars were chosen to be distributed across the LMC, with periods ranging from 3 to 100 days. They were also selected to be relatively isolated so that crowding effects would be minimized. The sample also does not contain overtone pulsators.  On average, 22 phase points were obtained at each wavelength for each star. The dispersions in the infrared PL, PLC, and extinction-free period-Wesenheit relations were found to be very small, amounting to less than $\\pm$0.14~mag, or 7\\% in distance.  There are a number of advantages to obtaining observations of Cepheids at long wavelengths (see, for example, the reviews by Madore \\& Freedman (1991, 1998) and references therein): (1) The sensitivity of surface brightness to temperature is a steeply declining function of wavelength. (2) The interstellar extinction curve decreases as a function of increasing wavelength (being almost linear with 1/$\\lambda$ at optical and near-infrared wavelengths). (3) At the temperatures typical of Cepheids, metallicity effects predominate in the UV, blue and visual parts of the spectrum where most of the line transitions occur, with declining effects at longer wavelengths (Bono {\\it et al.} 1999). The overall insensitivity of infrared magnitudes of Cepheids to each of these effects results in decreased amplitudes for individual Cepheids, as well as a decreased scatter in the apparent PL relations (first noted by Wisniewski \\& Johnson 1968, and by McGonegal {\\it et al.} 1982, respectively).  Given the pivotal role of the LMC distance in calibrating the extragalactic distance scale, developing techniques to further reduce systematic errors in that path remains a critical goal. The mid-infrared capability of Spitzer offers a new opportunity to provide a distance modulus to the LMC completely freed from the effects of reddening and with decreased sensitivity to metallicity.  As we discuss below, beyond about 4.5~$\\mu$m, the resulting extinction is at least a factor of 25 times smaller than the corresponding extinction in the B band.  We report here on single-epoch, mid-infrared data for a subset of the Persson {\\it et al.} (2004) LMC sample, which were observed with Spitzer at 3.6, 4.5, 5.8 and 8.0~$\\mu$m. Data for 70 of the 92 Persson et al. Cepheids were serendipitously obtained in the course of the SAGE project (Meixner {\\it et al.} 2006a), a large-area study of the interstellar medium and current star formation in the LMC. Since Cepheids are relatively young stars, large numbers of Cepheids were also observed in these fields, and fortunately, as we show here, they were observed with sufficient signal-to-noise to be suitable for our purposes also. The single-epoch, mid-IR Cepheid PL relations resulting from these observations have a total scatter of only $\\pm$0.16~mag (8\\% in distance) for a single Cepheid with less than 1\\% (statistical) uncertainty in the mean for this sample. We discuss the prospects for reducing the scatter even further. ",
        "conclusions": ""
    },
    "0801/0801.1510_arXiv.txt": {
        "abstract": "We demonstrate that gravitational lensing can be used to discover and study planets in the habitable zones of nearby dwarf stars.  If appropriate software is developed, a new generation of monitoring programs will automatically conduct a census of nearby planets in the habitable zones of dwarf stars. In addition, individual nearby dwarf stars can produce lensing events at predictable times; careful monitoring of these events can discover any planets located in the zone of habitability.  Because lensing can discover planets (1) in face-on orbits, and (2) in orbit around the dimmest stars, lensing techniques will provide complementary information to that gleaned through Doppler and/or transit investigations.  The ultimate result will be a comprehensive understanding of the variety of systems with conditions similar to those that gave rise to life on Earth. ",
        "introduction": "Microlensing has been considered as a method best suited to discover distant planets in, e.g., in the Magellanic Clouds, M31, the Galactic Bulge, Halo, and several kiloparsecs from us in the Galatic disk. Recently, it has been pointed out that lensing can provide a way to discover and study nearby (within $\\approx$ 1 kpc) planetary systems, providing complementary information to that obtained from Doppler and/or transit studies \\citep{DiStefano2007}.  Nearby planets are intriguing because detailed observations are possible.  If life exists on other planets, we will find it first and learn most about it on planets within a few hundred parsecs.  The importance of water to Earth life has suggested that we define a zone of habitability around each star to be the region in which liquid water can exist on the surface of a planet that has an atmosphere. The literature suggests that we can express $a,$ the separation between the central star of mass $M$, and a planet in the habitable zone as  \\begin{equation} a = 0.13\\, {\\rm AU}\\, h\\, \\Big({{M}\\over{0.25\\, {\\rm M_\\odot}}}\\Big)^{{3}\\over{2}}, \\end{equation}  \\noindent where $h$ is approximately equal to unity in the middle of the habitable zone, and to $2/3, 4/3$ at its inner and outer edge, respectively \\citep{Charbonneau2007, Tarter2007, Scalo2007}. There is presently significant uncertainty in this equation. The nearby dwarf Gliese 581, with a stellar mass of $0.31\\, {\\rm M_\\odot},$ has three known planets, one of which has an $h$ value of 1.39 and is generally considered to be possibly habitable. Another of its planets has an $h$ value of 0.41, and is generally considered not to be habitable \\citep{vonBloh2007, Selsis2007}, though mitigating effects have been proposed \\citep{Chylek2007}.  The idea motivating our work is to explore the action as a gravitational lens of a nearby dwarf star that hosts a planet in its zone of habitability. The deflection of light from a distant star by an intervening mass causes the image of the source to be split, distorted, and magnified.  When the lens is located in our Galaxy, and has mass comparable to that of a star or planet, the images of the source are separated by too small an angle to resolve.  We can, however, detect the increase in the amount of light received from the source when its position, $u,$ projected onto the lens plane is comparable to the Einstein radius, $R_E$.  \\begin{equation} R_E = 0.13\\, {\\rm AU}\\, \\Bigg[\\Big({{M}\\over{0.25\\, {\\rm M_\\odot}}}\\Big)\\, \\Big({{D_L}\\over{10\\, {\\rm pc}}}\\Big) \\Big(1-{{D_L}\\over{D_S}}\\Big)\\Bigg]^{{1}\\over{2}}, \\end{equation}  \\noindent where $D_L$ is the distance to the lens and $D_S$ is the distance to the source star. We have assumed that the mass of the planet is a small fraction of the mass of the star. Expressing $u$ in units of $R_E,$ the magnification is $34\\%, 6\\%, 2\\%, 1\\%$ when $u = 1, 2, 3, 3.5,$ respectively.  \\bigskip  Comparing equations (1) and (2) reveals that the spatial scale that defines the habitable zone is compatible with the size of the Einstein ring for a wide range of stellar masses and distances, $D_L.$ This is fortuitous, because the chances of detecting a planet around a lens star are greatest when the orbital distance is comparable in size to the Einstein radius.  In this Letter we focus on planetary systems located within about a kiloparsec. These systems are near enought that we can hope to conduct follow-up observations to learn more about the planet and perhaps eventually test for the presence of life.  It is convenient to define the parameter $\\alpha$ to be the ratio between the semimajor axis and the value of the Einstein radius: $a = \\alpha\\, R_E.$ Figure 1 shows values of $\\alpha$ as a function of distance, for planets in the zones of habitability around their stars.  Depending on the stellar mass, planets in the habitable zone are detectable through their action as lenses for planetary systems as close to us as a parsec (for the lowest-mass dwarf stars), and out to distances as far as a kiloparsec or more when the stellar mass is $0.5\\, {\\rm M_\\odot}$ or larger. ",
        "conclusions": "Until now, searches for habitable planets have relied on very successful techniques based on Doppler shifts and on transits.  We point out that gravitational lensing is a potentially powerful complement to these proven methods. Because it relies on light from background sources, gravitational lensing can discover planets in the habitable zones of even the dimmest stars, and can be successful even if spectral lines and/or the measured flux exhibit erratic variations. Because it relies only on the distribution of the lens-system's mass, even face-on systems, inaccessible to Doppler and transit techniques, can be studied.  The challenge faced by lensing studies is that the lens must pass in front of a background source of light. This requires that the planetary system lie in front of a dense source field, or else that its path happens to bring its projected position close to a background field star. Predictions \\citep{DiStefano2005}, bolstered by the discovery of events caused by nearby lenses \\citep{Nguyen2004,Kallivayalil2006}, indicate that both types of events are common.  Monitoring observations of the present and near future, including OGLE, Pan-STARRS, and LSST can discover events caused by planets in habitable zones of nearby stars. In most cases, the presence of the star and its proper motion will already have been catalogued by the same monitoring data. Software must be developed to recognize even subtle lensing-like deviations in real time, so that when they occur at the location of a known possible lens, a separate program of frequent follow-up observations can be initiated.  Targeted lensing observations are a new frontier. As has been the case for Doppler and transit studies, it will take time for them to yield results, but the results are well worth it. For any given nearby dwarf star, we can discover the presence of a planet, or place limits on the existence of one. If a planet is discovered, the quantities that can be derived under ideal circumstances are the planet's mass, its separation from the star, and its orbital period.  \\bigskip \\noindent The authors are grateful to Charles Alcock, David Charbonneau, Matthew Holman, Penny Sackett, Dimitar Sasselov for conversations. This work was supported in part by AST-0708924.   \\begin{table}[t] {\\scriptsize \\begin{tabular}{ccc|cccc|cccc|cccc|c} & & & \\multicolumn{4}{c|}{$\\Delta = 0.003$} & \\multicolumn{4}{c|}{$\\Delta = 0.01$} & \\multicolumn{4}{c|}{$\\Delta = 0.03$} & \\\\ \\hline $D_L$ & $\\alpha$ & $\\tau_E/P$ & $R_0$ & $R_1$ & $R_6$ & $R_{24}$ & $R_0$ & $R_1$ & $R_6$ & $R_{24}$ & $R_0$ & $R_1$ & $R_6$ & $R_{24}$ & $R_{cc}$ \\\\ \\hline 10 & 0.701 & 0.428 & 0.80 & 0.79 & 0.78 & 0.64 & 0.58 & 0.59 & 0.57 & 0.44 & 0.45 & 0.45 & 0.42 & 0.15 & 0.01 \\\\ 10 & 0.701 & 0.428 & 0.68 & 0.68 & 0.68 & 0.64 & 0.52 & 0.51 & 0.51 & 0.44 & 0.41 & 0.41 & 0.40 & 0.22 & 0.01 \\\\ \\hline 20 & 0.496 & 0.604 & 1.48 & 1.50 & 1.48 & 0.96 & 0.90 & 0.89 & 0.79 & 0.22 & 0.37 & 0.37 & 0.29 & 0.04 & 0.001 \\\\ 20 & 0.496 & 0.604 & 0.80 & 0.80 & 0.80 & 0.77 & 0.43 & 0.43 & 0.42 & 0.32 & 0.23 & 0.23 & 0.21 & 0.07 & 0.001 \\\\ \\hline 50 & 0.314 & 0.954 & 1.50 & 1.52 & 1.13 & 0.20 & 0.56 & 0.52 & 0.09 & 0.06 & 0.29 & 0.21 & 0.05 & 0.01 & 0.000 \\\\ 50 & 0.314 & 0.954 & 0.40 & 0.40 & 0.40 & 0.36 & 0.16 & 0.17 & 0.16 & 0.06 & 0.08 & 0.09 & 0.07 & 0.01 & 0.000 \\\\ \\hline 100 & 0.223 & 1.345 & 1.03 & 1.00 & 0.12 & 0.08 & 0.52 & 0.30 & 0.05 & 0.03 & 0.24 & 0.07 & 0.03 & 0.00 & 0.000 \\\\ 100 & 0.223 & 1.345 & 0.19 & 0.19 & 0.18 & 0.13 & 0.09 & 0.09 & 0.08 & 0.02 & 0.04 & 0.04 & 0.03 & 0.00 & 0.000 \\\\ \\hline 200 & 0.159 & 1.891 & 0.96 & 0.26 & 0.06 & 0.04 & 0.52 & 0.06 & 0.03 & 0.01 & 0.09 & 0.02 & 0.02 & 0.00 & 0.000 \\\\ 200 & 0.159 & 1.891 & 0.10 & 0.08 & 0.07 & 0.04 & 0.05 & 0.04 & 0.03 & 0.01 & 0.02 & 0.02 & 0.01 & 0.00 & 0.000 \\\\ \\end{tabular} } \\caption{ Rates for deviations of various strengths, as determined by simulation. For $M = 0.25 {\\rm M_\\odot}$, $M_{PL} = 0.26 M_{Jupiter}$, and $a = 0.1AU$ held fixed, a face-on system was placed at five different distances (as shown on Figure 1). For each value of $D_L$ there are two rows, the top including rotation, and the bottom for (unrealistic) static systems.  Rates are normalized to rate of events for which $A_{max} > 1.34$. If the number in the table is 1.00, then there will be as many events with deviations at the given level as there are events with a maximum magnification greater than 1.34. The rate of detectable deviations depends on two factors: the size of deviation $\\Delta$ needed for detection, and the duration. For a given $\\Delta$, we define $R_k$ to be the rate of deviations lasting at least $k$ hours. For instance, if deviations must have a strength of $1\\%$ and last 6 hours to be detectable, their rate may be found in the $R_6$ column under $\\Delta = 0.01$. Although the rotation period is the same ($P = 23.1$ days) for all systems considered here, each value of $D_L$ corresponds to a different Einstein radius $R_E$ and thus a different event duration $\\tau_E$. Therefore rotation is expected to play a larger role for more distant systems. This affects the event rate by slightly reducing the rate of long-lived deviations, but greatly increasing the rate of short-lived deviations.  $R_0$ is the theoretical maximum rate of deviations, if infinitely fine sampling could be achieved. Values of $R_0$ were determined not by simulation but by semi-analytic calculation. They should correspond to the maximum possible rate of deviations, but may be slightly lower than $R_1$ due to rounding errors.  } \\end{table}  \\begin{figure} \\psfig{file=avsdl.ps,height=11cm} \\caption{ $\\alpha$ vs $D_L$ for the habitable zone of low-mass stars. Each colored bar represents a star with a given mass: $M = 0.1\\, {\\rm M_\\odot}$ on the lower left, increasing by a factor of 2.5 for each subsequent bar.  The lower (upper) part of each bar corresponds to the inner (outer) edge of the habitable zone for a star of that mass.  The upper horizontal dashed line at $\\alpha = 1.6$ marks the approximate boundary between ``wide'' systems, in which the planet and star act as independent lenses \\citep{DiStefanoScalzo1, DiStefanoScalzo2}, and ``close'' systems in which distinctive non-linear effects, such as caustic crossings provide evidence of the planet \\citep{MaoPaczynski, GouldLoeb}. All of the planets detected so far have model fits with $\\alpha$ lying between $0.7$ and $1.6.$ In this range, the effects of caustics are the most pronounced. As $\\alpha$ decreases, the effect of the planet on the lensing light curve becomes more difficult to discern; the horizontal dashed line at $\\alpha = 0.2$ is an estimate of a lower limit. The probability of detecting the planet in such close systems ($\\alpha \\le 0.5$) is increased by the orbital motion. } \\end{figure}   \\begin{figure} \\begin{center} \\epsfig{file=contour.ps,width=14cm} \\end{center} \\caption{ Regions in the lens plane that correspond to deviations from the point-lens signature. All distance are measured in units of the Einstein radius. This lens has $M_{planet}/M = 0.001$ and separation $\\alpha$ = 0.314. If the star is 50pc away and has a mass of 0.25 ${\\rm M_\\odot}$, the planet would be in the habitable zone, at an orbital distance of 0.1 AU. The center of mass is placed at the origin, and the star is very slightly offset from this. The inner circle shows the planet's orbit.  At the instant of time depicted, the planet is located to the right of the star. The outer circle shows the Einstein radius; if the source passes through this circle, the peak magnification of the event will exceed 1.34. When the source is within a colored region, the difference $\\Delta$ between the actual magnification and the expected point-lens magnification exceeds a certain amount.  Blue (outermost): $\\Delta > 0.1\\%$.  Green: $\\Delta > 0.3\\%$.  Red: $\\Delta > 1\\%$.  Yellow (innermost): $\\Delta > 3\\%$. Note that the sizes of all of these regions are much larger than the caustic structures, which are too small to be shown on this plot. The regions of deviation rotate around the center of mass as the planet orbits. Since the outer region is much farther from the center, it moves much faster than the planet does. This effect increases the probability that the source will pass through a region in which $\\Delta$ is detectable. } \\end{figure}  \\begin{figure} \\psfig{file=lightcurve.ps,height=7cm} \\vspace{.5 true in} \\caption{ Typical deviations from the point-lens form, as found with our simulation. Plotted in each panel is the residual magnification $\\Delta(t)$ due to the planet. $\\Delta(t)$ is the magnification of the source star with respect to time, minus the point-lens amplification that would occur if the planet were not present. The x-axis shows time from the maximum of the deviation in days. For each of the five systems considered in our simulation, five deviations are shown, chosen randomly from the set of events with a deviation of $\\Delta > 0.01$ lasting at least 6 hours. Deviations include a mixture of close approaches (for which $A_{max}$ is large) and distant approaches (for which $A_{max}$ is small). This is because when $\\alpha$ is small, there are two general regions of deviation in the lens plane: one near to the lens and one far from it, as shown in Figure 2. Deviations on near approaches are due to crossing the region near the lens, and deviations on far approaches are due to crossing the region far from the lens. As $\\alpha$ approaches 1 (i.e., for small $D_L$), these two regions coalesce, and so for $D_L$ small, there are not two distinct deviation types, but rather a continuum. For the smallest values of $\\alpha$ (i.e., the largest values of $D_L$, for a given lens mass), deviations experienced in the outer region tend to be short-lived, so that the longest deviations occur when the source track traverses the central region.} \\end{figure}"
    },
    "0801/0801.4103_arXiv.txt": {
        "abstract": "We argue that Einstein gravity coupled to a Born-Infeld theory provides an attractive candidate to represent dark matter and dark energy. For cosmological models, the Born-Infeld field has an equation of state which interpolates between matter, $\\omega=0$ (small times), and a cosmological constant $\\omega=-1$ (large times). On galactic scales, the Born-Infeld theory predicts asymptotically flat rotation curves. ",
        "introduction": "The existence of dark matter and dark energy is now firmly established phenomenologically \\cite{Bradac,Spergel} but the theoretical understanding is far from complete.  Einstein equations require ``exotic\" components in the right hand side corresponding  to about $\\%96$ of the total energy density today. Understanding the microscopic nature of these extra components is one of the most challenging and important problems faced by theoretical physics at present.  Supersymmetric and exotic particles in the standard model are the best candidates for dark matter (for a review see \\cite{Jungman}).  Several experiments are now being devised for a direct detection of these particles. Alternative descriptions based on modifications to gravity have also been explored with interesting results. See \\cite{Sanders,Bekenstein,Moffat,Ferreira,Ferreira2,Carroll}, and references quoted therein, for some of these efforts and its consequences. For a recent review of the Einstein aether theory see \\cite{Jacobson}.  The problem of dark energy is somehow more recent, although the issue of the cosmological constant has been around for a long time. The discovery of an accelerating Universe \\cite{Expansion} resulted in deep changes in cosmology. The simplest explanation for this phenomena is a small positive cosmological constant, but many other possibilities have been explored (see \\cite{Carroll2,Lima} for recent reviews).  ~  In this paper we consider an action for general relativity coupled to a Born-Infeld theory. The Born-Infeld theory has as fundamental variable a symmetric connection  $C^\\rho_{\\ \\mu\\nu}(x)$. $C^{\\mu}_{\\ \\nu\\rho}$ has the same symmetries and transformation properties of the Christoffel symbol but is independent from it. The action is \\begin{equation}\\label{I} I[g_{\\mu\\nu},C^{\\mu}_{\\ \\nu\\rho},\\Psi]  = {1 \\over 16 \\pi G} \\int d^4x \\left[ \\sqrt{|g_{\\mu\\nu}|} R  + {2 \\over \\alpha l^2}\\sqrt{\\left|g_{\\mu\\nu}-l^2 K_{(\\mu\\nu)}\\right| } \\right] + \\int d^4x\\, {\\cal L}_m(\\Psi,g_{\\mu\\nu}), \\end{equation} where $|A_{\\mu\\nu}|$, for any $A_{\\mu\\nu}$, denotes the absolute value of the determinant of $A_{\\mu\\nu}$. $K_{\\mu\\nu}$ is the ``Ricci\" curvature associated to $C^{\\mu}_{\\ \\nu\\rho }(x)$, \\begin{equation}\\label{Kmn} K_{\\mu\\nu} \\equiv K^{\\alpha}_{\\ \\mu\\alpha\\nu } \\ \\ \\ \\   \\ \\  (K^{\\mu}_{\\ \\nu\\, \\alpha\\beta} = C ^{\\mu}_{\\ \\nu\\beta,\\alpha} + C^{\\mu}_{\\ \\sigma\\alpha} C^{\\sigma}_{ \\ \\nu\\beta} - [\\alpha \\leftrightarrow \\beta]). \\end{equation} Besides Newton's constant, the action (\\ref{I}) has two extra parameters: $l$ is a length and $\\alpha$ is dimensionless.  $\\Psi$ denotes all baryonic fields and ${\\cal L}_{{ m}}$ the baryonic Lagrangian.  The action (\\ref{I}) is similar in spirit although different in interpretation to the Born-Infeld gravity action proposed by Deser and Gibbons \\cite{Deser-Gibbons}, \\begin{equation}\\label{IDG} I[g_{\\mu\\nu}]=\\int \\sqrt{|g_{\\mu\\nu} - l^2 R_{\\mu\\nu} + X_{\\mu\\nu}(R)|} \\end{equation} and elaborated in \\cite{DG2}.  As discussed in \\cite{Deser-Gibbons}, the term $X_{\\mu\\nu}(R)$ must be chosen such that the action is free of ghost, and free of Schwarzschild-like singularities. The action (\\ref{IDG}) is an action for pure gravity, and it can be seen as a natural extension to spin two of the scalar $\\sqrt{|g_{\\mu\\nu} + \\partial_\\mu \\phi \\partial_\\nu\\phi| }$ and vector $\\sqrt{|g_{\\mu\\nu} + F_{\\mu\\nu}|}$ Born-Infeld (BI) theories.  For the scalar and vector BI theories the equations of motion are of second order.  For the spin two theory this is not automatic and requires the addition of $X_{\\mu\\nu}$.  The action (\\ref{I}), on the other hand, gives rise to second order equations because $K_{\\mu\\nu}(C)$ depends on first derivatives of the field $C^{\\mu}_{\\ \\nu\\rho }$. This action, however, is not an action for pure gravity but gravity coupled to $C^{\\mu}_{\\ \\nu\\rho }$. The equations of motion are discussed in the appendix and in Sec. \\ref{Eq/Sec} below.  It is known (e.g. \\cite{Fradkin-T}) that general relativity with cosmological constant is dual to Eddington's action \\cite{Eddington} $I[C]\\sim\\int \\sqrt{|K_{\\mu\\nu}|}$. The action (\\ref{I}) can then be interpreted as general relativity interacting with its own dual field theory.  The action (\\ref{I}) can also be motivated by looking at general relativity without metric \\cite{B}. This interpretation will be discussed in Sec. \\ref{Sec/g=0}.  Our main goal in this paper is to argue that the field $C^{\\mu}_{\\ \\nu\\rho}$ has good properties to represent dark matter and dark energy. We shall study the equations of motion following from (\\ref{I}) and prove the following properties. \\begin{enumerate} \\item For a cosmological model, there exist solutions where the expansion factor $a(t)$ behaves as $a(t) \\sim e^{Ht}$ for large $t$, and as $a(t)\\sim t^{2/3}$ for small $t$. The equation of state for the fluid interpolates between $p=0$ and $p=-\\rho$.    The parameters in the solution can be adjusted such that this field contributes to $\\sim 23\\%$ of the total matter energy density and $\\sim \\%73$ of vacuum energy density, as required by observations \\footnote{Couplings between dark matter and energy have appeared in \\cite{Comelli}, and in \\cite{Bertolami} involving a Chapligyn gas.}.  \\item For a spherically symmetric configurations, the action (\\ref{I}) predicts asymptotically flat rotation curves, as required by galactic dynamics.  The parameters involved in this solution can also be adjusted to deal with realistic situations. \\end{enumerate}   We would like to stress the simplicity of this proposal. The ``Born-Infeld\" term is all we need to account for both dark energy and dark matter, at least for the problems described above.  More complicated tests, like lensing, fluctuations, and others will be discussed elsewhere \\cite{BFS,BRR}. See also \\cite{Davi}. ",
        "conclusions": "Dark matter and dark energy have unique properties and their understanding in one of the most crucial problems faced by theoretical physics today.   Dark matter does not interact with normal matter and this property has motivated us to look for fields which have this property somehow ``built in\". We have explored gravity coupled to connection $C^{\\mu}_{\\ \\nu\\alpha}$ field with a Born-Infeld action.  This theory comply with the main background properties normally attributed to dark matter and dark energy. First, the evolution of the scale factor in cosmological models has the right time dependence interpolating between pressureless matter and a cosmological constant.  At galactic scales dark energy is less relevant but dark matter still plays an important role. By an approximation valid for distances much smaller to the Hubble radius we have solved the equations of motion for spherical objects and find the expected rotation curves.  These curves satisfy the basic asymptotic flatness observed in galaxies providing new support for this proposal.  We have left several topics for the future. The stability of this theory and the study of primordial fluctuations are important to determine the CMB anisotropies. This will be reported in \\cite{BFS}.  On galactic scales a systematic fit with observational curves is necessary. This issue is presently under study and will be reported in \\cite{BRR}."
    },
    "0801/0801.3665_arXiv.txt": {
        "abstract": "We determine the age of 1,104 early-type galaxies in eight rich clusters ($z = 0.0046$ to $0.175$) using a new continuum color technique.  We find that galaxies in clusters divide into two populations, an old population with a mean age similar to the age of the Universe (12 Gyrs) and a younger population with a mean age of 9 Gyrs.  The older population follows the expected relations for mass and metallicity that imply a classic monolithic collapse origin.  Although total galaxy metallicity is correlated with galaxy mass, it is uncorrelated with age.  It is impossible, with the current data, to distinguish between a later epoch of star formation, longer duration of star formation or late bursts of star formation to explain the difference between the old and young populations.  However, the global properties of this younger population are correlated with cluster environmental factors, which implies secondary processes, post-formation epoch, operate on the internal stellar population of a significant fraction of cluster galaxies.  In addition, the mean age of the oldest galaxies in a cluster are correlated with cluster velocity dispersion implying that galaxy formation in massive clusters begins at earlier epochs than less massive clusters. ",
        "introduction": "The age of the stars in a galaxy is a key parameter not only to cosmological models (i.e. the epoch of initial star formation being a lower limit to the age of the Universe), but also to our understanding of the star formation history of galaxies.  The age of stellar populations in galaxies guides our understanding of likely formation scenarios, e.g. monolithic versus hierarchical models (Kauffmann, White \\& Guiderdoni 1993).  However, even if a galaxy's stellar population age is decoupled from the formation epoch, the history of star formation and chemical evolution are determined primarily by the age of the stars and this value still serves as one of the most important variables in understanding galaxy evolution (Tinsley 1980).  Of course, since Nature is cruel (de Sade 1797), by far the hardest parameter to extract from a galaxy's spectral energy distribution is the age of the underlying stellar population.  The reason is mostly technical, in that there is no single color or spectral index that uniquely measures age (as derived from the turnoff point for the stellar population).  Thus, one is required to deduce a mean age based on a series of observables, such as the combined $<$Fe$>$ versus $H\\beta$ index (Worthey 1994) or $UV$ versus optical colors (Kaviraj \\etal 2006), guided by synthetic models based on stellar physics.  These models have only recently achieved a necessary level of success to transform from isochrone fits to synthetic colors, i.e. all the relevant stellar physics is correctly incorporated (Yi \\etal 2003, Schiavon 2007).  Techniques to measure metallicity and age have divided into two paths, use of line indices (e.g. the Lick/IDS system, Gonzalez 1993, Trager \\etal 2000) versus continuum colors (Rakos \\& Schombert 2005).  These two techniques are, effectively, attempting to determine age and metallicity through slightly different astrophysical phenomenon.  Line indices can trace the metallicity of a system by directly measuring the abundance of various element lines (although still influenced by the age of the system, Trager \\etal 2000).  Early work (Burstein \\etal 1984) focused on easily measurable lines such as Mg$_2$ and Ca $H$ and $K$ (Faber 1973).  These individual lines contain a wealth of information on the chemical history of stellar populations and the process of nucleosynthesis (Pagel 2001). However, extracting the mean metallicity of a stellar population, as expressed by [Fe/H], can be problematic if the ratio of these elements to Fe varies.  For example, Type Ia SN produce more Fe to $\\alpha$ elements than Type II, but require over 1 Gyr to evolve to the detonation stage such that the $\\alpha$/Fe ratio decreases with time (Maraston \\etal 2003). Thus, models which do not include these ratio variations risk incorrect estimates for color and line indice changes.  Improved technology has meant that, for a large number of galaxies, the relevant metallicity value, e.g. [Fe/H], can be determined directly from Fe lines themselves (e.g. Fe4383 and Fe5015, Gonzalez 1993, Trager \\etal 2000, S\\'{a}nchez-Bl\\'{a}zquez \\etal 2006).  As Fe is the primary source of electrons in stellar atmospheres, it is a trivial conversion from [Fe/H] to total metallicity, $Z$ or [M/H].  Age is determined in line indice studies, primarily, through the $H\\beta$ index (see Schiavon 2007 for a detailed discussion).  While this age determination is still model dependent, the separation in parameter space is clear (e.g. MgFe versus $H\\beta$).  Notably, in contrast to our expectation of old stellar populations from interpretation of the color-magnitude relation (Faber 1973), line indices studies find a high fraction of early-type galaxies with young ages (less than 7 Gyrs) in agreement with certain CDM frameworks for galaxy formation (Kauffmann, White \\& Guiderdoni 1993) that predict extended star formation histories (Kuntschner \\& Davies 1998, Trager \\etal 2000).  Later re-analysis, with refined spectral energy distribution (SED) models, finds that only 25\\% of the galaxies in previous samples have ages less than 7 Gyrs (Schiavon 2007).  However, even this fraction is unexpectedly high given the tightness of the color-magnitude relation and the expected increased scatter due to the influence of age on color (Bower \\etal 1992).  Narrow band continuum colors, used in this work, approach the age and metallicity determination problem from a different direction.  They measure the effects of metallicity, in composite stellar populations, by the change in color of the red giant branch (RGB) with [Fe/H].  Age effects are measured by color changes produced in the shifting of the turnoff point (Tinsley 1980).  This type of data collection (i.e. imaging) has the advantage of being faster (therefore, requiring less observing time or capable of penetrating to fainter luminosities than spectroscopy) and provides spatial information.  However, as with line indice work, this technique is also dependent on the use of SED models to interpret the resulting colors (see Steindling, Brosch \\& Rakos 2001 for an in-depth analysis of the technique).  We have recently adopted a more reliable principal component (PC) technique that is tied to the galactic globular age and metallicity scale (Rakos \\& Schombert 2005) which has resulted in a higher degree of accuracy in age determination than our previous papers.  The goal of this paper is three-fold.  First, we briefly review our technique, its strengths and weaknesses, our multi-metallicity models and, most importantly, where the key uncertainties lie.  Second, we will present our most recent compendium of ages and metallicities for cluster galaxies based on recent observations of four new clusters combined with data from four clusters previously published in the literature, now re-analyzed with our new techniques.  Third, using the age and [Fe/H] values extracted from our colors, we will explore the trends and correlations as constrained by the limitations of our technique.  Throughout this paper we use the Benchmark cosmological parameters of $\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$ and $H_o=75$. ",
        "conclusions": "\\subsection{Galaxy Age from Narrow Band Colors}  A graphical summary of our deduced mean ages and [Fe/H] for 1,104 galaxies in eight clusters (in order of increasing redshift: Fornax, A779, Coma, A400, A539, A1185, A119, A2218) is shown in Figures 2, 3 and 4, plots of log $\\tau$ (Gyrs), [Fe/H] and stellar mass (log $M_*/M_{\\sun}$).  Distance is derived using the Benchmark cosmology and stellar mass is calculated from the $M_{5500}$ luminosity assuming a M/L of 3 (Bender, Burstein \\& Faber 1992).  Since our sample is composed primarily of early-type galaxies, the amount of gas is minimal and, therefore, the stellar mass reflects the total baryonic mass of the galaxy.  Typical uncertainties, defined mostly by the SED models, are 0.5 Gyrs in age, 0.2 dex in [Fe/H] (although this will increase at lower luminosities due to a combination of higher errors in the observations and running against the limits of the SED models).  The uncertainties are represented by the greyscale image in each plot where each data point is taken to be a 2D gaussian of width given by the errors above.  Each datapoint's contribution per bin is summed to produce the greyscale intensity, normalized to the total sample size.  \\begin{figure} \\centering \\includegraphics[scale=0.95]{f2.eps} \\caption{Mean galaxy age versus galaxy stellar mass (in solar units) for all 1,104 galaxies in our cluster sample.  Cluster galaxies divide into two distinct populations, an old population with mean ages of 11.5 Gyrs and a younger population (the gap population) with a mean age of 9 Gyrs.  There is a slight tendency for the oldest galaxies to decrease in age with lower mass.  The statistically significant decrease in age with stellar mass for the gap population is driven by the increasing number of young galaxies ($\\tau <$ 7 Gyrs) for masses less than $10^9 M_{\\sun}$).  The green line displays the mean relation of age versus stellar mass from Gallazzi \\etal (2005), the red line displays the 1$\\sigma$ range in the SDSS sample of age per mass bin.  As is typical for line indices studies, the Gallazzi \\etal study finds sharply younger galaxy ages for decreasing mass, although we note that the SDSS sample includes all galaxy types as compared to our strict cluster sample. } \\end{figure}  All the key features noted over the decades relating to galaxy evolution are visible.  For example, the mass-metallicity relation is evident in Figure 3, although not the strictly linear relation as indicated by the color-magnitude diagram for galaxies (Bower \\etal 1992, Vazdekis \\etal 2001).  There is some suggestion in this diagram of two populations, a high mass, solar metallicity population with a strong correlation between mass and $<$Fe/H$>$ plus a lower mass population with lower metallicities and a much higher scatter.  Age does not appear to be a factor in the trend found in Figure 3 as there are numerous old galaxies with [Fe/H] values around $-$1 and plenty of younger galaxies with [Fe/H] near solar (although see Figure 13 for the effect of age on the scatter in the color-magnitude relation).  The only consistent trend with age is that all galaxies with $<$Fe/H$>$ values less than $-$1.5 have ages less than 10 Gyrs.  It would appear, since there is no evidence in Figures 3 or 4 that young galaxy age leads to high metallicity values, that the epoch of star formation, or its duration, is not a significant factor in determining the mean metallicity of a galaxy.  Mass is the primary driver of metallicity as predicted by infall models of chemical evolution (Pipino \\& Matteucci 2004), although probabily not the sole variable (see \\S3.7).  Determination of the age of the stellar populations in galaxies is the primary goal of this paper and a comparison in the range of galaxy ages we find versus the range in ages by two recent line indice studies (S\\'{a}nchez-Bl\\'{a}zquez \\etal 2006 and Thomas \\etal 2005) is found in Figure 5.  This comparison is key to our discussion as the S\\'{a}nchez-Bl\\'{a}zquez \\etal and Thomas \\etal studies use the Lick/IDS system of age determination (i.e. the model grid interpretation of the H$\\beta$ versus $<$Fe$>$ or MgFe index).  As can be seen from Figure 5, the distribution of galaxy ages differs dramatically between the Lick/IDS technique and our method.  Both the S\\'{a}nchez-Bl\\'{a}zquez \\etal and Thomas \\etal samples (only the high density sample is shown) find significant numbers of galaxies with ages less than 7 Gyrs.  Our technique is unable to calculate ages less than 3 Gyrs, however, those objects still signal their young ages with star forming colors and are rare in our cluster samples (Rakos, Maindl \\& Schombert 1996).  We also note that our technique finds that a majority of cluster galaxies have ages greater than 10 Gyrs, but old galaxies are in the minority in the samples that determine age by the Lick/IDS system.  For example, in Coma, we have three galaxies in common with the Thomas \\etal sample, NGC 4839, 4874 and 4889.  The Thomas \\etal ages are 15.8, 2.4 and 2.8 Gyrs respectfully.  Our continuum color ages are 11.6, 12.5 and 10.9 Gyrs for the same galaxies. Our ages are more in agreement with the expectations of stellar population evolution from intermediate redshift clusters (Rakos \\& Schombert 1995), where the Thomas \\etal values would predict large numbers of bright cluster ellipticals with star-forming colors, which is not seen.  \\begin{figure} \\centering \\includegraphics[scale=0.95]{f3.eps} \\caption{Average metallicity, $<$Fe/H$>$, versus galaxy stellar mass.  The average metallicity is the luminosity weighted metallicity from our PC technique adjusted for the expected underlying average metallicity (Schombert \\& Rakos 2008).  The mass-metallicity relation is evident, although, as noted by many previous studies, the scatter is higher than observational uncertainties.  There are statistically significant numbers of low mass galaxies with high metallicities and high mass galaxies with intrinsically low [Fe/H] values.  Thus, while galaxy mass (depth of its potential well) is the primary parameter in a galaxy's chemical evolution, it is not its sole determinate.  The green line displays the mean relation of [Fe/H] versus stellar mass from Gallazzi \\etal (2005), the red line displays the 1$\\sigma$ range in the SDSS sample of metallicity per mass bin.  The SDSS metallicity relation is in excellent agreement with our estimates. } \\end{figure}  The differences between the Lick/IDS system method and other techniques of age determination has been noted before (e.g. Bregman, Temi \\& Bergman 2006).  In order to confirm the legitimacy of our age values, we have turned to the third method mentioned in the Introduction, the use of longer baseline broadband filters to exploit the sensitivity of near-IR colors to the position of the RGB.  In this case, we can combine our narrow band colors in the Coma cluster (Odell, Rakos \\& Schombert 2002) with recent $K$ band data (Eisenhardt \\etal 2007).  There are 16 galaxies in common with their sample and ours, for which our near-blue color $vz-yz$ is plotted against their near-IR color $V-K$ in Figure 6.  We label each galaxy data point by its estimated age, using our technique, rounded to the nearest Gyr.  In this Figure it is clear that galaxies with the oldest ages ($\\tau > 10$ Gyrs) separate from the younger systems ($7 < \\tau < 10$ Gyrs). Thus, the Coma galaxies with near-IR colors separate nicely in this optical/near-IR color plane.  There is some overlap in ages, but within our expected errors of 0.5 Gyrs.  This Figure also allows an opportunity to check various SED models tracks to our age estimates and the absolute value of our age estimates (as our technique is tied to the galactic globular cluster ages).  In this case, we have chosen the most recent SED models from Bruzual \\& Charlot (2003, hereafter BC03).  Figure 6 displays the model color tracks over a range of metallicities ([Fe/H] from $-$0.8 at $V-K=2.5$ to $+$0.3 at $V-K=3.5$ for all the models) for a population of 8 and 12 Gyrs.  The dashed lines are the raw SSP models from BC03, which do not match the colors of the Coma sample.  This is not unexpected as galaxies are not composed of a single metallicity population, as represented by the SSP's.  As discussed in Schombert \\& Rakos (2008), the color-magnitude relation and the line indice versus color relations are best matched by a multi-metallicity model that incorporates a modified infall model of chemical evolution.  These models, referred to as the `push' models in Schombert \\& Rakos (2008), are shown as solid lines in Figure 6 and are an excellent match to the relative age differences.  Thus, we conclude that our absolute and relative age estimates conform to expectations from SED models.  \\subsection{Cluster Ellipticals Ages}  There are several features to the age and metallicity distributions to our sample that are common to all the clusters and visible in Figures 2, 3 and 4.  The first is that there clearly exist galaxies with old ages ($\\tau \\approx$ 12 Gyrs) in all clusters.  This mean value is comparable to the age of the Universe (i.e. these are the oldest galaxies) and this population makes up the majority of cluster galaxies.  \\begin{figure} \\centering \\includegraphics[scale=0.95]{f4.eps} \\caption{Mean galaxy age versus mean metallicity.  The two populations are still evident in this figure.  The oldest galaxies display a range of mean metallicities (from $-$1.5 to $+$0.3).  The gap population displays a correlation of decreasing metallicity with age. } \\end{figure}  Second, the red cluster galaxies divided neatly into two populations based on age, with a distinct gap in age between the populations.  The older population has a mean of 11.5 Gyrs, the younger population a mean age of 9.5 Gyrs.  The gap is wider than our estimated error of 0.5 Gyrs, therefore, appears to be a real feature.  This gap is not seen in studies that use the Lick/IDS, such as Thomas \\etal (2005) or S\\'{a}nchez-Bl\\'{a}zquez \\etal (2006), but this may be due to small number statistics.  The transition in age with mass is sharp in the Gallazzi \\etal (2005) work (see their Figure 9, also drawn in our Figure 2), which involves over 175,000 galaxies from the SDSS database, but does not take on a bimodality distribution as suggested by our data sample.  We note that the SDSS sample includes all galaxy types, both field and cluster, whereas our sample is a strict, high density, cluster population.  The inclusion of star-forming disk systems would clearly deviate the correlation of age and mass to younger ages.  \\begin{figure} \\centering \\includegraphics[scale=0.95]{f5.eps} \\caption{A comparison between galaxy ages derived by continuum colors (this work) and those derived from two Lick/IDS studies (Thomas \\etal 2005, high density sample, and S\\'{a}nchez-Bl\\'{a}zquez \\etal 2006).  Typical for spectroscopic studies, the Lick/IDS work finds a higher fraction of young galaxies ($\\tau <$ 7 Gyrs) compared to studies that use colors (Bregman, Temi \\& Bergman 2006). } \\end{figure}  Lastly, there are a couple trends (i.e. weak correlations) with age.  In general, brighter luminosity (higher mass) systems tend to be older.  Older galaxies also tend to be more metal-rich (as measured by [Fe/H]).  However, neither of these trends is strong suggesting that the path that a galaxy evolves through to reach these global values of age and metallicity is influenced by several factors not directly related to its total mass.  In addition, the lack of sharp correlations indicates that the global values we measure, such as luminosity-weighted age, may be composed of a more convoluted underlying stellar population.  We will examine each of these features, and combine our results here with previous results, in an attempt to untangle our observations in the following sections.  \\subsection{The Oldest Galaxies}  The oldest galaxies in our sample have mean ages between 11 and 12 Gyrs, which is consistent with the assumption of a primordial epoch of initial star formation and the age of the Universe under the Benchmark model (12.6 Gyrs).  While there exist galaxies with calculated ages of 13 Gyrs (only 1\\% of the sample galaxies have ages greater than 13 Gyrs, none older than 13.5 Gyrs), they are consistent with the age of the Universe within the errors of the age determination method.  Extremely old stellar populations in galaxies are not unexpected under certain galaxy formation scenarios (e.g. monolithic collapse, Eggen, Lynden-Bell \\& Sandage 1962, Larson 1974, Arimoto \\& Yoshii 1987), it is somewhat surprising to find large numbers of purely old galaxies in our sample as star formation has been observed in ellipticals between redshifts of zero and 1 (Stanford \\etal 2004).  Later star formation would push our integrated age estimates to younger values (see \\S3.4).  Two possible scenarios can resolve this dilemma.  The first is that the observed high redshift star formation events may involve a very small fraction of the total stellar population mass and, thus, their anomalous colors have since faded below detection, overwhelmed by the colors of an older stellar population.  The second is that the number of galaxies undergoing these recent bursts have been overestimated.  There certainly exist galaxies in our cluster samples with younger mean ages (see gap population discussion below) and they may be the ancestors of these higher redshift objects.  The existence of galaxies with ages nearly equal to the age of the Universe places a series of important constraints on scenarios for galaxy formation. While there are clearly evolutionary processes in play within rich clusters (i.e., mergers and tidal collisions which induce later star formation, Moore \\etal 1996), it appears, from the properties of our reddest galaxies, that there must exist a sizable fraction of cluster ellipticals whose dominant stellar population dates back to the epoch of cluster formation itself.  These are valuable inputs to our galaxy formation models as these ages allow for no room in the time budget for later epochs of star formation (which would lower the deduced mean age by our method) nor for a very long duration of initial star formation after the time of galaxy formation (which would also decrease the deduced mean age).  We note that that Thomas, Maraston \\& Bender (2002) find that the oldest galaxies in their sample have the highest $\\alpha$/Fe ratios, implying very short initial durations of star formation in agreement with their old ages. Short durations of star formation imply that the resulting chemical evolution is direct, in agreement with our determination that the internal metallicity distribution in ellipticals are bested described by simple infall models (Schombert \\& Rakos 2008).  At the very least, our oldest galaxies are a challenge to hierarchical models where galaxies are assumed to assemble gradually with time by mergers (Baugh, Cole \\& Frenk 1996).  All the stars in these merging systems would be required to have the same ages and chemical history to explain the mass-metallicity relation (although perhaps the scatter in the mass-metallicity diagram is a signature of past stochastic mergers, see de Lucia \\etal 2006 and Kaviraj \\etal 2005), $\\alpha$/Fe increases with mass and the age-mass relation (see \\S3.7).  While it is plausible to have all the galaxy components with the same epoch of initial star formation, our models of chemical enrichment require a varying mass to produce varying metallicities.  Thus, hierarchical formation is an inadequate explanation for these oldest cluster ellipticals.  \\subsection{Gap Population}  Common to seven of our eight clusters is a second, younger population (the exception is A119).  This second population is distinguished by a clear gap, varying in width of between one to two Gyrs, separating the younger population (which we will refer hereafter as the gap population) and the oldest galaxies.  While this gap population is not as numerous as the oldest galaxies, it is significant.  When we divide all galaxies with metallicities greater than [Fe/H]=-1.7 (the limit of our models) into three bins of greater than 10 Gyrs (591 galaxies, 53\\%), between 7 and 10 Gyrs (431 galaxies, 39\\%) and less than 7 Gyrs (88 galaxies, 8\\%), we see that the gap population constitutes about 1/3 of all cluster galaxies.  The oldest galaxies, having ages near the age of the Universe, must be objects with simple star formation histories, i.e. a short initial burst followed by a quiescent history.  However the gap population, with a mean age between 8 and 9 Gyrs, can reach this stage through various paths. First, it must be remembered that the calculated age is a luminosity weighted mean value.  For example, a mean age of 8 Gyrs can be obtained by a late epoch of star formation (5 Gyrs later than the oldest galaxies) or by an extended duration of star formation, i.e. going from 12 Gyrs ago up to 4 Gyrs from the current epoch to average as 8 Gyrs).  A comparison of just those two exact scenarios is shown in Figure 6, where we linearly combine SED models from 4 to 12 Gyrs (assuming equal mass for each age, green line) and compare them to a single age 8 Gyr model.  As can be seen in Figure 6, it is impossible, in practice, to distinguish these two paths of star formation history simply by colors.  A secondary clock to determine the duration of initial star formation in a galaxy can be estimated by the ratio of $\\alpha$ elements to Fe (Greggio \\& Renzini 1983, Thomas, Greggio \\& Bender 1998), since delayed enrichment by Type Ia supernova raises the abundance of Fe relative to $\\alpha$ elements (such as Mg and Ca).  An increasing $\\alpha$/Fe ratio with galaxy mass is found by Thomas, Maraston \\& Bender (2002), which is interpreted as increasing initial star formation duration with decreasing galaxy mass ranging from less than 0.5 Gyrs duration for the highest mass galaxies to 6 Gyrs in duration for galaxies around $10^9 M_{\\sun}$.  These duration values would produce mean galaxy ages consistent with the gap population ages, although we note that while there is a slight decrease in age with lower mass for the gap population (Figure 4), this trend is not as strong as would be expected from the change in $\\alpha$/Fe with mass (Gallazzi \\etal 2006).  \\begin{figure} \\centering \\includegraphics[scale=0.95]{f6.eps} \\caption{Near-IR colors versus our narrow band color, $vz-yz$, for sixteen Coma ellipticals.  Galaxies with ages less than 10 Gyrs are labeled by age in blue, galaxies older than 10 Gyrs are labeled in red.  Two sets of 8 and 12 Gyrs tracks are shown from Bruzual \\& Charlot (2003) SED models.  The dotted lines are for SSP models, the solid lines are for our multi-metallicity 'push' models (Schombert \\& Rakos 2008) which blend in a range of stellar metallicity populations using a simple infall model of chemical evolution.  The 'push' models are an excellent match to the data and the separation of age, as given by our continuum color technique, is clear.  A blend model is shown in green, a sum of 16 models with ages between 4 and 12 Gyrs in 0.5 steps to demonstrate that it is impossible to distinguish between a galaxy with a singular age of 8 Gyrs versus the sum of ages from 4 to 12 Gyrs. } \\end{figure}  In addition to varying durations of star formation, an alternative explanation to the younger age for the gap population is to consider the effect of a more recent star formation event averaged with an old (12 Gyrs) stellar population.  Recent burst models (referred to as frosting models, Trager \\etal 2000) are parameterized by the fraction of the mass of the galaxy involved in the burst and when the burst occurred.  Obviously, the effect of a star formation burst on the integrated color of a galaxy, will be less if the burst involves a small fraction of the galaxy mass and/or occurs farther in the past (the population having time to redden and blend in with the older stars).  In order to constrain the magnitude of a frosting star formation event, we can ignore any burst that would still leave a signature of stars hotter than A type at the present day.  For these stars, in sufficient numbers, would dominate the near-blue colors in early-type galaxies, which have been known for decades to have the continuum shape represented by K giants (O'Connell 1980).  However, a population that is over one Gyr old has already taken on the colors close to a metal-rich, old population and is only identifiable when mixed with an older population using hyper accurate measurements, at least in the optical region of the spectrum (Rose \\etal 1994).  \\begin{figure} \\centering \\includegraphics[scale=0.95]{f7.eps} \\caption{Same as Figure 6, near-IR versus narrow band colors.  The same infall models from Figure 6 are shown.  In addition, two 'frosting' models, where a fraction of the galaxy mass is a population of 1 Gyr mixed with a 12 Gyrs population.  A 1\\% frosting burst is nearly identical to a standard 8 Gyrs population. } \\end{figure}  To test the effect of an old, weak burst on the measured mean age of our galaxies, we have mixed a 1 Gyr population of varying fractions to an old 12 Gyr population using the BC03 models.  Returning to the optical/near-IR diagram we used in Figure 6, we have plotted a 1\\% burst (involving 1\\% the mass of a galaxy) and a 5\\% burst for a range of mean metallicities (again using a metallicity distribution calculation as given by Schombert \\& Rakos 2008).  These tracks are shown in Figure 7, again along with our Coma data. From these tracks, it is clear to see that even a small burst of 1\\% of the mass of the galaxy, after 1 Gyr, has produced a trend of color which closely resembles a composite 8 Gyrs population.  Stronger bursts, of up to 5\\%, can be ruled out by the $vz-yz$ and $bz-yz$ narrow band colors (as seen by the dotted line in Figure 7).  We conclude it is impossible to distinguish between a small (less than 1\\%), recent bursts or a composite 8 Gyrs model for the gap population.  \\subsection{Age-Density Relation}  A sharp separation in age for cluster early-type galaxies suggests either a formation process which delays or extends star formation for a subset of galaxies (i.e. a nurture scenario) versus one where some environmental process has operated on a subset of galaxies after the formation epoch to induce later star formation (i.e. a nature scenario).  Environmental processes are typically tested by looking for correlations with galaxy density, on the assumption that the operating process is tied to the dynamical interactions either with the cluster tidal field (galaxy harassment, Moore \\etal 1996) or by galaxy interactions.  While this environmental effect may be recent, as suggested by the frosting models, this may also be a remnant of a primordial event.  For example, one of the key predictions for hierarchical galaxy formation models is a difference between field and cluster galaxy ages (where cluster systems form first and, therefore, have older stellar populations).  Our sample is strictly a sample of rich cluster galaxies, however, we can test for any correlation with local density or with cluster radius, a crude measure of local density.  Just such a test is found in Figure 8, a plot of galaxy density as a function of cluster radius in kpc.  \\begin{figure} \\centering \\includegraphics[scale=0.95]{f8.eps} \\caption{Galaxy density as a function of cluster radius for galaxies older than 10 Gyrs (red), between 6 and 10 Gyrs (green) and younger than 7 Gyrs (blue).  Young galaxies clearly avoid the cluster center, however, even intermediate aged galaxies are less clustered than the oldest galaxies. Older galaxies are the most concentrated to the cluster center (higher densities), all these trends suggest an environmental influence on mean galaxy age. } \\end{figure}  To produce this plot, we have binned all the clusters into one geography (i.e. no correction for cluster size).  The cluster center is determined as either as the position of the brightest cluster member, or the geometric center of the cluster weighted by galaxy mass.  A119 was excluded in the calculations as it does not contain a gap population and A1185 was excluded since its irregular structure made the determination of the cluster center problematic.  The resulting sample of 652 galaxies were binned in 100 kpc radii and the number of galaxies per annulus area were summed for galaxies with ages greater than 10 Grys (above the gap), between 7 and 10 Gyrs (below the gap) and less than 7 Gyrs.  Immediately obvious from Figure 8 is that galaxies with younger mean age avoid the denser regions of the clusters and the older galaxies are more concentrated.  It is not surprising to find the very youngest galaxies located on the cluster outskirts as these objects are structurally similar to disk galaxies (see \\S3.6) and have similar properties to the Butcher-Oemler population.  However, even the intermediate aged galaxies are less concentrated than the oldest cluster members.  This effect, that younger galaxies avoid the cluster core, is slightly correlated with cluster type, the more evolved cluster types display a stronger separation by age than clusters that are dynamically young (as expressed by the BM and RS types), leading to the obvious conclusion that dynamical processes influence a galaxy's star formation history.  It is also notable that the highest redshift cluster in our sample (A2218) displays this separation to the greatest effect suggesting a process that has evolved recently in lookback time.  \\begin{figure} \\centering \\includegraphics[scale=0.95]{f9.eps} \\caption{Cluster velocity dispersion versus the mean age of the oldest galaxies in each cluster.  The clusters with the highest velocity dispersions (i.e. highest mass) have the oldest galaxies, implying earlier cluster formation epochs. } \\end{figure}  Environmental effects are not limited to the gap population.  Figure 9 displays the average age of the oldest galaxies (the 1st population) as a function of cluster velocity dispersion.  Although the sample size is small, there is a clear correlation between the mean age of the oldest galaxies and dynamical state of the cluster as represented by velocity dispersion.  The correlation here is such that high mass, high velocity dispersion clusters have the oldest 1st populations, implying that these clusters started galaxy formation (and the initial epoch of star formation) before their lower mass counterparts by approximately 1 Gyr.  \\subsection{Age-Morphology Relation}  A correlation between galaxy age and local density would imply a correlation between galaxy type (morphology) and age by way of the cluster density-morphology relationship (Ball \\etal 2006).  To test this hypothesis we extracted morphological types and structural parameters for the Coma sample (see Odell, Rakos \\& Schombert 2002).  Visual morphology was divided into four types, $r^{1/4}$ objects (ellipticals), objects with distinct bulge and disk components (S0's), objects with spiral or irregular structure (Sp/Irr) and objects with lenticular or dwarf appearance (dE's). Structural morphology was based on surface photometric fits to objects brighter than $-$18.  These fits are divided into three types, bulge objects ($r^{1/4}$), bulge+disk objects (S0's) and exponential disks (lenticulars).  \\begin{figure} \\centering \\includegraphics[scale=0.95]{f10.eps} \\caption{Histograms of the ages of various morphological and structural types in Coma.  Visual morphology was divided into four types, $r^{1/4}$ objects (ellipticals), objects with distinct bulge and disk components (S0's), objects with spiral or irregular structure (Sp/Irr) and objects with lenticular or dwarf appearance (dE's).  Structural morphology was based on surface photometric fits to objects brighter than $-$18.  These fits are divided into three types, bulge objects ($r^{1/4}$), bulge+disk objects (S0's) and exponential disks.  Aside from a tendency for young galaxies to be lower in luminosity (therefore, a lenticular or dwarf elliptical appearance), various ages are found for all morphological and structural types.   Surprisingly, S0's do not have a different age distribution from ellipticals despite the expectation of quenching at intermediate redshifts. } \\end{figure}  Histograms of both visual and structural morphology are shown in Figure 10. Surprisingly, there is no clear trend in morphology with age, such that the percentage of galaxies classed E:S0:dE:S/Irr less than 10 Gyrs in age is 12\\%:14\\%:65\\%:9\\% versus greater than 10 Gyrs of 29\\%:34\\%:34\\%:3\\%. Elliptical and S0 type galaxies have the same distribution in age, a majority being old ($\\tau$ = 12 Gyrs), but also form a significant fraction in the gap population.  Star forming spirals and irregulars tend to be young, but a majority of the young galaxies in Coma are in the lenticular or dwarf category.  These morphological types are simply low luminosity ellipticals with exponential shapes characteristic of ellipticals less than $-$18 (Schombert 1987).  Their young age reflects the age-mass relation discussed in \\S3.1.  In terms of structure, the fraction of young to old galaxies are similar for galaxies with bulges (ellipticals), bulge+disks (S0's) or pure disk systems (lenticulars).  The fractions B:B+D:D with ages less than 10 Gyrs is 37\\%:21\\%:42\\% versus the fraction for galaxies with ages greater than 10 Gyrs of 44\\%:24\\%:32\\%, i.e. slightly more r$^{1/4}$ objects over pure disk systems for older galaxies.  Since the structural sample is composed solely of the brightest galaxies in Coma, this histogram reflects the lack of any morphological distinction in age for the top of the cluster luminosity function.  Whatever processes determines a galaxy's age is decoupled from its morphological evolution (outside of current star formation).  The lack of an age difference between ellipticals and S0's (or bulge objects versus disk lenticulars) is surprising as there is an expectation that S0 disks are only recently quenched (i.e. stripped of their gas) and should display a younger age (Bedregal \\etal 2006, Burstein \\etal 2005). Although we note that the largest B/D S0's are the oldest such that even the oldest S0's may have young disks buried in the light of their dominant bulges.  \\subsection{Age Correlations with Mass and Metallicity}  The existence of correlations between a galaxy's stellar population age and mass or mean metallicity has been debated in the literature (Conselice 2007), but several recent large studies have converged on a consistent picture (Thomas \\etal 2005, Gallazzi \\etal 2006, Jimenez \\etal 2006, Annibali \\etal 2007).  While differing in the details, all these studies find that galaxy age increased with stellar mass (luminosity) and metallicity ([Fe/H]), although the deduced correlations contain a large amount of uncertainty.  And, common to all these studies, the scatter in the correlations exceed the observational errors, implying secondary processes are significant.  For this study, noting again that we measure age and metallicity based on a different technique than line indices, we find the same general correlations between stellar mass (luminosity) and metallicity as previous work (Figures 2, 3 and 4).  While the correlation between stellar mass and metallicity ($<$Fe/H$>$) is evident, the scatter, like the line indices studies, is clearly greater than the observational uncertainties and, therefore, the depth of the gravitation well of a galaxy is not the sole parameter that drives the chemical evolution of a galaxy (as predicted by galactic wind infall models, Pipino \\& Matteucci 2006).  Dividing the sample by age (for example, considering the mass-metallicity relation solely for galaxies with ages greater than 10 Gyrs), we find that the oldest systems display a tighter mass-metallicity relation, however continue to display scatter larger than observational errors.  In terms of age, there is a distinct separation by mass between the old galaxies and the gap population.  For galaxies brighter than $M_{5500} = -18$, the breakdown between the old galaxies and the gap population is 68\\% to 32\\%.  For fainter galaxies, this division reduces to 49\\% to 51\\%.  But while there is a tendency to find the youngest galaxies to be of the lowest masses, there is no clear linear correlation.  This is due, primarily, to the fact that the division into two populations is so distinct that each population should be analyzed separately.  Considering the oldest galaxies first, this population has a mean age of 11.5 Gyrs.  To within the uncertainties, the slope of the old population with respect to stellar mass is very shallow with a change of less than 1 Gyr over the full range in mass.  For the younger gap population, there are larger numbers of young galaxies at low masses ($M < 10^9 M_{\\sun}$).  This agrees with the shape and scatter of the color-magnitude relation (Odell, Rakos \\& Schombert 2002), however, some of these galaxies probably represent the low mass Butcher-Oemler objects (Rakos \\etal 1996) with traces of recent star formation (see Figure 13).  Likewise, the small number of high mass galaxies ($M > 10^{10} M_{\\sun}$) with ages less than 8 Gyrs could easily represent cluster galaxies which have undergone recent mergers of disk galaxies with younger stellar populations or induced star formation.  The age versus metallicity diagram (Figure 4) displays a wealth of buried information.  The old population of galaxies shows a slight trend for an increase in $<$Fe/H$>$ with age, certainly the oldest galaxies ($\\tau$ $>$ 12 Gyrs) all display greater than solar metallicities.  But the trend is weak and galaxies with an age of 11 Gyrs have a range of [Fe/H] from $-$1.2 to $+$0.3.  The implies that the age of the stellar population does not have a large influence on chemical evolution and that the final metallicity of a galaxy is dominated by its mass.  The situation differs for the younger gap population.  Here we find a fairly good correlation between age and metallicity such that the older gap galaxies ($\\tau = 9$ Gyrs) have the highest metallicities (near solar). The younger gap galaxies ($\\tau = 7$ Gyrs) have lower metallicities, near [Fe/H]$=-1$.  We note that this trend of age and metallicity for the gap population is opposite to the correlation of age and metallicity (i.e. younger galaxies having higher metallicities) demonstrated by Trager \\etal (1998), Jorgensen (1999), Ferreras \\etal (1999), Terlevich \\& Forbes (2002) and S\\'{a}nchez-Bl\\'{a}zquez \\etal (2006).  However, an extensive study of the correlated errors by S\\'{a}nchez-Bl\\'{a}zquez \\etal finds no such relationship in high density regions.  As our cluster sample is extracted from the very highest density regions in the Universe, it is no surprise that our oldest galaxies display no correlation between age and metallicity.  Our gap population displays the opposite correlation as in found for field galaxies by the S\\'{a}nchez-Bl\\'{a}zquez \\etal work, which we interpret to mean that the measured age from these galaxies is determined by post formation processes unique to the cluster environment.  If the younger age of these galaxies is due to longer durations of initial star formation, then we would expect the opposite correlation as more star formation pushes towards higher metallicities.  Instead, the effect we find would support a frosting scenario such that younger age correlates with mass (smaller galaxies being more susceptible to a burst event), but that chemical evolution is primordial and independent of measured age.  \\subsection{Age Gradients}  One of the most common criticisms of the Lick/IDS technique is that it finds a high fraction of galaxies with line indice determined ages of less than 7 Gyrs, yet very red optical to near-IR colors. This is inconsistent with other observations of ellipticals from the Fundamental Plane and at intermediate redshifts (Rakos \\& Schombert 1995).  A re-analysis of the Trager \\etal dataset by Schiavon (2007) finds that the number of young ellipticals ($\\tau <$ 7 Gyrs) had been overestimated.  Using an improved set of model tracks, Schiavon finds 3/4's of the Trager \\etal sample to have mean ages greater than 7 Gyrs and only 1/4 younger.  This study finds only 7\\% of cluster galaxies younger than 7 Gyrs, so even the re-calibrated Lick/IDS method still finds a higher fraction of young systems.  However, the Trager \\etal sample covers a range of galaxy environments, whereas our data is strictly a cluster sample.  There are many studies suggesting an age difference between the field and cluster environments (Thomas \\etal 2005, S\\'{a}nchez-Bl\\'{a}zquez \\etal 2006, Annibali \\etal 2007) in the direction of younger ages in their field samples.  But even considering only their high density samples, we still do not find the wide range of ages suggested by studies using the Lick/IDS system.  Another comparison to the age and metallicity values extracted for the Lick/IDS system, with respect to our continuum color technique, can be made by examining the SDSS averaged spectra produced by Eisenstein \\etal (2003). These spectra consist of thousands of SDSS spectra selected by color and morphology.  These averaged spectra are divided into four luminosity bins and, for our comparison, we use the spectra taken over the full range of environmental densities.  Eisenstein \\etal publish their Lick/IDS values (which can be converted into age and metallicity through the MgFe and H$\\beta$ indices using BC03 models) and their spectra were re-reduced by Schiavon (2007) using a different technique to derive the same Lick/IDS indices.  Figure 11 displays the age-metallicity-stellar mass plane for our sample and the Eisenstein SDSS spectra.  Our data are averaged in galaxy mass bins of 0.5, shown by the black symbols in Figure 11.  Each data point is the average of the bin, the error bars display the dispersion within each bin (i.e. not uncertainties).  Note that the separation into two galaxy populations by age is blurred by this type of analysis, but does reproduce many of the correlations seen in line indices study (e.g. increasing age with galaxy mass).  The blue diamonds display the original Eisenstein \\etal Lick/IDS values converted into age and metallicity.  The red diamonds are the re-analysis values by Schiavon.  The green diamonds are the spectra convolved through our narrow band filters and converted into age and metallicity by comparison to our galaxy sample (the $uz$ filter is estimated from extrapolation of the near-UV portion of the SDSS spectra).  \\begin{figure} \\centering \\includegraphics[scale=0.95]{f11.eps} \\caption{Figures 2-4 redrawn with averaged values for our cluster sample. Error bars display the dispersion in the averages, not uncertainties.  This analysis blurs the dual nature of galaxy populations, but reproduces the correlations seen in other studies.  Blue diamonds display the model extracted values from the Eisenstein \\etal SDSS averaged spectra.  The red diamonds are Schiavon's (2007) re-analysis of those spectra using a different Lick/IDS technique.  The green diamonds are the ages and metallicities deduced from our continuum color analysis to the same spectra (see text for discussion). } \\end{figure}  Figure 11 displays some of the more common problems with the interpretation of Lick/IDS values.  For example, the raw Eisenstein \\etal values are in fair agreement with our age values as a function of mass (blue diamonds), but metal poorer than our average [Fe/H] values.  The Schiavon re-analysis recovers metallicity, but now the calculated ages are several Gyrs younger than our mean ages.  This is the same behavior discovered by Monte Carlo simulations of correlated errors between age and metallicity by Thomas \\etal (2005).  When we reduce the Eisenstein spectra through our continuum color scheme we recover the mean ages and metallicities of our cluster sample (although slightly younger and more metal poor due to the mixture of field and cluster spectra in the sample).  While the claim is made that the Lick/IDS H$\\beta$ index is age sensitive independent of abundance changes, the comparison in Figure 11 appears to indicate otherwise.  The source of the difference between the Eisenstein \\etal Lick/IDS values and the Schiavon analysis is the amount of in-fill corrections made to the spectra. Schiavon himself notes that, without these corrections, the H$\\beta$ ages would be approximately 14 Gyrs, closer to our values.  Secondary age estimates from the Lick/IDS system (such as the H$\\delta$ index) produce even younger mean ages, however, these indices have serious technical difficulties in their application (Prochaska \\etal 2007).  In section \\S3.4 we showed that it is impossible to distinguish, using our colors, between a galaxy with a mean age of 8 Gyrs and one with an old population (13 Gyrs) mixed with a 1\\% 1 Gyr burst.  However, the former scenario would produce ages from the Lick/IDS system in concordance with our values and the frosting scenario would be an obvious candidate to explain the discrepancies between our age estimates and ones produced by the Lick/IDS technique (Tantalo \\& Chiosi 2004, James \\etal 2005). Schiavon demonstrates that Eisenstein \\etal values could be reproduced by assuming a 0.5 to 1\\% burst with a 1 Gyr population, and exact match to the effects of a mild burst on our continuum colors (see Figure 7).  To see how a frosting scenario would resolve age discrepancy, consider the primary differences between the results derived from spectroscopy (i.e.  the Lick/IDS system) and our continuum color technique: 1) spectroscopic values for [Fe/H] tend to higher than our values and 2) spectroscopic ages tend to be younger than our calculated ages (see red versus green diamonds for [Fe/H] in Figure 11 and the age histograms in Figure 5). In our most recent study of the metallicity values for ellipticals (Schombert \\& Rakos 2008), we performed a detailed study of the differences in mean [Fe/H] as determined by the Lick/IDS system (i.e. a direct measure of Fe lines) versus our technique (which determines [Fe/H] by model interpretation of the colors as driven by the position of the RGB).  The correspondence between the two metallicities derived by spectroscopy and colors was good, however, the Lick/IDS values were consistently higher than our estimates by approximately 0.2 dex.  Examination into how the Lick/IDS measurements are made revealed the source of the difference in our two metallicities.  Being a spectroscopic technique, the Lick/IDS values are extracted by slit observations of typically round galaxies.  This means the light observed by spectroscopic methods is light extracted from a slice through the galaxy's core.  For galaxies in the redshift range of the Trager \\etal study (between 1,000 and 5,000 km/sec), this means that the light through the slit represents only 1/10 to 1/3 the light of a galaxy as compared to our total luminosities or aperture colors (for standard slit widths of between 2 to 5 arcsecs).  In addition, the light obtained by spectroscopy is highly weighted by core luminosity.  In the typical galaxy observed at the distance of Coma, 70\\% of the luminosity measured by slit spectroscopy is due to stars within 1 kpc of the center (versus only 20\\% for imaging).  For fiber studies, this effect is magnified as the data will only encompass the core of a galaxy with no halo component.  When combining this bias with typical color gradients, Schombert \\& Rakos (2008) estimated that the resulting Lick/IDS metallicity values would be approximately 0.2 dex higher than values deduced from all the galaxy light (using a simple model of internal metallicity distribution), exactly in agreement with the color deduced values.  Whether the difference in age measurements made by the Lick/IDS system, versus our continuum color technique, is due primarily to an aperture effect will depend on the degree in which age gradients are found in ellipticals.  Color gradient work finds that the measured gradients in ellipticals are fully explained by changes in metallicity from galaxy core to halo (Thomas \\etal 2005).  This is also supported by a study of color gradients as a function of redshift (Tamura \\etal 2000) which found no evidence for changes in color gradients due to lookback time out to a redshift of 1.  On the other hand, several line indice studies have detected significant age gradients in the H$\\beta$ index (Tantalo \\etal 1998, S\\'{a}nchez-Bl\\'{a}zquez \\etal 2006b).  Both these studies found younger stellar populations in their sample galaxy's cores, as well as finding the cores to be more more metal rich and enhanced in $\\alpha$ elements (which would imply star formation from a simple burst rather than a prolonged event).  \\begin{figure} \\centering \\includegraphics[scale=0.85]{f12.eps} \\caption{Age gradients for eight galaxies in A1185.  Shown are age estimates for various integrated apertures listed by log radius. Decreasing age gradients are shown in blue, increasing age gradients are shown in red, neutral in green.  While galaxies with younger ages are found in A1185, their gradients are insufficient to explain the young galaxy ages measured by line indices studies. } \\end{figure}  Thus, the discrepancy between ages determined by our continuum colors, which measure the integrated light from a galaxy, can be resolved if one considers a burst model.  As bursts of star formation are concentrated in galaxy cores (Thomas 1999), then this will produce mean ages from spectroscopy skewed towards the burst mean age.  In addition, the burst population, being formed later in a galaxy's life, will be enriched in metallicity (agreeing with metallicity gradients).  To consider age gradients in our own continuum color sample, we examined the change in age with radius for a sub-sample of galaxies in A1185.  These age values are displayed in Figure 12, where the aperture radius is used and the calculated age values are deduced from the integrated colors of each aperture (versus a differential color).  The observed gradient are small, typically only a few Gyrs.  There are equal number of positive and negative gradients in Figure 12 and certainly nothing to indicate that spectroscopic studies are being influenced by an extremely young core population.  However, as discussed in \\S3.4, our intermediate ages may be the combination of a young (less than 1 Gyr population) and an old, 12 Gyr population.  Spectroscopic signatures of a burst population (i.e. the Balmer lines) may dominate the our continuum colors (Serra \\& Trager 2007). If true, then this would mean that the Lick/IDS technique is a powerful method of detecting small, medium aged bursts, but 1) produces an age measurement for the entire galaxy stellar population that is biased towards young values, 2) overestimates the metallicity of the galaxy (instead it measures the core [Fe/H] which should be a maximal value for the galaxy) and 3) deduces trends in $\\alpha$/Fe which reflect the chemical evolution of the burst population, not the global values for the galaxy as a whole. All consistent with the differences we observe between color and line indice estimates of age and metallicity.  Lastly, the age differences between colors and line indices may also be due to other hot star contributions to a galaxy's color and spectrum.  For example, metal-poor horizontal branch (HB) stars display spectral features similar to new star formation signatures, particularly to the $H\\beta$ index and near-blue colors.  However, in our examination of the internal metallicity distribution for ellipticals (Schombert \\& Rakos 2008), we estimated the contribution from an old, metal-poor component with a multi-metallicity model and measured its effect on our colors.  We found that, while metallicity estimates must take an old population into account, the metallicity distribution derived is similar to that proposed by infall models of chemical evolution and that the calculated mean age is insensitive to the metal-poor components (Schombert \\& Rakos 2008).  This is a recent change to our calculations due to the fact that the newest SED models (Schulz \\etal 2002) use isochrone tracks and properly fit the advanced stages of stellar evolution (i.e. HB populations).  This result is also in agreement with a similar analysis by Trager \\etal (2005) where they find that the high H$\\beta$ values for Coma galaxies has a small component from old, metal-poor HB stars, but that a younger burst population must be mixed with an older, primordial population to account for the variation in metallicity and age indices."
    },
    "0801/0801.4759_arXiv.txt": {
        "abstract": "We show that Gamma Ray Bursts (GRBs) of known redshift and rest frame optical extinction detected by the {\\it Swift} satellite fully confirm earlier results concerning the distribution of the optical afterglow luminosity at 12 hours after trigger (rest frame time). This distribution is bimodal and relatively narrow, especially for the high luminosity branch. This is intriguing, given that {\\it Swift} GRBs have, on average, a redshift larger than pre--{\\it Swift} ones, and is unexpected in the common scenario explaining the GRB afterglow. We investigate if the observed distribution can be the result of selection effects affecting a unimodal parent luminosity distribution, and find that either the distribution is intrinsically bimodal, or most (60 per cent) of the bursts are absorbed by a substantial amount of grey dust. In both cases we suggest that most dark bursts should belong to the underluminous optical family. ",
        "introduction": "After 3 years since the launch of the {\\it Swift} satellite (Gehrels et al. 2004), the number of long GRBs with known redshift strongly increased. The optical multiband follow up of these events allowed also the analysis of the spectral energy distribution for a large fraction of their optical afterglows.  In Nardini et al. (2006a) we analysed the optical (in the $R$ band) luminosity distribution after 12h (rest frame time) of all the 24 pre--{\\it Swift}  long GRBs with known redshift and a published estimate of the host galaxy absorption $A_V^{host}$.  Most of them (i.e. 21/24) had optical luminosities $\\log{L_{\\nu_R}}$ that lie in a very narrow range with mean value $\\langle \\log{L_{\\nu_R}}\\rangle =30.65$ (monochromatic luminosities in units of erg s$^{-1}$ Hz$^{-1}$), with a dispersion $\\sigma=0.28$.  The remaining 3 events were about 15 times fainter than the main group. The clustering of the observed optical afterglow luminosities for most GRBs, and the hint of a bimodality, with a few underluminous events (3.6 $\\sigma$ smaller luminosities) have been found independently by Liang \\& Zhang (2006) and confirmed also in Nardini et al. (2006b) who considered a small number of GRBs detected by {\\it Swift} (but with their intrinsic optical extinction still unpublished).  In Nardini, Ghisellini \\& Ghirlanda (2008, hereafter NGG08) we tested whether the observed clustering and bimodality of the optical afterglow luminosities could be due to possible intervening selection effects. Our simulations showed that the observed distribution could be obtained either if there are indeed two intrinsically separated GRB families, or if there is a large amount (more than 1.5 magnitudes) of unrecognised achromatic (grey) absorption affecting most (but not all) sources. The existence of an underluminous family, or the presence of grey dust, could explain why a sizeable fraction of GRBs are optically dark.  The observed underluminous events should represent the ``tip of the iceberg'' of a much more populated optically fainter family that includes the undetected dark GRBs. J\\'ohannesson Bj\\\"ornsson \\& Gudmundsson (2007) found that the observed bimodality of the optical luminosity distribution can be reproduced in the standard fireball model only assuming a bimodality in the intrinsic model parameters (i.e. two families with different mean isotropic kinetic energy $E_0$). However,  the latter result seems not to agree with the X--ray data.   In this work we consider all GRBs detected after the launch of {\\it Swift} with known redshift and $A_V^{host}$. The obtained optical luminosity distribution confirms the results obtained in Nardini et al. (2006a) for what concerns both the clustering of the luminous family and the bimodal nature of the distribution. We then repeat the exercise of NGG08 on this larger sample of events to test whether the clustering and bimodal luminosity distribution is intrinsic or due to some selection effect. To this aim we collected all upper limits of the optically dark GRBs detected in the same {\\it Swift} epoch.  While we were completing our study, Kann et al. (2008) confirmed the bimodality of the intrinsic optical luminosities considering a slightly differently selected sample of events, and noted also a strong similarity between the pre--{\\it Swift} and the {\\it Swift} luminosity distributions.   We use $H_0=70$ km s$^{-1}$ Mpc$^{-1}$ and $\\Omega_\\Lambda=0.7$, $\\Omega_{\\rm M}=0.3$. ",
        "conclusions": "We have shown that {\\it Swift} bursts confirm the distribution of the luminosities of the optical afterglows observed at a fixed time (12h) in the rest frame of the source: there are two families, both contained in a narrow luminosity range, and with a gap between the two. The ratio of the averaged luminosities of the two families is about 25 (i.e. $\\mu_{faint}$=29.3, $\\mu_{bright}$=30.7). We proved that this observed dichotomy is not due to some simple intervening observational selection effects, but it must corresponds to an intrinsic bimodality either of the afterglow luminosity function itself or of the distribution of the absorption, with half of the burst affected only by moderate ``normal\" (i.e. chromatic) extinction, and the other half dimmed by a further 1.5--2 mag of ``grey\" dust absorption. The first possibility suggests a dichotomy of the intrinsic properties of the burst, while the second suggests a dichotomy of the properties of the GRB environment. We cannot (yet) distinguish between the two possibilities, but an increase of the number of GRBs (say, twice as many as we have now, with redshift, well monitored optical afterglow and estimate of the ``normal\", chromatic, host extinction) will make it possible to well constrain either the slope of the luminosity function or the shape of the grey absorption distribution. Other crucial information will come from the analysis of the optical to X--ray afterglow evolution. This will tell us if they belong or not to the same component (see e.g. Panaitescu 2007), with important consequences on the sometimes puzzling connection between these bands (e.g. Stratta et al. 2005, Urata et al. 2007, Troja et al. 2007, li, li \\& Wei 2008, Perley et al. 2008) and the nature of the host galaxy dust absorption.  Also for the X--ray luminosities there is some evidence of clustering and dichotomy, but there is no simple connection between this bimodality and the one observed in optical (Gendre \\& Bo\\\"er 2005; Gendre, Galli \\& Bo\\\"er 2008). Also the energetics of the prompt emission seem unrelated to the afterglow luminosities: the events with the lower bolometric isotropic energy $E_{\\gamma, iso}$ are also members of the optically fainter family (e.g. GRB 050416, GRB 060614, GRB 030329), but there are optically faint events with high  $E_{\\gamma, iso}$ (e.g. GRB 061007, GRB 061126). We therefore confirm what found in Nardini et al. 2006a, i.e. there is no evident correlation between $E_{\\gamma, iso}$ and the optical luminosities (see also Kann et al. 2008).  The observed optical luminosity distribution has no convincing explanation yet, but it can shed new light for understanding the existence of the optically dark long GRBs. During the past 3 years after the launch of {\\it Swift}, there have been 167 GRBs with at least one optical afterglow detection. If we define dark GRBs all the events observed, but not detected in the optical, and for which there is an optical upper limit, we find 126 events in these 3 years, i.e. about 40\\% of all long GRBs in the {\\it Swift} epoch. In our simulations about 2/3 of the $z<5$ events are undetectable and most of them are members of the low luminosity family. This overestimate of the number of the simulated dark burst is probably due to the assumption (made for simplicity) that the faint and the bright families have the same shape (even if with different normalisations). With this caveat in mind, we suggest that dark GBRs preferentially are optically underluminous GRBs."
    },
    "0801/0801.4273_arXiv.txt": {
        "abstract": "{Thermal wind emission in the form of free-free and free-bound emission is known to show up in the infrared and radio continuum of hot and massive stars. For OB supergiants with moderate mass loss rates and a wind velocity distribution with $\\beta\\simeq 0.8\\ldots 1.0$, no influence of the wind to the optical continuum, i.e. for $\\lambda \\la 1.0\\,\\mu$m, is expected. Investigations of stellar and wind parameters of OB supergiants over the last few years suggest, however, that for many objects $\\beta$ is much higher than 1.0, reaching values up to 3.5.} {We investigate the influence of the free-free and free-bound emission on the emerging radiation, especially at optical wavelengths, from OB supergiants having wind velocity distributions with $\\beta \\ge 1.0$.} {For the case of a spherically symmetric, isothermal wind in local thermodynamical equilibrium (LTE) we calculate the free-free and free-bound processes and the emerging wind and total continuum spectra. We localize the generation region of the optical wind continuum and especially focus on the influence of a $\\beta$-type wind velocity distribution with $\\beta > 1$ on the formation of the wind continuum at optical wavelengths.} {The optical wind continuum is found to be generated within about $2\\,R_{*}$ which is exactly the wind region where $\\beta$ strongly influences the density distribution. We find that for $\\beta > 1$, the continuum of a typical OB supergiant can indeed be contaminated with thermal wind emission, \\protect{\\it even at optical wavelengths}. The strong increase in the optical wind emission is dominantly produced by free-bound processes.} {} ",
        "introduction": "It is well established that for massive stars the appearance of (thermal) excess emission at infrared (IR) and radio wavelengths is caused by free-free and (to a small fraction also by) free-bound emission generated in their winds (see e.g. Panagia \\& Felli \\cite{PanagiaFelli}; Olnon \\cite{Olnon}). At radio wavelengths, the free-free excess emission is usually used to derive the mass loss rates of hot and massive stars (see e.g. Lamers \\& Leitherer \\cite{LamersLeitherer}; Puls et al. \\cite{Puls96}).  Waters \\& Lamers (\\cite{WatersLamers}) have investigated this excess emission in detail, especially in the near-IR region. These authors studied the influence of free-free and free-bound emission to the total continuum and pointed already to the importance of the wind velocity (and hence the wind density) distribution that can severely alter the wind contribution in the IR.  However, the investigations of Waters \\& Lamers (\\cite{WatersLamers}) were restricted to the IR and radio range, i.e. they calculated the wind contribution for $\\lambda \\ga 1\\,\\mu$m only, while during earlier studies Brussaard \\& van de Hulst (\\cite{Brussaard}) had noted that free-bound processes might become very important in the optical and UV range for temperatures typically found in the winds of hot stars and supergiants.  Whether the free-bound emission in the wind indeed influences the optical continuum, depends severely on the density distribution. For line-driven winds, the density follows from the equation of mass continuity, i.e. it is proportional to the mass loss rate, and inversely proportional to the wind velocity. A high wind density can therefore be reached by either a high mass loss rate, or a low wind velocity.  For massive stars with pronounced high mass loss rates like Wolf-Rayet stars, Luminous Blue Variables, or the group of B[e] stars, it is well known that the wind not only influences, but even dominates, the optical spectrum. In these objects, the wind is usually optically thick even in the visual range and thus completely hides the stellar spectrum. Some recent examples of the thermal wind influence in the form of free-free and free-bound emission at optical wavelengths have been published, e.g. by Guo \\& Li (\\cite{Guo}) for the case of Luminous Blue Variables, and by Kraus et al. (\\cite{Kraus07}) for a Magellanic Cloud B[e] supergiant.  The second density triggering parameter is the wind velocity distribution. For OB-type stars with line-driven winds, the wind velocity is very often approximated by a so-called $\\beta$-law, where $\\beta$ is in the range of $0.8\\ldots 1.0$. Such a $\\beta$ value causes a rather fast wind acceleration at the base of the wind, and the wind reaches its terminal value within a few stellar radii (see e.g. Lamers \\& Cassinelli \\cite{LamersCassinelli}). Consequently, the region of very high density that might cause enhanced free-bound emission is restricted to an extremely small volume around the stellar surface. We can, therefore, expect that for OB-type stars, which have only moderate mass loss rates and whose wind velocity distributions have $\\beta$ values between 0.8 and 1.0, there will be no noticable influence of the wind on their optical continuum emission.  During the last few years, huge effort has been made to determine precisely the stellar and wind parameters of OB-type supergiants in the Galaxy (e.g. Kudritzki et al. \\cite{Kudritzki}; Markova et al.\\,\\cite{Markova}; Fullerton et al. \\cite{Fullerton}; Crowther et al. \\cite{Crowther}; Puls et al. \\cite{Puls}), in the Magellanic Clouds (see e.g. Evans et al. \\cite{Evans}; Trundle et al. \\cite{Trundle}; Trundle \\& Lennon \\cite{TrundleLennon}), and beyond, e.g., in M\\,31 (e.g. Bresolin et al. \\cite{Bresolin}). Interestingly, many OB supergiants are found to have rather high $\\beta$ values (up to $\\beta = 3.5$, see Sect.\\,\\ref{highbeta}). Such high $\\beta$-values, resulting in a strong density increase close to the stellar surface due to a much slower wind acceleration, should have noticable effects on the thermal wind emission not even in the near-IR, but extending also to optical wavelengths. In this paper, we, therefore, aim to investigate and discuss in detail the influence of high $\\beta$ values on the wind emission of OB supergiants, especially at optical wavelengths. ",
        "conclusions": "We investigated the influence of the thermal wind emission produced by free-free and free-bound processes to the total continuum of normal OB supergiants. While for winds with a velocity distribution following a $\\beta$-law with $\\beta$ typically in the range of 0.8 to 1.0, no influence of the wind at optical wavelengths is expected, the situation can be different when $\\beta$ exceeds 1.0. High $\\beta$ values reaching even 3.5 have recently been found for many OB supergiants. Our investigations therefore concentrated on such high beta values and their influence on the wind continuum emission especially at optical wavelengths.  We found that the wind emission in the optical is generated within $2\\,R_{*}$, only. This region is exactly the region where $\\beta$ has its highest influence on the wind density structure. Since with increasing $\\beta$ the wind is accelerated much more slowly than for $\\beta = 1.0$, the density close to the stellar surface is strongly enhanced, leading to an enhanced production of especially free-bound emission at optical wavelengths. At the same time the stellar emission, which passes through the wind on its way to the observer, is absorbed. These effects increase with increasing $\\beta$. The total continuum of OB supergiants, for which $\\beta$ values higher than 1.0 are found, can thus contain {\\it non-negligible contributions of wind emission even at optical wavelengths}."
    },
    "0801/0801.1660_arXiv.txt": {
        "abstract": "The prompt $\\nu_e$ burst from a core-collapse supernova (SN) is subject to both matter-induced flavor conversions and strong neutrino-neutrino refractive effects. For the lowest-mass progenitors, leading to O-Ne-Mg core SNe, the matter density profile can be so steep that the usual MSW matter effects occur within the dense-neutrino region close to the neutrino sphere. In this case a ``split'' occurs in the emerging spectrum, i.e., the $\\nu_e$ flavor survival probability shows a step-like feature. We explain this feature analytically as a ``MSW prepared spectral split.'' In a three-flavor treatment, the step-like feature actually consists of two narrowly spaced splits. They are determined by two combinations of flavor-lepton numbers that are conserved under collective oscillations. ",
        "introduction": "\\label{sec:introduction}  The neutrino flux streaming off a collapsed supernova (SN) core is an intriguing astrophysical case where flavor transformations depend sensitively on some of the unknown elements of the leptonic mixing matrix~\\cite{Dighe:1999bi,Dighe:2007ks}. Since the conversion probabilities depend also on the time-dependent matter profile, a high-statistics SN neutrino observation may also reveal, for example, signatures for shock-wave propagation~\\cite{Schi, Foglish, Tomas, FogliMega, Dasgupta:2005wn, Fogli:2006xy, Friedland:2006ta}. While galactic SNe are rare, various ongoing and future experimental programmes depend on large detectors that, besides their main purpose, are also sensitive to SN neutrinos~\\cite{Autiero:2007zj}. Therefore, understanding the flavor evolution of a SN neutrino signal remains of topical interest.  The flavor transformation probabilities not only depend on the matter background, but also on the neutrino fluxes themselves: neutrino-neutrino interactions provide a nonlinear term in the equations of motion~\\cite{Pantaleone:1992eq, Sigl:1992fn} that causes collective flavor transformations~\\cite{Samuel:1993uw, Kostelecky:1993dm, Kostelecky:1995dt, Samuel:1996ri, Pastor:2001iu, Wong:2002fa, Abazajian:2002qx, Pastor:2002we, Sawyer:2004ai, Sawyer:2005jk}. Only recently has it been fully appreciated that in the SN context these collective effects give rise to qualitatively new phenomena~\\cite{Duan:2005cp, Duan:2006an, Hannestad:2006nj, Duan:2007mv, Raffelt:2007yz, EstebanPretel:2007ec, Raffelt:2007cb, Raffelt:2007xt, Duan:2007fw, Fogli:2007bk, Duan:2007bt, Duan:2007sh, EstebanPretel:2007yq, Dasgupta07}.  One peculiar aspect of the expected SN neutrino fluxes is the hierarchy $F_{\\nu_e} > F_{\\bar\\nu_e} > F_{\\nu_\\mu} = F_{\\bar\\nu_\\mu} =F_{\\nu_\\tau} = F_{\\bar\\nu_\\tau}$ so that there is an excess flux of $\\nu_e\\bar\\nu_e$ pairs over those of the other flavors. The nonlinear terms cause a collective transformation $\\nu_e\\bar\\nu_e\\to\\nu_x\\bar\\nu_x$, where $\\nu_x$ is a specific linear combination of $\\nu_\\mu$ and $\\nu_\\tau$. The detailed dynamics of this transition is complicated and several important aspects are only numerically observed, not analytically understood. Still, the most crucial point is that the pair transformation $\\nu_e\\bar\\nu_e\\to\\nu_x\\bar\\nu_x$ proceeds collectively much faster than ordinary pair annihilation, so we have to contend with a ``speed-up phenomenon'' \\cite{Sawyer:2004ai, Sawyer:2005jk}. The pair process does not violate flavor-lepton number. Being an instability in flavor space, it proceeds efficiently even for a very small mixing angle.  \\begin{figure}[b] \\includegraphics[angle=0,width=0.8\\columnwidth]{fig1.eps} \\caption{Profile of the matter potential $\\lambda$ and the effective neutrino-neutrino interaction potential $\\mu$ for an O-Ne-Mg core collapse SN~\\cite{Nomoto1984, Nomoto1987, Kitaura:2005bt,Janka:2007di}.\\label{fig:profile}} \\end{figure}  \\begin{figure*} \\includegraphics[angle=0,width=0.6\\textwidth]{fig2.eps} \\caption{Mass eigenstate fractions $P_{ii}$ as well as the $\\nu_e$ survival probabilities far away from the star, numerically computed using the SN model of Fig.~\\ref{fig:profile} and an initial flux of pure $\\nu_e$. \\label{fig:numericalsplit}} \\end{figure*}  A very different situation prevails in the interior of a SN core where the $\\nu_e$ distribution is determined by a large chemical potential that enhances the $\\nu_e$ density and suppresses the $\\bar\\nu_e$ density relative to that of $\\nu_x$ and $\\bar\\nu_x$ so that collective pair conversions are not possible. Neutrino-neutrino interactions are strong, but their only impact is to synchronize the flavor oscillations (``self-maintained coherence''). Significant flavor transformation here requires a violation of flavor-lepton number. However, the high density of ordinary matter suppresses the effective mixing angles between $\\nu_e$ and the other flavors. Therefore, in the interior of a SN core the individual flavor-lepton numbers are almost perfectly conserved~\\cite{Hannestad:1999zy}.  Immediately after collapse, the SN emits a prompt $\\nu_e$ burst that arises from the de-leptonization (neutronization) of the outer layers of the collapsed core. Once more we have a strongly enhanced $\\nu_e$ and a suppressed $\\bar\\nu_e$ flux relative to the other flavors~\\cite{Kachelriess:2004ds}. Once more, collective pair transformations are not possible: efficient flavor conversion requires a large violation of flavor-lepton number and is not possible if the ordinary matter density is large. At some distance from the neutrino sphere, the $\\nu_e$ flux encounters the usual MSW level crossings driven by the atmospheric neutrino mass difference $\\Delta m^2_{\\rm atm}$ (H~crossing) and by the solar mass difference $\\Delta m^2_{\\odot}$ (L~crossing), leading to the usual resonant transformations~\\cite{Dighe:1999bi,Kachelriess:2004ds}.  An interesting new case is motivated by the insight that SNe with the lowest progenitor masses of 8--10~$M_\\odot$, encompassing perhaps 30\\% of all cases, collapse before forming an iron core, the class of O-Ne-Mg core collapse SNe~\\cite{Nomoto1984, Nomoto1987, Kitaura:2005bt, Janka:2007di}. In state-of-the-art numerical simulations these SNe explode even in a spherically symmetric treatment (convection plays no role), largely because their envelope mass is very small. By the same token, the matter density profile above the core is very steep even at the time of core bounce.  In this case the H and L level crossings occur very close to the neutrino sphere and may well lie deeply within the collective neutrino region. This is illustrated in Fig.~\\ref{fig:profile} where we show $\\lambda(r)=\\sqrt{2}G_{\\rm F}n_e(r)$ of an O-Ne-Mg core progenitor star~\\cite{Nomoto1984, Nomoto1987, Kitaura:2005bt}. We also show $\\omega_{\\rm H}=\\langle\\Delta m^2_{\\rm atm}/2E\\rangle$ and $\\omega_{\\rm L}=\\langle\\Delta m^2_{\\odot}/2E\\rangle$ as horizontal lines, where the average is over the Fermi-Dirac spectrum of neutrino energies described below.  The intersection of $\\lambda(r)$ with these lines indicates the locations of the H and L level crossings.  In Fig.~\\ref{fig:profile} we also show the effective neutrino-neutrino interaction potential $\\mu=\\sqrt{2}G_{\\rm F}F_{\\nu_e}\\langle 1-\\cos\\theta\\rangle_{\\rm eff}$, where $\\theta$ is the angle between different neutrino trajectories and $\\langle\\ldots\\rangle_{\\rm eff}$ stands for a suitable average.  At large distances, $\\mu$~scales approximately as $r^{-4}$.  Collective neutrino effects driven by $\\Delta m^2_{\\rm atm}$ are important for $\\mu(r)\\agt\\omega_{\\rm H}$ and driven by $\\Delta m^2_{\\odot}$ for $\\mu(r)\\agt\\omega_{\\rm L}$.   Duan et al.~\\cite{Duan:2007sh} have recently shown that in this case the interplay of ordinary MSW conversions with collective oscillations leads to interesting effects. We start with a pure $\\nu_e$ flux with a Fermi-Dirac spectrum ($\\langle E_{\\nu_e}\\rangle=11$~MeV, degeneracy parameter $\\eta=3$), and numerically calculate the mass eigenstate fractions $P_{ii}$ and the $\\nu_e$ survival probabilities $P_{ee}$ far away from the star, as shown in Fig.~\\ref{fig:numericalsplit}. These plots are in qualitative agreement with the corresponding curves in Fig.~2 of  Ref.~\\cite{Duan:2007sh}. However, our $P_{ee}$ is constructed as an incoherent sum of the mass fractions, thus representing the physical situation far away from the star, where the oscillatory features seen in Duan et al.'s $P_{ee}$ have disappeared.  In inverted hierarchy, one observes that the neutrinos emerging from the star are in the $\\nu_2$ state at low energies and in the $\\nu_1$ state at high energies, with the transition taking place around $E \\approx 12$ MeV. This results in a step function in energy for $P_{ee}$.  In the normal hierarchy, the neutrinos emerging from the star are in $\\nu_1$ state for $E \\gtrsim 17$ MeV, in $\\nu_2$ state for $15 \\, {\\rm MeV} \\; \\lesssim E \\lesssim 17$ MeV, and in the $\\nu_3$ state for $E \\lesssim 15$ MeV. The bump seen around $5$ MeV is due to an abrupt change in the matter density profile used for the computation (see \\cite{Duan:2007sh} for details), and we do not address it here. The transition at $E\\approx 15$ MeV is rather sharp, however the one at $E \\approx 17$ MeV is not as abrupt. This results in a two-step function for $P_{ee}$, with the step at $E \\approx 17$ MeV somewhat smoothened out.   Broadly, this is an example of a ``MSW prepared spectral split.'' In a two-flavor language, it is explained as follows. The strong neutrino-neutrino interactions lead to a synchronization of the neutrino oscillations. The flavor polarization vector of the ensemble begins at high density essentially aligned with the $\\nu_e$ direction in flavor space. After passing the MSW region, the polarization vector emerges with a significant transverse component relative to the mass direction because the MSW transition is not fully adiabatic. Subsequently this MSW-prepared initial condition is subject to collective effects only. As the effective neutrino-neutrino interaction becomes weaker, the modes above a certain energy $E_{\\rm split}$ orient themselves along the mass direction, those with smaller energies in the opposite direction. This is precisely the ``neutrino only'' case studied in Refs.~\\cite{Raffelt:2007cb, Raffelt:2007xt} as a generic case for a spectral split, a case that requires one to prepare the polarization vector with a large transverse component. (For neutrinos plus antineutrinos, the MSW preparation is not necessary because the collective pair transformations alone engineer a split.)  Our main goal here is to use the picture of a MSW-prepared spectral split to derive analytically the main features seen in Duan et al.'s numerical study~\\cite{Duan:2007sh}, viz. the existence and the positions of the spectral splits.  The numerical model in Fig.~\\ref{fig:profile} shows that the MSW resonances are within the collective neutrino region, but not deeply inside.  Therefore it is not a priori obvious if the ``synchronized MSW preparation'' and the subsequent ``split'' can be clearly separated. For our discussion we therefore adopt a more schematic model. We artificially increase the neutrino-neutrino interaction strength (raise the $\\mu$-profile in Fig.~\\ref{fig:profile}) such that the MSW region and the spectral split regions are clearly separate. In this framework we first calculate the MSW preparation analytically and then study a three-flavor treatment of the spectral split, based on the machinery recently developed by two of us~\\cite{Dasgupta07}. We find that our analytic treatment reproduces the numerical results surprisingly well. It also reproduces all relevant features of the realistic case, i.e. with the $\\mu$-profile of Fig.~\\ref{fig:profile}, justifying our simplifying assumptions and verifying our general interpretation.  We first set up in Sec.~\\ref{sec:SNmodel} our schematic SN model that captures the features relevant for our treatment and continue in Sec.~\\ref{sec:eoms} with the equations of motion. In Sec.~\\ref{sec:MSW} we derive analytically the MSW-prepared three-flavor state that serves as input for our three-flavor spectral split study in Sec.~\\ref{sec:splitting}. We conclude in Sec.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have studied the three-flavor evolution of a $\\nu_e$ burst that first undergoes two MSW level crossings driven by $\\Delta m^2_{\\rm atm}$ and $\\Delta m^2_\\odot$, respectively, and then undergoes spectral splits by the adiabatically decreasing strength of the neutrino-neutrino interaction. This case study of an MSW-prepared spectral split serves as a proxy for the recent numerical study of the prompt $\\nu_e$ burst in an O-Ne-Mg core collapse SN. Here, the matter density profile is so steep that the sequence between MSW crossings and collective neutrino oscillations is reversed from what would be expected in a traditional iron-core SN.  First we have analytically estimated the population of the propagation eigenstates after the MSW transformations.  Due to the sharply falling matter density at the edge of the core of the star, the MSW transitions are not completely adiabatic, which helps in satisfying a precondition for the spectral splits to take place.  We have used the well-known analytic double-exponential formula for calculating the jump probabilities.  We have studied numerically the subsequent spectral splits. We have analytically determined the split frequencies on the basis of two conserved flavor-lepton number combinations that supersede the single conservation law encountered in a two-flavor situation. The neutrino mass-gap hierarchy allows for a factorization of an H split and an L split, similar to the factorization of the MSW effect into an H and an L crossing. The  dynamics of the split evolution is clearly seen to be a two-step process.  The dynamics of the two spectral splits can be understood in terms of the motion of the neutrino state in the ${\\bf e}_3$--${\\bf e}_8$ triangle diagram, which can explain many of the features of neutrino evolution qualitatively. The number of possible splits can be deduced by the location of the neutrino state inside the triangle. We have also shown how the positions of the splits can be calculated accurately given the initial neutrino spectra, and calculated the $\\nu_e$ survival probability analytically, that matches the numerical computations.  Our analytic treatment accounts very nicely for the numerical findings of Duan et al.~\\cite{Duan:2007sh}. In their case the MSW conversion and the spectral splits are spatially very close so that it is not a priori obvious that our schematic model would be a good representation. In our treatment we have enforced a clear separation between the MSW region and the spectral-split region by assuming the limit of large neutrino-neutrino interactions, a limit that also ensures that the MSW transition is perfectly synchronized and that the spectral split is perfectly adiabatic. A posteriori, however, our interpretation as a MSW prepared spectral split appears nicely justified and quantitatively appropriate. Our treatment is also a useful application of the three-flavor oscillation machinery developed by two of us recently~\\cite{Dasgupta07}.  It appears that the impact of collective neutrino oscillations on the propagation of the prompt $\\nu_e$ burst is conceptually and quantitatively well under control. Characteristic signatures of these flavor transitions in large underground detectors have also been recently investigated~\\cite{Lunardini:2007vn}. What remains is to observe these features in the neutrino signal of the next galactic~SN."
    },
    "0801/0801.4874_arXiv.txt": {
        "abstract": "{ VIVA H{\\sc i} observations of the Virgo spiral galaxy NGC~4501 are presented. The H{\\sc i} disk is sharply truncated to the southwest, well within the stellar disk. A region of low surface-density gas, which is more extended than the main H{\\sc i} disk, is discovered northeast of the galaxy center. These data are compared to existing 6~cm polarized radio continuum emission, H$\\alpha$, and optical broad band images. We observe a coincidence between the western H{\\sc i} and polarized emission edges, on the one hand, and a faint H$\\alpha$ emission ridge, on the other. The polarized emission maxima are located within the gaps between the spiral arms and the faint H$\\alpha$ ridge. Based on the comparison of these observations with a sample of dynamical simulations with different values for maximum ram pressure and different inclination angles between the disk and the orbital plane, we conclude that ram pressure stripping can account for the main observed characteristics. NGC~4501 is stripped nearly edge-on, is heading southwest, and is $\\sim 200-300$~Myr before peak ram pressure, i.e. its closest approach to M87. The southwestern ridge of enhanced gas surface density and enhanced polarized radio-continuum emission is due to ram pressure compression. It is argued that the faint western H$\\alpha$ emission ridge is induced by nearly edge-on ram pressure stripping. NGC~4501 represents an especially clear example of early stage ram pressure stripping of a large cluster-spiral galaxy. ",
        "introduction": "}  The Virgo cluster is dynamically young and spiral-rich making it an ideal laboratory for studying the influence of the cluster environment on spiral galaxies. Most of the Virgo spiral galaxies are H{\\sc i} deficient, i.e. they have lost a significant amount of their ISM (Chamaraux et al. 1980, Giovanelli \\& Haynes 1983). Imaging H{\\sc i} observations have shown that these galaxies have truncated H{\\sc i} disks (Cayatte et al. 1990). Thus, the cluster environment changes the H{\\sc i} content and morphology of Virgo cluster spiral galaxies. However, it remains unclear at which cluster distance the spiral galaxies lose their external H{\\sc i} disk. Based on deep Effelsberg H{\\sc i} observations around 6 Virgo galaxies and a balance of previous detections of extraplanar gas in targeted galaxies Vollmer \\& Huchtmeier (2007) propose a global picture where the outer gas disk (beyond the optical radius $R_{25}$) is removed much earlier than expected by the classical ram pressure criterion. They argue that the vulnerable outer disk is removed much more rapidly than predicted by the Gunn \\& Gott criterion. Furthermore, Chung et al. (2007) found in a new H{\\sc i} imaging survey of $\\sim 50$  Virgo spiral galaxies 7 galaxies with long H{\\sc i} tails pointing away from M87. These galaxies are located at intermediate distances from the cluster center (0.6--1~Mpc). They conclude that these one-sided H{\\sc i} tail galaxies have recently arrived in the cluster, falling in on highly radial orbits. It appears that galaxies begin to lose their gas already at intermediate distances from the cluster center through ram-pressure or turbulent viscous stripping (Nulsen 1982) and tidal interactions with neighbours, or a combination of both. The outer gas disks of spiral galaxies at projected distances smaller than $\\sim1$~Mpc are thus affected by ram pressure. In this article we show that this is also the case for the massive spiral galaxy NGC~4501 (for its properties see table~\\ref{tab:parameters}).  NGC~4501 is located at the lower end of the distance interval defined by Chung et al. (2007). Moreover, Cayatte et al. (1990, 1994) and B\\\"{o}hringer et al. (1997) proposed that it approaches the center of the Virgo cluster (M87) on a highly eccentric orbit. NGC~4501 is moderately H{\\sc i} deficient and has an H{\\sc i} disk whose southwestern side is truncated well within optical radius ($R_{25}$). The H{\\sc i} surface density distribution within the disk is asymmetric with an overdense region to the southwest showing a sharp outer edge (Cayatte et al. 1990). Since the direction of these features is pointing toward the Virgo cluster center, Cayatte et al. (1990, 1994) and B\\\"{o}hringer et al. (1997) speculate that they are due to ram pressure compression (Gunn \\& Gott 1972). Since ram pressure only affects the interstellar medium of the galaxy and not its stars, one expects a symmetric stellar disk. However, the DSS B band morphology is asymmetric, the southwestern side of the major axis being brighter than the northeastern side, rising the question if NGC~4501 is exclusively experiencing an interaction with the intracluster medium and if a past gravitational interaction can be excluded.  In this article we investigate the nature of the interaction of NGC~4501 with the cluster environment using VIVA (VLA Imaging of Virgo galaxies in Atomic gas, Chung et al. in prep.) H{\\sc i} data, VLA 6~cm polarized radio emission data, and a detailed dynamical model including MHD calculations. Whereas the H{\\sc i} distribution and kinematics show a distorted gas disk where some of the gas is pushed to the north-east, the polarized radio-continuum emission represents an ideal tracer of ram pressure compression at the opposite side of the disk. The comparison between observations and simulations shows that NGC~4501 is indeed approaching the cluster center. Ram pressure stripping can account for the main properties of this Virgo spiral galaxy.  We describe the H{\\sc i} and 6~cm polarized radio-continuum observations and their results in Sect.~\\ref{sec:observations} and \\ref{sec:results}. Additional H$\\alpha$ data are discussed in Sect.~\\ref{sec:starform}. After a description of the dynamical model, the best fitting model and following MHD calculations are presented (Sect.~\\ref{sec:model}). Our models are compared to the observational data in Sect.~\\ref{sec:comparison} followed by a discussion (Sect.~\\ref{sec:discussion}) and our conclusions (Sect.~\\ref{sec:conclusions}).  \\begin{table} \\caption{Physical Parameters of NGC~4501} \\label{tab:parameters} \\[ \\begin{array}{lr} \\hline \\noalign{\\smallskip} {\\rm Other\\ names} &  {\\rm M~88} \\\\ & {\\rm VCC~1401} \\\\ & {\\rm UGC~7675}  \\\\ $$\\alpha$$\\ (2000)$$^{\\rm a}$$ &  12$$^{\\rm h}31^{\\rm m}59.2^{\\rm s}$$\\\\ $$\\delta$$\\ (2000)$$^{\\rm a}$$ &  14$$^{\\rm o}25'14''$$\\\\ {\\rm Morphological\\ type}$$^{\\rm a}$$ & {\\rm Sb} \\\\ {\\rm Distance\\ to\\ the\\ cluster\\ center} & 2.0$$^{\\rm o},\\ 0.6~{\\rm Mpc}$$\\\\ {\\rm Optical\\ diameter\\ D}_{25}$$^{\\rm a}$$ & 6.9$$'$$,\\ 34~{\\rm kpc} \\\\ {\\rm B}$$_{T}^{0}$$$$^{\\rm a}$$ & 9.86\\\\ {\\rm Systemic\\ heliocentric\\ velocity}$$^{\\rm a}$$\\ {\\rm (km\\,s}$$^{-1}$$)\\ & 2280$$\\pm$$4\\\\ {\\rm Velocity\\ with\\ respect\\ to\\ Virgo\\ mean}\\  {\\rm (km\\,s}$$^{-1}$$)\\ & 1130\\\\ {\\rm Distance\\ D\\ (Mpc)} & 17 \\\\ {\\rm Vrot}$$_{\\rm max}\\ {\\rm (km\\,s}$$^{-1}$$) & 300$$^{\\rm b}$$ \\\\ {\\rm PA} & 142$$^{\\rm o}$$\\ $$^{\\rm b}$$\\\\ {\\rm Inclination\\ angle} & 56$$^{\\rm o}$$\\ $$^{\\rm b}$$ \\\\ {\\rm HI\\ deficiency}^{\\rm c}$$ &  0.56$$\\pm$$0.2\\\\ \\noalign{\\smallskip} \\hline \\end{array} \\] \\begin{list}{}{} \\item[$^{\\rm{a}}$] RC3, de Vaucouleurs et al. (1991) \\item[$^{\\rm{b}}$] Guharthakurta et al. (1988) \\item[$^{\\rm{c}}$] Cayatte et al. (1994) \\end{list} \\end{table} ",
        "conclusions": "}  VIVA H{\\sc i} observations of the Virgo spiral galaxy NGC~4501 are presented. The outer edge of the H{\\sc i} distribution follows one optical isophot contour except in the northeast, where a region of low surface-density gas is discovered. Whereas the B band image is asymmetric, the lowest B band isophot contour and the H band isophot contours are symmetric. The observed southwestern ridge of 6~cm polarized radio-continuum emission (Vollmer et al. 2007) is located at the outer edge of the H{\\sc i} distribution. Both show the same rapid increase at the edge, which is expected in a ram pressure scenario. Along the gas ridge, a faint H$\\alpha$ emission ridge is present. The maximum of the polarized radio-continuum emission are located within the gap between the H$\\alpha$ emission of the spiral arms and the faint emission ridge.  We compare these observational characteristics with detailed dynamical models including MHD calculations. Models with different galaxy orbits (high and low ram pressure) and disk-wind angles between the disk and the orbital plane are presented. Model snapshots at different times from peak ram pressure are compared to the multi-frequency observations. For the first time we present model maps of massive star formation. We qualitatively reproduce the gas distribution, kinematics, star formation distribution, and polarized radio-continuum emission distribution. We conclude that ram pressure stripping can account for the main characteristic. The western ridge of enhanced gas surface density and enhanced polarized radio-continuum emission is due to ram pressure compression. Assuming radial orbits in a static intracluster medium of the Virgo cluster (Vollmer et al. 2001), NGC~4501 will have its nearest approach to the cluster core, i.e. peak ram pressure, in $\\sim 200-300$~Myr (Fig.~\\ref{fig:dessin1}). It is argued that the faint western H$\\alpha$ emission ridge is induced by nearly edge-on ram pressure stripping."
    },
    "0801/0801.0982_arXiv.txt": {
        "abstract": "Recent observations of high energy ($> 20$ keV) X-ray emission in a few clusters of galaxies broaden our knowledge of physical phenomena in the intracluster space. This emission is likely to be nonthermal, probably resulting from Compton scattering of relativistic electrons by the cosmic microwave background (CMB) radiation. Direct evidence for the presence of relativistic electrons in some $50$ clusters comes from measurements of extended radio emission in their central regions. We briefly review the main results from observations of extended regions of radio emission, and Faraday rotation measurements of background and cluster radio sources. The main focus of the review are searches for nonthermal X-ray emission conducted with past and currently operating satellites, which yielded appreciable evidence for nonthermal emission components in the spectra of a few clusters. This evidence is clearly not unequivocal, due to substantial observational and systematic uncertainties, in addition to virtually complete lack of spatial information. If indeed the emission has its origin in Compton scattering of relativistic electrons by the CMB, then the mean magnetic field strength and density of relativistic electrons in the cluster can be directly determined. Knowledge of these basic nonthermal quantities is valuable for the detailed description of processes in intracluster gas and for the origin of magnetic fields. ",
        "introduction": "Clusters of galaxies are the largest bound systems and the most important link to the large scale structure (LSS) of the Universe. The detailed properties of clusters, such as the distributions of their various mass constituents - dark matter, hot intracluster (IC) gas, and galaxies - dynamics, and thermal structure, are of much interest both intrinsically, and for the understanding of the formation and evolution of the LSS. Moreover, detailed {\\sl astrophysical} knowledge of clusters is essential for their use as {\\sl precise} cosmological probes to measure global parameters, such as $H_0$, $\\Omega_{\\mathrm M}$, \\& $\\Omega_{\\Lambda}$, and parameters characterising the primordial density fluctuation field.  Recent observations of many clusters of galaxies with the {\\sl Chandra} and {\\sl XMM} satellites at energies $\\epsilon \\leq 10$ keV, have significantly advanced our knowledge of the morphology and thermal structure of hot IC gas, the source of the cluster thermal Bremsstrahlung emission. The improved determinations of the gas temperature, density, and metal abundances from these observations significantly improve estimates of such important quantities as the total cluster mass and its gaseous and baryonic mass fractions.  As has been the case in galaxies, in clusters too a more physically complete understanding of these systems necessitates knowledge also of non-thermal (NT) quantities and phenomena in the IC space. Observational evidence for the relevance of these phenomena in clusters comes mostly from measurements of extended regions of radio emission, and from Faraday rotation (FR) of the plane of polarisation of radio sources seen through (or inside) clusters. Since the observed radio emission is clearly synchrotron-produced, its level and spectrum yield direct information on IC relativistic electrons and magnetic fields. Information on cluster magnetic fields (separately from relativistic electron properties) is obtained also from FR measurements. Compton scattering of cosmic microwave background (CMB) photons by the radio-emitting relativistic electrons boosts photon energies to the X-and-$\\gamma$ regions (e.g., \\citealt{rephaeli1979}). The search for cluster NT X-ray emission has begun long ago \\citep{rephaeligruber1987}, but first clear indications for emission at energies $\\epsilon \\geq 20$ keV came only after deep dedicated observations of a few clusters with the {\\sl{RXTE}} and {\\sl{BeppoSAX}} satellites (beginning with analyses of observations of the Coma cluster \\citep{rephaeli1999,fuscofemiano1999}. Radio and NT X-ray observations provide quantitative measures of very appreciable magnetic fields and relativistic electron densities in the observed clusters. These results open a new dimension in the study of clusters.  This is a review of cluster NT X-ray observations and their direct implications, including prospects for the detection of $\\gamma$-ray emission. The literature on cluster NT phenomena is (perhaps somewhat surprisingly) quite extensive, including several reviews of radio emission (e.g., \\citealt{govoni2004}) and cluster magnetic fields (e.g., \\citealt{carilli2002}), and a review of the current status of radio observations by \\citealt{ferrari2008} - Chapter 6, this volume. In order to properly address the comparison between magnetic field values deduced from radio observations and jointly from NT X-ray and radio measurements, we include here a brief summary of cluster radio observations. Measurements of EUV emission in a few clusters, and claims that this emission is by energetic electrons, are reviewed by \\citealt{durret2008} - Chapter 4, this volume. NT radiation processes are reviewed by \\citealt{petrosian2008a} - Chapter 10, this volume, and relevant aspects of particle acceleration mechanisms are reviewed by \\citealt{petrosian2008b} - Chapter 11, this volume. ",
        "conclusions": "It has been known all along that definite detection of NT X-ray emission in clusters is a challenging task. This is due to several factors, major among which are its intrinsically weak level - swamped as it is by the intense primary thermal emission - and the difficult to achieve requisite high sensitivity and low detector background in the $10-100$ keV band. As has been emphasised, results of the search are not unequivocal, even for Coma and A~2256 for which detection (by {\\sl{RXTE}} and {\\sl{BeppoSAX}}) was claimed to be at moderately high level of statistical significance. This is mainly due to source confusion and a complete lack of spatial information. The search should continue (with {\\sl{Suzaku}} and future satellites), spurred by its sound physical basis - the ubiquity of the CMB, and the observed radio synchrotron emission from relativistic electrons.  The expected level of NT emission of Compton origin depends steeply on the mean magnetic field strength in the central cluster region. Clearly, if this is typically a fraction of a $\\mu$G, as has been deduced in the analyses described in the previous section, then prospects for more definite detection of NT emission - perhaps already with deep {\\sl{Suzaku}} measurements - are indeed good. However, significantly higher $B_{\\mathrm{fr}}$ values - a few $\\mu$G - are deduced from FR measurements; had these been meaningful estimates of the {\\sl volume-averaged} field strength, then definite detection of NT emission would have been seriously questioned.  As we have noted already (in Sect.~2.2), the apparent discrepancy between deduced values of $B_{\\mathrm{rx}}$ and $B_{\\mathrm{fr}}$ has been investigated at some length. These two field measures are quite different; the former is essentially a volume average of the relativistic electron density and (roughly) the square of the field, the latter is an average along the line of sight of the product of the field and gas density. All these quantities vary considerably across the cluster; in addition, the field is very likely tangled, with a wide range of coherence scales which can only be roughly estimated. These make the determination of the field by both methods considerably uncertain. Thus, the unsatisfactory observational status (stemming mainly from lack of spatial information) and the intrinsic difference between $B_{\\mathrm{rx}}$ and $B_{\\mathrm{fr}}$, do not allow a simple comparison of these quantities. Even if the large observational and systematic uncertainties, the different spatial dependences of the fields, relativistic electron density, and thermal electron density, already imply that $B_{\\mathrm{rx}}$ and $B_{\\mathrm{fr}}$ will in general be quite different. This was specifically shown by \\citet{goldshmidt1993} in the context of reasonable assumptions for the field morphology, and the known range of IC gas density profiles. They concluded that $B_{\\mathrm{rx}}$ is indeed generally smaller than $B_{\\mathrm{fr}}$. The implication is that prospects for detection of NT emission should not be based on the relatively high deduced values of $B_{\\mathrm{fr}}$.  In conclusion, {\\sl{RXTE}} and {\\sl{BeppoSAX}} measurements have yielded appreciable evidence for power-law X-ray emission in four clusters. These results motivate further measurements and theoretical studies of NT phenomena on cluster and cosmological scales."
    },
    "0801/0801.2380_arXiv.txt": {
        "abstract": "We have conducted a systematic survey for intervening O\\,VI absorbers in available echelle spectra of 16 QSOs at $z_{\\rm QSO}=0.17-0.57$.  These spectra were obtained using HST/STIS with the E140M grating.  Our search uncovered a total of \\nabs\\ foreground \\OVI\\ absorbers with rest-frame absorption equivalent width $W_r(1031)\\gtrsim 25\\,\\mA$.  Ten of these QSOs exhibit strong O\\,VI absorbers in their vicinity.  Our O\\,VI survey does not require the known presence of \\lya, and the echelle resolution allows us to identify the O\\,VI absorption doublet based on their common line centroid and known flux ratio.  We estimate the total redshift survey path, $\\Delta\\,z$, using a series of Monte-Carlo simulations, and find that $\\Delta\\,z=1.66, 2.18$, and 2.42 for absorbers of strength $W_r=30, 50$ and $80$ m\\AA, respectively, leading to a number density of $d\\,{\\cal N}(W\\ge 50\\,\\mA)/dz = 6.7\\pm 1.7 $ and $d\\,{\\cal N}(W\\ge 30\\,\\mA)/dz = 10.4 \\pm 2.2$. In contrast, we also measure $d\\,{\\cal N}/dz=27\\pm 9$ for O\\,VI absorbers of $W_r>50\\ \\mA$ at $|\\Delta\\,v|< 5000$ \\kms\\ from the background QSOs.  Using the random sample of O\\,VI absorbers with well characterized survey completeness, we estimate a mean cosmological mass density of the O\\,VI gas $\\Omega({\\rm O}^{5+})\\,h = \\OmegaOVInum$.  In addition, we show that $<5$\\% of O\\,VI absorbers originate in underdense regions that do not show a significant trace of H\\,I. Furthermore, we show that the neutral gas column $N(\\hI)$ associated with these O\\,VI absorbers spans nearly five orders of magnitude, and shows moderate correlation with $N({\\rm O\\,VI})$. Finally, while the number density of O\\,VI absorbers varies substantially from one sightline to another, it also appears to be inversely correlated with the number density of \\hI\\ absorbers along individual lines of sight. ",
        "introduction": "Uncovering the missing baryons in the present-day universe is one of the central issues in observational cosmology, because an accurate accounting of the baryon budget at different locations sets an important constraint in modeling galaxy formation and evolution.  Over the past few years, measurements of the baryonic mass density from independent methods have approached a consistent value with high precision, $\\Omega_b\\,h^2 = 0.020 \\pm 0.002$ \\citep{burles-etal-01, spergel-etal-03, wang-etal-03, omeara-etal-06}.  But while $> 95$\\% of all baryons can be accounted for in the \\lya\\ forest at redshift $z\\sim 3$ \\citep{rauch-etal-97}, only $1/3$ have been identified at redshift $z=0$ in known components, such as stars and neutral gas in galaxies or hot plasma in galaxy clusters \\citep[e.g.][]{persic-salucci-92, fukugita-etal-98}.  Cosmological simulations indicate that $\\approx 30-40$\\% of all baryons at the present epoch reside in a diffuse warm-hot intergalactic medium (WHIM), with temperatures in the range $10^5 < T < 10^7$ K, rather than in virialized halos where galaxies or galaxy groups lie \\citep{cen-ostriker-99-missing-baryons,dave-etal-01-WHIM-baryons}.  The simulations also show that this WHIM is shock-heated as tenuous gas accretes onto large-scale filamentary structures and remains hot owing to the low gas density and inefficient cooling.  Locating the WHIM in order to verify these models is difficult, because the temperature of the gas is too low to be easily detected with the current generation of x-ray satellites and is too hot to be observed from the ground using optical-IR facilities.  Simple ionization modeling of hot gas at $T\\sim 3\\times 10^5$ K and metallicity 0.1 solar has shown that the gaseous clouds would imprint prominent absorption features produced by ionic transitions such as O\\,VI $\\lambda\\lambda$ 1031, 1037 and Ne\\,VIII $\\lambda\\lambda$ 770, 780 in the spectra of background QSOs \\citep[e.g.][]{verner-etal-94, mulchaey-etal-96}.  Surveys for the WHIM based on the presence of these ions at $z<1$ therefore require a high-resolution UV spectrograph in space.  \\citet{tripp-etal-00-H1821+643} searched for intervening O\\,VI absorbers along the sightline toward H1821$+$643 ($z_{\\rm QSO}=0.297$) in echelle spectra ($R\\equiv\\lambda/\\Delta\\lambda\\approx 43,000$) obtained using the Space Telescope Imaging Spectrograph (STIS) on board the Hubble Space Telescope (HST).  The search indicated that intervening O\\,VI absorption systems found in QSO spectra may indeed contain a significant fraction of baryons that had been missed by surveys based on the direct detection of photons emitted by known objects.  Subsequent searches along individual sightlines have been carried out toward PG0953$+$415 at $z_{\\rm QSO}=0.239$ \\citep{savage-etal-02-PG0953+415}, PG1259$+$593 at $z_{\\rm QSO}=0.478$ \\citep{richter-etal-04-PG1259+593}; PG1116$+$215 at $z_{\\rm QSO}=0.1765$ \\citep{sembach-etal-04-PG1116+215}, PKS0405$-$123 at $z_{\\rm QSO}=0.573$ \\citep{prochaska-etal-04-PKS0405-12}, and HE0226$-$4110 at $z_{\\rm QSO}=0.495$ \\citep{lehner-etal-06-HE0226-4110}.  The measured line density of intervening O\\,VI absorbers with rest-frame absorption equivalent width of the 1031 transition $W_r \\ge 50$ m\\AA\\ ranges from $d{\\cal N}({\\rm O\\,VI})/dz=11$ to 36\\footnote{\\citet{tripp-etal-00-H1821+643} reported $d{\\cal N}({\\rm O\\,VI})/dz\\sim 48$ based on four O\\,VI absorbers found with $W_r > 30$ m\\AA\\ toward H1821$+$643.  Three of the four absorbers have $W_r > 50$ m\\AA.  We therefore scale $d{\\cal N}({\\rm O\\,VI})/dz$ accordingly and work out $d{\\cal N}({\\rm O\\,VI})/dz=36$ for O\\,VI absorbers stronger than $W_r=50$ m\\AA.} per line of sight at $z_{\\rm O\\,VI}=[0,0.5]$.  The large scatter in the observed number density of O\\,VI absorbers suggests a significant variation in the spatial distribution of O$^{5+}$ ions across different sightlines.  It also underscores a potential systematic bias in evaluating the cosmological mass density of warm-hot gas based on selective quasar sightlines.  The first blind search of O\\,VI absorbers was carried out at $z_{\\rm O\\,VI}=[0.5,2.1]$ toward 11 background quasars by \\citet{burles-tytler-96-OVI-Omega_b} using medium resolution ($R\\approx 1300$) and $S/N\\approx 20$ Faint Object Spectrograph spectra.  They reported $d{\\cal N}({\\rm O\\,VI})/dz=1.0\\pm 0.6$ per line of sight at a mean redshift $\\langle z_{\\rm O\\,VI}\\rangle=0.9$ with $W_r \\ge 210$ m\\AA.  More recently, \\citet{danforth-etal-05-OVI-baryon-census} conducted a search of O\\,VI absorption features in {\\it known} \\lya\\ absorbers at $z_{\\rm abs}<0.15$, using Far Ultraviolet Spectroscopic Explorer (FUSE) spectra ($R\\approx 15,000$) of 31 active galactic nuclei (AGN).  Their search yielded $d{\\cal N}({\\rm O\\,VI})/dz=9\\pm 2$ and $17\\pm 3$ per line of sight with $W_r \\ge 50$ and $W_r\\ge 30$ m\\AA, respectively, at $z_{\\rm O\\,VI} < 0.15$, which is lower than measurements from individual lines of sight discussed earlier.  Despite the large scatter in the number density of intervening O\\,VI absorbers, these measurements imply that at least $\\sim 10\\%$ of the total baryons reside in the gaseous clouds probed by the O\\,VI transition.  Assuming a mean metallicity of 1/10 solar and an ionization fraction of 20\\% for the O\\,VI absorbers, Tripp \\etal\\ (2000) derived a lower limit to the cosmological mass density in the ionized gas of $\\Omega_b({\\rm O\\,VI})\\gtrsim 0.003\\,h^{-1}$ and Danforth \\& Shull (2005) derived $\\Omega_b({\\rm O\\,VI})\\approx 0.002\\,h^{-1}$.  This general inference of O\\,VI absorbers being a substantial reservoir of missing baryons is supported by numerical simulations that result in a mean metallicity of $1/10$ solar for the WHIM \\citep{cen-etal-01, fang-bryan-01-H1821+643_IGM}. However, additional uncertainties arise due to the unknown ionization state and metallicity of the gas.  We have conducted a systematic survey of intervening O\\,VI absorbers in available echelle spectra obtained using HST/STIS with the E140M grating.  A search of the HST data archive yielded high-resolution ($R\\approx 45,000$) spectra of 16 distant quasars at $z_{\\rm QSO}=0.17-0.57$ with a mean $S/N > 7$ per resolution element.  These quasars have been selected by different authors for different studies and have bright UV fluxes that are suitable for STIS echelle observations. Therefore, they represent a random, flux-limited sample of quasar spectra for an unbiased search of intervening absorbers.  The entire STIS echelle spectroscopic sample of quasars summarized in Table 1 defines the random sightlines along which our blind search of intervening O\\,VI absorbers was conducted.  The primary objectives of the blind O\\,VI survey are: (1) to obtain a statistically representative sample of O\\,VI absorbers along random lines of sight for an accurate measurement of the cosmological mass density of warm-hot gas at low redshift; and (2) to study the incidence of O\\,VI absorbers along individual sightlines for quantifying the spatial variation of O\\,VI absorbing gas. These O\\,VI absorbers identified along random sightlines also form a uniform sample for studying the physical nature of tenuous gas probed by the presence of O$^{5+}$ and for investigating their large-scale galactic environment.  Here we focus our discussion on the statistical properties of O\\,VI absorbers.  Studies of the gas properties and discussions of individual absorbers are given in an accompanying paper (Thom \\& Chen, 2008; hereafter paper~II), and studies of the surrounding galaxies will be presented in future papers.  This paper is organized as follows.  In Section 2, we summarize the available quasar spectra and the data quality.  In Section 3, we describe our approach for identifying an O\\,VI $\\lambda\\lambda$ 1031, 1037 absorption doublet with no prior knowledge of the presence of H\\,I or other ionic transitions.  A Monte-Carlo analysis was performed for evaluating the completeness of the survey. In Section 4, we present a statistical sample of O\\,VI absorbers identified along the sightlines toward the 16 background quasars.  We measure different statistical quantities of the ionized gas. In Section 5, we compare our search results with others and discuss the implications in Section 6. We adopt a $\\Lambda$ cosmology, $\\Omega_{\\rm M}=0.3$ and $\\Omega_\\Lambda = 0.7$, with a dimensionless Hubble constant $h = H_0/(100 \\ {\\rm km} \\ {\\rm s}^{-1}\\ {\\rm Mpc}^{-1})$ throughout the paper. ",
        "conclusions": "Summary of the Blind O\\,VI Survey and the Survey Limits} \\tablewidth{0pt} \\tablehead{ & \\colhead{} & \\colhead{$W_r^{\\rm min}$}  & \\colhead{Associated} & \\multicolumn{2}{c}{$\\Delta\\,z$ ($W\\ge W_*$)} & & \\colhead{\\dndz} \\\\ \\cline{5-6} \\cline{8-8} \\colhead{QSO} & \\colhead{$N_{\\mbox{abs}}^{\\rm random}$} & \\colhead{$(\\mA)$} & \\colhead{Absorbers} & \\colhead{$(W_* = 50\\,\\mA)$} & \\colhead{$(W_* = 30\\,\\mA)$} & & \\colhead{$(W_* = 30\\,\\mA)$} \\\\ \\colhead{(1)} & \\colhead{(2)} & \\colhead{(3)} & \\colhead{(4)} & \\colhead{(5)} & \\colhead{(6)} & & \\colhead{(7)} } \\startdata HE\\,0226$-$4110\\tablenotemark{1,2,3} \\dotfill & 4 &  43 & Y & 0.3055 & 0.2593 & & $13\\pm 7$ \\\\ PKS\\,0312$-$77                       \\dotfill & 1 & 655 & N & 0.0375 & 0.0014 & & $18\\pm 18$ \\\\ PKS\\,0405$-$12\\tablenotemark{4,5}    \\dotfill & 4 &  30 & N & 0.3542 & 0.2780 & & $12\\pm 6$ \\\\ HS\\,0624$+$6907\\tablenotemark{6}     \\dotfill & 2 &  27 & Y & 0.1681 & 0.1458 & & $6\\pm 6$ \\\\ PG\\,0953$+$415\\tablenotemark{7,8}    \\dotfill & 1 & 121 & Y & 0.0698 & 0.0434 & & $11\\pm 11$ \\\\ 3C\\,249.1                            \\dotfill & 1 &  72 & Y & 0.1206 & 0.0826 & & $8\\pm 8$ \\\\ PG\\,1116$+$215\\tablenotemark{9,10}   \\dotfill & 1 &  76 & Y & 0.0307 & 0.0053 & & $28\\pm 28$ \\\\ PG\\,1216$+$069                       \\dotfill & 1 &  26 & N & 0.1339 & 0.1047 & & $0\\pm 10$ \\\\ 3C\\,273\\tablenotemark{11}            \\dotfill & 1 &  22 & N & 0.0271 & 0.0195 & & $0\\pm 51$ \\\\ PG\\,1259$+$593\\tablenotemark{12}     \\dotfill & 2 &  77 & N & 0.2851 & 0.2680 & & $7\\pm 5$ \\\\ PKS\\,1302$-$102                      \\dotfill & 2 &  30 & N & 0.0992 & 0.0596 & & $26\\pm 19$ \\\\ PG\\,1444$+$407                       \\dotfill & 0 & ... & Y & 0.0927 & 0.0478 & & $0\\pm 21$ \\\\ 3C\\,351.0\\tablenotemark{13}          \\dotfill & 1 & 232 & Y & 0.1575 & 0.1313 & & $6\\pm 6$ \\\\ PHL\\,1811\\tablenotemark{14}          \\dotfill & 1 &  64 & Y & 0.0388 & 0.0133 & & $23\\pm 23$ \\\\ H\\,1821$+$643\\tablenotemark{15}      \\dotfill & 4 &  30 & Y & 0.1226 & 0.0999 & & $34\\pm 17$ \\\\ Ton\\,28                              \\dotfill & 1 &  26 & Y & 0.1396 & 0.0991 & & $0 \\pm 10$\\\\ \\hline total (16 sightlines)                         &27 &     & 10&  2.18  &  1.66  & & $10\\pm 2$       \\\\ \\enddata \\tablecomments{See also: (1) \\citet{lehner-etal-06-HE0226-4110}; (2) \\citet{savage-etal-05-HE0226-4110-NeVIII}; (3) \\citet{ganguly-etal-06-HE0226-4110_qso_host_lines}; (4) \\citet{chen-prochaska-00-PKS0405-12_z0.167}; (5) \\citet{prochaska-etal-04-PKS0405-12}; (6) \\citet{aracil-etal-06a-metals}; (7) \\citet{tripp-savage-00-PG0953+415_z0.1423}; (8) \\citet{savage-etal-02-PG0953+415}; (9) \\citet{sembach-etal-04-PG1116+215}; (10) \\citet{richter-etal-06-BLAs}; (11) \\citet{heap-etal-02}; (12) \\citet{richter-etal-04-PG1259+593}; (13) \\citet{yuan-etal-02-3C351-qso-lines}; (14) \\citet{jenkins-etal-03}; (15) \\citet{tripp-etal-00-H1821+643}.} \\end{deluxetable*}  We detect a total of \\nabs\\ random \\OVI\\ doublets along 16 QSO lines of sight over a redshift range from $z_{\\rm O\\,VI}=0.12$ to $z_{\\rm O\\,VI}=0.495$ with a mean redshift of $\\langle z\\rangle = 0.25$.  Furthermore, we uncover associated absorbers toward nine QSOs at $|\\Delta\\,v|< 5000$ \\kms\\ from the background QSOs.  The random O\\,VI absorbers form a statistically representative sample for constraining their global mean properties.  The associated O\\,VI absorbers form a uniform sample for studying the gaseous halos of distant QSO hosts.  Here we consider the statistical properties of the two samples separately.  Studies of the physical nature of individual absorbers will be presented in a separate paper (Thom \\& Chen 2008; Paper II).  \\subsection{The Incidence of Random O\\,VI Absorbers}  The Monte-Carlo simulations presented in \\S\\ 3.2 allow us to estimate the total redshift path length as a function of absorption equivalent width, from which we can then derive an accurate estimate of the number density of O\\,VI absorbers per line of sight, using individual sightlines separately or the entire QSO echelle sample collectively.  For each absorber of rest-frame O\\,VI 1031 absorption equivalent width $W_r^i>W_*$ identified in our search, the contribution to the total number density per line of sight $d{\\cal N}/dz$ is $1/g(W_r^i)$.  Over a statistical sample, we can therefore derive for a given equivalent width limit $W_*$, \\begin{equation} d\\,{\\cal N}(W \\ge W_*)/dz = \\sum_{i}\\,\\frac{1}{\\Delta\\,z (W_r^i)}\\,\\times\\,H(W_r^i-W_*) \\end{equation} where the sum extends over all systems found in a random sample and $H$ is the Heaviside step function.  Note that $\\Delta\\,z\\,(W_r^i)$ varies with $W_r^i$.  A stronger absorber may be detected in lower $S/N$ spectra, allowing a larger survey path length for a given QSO echelle sample. Equation (3) is therefore not the total number of absorbers, $N_{\\rm abs}^{\\rm random}$, divided by a fixed path length, $(\\Delta\\,z)$.  The variance of $\\dndz$ is evaluated according to \\begin{equation} \\sigma_{d{\\cal N}/dz}^2 = \\sum_{i}\\,\\frac{1}{[\\Delta\\,z (W_r^i)]^2}\\,\\times\\,H(W_r^i-W_*). \\end{equation}  We first evaluate \\dndz\\ separately for individual lines of sight, adopting a common threshold for all QSO spectra.  The results for $W_*=30$ \\mA\\ are summarized in column (7) of Table 2, and presented in panel (c) of Fig.~\\ref{fig: gz}.  Despite the relatively short path length per line of sight, it is immediately apparent from Fig.~\\ref{fig: gz} that there is a substantial variation in \\dndz\\ from one line-of-sight to another.  In contrast to the line-of-sight toward H\\,1821$+$643, along which four O\\,VI absorbers are found over $\\Delta\\,z\\lesssim 0.1$ \\citep{tripp-etal-00-H1821+643}, we note that the sightlines toward HS\\,0624$+$6907, PG1259$+$593, and 3C\\,351.0 allow longer survey paths but exhibit fewer absorbers (between 1 and 2).  The large scatter in \\dndz\\ between random sightlines underscores the necessity of a statistically significant sample of O\\,VI absorbers for an unbiased estimate of their incidence.  Combining all 16 QSO lines of sight together, we derive $d\\,{\\cal N}(W\\ge 50\\,\\mA)/dz = 6 \\pm 3$ over $\\Delta\\,z\\,(W_*=50\\,\\mA)=2.18$, and $d\\,{\\cal N}(W\\ge 30\\,\\mA)/dz = 9 \\pm 4$ over $\\Delta\\,z\\,(W_*=30\\,\\mA)=1.66$.  In \\S\\ 5 and panel (b) of Figure 5, we  show that our number density measurements can be explained well by recent numerical simulations that do not require super starburst outflow, but account for non-equilibrium ionization conditions \\citep{cen-fang-06}.   \\subsection{The Cosmological Mass Density of Ionized Gas}  A particularly interesting quantity to determine, based on a statistically representative O\\,VI absorber sample, is the cosmological mean mass density contained in the O$^{5+}$ ions, $\\Omega_{{\\rm O}^{5+}}$.  In principle, we can then derive the cosmological mean mass density of warm baryons, $\\Omega_{\\rm b}$, probed by the O\\,VI transition according to  \\begin{equation} \\Omega_{\\rm b}=\\frac{1}{f}\\,\\frac{\\mu}{Z\\times ({\\rm O}/{\\rm H})_\\odot}\\,\\frac{m_{\\rm H}}{m_{\\rm O}}\\,\\Omega_{{\\rm O}^{5+}}, \\end{equation} where $f\\equiv N({\\rm O}^{5+})/N({\\rm O})$ is the fiducial ionization fraction of the O\\,VI absorbing gas, $Z\\equiv ({\\rm O}/{\\rm H})/({\\rm O}/{\\rm H})_\\odot$ represents its mean metallicity, $({\\rm O}/{\\rm H})_\\odot$ is the solar oxygen to hydrogen ratio by number, $\\mu$ is the mean atomic mass of the gas ($\\mu=1.3$ for gas of solar composition), and $m_{\\rm H}$ and $m_{\\rm O}$ are the atomic weight of hydrogen and oxygen, respectively.  In the following calculation, we adopt $({\\rm O}/{\\rm H})_\\odot=4.57\\times 10^{-4}$ (Asplund \\etal\\ 2004).  Following \\citet{lanzetta-etal-01}, we calculate $\\Omega_{{\\rm O}^{5+}}$ according to \\begin{eqnarray} \\Omega_{{\\rm O}^{5+}} &=& \\frac{H_0\\,m_{\\rm O}}{c\\,\\rho_c} \\int N({\\rm O\\,VI})\\,f(N({\\rm O\\,VI}))\\,d\\,N({\\rm O\\,VI}) \\\\ &=& \\frac{H_0\\,m_{\\rm O}}{c\\,\\rho_c} \\sum_i\\,\\frac{N_i({\\rm O\\,VI})}{\\Delta\\,X_i} \\end{eqnarray} where $\\rho_c$ is the present-day critical density, $f(N({\\rm O\\,VI}))$ is the frequency distribution function defined as the number of O\\,VI absorbers per column density interval per unit absorption path-length, $N_i$ is the corresponding O\\,VI absorption column density for absorber $i$ with rest-frame absorption equivalent width $W_r^i$ and $\\Delta\\,X_i$ is the co-moving absorption path length covered by the entire echelle sample at the absorption strength $W_r^i$.  The co-moving absorption path length $\\Delta\\,X_i$ for absorber $i$ is related to the total redshift path length $\\Delta\\,z$ evaluated from Equation (2) at $W_r^i$ according to \\begin{equation} \\frac{d\\,X}{d\\,z} = \\frac{(1+z)^2}{[\\Omega_M\\,(1+z)^3 - (\\Omega_M + \\Omega_{\\Lambda} - 1)(1+z)^2 + \\Omega_{\\Lambda}]^{1/2}}. \\end{equation} The corresponding 1-$\\sigma$ error is evaluated according to \\begin{equation} \\sigma_{\\Omega_{{\\rm O}^{5+}}} = \\frac{H_0\\,m_{\\rm O}}{c\\,\\rho_c} \\left[\\sum_i\\,\\frac{\\sigma_{N_i({\\rm O\\,VI})}^2}{(\\Delta\\,X_i)^2}\\right]^{1/2}. \\end{equation} For each uncovered absorber of $W_r^i$, we perform a Voigt profile fit to determine the corresponding $N_i$.  We present the frequency distribution function determined based on the random O\\,VI absorber sample in panel (a) of Figure 5.  The bin size is selected to contain roughly the same number of absorbers per bin.  Error bars represent 1-$\\sigma$ Poisson fluctuations.  For the sample of 23 random O\\,VI absorbers with $W_r > 30\\mA$, we therefore derive $\\Omega({\\rm O}^{5+})\\,h = \\OmegaOVInum$. Assuming a $Z=1/10\\,Z_\\odot$ mean metallicity and $f=0.2$ ionization fraction for the O\\,VI gas, we further derive $\\Omega_{\\rm b}h\\approx 0.0014$.  Bearing in mind the large uncertainties in the metallicity and ionization fraction of the O$^{5+}$ bearing gas, our derived mean mass density of ionized gas is less than 10\\% of the total baryon mass density (e.g.\\ Burles \\etal\\ 2001).  \\subsection{Associated O\\,VI Absorbers}  \\begin{figure*} \\epsscale{1.0} \\plotone{f4.ps} \\caption{Absorption profiles of associated O\\,VI absorbers identified in our STIS echelle sample.} \\label{fig: qso} \\end{figure*}  We have excluded spectral regions within $|\\Delta\\,v|=0-5000$ \\kms\\ from the background QSOs for establishing the statistical sample of random foreground absorbers in the previous discussion, due to potential contamination features from the QSO host environment.  However, the blind O\\,VI search has been carried out all the way through the emission redshifts of these QSOs.  The O\\,VI absorbers found in the proximity of QSOs form a sample of associated absorbers that offers important constraints on the halo gas in QSO hosts (Chelouche \\etal\\ 2007, in prep).  Ten of the 16 QSOs in our STIS echelle sample exhibit strong O\\,VI absorbers $|\\Delta\\,v|\\le 5000$ \\kms\\ blue-shifted from their emission redshifts (indicated in Table 2).  We present the absorption profiles of these O\\,VI absorbers in Figure 4.  The rest-frame absorption equivalent width of the 1031 member ranges from $W_r=60\\ \\mA$ (the absorber at $\\Delta\\,v=-480$ \\kms\\ from Ton\\,28) to $W_r=3$ \\AA\\ (the multi-component absorber at $\\Delta\\,v\\approx -1600$ \\kms\\ from 3C\\,351.0).  The multi-component absorber found in the vicinity of 3C\\,351.0 has complex velocity profiles, extending from $\\Delta\\,v\\approx 0$ \\kms\\ to $\\Delta\\,v\\approx -2200$ \\kms.  It exhibits evidence of partial covering of the O$^{5+}$ ions and is understood to originate in the outflow gas from the background QSO \\citep{yuan-etal-02-3C351-qso-lines}.  Excluding the complex sightline toward 3C351.0, the total redshift path length defined within $|\\Delta\\,v|=0-5000$ \\kms\\ blue-ward of the emission redshifts of the remaining 15 QSOs in our echelle sample is $\\Delta\\,z=0.33$.  At the equivalent width limit of our search $W_*=50\\ \\mA$, we found nine O\\,VI absorbers.  We measure $d\\,{\\cal N}/dz=27\\pm 9$ for O\\,VI absorbers of $W_r(1031)>50\\ \\mA$ in the vicinity of background QSOs.  This is significantly higher than the incidence of random absorbers in the foreground.   \\label{sec: discussion}  \\subsection{The Ionization State of the O\\,VI Gas}   \\begin{figure}[t] \\epsscale{1.0} \\plotone{f5.eps} \\caption{(a) The frequency distribution function of O\\,VI absorbers from our random sample.  The bin size is chosen for each bin to contain roughly the same number of absorbers.  Error bars represents 1-$\\sigma$ Poisson fluctuations.  (b) Observed incidence of O\\,VI absorbers from different studies. Star points are from this work; open circles are from \\citet{danforth-etal-05-OVI-baryon-census} ; and open squares are from \\citet{savage-etal-02-PG0953+415, prochaska-etal-04-PKS0405-12, tripp-etal-06-conf, cooksey-etal-08}. Note that we have slightly offset the data points in the horizontal direction for clarification.  The curves are recent WHIM predictions based on numerical simulations \\citep{cen-fang-06} that account for non-equilibrium ionization conditions.  The top curve further includes Galactic wind feedback.} \\label{fig: dndz} \\end{figure}  Various studies of individual O\\,VI absorbers that consider both kinematic signatures of the absorption profiles and abundance ratios between different ions have concluded that a large fraction of these absorbers can be explained by a simple photo-ionization model \\citep[e.g.][]{tripp-etal-00-H1821+643, prochaska-etal-04-PKS0405-12, sembach-etal-04-PG1116+215, lehner-etal-06-HE0226-4110}.  These results are in contradiction to the expectation of a WHIM origin, under which the bulk of the O$^{5+}$ ions originate in the shocks created as the gas accretes onto the large filamentary structures.  However, it is often difficult to rule out a collisional ionization scenario for these absorbers.  Additional uncertainties arise, if O\\,VI absorbers arise in post-shock gas that is undergoing radiative cooling and a simple ionization equilibrium assumption does not apply (Gnat \\& Sternberg 2007).  The comparison of the absorption column densities in Figure 6 between \\lya\\ and O\\,VI transitions for the O\\,VI selected sample (solid points) demonstrates that the $N(\\hI)$ associated with these O\\,VI absorbers spans nearly five orders of magnitude.  In addition, stronger O\\,VI transitions have on average stronger associated \\lya\\ transitions.  The null hypothesis in which \\lya\\ and O\\,VI column densities are randomly distributed among each other can be ruled out at $>97$\\% confidence level based on a generalized Kendall test that considers the non-detection of \\lya\\ for the O\\,VI absorber at $z=0.32639$ toward HE\\,0226$-$4110.  We note that the STIS echelle sample is sensitive to O\\,VI absorbers of $W_r>80$ \\mA\\ over the entire redshift path-length allowed by the data, and O\\,VI absorbers of $W_r>50$ \\mA\\ over $\\approx 80$\\% of the sightlines.  The echelle spectra therefore allow us to uncover O\\,VI absorbers of $\\log\\,N({\\rm O\\,VI}) > 13.5$ over the majority of the sightlines.  The correlation is therefore unlikely due to a selection bias.  Including measurements from known O\\,VI absorbers at low redshift (open squares and crosses), we find that the correlation remains at $>80$\\% significance level.  The broad range of $N(\\hI)$ spanned by these O\\,VI absorbers also makes it difficult to apply a simple collisional ionization equilibrium model for explaining their origin.  Furthermore, the correlation between $N({\\rm O\\,VI})$ and $N(\\hI)$ indicates that the O\\,VI absorbers appears to track closely H\\,I gas in overdense regions, while the opposite does not apply \\citep[e.g.][]{danforth-etal-05-OVI-baryon-census}.  This trend can be understood under a simple photo-ionization scenario, and the slope in $N({\\rm O\\,VI})$ versus $N(\\hI)$ implies a declining ionization parameter (defined as the number of ionizing photons per gas particle) with increasing gas column density (see Figure 31 of Paper II).  \\subsection{O\\,VI Absorbers without Associated H\\,I}  A unique feature of our blind search is that the finding of O\\,VI doublet does not rely on known presence of \\lya\\ transitions.  The sample therefore allows us to constrain the fraction of O\\,VI absorbers that do not show significant traces of H\\,I.  These absorbers may represent an extreme case for outflow gas that has been ejected into underdense regions by strong galactic winds \\citep[e.g.][]{kawata-rauch-07}.  In Figure 6, we find that 26 of the 27 O\\,VI absorbers in the random sample show associated H\\,I absorption features.  The only one possible \\lya-free system is found to have $N({\\rm O\\,VI})=13.6\\pm 0.1$ at $z_{\\rm O\\,VI}=0.3264$ toward HE\\,0226$-$4110 (the arrow in Figure 6). The comparison indicates that $<5$\\% of O\\,VI absorbers originate in underdense regions of $N(\\hI)\\lesssim 10^{13}$ \\cmjj\\ that are heavily metal-enriched by starburst outflows  \\subsection{Large-scale Variation in Absorber Number Density}  \\begin{figure*} \\epsscale{0.6} \\plotone{f6.eps} \\caption{Comparison between $N({\\rm O\\,VI})$ and $N(\\hI)$ for the 27 O\\,VI absorbers identified in our blind search (solid points.  All absorbers but one have an associated \\lya\\ absorption line.  We have also included published absorbers at low redshifts.  Open squares are $z<0.12$ absorbers presented in \\citet{savage-etal-02-PG0953+415} for PG\\,0953$+$415, in \\citet{sembach-etal-04-PG1116+215} for PG\\,1116$+$215, in \\citet{richter-etal-04-PG1259+593} fo PG\\,1259$+$593, in \\citet{prochaska-etal-04-PKS0405-12} for PKS\\,0405$-$123, and in \\citet{cooksey-etal-08} for PKS\\,1302$-$102.  Crosses indicate measurements from a FUSE survey by \\citet{danforth-etal-05-OVI-baryon-census}.} \\label{fig: ncomp} \\end{figure*}  Cosmological simulations also show that O\\,VI absorbers originating in the WHIM reside in regions of over-density $\\delta\\equiv \\rho/\\langle\\rho\\rangle=10-100$ at $z\\sim 0$ \\citep{dave-etal-01-WHIM-baryons, cen-etal-01, fang-bryan-01-H1821+643_IGM}.  This is similar to what is found in simulation studies of \\lya\\ forest absorbers, according to which by $z=0$ the majority of \\lya\\ absorbers are associated with $\\delta = 10-100$.  It is therefore interesting to examine the relation between \\lya\\ and O\\,VI absorbers.  Of the 16 QSO sightlines in our echelle sample, seven have been studied extensively for the line-of-sight properties of \\lya\\ absorbers \\citep{sembach-etal-04-PG1116+215, williger-etal-06, lehner-etal-06-HE0226-4110}.  We present in Figure 7 a comparison of the observed number densities between \\lya\\ absorbers and O\\,VI absorbers measured along individual sightlines.  These include H\\,1821$+$643, HE\\,0226$-$4110, HS\\,0624$+$6907, PG\\,0953$+$415, PG\\,1116$+$215, PG\\,1259$+$593, and PKS\\,0405$-$123.  The statistics of \\lya\\ absorbers are taken from the compilation of \\citet{lehner-etal-06-HE0226-4110}.  We have included the corrected \\lya\\ absorber catalog for the PKS\\,0405$-$123 sightline.  The error bars represent 1-$\\sigma$ uncertainties.  Considering all seven sightlines together, we find a mean number density of $\\langle\\,d\\,{\\cal N}/dz\\rangle=66$ for the \\lya\\ absorbers with an r.m.s.\\ dispersion of 15.  In contrast, we find $\\langle\\,d\\,{\\cal N}/dz\\rangle=12$ for the O\\,VI absorbers with an r.m.s.\\ dispersion of 10. Examining the number density variation from one sightline to another, we find that the line of sight toward H\\,1821$+$643 with the highest incidence of O\\,VI absorbers has the lowest incidence of \\lya\\ absorbers.  Figure 7 shows a possible inverse correlation between the number density of \\hI\\ absorbers and the number density of O\\,VI absorbers.  While \\lya\\ absorbers appear to be relatively more abundant than O\\,VI along each sightline, the number density of O\\,VI absorbers do not scale proportionally with the number density of \\lya\\ absorbers along individual lines of sight.  This anti-correlation between $[d{\\cal N}/dz]_{\\rm O\\,VI}$ and $[d{\\cal N}/dz]_{\\rm \\lya}$ may be understood, if those sightlines with a higher incidence of O\\,VI absorbers probe a larger fraction of more highly ionized regions.  However, we note the result from \\S\\ 6.1, where we find nearly every one of the O\\,VI absorbers in our statistical sample shows an associated \\lya\\ transition.  On the other hand, if O\\,VI absorbers originate preferentially in underdense regions, then a higher incidence of O\\,VI implies a substantial fraction of the sightline probe underdense regions and hence a lower incidence of \\lya\\ absorbers.  This is expected if heavy elements from starburst outflows are ejected primarily into underdense regions \\citep[e.g.][]{kawata-rauch-07}.  Further insights can be obtained from a detailed examination of the large-scale galaxy environment of the O\\,VI absorbers.  A larger sample of QSO echelle spectra will also be valuable for a more quantitative study to confirm this result.  \\begin{figure*} \\epsscale{0.6} \\plotone{f7.ps} \\caption{Comparison of the line density variation from one line of sight to another for hydrogen \\lya\\ absorbers and O\\,VI absorbers. The number density of \\lya\\ absorbers measured for individual sightlines are taken from Lehner \\etal\\ (2007) for \\lya\\ absorbers of rest-frame equivalent width $W_r(\\lya) > 90 \\mA$ and the number density of O\\,VI absorbers measured for individual sightlines are from the blind search for O\\,VI absorbers of rest-frame equivalent width $W_r({\\rm O\\,VI}) > 30 \\mA$ presented in this paper.  We find a possible inverse correlation between the number density of \\hI\\ absorbers and the number density of O\\,VI absorbers.} \\label{fig: lden} \\end{figure*}"
    },
    "0801/0801.2039_arXiv.txt": {
        "abstract": "Dedicated to spectroscopic and imaging observations of the ultraviolet sky, the World Space Observatory for Ultraviolet Project is a Russia led international collaboration presently involving also China, Germany, Italy, Spain and Ukraine. The mission consists of a 1.7m telescope able to perform: a) high resolution (R$\\ge$60\\,000) spectroscopy by means of two echelle spectrographs covering the 103-310 nm range; b) long slit (1$\\times$75 arcsec) low resolution (R$\\sim$1500-2500) spectroscopy using a near-UV channel and a far-UV channel to cover the 102-310nm range; c) deep UV and diffraction limited UV and optical imaging (from 115 to 700 nm). Overall information on the project and on its science objectives are given by other two papers in these proceedings. Here we present the WSO-UV focal plane instruments, their status of implementation, and the expected performances. ",
        "introduction": "\\label{sec:1} The World Space Observatory-UV is an international collaboration led by Russia to build a space telescope dedicated mainly to UV astrophysics. The telescope is a Ritchey-Chretien with a diameter of 1.7 m, a F/10  focal ratio and a corrected field of view of 0.5 degrees. The primary wavelength range is 100-350 nm with an extension into the visible range. Specific data on the WSO-UV project (telescope, satellite, orbit, launcher, ground segment etc.) are given by Sachkov et al. (these proceedings). The WSO-UV telescope feeds in its focal plane a set of instruments for spectroscopy and imaging.  \\begin{table} \\centering \\caption{Characteristics of the WSO-UV spectrographs} \\begin{tabular}{lcccc} \\hline Spectrograph & Range & Spectral & Detector & Minimum\\\\ & nm & Resolution & Sensitivity$^a$\\\\ \\hline HIRDES-VUVES & 102--176 & $>$60\\,000 & MCP & 16 mag\\\\ HIRDES-UVES & 174--310 & $>$60,000 & MCP &18 mag\\\\ LSS & 102--320 & 1\\,500--2\\,500 & MCP & \\\\ \\hline \\multicolumn{5}{l}{SNR=10 in 10 h}\\\\ \\end{tabular} \\label{tab:1} \\end{table}  \\begin{description}  \\item [{\\bf HIRDES}:] The High Resolution Double Echelle Spectrograph (Funding Agency: DLR -- Principal Investigator: Prof. K. Werner, University of T\\\"ubingen, Germany -- Industrial Contractor: Kayser Threde), that delivers spectra at R$\\ge$60\\,000 in the 102-176 nm (VUVES), and in the 174-310 nm  (UVES) ranges. \\item [{\\bf LSS}:] The Long Slit Spectrograph (Funding Agency: CNSA - Project Director: Prof. Gang Zhao, National Astronomical Observatories of Chinese Academy of Science), that delivers spectra in the range 102-320 nm, having  spectral resolution R$\\sim$1500-2500, and allowing the use of a long slit of 1$\\times$75 arcsec. \\item [{\\bf FCU}:] The Field Camera Unit (Funding Agency: ASI -- Principal Investigator: Dr. I. Pagano, INAF, Catania Astrophysical Observatory -- Industrial contractors: Galileo Avionica and Thales Alenia Space IT-MI), that provides images, low resolution field spectroscopy, and polarimetry in three channels: FUV (115-190 nm), NUV (150-280 nm), and UVO (200-700 nm) with spatial resolution of 0.2, 0.03, and 0.07 arcsec pixel$^{-1}$ and field of view of 6.6$\\times$6.6, 1$\\times$1, and 4.7$\\times$4.7 arcmin$^{2}$, respectively. \\end{description}  In the next sections we describe the characteristics and performances of these science instruments.   \\begin{table} \\centering \\caption{Characteristics of the FCU imager} \\begin{tabular}{lccccccc} \\hline Channel & Range & Spatial & FoV &Detector & Array & Limiting\\\\ &  & Resolution$^a$ &  &  & Size&  Flux$^b$\\\\ & $nm$ &$arcsec$  & $arcmin^2$& & & erg\\,cm$^{-2}$\\,s$^{-1}$\\,\\AA$^{-1}$\\\\ \\hline FUV & 115--190 & 0.2~ & 6.6$\\times$6.6 & MCP(CsI) & 2k$\\times$2k & 5.5$\\times$10$^{-16}$\\\\ NUV & 150--280 & 0.06 & 1$\\times$1 & MCP(CsTe) &2k$\\times$2k & 5.0$\\times$10$^{-16}$\\\\ UVO & 200--700 & 0.07 &  4.7$\\times$4.7 & CCD$^c$ & 4k$\\times$4k & 1.8$\\times$10$^{-16}$\\\\ \\hline \\multicolumn{8}{l}{a) per pixel element}\\\\ \\multicolumn{8}{l}{b) SNR=10 in 1 h @[150, 255, 550]nm for [FUV,NUV,UVO], respectively}\\\\ \\multicolumn{8}{l}{c) UV Optimized}\\\\ \\end{tabular} \\label{tab:2} \\end{table} ",
        "conclusions": ""
    },
    "0801/0801.0819_arXiv.txt": {
        "abstract": "OH absorption against \\pulsar\\ at $(l,b) =351\\fdg688, +0\\fdg671$ has been discovered at 1665 and 1667\\,MHz using the Green Bank Telescope.  The absorption appears to arise at the interface of an \\hii\\ region and a molecular cloud which are likely associated with the high mass star forming region \\ngc. Beam dilution is found to be the cause of differences in the opacity of the OH against the Galactic background continuum emission and against the pulsar. The OH cloud is approximately 3 by 1.3\\,pc and is located behind the \\hii\\ region. ",
        "introduction": "\\hi\\ absorption measurements against pulsars have provided a wealth of information on the structure of the Interstellar Medium (ISM). They provide one of the best means of obtaining distance estimates to a large number of pulsars \\citep[e.g.][]{frail}.  Combining the distances with the pulsar's dispersion measure (DM) allows the average electron density along the line of sight to the pulsar to be determined. This also allows models of the Galactic electron density to be computed \\citep[e.g.][]{cordes,lazio}.  The \\hi\\ absorption measurements essentially provide the means by which the electron density model and the atomic neutral gas models of the Galaxy are tied together.  Comparison of the \\hi\\ emission when the pulsar is ``off'' with the \\hi\\ absorption against the pulsar can provide information on the density and spin temperature of the absorbing \\hi\\ gas \\citep[e.g.][]{balser}. \\hi\\ absorption measurements against pulsars have been used to search for structure in the neutral atomic gas on very small scales \\citep[e.g.][]{balser, tsas}. These measurements also play an important role in the study of pulsar population statistics, providing the spatial distribution of pulsars throughout the galaxy \\citep[e.g.][]{population}.  It is only natural to attempt to extend absorption measurements against pulsars to molecular gas.  Since pulsars have steep spectral indices, it is necessary to look for molecules which have transitions at lower frequencies and which have relatively strong line strengths. This makes the ${\\rm\\bf ^2\\pi_{3/2}(J=3/2)}$ transitions of the hydroxyl radical (OH) at 1612, 1665, 1667 and 1720\\,MHz ideal for looking for absorption from molecular material against pulsars.  Previously, only two pulsars have been found which have measurable OH absorption, \\arecibo\\ discovered by \\citet{snez} using the Arecibo radio telescope and \\parkes\\ discovered by \\citet{joel} using Parkes radio telescope.  Both \\citet{snez} and \\citet{joel} observed numerous bright pulsars, 25 in total.  In both cases, the OH absorption seen toward the pulsars have unique properties.  The OH absorption against the pulsars have narrower linewidths and deeper absorption (larger opacities) than is observed when the pulsar is ``off.'' This has been attributed to the presence of small scale structure in the molecular material along the lines of sight to the pulsars.  There are a large number of pulsars which have not been observed for OH absorption.  In this paper I present observations of sixteen pulsars (\\srca, \\srcb, \\srcc, \\srcd, \\srce, \\srcf, \\srcg, \\srch, \\srci, \\srcj, \\srck, \\srcl, \\srcm, \\srcn, \\srco, and \\pulsar ) looking for OH absorption.  OH absorption was only detected against \\pulsar. ",
        "conclusions": "\\pulsar\\ is only the third pulsar found to have OH absorption.  In all three cases the OH absorption arises in the interaction of a MC with a SNR or an \\hii\\ region. Forty-seven (47) pulsars have been searched for OH absorption.  This suggests that OH absorption against pulsars is either quite rare, or that the absorption in most MCs is very weak and below the detection capabilities of current telescopes without very deep searches.  That all known cases involve a MC that is in the process of being destroyed by a SNR or \\hii\\ region and that the detection rate is only $\\sim 6$\\% implies that OH absorption against pulsars is a rare occurrence.  The detection rate using pulsar brightness is about $\\sim 6$\\%.  I found a detection rate of 17\\% using scintillation strength for the selection criteria.  I also had a detection rate of 0\\% using detected OH emission that could be between the Sun and the pulsar.  Since the \\pulsar\\ OH detection was found in a search of six pulsars based on their strong scattering, and \\pulsar\\ is the 7th most strongly scattered pulsar, some credence can be given to the hypothesis of \\citet{boldyrev} but it by no means provides conclusive proof for their idea that scattering originates in PDRs at the boundaries of \\hii\\ regions and MCs.  In all likelihood, the continuum object along the line of sight toward \\pulsar\\ is an \\hii\\ region associated with the \\ngc\\ complex.  The spectral index of the source is consistent with what is expected for an \\hii\\ region. The morphology (see Figure~\\ref{fig:spitzer}) of cold dust surrounding warm dust is not consistent with the source being a SNR.  There is some H$_\\alpha$ emission associated with the continuum source (see Figure~\\ref{fig:spitzer}). Unfortunately, the SHASSA survey \\citep{shassa} does not provide velocity information for the H$_\\alpha$.  Although this line of sight was observed as part of the WHAM survey \\citep{wham}, the 1$^\\circ$ resolution convolves the emission from the \\hii\\ region toward \\pulsar\\ with other \\hii\\ regions in \\ngc. Thus it is not known if the H$_\\alpha$ arises at the same velocities as the OH absorption against \\pulsar.  Hydrogen radio recombination lines or higher spatial resolution H$_\\alpha$ with velocity information observations should be performed to confirm that the continuum source is an \\hii\\ region.  It was found that the OH cloud along the line of sight of \\pulsar\\ is likely behind the \\hii\\ region.  The OH cloud has a broad linewidth component and a narrow linewidth component.  Only the narrow linewidth component is seen in absorption against \\pulsar.  A similar situation of broad and narrow linewidth components with only the narrow component seen in absorption against the pulsar was observed for \\arecibo\\ \\citep{snez}. This is also the case for \\parkes\\ \\citep{joel}. The opacities of the OH absorption against \\pulsar\\ and against the continuum background (pulsar ``off'' spectrum) can be reconciled via consideration of the beam-filling factor of the OH cloud.  This should be investigated further for the \\arecibo\\ and \\parkes\\ OH absorption."
    },
    "0801/0801.0971.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {Surprisingly high amounts of $\\rm{H_2O}$ have recently been reported in the circumstellar envelope around the M-type asymptotic giant branch star W~Hya. This has lead to the speculation that evaporation of icy cometary or planetary bodies might be an effective ongoing mechanism in such systems. However, substantial uncertainties remain, as the required radiative transfer modelling is difficult due to high optical depths, sub-thermal excitation and the sensitivity to the combined radiation field from the central star and  dust grains.} %aims heading (mandatory) {By performing a detailed radiative transfer analysis, we determine fractional abundances of circumstellar $\\rm{H_2O}$ in the envelopes around six M-type asymptotic giant branch stars. The models are also used to predict $\\rm{H_2O}$ spectral line emission for the upcoming Herschel/HIFI mission.} %methods heading (mandatory) {We use Infrared Space Observatory Long Wavelength Spectrometer spectra to constrain the circumstellar fractional abundance distribution of ortho-$\\rm{H_2O}$, using a non-local thermal equilibrium, and non-local, radiative transfer code based on the accelerated lambda iteration formalism. The mass-loss rates and kinetic temperature structures for the sample stars are determined through radiative transfer modelling of CO line emission based on the Monte-Carlo method. The density and temperature profiles of the circumstellar dust grains are determined through spectral energy distribution modelling using the publicly available code Dusty.} %results heading (mandatory) {The determined ortho-$\\rm{H_2O}$ abundances lie between $2\\times10^{-4}$ and $1.5\\times10^{-3}$ relative to $\\rm{H_2}$, with the exception of WX~Psc, which has a much lower estimated ortho-$\\rm{H_2O}$ abundance of only $2\\times10^{-6}$, possibly indicating $\\rm{H_2O}$ adsorption onto dust grains or recent mass-loss-rate modulations. The estimated abundances are uncertain by, at best, a factor of a few.} % conclusions heading (optional), leave it empty if necessary {The high water abundance found for the majority of the sources suggests that either the `normal' chemical processes are very effective in producing $\\rm{H_2O}$, or else non-local thermal equilibrium atmospheric chemistry, grain surface reactions, or a release of $\\rm{H_2O}$ (e.g. from icy bodies like Kuiper belt objects) play a role. However, more detailed information on the physical structure and the velocity field of the region where the water vapour lines are formed is required to improve abundance estimates. We provide predictions for ortho-$\\rm{H_2O}$ lines in the spectral window of Herschel/HIFI. These spectrally resolved lines cover a wide range of excitation conditions and will provide  valuable additional information on the physical and chemical properties of the inner stellar wind where $\\rm{H_2O}$ is abundant.} ",
        "introduction": "\\label{intro}  Asymptotic giant branch stars (AGB stars) are chemically characterised according to their relative abundances of carbon and oxygen. Carbon stars have a photospheric C/O ratio $>$\\,1, whereas M-type (`oxygen rich') AGB stars have C/O $<$\\,1 (e.g., Russel~\\cite{russel}; Beck et al.~\\cite{becketal}). It is believed that M-type AGB stars evolve to carbon stars by the dredging up of nucleosynthesised carbon during the thermally pulsing AGB (TP-AGB) (e.g., Iben~\\cite{iben}; Boothroyd \\& Sackmann~\\cite{boothroydsackmann}). The evolution along the AGB is dominated by the mass loss, starting with a low mass-loss rate $\\sim$\\,$10^{-8}$\\,M$_{\\odot}$\\,yr$^{-1}$ and ending in an intense wind with rates up to $10^{-4}$\\,M$_{\\odot}$\\,yr$^{-1}$. The photosphere and expanding circumstellar envelope (CSE) are effective producers of a variety of molecular species and microscopic dust particles, providing up to 80\\% of the dust in galaxies (Whittet~\\cite{whittet}). When the star reaches the end of the AGB, most of the initial mass has been lost, leaving behind only a C/O core that ionises the surrounding material and, for a short while, creates a planetary nebula (Habing~\\cite{habing}). The lifetime on the AGB is estimated to be a few $\\rm{\\sim10^6\\,yr}$ (Marigo \\& Girardi~\\cite{marigoco07}).  To date, about 70 circumstellar molecular species have been detected, most of them through transitions at radio wavelengths (Olofsson~\\cite{olofsson07}). $\\rm{H_2O}$ is of special importance in the case of M-type AGB stars, as it is possibly the dominant coolant in the inner part of the stellar wind (where it is abundant) with a large number of far-infrared (FIR) transitions (Truong-Bach et al.~\\cite{truongbachetal}). $\\rm{H_2O}$ line emission is also a potentially important probe of the physical and chemical properties of the inner regions of CSEs. $\\rm{H_2O}$ is the main reservoir of circumstellar O, and whereas CO is abundant in all AGB stars, chemical models predict $\\rm{H_2O}$ to be one of the most abundant species only in O-rich stars (e.g., Cherchneff 2006\\nocite{cherchneff}). However, $\\rm{H_2O}$ has also been detected in the nearby C-rich AGB star IRC+10216 (Melnick et al.~\\cite{melnicketal01}; Hasegawa et al.~\\cite{hasegawaetal}), but here the estimated abundance is comparatively low, and the origin of $\\rm{H_2O}$ uncertain.  The atmosphere of the Earth is mostly opaque at $\\rm{H_2O}$ line wavelengths and it is therefore difficult to observe $\\rm{H_2O}$ from the ground, in particular in its ground vibrational state. However, Menten et al. (\\cite{mentenetal}) managed to observe two vibrationally excited lines in the $\\rm{\\nu_2\\,=\\,1}$ bending mode ($\\rm{5_{23}-6_{16}}$ at 336.2 GHz and $\\rm{6_{61}-7_{52}}$ at 293.6 GHz) towards the red supergiant VY CMa. $\\rm{H_2O}$ masers from the ground and excited vibrational states can also be detected from the ground (Menten \\& Melnick~\\cite{mentenmelnick}). Observations from the \\emph{Infrared Space Observatory} (ISO) are not hindered by the atmosphere, and the Short- and Long-Wavelength Spectrometer (SWS and LWS) data show a rich $\\rm{H_2O}$ spectrum in M-type AGB stars  ({e.g., Truong-Bach et al.~\\cite{truongbachetal}; Barlow et al.~\\cite{barlowetal}; Neufeld et al.~\\cite{neufeldetal}; Justtanont et al.~\\cite{justtanontetal}). The ortho-$\\rm{H_2O}$ ground state line at 557\\,GHz was also observed by the SWAS (Melnick et al.~\\cite{melnicketal}) and Odin satellites (Nordh et al.~\\cite{nordhetal}). In contrast to the CO radio lines, the $\\rm{H_2O}$ line emission originates from closer to the star, providing information on the warm, high-density innermost layers of the CSE (Truong-Bach et al.~\\cite{truongbachetal}). However, the high optical depths and the subthermal excitation of $\\rm{H_2O}$ in CSEs make it a challenge to accurately model the observed $\\rm{H_2O}$ spectral lines (e.g., Justtanont et al.~\\cite{justtanontetal}).  We present here the results of radiative transfer modelling of the ortho-$\\rm{H_2O}$ line emission in six M-type AGB stars with mass-loss rates between $\\rm{1\\times10^{-7}\\,M_{\\odot}\\,yr^{-1}}$ and $\\rm{4\\times10^{-5}\\,M_{\\odot}\\,yr^{-1}}$. Detailed radiative transfer models are created using the accelerated lambda iteration (ALI) method (Rybicki \\& Hummer~\\cite{rybickihummer91}, \\cite{rybickihummer92}). The model results are fit to LWS ISO observations between 43 and 197 $\\rm{\\mu m}$. The mass-loss rate and kinetic temperature structure are determined through CO radio line modelling, and the dust temperature and dust density structure are determined using Dusty (Ivezi\\'c et al.~\\cite{ivezicetal}). The results show the ability of ALI to model circumstellar $\\rm{H_2O}$ lines. We also make predictions for $\\rm{H_2O}$ lines in the frequency range of the upcoming Herschel/HIFI mission. These observations will include information on the line profiles and can be used to set firmer constraints on the physical structure of the inner CSE and the $\\rm{H_2O}$ abundance. Previous attempts to determine the water vapour content in stars included in our sample have been made for W~Hya (Barlow et al.~\\cite{barlowetal}; Zubko \\& Elitzur~\\cite{zubkoelitzur}; Justtanont et al.~\\cite{justtanontetal}) and R~Cas (Truong-Bach et al.~\\cite{truongbachetal}). Gonz\\'alez-Alfonso \\& Cernicharo (\\cite{gonzalezcernicharo}) used the large-velocity-gradient (LVG) method to model the 183 GHz maser line in O-rich evolved stars, to study the line intensity dependence on source properties and physical conditions. Ryde \\& Eriksson (\\cite{rydeeriksson}) model the $2.6-3.6\\,\\rm{\\mu m}$ spectrum for R~Dor observed with the SWS on ISO, showing the dominance of water vapour in the spectrum.  Determining the para-$\\rm{H_2O}$ abundance would be a straightforward task. However, considering the uncertainty in the abundance estimates, this would not give any firm constraints on the ortho-to-para ratio or the total $\\rm{H_2O}$ abundance.  In Sect.~\\ref{obs} the ISO observations are briefly described. Section~\\ref{models} contains a description of the dust and CO emission line modelling. In Sect.~\\ref{cslm} we describe the $\\rm{H_2O}$ line model, and in Sects.~\\ref{res} and~\\ref{disc} we present and discuss the results, respectively. Appendix~\\ref{ali} contains a description of the accelerated lambda technique, adopted in our radiative transfer model, in more detail and benchmark tests are presented. All abundances mentioned refer to the fractional abundance compared to $\\rm{H_2}$, unless stated otherwise.  %________________________________________________________________ %________________________________________________________________ ",
        "conclusions": "\\label{conc} % We have demonstrated the ability of our accelerated lambda iteration (ALI) code to model circumstellar $\\rm{H_2O}$ emission lines in M-type AGB stars at optical depths $>$\\,1000, and ortho-$\\rm{H_2O}$ abundances are estimated. The models include the ground and first excited vibrational states, collisions within the ground vibrational state, the first excited vibrational state, and between these states, and excitation by the stellar and dust radiation fields. The sample of stars modelled spans one order of magnitude in mass-loss rate. The ortho-$\\rm{H_2O}$ abundances determined, in the range (0.02--15)$\\times$10$^{-4}$, are estimated to be accurate to within 50\\%, within the adopted circumstellar model, while the absolute error (including model-dependent assumptions) is within a factor of 5. The estimated $\\rm{H_2O}$ abundances are high in comparison to what can be produced by chemical models, but we explain that the estimates are rather uncertain.  The outer radius of the $\\rm{H_2O}$ envelope is better determined for the low-mass-loss-rate objects and corresponds to within 1$\\sigma$ to the radius given by the theoretical estimate in Eq.~\\ref{outrad}.  Including excitation of the ground state through the vibrationally excited state affects the low-mass-loss-rate objects, increasing the model line fluxes by $\\approx\\,40\\%$. This results in a lower abundance for these objects compared to models that do not include the excited vibrational state (Justtanont et al.~\\cite{justtanontetal}). For higher mass-loss-rate objects the vibrational excitation is less important.  Several formation processes for the observed $\\rm{H_2O}$ abundances are possible. LTE or NLTE models may explain the formation of $\\rm{H_2O}$ in the innermost layers of the CSE, whereas Fischer-Tropsch catalysis or the evaporation of Kuiper-belt like objects are possible sources of $\\rm{H_2O}$ in the outer parts of the envelope. Future observations with the Herschel telescope will help to differentiate between these scenarios. The HIFI lines are generally less saturated and will therefor give a better estimate of the circumstellar $\\rm{H_2O}$ abundance. In particular, optically thin lines, such as the $8_{45}-7_{52}$ transition, are sensitive to the $\\rm{H_2O}$ abundance, and will help constrain this parameter. The ground state line at 557 GHz, on the other hand, is excited throughout the envelope and will help in setting constraints on the size of the $\\rm{H_2O}$ envelope.  For WX~Psc, we suggest that sticking of $\\rm{H_2O}$ onto dust grains may explain the observed low amount of $\\rm{H_2O}$. Alternatively, the present day mass-loss rate is low."
    },
    "0801/0801.0358_arXiv.txt": {
        "abstract": "\\noindent The thermal evolution of neutron stars is coupled to their spin down and the resulting changes in structure and chemical composition. This coupling correlates stellar surface temperatures with rotational state as well as time. We report an extensive investigation of the coupling between spin down and cooling for hybrid stars which undergo a phase transition to deconfined quark matter at the high densities present in stars at low rotation frequencies. The thermal balance of neutron stars is re-analyzed to incorporate phase transitions and the related latent heat self-consistently, and numerical calculations are undertaken to simultaneously evolve the stellar structure and temperature distribution. We find that the changes in stellar structure and chemical composition with the introduction of a pure quark matter phase in the core delay the cooling and produce a period of increasing surface temperature for strongly superfluid stars of strong and intermediate magnetic field strength. The latent heat of deconfinement is found to reinforce this signature if quark matter is superfluid and it can dominate the thermal balance during the formation of a pure quark matter core. At other times it is less important and does not significantly change the thermal evolution. ",
        "introduction": "The chemical composition of neutron stars at densities beyond nuclear saturation remains uncertain with alternatives ranging from purely nucleonic compositions through hyperon or meson condensates to deconfined quark matter -- see e.g. \\cite{Weber:2005} and \\cite{Page:2006} for recent reviews with emphasis on quark matter. A future understanding of neutron star structure gained through confrontation of theoretical models with the now steadily growing body of observational facts will therefore simultaneously constrain fundamental elements of particle and nuclear physics \\citep{Lattimer:2007}. The interlinked processes of spin down and thermal cooling present intriguing prospects of gaining insight in the properties of matter in neutron star cores by confrontation with soft X-ray observations of thermal radiation from neutron star surfaces as they both depend sensitively on and to some extent determine the chemical composition.  As neutron stars spin down and contract, their structure and composition change with the increasing density -- drastically if phase boundaries are crossed so new forms of matter become possible. We shall here investigate how this influences the thermal evolution of hybrid stars which contain large amounts of deconfined quark matter. The increasing density and changing chemical composition further imply additional entropy production in bulk and the release of latent heat as particles cross any phase boundaries present. We therefore re-analyze the thermal equilibrium of compact stars to show how mixed phases may be incorporated. We thus arrive at a natural description of the latent heat of phase transitions in compact stars, but also find -- through direct numerical calculations -- that unless the stellar structure changes very rapidly the effects of latent heat on the thermal evolution are insignificant when compared to those of the changing chemical composition and the surface area reduction.\\\\   Neutron stars are extremely compact objects and densities in their cores reach well in excess of the nuclear saturation density. At such densities the distance between particles is on the order of the characteristic range of the nuclear forces. Therefore, as was stressed recently by \\cite{Baym:2006}, perturbative treatments in terms of few- or even many-body forces -- although highly successful in describing the properties of matter below nuclear saturation -- are no longer well defined in the core of neutron stars. Further, the relevant degrees of freedom should include the appearance of hyperons and possibly deconfined quarks, so treatments in terms of nucleons alone must also be seen as approximative. Hyperons are expected to appear at densities around $2-3\\rho_0$, where $\\rho_0 = 0.153$ fm$^{-3}$ is the baryon density at nuclear saturation. The appearance of hyperons so softens the equation of state that purely hadronic equations of state may not allow stable models compatible with the accurately measured masses of neutron stars in binary systems \\citep{Baldo:2003,Schulze:2006}. \\cite{Schulze:2006} further demonstrate that this conclusion is highly robust with respect to different assumptions about hyperon interactions and the nucleonic equation of state. At best these studies are strong arguments against purely hadronic compositions and they certainly do show the relevance of considering alternatives such as a transition at high densities to deconfined quark matter.\\\\  Quark matter represents an entirely new type of matter -- as opposed to just an additional degree of freedom as in the case of hyperons in the hadronic phase -- and it cannot be assumed to soften the equation of state to the same extent \\citep{Alford:2007a,Schaffner-Bielich:2007}. Hybrid stars can be be consistent with the high masses and radii indicated by recent observations (e.g. \\cite{Ozel:2006}, \\cite{Freire:2008,Freire:2008a,Freire:2007a}), and the quark matter equation of state can fulfill the constraints imposed by heavy-ion collision transverse flow data and $K^+$ production (which nucleonic equations of state do not). Further, \\cite{Drago:2008} found that only stars with a quark matter component can rotate stably without losing angular momentum by emission of gravitational waves through the r-mode instability at the 1122 Hz indicated by recent observations for the X-ray transient XTE J1739-285 \\citep{Kaaret:2007}. This conclusion is tentative as the observed neutron star spin frequency awaits confirmation, but it again shows how the composition of neutron stars is an open question to be determined by a confrontation of theory and observation, and that a deconfined quark matter phase remains a viable alternative.\\\\  For definiteness we shall work with the equation of state suggested by \\cite{Glendenning:1992}; hereafter the G$^{300}_{180}$ equation of state. While this equation of state is not sophisticated in its treatment of the quark matter phase, it is illustrative in that it allows a very large pure quark matter core with a transition through a mixed phase at relatively low densities around 2-5$\\rho_0$. For our purposes it is interesting in that it implies drastic changes in composition with spin down -- the pure quark matter core disappears at high spin frequencies for instance -- and it therefore represents something close to a limiting case. We shall then be able to investigate both the effects of steady conversion of hadronic matter to quark matter and the sudden appearance of a pure quark matter core in the hadronic phase.  In the so-called minimal scenario, which excludes exotic and very rapid neutrino processes \\citep{Page:2004}, the internal temperature of neutron stars drops from beyond $10^{10}$~K to around $10^9$~K within a few minutes after their birth, and neutrino cooling continues to dominate for at least the next few thousand years until the internal temperature has dropped below $10^8$~K and photon cooling takes over (see, e.g. \\cite{Page:2004, Yakovlev:2004a} for reviews). If the highly efficient direct Urca neutrino emission (i.e., essentially beta decays, $n\\rightarrow p+e^-+\\bar{\\nu}_e$ and related processes in quark matter) is active, the stars may cool very rapidly and reach very low temperatures on a timescale of a few hundred years. As we shall see the nuclear direct Urca process is active in the mixed phase of hybrid stars, and may thus control the thermal evolution. These processes may be suppressed by pairing of the participating particles however, and further the extent of the mixed phase depends strongly on the rotation frequency thus giving rise to a diverse range of possible cooling paths.  The latent heat of the phase transition and the related release of entropy in bulk with changing density are generally found to be of smaller importance than the changing structure and chemical composition. It does not significantly delay or enhance the cooling, although it does briefly balance or even dominate the cooling terms when the pure quark matter core is first formed for certain choices of stellar parameters. The latent heat of the quark matter phase transition was previously considered by \\cite{Miao:2007, Miao:2007a} and, \\cite{Xiaoping:2008}. These authors took a different approach to calculate the latent heat than what is discussed below and found very significant heating terms which we do not recover.  The work of \\cite{Reisenegger:1995} and lately \\cite{Fernandez:2005} considered the related process of roto-chemical heating for neutron stars in which weak reactions are driven out of equilibrium by the changing density. The resulting release of energy was found able to maintain old stars at relatively high temperatures determined by their rate of spin down. This term naturally appears in our equations for the thermal balance and should be included in a full model. It is to some extent complementary to the effects discussed here, but the treatment of weak reactions beyond equilibrium is beyond our scope in this work, and we shall assume chemical equilibrium throughout. We return briefly to discuss this issue in Sect.~\\ref{discuss}. \\\\  The link between the thermal evolution of neutron stars and their spin down has become a subject of some interest as an observational correlation between the temperatures and inferred magnetic fields of neutron stars was recently discovered by \\cite{Pons:2007}. This was interpreted as evidence for magnetic field decay and detailed work by \\cite{Aguilera:2007, Aguilera:2008} strongly supports this conclusion. An alternative interpretation was offered recently by \\cite{Niebergal:2007} in terms of magnetic flux expulsion from color-flavor locked quark stars (a hypothetical class of stars consisting entirely of quark matter which in this case is assumed to be absolutely stable). The previously mentioned work of \\cite{Drago:2008} also indicates a link between the spin down and cooling of hybrid stars. These correlations -- if they can indeed be shown to exist -- complement the traditional cooling calculations which relate only temperature and age, and they may help break the degeneracy seen between such calculations with different underlying assumptions about the state of matter at very high density.\\\\  In the following, we shall first revisit the equations of thermal balance for compact stars in Sect.~\\ref{thequilibrium} to show the effects of a time-dependent density and discuss how the presence of phase transitions may be included. In Sect.~\\ref{eosandmodel}, we discuss the G$_{180}^{300}$ equation of state, the resulting stellar models and some simple estimates for the additional terms in the thermal balance equations in more detail. In Sect.~\\ref{calculations} we show the results of including these terms in spherical isothermal cooling calculations. We conclude with a discussion in Sect.~\\ref{discuss}. ",
        "conclusions": "\\label{discuss} Our intention with the present work was to explore a possible connection between the spin down and thermal cooling of hybrid stars with a deconfined quark matter core -- the appearance of which might be expected to influence theoretical cooling tracks. Specifically we were interested in the latent heat of the phase transition and the drastic changes in structure and chemical composition resulting from the increasing density with spin down. The general formalism worked out in Sect.~\\ref{thequilibrium} would also be relevant to models with no phase transition, however, as the bulk terms in the spin down derived entropy production might also be important in such cases. While estimates of the importance of the spin down derived entropy production did not give cause to expect strong signals of deconfinement in the thermal evolution, numerical calculations revealed clear effects of the changes caused by the appearance of a pure quark matter core. This signature is related to changes in radius, chemical composition, structure and under certain circumstances the additional entropy production both in bulk and in the form of latent heat.\\\\  Our numerical work was carried out using a sophisticated and reliable code with respect to the stellar structure solving Hartles perturbative equations for the structure of rotating compact stars self-consistently. The thermal evolution, however, was treated in a spherical isothermal approximation with a somewhat ad hoc approach to the treatment of superfluid pairing in the quark matter phase. Still we believe our work is sufficiently detailed to demonstrate the general point that important signals may be derived from the interplay between spin down and cooling of compact stars, which could in the long run help answer pressing questions about the state of matter at high densities. The signals we have found do not dominate the general thermal evolution of hybrid stars but they do complement the standard picture. By correlating temperature with magnetic field strength and spin frequency they may also help break the degeneracy between models relating only temperature and time.\\\\   The signature of deconfinement found here is below the present observational sensitivity and not of sufficient strength to set apart the cooling curves with temperature versus time for hybrid stars. The correlation between temperature and magnetic field strength provides a possible alternative. To test such a correlation it would be necessary to obtain accurate temperatures for a number of stars of approximately the same age for which the spin down could also be detected by the emission of pulses in either radio or X-ray. It might then be possible to test for features in the relation between temperature and magnetic field -- or equally interesting, the rotation frequency itself -- which as demonstrated here could result from a strong phase transition if the initial rotation frequency was sufficiently high to have spun out the high density phase.  Contemplating these prospects it is important to consider that alternative effects not treated here might determine the relation between spin down and thermal evolution. Most pressingly the heat from magnetic field decay was employed by \\cite{Pons:2007} to explain the observed correlation between temperature and magnetic field and it may well drown out any other effects -- although as noted by the authors the magnetic field decay itself has not yet been independently demonstrated. Recent work by \\cite{Aguilera:2007} and \\cite{Aguilera:2008} on the cooling of magnetized neutron stars in two dimensions also highlighted the importance of the magnetic field strength, geometry and possible decay for the thermal evolution of neutron stars. These authors find strongly anisotropic surface temperature distributions and possibly an inverted temperature distribution with hot equatorial regions for middle aged stars. They further show that the effects of magnetic field decay and Joule heating can dominate the thermal evolution at strong and intermediate field strengths -- even to the point that the effects of the direct Urca process may be hidden by this heating term in magnetars.   The work of \\cite{Reisenegger:1995} and \\cite{Fernandez:2005} is also most important. These authors consider the second term on the right hand side of Eq.~(\\ref{balance2}) which is shown to give rise to so-called roto-chemical heating as the weak interactions required to maintain chemical equilibrium are unable to keep pace with the increasing density. \\cite{Fernandez:2005} found that old millisecond pulsars reach a quasi-equilibrium in which the photon luminosity is determined entirely by the spin down power and remains much higher than otherwise attainable. Their work considered only nucleonic equations of state but similar results should clearly be expected for hybrid stars just as part of the effects demonstrated here should be expected to show up in nucleonic stars.\\\\   Given the possible importance of such alternatives and the shortcomings of the present study discussed above, a more complete treatment of the interplay between the spin down and thermal evolution of neutron stars seems desirable before strong assertions can be made concerning the specific shape of cooling tracks. As well as including all possible energy sources and sinks this should be extended to consider a wider range of stellar masses, equations of state, phase transitions and pairing regimes. A two dimensional code could further treat the influence of the magnetic field on heat flows in the inner crust and the effects of nonspherical heat flows for rotationally perturbed stars -- aspects which are all independently well described in the literature.   It should further be noted that the rotation frequency at which the pure quark matter core appears depends strongly on the stellar baryon number and the equation of state, and it may well be lower than discussed here thus changing the timing of spin down related effects. The strongest signature of deconfinement found here depends entirely on the formation of a pure quark matter core where none existed previously. If, on the other hand, a pure quark matter phase is present in the core of stars at even the highest rotation frequencies -- as is the case for other equations of state -- we would not expect equally strong signals from deconfinement with spin down compression.  It is possible that circumstances not considered here could provide a more pronounced link between rotation and cooling. If for instance the magnetic field varies over time -- and as discussed by \\cite{Geppert:2006} this may well be the case -- the timing of any strengthening or decay of the magnetic field may provide entirely different spin down histories. Alternatively the temperature of stars being spun up in accreting binaries should to some extent be determined by the changing chemical composition and this could provide an alternative approach.  Additional energy release would also be expected if the phase transition cannot maintain thermodynamic or hydrostatic equilibrium as was assumed here. As the most extreme example of this, we have entirely ignored the instabilities and corequakes which might accompany the appearance of a pure quark matter phase in a star far from equilibrium \\citep{Zdunik:2006}. Such events could release large amounts of energy and significantly change the thermal evolution by reheating the star.\\\\  While many essential issues thus remain poorly explored we hope to have demonstrated the possible usefulness of considering the connection between the spin down and thermal evolution of neutron stars. Although the results of such work may not be immediately applicable to observational tests, we believe they could prove most valuable in the long run.\\\\   We wish to acknowledge helpful conversation with Sanjay Reddy on the hadronic pair breaking and formation process and to thank Dany Page for the use of his compilation of observational data. We are also grateful to Dima Yakovlev and Oleg Gnedin for providing us with an earlier version of their cooling code, which we took advantage of when developing the code used for this study.\\\\ M. Stejner wishes to thank San Diego State University for its hospitality during part of this work.\\\\ The research of F. Weber is supported by the National Science Foundation under Grant PHY-0457329, and by the Research Corporation.\\\\ J. Madsen is supported by the Danish Natural Science Research Council.  \\appendix"
    },
    "0801/0801.1442_arXiv.txt": {
        "abstract": "This is a brief review on the history of the Bose-Einstein condensate (BEC) or boson star model of galactic dark matter halos, where ultra-light scalar dark matter particles condense in a single BEC quantum state. The halos can be described as a self-gravitating, possibly self-interacting, coherent scalar field. On a scale larger than galaxies, dark matter behaves like cold dark matter while below that scale the quantum mechanical nature suppresses the dark matter structure formation due to the minimum length scale determined by the mass  $m\\st{>}{\\sim}10^{-24} eV$ and the self-interaction of  the particles. This property could alleviate the cusp problem and missing satellite problems  of the $\\Lambda$CDM model. Furthermore, this model well  reproduces the observed rotation curves of spiral and dwarf galaxies, which makes the model promising. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.0742.txt": {
        "abstract": "If inflation is to be considered in an unbiased way, as possibly originating from one of a wide range of underlying theories, then observations need not be simply applied to reconstructing the inflaton potential, $V(\\phi)$, or a specific kinetic term, as in DBI inflation, but rather to reconstruct the inflationary action in its entirety.  We discuss the constraints that can be placed on a general single field action from measurements of the primordial scalar and tensor fluctuation power spectra and non-Gaussianities. We also present the flow equation formalism for reconstructing a general inflationary Lagrangian, $\\L(X,\\phi)$, with $X=\\frac{1}{2}\\partial_\\mu\\phi\\partial^\\mu\\phi$, in a general gauge, that reduces to canonical and DBI inflation in the specific gauge $\\Lx = c_s^{-1}$. ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% The Cosmic Microwave Background (CMB) \\citep{Page:2006hz,Hinshaw:2006ia,Jarosik:2006ib,Kuo:2006ya} is now measured with exquisite precision from horizon scales down to a few arc minutes angular resolution. In combination with large scale structure surveys \\cite{Tegmark:2003uf,Cole:2005sx,Tegmark:2006az}, this allows the primordial spectrum of fluctuations to be characterized in fine detail \\cite{Spergel:2003cb,Peiris:2003ff,Bridle:2003sa,Spergel:2006hy}.  There has been significant effort to relate the observed primordial spectrum of fluctuations to the underlying theory that seeded them. In the context of slow roll inflation \\citep{Guth:1980zm,Linde:1981mu,Albrecht:1982wi}, taking a specific potential and comparing it to data is a sensible approach to assess if the theory is consistent (see for example \\citep{Spergel:2003cb,Kinney:2006qm}). An alternative application of the data, however, is to invert this process, and \u00d4reconstruct\u00d5 what we can know about the underlying theory \\citep{Copeland:1993ie,Copeland:1993zn,Adams:1994qp,Turner:1995ge,Hoffman:2000ue, Easther:2002rw,Kinney:2003uw,Peiris:2006ug,Boyanovsky:2007ry, Powell:2007gu,Lesgourgues:2007gp,Cline:2006db,Ballesteros:2005eg,Cortes:2006ap, Chung:2003iu,Chung:2005hn,Spalinski:2007ef,Ballesteros:2007te,Destri:2007pv}.  With the introduction of a broader array of inflationary theories, for example arising out of the Dirac-Born-Infeld action \\citep{Dvali:1998pa} or from k-inflation \\citep{ArmendarizPicon:1999rj, Garriga:1999vw}, the interpretation of observations has necessarily extended beyond focusing only on the inflationary potential to including information about the form of the Lagrangian kinetic term.  Recently, cosmological constraints on brane inflation models have been studied both in the context of specific models \\citep{Bean:2007hc,Easson:2007dh,Bean:2007eh} and in more model-independent studies \\citep{Peiris:2007gz}.  If inflation is to be considered without theoretical bias, then the objective must not be to simply reconstruct the inflaton potential, or a specific kinetic term, but rather to reconstruct what observations tell us quantitatively about the effective inflaton action in its entirety. In this paper, we develop a formalism for such a general inflationary reconstruction in the context of single field models, and present explicit analytic techniques for action reconstruction.  In the usual potential reconstruction formalism, the inflationary observables as a function of scale can be mapped to the behavior of the inflationary potential $V(\\phi)$ as a function of the inflaton field $\\phi$.  If for example the scalar spectral index $n_s(k)$ can be extracted exactly from data, the shape of the potential $V(\\phi)$ can be deduced in the usual formalism if a reheating scenario is fixed.  However, the reconstruction of the entire action, including the possibility of non-minimal kinetic terms, is harder as the action is now a functional of two independent functions, $X\\equiv (\\partial \\phi)^2/2$ and $\\phi$.  To fix this new functional degree of freedom, in principle a continuous set of independent data analogous to $n_s(k)$ is needed (i.e. an infinite number of observables).  Although daunting, the formalism that we present may be used as a starting point to connect cosmological data to high energy theories which may have other phenomenological, theoretical, and aesthetic constraints.  The analytic form of the non-minimal kinetic actions consistent with data can be written in a surprisingly simple form given in section \\ref{srrecon} by Eq.~(\\ref{generalaction}).  Each consistent action is simply a manifold parameterized by $X$ and $\\phi$ satisfying certain derivative conditions on a one dimensional submanifold which represents the data.  Furthermore, using the Hamilton-Jacobi formalism, we extend the inflationary flow parameter approach to describe the evolutionary trajectories of general actions. This involves introducing three hierarchies of flow parameters to describe the evolution of a general action without using the specific restriction of field redefinition used in canonical and DBI inflation. These equations hold for all single field inflationary scenarios, whether or not slow roll conditions are met.  The importance of including kinetic terms in the inflationary reconstruction program cannot be overemphasized in light of recent theoretical and expected experimental advances.  Inflationary models with non-minimal kinetic terms are able to produce large non-Gaussian behavior for the curvature perturbations without ruining other inflationary observables \\citep{Alishahiha:2004eh} and predictions of non-Gaussian signatures for specific models have been established, for example in DBI inflation \\cite{Creminelli:2004yq,Seery:2005wm,Chen:2006nt,Maldacena:2002vr,Bean:2007eh}. Indeed, the search for such non-Gaussian effects is one of the primary current activities in observational cosmology, for example \\citep{Komatsu:2003fd,Spergel:2006hy,Yadav:2007yy}.  Non-Gaussianity detections open up the possibility of establishing which non-minimal kinetic terms may exist for inflationary models. The formalism that we present here will be useful for this purpose.  The explanation of any future (or current) observations of non-Gaussianities can also be checked in the context of single field inflation through the attendant modification of the tensor spectral index consistency relationship \\cite{Garriga:1999vw}.  The latter can be deduced experimentally from observations of tensor perturbations implied by CMB B-mode polarization measurements.  However, one advantage of emphasizing the non-Gaussianity connection with non-minimal kinetic terms is that the possibility of a large non-Gaussian contribution is generically independent of the single field paradigm.  On the other hand, the tensor spectral index consistency relationship changes for multifield inflationary models.  The order of our presentation will be as follows.  In section \\ref{kinetic}, we review and clarify the physics of how non-minimal kinetic terms contribute to non-Gaussianities, whose possible future observation is one of the strongest motivations for developing the action reconstruction formalism.  In \\ref{backg} we outline the general equations for the background evolution. We discuss the conditions for slow roll inflation in a general action in section \\ref{sr} . In \\ref{pert} we summarize how the generalized flow parameters relate to the properties of the primordial power spectrum and discuss how properties of a general action can be distinguished from canonical inflation using cosmological observations. In section \\ref{srrecon} we establish how the general action can be reconstructed from measurements of the lowest order flow parameters in the slow roll regime, and in \\ref{flow} we extend the inflationary flow parameters \\citep{Liddle:2003py} to describe a general inflationary action. In \\ref{conc} we summarize our findings and discuss their implications.  Throughout this paper with the exception of section \\ref{srrecon}, we use the usual reduced Planck scale conventions of $\\mpl^2 = (8\\pi G)^{-1}\\approx (2.4\\times 10^{18} \\mbox{GeV})^2$.  In section VI, we will use geometricized units and set $M_{pl}=1$ for simplicity in notation.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{conc} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  Complementary CMB and large scale structure measurements over scales spanning four orders of magnitude have driven impressive improvements in the measurement of the primordial scalar power spectrum. In addition improvements in non-Gaussianity constraints are expected from the PLANCK satellite and there is the exciting prospect of tensor mode measurements in the near future, with a number of CMB surveys being developed to target B-mode polarization.  With the hope of connecting this to high energy theory, there has been significant interest in establishing what current and planned observations might elucidate about the primordial spectrum of fluctuations from inflation, in terms of potential reconstruction for canonical inflation and action reconstruction in specific theories such as DBI inflation.  In this paper we extend these considerations to address what we can maximally learn about the inflationary action without making any assumptions, a priori, about its form.  We establish how observational constraints on the inflationary slow roll parameters could be successfully applied to reconstruct the general action over observable scales in the context of single field inflationary models.  Under the assumption of slow roll inflation, we have demonstrated that in an idealized case in which $\\{H, c_s, \\ep,\\obseta,\\obskappa\\}$ are measured over a finite range of scales, we analytically obtain the trajectory of the general action $\\L, X\\Lx, X^2\\Lxx$,  independent of the scalar field definition, with respect to some reference point. With the specification of a gauge, the measurement of the first level of flow parameters enables trajectories of $X, \\Lx$ and $\\Lxx$ and information about $\\Lp$ and $\\Lxp$ to be established.  Using the Hamilton-Jacobi formalism, we extend the inflationary flow parameter approach to describe the evolutionary trajectories of general actions. This involves introducing three hierarchies of flow parameters to describe the evolution of a general action without using the specific gauge, $\\Lx=c_s^{-1}$, used in canonical and DBI inflation. These equations hold for all single field inflationary scenarios, whether or not slow roll conditions are met.  Observations promise to allow us to reconstruct a wealth of information about the general action including powerful insights into the form of the Lagrangian kinetic, potential and hybrid terms and the relative importance of kinetic and potential components over the course of the trajectory. As it is difficult to obtain large observable non-Gaussianities without non-minimal kinetic terms (and/or resorting to curvaton scenarios), future detection of non-Gaussianities would make formalisms such as the one we present here indispensable to understand what kind of high energy theories are compatible with cosmological data. This is good news for PLANCK and other future CMB experiments which are certain to obtain increasingly precise data regarding non-Gaussianities, and it is also good news for high energy theorists looking for distinctive clues to the identity of the inflaton. In work in preparation, we are investigating the observational constraints on the general inflationary action using this formalism."
    },
    "0801/0801.4926_arXiv.txt": {
        "abstract": "This paper presents a mechanism that may modify the extinction law for SNIa observed at higher redshift. Starting from the observations that (1) SNIa occur predominantly in spiral galaxies, (2) star-formation ejects ISM out of the plane of spirals, (3) star-formation alters the extinction properties of the dust in the ISM, and (4) there is substantially more star-formation at higher redshift, I propose that spiral galaxies have a dustier halo in the past than they do now. The ejected material's lower value of $R_V$ will lead to a lower average value ($\\bar{R}_V$) for SNIa observed at higher redshift.   Two relations in SNIa observations indicate evolution of the average $R_V$: the relation of observed $R_V$ with inclination of the host galaxy at low redshift and the matching of the distribution of extinction values ($A_V$) for SNIa in different redshift intervals. The inclination effect does point to a halo with lower $R_V$ values. In contrast, the distributions of $A_V$ values match best for a $\\bar{R}_V(z)$ evolution that mimics the relation of SNIa dimming with redshift attributed to the cosmological constant. However, even in the worse case scenario, the evolution $\\bar{R}_V$ can not fully explain the dimming of SNIa: host galaxy extinction law evolution is not a viable alternative to account for the dimming of SNIa.  Future observations of SNIa --multi-color lightcurves and spectra-- will solve separately for values of $A_V$ and $R_V$ for each SNIa . Solving for evolution of $\\bar{R}_V$ (and $A_V$) with redshift will be important for the coming generation of cosmological SNIa measurements and has the bonus science of insight into the distribution of dust-rich ISM in the host galaxies in the distant past. ",
        "introduction": "Introduction}  Supernova type 1a (SNIa) distance modulus measurements have grown into a powerful measurement of the equation of state of our Universe. Accurate cosmological distances combined with redshift measurements allow for a precise characterisation of the Hubble flow as well as the additional acceleration attributed to the cosmological constant \\citep{Riess98,Perlmutter99}. The increasing statistics of SNIa measurements together with more information for each separate supernova event have progressively lowered the observational uncertainties: dust extinction, photometric error and lightcurve characterisation \\citep[See e.g.,][]{Tonry03, Knop03, Barris04, Conley06, Astier06}.  Extinction by dust remains a problematic systematic uncertainty in SNIa observations because the applicable extinction law remains poorly understood. This will need to be addressed in order to use SNIa in the next step in accuracy for a cosmological probe.  Dust attenuation affects both the observed SNIa rate and the distance modulus. Extinction by dust occurs in three instances before observation of the SNIa light: % (1) in our own Milky way, (2) in intergalactic space, and (3) in the host galaxy of the SNIa.  Galactic extinction (1) is a well-studied problem, because it is an ubiquitous one. \\cite{Burstein84} produced a  map of Galactic extinction based on HI maps and distant galaxy counts. It was superseded by the map of \\cite{Schlegel98}, based on COBE and IRAS far-infrared and sub-mm maps. The latter map is now generally used, together with the Galactic Extinction Law to correct extragalactic sources \\citep{Cardelli89, Fitzpatrick99}; the inferred extinction along a line-of-sight is proportional to the reddening: $A_V = R_V \\times E(B-V)$. The canonical value of $R_V$ is 3.1 with occasionally lower values towards the Galactic Center \\citep{Udalski03} and higher elsewhere \\citep{Cardelli89, Fitzpatrick99}.  Extinction by dust in intergalactic space (2) has been proposed as an alternative explanation for SNIa dimming, which is generally attributed to the cosmological acceleration \\citep{Aguirre99b, Aguirre99a, Aguirre00}. The resulting extinction law of this dust would effectively be grey\\footnote{The term grey extinction is used to denote that there is no relation between the reddening of an object and its extinction.} because the attenuating dust would be spread over all redshifts along every line-of-sight. The coincidence of both uniform distribution in both sky and redshift space does give this explanation a somewhat contrived appearance. The injection of dust into the intergalactic medium would have to be constant and substantial. Models of a dusty universe \\citep{Goobar02,Robaina07} find this grey dust explanation increasingly inconsistent with observational data \\citep[See also e.g., ][]{Riess07}.  Extinction within the SNIa's host galaxy (3), dust in the immediate surroundings and any disk, ring or spiral arm the line-of-sight passes through in projection, is an observationally evident, yet not fully constrained uncertainty. The Dark Energy Task Group Report \\citep{DETF} notes this as a primary source of uncertainty for SNIa measurements.  Three characteristics of the host galaxy's dust could --and are expected to-- change over the history of the universe: (a) total dust content or mass, (b) dust distribution within the host galaxy and (c) dust composition. The overall effect on the effective extinction law is of interest for the distance determination from SNIa light curves.  Dust mass (a) is a variable in several Spectral Energy Distribution (SED) studies of distant galaxies. \\cite{Calzetti99} modeled the overall dust content of galaxies over time and found that a maximum occurred either at z=1 or at z=3. \\cite{RR03} modeled the SED's from distant galaxies and similarly found a maximum dust content at z=1. \\cite{IP07} find a steady increase of dust mass with time from the UV-IR SED of galaxies. The typical dust mass found in distant galaxies is very much a function of the selected sample. Far-infrared selected samples point to dust-rich galaxies, similar to Arp 220 \\citep{RR05}, optical/UV selected samples point to disks very similar to the local ones \\citep{Sajina06} and Lyman-$\\alpha$ galaxies point to low-extinction disks \\citep{Nilsson07}. However, more dust mass should not affect the SNIa distance determinations if the extinction law remains the same for nearby and distant galaxies. More dust in the distant galaxies will predominantly affect the observed SNIa {\\it rate} \\citep{Hatano98, Cappellaro99, Goobar02, Riello05, Mannucci07}, as heavily obscured SNIa drop from observed samples.  The dust distribution in host galaxies (b) is commonly modeled as a double exponential, one radial and one vertical, sometimes with rings to mimic spiral arms. The radial scale of the dust distribution is assumed to be similar to the stellar one and the vertical dust scale is supposed to be much smaller than the stellar one. Previous observations of the scale-height in nearby (edge-on) galaxies appeared to corroborate the small scale-height \\citep[e.g,][]{Xilouris99,Holwerda05b,Bianchi07} but recent observations indicate a much higher scale-height for the dust \\citep{Seth05,Kamphuis07}, similar to the stellar scale. If the average scale-height of the dust distribution was higher in the past, then SNIa in the plane of the disk will encounter more extinction, especially when viewed in an inclined system. The different distribution has a similar effect as variation the dust content of galaxies and only the observed SNIa rate will be affected, unless dust composition and distribution are related.  The dust composition (c), notably the ratio of small to large grains, directly affects the observed extinction law \\citep[See the  review by][]{Draine03}. Evolution in the average extinction law is the most troubling possibility, as this would affect the extinction correction of SNIa and indirectly the measured Hubble flow and acceleration. For local SNIa, variations in the extinction law have been observed \\citep[e.g.,][]{Riess96b,Jha07}. \\cite{Wang05} explained the different observed extinction law for some SNIa as the effect of circumstellar material around the SNIa progenitor. This is a plausible scenario as the massive-star supernovae are very efficient dust producers \\citep{Sugerman06}.  \\cite{Patat07} and \\cite{Wang07} report observations of such material.  Alternatively, there is substantial evidence for a link between star-formation and extinction law in star-forming galaxies \\citep[See the review in][and the references therein.]{Calzetti01}. Star-formation produces more small grains and the intense UV fields alter grain composition.  Since it is not unreasonable to suppose that star-formation affects both the distribution and the extinction characteristic of dust in spiral disks, I propose a simple model for the evolution of the extinction law applicable to SNIa measurements. The aim of this paper is to explore this link between the extinction law for SNIa and star-formation of host galaxies and to investigate how much extinction law evolution could reasonably influence SNIa distance measurements. The paper is organized in a brief review of current observations of evolution in host galaxy dust and $R_V$ (\\S 2), a description of the model (\\S 3), two tests based on current SNIa observations (\\S 4), and discussion and conclusions (\\S 5). ",
        "conclusions": ""
    },
    "0801/0801.0608_arXiv.txt": {
        "abstract": "We used the Submillimeter Array to map the angular distribution of the H30$\\alpha$ recombination line (231.9 GHz) in the circumstellar region of the peculiar star MWC349A. The resolution was $1\\farcs2$, but because of high signal-to-noise ratio we measured the positions of all maser components to accuracies better than $0\\farcs01$, at a velocity resolution of $1~\\kms$. The two strongest maser components (called high velocity components) at velocities near $-14$ and $32~\\kms$ are separated by $0\\farcs048 \\pm 0\\farcs001$ (60 AU) along a position angle of $102 \\pm 1\\arcdeg$.  The distribution of maser emission at velocities between and beyond these two strongest components were also provided. The continuum emission lies at the center of the maser distribution to within 10 mas. The masers appear to trace a nearly edge-on rotating disk structure, reminiscent of the water masers in Keplerian rotation in the nuclear accretion disk of the galaxy NGC4258.  However, the maser components in MWC349A do not follow a simple Keplerian kinematic prescription with $v \\sim r^{-1/2}$, but have a larger power law index.  We explore the possibility that the high velocity masers trace spiral density or shock waves.  We also emphasize caution in the interpretation of relative centroid maser positions where the maser is not clearly resolved in position or velocity, and we present simulations that illustrate the range of applicability of the centroiding method. ",
        "introduction": "The hydrogen recombination line emission from the peculiar B[e] star MWC349A is double-peaked at frequencies above 100~GHz, and is thought to be due to maser action in a gaseous rotating circumstellar disk.  Thermal continuum emission from radio to IR wavelengths reveals that the star also has an ionized outflow.  Maser action occurs in recombination lines from about 100~GHz (H39$\\alpha$) to 1800~GHz (H15$\\alpha$) and beyond \\citep{str96a, str96b}.  A physical model developed by \\citet{thu94a,thu94b} suggests that the maser spots originate on the ionized surface of a rotating disk that is nearly edge-on in orientation.  \\citet{pon94} performed detailed radiative transfer modeling and concluded that the masers are probably unsaturated with a maximum negative opacity (gain) of about five.  High resolution VLA continuum images of MWC349A \\citep[e.g.,][]{whi85, tafo04} show an hourglass-shaped image whose angular size scales with frequency $\\nu$ as $\\nu^{-0.7}$ and flux density as  $\\nu^{0.6}$, as expected for a dense ionized outflow.  The star is located in the waist of an hourglass-shaped gas structure, with the north-south lobes corresponding to outflows, and the putative rotating disk presumed to  be located  around the star and viewed nearly edge-on.  It is interesting to note that VLA continuum images spanning the period of 1996--2006 \\citep{rod07} show that the continuum morphology of the object has evolved from an hourglass structure to a more of a rectangular projected structure (i.e., the waist of the hour glass structure is no longer prominent). The spectrum of the maser emission is also variable \\citep[see][]{mar89,gor01}.  \\citet{pla92} observed the H$30\\alpha$ maser emission and showed that its dominant characteristic is two features (aka ``spots'') at velocities of $-14$ and $32~\\kms$ separated by about $0\\farcs065 \\pm 0\\farcs005$, along a position angle of about $107\\arcdeg \\pm 7\\arcdeg$.  The close correspondence between the position angle of the spots and that of the waist of the hourglass continuum map suggests that the maser spots are located in the waist of the continuum.  For a distance of 1200 pc \\citep{coh85}, the velocity and spatial separation between the two maser peaks of $47~\\kms$ and 120 AU respectively, and a central mass of about $30~\\msun$, can be inferred from Kepler's third law  \\citep[e.g.,][]{str96a}.  Infrared continuum images at 1.65, 2.27, and $3.08~\\micron$ were made by \\citet{dan01} with the Keck telescope using an aperture masking technique.  The images are elliptical with major axes of 36, 47 and 62 mas, respectively. The position angles of the images are about $100\\arcdeg$.  These images trace the dust and the neutral gas that is probably undergoing photoevaporation, giving rise to the whole hourglass ionized region \\citep{hol94}. The structures of the radio continuum at 7~mm \\citep{tafo04}, the infrared at $2.27~\\micron$, and the maser emission from OVRO \\citep{pla92} are shown in a composite image in Figure \\ref{danchicont}.  \\begin{figure} \\begin{center} \\includegraphics[angle=-90,scale=0.8]{f1.ps} \\caption{Superposition of the 7--mm VLA continuum image (black contours) made by \\citet{tafo04}, $2.27~\\micron$ Keck aperture-masked image (green contours) made by \\citet{dan01}, and two maser components (red dot for 32~$\\kms$ and blue dot for $-16~\\kms$) measured by \\citet{pla92}. Relative astrometric accuracy among the wavelength bands was poorer than 0$\\farcs$1, we have co-registered the images to maximize the probable symmetry of the source.  This figure is an updated version of a similar figure published by \\citet{dan01}. The reference position for the image is RA$ = 20^h 32^m 45.528^s$ and DEC$ = 40\\arcdeg 39\\arcmin 36\\farcs622$ (2000), taken from the VLA observations. \\label{danchicont}} \\end{center} \\end{figure}  Here we present interferometric images of the maser emission in the H$30\\alpha$ radio recombination line, which have much higher signal-to-noise ratios than those of previous investigations.  This allows for the study of the distribution of maser emission in significantly greater detail.  We have additional data from the SMA in the H$26\\alpha$ (353 GHz) and H$21\\alpha$ (652 GHz) transitions.  These results have lower signal-to-noise ratio than the ones presented here. They will be discussed in a later publication. ",
        "conclusions": "We have presented a new high resolution image of the maser features in the inner ionized envelop of MWC349A.  The masers act as tracers of the velocity distribution in the disk.  We used the PCA method to determine the positions of the maser spots to an accuracy much finer than the resolution of the interferometer. Based on our simulations we find that there are two plausible interpretations to our data. In both models the linear distribution of masers in the p-v diagram between $-10$ and $25~\\kms$ suggests the presence of an annular ring of maser emission with a radius of about $0\\farcs025$. The scatter is due to measurement noise, intrinsic thickness in the annular masing region, and inclination of the ring.  If the high velocity maser components have linewidths of about $5~\\kms$, then the simulations of the PCA method indicate that they may lie along the midline. If, on the other hand, these components have much narrower intrinsic linewidths because of higher maser amplification, then they may be distributed along spiral structures.  If the maser components are the result of the unsaturated amplification, then their brightness temperatures will be much greater than the thermal temperature of the gas by a factor of exp(G), where G is the maser gain factor. Thus, these masers may be observable at much higher resolution. ALMA, in its extended configuration of 10 km, could be used to observe MWC349 (which will have a maximum elevation angle of about 30$\\arcdeg$) at a resolution of about $0\\farcs 02$ at 230 GHz, and could distinguish between the models we have proposed."
    },
    "0801/0801.2327_arXiv.txt": {
        "abstract": "Shell galaxies are considered the debris of recent accretion/merging episodes. Their high frequency in low density environments suggest that such episodes could drive the secular evolution for  at least some fraction of the early-type galaxy population. We present here the preliminary results of ultraviolet and X-ray data for a sample of  three shell galaxies, namely NGC~474, NGC~7070A and ESO~2400100.  The Far UV morphology and photometry are derived using the  observations obtained with the {\\it Galaxy Evolution Explorer} and the XMM-{\\it Newton} Optical Monitor. We aim at investigating the {\\it rejuvenation} processes in the stellar population using the UV information as well as at gaining information about the possible evolution with time of the X-ray emission due interaction/merging processes. ",
        "introduction": "In a hierarchical evolutionary scenario, galaxies experience accretion/merging events during their lifetime. While early-type  galaxies in nearby clusters appear (homogeneously) old, the field early-type galaxy population seems to contain genuinely, recently {\\it rejuvenated} objects \\citep[see e.g.][]{Clemens06}. Early-type galaxies showing fine structure, like shells, occupy a special position since they are believed to fill the gap between ongoing mergers and normal elliptical galaxies.  Shells are faint, sharp-edged stellar features  characterizing $\\approx$ 16.5\\% of the field early-type galaxy population and avoiding  cluster environments \\citep[e.g.][]{ Malin83, Reduzzi96, Colbert01}. % Two major scenarios for their origin have been proposed: {\\em a)} a weak interaction between galaxies \\citep{Thomson90, Thomson91}, {\\em b)} merging/accretion events \\citep{Dupraz86, Hernquist87a, Hernquist87b}. In the former scenario, weak interaction can form long lasting azimuthally distributed shells through the interference of density waves producing a thick disc population of dynamically cold stars.  However, this requires a cold thick disc, not found  in ellipticals. In the  merging/accretion events between galaxies of different masses (mass ratios typically 1/10 - 1/100, {\\em b)} scenario above) shells are density waves formed from infalling stars from the companion during a minor merger. Major merger can also  produce shells \\citep{Barnes92, Hernquist92, Hernquist95}.  Whatever the mechanism is, interaction/merging events seem to have played a significant role in the formation/evolution of the early-type class in the field and  activated a new star formation event \\citep{Longhetti99,Longhetti00}.  Whether there is a link between shell galaxies, the early phases of merging processes, and the class of 'normal' early-type galaxies is still an open question. To test the depth of the link, a multiwavelength approach is of paramount importance. The UV  data probe the ongoing/recent star formation processes and study their distribution across the galaxy while the X-ray emission is connected to the past star formation and metal enrichment history  of the bulk of the galaxy and may disclose hidden AGN activity.  We present  {\\it Galaxy Evolution Explorer} ({\\it GALEX}), XMM-{\\it Newton} Optical Monitor (OM) and XMM-EPIC observations of  NGC 474, NGC 7070A, and ESO 2400100, three shell galaxies in the \\citet{Malin83} catalogue.  \\begin{table*} \\tabletypesize{\\scriptsize} \\caption{Journal of the XMM-{\\it Newton} EPIC and  Optical Monitor observations} \\begin{tabular}{llllllllll} \\hline \\multicolumn{1}{l}{Galaxy}& \\multicolumn{1}{c}{EPIC}  & \\multicolumn{1}{c}{EPIC}  & \\multicolumn{1}{c}{OM}  & \\multicolumn{1}{c}{OM}  & \\multicolumn{1}{c}{OM}  & \\multicolumn{1}{c}{OM}  & \\multicolumn{1}{l}{observing} & \\multicolumn{1}{l}{observation}\\\\ \\multicolumn{1}{l}{}& \\multicolumn{1}{c}{PN}  & \\multicolumn{1}{c}{MOS}  & \\multicolumn{1}{c}{UVM2}  & \\multicolumn{1}{c}{UVW1}  & \\multicolumn{1}{c}{U}  & \\multicolumn{1}{c}{B}  & \\multicolumn{1}{l}{date} & \\multicolumn{1}{l}{ident.}\\\\ \\multicolumn{1}{c}{}& \\multicolumn{1}{c}{ [sec]} & \\multicolumn{1}{c}{ [sec]} & \\multicolumn{1}{c}{ [sec]} & \\multicolumn{1}{c}{ [sec]} & \\multicolumn{1}{c}{ [sec]} & \\multicolumn{1}{c}{ [sec]} & \\multicolumn{1}{c} {}& \\multicolumn{1}{c}{}\\\\ \\hline NGC~474        &   4300 & 11400  &   5000 &  5000  &  5000 &  & 2004-01-24 & 0200780101 \\\\ NGC~7070A    &  26250 & 30840  &  4400  &          &         & 4000  &2004-10-28  & 0200780301  \\\\ ESO 2400100  & 19970  & 26330  & 4400    &  4400  & 4400 &  &2004-05-11  & 0200780201 \\\\ \\hline \\end{tabular}  \\medskip {XMM-{\\it Newton} observations of NGC 474 are fully discussed in \\citet{Rampazzo06} } \\label{table1} \\end{table*} ",
        "conclusions": ""
    },
    "0801/0801.2782_arXiv.txt": {
        "abstract": "The globular cluster $\\omega$ Centauri is one of the largest and most massive members of the galactic system. However, its classification as a globular cluster has been challenged making it a candidate for being the stripped core of an accreted dwarf galaxy; this together with the fact that it has one of the largest velocity dispersions for star clusters in our galaxy makes it an interesting candidate for harboring an intermediate mass black hole. We measure the surface brightness profile from integrated light on an $\\it{HST}$/ACS image of the center, and find a central power-law cusp of logarithmic slope -0.08. We also analyze Gemini GMOS-IFU kinematic data for a 5x5\\arcsec~field centered on the nucleus of the cluster, as well as for a field 14\\arcsec~away.  We detect a clear rise in the velocity dispersion from 18.6~\\kms\\ at 14\\arcsec~to 23 \\kms\\ in the center. A rise in the velocity dispersion could be due to a central black hole, a central concentration of stellar remnants, or a central orbital structure that is radially biased. We discuss each of these possibilities. An isotropic, spherical dynamical model implies a black hole mass of $4.0^{+0.75}_{-1.0}\\times10^4 M_\\odot$, and excludes the no black hole case at greater than 99\\% significance. We have also run flattened, orbit-based models and find similar results. While our preferred model is the existence of a central black hole, detailed numerical simulations are required to confidently rule out the other possibilities. ",
        "introduction": "\\vspace{10pt}  The globular cluster $\\omega$ Centauri (NGC~5139) is regarded to be the largest and most massive member of the Galactic cluster system with a tidal radius of 69 parsecs \\citep{har96}, an estimated mass of $5.1\\times10^6 M_\\odot$, and a measured central velocity dispersion of $22 \\pm 4$ \\kms\\ \\citep {mey95}. The cluster presents a large scale global rotation, measured with radial velocities, of 8 \\kms\\ at a radius of 11pc from the center \\citep{mer97} and confirmed with proper motions \\citep{van00}, which makes it one of the most flattened galactic globular clusters \\citep{whi87}. A rotating flattened model including proper motion and radial velocity datasets by \\citet{ven06} calculate a lower total mass of $2.5\\times10^6 M_\\odot$ and confirm the central line of sight velocity dispersion value of 20 \\kms. They measure a dynamical distance of $4.8\\pm0.3$ kpc (which we adopt for this paper). $\\omega$ Cen has a peculiar highly bound retrograde orbit \\citep{din01}. It also has a stellar population that makes it stand out from the rest of the Galactic globular clusters due to its complexity. It shows a broad metallicity distribution \\citep{bed04,nor96}, as well as a kinematical and spatial separation between the different subpopulations \\citep{pan03, nor97}.  All the above results have led to the hypothesis that $\\omega$ Cen is not a classical globular cluster, but instead is the nucleus of an accreted galaxy \\citep{fre03,bek03,mez05}. The scenarios of it being the product a merger of two globular clusters \\citep{ick88} and of self-enrichment \\citep{iku00} have also been proposed to explain the stellar populations.  The high measured velocity dispersion together with the possibility of being a stripped galaxy make $\\omega$ Cen an interesting candidate for harboring a black hole in its center. The extrapolation of $M_\\bullet-\\sigma$ relation for galaxies \\citep{geb00a,fer00,tre02} predicts a $1.3\\times10^4 M_{\\odot}$ black hole for this cluster. The sphere of influence of such black hole for a star cluster at the distance of $\\omega$ Cen with a velocity dispersion of 20 \\kms\\ is $\\sim$ 5\\arcsec, making it an excellent target for ground-based observations.  Two globular clusters have been suggested for harboring an intermediate mass black hole in their nucleus. One is the galactic cluster M15 \\citep{geb00,ger02,ger03} and the other is G1, a giant globular cluster around M31 \\citep{geb02,geb05}. M15 is assumed to be in a post-core collapse state, therefore its dynamical state has been debated between harboring a black hole or containing a large number of compact remnants in its center \\citep{bau03a,bau03b}. Unfortunately, observational constraints between these two hypothesis remain inconclusive \\citep{bos06}. G1 on the other hand, has a core with characteristics closer to those of $\\omega$ Cen, and observations support the black hole interpretation for G1. \\citet{bau03c} propose an alternative interpretation for G1 in which they match the observations with a model of two colliding globular clusters. The G1 black hole models are preferred since the M/L profile is expected to be flat in its core, so any rise in the velocity dispersion is unlikely to be due to remnants that concentrated there from mass segregation. The black hole interpretation for G1 is strongly supported by radio \\citep{ulv07} and x-ray \\citep {poo06} detections centered on the nucleus. $\\omega$ Cen and G1 have similar properties in both their photometric and kinematic profiles. In this paper we report photometric and kinematical measurements that suggest the presence of a central black hole in $\\omega$ Cen.  \\begin{figure} \\plotone{f1.eps} \\caption{Surface brightness profiles for $\\omega$ Cen. The circles show our measured photometric points from the ACS (filled) and H-alpha (open) images. The triangles show photometric points obtained from ground based images by Trager et al. The dashed line is Trager's Chebychev fit. The solid line is our smooth fit to the combination of the ACS points inside 40\\arcsec\\ and Trager's Chebychev fit outside.} \\label{sb} \\end{figure} ",
        "conclusions": "\\vspace{10pt}  We measure the surface brightness profile for the globular cluster $\\omega$ Centauri (NGC~5139) from an ACS image in the central 40\\arcsec. The profile shows a continuous rise toward the center with a logarithmic slope of $-0.08\\pm0.03$, in contrast with previous measurements which found a flat core. The shape of the profile is similar to that obtained from numerical models of star clusters containing black holes in their centers. We measure a line of sight velocity dispersion for two 5\\arcsec$\\times$5\\arcsec~regions, one at the center of the cluster and the other 14\\arcsec~away. We detect a rise in velocity dispersion from 18.6 \\kms\\ for the outer field to 23 \\kms\\ for the central one. We combine these two measurements with previously measured velocity dispersion at larger radii.  When we compare the observed velocity dispersion profile with a series of isotropic models containing black holes of various masses, we find that a black hole of $4.0 ^{+0.75}_{-1.0}\\times~10^4 M_\\odot$ is necessary to match the observations. We explore alternative explanations for the observed rise in our central velocity dispersion measurements. First we consider the possibility that the observed $M/L$ rise is due to the presence of an extended component composed of dark remnants such as neutron stars or faint white dwarfs. The density profile of the dark component is required to be extremely concentrated toward the center, with a configuration practically decoupled from the luminous component. $\\omega$ Cen has a weak cusp in the central luminosity density profile, implying that the gravitational potential is very shallow inside the core and therefore mass segregation is only a weak effect. \\citet{fer06} confirm the lack of segregation by measuring the radial distribution of blue straggler stars, which are heavy stars and should sink to the center of the cluster if there is mass segregation. They find a flat radial distribution of blue stragglers with respect to lighter stellar populations. Also, the formation channel of the blue stragglers is not collisional, as it would be expected if there was a considerable amount of mass segregation. With this evidence in hand, there is no reason to expect a large variation of $M/L$ inside the core due to stellar content, so a detected rise in $M/L$ is likely to come from the presence of a concentrated massive object.   \\begin{figure} \\plotone{f9.eps} \\caption{$M_\\bullet-\\sigma_{vel}$ relation for elliptical galaxies and bulges. The solid line is the relation in \\citet{tre02}. $\\omega$ Cen lies on the low mass extrapolation and suggests a similarity between it and the galaxies. Different types of systems such as star clusters and low luminosity AGN appear to populate the low mass end of the diagram.} \\label{bhsigma} \\end{figure}  There is also a stability argument against a dense compact cluster of dark remnants. The central density as measured from the 23~\\kms\\ dispersion estimate at 1.8\\arcsec\\ is $5.6\\times10^7 M_\\odot/pc^3$. This is the largest measured for a globular cluster and it would be difficult to maintain using stellar remnants. Obviously, if the density is due to solar mass remnants, over $10^4$ remnants would be required inside of 0.05 parsecs. Using the arguments of Maoz (1998) and Miller (2006), this mass and density makes $\\omega$~Cen one of the better examples where stellar remnants can be ruled out due to evaporation. Maoz estimates that for these numbers, any cluster of remnants will have evaporated within the age of the cluster.  An observed velocity dispersion rise toward the center of a cluster can also occur if a degree of anisotropy is present. If more radial orbits are present, those stars pass near the center at higher velocities than they would in an isotropic case. This possibility is evaluated with our orbit-based models. Our results agree with the model by van den Ven et al. (2006) in showing no anisotropy in the central 10 arcmin. Their models show a degree of tangential anisotropy at large radius, but no radial anisotropy. Models without a central black hole but having a large degree of anisotropy inside 28\\arcsec\\ are not as as good as models including a black hole. Furthermore, the high degree of radial anisotropy is highly unstable in dense systems like $\\omega$~Cen and therefore it is an unlikely explanation for the observed kinematics.  Figure \\ref{bhsigma} shows the known $M_\\bullet-\\sigma_v$ relation for black holes in elliptical galaxies and bulges \\citep{geb00a,fer00}. The galaxies used to determine the relation \\citep{tre02} are plotted along with objects containing smaller black holes in low luminosity quasars \\citep{bar05}, two nearby low luminosity AGN (NGC~4595 and Pox~52), and three globular clusters (G1, M15 and $\\omega$ Cen). We also plot the upper limit for the black hole mass in the nucleus of M33 \\citep{geb01}, which does not lie on the correlation. The black hole in $\\omega$ Cen lies above the relation, but it is consistent with the scatter observed a larger masses.  The measured black hole mass is 1.6\\% of the total mass of the cluster, which is much larger than the canonical value of $\\sim0.3$\\% for larger spheroids \\citep{mag98}. If $\\omega$ Cen is indeed the nucleus of an accreted galaxy it is expected that it's original mass was considerably larger than what we measure now. \\citet{bek03} reproduce the current mass and orbital characteristics of $\\omega$ Cen with a model of an accreted $10^7M_\\odot$ dwarf galaxy. A mass of $4\\times10^7M_\\odot$ for the original spheroid would put the black hole near the 0.3\\% value.  The two pieces of observational evidence that $\\omega$ Cen could harbor a central black hole come from the photometry and the kinematics. From the HST image of $\\omega$ Cen, we measure a central logarithmic surface brightness slope of $-0.08\\pm0.03$. This value is very similar to that claimed by the N-body simulations of Baumgardt et al. (2005) that are most likely explained by a central black hole. Standard core-collapse does not lead to such a large core with a shallow central slope. The black hole tends to prevent core collapse while leaving an imprint of a shallow cusp. It will be important to run models tailored to $\\omega$~Cen to see if one can cause and maintain a shallow cusp without invoking a central black hole. However, the main observational evidence for the central mass comes from the increase in the central velocity dispersion, where we detect a rise from 18.6 to 23~\\kms\\ from radii of 14 to 2.5\\arcsec. In fact, even excluding the Gemini data presented here, the previous ground-based data suggest a central mass concentration as well. The core of $\\omega$~Cen is around 155 \\arcsec\\ (about 2.5 \\arcmin), so the dispersion rise is seen well within the core."
    },
    "0801/0801.0264_arXiv.txt": {
        "abstract": "Density is the turbulence statistics that is most readily available from observations. Different regimes of turbulence correspond to different density spectra. For instance, the viscosity-damped regime of MHD turbulence relevant, for instance, to partially ionized gas, can be characterized by shallow and very anisotropic spectrum of density. This spectrum can result in substantial variations of the column densities. Addressing MHD turbulence in the regime when viscosity is not important over the inertial range, we demonstrate with numerical simulations that it is possible to reproduce both the observed Kolmogorov spectrum of density fluctuations observed in ionized gas by measuring scintillations and more shallow spectra that are obtained through the emission measurements. We show that in supersonic turbulence the high density peaks dominate shallow isotropic spectrum, while the small-scale underlying turbulence that fills most of the volume has the Kolmogorov spectrum and demonstrates scale-dependent anisotropy. The limitations of the spectrum in studying turbulence induce searches of alternative statistics. We demonstrate that a measure called \"bispectrum\" may be a promising tool. Unlike spectrum, the bispectrum preserves the information about wave phases. ",
        "introduction": "The paradigm of interstellar medium has undergone substantial changes recently. Instead of quiescent medium with hanging and slowly evolving clouds a turbulent picture emerged (see review by Ballesteros-Parredes et al. 2006, McKee \\& Ostriker 2007 and ref. therein). With magnetic field being dynamically important and dominating the gas pressure in molecular clouds, this calls for studies of compressible magnetohydrodynamic (MHD) turbulence.  Key ideas in describing MHD turbulence can be traced back to Iroshnikov (1963) and Kraichnan (1965) work and the classical work that followed (see Montgometry \\& Turner 1981, Higdon 1984, Montgomery, Brown \\& Matthaeus 1987). A more recent model by Goldreich \\& Sridhar (1995, henceforth GS95) has been successfully tested by numerical simulations (Cho \\& Vishniac 2000, Maron \\& Goldreich 2001, Cho, Lazarian \\& Vishniac 2002, Cho \\& Lazarian 2002, 2003, henceforth CL03).  Density statistics is the easiest to infer from observations. In comparison, velocity spectra require elaborate techniques to be used (see Lazarian 2006 and ref. therein). How informative can be density fluctuations for understanding of MHD turbulence is the subject of the present review. However, we should mention that for many important astrophysical applications, e.g. for interstellar chemistry, for star formation, for propagation of radiation etc. the density fluctuations themselves are crucially important.  Subsonic compressible MHD is rather well studied topic today. It is  suggestive that there may be an analogy between the subsonic MHD turbulence and its incompressible counterpart, namely, GS95 model. Therefore the correspondence between the the two revealed in CL03 is expected.  It could be easily seen, that in the low-beta case density is perturbed mainly due to the slow mode (CL03). Slow modes are sheared by Alfv\\'en turbulence, therefore they exhibit Kolmogorov scaling and GS95 anisotropy for low Mach numbers. However, for high Mach numbers we expect shocks to develop. Density will be perturbed mainly by those shocks.  One can also approach the problem from the point of view of underlying hydrodynamic equations. It is well known that there is a multiplicative symmetry with respect to density in the ideal flow equations for an isothermal fluid (see e.g. Passot \\& Vazquez-Semadeni, 1998). This assume that if there is some stochastic process disturbing the density it should be a multiplicative process with respect to density, rather than additive, and the distribution for density values should be lognormal, rather that normal. The aforementioned work shows that for 1D numerical simulations of high-Mach hydrodynamics the distribution is approximately lognormal, having power-law tails in case of $\\gamma\\neq 1$.  In MHD the above described symmetry is broken by the magnetic tension. However, numerical studies in Beresnyak, Lazarian \\& Cho (2005, henceforth BLC05) show an approximate correspondence of the PDFs obtained with MHD simulations to the lognormal scaling.  With a high sonic Mach we expect a considerable amount of shocks arise. In a sub-Alfv\\'enic case, however, we expect oblique shocks be disrupted by Alfv\\'enic shearing, and, as most of the shocks are generated randomly by driving, almost all of them will be sheared to smaller shocks. The evolution of the weak shocks will be again governed by the sonic speed, and structures from shearing as in low Mach case should arise.  We also note that shearing will not affect probability density function (PDF) of the density, but have to affect its spectra and structure function (SF) scaling. In other words, we deal with two distinct physical processes, one of which, random multiplication or division of density in presence of shocks, affect PDF, while the other, Alfv\\'enic shearing has to affect anisotropy and scaling of the structure function of the density. The numerical studies that we describe below confirm the theoretical considerations above. ",
        "conclusions": "Above we discussed the statistics of density in compressible MHD turbulence. We analyzed spectra, structure functions for experiments with different sonic ${\\cal M}_s$ and Alfv\\'{e}n ${\\cal M}_A$ Mach numbers. Our results are as follows:  $\\bullet$ The viscosity-damped regime of MHD turbulence relevant to partially ionized gas can be characterized by shallow and very anisotropic spectrum of density. This spectrum can result in large variations of the column densities.  $\\bullet$ The amplitude of density fluctuations strongly depends on ${\\cal M}_s$ both in weakly and strongly magnetized turbulent plasmas.  $\\bullet$ The flattening observed in density spectra is due to the contribution of the highly dense small-scale structures generated in the supersonic turbulence.  $\\bullet$ Fluctuations of the logarithm of density are much more regular than those of density. The logarithm of density exhibits the GS95 scalings and anisotropies.  $\\bullet$ Bispectra show strong correlations for supersonic models than for subsonic ones. It suggests the importance of shocks in mode correlations.  {\\footnotesize {\\bf Acknowledgement} of the support of the NSF Center for Magnetic Self-Organization in Laboratory and Astrophysical Plasma.}"
    },
    "0801/0801.0863_arXiv.txt": {
        "abstract": "{} {Despite photometry and spectroscopy of its oscillations obtained over the past 25 years, the pulsation frequency spectrum of the rapidly oscillating Ap (roAp) star $\\gamma$ Equ has remained poorly understood. Better time-series photometry, combined with recent advances to incorporate interior magnetic field geometry into pulsational models, enable us to perform improved asteroseismology of this roAp star.} {We obtained 19 days of continuous high-precision photometry of \\object{$\\gamma$ Equ} with the {{\\sc Most}} (Microvariability \\& Oscillations of STars) satellite.  The data were reduced with two different reduction techniques and significant frequencies were identified.  Those frequencies were fitted by interpolating a grid of pulsation models that include dipole magnetic fields of various polar strengths.} {We identify 7 frequencies in $\\gamma$ Equ that we associate with 5 high-overtone p-modes and 1st and 2nd harmonics of the dominant p-mode. One of the modes and both harmonics are new discoveries for this star. Our best model solution (1.8 $M_{\\sun}$, log $T_{\\rm eff} \\sim 3.882$; polar field strength $\\sim$ 8.1 kG) leads to unique mode identifications for these frequencies ($\\ell = $ 0, 1, 2 and 4).  This is the first purely asteroseismic fit to a grid of magnetic models.  We measure amplitude and phase modulation of the primary frequency due to beating with a closely spaced frequency that had never been resolved.  This casts doubts on theories that such modulation -- unrelated to the rotation of the star -- is due to a stochastic excitation mechanism.} {} ",
        "introduction": "\\label{intro}  {Rapidly oscillating Ap (roAp) stars form a class of variables consisting of cool magnetic Ap stars with spectral types ranging from A-F, luminosity class V, and were discovered by Kurtz (1982). Photometric and spectroscopic observations during the last 2 decades characterize roAp stars with low amplitude pulsations ($<$ 13\\,mmag) and periods between 5 to 21 minutes. It is now widely accepted that the oscillations are due to high-overtone ($n >15$) non-radial p-mode pulsations. The observed frequencies can be described in a first-order approximation by the oblique pulsator model as described by \\citet{kurtz82}. In this model the pulsation axis is aligned with the magnetic axis, which itself is oblique to the rotation axis of the star. It has since been refined and improved to include additional effects like the coriolis force \\citep{bigotdziem02}.  The oscillations in roAp stars are most probably excited by the $\\kappa$ mechanism acting in the hydrogen ionization zone \\citep{dziemgo}, while the prime candidate for the selection mechanism of the observed pulsation modes is the magnetic field \\citep[e.g.][]{bigotetal00, cunhagough, balmf, bigotdziem02,  bigotdziem03, saio, cunha06}. Similar to non-oscillating Ap (noAp) stars, the group of roAp stars also shows a distinctive chemical pattern in their atmospheres, in particular an overabundance of rare earth elements, which points to stratification and vertical abundance gradients. Tackling the question of what drives the roAp pulsation, it is important to note that apart from the similar chemical composition, noAp stars show comparable rotation periods and magnetic fields but higher effective temperatures \\citep{tanya04}. Thus, their hydrogen ionization zone lies further out in the stellar envelope, which prevents efficient driving of pulsation. All together, roAp stars have given new impulse to asteroseismology and 2D mapping of pulsation \\citep{oleg04apj}, as well as detailed 3D mapping of (magnetic) stellar atmospheres, has become possible. For reviews on primarily observational aspects of roAp stars we refer to \\citet{kurtz00} and \\citet{kurtz04} and on theoretical aspects to \\citet{shiba03} and \\citet{cunha05}.  $\\gamma$ Equ, (HD\\,201601, HR\\,8097, A9p, $m_{\\rm{V}} = 4.7$) was the 6th discovery as a roAp star. It was long known to have a significant magnetic field \\citep{babcock} which is variable with a period of $\\sim$\\,72 years \\citep{bonsack, scholz79, leroy}.  Rotation, a solar-like magnetic cycle or precession of the star's rotation axis due to its binary companion are candidate mechanisms for this cyclic variation. The binary hypothesis was supported by \\citet{scholz97}, who observed a temporary drop in radial velocity for $\\gamma$ Equ during 4 consecutive nights. Their results are however contradicted by \\citet{mkrtich98, mkrtich99}. Magnetic field data spanning more than 58 years indicate that the variation is most likely due to rotation and therefore $P_{\\rm{mag}} = P_{\\rm{rot}}$ \\citep{bychkov06}. This explanation is also supported by many spectroscopic observations showing very sharp absorption lines \\citep[e.g.][]{kanaan} which are typical for a slow rotator (or a star seen pole-on).  \\citet{kurtz83} was the first to detect a pulsation period of 12.5\\,min (1.339\\,mHz) with an amplitude varying between 0.32 and 1.43\\,mmag and speculated about a rotation period of 38 days, inconsistent with the magnetic field measurements mentioned above.  Aside from rotation, beating with a closely spaced frequency has been proposed as a cause, for which the present paper gives further evidence. The pulsation was confirmed shortly after Kurtz's study by \\citet{weiss83}. \\citet{bychkov87} reported first evidence for radial velocity (RV) variations, but it took two more years for a clear detection \\citep{libbrecht}. Although unable to detect the 1.339 mHz oscillations reported by Kurtz (1983), Libbrecht discovered three frequencies at 1.365\\,mHz, 1.369\\,mHz and 1.427\\,mHz.  He suggested that the amplitude modulation observed in the spectra of roAp stars may not be due to closely spaced frequencies, but rather caused by short mode lifetimes in the order of $\\sim$ 1\\,d. He concluded that both peaks therefore belong to a single p-mode oscillation. The follow-up study by \\citet{weiss89} aimed at performing simultaneous spectroscopic and photometric studies but failed to confirm $\\gamma$ Equ's pulsation. The latter authors applied a correlation technique using more than 100\\,\\AA\\ wide spectral ranges in order to improve the S/N ratio relative to previous techniques, based on individual lines. This approach turned out to be inapplicable, considering the peculiar abundance structure of roAp star atmospheres which is now well known. \\citet{matthews} then claimed to have detected variability with rather large radial velocity amplitudes but their results were inconsistent with previous findings.  Detailed photometric observations were conducted by \\citet{martinez} using a multi-site campaign in 1992, spanning a total of 26 nights. Their results also suggested limited life times of pulsation modes, because additional frequencies appeared in their analysis of individual nights. However, all three frequencies detected so far were confirmed, and prewhitening of 1.366\\,mHz revealed a fourth eigenfrequency at 1.397\\,mHz. Using these four frequencies they concluded a spacing of consecutive radial overtones of $\\Delta\\nu\\sim 30\\,\\mu$Hz based on the asymptotic theory for low-degree, high-overtone p-mode pulsations \\citep{tassoul}.  They did not detect any evidence for non-linear behavior of the eigenfrequencies (i.e. harmonics) which, as will be shown in this paper, is indeed present in the {\\sc Most} data. So far their observing campaign was the most recent photometric investigation of $\\gamma$ Equ. A period of spectroscopic studies focusing on pulsation, abundance and stratification analyses followed.  \\citet{kanaan} reported on RV amplitudes for chromium and titanium and speculated that this is due to their concentration close to the magnetic -- hence also pulsation -- poles, while other elements (e.g. iron) showing no RV variations may be concentrated at the magnetic equator. \\citet{malanush} independently arrived at a similar conclusion and they are the first who identified the rare earth elements (REE) Pr{\\,\\sc iii} and Nd{\\,\\sc iii} to show the largest variations (up to 800\\,$\\rm ms^{-1}$). Most other atomic species had very low or non-measurable RV variations and the authors concluded that previous spectroscopic observations failed to detect significant radial velocity variations due to the chosen spectral regions and/or low spectral resolution. \\citet{savanov99} further clarified the situation by drawing attention to a possible incorrect identification of spectral lines in \\citet{kanaan} and concluded that Pr{\\,\\sc iii} and Nd{\\,\\sc iii} were responsible for the large RV amplitudes discussed in their analysis.  Further investigations by \\citet{oleg01} and \\citet{tanya02} impressively explained amplitude modulations of different elements and even ions of the same element by lines being formed at different atmospheric depths. The authors also tried a first mode identification based on rather short data sets of line profile variations (LPVs) and they argue for $\\ell = 2$ or $3$, and $m = -\\ell$ or $-\\ell + 1$ modes. \\citet{shiba04} disagreed with this analysis and suggested a shock wave causing the observed LPVs which in turn was questioned by \\citet{oleg07} who argued that the pulsational velocity should not exceed the local sound speed. The latter authors propose a modified oblique pulsator model where LPVs are caused by pulsation velocity fields superposed by sinusoidal line width changes due to convection, but which previously was thought to be suppressed. In any case, they agree with Shibahashi et al. that the identification of the primary frequency as an $\\ell=1, m=0$ mode is still the most likely explanation, which is also supported by this paper. The issue is still being debated though \\citep{shiba07}.  A still unsolved issue are magnetic field variations synchronized with pulsation. \\citet{leone} showed evidence for this, but subsequent investigations \\citep{olegetal04, bychkov05, savanov06} could not confirm  their findings.  \\citet{hubrig} corroborate this null detection for $\\gamma$ Equ with ESO-FORS1 data.  Despite a history of more than 25 years of observations, the pulsation frequency spectrum of $\\gamma$ Equ remains poorly understood. Several frequencies have been published, but up to now these have never been observed simultaneously. This is where {\\sc Most} steps in, providing the most complete high precision photometric campaign ever conducted for $\\gamma$ Equ and covering continuously a timespan of 19 days.} ",
        "conclusions": "\\label{concl} {We have demonstrated that our models of roAp stars reproduce the observations in great detail. The temperature and luminosity parameters of our best fit are in good agreement with those expected from previous observations \\citep{olegbagnulo}, although we find a slightly higher mass. According to \\citet{tanya97} the mean magnetic field modulus of $\\gamma$ Equ is $\\rm \\simeq 4$\\,kG, which is half of the polar magnetic field strength of 8\\,kG as suggested by our best model. At present we can only speculate about a difference of internal magnetic field strengths, to which pulsation is sensitive, relative to surface magnetic fields, which are accessible to  spectroscopy. Furthermore, our models assume a simple, axisymmetric magnetic dipole, which may not sufficiently reflect a more complicated reality. The factor of 2 between (surface) magnetic field modulus and best fitting magnetic pulsation model appears again for the roAp star \\object{10\\,Aql} \\citep{huber}, another {\\sc Most} primary target. Certainly, more investigations of this sort have to be conducted in order to solve this problem.  Pulsation amplitude changes for roAp stars were discussed in the literature as a consequence of limited mode life time or beating frequencies. For $\\gamma$ Equ we can clearly identify the beating frequencies, which seriously  questions excitation mechanisms for roAp stars based on stochastic processes.  Finally, we could identify the modes of all 7 frequencies detected so far in $\\gamma$ Equ  as $\\ell = 0$, $1$, $2$ and $4$, with $f_{\\rm{6}}$ and $f_{\\rm{7}}$ being the harmonics of $f_{\\rm{1}}$ produced by non-linear pulsation. Since we are not able to calculate non-axisymmetric modes, all of these frequencies are assumed to be ($m=0$) p-modes. While we cannot rule out the possibility of excitation of frequencies with ($m\\neq0$), axisymmetric modes are most probable, because of the extremely long rotation period of $\\gamma$ Equ and the assumption of a magnetic dipole. The assignment of any $\\ell$-value to these modes, however, has to be understood as a convenient simplification. In the presence of strong magnetic fields a mode's oscillation behavior cannot be described by a single spherical harmonic. Consequently, we alert the reader that analysis techniques using this assumption, e.g. mode identification based on LPVs, should be adjusted for magnetic effects on the pulsation geometry.}"
    },
    "0801/0801.3267_arXiv.txt": {
        "abstract": "We present a three-year series of observations at 24$\\mu$m with the {\\it Spitzer} Space Telescope of the interstellar material in a 200 x 200 arcmin square area centered on Cassiopeia A. Interstellar dust heated by the outward light pulse from the supernova explosion emits in the form of compact, moving features. Their sequential outward movements allow us to study the complicated {\\it three-dimensional} structure of the interstellar medium (ISM) behind and near Cassiopeia A. The ISM consists of sheets and filaments, with many structures on a scale of a parsec or less. The spatial power spectrum of the ISM appears to be similar to that of fractals with a spectral index of 3.5. The filling factor for the small structures above the spatial wavenumber \\textit{k} $\\sim $ 0.5 cycles pc$^{-1}$ is only $\\sim $ 0.4\\%. ",
        "introduction": "The composition and structure of the interstellar medium (ISM) are critical to our understanding of the evolution of stars and galaxies. The gas is shaped by turbulence and magnetic fields and is highly structured (Verschuur 1995). Instabilities cause the gas to collect into compact filaments and sheets that collapse into stars (e.g., V\\'azquez-Semadeni et al. 2007 and references therein; Boss 2005). The interstellar grains account for only 1\\% of the total mass of the ISM, but their interaction with stellar photons is the dominant source of opacity and hence strongly influences star formation and evolution. There have been several observations that give us some clues to the nature of the dust in the ISM: for example, the behavior of Polycyclic Aromatic Hydrocarbons (PAHs) in nearby galaxies (Smith et al. 2007), and the mass or size distribution of dust particles in the ISM (Kim \\& Martin 1996; Mathis et al. 1977), to just name a few (see review by Draine 2003). However, as yet, no fully developed model has been developed for the detailed morphology of the ISM.  An understanding of the ISM structure is necessary to understand the forces and mechanisms that shape it and set the preconditions for star formation and other widely studied processes (e.g., Carlqvist \\& Gahm 1992; Verschuur 1995; Fiege \\& Pudritz 2000). The ISM structure is also a key input to studies of radiative transfer, and therefore is critical to understanding the properties of dust-embedded sources (e.g., Bruzual et al. 1988; Witt \\& Gordon 1996; Wolf et al. 1999; Hegmann \\& Kegel 2003). Specific examples include young stellar object clumpy circumstellar media (e.g., Indebetouw et al. 2006), HII regions (e.g., Wood et al. 2005), and galaxy disks (e.g., Inoue 2005) and references in these works. It has been shown that the assumed structure can affect the wavelength dependence of the attenuation (e.g., Fischera \\& Dopita 2005) and can significantly modify the observed spectral energy distribution of a source (Indebetouw et al. 2006). Clumping can also allow self-shielding of the ISM gas and therefore can alter observed line strengths (e.g. Burton et al. 1990).  The structure of the ISM is conventionally measured using HI observations. Verschuur (1995) reviews these results, emphasizing the prevalence of filamentary structures and the possible causes for them. HI data can identify such features through coherent velocity structures, but it is difficult to relate them accurately to one another along the line of sight. Braun \\& Kanekar (2005) have recently used HI measurements to discover tiny interstellar clouds in the local ISM, with dimensions of 0.1 pc or even less. There have been several efforts to study the fractal structures of the ISM (Elmegreen \\& Falgarone 1996) on subparsec scales (Ingalls et al. 2004; Falgarone et al. 2004, and references therein). Small scale structures have also been probed by observing the scintillation of signals from compact radio sources, indicating that the density fluctuation spectrum is consistent with Kolmorgov turbulence (e.g., Stinebring et al. 2000). More closely related to our study are three-dimensional maps of the dust structures in the foreground of supernova (SN) 1987A (Xu et al. 1995) and of type Ia SN in general (Patat 2005). These studies are based on the highly complex physical characteristics of interstellar dust particles, such as three-dimensional multiple scattering or polarization, with dependencies on both the dust parameters and the geometry. Thus, they must introduce additional parameters that are difficult to constrain and can increase the uncertainties in the results.  Here we present a simple and straightforward method to study the structure of the ISM surrounding the supernova remnant (SNR) Cassiopeia A (Cas A) at a resolution of $\\sim$ 0.1pc. Cas A exploded approximately 10,000 years ago, and the light coming directly from the explosion reached Earth about 325 years ago (Thorstensen et al. 2001). In previous work (Krause et al. 2005), we discovered infrared (IR) features at 24$\\mu$m up to about 25 arcmin from the center of the SNR. These features were apparently moving at superluminal velocities, much faster than the velocity of fast-moving knots around the SNR. Although initially some features appeared to be associated with a more recent event in the core of the SNR, further tracking of their motions now makes this explanation appear unlikely. We believe that all the observed IR features are the echoed reradiation of thermal energy as the light from the Cas A supernova passes through its surrounding ISM and heats the dust. Because of the indirect path to reach us and the resulting increased propagation time, we see these features now. These IR-radiating features, called IR echoes, are a novel approach to explore the distribution and the morphological structure of the ISM in three dimensions, and thus they solve the limitations of the previous integrated measures of the ISM structure in two dimensional, projected images. In this paper, we describe the geometry of the echoes in \\S 2, present observations of them over nearly a three year interval in \\S 3, develop and apply the tools to analyze the data and construct images of the ISM in \\S 4, discuss the results in \\S5, and summarize our conclusions in \\S 6. ",
        "conclusions": "  $ \\bullet $ The detailed behavior of the IR light echoes confirms that they are re-radiation of the absorbed outward-moving SN light pulse by the ISM surrounding Cas A.  $ \\bullet $ The three-dimensional display of the echoes shows different sizes of clumps connected with filaments that trace sheet-like patterns.  $ \\bullet $ The surface filling factor of the heated dust grains is about 0.4\\%, for fine structures (spatial frequencies $\\ge$ 0.5 cycles pc$^{-1}$).  $ \\bullet $ The structures have a power-law power spectrum k$^{-\\beta}$ with index $\\beta$ $\\sim $ 3.5, consistent with fractal structure.  $ \\bullet $ The average ISM density above 0.5 cycles pc$^{-1}$ and below 5 cycles pc$^{-1}$ in the region we have studied is $\\sim 3 \\times 10^{-25}$ kg m$^{-3}$. If we integrate the power spectrum down to $\\sim $ 0.01 cycles pc$^{-1}$, we estimate that the ISM surrounding Cas A is $\\sim 6 \\times 10^{-23}$ kg m$^{-3}$, more dense than the general diffuse ISM by a factor of a few."
    },
    "0801/0801.1106.txt": {
        "abstract": "E+A galaxies have been interpreted as post-starburst galaxies based on the presence of strong Balmer absorption lines combined with the absence of major emission lines ([OII] nor H$\\alpha$). As a population of galaxies in the midst of the transformation, E+A galaxies has been subject to an intense research activity. It has been, however, difficult to investigate E+A galaxies statistically since E+A galaxies are an extremely rare population of galaxies ($<$1\\% of all galaxies in the local Universe).  Here, we present a large catalog of 564 E+A (post-starburst) galaxies carefully selected from half million spectra of the Sloan Digital Sky Survey Data Release 5. We define E+A galaxies as those with H$\\delta$ equivalent width $>$5\\AA\\  and no detectable emission in [OII] $and$ H$\\alpha$. The catalog contains 564 E+A galaxies, and thus, is one of the largest of the kind to date. In addition, we have included the H$\\alpha$ line in the  selection to remove dusty star-forming galaxies, which could have contaminated previous [OII]-based samples of E+A galaxies up to 52\\%. Thus, the catalog is one of the most homogeneous, let alone its size. The catalog presented here can be used for follow-up observations and statistical analyses of this rare population of galaxies. ",
        "introduction": "%\\subsection{post-starburst  interpretation of E+A galaxies} %\u00a3\u00c5\u00a1\u00dc\u00a3\u00c1\u00b6\u00e4\u00b2\u00cf\u00a4\u00cf\u00a5\u00dd\u00a5\u00b9\u00a5\u00c8\u00a5\u00b9\u00a5\u00bf\u00a1\u00bc\u00a5\u00d0\u00a1\u00bc\u00a5\u00b9\u00a5\u00c8\u00a4\u00c7\u00a4\u00a2\u00a4\u00eb\u00a1\u00a3 \u00a4\u00c8\u00b2\u00f2\u00bc\u00e1\u00a4\u00b5\u00a4\u00ec\u00a4\u00c6\u00c3\u00ed\u00cc\u00dc\u00a4\u00f2\u00cd\u00e1\u00a4\u00d3\u00a4\u00c6\u00a4\u00ad\u00a4\u00bf\u00a1\u00a3 \\citet{1983ApJ...270....7D,1992ApJS...78....1D} found galaxies with mysterious spectra while investigating high redshift cluster galaxies. The galaxies had strong Balmer absorption lines with no emission in [OII]. These galaxies are called  ``E+A'' galaxies since their spectra looked like a superposition of that of elliptical galaxies (Mg$_{5175}$, Fe$_{5270}$ and Ca$_{3934,3468}$ absorption lines) and that of A-type stars (strong Balmer absorption lines) \\footnote{Since some of E+A galaxies are found to have disk-like morphology \\citep{1994ApJ...430..121C,1994ApJ...430..107D,1997AJ....113..492C,1999ApJS..122...51D}, these galaxies are sometimes called ``K+A'' galaxies.  However, \\citet{2003PhDT.........2G} found that their E+A sample with higher completeness has early-type morphology. Following this discovery, we call them as ``E+A'' throughout this work.}. The existence of strong Balmer absorption lines shows that these galaxies have experienced starburst recently \\citep[within a gigayear;][]{2004A&A...427..125G}.  However, these galaxies do not show any sign of on-going star formation as non-detection in the [OII] emission line indicates. Therefore, E+A galaxies have been interpreted as a post-starburst galaxy, that is,  a galaxy which truncated starburst suddenly \\citep{1983ApJ...270....7D,1992ApJS...78....1D,1987MNRAS.229..423C,1988MNRAS.230..249M,1990ApJ...350..585N,1991ApJ...381...33F,1996ApJ...471..694A,1999ApJS..122...51D,1999ApJ...518..576P,2004ApJ...617..867D,2003PASJ...55..771G,2004A&A...427..125G,2005MNRAS.357..937G,goto2006}. However, the reason why they started starburst, and why they abruptly stopped starburst remain one of the mysteries in galaxy evolution.   %%\u00a4\u00b3\u00a4\u00b3\u00a4\u00cb\u00a6\u00c1\u00a4\u00c7\u00a4\u00ef\u00a4\u00ab\u00a4\u00c3\u00a4\u00bf\u00a4\u00b3\u00a4\u00c8\u00a4\u00e2\u00bd\u00f1\u00a4\u00af\u00a1\u00a3\u00a4\u00b3\u00a4\u00b3\u00a4\u00cf\u00a5\u00ec\u00a5\u00d5\u00a5\u00ea\u00a1\u00bc\u00a4\u00cb\u00c4\u00b9\u00a4\u00a4\u00a4\u00c8\u00b8\u00c0\u00a4\u00ef\u00a4\u00ec\u00a4\u00bf\u00a4\u00ce\u00a4\u00c7\u00ba\u00ef\u00bd\u00fc\u00a1\u00a3 %Recent study found that E+A galaxies have $\\alpha$-element excess \\citep{goto2006}, which also strongly supported  the post-starburst interpretation of E+A galaxies. %%\u00a6\u00c1\u00a4\u00ce\u00b2\u00f2\u00c0\u00e2\u00a1\u00a3 %% The abundance of  $\\alpha$-element  can carry crucial information about the star formation time-scales. %$\\alpha$-elements such as Magnesium\\footnote{$\\alpha$-element includes O, Mg, Si, S, Ca and T.} are formed in the explosion of type II supernovae which occur rapidly after a burst of star formation (time scales of $10^6-10^7$ yr) whereas the iron peak elements originate in Type Ia supernovae which lag behind by at least 1 gigayear \\citep{1984ApJ...286..644N,1995ApJS..101..181W}. %Therefore, a super-solar $\\alpha$-element to Fe ratio occurs when an initial star burst formed Type II supernova abundance pattern and then the star formation ceases before type Ia SNe create Fe. %, perhaps through the onset of galactic winds sweeping away the gas material (Greggio \\& Renzini 1983; Matteucci \\& greggio 1986; Thomas et al. 1998). % Therefore, observed $\\alpha$-element excess \\citep{goto2006} can be interpreted yet another evidence of E+A galaxies being in the post-starburst phase. %% If E+A galaxies are truly in the post-starburst phase as has been claimed before \\citep[e.g.,][]{1999ApJ...518..576P}, E+A galaxies are expected to have an enhanced $\\alpha$-element ratio, where Type II SNe produced $\\alpha$-elements such as Mg, but the Type I SNe has not yet exploded to produce enough Fe. % % Taking all together, there seem to be abundant evidence that E+A galaxies are truly in the post-starburst phase. %  Nevertheless, the reason why they started starburst, and why they abruptly stopped starburst remain one of the mysteries in galaxy evolution. %%  Nevertheless, the reason why they started starburst, and why they abruptly stopped starburst remain one of the mysteries in galaxy evolution. %As a population of galaxies that is undergoing the galaxy transformation from starburst galaxies to quiescent elliptical galaxies, E+A galaxies have been a subject of extensitve studies. %%\u00b8\u00bd\u00ba\u00df\u00a4\u00de\u00a4\u00c7\u00a4\u00ce\u00b8\u00a6\u00b5\u00e6\u00a4\u00ce\u00a4\u00de\u00a4\u00c8\u00a4\u00e1\u00a4\u00f2\u00a4\u00b3\u00a4\u00b3\u00a4\u00cb\u00bd\u00f1\u00a4\u00af\u00a1\u00a3\u00a4\u00b3\u00a4\u00b3\u00a4\u00cf\u00c4\u00b9\u00a4\u00a4review.  %\\subsection{Are E+As created in galaxy clusters?} %%\u00ba\u00c7\u00bd\u00e9\u00a4\u00cf\u00b6\u00e4\u00b2\u00cf\u00c3\u00c4\u00b5\u00af\u00b8\u00bb\u00a4\u00c8\u00b9\u00cd\u00a4\u00a8\u00a4\u00e9\u00a4\u00ec\u00a4\u00c6\u00a4\u00a4\u00a4\u00bf\u00a1\u00a3 %  At first, E+A galaxies are found in cluster regions, both in low %  redshift clusters \\citep{1983ApJ...270....7D,1993AJ....106..473C,1996AJ....111...78C,1997AJ....113..492C,2001AJ....121.2331C,2001AJ....121..793R} and high redshift   clusters \\citep{1985MNRAS.212..687S,1986ApJ...304L...5L,1987MNRAS.229..423C,1988MNRAS.235..827B,1991ApJ...381...33F,1995A&A...297...61B,1996MNRAS.279....1B,1998ApJ...498..195F,1998ApJ...507...84M,1998ApJ...497..188C,1999ApJS..122...51D,2003ApJ...599..865T}. Therefore a %  cluster specific phenomenon was thought to be responsible for the %  violent star formation history of E+A galaxies. % A ram-pressure stripping model % %\\citep{1951ApJ...113..413S,1972ApJ...176....1G,1980ApJ...241..928F,1981ApJ...245..805K,1999MNRAS.308..947A,1999PASJ...51L...1F,2000Sci...288.1617Q,2003ApJ...584..190F,2004PASJ...56..621F}%\u00a4\u00b3\u00a4\u00ec\u00a4\u00e9\u00a4\u00cf\u00ba\u00c7\u00bd\u00e9\u00a4\u00cestarburst\u00a4\u00f2\u00b5\u00c4\u00cf\u00c0\u00a4\u00b7\u00a4\u00c6\u00a4\u00a4\u00a4\u00ca\u00a4\u00a4\u00a1\u00a3 %\\citep{1983ApJ...270....7D,2003ApJ...596L..13B} %  can first accelerate star formation of cluster galaxies and later turn it off. %Galaxy-galaxy merger/interaction in cluster regions is also a good candidate to %  explain E+A phenomena \\citep{1984ApJ...287...95L,1987AJ.....93.1011K,1988ApJ...325...74S,1988ApJ...324..112T,1988ApJ...330..596L,1994ApJ...426..524L,1992AJ....104.2067L}  although several authors pointed out %  that merging/interaction is difficult to %  happen in cluster core regions since relative velocities of galaxies are too %  high \\citep{1980ComAp...8..177O,1987gady.book.....B,1992ApJ...401L...3M,1997ApJ...481...83M}. %   Other candidate mechanisms include galaxy %  harassment \\citep{1996Natur.379..613M,1999MNRAS.304..465M}, cluster tidal forces \\citep{1990ApJ...350...89B,1993ApJ...408...57V,1998ApJ...509..587F}, removal and consumption of the  gas  \\citep[strangulation;][]{1980ApJ...237..692L,2001ApJ...547L..17B,2002MNRAS.333..768M,2003MNRAS.342..664O}, and evaporation of the cold gas in disk galaxies via heat %  conduction from the surrounding hot intra-cluster-medium \\citep[ICM;][]{1977Natur.266..501C,2003ApJ...584..190F,2004PASJ...56..621F}. %\\citet{1999PASJ...51L...1F} discussed that a cluster-cluster %  merger can increase the ram-pressure, and thus, can increase the %  fraction of post-starburst galaxies such as E+As. % %%\u00a4\u00bf\u00a4\u00c0truncate \u00a4\u00b9\u00a4\u00eb\u00a4\u00c0\u00a4\u00b1\u00a4\u00c7\u00a4\u00cfE+A\u00a4\u00cb\u00a4\u00ca\u00a4\u00e9\u00a4\u00ca\u00a4\u00a4\u00a4\u00ab\u00a4\u00e9BO\u00a4\u00cf\u00b4\u00d8\u00b7\u00b8\u00a4\u00ca\u00a4\u00a4\u00a4\u00c8\u00a5\u00ec\u00a5\u00d5\u00a5\u00ea\u00a1\u00bc\u00a4\u00cb\u00b8\u00c0\u00a4\u00ef\u00a4\u00ec\u00a4\u00bf\u00a4\u00ce\u00a4\u00c7\u00b2\u00bc\u00a4\u00cf\u00ba\u00ef\u00bd\u00fc\u00a4\u00b9\u00a4\u00eb\u00a1\u00a32007/05/07 %%    Observationally strong evolution in colors of cluster galaxies has been discovered, in a  sense that there are more %%  blue galaxies in higher redshift clusters \\citep[the so-called %%  Butcher-Oemler effect;][]{1978ApJ...226..559B,1987MNRAS.229..423C,2000ApJ...540..715R,1994ApJ...430..121C,1998ApJ...497..188C,2000AJ....119.1562M,2001ApJ...548L.143M,2001ApJ...547..609E,2001MNRAS.321...18K,2003PASJ...55..739G}. %%  As a population of galaxies in the middle of the transformation, E+A galaxies appeared %%  to fit in the Butcher-Oemler effect naturally, as blue, star-forming %%  galaxies at high redshifts turn off their star formation abruptly, %%  changing their spectra to E+A type, and they become red, %%  non-star-forming galaxies at low redshifts. In addition, % If E+As are cluster %  related, it is   of considerable interest to elucidate the connection to the spiral-to-S0 morphological  evolution of cluster galaxies recently observed \\citep{1984ApJ...281...95P,1997ApJ...490..577D,2000ApJ...542..673F,2005ApJ...623..721P,2003PASJ...55..739G,2005ApJ...621..188G}. %   E+A galaxies might be a critical link in a galaxy evolution in %  which a blue, star-forming disk galaxy evolves into a red S0 or %  elliptical galaxy. % %\u00a4\u00b7\u00a4\u00ab\u00a4\u00b7\u00b6\u00e4\u00b2\u00cf\u00c3\u00c4\u00a4\u00c7\u00a4\u00cf\u00a4\u00ca\u00a4\u00a4\u00a4\u00c8\u00a4\u00ef\u00a4\u00ab\u00a4\u00c3\u00a4\u00bf\u00a4\u00c8\u00a4\u00b3\u00a4\u00b3\u00a4\u00cb\u00bd\u00f1\u00a4\u00af\u00a1\u00a3 % %However, recently, \\citet{2004MNRAS.355..713B} found that low redshift %E+A systems are located predominantly in the field based on the 56 E+A %galaxies selected from the 2dF Galaxy Redshift Survey (2dFGRS). %\\citet{2005MNRAS.357..937G} showed that the environment of E+A galaxies is consistent with that of the field galaxies based on a larger sample of 266 E+A galaxies carefully selected from the SDSS, i.e., E+A galaxies do not preferentially exist in the galaxy cluster environment in the local Universe. %%This work highlights the importance to use a large and homogeneous sample of E+A galaxies, especially with the H$\\alpha$ in the selection criteria. %  In the local Universe,  majority of E+A galaxies in the field regions cannot be explained by the cluster related physical mechanisms such as the ram-pressure stripping, or tidal interaction with the cluster potential. % % However, the situation is different in higher redshift Universe, especially in the cluster environments. %  \\citet{2003ApJ...599..865T}  found 7-13\\% of E+A galaxies in three high redshift clusters at z=0.33,0.58, and 0.83, claiming that $>30$\\% of E+S0 members may have undergone the E+A phase if the effects of E+A downsizing and increasing E+A fraction as a function of redshift are considered. In their search for field E+A galaxies among 800 spectra, \\citet{2004ApJ...609..683T}   measured the E+A fraction at 0.3 $< z <$ 1 to be 2.7\\% $\\pm$ 1.1\\%, a value lower than that in galaxy clusters at comparable redshifts. % MORPHS XXX here. % However note that  based on 1823 galaxies in the 15 X-ray luminous clusters at 0.18 $< z <$ 0.55, on the contrary, \\citet{1999ApJ...527...54B} found that the fraction of K+A galaxies is 1.5\\% $\\pm$ 0.8\\% in the cluster and 1.2\\% $\\pm$ 0.8\\% in the field, i.e., they found no difference. % Taken together with the much smaller fraction of E+A galaxies in the local field, the sudden increase of the fractions in distant clusters over all other environments suggests that there is something unique to this time and the environment that gives rise to so many E+A galaxies. % It is also important to note that E+A galaxies may have heterogeneous origin: in the local Universe, E+As are preferentially found in the rarefied field regions, whereas at higher redshift, E+As are found in the cluster environment. Considering the large difference in time and the environment, E+As in the local field, and high-z clusters may have been created by different physical mechanisms. % % % \\subsection{Are E+As dusty-starburst galaxies? } % % %%\u00a5\u00c0\u00a5\u00b9\u00a5\u00c6\u00a5\u00a3\u00a5\u00b9\u00a5\u00bf\u00a1\u00bc\u00a5\u00d0\u00a1\u00bc\u00a5\u00b9\u00a5\u00c8\u00a4\u00ab\u00a4\u00e2\u00a4\u00b7\u00a4\u00ec\u00a4\u00ca\u00a4\u00a4\u00a1\u00a3 % % However, there remain other possibilities: %%    E+A galaxies could also  be explained as dusty-starburst galaxies. In this scenario, E+A galaxies are not post-starburst galaxies, but in reality star-forming galaxies whose emission lines are invisible in optical wavelengths due to the heavy obscuration by dust. %%   \\citet{1999ApJ...525..609S}  performed a radio observation to investigate the star formation activity hidden by the dust, and detected 5 galaxies in radio (out of 8), which have strong Balmer absorption with no detection in [OII]. %%    \\citet{1999AJ....118..633O} investigated the radio properties of galaxies in %%   a rich cluster at $z\\sim$0.25 (A2125) and found that optical line %%   luminosities (e.g., H$\\alpha$+[NII]) were often weaker than one would %%   expect for the star formation rates (SFRs) implied by the radio emission. %%     As a variant,  \\citet{2000ApJ...529..157P} presented the selective dust extinction %%    hypothesis, where dust extinction is dependent on stellar age since %%     youngest stars inhabit very dusty star-forming HII regions while older %%     stars have had time to migrate out of such dusty regions. %%      If O, B-type stars in E+A galaxies are embedded in dusty regions and %%     only A-type stars have long enough lifetimes ($\\sim$ 1 gigayear) to move out from such regions, %%     this scenario can naturally explain the E+A phenomena. %%    \\citet{2006ApJ...649..163B} detected large amount of neutral hydrogen ($10^9 M_{\\odot}$)in four nearby E+A galaxies, suggesting that some E+A galaxies are observed in the inactive phase of their star formation duty cycle. %% % Therefore, there still remains a possibility that E+A galaxies may not be post-starburst galaxies. %%  Accordingly, an independent evidence of the post-starburst event in addition to the strong Balmer absorption has been awaited. % %%\u00a5\u00c0\u00a5\u00b9\u00a5\u00c6\u00a5\u00a3\u00a1\u00bc \u00c0\u00e2\u00a4\u00ce2\u00c8\u00d6\u00cc\u00dc. %    Another possible  explanation for E+A phenomena is dust enshrouded star formation, where %    E+A galaxies are actually star-forming, but emission lines are invisible in optical wavelengths due to the heavy obscuration by dust. %   %This effect is particularly important at high redshifts where %  %$UV$ light is often used to select sample galaxies. %   \\citet{1999ApJ...518..576P} pointed out that galaxies with strong H$\\delta$ and weak %  [OII] emissions (e(a) galaxies in their classification) might be a %  dust-obscured starbursting galaxy. % %%  $(iv)\\ Selective\\ Extinction$. %     As a variation of the dust enshrouded star-forming scenario, %   \\citet{2000ApJ...529..157P} presented a selective dust extinction %   hypothesis, where dust extinction depends on stellar age since %the %    youngest stars inhabit very dusty star-forming HII regions while older %    stars have had time to migrate out of such dusty regions \\citep[also see ][]{1994ApJ...429..582C,2001ApJ...550..195P}. %%     Poggianti et al. (2001) performed spectroscopy of very luminous %%    infrared galaxies and found the evidence for the selective dust %%    extinction. %   %, where dust extinction is dependent on stellar age since the %   % youngest stars inhabit very dusty star-forming HII regions while older %   % stars have had time to migrate out of such dusty regions (also see %   % Calzetti, Kinney, \\& Storchi-Bergmann 1994; Poggianti \\& Wu 2000). %     If O, B-type stars in E+A galaxies are embedded in dusty regions and %    only A-type stars have long enough lifetimes (10$^7\\sim 1.5\\times %    10^9$ yr) to move out from such regions, %    this scenario can naturally explain the E+A phenomena. However, this %    scenario has to explain why the selective extinction happens only in %    E+A galaxies since there exist many star-forming galaxies with %    detectable [OII] and H$\\alpha$ emissions. % %   A straightforward test for this scenario is to observe in radio wavelengths %  where dust obscuration is negligible \\citep[or in far-infrared/sub-millimeter to detect emission from dust;][]{2005MNRAS.360..322G}. At radio wavelengths, the synchrotron radiation   from electrons accelerated by supernovae can be observed. Therefore, %  in the absence of a radio-loud active nucleus, the radio flux of a %  star-forming galaxy can be used to estimate its current massive star %  formation rate \\citep{1992ARA&A..30..575C,2003ApJ...599..971H}. %   \\citet{1999ApJ...525..609S}  performed such a radio observation and found that %  among 8 galaxies detected in radio, 5 galaxies have strong Balmer %  absorption with no detection in [OII]. % %showed %  %a post-starburst %  %signature in optical spectra. %  They concluded that massive stars are  currently forming in these 5 galaxies. %   \\citet{1999AJ....118..633O} investigated the radio properties of galaxies in %  a rich cluster at $z\\sim$0.25 (A2125) and found that optical line %  luminosities (e.g., H$\\alpha$+[NII]) were often weaker than one would %  expect for the star formation rates (SFRs) implied by the radio emission. % %  On the other hand, %\\citet{2001AJ....121.1965C} detected none of 5 nearby E+A galaxies in radio %  continuum using VLA and excluded the possibility that their E+As are %  dust-enshrouded massive starburst galaxies. %    \\citet{2002AJ....124.2453M} observed radio continua of 15 E+A galaxies found in LCRS   and detected moderate levels of star formation %  in two of them. The star formation rates of these two galaxies, however, are %  5.9 and 2.2 M$_{\\odot}$ yr$^{-1}$, which are an order of magnitude %  smaller than those in \\citet{1999ApJ...525..609S}, and consistent with normal to low SFR instead of starburst. %    Recently, \\citet{2004A&A...427..125G} performed 20-cm radio continuum observations of 36 E+A galaxies using the VLA, and found that none of their E+A galaxies are detected at 20 cm, suggesting that E+A galaxies are not dusty-starburst galaxies. % %   These studies warn us that it is important to acquire %  data in a broad wavelength range. %Although we do not have radio information, %  Based on the past observations, E+A galaxies selected with the H$\\alpha$ in the selection criteria are not detected in the radio continuum, and do not seem to be dusty-starburst galaxies. % % % % % %\\subsection{Are E+As the progenitor of the present-day elliptical galaxies?} %%\u00c2\u00ca\u00b1\u00df\u00b6\u00e4\u00b2\u00cf\u00a4\u00ce\u00c1\u00c4\u00c0\u00e8\u00a4\u00ab\u00a4\u00e2\u00a4\u00b7\u00a4\u00ec\u00a4\u00ca\u00a4\u00a4\u00a1\u00a3 % In the mean time, there has been a claim that E+A galaxies could be a progenitor of present day elliptical galaxies \\citep{2004ApJ...602..190Q,2006ApJ...650..763H}. %%It has been well-established that local elliptical galaxies shows $alpha$-element\\footnote{$\\alpha$ includes O, Mg, Si, S, Ca, and Ti.} enhancement compared to the theoretical expectation, indicating they might have experienced the truncation of starburst. % probably because of the dominant supernovae (SNe) over SNe Ia enrichment in these galaxies (Gonzales 1993; Worthey et al. 1992). % It has been found that E+A galaxies have early-type morphological appearance \\citep{2005MNRAS.357..937G,1991ApJ...381L...9O,2005MNRAS.359.1557Y,2005MNRAS.359.1421P,2006ApJ...650..763H}. %The metallicity and the $\\alpha$-element abundance of E+A galaxies are similar to those of present day elliptical galaxies \\citep{goto2006}. Since it is important to understand the evolutionary path of elliptical galaxies abundant in the present Universe, it is important to investigate how E+A galaxies contribute to the overall evolution of galaxies in the whole Universe. % % %\\subsection{Merger/interaction origin of E+A galaxies}\\label{merger} %Recently, \\citet{2005MNRAS.357..937G} has shown that E+A galaxies have more close companion galaxies than average galaxies, showing that the dynamical merger/interaction could be the physical origin of E+A galaxies. Since the dynamical computer simulations have shown that the dynamical merger can create an elliptical galaxy \\citep{1978MNRAS.184..185W,1992ApJ...393..484B,2001ApJ...547L..17B,2001ApJ...558...42S}, E+A galaxies are a perfect candidate of the progenitor of present-day elliptical galaxies, although two orders of magnitude of difference in the number density needs to be carefully considered. % %%\u00b5\u00af\u00b8\u00bb\u00a4\u00cb\u00a4\u00c4\u00a4\u00a4\u00a4\u00c6\u00a4\u00ce\u00cf\u00c3\u00a1\u00a3 % Recently, important observational evidence on the physical origins of E+A galaxies have been collected. \\citet{2005MNRAS.357..937G} showed that E+A galaxies has excess number of close companion galaxies within 100 kpc, suggesting that the dynamical merger/interaction with close companion galaxies may be the cause of the post-starburst phenomena. In addition to this, it has been reported that some E+A galaxies have disturbed appearance \\citep{2006astro.ph.12053L}. % \\citet{2006ApJ...642..152Y,2006AJ....131.2050Y} showed that post-starburst region in the E+A galaxies were extended more than 5kpc from the centre, being consistent with the merger/interaction origin. \\citet{2005MNRAS.359.1557Y} found that irregular `colour morphologies'- asymmetrical and clumpy patterns - at the centre of $g-r$ and $r-i$ 2D colour maps of E+A galaxies with positive slopes of colour gradient. %\\citet{2006astro.ph.12053L}  presented that a text-book example of an end stage of galaxy major-merger system G515 (J152426.55+080906.7), is a $\\sim 1$ gigayear old post-merger, post-starburst system \\citep[see also ][]{1991ApJ...381L...9O}. %Their results favour the hypothesis that E+A galaxies are post-starburst galaxies caused by merger/interaction, having undergone a centralized violent starburst. %  %\\subsection{A need for large and clean E+A catalog} In all the above cases, One of the major difficulties in investigating E+A galaxies has been their rarity: less than 1\\% galaxies are in the E+A phase in the present Universe \\citep{2003PASJ...55..771G}. In this work, we solve this problem by creating a large catalog of 564 E+A galaxies carefully selected from the fifth public data release of the Sloan Digital Sky Survey \\citep[SDSS;][]{2006ApJS..162...38A}. In addition to its size, due to the availability of the information on the H$\\alpha$ line, this  catalog is one of the most homogeneous of the kind, nicely removing the contanimation from the H$\\alpha$ emitting galaxies. The catalog is publicly released\\footnote{ http://www.ir.isas.jaxa.jp/$^{\\sim}$tomo/ea5} so that world-wide researchers can utilize for statistical analyses and follow-up observations in various wavelength. Unless otherwise stated, we adopt the best-fit WMAP cosmology: $(h,\\Omega_m,\\Omega_L) = (0.71,0.27,0.73)$ \\citep{2003ApJS..148....1B}.   %-Blake et al. %-Tran et al.2003,2004 %-Liu et al. %- redshift histogram. recommend z>0.05 cut. aperture bias plot. %- Hdelta histogram. %- fraction as a function of z %- aperture bias %- pretty pictures %- hd vs ha ",
        "conclusions": "\\subsection{Redshift Distribution}  \\begin{figure} \\begin{center} \\includegraphics[scale=0.5]{070124_ea_zhist.ps} \\end{center} \\caption{ Redshift distribution of the E+A galaxies. }\\label{fig:zhist} \\end{figure}    \\begin{figure} \\begin{center} \\includegraphics[scale=0.5]{070124_fiber_physical_ratio.ps} \\end{center} \\caption{ The ratio of the physical galaxy size (taken from the petrosian 50\\% light radius) to the SDSS fiber size is shown as a function of redshift for the E+A galaxies. }\\label{fig:070124_fiber_physical_ratio.ps} \\end{figure}   In Figure \\ref{fig:zhist}, we show the redshift distribution of the E+A galaxies. The median redshift is z=0.138.  There is a possible peak at the lowest redshift bin ($z\\sim 0.04$). The reason for the peak is not certain. We have checked all the images of $z<0.05$ E+A galaxies by eye to find that all of them show strong H$\\delta$ absorption, confirming that our selection criteria work properly. We, however, suspect the aperture bias may be one of the causes; the SDSS fiber spectrograph only samples the light within the inner 3'' \\citep{2002AJ....124.1810S}. This does not bring a large bias into the distant E+A galaxies, whose sizes are smaller (the median Petrosian 50\\% radius of our sample is 1.4''), but it does bring a large aperture bias to nearby E+A galaxies due to their large apparent size. At $z\\sim 0.04$, the median physical size (taken from the Petrosian 50\\% light radius) of E+A galaxies is 2.0'', whereas the 3''-fiber only collects the light from the inner 1.5'' of radius. In Figure \\ref{fig:070124_fiber_physical_ratio.ps}, we plot the ratio of Petrosian 50\\% light radius to the fiber radius (1.5'') as a function of redshift. The figure shows that at $z\\sim 0.04$, there are significant number of E+A galaxies a few times larger than the fiber size. Therefore, we recommend to use a low redshift cut \\citep[e.g., $z>0.05$; see also ][]{2003ApJ...584..210G,2003MNRAS.346..601G} when one performs statistical analysis on the sample.  These nearby E+A galaxies may have remaining star-formation activity outside of the 3'' fiber. Nevertheless, it is also true that these galaxies have a post-starburst stellar population within the 3'' of radius. It is still important to investigate what caused the central post-starburst in these galaxies, and thus, we keep these nearby E+A galaxies in our sample. Due to their larger apparent size, these E+A galaxies are suitable targets for spatially-resolved observations \\citep[e.g., ][]{2006ApJ...642..152Y,2006AJ....131.2050Y}.  In Figure \\ref{fig:rad_spectra}, we show four example spectra of E+A galaxies with a large apparent size of Petrosian 90\\% radius of $>14''$. These E+A galaxies are large enough for detailed morphological/substructure studies.  The corresponding $g,r,i$-composite images are shown in Figure \\ref{fig:rad_images}. There is a hint of possible dynamical disturbances in all four galaxies shown here.   % \\subsection{Aperture bias} %Yagi et al. %We caution readers that      \\begin{figure*} \\begin{center} \\includegraphics[scale=0.45]{SDSSJ161330.2+510336_z=0.034.ps} \\includegraphics[scale=0.45]{SDSSJ121333.0+142900_z=0.064.ps} \\includegraphics[scale=0.45]{SDSSJ155435.5+291320_z=0.094.ps} \\includegraphics[scale=0.45]{SDSSJ123936.0+122620_z=0.041.ps} %\\includegraphics[scale=0.45]{spectra_nearby/SDSSJ084416.6+263303_z=0.079.ps} %\\includegraphics[scale=0.45]{spectra_nearby/SDSSJ163925.0+303710_z=0.106.ps} \\end{center} \\caption{Example spectra of 4 E+As with largest Petrosian 90\\% light radii ($>14''$).  Each spectrum is shifted to the restframe wavelength and smoothed using a 20 \\AA\\ box. The corresponding images are shown in Fig. \\ref{fig:rad_images}. (The image of the same galaxy can be found in the same column/row panel of Fig.\\ref{fig:rad_images}) }\\label{fig:rad_spectra} \\end{figure*}    \\begin{figure*} \\begin{center} \\includegraphics[scale=0.39]{J161330.2+510336_z=0.034.ps} \\includegraphics[scale=0.39]{J121333.0+142900_z=0.064.ps} \\includegraphics[scale=0.39]{J155435.5+291320_z=0.094.ps} \\includegraphics[scale=0.39]{J123936.0+122620_z=0.041.ps} %\\includegraphics[scale=0.39]{images_nearby/J084416.6+263303_z=0.079.ps} %\\includegraphics[scale=0.39]{images_nearby/J163925.0+303710_z=0.106.ps} \\end{center} \\caption{Examples of $g,r,i$-composite images of E+A galaxies with largest Petrosian 90\\% light radii ($>14''$). The corresponding spectra with name and redshift are presented in Fig. \\ref{fig:rad_spectra}. (The spectrum of the same galaxy can be found in the same column/row panel of Fig.\\ref{fig:rad_spectra}) }\\label{fig:rad_images} \\end{figure*}    \\subsection{E+As in the high redshift Universe}  Our catalog of E+A galaxies provides us with a useful benchmark in studying the evolution of E+A galaxies in higher redshift Universe. %Assuming that the lifetime of the E+A phase is $\\sim$1 gigayear based on %the lifetime of A-type stars, only 2-3\\% of galaxies have been through %the E+A phase in the Universe of 13.7 gigayears of age, if the fraction %of the E+A galaxies has been constant at $\\sim$0.2 \\%. This, however, is %a too simple assumption. It has been known that cosmic star formation density was higher at $z>1$ \\citep{1996MNRAS.283.1388M}. According to \\citet{2006ApJ...642...48L}, the fractions of the post-starburst galaxies was larger by a factor of 2 at $z\\sim 1$ \\citep[also see ][]{2006MNRAS.373..349R}.  By comparing the fraction of E+A galaxies in different redshift, we can trace the cosmic star formation history in terms of the post-starburst galaxies.   %If we loosen our criteria to select post-starburst galaxies to H$\\delta$\\_EW $>$3\\AA, the fraction increase by a factor of $\\sim$3 \\citep{2003PASJ...55..771G}. Therefore, an estimate on the high side yield $\\sim$15\\% of galaxies may have been through the E+A phase. However, this is still a small fraction of galaxies compared to the total number of galaxies and suggests that not all the elliptical galaxies have not been through the E+A phase. % Therefore, it is likely that E+A galaxies are only one of the multiple progenitors of elliptical galaxies. % If star-forming galaxies decrease their star formation rate more slowly, for example, exponentially with $\\tau\\sim$1 gigayear, the galaxies do not experience the E+A phase before they become passive galaxies \\citep{2003PhDT.........2G}. %It is important to investigate what special physical mechanism is needed for a galaxy to become an E+A galaxy.  In high redshift cluster environments, E+A galaxies are much more numerous; pioneering work were done by \\citet{1999ApJS..122...51D,1999ApJ...518..576P}, who  found E+A galaxies (H$\\delta$ EW$>$3\\AA and not detectable emission in [OII]) are significantly more common in 10 clusters at 0.37$<z<$0.56 than in the field (21$\\pm$2\\% compared with 6$\\pm$3\\%). Later, \\citet{2003ApJ...599..865T}  found 7-13\\% of E+A galaxies in three high redshift clusters at z=0.33,0.58, and 0.83, claiming that $>30$\\% of E+S0 members may have undergone the E+A phase if the effects of E+A downsizing and increasing E+A fraction as a function of redshift are considered (their selection criteria is $\\frac{H\\delta EW+H\\gamma EW}{2} >4\\AA$ and  [OII] EW $>$-5\\AA ). In their search for field E+A galaxies among 800 spectra, \\citet{2004ApJ...609..683T}   measured the E+A fraction at 0.3 $< z <$ 1 to be 2.7\\% $\\pm$ 1.1\\%, a value lower than that in galaxy clusters at comparable redshifts. \\footnote{Note that  \\citet{1999ApJ...527...54B} found smaller E+A fractions of  1.5\\% $\\pm$ 0.8\\% in high redshift clusters ($z\\sim$0.25).  However, it has been controversial whether there is a discrepany on the fraction of E+A galaxies in the high redshift clusters. See Section 5.1 of \\citet{2004ApJ...617..867D} for detailed discussion.}  % Note that there is an exception; Based on 1823 galaxies in the 15 X-ray % luminous clusters at 0.18 $< z <$ 0.55, on the contrary, % \\citet{1999ApJ...527...54B} found that the fraction of K+A galaxies is % 1.5\\% $\\pm$ 0.8\\% in the cluster and 1.2\\% $\\pm$ 0.8\\% in the field, % i.e., they found no difference. % However, it is not clear if there is a discrepany exists on the fraction of E+A galaxies in the high redshift cluster environment. % Possible causes include difference in cluster sample (differt X-ray luminosity and different redshift), different selection criteria (H$\\delta$ EW of 3 and 5A\\\\ ), and differnt method used in measuring H$\\delta$ EW. See Section 5.1 of \\citet{2004ApJ...617..867D} for more detailed discussion on the possible causes in the difference of E+A fractions in the higher redshift cluster % environment.   Taken together with the much smaller fraction of E+A galaxies in the local field, the sudden increase of the fractions in distant clusters over all other environments suggests that there is something unique to this time and the environment that gives rise to so many E+A galaxies. It is also important to note that E+A galaxies may have heterogeneous origins; in the local Universe, E+As are preferentially found in the rarefied field regions, whereas at higher redshift, E+As are found in the cluster environment. Considering the large difference in time and the environment, E+As in the local field and high-z clusters may have been created by different physical mechanisms. For example, \\citet{2005MNRAS.357..937G} has shown that local (z$\\sim$0.1) E+A galaxies have more close companion galaxies than average galaxies, showing that the dynamical merger/interaction could be the physical origin of E+A galaxies. However, most E+A galaxies in the high redshift clusters are known not to be major mergers \\citep{1999ApJS..122...51D}. %It is important to create a large sample of E+A galaxies at higher redshift spanning a wide range of environment to investigate the evolution of E+A galaxies and its effect on the overall galaxy evolution in the Universe.  %Both of these work did not use H$\\alpha$ line in the selection criteria, and thus, could suffer from the up to 52\\% of the contamination \\citep{2003PASJ...55..771G}. More statistically significant analysis based on a larger sample is needed to shed light on the galaxy evolution in the high redshift cluster environment. % In the low redshift cluster environment, \\citep{2003PASJ...55..757G} identified that passive spiral galaxies, i.e.,  galaxies with no star formation but with spiral morphology, preferentially reside in the cluster infalling regions. These passive spiral galaxies may be a more important smoking-gun in galaxy evolution in cluster environment.  \\subsection{Selection criteria of E+A galaxies}  We used the selection criteria of H$\\delta$EW $>5\\AA$ and H$\\alpha$ EW $<-3.0\\AA$,[OII] EW$<-2.5\\AA$ to select 564 E+A galaxies (Section \\ref{catalog}). However, we caution readers that the number (or fraction) of the selected E+A galaxies strongly depends on the selection criteria. For example, if we loosen our criteria to H$\\delta$ EW$>4\\AA$, the number of E+A galaxies increases to 1062 (almost twice as numerous). With H$\\delta$ EW$>3\\AA$, 2298 E+A galaxies are selected (4 times more numerous). Therefore, when one compares samples of E+A galaxies selected from different data sets, it is important to synchronize the selection criteria. Especially, one compares to a sample at higher redshift, inconsistency in selection criteria can produce spurious  evolutionary effect."
    },
    "0801/0801.1372_arXiv.txt": {
        "abstract": "We discuss some implications of our recent detection of extragalactic \\hhhop: the location of the gas in M82, the origin of energetic radiation in M82, and the possible feedback effects of star formation on the cosmic ray flux in galaxies. ",
        "introduction": "Last year saw the first detection of the \\hhhop\\ molecule outside the Galaxy \\citep{vdtak:xgal}. Using the new 16-pixel HARP imaging spectrometer on the James Clerk Maxwell Telescope, line emission at a rest frequency of 364\\,GHz was detected towards the prototypical starburst galaxy M82 and the prototypical ultraluminous merger Arp~220. The derived \\hhhop\\ abundances and \\hhhop/\\hho\\ ratios imply very high ionization rates for the dense molecular gas in the nuclei of these two galaxies. Chemical models \\citep{meijerink:grid} indicate that the origin of this high ionization is irradiation by X-rays in the case of Arp~220, whereas for M82, a combination of ultraviolet light and cosmic rays is needed. The high ionization rates of these galactic nuclei make magnetic fields more effective in retarding their gravitational collapse and the formation of stars. The origin of the irradiation in M82 is the evolved starburst \\citep{schreiber:m82}, whereas for Arp~220, an AGN may be needed, as claimed before by \\citet{downes:arp220}. Naturally, some questions remain, of which this paper discusses a few. ",
        "conclusions": ""
    },
    "0801/0801.4761_arXiv.txt": {
        "abstract": "We examine the dynamics and X-ray spectrum of the young Type Ia supernova remnant \\of\\ in the context of the recent results obtained from the optical spectroscopy of its light echo. Our goal is to estimate the kinetic energy of the supernova explosion using \\ch\\ and \\xmm\\ observations of the supernova remnant, thus placing the birth event of \\of\\ in the sequence of dim to bright Type Ia supernovae. We base our analysis on a standard grid of one-dimensional delayed detonation explosion models, together with hydrodynamic and X-ray spectral calculations of the supernova remnant evolution. From the remnant dynamics and the properties of the O, Si, S, and Fe emission in its X-ray spectrum we conclude that \\of\\ was originated $\\sim 400$ years ago by a bright, highly energetic Type Ia explosion similar to SN 1991T. Our best model has a kinetic energy of $1.4 \\times 10^{51}$ erg and synthesizes $0.97 \\, \\mathrm{M_{\\odot}}$ of $^{56}$Ni. These results are in excellent agreement with the age estimate and spectroscopy from the light echo. We have thus established the first connection between a Type Ia supernova and its supernova remnant based on a detailed quantitative analysis of both objects. ",
        "introduction": "\\label{sec:Intro}  Astronomical observations can probe the material ejected by supernova (SN) explosions during two transient phases with very different time scales. The initial optical transient (the SN itself) lasts for several months, and the ejecta structure can be studied through the emission and absorption lines produced as the photosphere recedes into the exploded star. After the SN fades away, the ejected material starts to interact with its surroundings, and the supernova remnant (SNR) phase begins. In this phase, the ejecta structure is revealed by the X-ray emission of the material heated by the reverse shock on timescales of hundreds or thousands of years. Both approaches are valid in principle, but up to now the disparity of the timescales involved has made it impossible to verify their mutual consistency by applying them to the same object.  The serendipitous discovery of light echoes associated with the young SNRs \\of, 0519$-$69.0, and N103B in the Large Magellanic Cloud (LMC) by \\citet{rest05:LMC_light_echoes} (henceforth R05) has opened new possibilities for establishing connections between SNRs and their parent SNe. In particular, spectroscopy of the light echoes has the potential to confirm the type of the SN explosion in a straightforward way, avoiding the difficulties inherent to typing SNRs from their X-ray spectra \\citep[for an in-depth discussion of these difficulties, see][]{rakowski05:G337}. In a recent paper, \\citet{rest07:0509_light_echo} (henceforth R07) examined the light echo associated with SNR \\of, and established that this object originated in a Type Ia explosion, in agreement with the longstanding claims based on X-ray observations \\citep[][henceforth WH04]{hughes95:typing_SN_from_SNR,warren03:0509-67.5}. Furthermore, the quality of the light echo was such that it allowed for a meaningful comparison with the spectra of a variety of Type Ia SNe. Based on these comparisons, R07 concluded that the parent event of SNR \\of\\ belonged to the group of overluminous, highly energetic Type Ia SNe whose prototype is SN 1991T \\citep{benetti05:Ia_diversity}.  These results offer an excellent opportunity to revisit the X-ray spectrum of SNR \\of. Recent developments in modeling the thermal emission in young Type Ia SNRs using hydrodynamic calculations and nonequilibrium ionization simulations \\citep[HD+NEI models,][]{badenes03:xray,badenes05:xray} have led to a better understanding of the relationship between the structure of SN ejecta and the X-ray spectra of SNRs. By applying HD+NEI models to the Tycho SNR, \\citet{badenes06:tycho} were able to estimate the kinetic energy and nucleosynthetic yields of the explosion. We can now do the same for \\of, comparing the results of the HD+NEI models with the light echo spectroscopy. In particular, we want to examine to what extent is it possible to determine the brightness of a Type Ia SN by studying the X-ray emission from its SNR hundreds of years after the explosion.  This paper is organized as follows. In \\S~\\ref{sec:DDT}, we begin with a brief overview of delayed detonation models. In \\S~\\ref{sec:Dynamics}, we use these explosion models, together with the forward shock radius and velocity to constrain the dynamics of SNR \\of. In \\S~\\ref{sec:Observations}, we review the \\ch\\ and \\xmm\\ observations of \\of, and we calculate the values for the most relevant line centroids and line flux ratios. In \\S~\\ref{sec:Modeling} we describe the HD+NEI models, and we evaluate them against both the dynamic and spectral constraints posed by the observations of \\of. Finally, we discuss our results and present our conclusions in \\S~\\ref{sec:Disc-Concl}. ",
        "conclusions": "\\label{sec:Disc-Concl}  The light echo observations of R05 and R07 have turned \\of\\ into a unique object: the only X-ray bright SNR whose Type Ia origin is confirmed and whose explosion energy can be estimated based on spectroscopic information from its parent supernova. Motivated by these groundbreaking observations, we have re-examined the dynamics and X-ray spectrum of \\of, applying the HD+NEI modeling techniques introduced in \\citet{badenes03:xray} and \\citet{badenes05:xray} to \\ch\\ and \\xmm\\ observations. We have based our analysis on a grid of one-dimensional DDT explosions, taking advantage of the relationship between ejecta structure and $E_{k}$ in these models to place the birth event of \\of\\ in the sequence of dim to bright Type Ia SNe. Our conclusions are in excellent agreement with the light echo results of R07: SNR \\of\\ was originated by a bright, highly energetic Type Ia event of the subtype often referred to as SN 1991T-like. We thus present, for the first time, a scenario where the dynamics and X-ray spectrum of a Type Ia SNR form a consistent picture with the spectroscopy of its parent supernova.  Our preferred Type Ia explosion model, DDTa, has a kinetic energy of $1.40 \\times 10^{51}$ erg, and synthesizes $1.03 \\, \\mathrm{M_{\\odot}}$ of Fe, $0.09 \\, \\mathrm{M_{\\odot}}$ of Si, $0.07 \\, \\mathrm{M_{\\odot}}$ of S, $0.04 \\, \\mathrm{M_{\\odot}}$ of O, and $0.14 \\, \\mathrm{M_{\\odot}}$ of other elements (mostly Ar and Ca). The peak bolometric luminosity of this model is $\\log L_{\\mathrm{bol}} (\\mathrm{ergs~s^{-1}})=43.39$, and the amount of $^{56}$Ni in the ejecta before nuclear decays is $0.97\\,\\mathrm{M_{\\odot}}$. These values compare very well with the estimates for SN 1991T: $\\log L_{\\mathrm{bol}}=43.36$ \\citep{contardo00:bolometric_light_curves_Ia}; $\\mathrm{M_{^{56}Ni}}=0.87\\,\\mathrm{M_{\\odot}}$ \\citep[][UVOIR method]{stritzinger06:consistent_estimates_56Ni_Ias}, $0.94\\,\\mathrm{M_{\\odot}}$ \\citep[][B-band lightcurve method]{stritzinger06:consistent_estimates_56Ni_Ias,mazzali07:zorro}.  The ejecta density and chemical composition profiles in model DDTa ($\\rho_{AM}=10^{-24}\\,\\mathrm{g \\, cm^{-3}}$, $\\beta=0.02$) can reproduce very well both the dynamics (FS radius and velocity) and the fundamental properties of the X-ray spectrum (O, Si, S, and Fe emission) of \\of\\ at an age of 400 yr. Several dynamical, spectral, and historical arguments indicate that the SNR age cannot be very different from this value, which also coincides with the completely independent estimate derived by R05 from the geometry of the light echo.  These results, together with our previous work on the Tycho SNR \\citep{badenes06:tycho}, constitute firm evidence that the phenomenological one-dimensional DDT models, which have proved so successful in explaining the light curves and spectra of Type Ia SNe, are equally capable of reproducing the fundamental properties of the X-ray emission from young Type Ia SNRs. In particular, it is possible to use HD+NEI simulations based on these DDT models to estimate the kinetic energy (and ultimately the brightness) of a Type Ia SN explosion from the X-ray spectrum of its SNR. Hundreds of years after the SN explosion, the memory of the cataclysmic event persists in the X-rays from its SNR, opening a window onto ages past. This enables us to explore the relationship between individual dim/bright Type Ia SNe and their immediate surroundings (presence or absence of stellar formation, local metallicities, etc.) with much greater detail than is available to extragalactic studies \\citep[see][]{prieto07:SN_Progenitors_Metallicities}. In the context of SNR research, having a good estimate for $E_{k}$ can help to build better models for the impact of CR acceleration on the SNR dynamics. HD+NEI simulations have clearly become a powerful and flexible tool to study the relationship between SNRs and their parent SNe. For interested readers, the synthetic X-ray spectra from our models are available from the authors upon request.  We conclude with a reminder that much is left to do in the study of Type Ia SNe and their SNRs in general, and of \\of\\ in particular. Ongoing direct measurements of the ejecta and FS expansion using grating observations and proper motion studies should improve our knowledge of the dynamics of this SNR in the near future. Further work on the X-ray emission should take the asymmetries of the object into account, in particular the enhanced Fe knots found by WH04. The one-dimensional models presented here can only sample the average or bulk properties of the shocked ejecta. Observational evidence for moderate asymmetries \\citep{fesen07:SN1885,gerardy07:SNIa_midIR} and clumping \\citep{leonard05:spectropolarimetric_diversity_SNIa,wang07:SNIa_spectropolarimetry} in the ejecta of Type Ia SNe is widespread, and SNR studies offer a unique opportunity to study these effects. This might provide crucial constraints for the multi-dimensional simulations of physically-motivated DDT explosions currently under way \\citep{bravo06:PRD,plewa07:DFD,jordan07:GCD_3D}"
    },
    "0801/0801.0374_arXiv.txt": {
        "abstract": "We use synchronous movies from the Dutch Open Telescope sampling the G band, Ca\\,{\\sc ii}\\,H, and \\Halpha{} with five-wavelength profile sampling to study the response of the chromosphere to acoustic events in the underlying photosphere.  We first compare the visibility of the chromosphere in Ca\\,{\\sc ii}\\,H and \\Halpha, demonstrate that studying the chromosphere requires \\Halpha{} data, and summarize recent developments in understanding why this is so. We construct divergence and vorticity maps of the photospheric flow field from the G-band images and locate specific events through the appearance of bright Ca\\,{\\sc ii}\\,H grains.  The reaction of the \\Halpha{} chromosphere is diagnosed in terms of brightness and Doppler shift.  We show and discuss three particular cases in detail: a regular acoustic grain marking shock excitation by granular dynamics, a persistent flasher which probably marks magnetic-field concentration, and an exploding granule.  All three appear to buffet overlying fibrils, most clearly in Dopplergrams. Although our diagnostic displays to dissect these phenomena are unprecedentedly comprehensive, adding even more information (photospheric Doppler tomography and magnetograms, chromospheric imaging and Doppler mapping in the ultraviolet) is warranted. ",
        "introduction": "``Acoustic events'' (or ``seismic events'') is the term used by P.~Goode and coworkers (\\cite{1987SoPh..110..237S}; % \\cite{1992ApJ...387..707G}; % \\cite{1993ApJ...408L..57R}; % \\cite{1995ApJ...444L.119R}; % \\cite{1998ApJ...495L..27G}; % \\cite{2000ApJ...535.1000S}) % to describe localized small-scale happenings in the granulation that produce excessive amounts of upward-propagating acoustic waves in the upper photosphere, with the claim that these indicate the kinetic sources of the global $p$-modes, ``the smoke from the fire exciting the solar oscillations'' (\\cite{1995STIN...9610005G}). % Their technique was to sample the Doppler modulation of a suitable spectral line at different profile heights through Fabry-Perot imaging with rapid wavelength shifting, and then isolate the surface locations with the largest oscillatory amplitude at five-minute periodicity and upward propagating phase.  Their result was that this measure of ``acoustic flux'' tends to be maximum above intergranular lanes, in particular those in which a small granule has vanished in so-called granular collapse.  Acoustic excitation through granular dynamics was also addressed theoretically and through numerical hydrodynamics simulations by Rast (\\citeyear{1995ApJ...443..863R}, % \\citeyear{1999ApJ...524..462R}), % \\citet{2000ApJ...535..464S} % and \\citet{2000ApJ...541..468S}, % confirming the picture of small vanishing granules, especially at mesogranular downdraft boundaries, acting as ``collapsars'' to excite upward-propagating waves, in particular at three-minute periodicity corresponding to the photospheric acoustic cut-off frequency.  In this paper we address the reaction of the overlying chromosphere to such events. \\citet{2000ApJ...541..468S} % suggested that they may also cause \\CaII\\ \\HtwoV\\ and \\KtwoV\\ cell grains, an enigma during many decades (see review by \\cite{1991SoPh..134...15R}) % that was eventually solved by Carlsson and Stein (\\citeyear{Carlsson+Stein1994}, % \\citeyear{1997ApJ...481..500C}) % who identified them as marking upward-propagating acoustic shocks.  The announcement by \\citet{2002AAS...200.5305G} % that indeed acoustic events cause such acoustic grains in the chromosphere was followed up by \\citet{2002A&A...390..681H} % who found that, while extreme acoustic events do tend to correlate with the subsequent appearance of exceptionally bright \\CaII\\ \\KtwoV\\ internetwork grains, this correspondence is far from a one-to-one correlation.  \\begin{figure} \\centering \\includegraphics[width=\\textwidth]{\\figspath/small-mosaic-ca}\\\\[1mm] \\includegraphics[width=\\textwidth]{\\figspath/small-mosaic-ha} \\caption[]{ Simultaneous \\CaIIH\\ and \\Halpha\\ image mosaics taken with the Dutch Open Telescope on 4 October 2005.  The field of view is close to the limb (located off the top) and measures about $265\\times 143$~arcsec$^2$.  Spectral passbands: FWHM 1.4\\,\\AA\\ for \\CaIIH, 0.25\\,\\AA\\ for \\Halpha.  The same area was observed with the SST and described by \\citet{2006ApJ...648L..67V}. % We invite the reader to enlarge this and the following figures on-screen. }\\label{fig:mosaics} \\end{figure}   \\begin{figure} \\centering \\includegraphics[width=\\textwidth]{\\figspath/small-lites} \\caption[]{ Various diagnostic measures of \\CaIIH\\ plotted as space\\,-\\,ime ``timeslice'' evolution plots, assembled by B.W.~Lites from a spectral sequence taken by himself and W.~Kalkofen in 1984 with the Dunn Solar Telescope at the National Solar Observatory/Sacramento Peak. The waves in the network part (bright strip at the center) were analyzed by \\citet{1993ApJ...414..345L}. % The waves in the internetwork columns identified by numbers along the top were numerically simulated by \\citet{1997ApJ...481..500C}. % The spatial coordinate is measured along the slit of the spectrograph.  The first three panels have logarithmic greyscales to reduce the large contrast between network and internetwork. Horizontal greyish ``erasures'' are due to lesser seeing.  First panel: intensity integrated over a spectral band of 0.9\\,\\AA\\ width centered on line center.  Second panel: the same for 0.16\\,\\AA\\ bandpass.  Third panel: intensity of the core-profile minimum. Fourth panel: wavelength variation of the core-profile minimum. Fifth panel: Dopplergram ratio $(R-V)/(R+V)$ with $R$ and $V$ the intensities in 0.16\\,\\AA\\ passbands at \\HtwoR\\ and \\HtwoV. }\\label{fig:lites} \\end{figure}   We now turn to \\Halpha\\ rather than \\HK\\ as diagnostic of the chromospheric response to acoustic events.  It is important to note and to explain that \\Halpha\\ is a much better chromospheric diagnostic than the \\CaII\\ \\HK\\ lines are.  This is not the case for simple Saha-Boltzmann LTE partitioning in which \\HK\\ have larger opacity than \\Halpha\\ throughout the atmosphere (see Figure~6 of \\cite{2006A&A...449.1209L} % and the second student exercise at \\url{http://www.astro.uu.nl/~rutten/education/rjr-material/ssa}), and it is also not the case in the standard NLTE statistical and hydrostatic equilibrium VAL modeling of Vernazza, Avrett, and Loeser (1973, 1976, 1981) \\nocite{VALI} \\nocite{VALII} \\nocite{VALIII} in which the line center of \\CaIIK\\ is formed higher than the line center of \\Halpha\\ (see the celebrated Figure~1 of \\cite{VALIII}).  We use displays from older data here to argue that the real Sun does not conform.  Figure~\\ref{fig:mosaics} compares Dutch Open Telescope (DOT) images using the two diagnostics.  The upper image is taken in Ca\\,II\\,H and shows a mixture of mid- and upper-photospheric contributions (reversed granulation and acoustic grains) and chromospheric contributions (bright plage, network, and straws in hedge rows as\t described by \\cite{2007ASPC..368...27R}). % The lower panel shows the same area observed in H$\\alpha$ line center. Comparison of the two demonstrates unequivocally that in such filter imaging H$\\alpha$ presents a much more complete picture of the chromosphere.  In our opinion, the chromosphere is best defined as the collective of ubiquitous fibrils seen in this line.  Most of these are internetwork-spanning structures, which probably outline magnetic canopies.  Only truly quiet internetwork areas (there are some in the lower part of the image) are not masked off by such canopy fibrils but show much shorter, highly dynamic loops and grainy brightness patterns that are not connected to network (see also \\cite{2007ApJ...660L.169R}). % None of these structures is observed in the upper panel, except for the straws which correspond to bright \\Halpha\\ fibril endings.  What about the spectral passbands in this comparison?  The Doppler cores of \\HK\\ are much narrower than for \\Halpha, but \\HK\\ filter bandpasses are usually wider: FWHM 3\\,\\AA\\ for {\\em Hinode\\/}, 1.4\\,\\AA\\ for the DOT, 0.6\\,\\AA\\ for the Lyot filter at the German VTT (\\eg\\ \\cite{2007A&A...462..303T}), % 0.3\\,\\AA\\ for the Lockheed-Martin filter at the former SVST (\\cite{Brandt+Rutten+Shine+Trujillo1992}; \\cite{1999ApJ...517.1013L}), % and also 0.3\\,\\AA\\ for the Halle filter at the NSO/SP/DST (\\eg\\ \\cite{2004ApJ...604..936B}, % who assigned propagation speeds and mode conversions to chromospheric oscillations naively adopting the VAL3C line-center height difference).  Such \\HK\\ filter imaging adds considerable inner-wing contributions from the upper photosphere, including bright reversed granulation and yet brighter acoustic grains which may wash out the line-center-only contribution from chromospheric fibrils in the internetwork except where these are very bright, as in straws.  The narrowest-band images with high spatial resolution published so far are R.A.~Shine's 0.3\\,\\AA\\ ones in \\citet{1999ApJ...517.1013L}, % which only show reversed granulation and acoustic grains in the internetwork.  Narrow-band \\Kthree\\ spectroheliograms as the 0.1\\,\\AA\\ ones in the collection of \\citet{Title1966a} % have lower angular resolution but do not suggest the presence of similar masses of internetwork \\CaIIK\\ fibrils as seen in \\Halpha\\ at similar resolution, except near active network and plage.  Figure~\\ref{fig:lites} extends the \\CaIIH\\ passband comparisons in Figure~1 of \\citet{1999ASPC..183..383R} and Figure~9 of \\citet{2001A&A...379.1052K} % with additional spectral measures from the ancient but high-quality \\CaIIH\\ spectrogram sequence of \\citet{1993ApJ...414..345L} % also used by \\citet{1997ApJ...481..500C}. % The first three panels show that narrower \\CaIIH\\ passband produces larger network-to-internetwork brightness contrast, but does not significantly change the morphology of the observed scene.  The oscillation patterns in the internetwork, brightest at \\HtwoV, remain visible even in the line-core intensity in the third panel. The line-core shift in the fourth panel confirms the acoustic nature of these patterns.  Thus, even at full spectral resolution the \\CaIIH\\ core in this internetwork sample is dominated by radial three-minute waves, without obliteration by overlying fibrils or fibril-aligned motions. The final panel showing the \\HtwoR$-$\\HtwoV\\ Dopplergram ratio strengthens this conclusion.  This measure comes closest to exhibiting chromospheric fibrilar structure in the form of oscillatory branches jutting out from the network with time, but regular three-minute oscillation patterns dominate further away in the internetwork.   The lower-left panel of Figure~21 of \\citet{2001A&A...379.1052K} % demonstrated the close correspondence between \\CaIIH\\ line-center dynamics and underlying photospheric dynamics in the internetwork parts of these spectrograms in Fourier terms.  In fact, it should be strange if the line-center formation layer would not fully partake in the shock dynamics in these data because substantial in-phase upper-atmosphere downdraft, above the next upcoming grain-causing acoustic shock, is required to give the \\HtwoV\\ and \\KtwoV\\ grains their characteristic spectral asymmetry as shown in the formation breakdown diagrams of \\citet{1997ApJ...481..500C}. % They actually obtained their best sequence reproduction for near-network column 110 in Figure~\\ref{fig:lites}, suggesting that even there the field was too weak to upset \\Hthree\\ acoustic shock signatures.  Thus, these well-studied \\CaIIH\\ spectrograms contradict the presence of opaque chromospheric fibrils masking upper-photosphere (or ``clapotispheric'') internetwork dynamics in \\HK, whereas \\Halpha\\ canopy fibrils ubiquitously do so. The notion that the \\CaII\\ core generally forms higher than the \\Halpha\\ core, or that chromospheric structures should generally be more opaque in \\HK\\ than in \\Halpha, seems questionable.  The observations suggest instead that the fibrils that constitute the \\Halpha\\ chromosphere are either transparent in \\HK\\ or, when opaque, are either pummeled by shocks from below without hindrance (such as magnetic tension) or represent a featureless, very dark blanket located well above the clapotisphere.  In summary, we suspect that much of the \\Kthree\\ internetwork dynamics is dominated by shock modulation at heights around 1000~km, well below the VAL3C \\Kthree\\ formation height of 1800\\,--\\,2000\\,km and far below most \\Halpha\\ canopy fibrils.  This large discrepancy with equilibrium modeling may become understandable with the recent simulation of \\citet{2007A&A...473..625L} % implementing non-instantaneous hydrogen ionization and recombination. Three aspects combine in such an explanation.  Firstly, the simulated chromosphere is pervaded by shocks.  Secondly, the large difference between hydrogen ionization/recombination balancing speed in hot shocks and in cool post-shock gas implies that instantaneous statistical equilibrium does not apply to the latter. \\citet{2002ApJ...572..626C} % already explained why this is so.  The large 10~eV excitation energy of its $n\\is2$ level causes hydrogen to be ionized fast in shocks but to recombine slowly in the cool inter-shock phases so that the ionization degree remains high also in the latter.  Thirdly, the population of \\Halpha's lower $n\\is2$ level is closely coupled to the ion population (through the Balmer continuum; see also \\cite{1994IAUS..154..309R}) % and therefore remains similarly high at low post-shock temperature. Thus, the large $n\\is2$ excitation energy, being the very reason why \\Halpha\\ has very small cool-gas opacity in Boltzmann balancing, slows down the hydrogen recombination in post-shock cool gas so much that its \\Halpha\\ opacity exceeds the LTE value by many orders of magnitude.  Therefore, fibrils can be opaque in \\Halpha\\ whether they are shock-hot or post-shock cool, and they even may be more opaque in \\Halpha\\ than in \\HK\\ at any temperature -- in utter conflict with equilibrium modeling.  The chromospheric non-equilibrium balancing between \\CaII\\ (12~eV ionization energy) and \\CaIII\\ is probably faster (no top-heavy term diagram) but has not yet been analyzed in similar detail; it is of interest to evaluate the opacity ratio between \\HK\\ and \\Halpha\\ in dynamic conditions.  Another point of interest is that chromospheric fibrils seem to be more obvious in \\CaII~8542\\,\\AA\\ than in \\Hthree\\ and \\Kthree, possibly due to larger Doppler sensitivity (\\cite{2007arXiv0709.2417C}). %  Thus, \\Halpha\\ is the key diagnostic of the internetwork chromosphere. Recent analyses of the chromosphere observed at high cadence in \\Halpha\\ line-center with the Swedish 1-m Solar Telescope (SST) have shed new light on chromospheric fibrils, showing that near network and plage ``dynamic fibrils'' tend to have repetitive mass loading at three\\,--\\,five-minute shock periodicity (\\cite{2007ApJ...655..624D}) % whereas the little loops in quiet-sun internetwork are even more dynamic, presumably due to stronger buffeting by shocks where the magnetic field is weak (\\cite{2007ApJ...660L.169R}). %  In this paper we use image sequences from the Dutch Open Telescope (DOT) with lower cadence than \\Halpha\\ really requires (\\cite{2006ApJ...648L..67V}), % but with \\Halpha\\ profile sampling and synchronous co-spatial imaging in other wavelengths including \\CaIIH.  They permit us to isolate specific events in the photosphere and inspect the chromospheric response in the apparent brightness of \\CaIIH\\ and \\Halpha, and particularly in \\Halpha\\ Doppler modulation.  We do this here by presenting three exemplary cases: an acoustic grain, a persistent flasher, and an exploding granule.   \\begin{figure} \\centering \\includegraphics[width=100mm]{\\figspath/small-samples} \\caption[]{ Simultaneous sample images.  In row order: G band, \\CaIIH\\ line center, \\Haouter, \\Halpha\\ line center, and HaDouter\\ and \\HaDinner\\ Dopplergrams with bright implying downdraft.  The G-band image is Fourier filtered to pass only subsonic components.  The \\Halpha\\ images and Dopplergrams are filtered to pass only supersonic components.  The white boxes specify the three subfields analyzed in Figures \\ref{fig:grain}\\,--\\,\\ref{fig:exploder}, from left to right respectively.  The cutout boxes are larger for \\CaIIH\\ and \\Halpha\\ to show more spatial context in the timeslices of Figures \\ref{fig:grain}\\,--\\,\\ref{fig:exploder}. }\\label{fig:samples} \\end{figure} ",
        "conclusions": "Each example shows large morphological difference between the low-photosphere scene in the G-band, the high-photosphere scene in \\CaIIH, and the chromospheric scene in \\Halpha.  The \\Halpha\\ cutouts do not contain internetwork-spanning fibrils in this quiet region; each of the three events therefore shows chromospheric response that would otherwise have been blocked by overlying canopy fibrils.  The responses are not very striking; in no case may one locate the photospheric happening uniquely from its subsequent \\Halpha\\ signature.  The \\Halpha\\ chromosphere sampled by the image cutouts comes close to our tentative description above of the \\CaII\\ \\Hthree\\ and \\Kthree\\ internetwork chromosphere as a featureless opaque blanket pummeled from below.  They may indeed be about the same, as suggested by the \\Halpha\\ downdrafts in Figure~\\ref{fig:grain} above grains requiring \\Hthree\\ downdraft for \\HtwoV\\ profile asymmetry.  The much smaller temperature sensitivity of the \\HK\\ opacity then indeed diminishes the \\HK\\ feature contrast. Thanks to being an excited and narrower line, Ca\\,II\\,8542 combines larger temperature and Doppler sensitivity to such pummeling but at smaller blanket opacity.  Fibrilar structuring appears most clearly in the \\Halpha\\ Dopplergram timeslices.  Our impression is that indeed this quiet \\Halpha\\ chromosphere is continuously pummeled by internetwork waves, gaining much dynamics from these in the quietest areas where the magnetic fields have low tension but yet enough to impose slight fibrilarity. We suspect that slanted internal gravity waves contribute significantly in this pummeling (see Lighthill 1967, pp.~440\\,--\\,443), \\nocite{Lighthill1967} % and together with the primarily vertically growing acoustic shocks cause intricate and fast-changing interference patterns at granular to mesogranular scales.  Only the fiercest events punch through to be individually recognizable directly above their sources.  Our three examples whet the appetite for even more comprehensive multi-diagnostic solar-atmosphere tomography.  The complexity of the panel layout in Figures~\\ref{fig:grain}\\,--\\,\\ref{fig:exploder} already demonstrates that the dynamical coupling between photosphere and chromosphere can only be addressed holistically with diverse diagnostics.  The complexity of the solar scenes shown in these panels and their large differences between diagnostics strengthen this conclusion.  In fact, our data is yet incomplete: we lack the principal acoustic-event measure of ``acoustic flux'' in this analysis, requiring Fabry-Perot Doppler mapping of the photosphere at multiple heights.  Sensitive photospheric magnetic field mapping and ultraviolet chromospheric image and Doppler diagnostics would also be welcome -- preferably all at similar or better angular resolution, cadence, duration, and field size than used here.  In this paper we have limited the discussion to three hand-picked cases from an early set of tomographic DOT image sequences.  The large DOT database collected in the meantime (openly available at \\url{http://dotdb.phys.uu.nl/DOT}) now permits wider statistical study of such various types of events.  Rapid-cadence Fabry-Perot imaging at the SST promises unprecedented \\Halpha\\ diagnostics. Space-mission co-pointing adds EUV diagnostics of the transition region.  Pertinent numerical simulations including realistic formation of \\Halpha\\ are also coming into reach.    \\begin{acks} The DOT is owned by Utrecht University and located at the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrof{\\'{\\i}}sica de Canarias.  We are deeply indebted to V.~Gaizauskas concerning the DOT \\Halpha\\ filter.  We thank M.~Carlsson, \\O.~Langangen, S.E.M.~Keek, J.M.D.~Kruijssen and A.G. de Wijn for inspiring debates.  This research made much use of SST hospitality and of NASA's Astrophysics Data System.  R.J.~Rutten thanks the Leids Kerkhoven-Bosscha Fonds and the organizers of the SOHO\\,19/GONG\\,2007 meeting for travel support. \\end{acks}"
    },
    "0801/0801.2692_arXiv.txt": {
        "abstract": "\\noindent We evaluate the average expansion rate of a universe which contains a realistic evolving ensemble of non-linear structures. We use the peak model of structure formation to obtain the number density of structures, and take the individual structures to be spherical. The expansion rate increases relative to the FRW value on a timescale of 10--100 billion years, because the universe becomes dominated by fast-expanding voids. However, the increase is not rapid enough to correspond to acceleration. We discuss how to improve our treatment. We also consider various qualitative issues related to backreaction. ",
        "introduction": "\\label{sec:intro}  \\paragraph{Observations and acceleration.}  There is overwhelming observational evidence against homogeneous and isotropic models of the universe based on standard general relativity (i.e. the Einstein-Hilbert action in four dimensions) with ordinary matter (i.e. baryons, dark matter, neutrinos and photons). Observations are usually interpreted keeping to the linearly perturbed homogeneous and isotropic Friedmann-Robertson-Walker (FRW) models, in which case it is necessary to modify gravity or introduce matter with negative pressure. The \\LCDM model, which involves the simplest modification of gravity, the cosmological constant, or the equivalent form of matter, vacuum energy, is in agreement with various observations of cosmological distances. In FRW models, distance has a simple correspondence with the expansion rate (and spatial curvature).  However, linearly perturbed FRW models do not describe the non-linear structures present in the real universe. It is also important to recognise that the only evidence for modified gravity or exotic matter comes from observations of cosmological distances, interpreted in terms of the expansion rate. For example, deviations from general relativity have not been observed in the solar system \\cite{Reynaud:2008}. (The Pioneer anomaly is a possible exception. However, it does not agree with the predictions of the modified gravity models which have been developed to explain the cosmological observations.)  The situation is quite different from that of dark matter, for which there are several independent lines of evidence, such as the motions of stars in galaxies, the motions of clusters, the peak structure of the cosmic microwave background (CMB), the early formation of structures, gravitational lensing, as well as direct measures of the matter density combined with the baryon density given by Big Bang Nucleosynthesis. For this reason, constructing alternatives to dark matter requires resort to baroque models \\cite{teves}, if it is possible at all. In contrast, in order to explain the observations without modified gravity or exotic matter, it is only necessary to change the distance scale (or the expansion rate) as a function of redshift.  The measured quantities from which the expansion rate is inferred can be divided into background quantities and perturbations, though measurements of background quantities also involve perturbations, apart from the luminosities of Type Ia supernovae (SNe Ia). (In realistic models with non-linear structures, there is no simple division into background and perturbations.) Most of the information on the expansion rate comes from measures of the background geometry. While the expansion rate has been measured from different physical systems, such as the CMB \\cite{Spergel:2006}, large scale structure \\cite{BAO} and SNe Ia \\cite{SNobs}, they all probe essentially the same distance measure at different redshifts. (The luminosity distance and the angular diameter distance are directly related to each other. In FRW models, they can be expressed in terms of the proper distance and redshift. For the different distance measures in the context of FRW models, see \\cite{Hogg:1999, Cattoen}.)  The details of the expansion rate (or the distance scale) are not well known except in specific models. Model independent constraints are relatively weak. The CMB is mostly sensitive to the angular diameter distance to the last scattering surface, which is related to the position of the peaks in the angular power spectrum \\cite{Efstathiou:1998, Mukhanov:2003}. Type Ia supernovae provide a more direct measure of the expansion rate, but present data is not sufficiently accurate to give detailed information. Quoted constraints on, for example, the equation of state are often driven by the assumed parametrisation, which can be responsible for misleading apparent precision and artificially small confidence level contours. For the importance of parametrisation, see \\cite{Cattoen, trans, Shapiro:2005, Gong:2006, Elgaroy:2006, Linder:2007}; an analysis of the current situation in terms of a piecewise constant equation of state is given in \\cite{Zunckel:2007}. Additional systematic effects, such as metallicity \\cite{Podsiadlowski:2006}, changes in the treatment of dust \\cite{Conley:2007} or the light-curve fitting method \\cite{Seikel:2007} can further degrade the reliability of the SN Ia measurements. It is noteworthy that the quality of fit of the \\LCDM model has decreased with the introduction of each new SN Ia dataset up to and including the 'gold' sample and the ESSENCE data \\cite{Vishwakarma:2005}. This may hint at inadequacy of the \\LCDM description, or it can be indicative of underestimated systematic errors.  Observational constraints from perturbations are much weaker than those from background quantities. The Integrated Sachs-Wolfe (ISW) effect has been detected via cross-correlation of the CMB and matter tracers at different redshifts \\cite{Ho:2008}. This is interpreted as time evolution of the gravitational potential, which may be due to accelerating expansion or spatial curvature. The amplitude is slightly larger than expected in the \\LCDM model, though not significantly so (around 2$\\sigma$). Other constraints from the evolution of perturbations are rather weak \\cite{lineargrowth}. A measure which combines background evolution and the evolution of perturbations is the number count of clusters as a function of redshift: in an accelerating model the number density should rise sharply with increasing redshift. It has been earlier argued that this is not seen in the data, and that the observations instead prefer deceleration \\cite{Blanchard}. However, it seems that limited understanding of gas physics prevents drawing reliable conclusions from cluster counts at present \\cite{Ferramacho:2007}.  In the context of homogeneous and isotropic models, we can with confidence say that the expansion has accelerated within the last few billion years (though it is not necessarily accelerating at present), and determine the overall magnitude of the acceleration \\cite{Shapiro:2005, Gong:2006, Elgaroy:2006, Seikel:2007}. In order for a model to agree with the observations, it is not necessary to reproduce the expansion history of the \\LCDM model in detail, only to produce a roughly similar amount of acceleration in the same era.  \\para{Structure formation and the coincidence problem.}  While there is no evidence apart from the expansion rate for modified gravity or negative pressure matter, we know that the assumption of linearly perturbed homogeneity and isotropy breaks down in the universe at late times. One has to evaluate the effect of non-linear structure formation on the expansion rate (or the distance scale) before concluding that it is necessary to change either gravity or the matter content.  It was suggested in \\cite{Schwarz, Rasanen} that inhomogeneities related to structure formation could be responsible for accelerated expansion (the possibility had been earlier touched upon in \\cite{Buchert:2000, Tatekawa:2001}). This was discussed more concretely in \\cite{Rasanen:2006a, Rasanen:2006b} where it was demonstrated with a toy model how the formation of non-linear structures can lead to average acceleration, even when the expansion locally decelerates (see also \\cite{Kai:2006}). The physical reason is simply that the fraction of volume in faster expanding regions rises, so the average expansion rate can rise. The bigger is the difference between the slower and faster expanding regions, the more rapid is the change as the faster regions take over. It was pointed out in \\cite{Rasanen:2006a, Rasanen:2006b} that the growth of non-linear structures involves a growing variance in the expansion rate, and that structure formation has a preferred time around 10--100 billion years, near the observed acceleration era. The timescale emerges from the CDM power spectrum, essentially from the time of matter-radiation equality encoded in the change of slope of the CDM transfer function. This could solve the coincidence problem of why the acceleration has started recently in cosmic history, something that the \\LCDM model does not explain.  The link between the preferred time in structure formation and the change of the expansion rate was only conjectured in \\cite{Rasanen:2006a, Rasanen:2006b}. The toy model with acceleration involved only two regions instead of a realistic ensemble of structures, and had no link to the preferred time. In the present work, we remedy both shortcomings. We first discuss the effect of structures on the propagation of light and the expansion rate in section 2, and set up the Buchert backreaction formalism in section 3. In section 4 we present our model for a statistically homogeneous and isotropic universe containing an evolving ensemble of non-linear structures. We calculate the average expansion rate and find that it grows relative to the FRW value around 10--100 billion years, though the change is not rapid enough to correspond to acceleration. We explain the reason and consider how improve our treatment. In section 5 we discuss some observational and theoretical issues related to backreaction, and summarise our results. This work is a follow-up to \\cite{Rasanen:2006b}, and more discussion of these topics can be found there. However, we have repeated some material to make the presentation self-contained. ",
        "conclusions": "\\subsection{Observational issues} \\label{sec:obs}  \\para{The age of the universe.}  Observationally, the age of the universe is not known with precision in a model-independent manner. There is an important constraint from the ages of globular clusters, which give the lower limit $t_0\\geq11.2$ Gyr at 95\\% C.L. and a best-fit age of $t_0=13.4$ Gyr \\cite{Krauss:2003}. These results do not depend on the details of late-time cosmology, and should be valid as long as the expansion at redshifts $z>6$ is close to the Einstein-de Sitter case. For the Hubble parameter, the most accurate model-independent determination is from SNe Ia observed with the Hubble Space Telescope \\cite{Jackson:2007}. The usually quoted value is $h=0.72\\pm0.08$ \\cite{Freedman:2000} (in close agreement with $h=0.73\\pm0.06$ found in \\cite{Riess:2005}), while recent work with a different treatment of Cepheid metallicity gives $h=0.62\\pm0.05$  \\cite{Sandage:2006} (all 1$\\sigma$ limits). These estimates give a mean value of $H_0t_0\\approx0.99$ or $H_0t_0\\approx0.85$, respectively, using $t_0=13.4$ Gyr. Using the lower limit for $t_0$ and the mean value for $h$ gives $H_0t_0\\gtrsim0.83$ or $H_0t_0\\gtrsim0.71$, respectively. The treatment of Cepheids does not appear to be settled \\cite{cepheid, Jackson:2007}, but in any case, $H_0 t_0$ values in the lower range are not ruled out. In principle, $H_0t_0$ can provide an important constraint on backreaction models. In particular, a definitive measurement of $Ht>1$ would imply that exotic matter is needed, because in a dust universe $Ht\\leq1$, regardless of the structures present (assuming that vorticity can be neglected) \\cite{Rasanen:2005}.  The values of $H t$ today in \\fig{fig:Ht} are rather low for the BBKS transfer function, while the BD transfer function gives results more in agreement with the observations. The difference is mainly due to the fact that the change in the fraction of mass which is in structures is slower in the BBKS case. For the BBKS transfer function, only 30\\% of the mass is in structures at 15 Gyr, compared with 54\\% for the BD transfer function. In both cases, initially (i.e. in the limit $t\\rightarrow0$) the fraction is 10\\%.  \\para{The CMB peaks.}  As the exact expression \\re{RfirstQ} shows and \\fig{fig:omegas} demonstrates, the average spatial curvature is generically large in a dust universe with significant clumpiness. This raises the question of compatibility of the backreaction explanation for the acceleration with CMB observations. However, the fact that the CMB data is in good agreement with FRW models that have no spatial curvature does not imply that the CMB rules out models with spatial curvature. Conclusions drawn about spatial curvature from the CMB are model- and prior-dependent.  In physical terms, the CMB peak structure is mostly determined by the primordial spectrum of perturbations, by $\\Om h^2$ and $\\Ob h^2$, and by the shift parameter, which is a measure of the distance to the last scattering surface \\cite{Efstathiou:1998, Mukhanov:2003}. The CMB anisotropies will be the mostly the same in models with identical values of these parameters, apart from the low multipoles (though see \\cite{Elgaroy:2007, Corasaniti:2007}).  Even in the \\LCDM model with primordial perturbations given by a power law, the CMB alone does not constrain spatial curvature to be small \\cite{Spergel:2006}, though combining the CMB with a measurement of Hubble parameter does provide a strong constraint. In a FRW model with a time-varying equation of state, at least $\\Omega_K\\approx0.2$ is compatible with the correct CMB shift parameter (and fitting the SN Ia data and the $A$-parameter from baryon acoustic oscillations) \\cite{Ichikawa}.  In a backreaction model, there is no simple argument for obtaining the position of the CMB peaks. Both positively and negatively curved regions are present, and there may be more cancellation in the effect on the passage of light than would be expected based on a FRW model with the equivalent amount of spatial curvature. One has to redo the CMB analysis, considering the passage of light in a universe with realistic structures, as discussed in section \\ref{sec:prop}. It might even be that the change in the passage of light together with the increase in $Ht$ could explain the observations without acceleration, as discussed in \\cite{Mattsson:2007b}.  In any case, the contribution of spatial curvature is regionally not negligible in the real universe. (The following discussion neglects effects due to breakdown of the dust approximation.) For example, consider an overdense dust structure. For a shell that is turning around from expansion to collapse, the left-hand side of \\re{Hamloc} is zero. Neglecting vorticity (in the spherically symmetric case, it would be zero), the spatial curvature must be positive, and equal to the sum of the contributions of the energy density and shear, $\\sR=16\\pi\\GN\\rho+2\\sigma^2$. For a stabilised structure, we have $\\theta=\\dot\\theta=0$, so solving from \\re{Rayloc} and \\re{Hamloc} we obtain $\\sR=12\\pi\\GN\\rho$ and $\\omega^2=2\\pi\\GN\\rho+\\sigma^2$. Just as vorticity is needed to balance against the energy density and shear to have zero acceleration, zero expansion requires the spatial hypersurfaces to be positively curved. The energy density of stabilised structures can be much larger than the average energy density, so the same is true of their spatial curvature. (The growth in the relative spatial curvature is balanced by the fact that the volume occupied by stabilised regions shrinks by the same factor that their energy density increases, so the contribution to the average Hubble rate \\re{Ham} remains constant.)  Inside voids, the energy density is much smaller than the background value, while the expansion rate is larger, so the contribution of the spatial curvature is again larger than that of energy density. For example, consider a void on an Einstein-de Sitter background. Neglecting shear and vorticity, we have $|\\sR_{\\mathrm{void}}|>16\\pi\\GN |\\delta|\\av{\\rho}$, where $\\delta$ is the real (not linear) density contrast of the void (typically $\\delta\\approx-0.9$ for observed voids \\cite{voidobs}).  In principle, observations of structures correlated with the CMB at different redshifts could lead to useful constraints on the evolution of the spatial curvature with redshift, as with the ISW effect \\cite{Ho:2008}, once the effect of non-linear structures on the passage of light is properly understood.  \\para{Variance.}  A variance $|\\OQ|\\gtrsim0.2$ is required to explain the observed acceleration \\cite{Rasanen:2006b}. The directional variation of the expansion rate was studied in \\cite{McClure:2007} using SNe Ia (following the treatment of \\cite{Freedman:2000} rather than \\cite{Sandage:2006}). Typical variation was found at the 10--20\\% level, the maximum difference in the expansion rates being over 50\\%. Of course, angular variation is different from volume variance, and the systematics of SNe Ia are perhaps not completely understood, as discussed in section \\ref{sec:intro}. For directional analysis of SNe Ia, see also \\cite{Haugbolle:2006, Schwarz:2007}.  In \\cite{Li:2007b}, the directional variation of the expansion rate inferred from SNe Ia in \\cite{Freedman:2000} was interpreted as evidence for backreaction. However, the estimate of the variance in \\cite{Li:2007b} is unreliable (even apart from the applicability of linear perturbation theory). The magnitude of the the relevant term $B=\\av{\\pat_i(\\pat_i\\varphi\\nabla^2\\varphi)}-\\av{\\pat_i(\\pat_j\\varphi\\pat_i\\pat_j\\varphi)}-\\frac{2}{3}\\av{\\nabla^2\\varphi}^2$ was estimated by replacing spatial derivatives with the inverse of the averaging scale, and replacing the two powers of $\\varphi$ by the primordial amplitude of the power spectrum. However, spatial gradients are determined by the scale over which the perturbations vary significantly (as determined by the power spectrum and the transfer function), not by the averaging scale. For example, $\\av{\\delta^2}\\propto\\av{\\nabla^2\\varphi\\nabla^2\\varphi}$ is divergent for a scale-invariant spectrum (with no free-streaming cut-off), regardless of how large the averaging scale is, because there are perturbations on arbitrarily small scales. Conversely, with a free-streaming cut-off, the quantity is finite no matter how small the averaging scale is.  Note that the variance in the expansion rate evaluated from Newtonian simulations \\cite{Rauch:2005} is large enough to provide acceleration, were it not balanced by shear, as noted in \\cite{Rasanen:2006b}. As we discuss in  section \\ref{sec:Newton} below, such a cancellation is not in present in general relativity. If the variance in simulations were negligible, it would be less plausible that it is large in the real relativistic case. However, the Newtonian variance is large, so in order to have acceleration, the shear simply has to be smaller than in the Newtonian case.  It is sometimes argued that the dust shear would be of the order $\\sigma^2\\lesssim 10^{-10} H^2$ from the isotropy of the CMB \\cite{EGS} (see also \\cite{moreEGS}). Any real non-linear dust structure violates this bound, so it is clear that the bound does not apply to the real universe. (The calculation also gives a limit of $\\abs{\\nabla\\rho/\\rho}\\lesssim 10^{-5} H$ for the spatial scale of density variation, which is even more obviously violated.) However, if the shear really was negligible, then this would indicate that backreaction is important, as observations and simulations suggest a significant variance of the expansion rate \\cite{Rasanen:2006b}.  \\para{Coldness of the Hubble flow.}  We have noted that there is no evidence for a matter component with negative pressure apart from the cosmological observations of accelerated expansion. In contrast to this, it has been claimed that the influence of vacuum energy can be seen in the local dynamics within the local 10 Mpc or so from the 'coldness' of the local Hubble flow.  The coldness refers to two distinct phenomena\\footnote{The linearity of the nearby mean flow with distance is sometimes mentioned as a third aspect.}. First, the velocity dispersion in the local volume has been claimed to be anomalously low. Some authors find the dispersion within the nearest few Mpc to be $\\approx$ 40 km/s \\cite{lowsigma}, while others find a value of 88 km/s $\\pm$ 20 km/s \\cite{Maccio:2004}, or over 100 km/s \\cite{Whiting:2002}. The second aspect is that the expansion rate measured within the nearest 10 Mpc is said to agree quite well with the global Hubble rate. Given that the matter distribution is locally quite clumpy, and does not become homogeneous until around 100 \\mpc, this seems somewhat surprising. It has been suggested that both observations are explained by vacuum energy. It can 'cool' the local expansion rate, and if vacuum energy dominates also the global dynamics, it is natural that the local and global expansion rates agree.  The velocity dispersion argument was presented as evidence against a high density matter-dominated FRW model before the SN Ia observations supported accelerating expansion \\cite{Suto}. It has been argued that the observed low value of the velocity dispersion could be due to the known peculiarity of the local structures instead \\cite{vdW}, or even that such a small dispersion is typical \\cite{Leong:2003}. However, it seems that the typical velocity dispersion in simulations of the EdS model is indeed considerably higher, 300--700 km/s \\cite{Governato:1996}, while \\LCDM simulations can reproduce the small velocity dispersion \\cite{Klypin:2001}. In simulations constrained to reproduce the large scale structure of the local universe, the velocity dispersion around regions similar to the Local Group of galaxies has been found to be as cold in the open CDM model as in the \\LCDM model, given the same matter density \\cite{Hoffman:2007b}. This still leaves the question of how common such regions are. However, the lower range of the values 150--300 km/s found in unconstrained simulations of the open CDM model \\cite{Governato:1996} is not very dissonant with an observational value of over 100 km/s. It does not seem unreasonable that the effect of spatial curvature would be even stronger in a clumpy model, where the density parameter of spatial curvature can be larger.  Unlike for the velocity dispersion, the probability of being located in a region where the local expansion rate agrees with the global value has not been studied quantitatively. At any rate, as long as the measurement of the global and local Hubble parameters is uncertain, it seems difficult to draw strong conclusions. The notably different values $h=0.62$ and $h=0.72$ have both been used as evidence for the vacuum energy origin of the local expansion rate \\cite{Sandage:2006, Chernin, localde}.  It has also been argued that the influence of vacuum energy has been directly measured in the motions of nearby galaxies \\cite{Chernin, localde}. The idea is that for a spherically symmetric system with a large central mass, the repulsive gravity of vacuum energy dominates beyond a certain radius, estimated as 1--2 Mpc for the Local Group. However, the assumption of a spherically symmetric field generated by a point mass used in the analysis does not hold very well, and the observations are also consistent with domination by negative curvature \\cite{Hoffman:2007b}.  \\subsection{Non-Newtonian effects} \\label{sec:Newton}  \\para{Spatial curvature and Newtonian gravity.}  It is not clear how much the overdense regions slow down the expansion rate, but from the model we have discussed, it seems difficult to avoid the conclusion that underdense regions cause $Ht$ to increase significantly. The physical interpretation seems straightforward: the fraction of space in faster expanding regions grows. However, there is a subtlety involved. The evolution we took for the individual regions is purely Newtonian, and the peak statistics do not involve general relativity. Had we formulated the problem in Newtonian gravity instead of general relativity, the answer would have been the same. However, in Newtonian gravity, the backreaction variable $\\sQ$ is a boundary term \\cite{Buchert:1995}, so backreaction vanishes for periodic boundary conditions. (In particular, this is true for Newtonian simulations.)  While the universe presumably does not have periodic boundary conditions at the visual horizon, the result implies that backreaction in Newtonian gravity also vanishes for a statistically homogeneous and isotropic system. This can be seen in two ways. First, consider a volume which has periodic boundary conditions on a scale much larger than the homogeneity scale. By statistical homogeneity and isotropy, the value of $\\sQ$ evaluated within each homogeneity scale sized box in the volume is the same as the overall value, which is zero due to the boundary conditions. Because the local physics is independent of the boundary conditions on very large scales, this result should hold even if the boundary conditions are not periodic\\footnote{I am grateful to Christof Wetterich for this argument.}. A perhaps more physical argument, explained in \\cite{Notari:2005}, is that a boundary term can be viewed as a flux going from one region to another, and in a statistically homogeneous and isotropic region, there should be an equal flux in and out.  In Newtonian gravity, the average evolution of a statistically homogeneous and isotropic dust space does always follow the FRW equations (i.e. the Buchert equations with $\\sQ=0$). However, this is not true in general relativity. The Newtonian theory constraint that variance and shear in $\\sQ$ in \\re{Q} cancel up to a boundary term is not present in general relativity \\cite{Buchert:1999}. Viewed equivalently in terms of the spatial curvature via the integrability condition \\re{int}, this difference is a reflection of the absence of spatial curvature in Newtonian gravity.  In Newtonian space and time, the geometry of the spatial hypersurfaces is fixed, so the spatial curvature tensor and its trace, the spatial curvature scalar $\\sR$, do not exist \\cite{Ellis:1971}. Therefore the Newtonian equations of motion do not include an equation for the expansion rate such as the Hamiltonian constraint \\re{Hamloc} independent of the Raychaudhuri equation. The analogue of the relativistic Hubble equation \\re{Hamloc} emerges only as the first integral of the Raychaudhuri equation.  In Newtonian gravity, the backreaction variable $\\sQ$ is given by a boundary term which is in general non-zero. However, for an isolated system or one that is statistically homogeneous and isotropic, we have $\\sQ=0$. Then the the first integral of the Newtonian equivalent of the average Raychaudhuri equation \\re{Ray} gives an equation like the average Hamiltonian constraint \\re{Ham}, but with a term proportional to $a^{-2}$ in place of $\\av{\\sR}$, as \\re{firstQ} shows. The $a^{-2}$ term multiplied by $a^2$ can be interpreted as the conserved energy of the Newtonian system.  There is no such conserved quantity in the relativistic case. The average spatial curvature can evolve non-trivially, unlike the total energy of an isolated Newtonian system. The average spatial curvature is constrained by the integrability condition \\re{int} for the average Hamiltonian constraint and the average Raychaudhuri equation, instead of being proportional to $a^{-2}$. (In the FRW case, the relativistic spatial curvature term and the Newtonian energy term are mathematically identical, only the physical interpretation is different.)  In our model, the evolution of overdense and underdense regions is uncorrelated. We do not have a Newtonian constraint on their behaviour, and the overall average spatial curvature can evolve non-trivially, as in the general relativistic case. A Newtonian calculation for a statistically homogeneous and isotropic system would have to include the global constraint that the underdense and overdense regions add up in such a manner that the energy term (corresponding to the spatial curvature) is proportional to $a^{-2}$. (In a consistent treatment with a continuous distribution of matter, instead of isolated regions, this would follow automatically.)  There is a similar situation in the case of spherical symmetry. As noted in section \\ref{sec:setup}, the backreaction variable $\\sQ$ vanishes for spherical symmetry in Newtonian gravity \\cite{Sicka, Buchert:2000}, resulting in the spherical collapse/expansion model. In general relativity, spherical symmetry does not imply perfect cancellation between variance and shear, and the average expansion of a spherically symmetric dust solution can even accelerate \\cite{Chuang:2005, Kai:2006, Paranjape:2006a}.  \\para{The Newtonian limit of general relativity.}  In order for backreaction to be able to account for the observed acceleration, non-Newtonian aspects of gravity should be important already at the homogeneity scale of around $100$ \\mpc. It is sometimes argued that general relativistic effects in the present-day universe can only be important on super-Hubble scales or near neutron stars and black holes. However, this conclusion is based on linearly perturbed FRW or Minkowski universes or the Schwarzschild solution, and the general situation is not that simple. (The subtleties of the Newtonian limit of near-FRW cosmological models have been discussed in \\cite{Ellis:1994, vanElst:1998}; see also \\cite{Kofman:1995} regarding the magnetic part of the Weyl tensor.)  The Einstein equation has ten components, with four constraints, whereas in Newtonian gravity there is only one equation, the Poisson equation. There are two aspects to the additional equations. Some of them correspond to degrees of freedom which do not exist in Newtonian gravity (such as gravity waves and spatial curvature), while others provide new constraints on the existing degrees of freedom. The low-velocity, weak-field limit of general relativity is not Newtonian gravity, but Newtonian gravity with extra degrees of freedom and additional constraints.  General relativity reduces to Newtonian gravity in the limit of infinite speed of light. However, this limit is not smooth. Taking the speed of light to infinity removes some equations completely and makes Newtonian gravity qualitatively different from general relativity, where the speed of light is finite. Relativistic effects do not require velocities near the speed of light, only a finite speed of light.  For example, there exist solutions of Newtonian gravity which are not the limit of any general relativity solution, due to additional constraints in the relativistic case. There are Newtonian expanding dust solutions which have zero shear but non-zero vorticity. However, in general relativity, non-zero vorticity implies non-zero shear for expanding dust \\cite{Ellis:1967, Ellis:1971}. This does not require large velocities or strong gravitational fields; see \\cite{Senovilla:1997} for a clear comparison of the general relativistic and Newtonian cases.  In the case of backreaction, it is the presence of new degrees of freedom, namely spatial curvature, rather than extra constraints which is important. In perturbation theory around FRW universes, the leading terms in the backreaction variable $\\sQ$ are Newtonian, so they are total derivatives and do not contribute in a statistically homogeneous and isotropic system \\cite{Rasanen, Kolb:2004a, Notari:2005}. The non-Newtonian terms do not appear earlier than fourth order in perturbation theory. Their numerical coefficients have not been evaluated (it may be possible to do so using third order perturbation theory, instead of having to calculate to fourth order \\cite{Li:2007a}), but their form is known, and velocities near the speed of light are not required for them to be important.  \\para{Simulations and backreaction.}  If non-Newtonian aspects are important already at the homogeneity scale, one would expect them to show up also in quantities other than the expansion rate. We can ask whether there is any room for such effects, given the comparison of simulations of structure formation (which are completely Newtonian) with observations. In fact, while simulations reproduce many features of observations, there are some notable differences.  In \\cite{SylosLabini:2006} it was found that the homogeneity scale evaluated from simulations was only of the order 10 \\mpc, an order of magnitude smaller than the observed 70--100 \\mpc \\cite{Hogg:2004, Pietronero}. (The box size of the simulations studied in \\cite{SylosLabini:2006} was 141 \\mpc, so the correct homogeneity scale could not have been reliably recovered. However, the homogeneity scale that was found is much smaller than the box size, so it is not likely to be a finite size artifact.)  According to \\cite{Einasto}, the number of largest structures in the Millennium Simulation is underproduced compared to observations, by a factor of 10 for the most luminous superclusters. As for large underdense structures, the voids produced in simulations are not as empty as those observed. This `void phenomenon' has even been called a crisis of the \\LCDM model \\cite{Peebles}. (It has recently been argued that simulations coupled with semianalytical galaxy formation in fact agree with the void observations \\cite{vonBendaBeckmann:2007}.)  These discrepancies may have resolutions which have nothing to do with non-Newtonian physics or backreaction. The statistics for the largest structures are probing the tail of the distribution, which may be affected by non-linear corrections to the evolution of the power spectrum \\cite{Crocce:2007} or transients from initial conditions \\cite{Crocce:2006}. The magnitude of such effects has been found at the 10--30\\% level, and the discrepancy is an order of magnitude, but one should not be confident that there are no unaccounted for effects in simulations, for example related to discretisation \\cite{discrete}. The apparent emptiness of voids, in turn, might be due to bias or subtleties in galaxy formation and the definition of voids, rather than spatial curvature \\cite{Ostriker:2003, Furlanetto:2005}.  Nevertheless, the point is that we do not have confirmation that Newtonian gravity works well near the homogeneity scale, since the Newtonian results differ in some important respects from the observations.  \\subsection{Conclusion}  \\para{Summary.}  We have studied the effect of structure formation on the expansion rate in a statistically homogeneous and isotropic space. We first reviewed studies of the propagation of light in a space with non-linear structures, and discussed some qualitative issues related to the average expansion rate. We then calculated the average expansion rate in a model where the number density of structures is given by the peak model of structure formation with cold dark matter, and the individual structures are described with the spherical collapse model and its underdense equivalent. In the calculation, there are no adjustable free parameters in addition to those related to structure formation in a spatially flat matter-dominated FRW universe.  We find that the expansion rate increases relative to the FRW value at a time of tens of billions years, about the observed acceleration era, possibly offering a solution to the coincidence problem. The timescale has its origin in the change of slope of the CDM transfer function around the matter-radiation equality scale $\\keq$. This leads to a preferred time around $10^5\\teq\\approx 10^{10}$ years, when the volume occupied by structures and their size relative to the visual horizon saturate, and the impact of structures on the expansion rate becomes large.  However, while $H t$ increases in the model, it does not rise sufficiently rapidly to correspond to acceleration. The expansion rate increases because the relative volume of the faster expanding regions rises, as quantified by the Buchert equations. In order to have acceleration, it is likely that the average expansion rate would first have to be damped by slower expanding overdense regions before it is increased by shedding their contribution when voids become dominant. In the present model, this does not happen, because the volume occupied by the overdense regions is rather small. This is partly due to the fact that we have equal amounts of mass in overdense and underdense regions. Taking into account mass flow to the overdense regions would increase their impact.  Improving the treatment of the statistics of the peaks and the model used for individual structures would lead to a more realistic estimate of the expansion rate. The propagation of light in a universe with realistic, evolving structures also has to be further studied to see how the average expansion rate is related to observations.  If a realistic backreaction model does turn out to provide acceleration in agreement with the observations, one could say that it unifies three historically popular FRW models. There is only matter present and the initial value of the spatial curvature at the background level is zero as in the standard CDM model, the universe has negative curvature as in the open CDM model, and the expansion accelerates as in the \\LCDM model. Acceleration due to structure formation would probably be a transient phase on the way to negative curvature voids dominating the expansion of the universe, though the behaviour at very late times depends on how the universe is structured on scales which are currently far beyond the horizon, of which we have no theoretical or observational understanding.  \\ack  I thank Thomas Buchert and Ruth Durrer for discussions and comments on the manuscript, Henk van Elst, Troels Haugb{\\o}lle, Julien Lesgourgues, David Mota, Misao Sasaki, Douglas Shaw, Martin Sloth and Christof Wetterich for helpful discussions and correspondence, and the Department of Physics and Astronomy at the University of \\AA{}rhus for hospitality. This work was partly done at the CERN Theory Unit.\\\\  \\setcounter{secnumdepth}{0}"
    },
    "0801/0801.2976_arXiv.txt": {
        "abstract": "{ We report on the high-redshift blazar identification of a new gamma--ray source, Swift J1656.3$-$3302, detected with the BAT imager onboard the {\\it Swift} satellite and the IBIS instrument on the {\\it INTEGRAL} satellite. Follow-up optical spectroscopy has allowed us to identify the counterpart as an $R\\sim$ 19 mag source that shows broad Lyman-$\\alpha$, Si {\\sc iv}, He {\\sc ii}, C {\\sc iv}, and C {\\sc iii}] emission lines at redshift $z$ = 2.40$\\pm$0.01. Spectral evolution is observed in X--rays when the {\\it INTEGRAL}/IBIS data are compared to the {\\it Swift}/BAT results, with the spectrum steepening when the source gets fainter. The 0.7--200 keV X--ray continuum, observed with {\\it Swift}/XRT and {\\it INTEGRAL}/IBIS, shows the power law shape typical of radio loud (broad emission line) active galactic nuclei (with a photon index $\\Gamma \\sim$ 1.6) and a hint of spectral curvature below $\\sim$2 keV, possibly due to intrinsic absorption (N$_{\\rm H} \\sim$ 7$\\times$10$^{22}$ cm$^{-2}$) local to the source. Alternatively, a slope change ($\\Delta$$\\Gamma$ $\\sim$ 1) around 2.7 keV can describe the X--ray spectrum equally well. At this redshift, the observed 20--100 keV luminosity of the source is $\\sim$10$^{48}$ erg s$^{-1}$ (assuming isotropic emission), making Swift J1656.3$-$3302 one of the most X--ray luminous blazars. This source is yet another example of a distant gamma--ray loud quasar discovered above 20 keV. It is also the farthest object, among the previously unidentified {\\it INTEGRAL} sources, whose nature has been determined {\\it a posteriori} through optical spectroscopy. ",
        "introduction": "Blazars are distant and powerful active galactic nuclei (AGNs) which are oriented in such a way that a jet expelled from the central black hole is directed at small angles with respect to the observer's line of sight (for a recent review, see Padovani 2007). In the widely adopted scenario of blazars, a single population of high-energy electrons in a relativistic jet radiates over the entire electromagnetic spectrum via synchrotron and inverse Compton processes, the former dominating at low energies, the latter being relevant at high energies (Ghisellini et al. 1998). The ambient photons that are inverse Compton scattered can be either internal (synchrotron self-Compton) and/or external (external Compton scattering) to the jet. As a consequence, the spectral energy distribution (SED) of blazars shows a double-humped shape, with the synchrotron component peaking anywhere from infrared to X--rays and the inverse Compton emission extending up to GeV/TeV gamma rays.  To explain the various SED shapes observed in blazars, Fossati et al. (1998) proposed the so-called ``blazar sequence\", according to which a relation between peak energies and $\\gamma$-dominance (the luminosity ratio of the second to the first peak) is present as a function of the source total power. This means that more luminous sources have both synchrotron and inverse Compton peaks located at lower energies and are more gamma--ray dominated than their fainter (and generally lower redshift) analogues.  Within the blazar population, high-redshift objects are the most luminous and generally belong to the class of Flat-Spectrum Radio Quasars (FSRQ). Observations of high-luminosity blazars in the X--/gamma--ray band are particularly important (especially if available over a broad energy range) as they allow the characterization of the inverse Compton peak and related parameters. More specifically, a flattening in the spectral distribution of the seed photons producing X--rays via inverse Compton is often observed at low energies in the X--ray spectra of these objects and can be measured only with broad band data (see e.g. Tavecchio et al. 2007 and references therein).  Unfortunately, the situation is far more complex, as absorption intrinsic to the source can also reproduce the spectral curvature observed in the X--ray band (e.g., Page et al. 2005; Yuan et al. 2006); in this case, information on the absorption is useful to understand the source environment and its relation to the jet. Besides this, X--/gamma--ray observations can provide evidence for the existence of extreme blazars, i.e. those with the synchrotron peak lying at X--ray energies (Giommi et al. 2007; Bassani et al. 2007).  Here, we report detailed information on a new, powerful and hard X--ray selected blazar, Swift J1656.3$-$3302, recently discovered through high-energy observations made with {\\it Swift}/BAT and {\\it INTEGRAL}/IBIS. We present the results of our optical follow-up work, which has allowed the identification of the source with a blazar at redshift $z$ = 2.4, along with an accurate analysis of the available {\\it Swift}/XRT and {\\it INTEGRAL}/IBIS data. We also construct a SED for Swift J1656.3$-$3302 and discuss the characteristics of the source broad band emission.  The paper is structured as follows: Sect. 2 reports a collection of the main results available in the literature on this source; Sects. 3 and 4 illustrate the optical and high-energy observations, respectively; Sect. 5 contains the results of this observational campaign, while a discussion on them is given in Sect. 6. Conclusions are outlined in Sect. 7. Throughout the paper, and unless otherwise specified, uncertainties are given at the 90\\% confidence level. We also assume a cosmology with $H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_\\Lambda$ = 0.7 and $\\Omega_{\\rm m}$ = 0.3. ",
        "conclusions": "Through optical follow-up observations we have been able to identify the newly discovered hard X--ray source Swift J1656.3$-$3302 with a blazar at $z$ = 2.4. Spectral evolution is observed when comparing the Swift/BAT and {\\it INTEGRAL}/IBIS data, with the X--ray spectrum undergoing softening when the emission becomes fainter. The source broadband X--/gamma--ray spectrum is well described by a power law of index $\\Gamma \\sim$ 1.6. The source is extremely bright with a 20--100 keV observer's frame luminosity of $\\sim$10$^{48}$ erg s$^{-1}$, assuming isotropic emission.  The source SED is typical of high luminosity blazars, as it has two peaks: the lower (synchrotron) one likely located at infrared frequencies, and the higher (inverse Compton) one positioned above a few hundred keV. The spectral curvature detected in the X--ray spectrum of the object can either be intrinsic and due to the distribution of the emitting electrons, or associated with the presence of absorption local to the source and produced by a column density of a $\\sim$7$\\times$10$^{22}$ cm$^{-2}$, higher than that typically observed in high-$z$ blazars.  All of the observational evidence gathered so far therefore points to Swift J1656.3$-$3302 being another case of a distant luminous blazar selected at gamma--ray energies. It is also the farthest object, among the previously unidentified {\\it INTEGRAL} sources, whose nature has been determined {\\it a posteriori} through optical spectroscopy."
    },
    "0801/0801.2147_arXiv.txt": {
        "abstract": "Emission-line abundances have been uncertain for more than a decade due to unexplained discrepancies in the relative intensities of the forbidden lines and weak permitted recombination lines in planetary nebulae (PNe) and H II regions.  The observed intensities of forbidden and recombination lines originating from the same parent ion differ from their theoretical values by factors of more than an order of magnitude in some of these nebulae.  In this study we observe UV resonance line absorption in the central stars of PNe produced by the nebular gas, and from the same ions that emit optical forbidden lines. We then compare the derived absorption column densities with the emission measures determined from ground-based observations of the nebular forbidden lines.  We find for our sample of PNe that the collisionally excited forbidden lines yield column densities that are in basic agreement with the column densities derived for the same ions from the UV absorption lines.  A similar comparison involving recombination line column densities produces poorer agreement, although near the limits of the formal uncertainties of the analyses. An additional sample of objects with larger abundance discrepancy factors will need to be studied before a stronger statement can be made that recombination line abundances are not correct. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.1622_arXiv.txt": {
        "abstract": "The discovery of refractory grains amongst the particles collected from Comet 81P/Wild 2 by the Stardust spacecraft (Brownlee et al. 2006) provides the ground truth for large-scale transport of materials formed in high temperature regions close to the protosun outward to the comet-forming regions of the solar nebula. While accretion disk models driven by a generic turbulent viscosity have been invoked as a means to explain such large-scale transport, the detailed physics behind such an ``alpha'' viscosity remains unclear. We present here an alternative physical mechanism for large-scale transport in the solar nebula: gravitational torques associated with the transient spiral arms in a marginally gravitationally unstable disk, of the type that appears to be necessary to form gas giant planets. Three dimensional models are presented of the time evolution of self-gravitating disks, including radiative transfer and detailed equations of state, showing that small dust grains will be transported upstream and downstream (with respect to the mean inward flow of gas and dust being accreted by the central protostar) inside the disk on time scales of less than 1000 yr inside 10 AU. These models furthermore show that any initial spatial heterogeneities present (e.g., in short-lived isotopes such as $^{26}$Al) will be homogenized by disk mixing down to a level of $\\sim$ 10\\%, preserving the use of short-lived isotopes as accurate nebular chronometers, while simultaneously allowing for the spread of stable oxygen isotope ratios. This finite level of nebular spatial heterogeneity appears to be related to the coarse mixing achieved by spiral arms, with radial widths of order 1 AU, over time scales of $\\sim$ 1000 yrs. ",
        "introduction": "Comets have long been recognized as relatively pristine repositories of particles from the early solar system, and combined with primitive chondritic meteorites, these objects represent the best records we have of the events in the solar nebula some 4.57 Gyr ago. For this very reason, the Stardust Mission to Comet 81P/Wild 2 was conceived and launched to sample the particles in the coma of this Jupiter family comet (JFC). JFCs are believed to have formed in the Kuiper Belt, just beyond Neptune's orbit, and so Wild 2 is thought to be a sample of the outermost planet-forming region of the solar nebula. The discovery of refractory grains by Stardust (Brownlee et al. 2006) therefore provided proof that grains that formed in high temperature regions of the disk, presumably close to the protosun, were able to be transported outward tens of AUs to the JFC-forming region prior to their incorporation in comets (see also Nuth \\& Johnson 2006).  There are two possibilities for the outward transport of the high temperature minerals found by Stardust. First, refractory grains formed close to the protosun may have been carried upward by the protosun's bipolar outflow and lofted onto trajectories that would return them to the surface of the solar nebula at much greater distances (Shu et al. 2001). Bipolar outflows are known to occur, but the lofting of grains and their return to the outer disk are processes that are not constrained by astronomical observations, though the Stardust discoveries can be interpreted as proof of this possibility. Alternatively, refractory solids might have been transported outward to comet-forming distances by mixing processes within the solar nebula, such as generic accretion disk turbulence (Gail 2002; Ciesla \\& Cuzzi 2006; Tscharnuter \\& Gail 2007; Ciesla 2007), or the actions of spiral arms in a marginally gravitationally unstable nebula (Boss 2004, 2007). The magnetorotational instability (MRI) in an ionized disk is often considered to be the physical mechanism responsible for disk evolution and the source of the disk's effective ``alpha'' viscosity. However, MRI effects are limited to the disk's surfaces and its innermost and outermost regions by the need for appreciable fractional ionization. The ionized surface layers are thought to be dominated by the inward flow of gas accreting onto the protostar. The disk's midplane is essentially neutral in the planet-forming region from $\\sim 1$ to $\\sim 15$ AU, preventing the MRI from serving as a driver of midplane disk evolution in this key region (e.g., Matsumura \\& Pudritz 2006). Gravitational torques in a marginally gravitationally unstable disk, however, are able to drive global angular momentum transport throughout the dead zone, and in fact periods of gravitational instability may well be regulated by the pile-up of disk gas on the edge of the dead zone by the action of the MRI in the more ionized regions of the disk.  Astronomical observations of disks around young stars often find evidence for crystalline silicate grains at distances ranging from inside 3 AU to beyond 5 AU, in both the disk's midplane and its surface layers (Mer\\'in et al. 2007). These grains could have been produced through thermal annealing of amorphous grains by hot disk temperatures reached only within the innermost disk, well inside 1 AU. Crystalline and amorphous silicate grains thus appear to require large-scale transport as well. In addition, isotopic evidence suggests (e.g., Bizzarro, Baker \\& Haack 2004) that some Allende chondrules formed with an $^{26}$Al/$^{27}$Al ratio similar to that of calcium, aluminum-rich inclusions (CAIs). Chondrule formation thus appears to have begun shortly after CAI formation, and to have lasted for $\\sim$ 1 Myr to 4 Myr (Amelin et al. 2002). While chondrules are generally believed to have been melted by shock-front heating at asteroidal distances (e.g., Desch \\& Connolly 2002), CAIs are thought to have formed much closer to the protosun, again because of their refractory compositions. Transport of solids from the inner solar nebula out to asteroidal distances thus seems to be required in order to assemble chondrules, CAIs, and matrix grains into the chondritic meteorites (Boss \\& Durisen 2005).  While the evidence for large-scale transport is clear, the degree of mixing of the nebula gas and solid has been less clear for some time. Isotopic evidence has been presented for both homogeneity (e.g., Hsu, Huss \\& Wasserburg 2003) and heterogeneity (e.g., Simon et al. 1998) of short-lived radioisotopes (SLRI) such as $^{26}$Al and $^{60}$Fe in the solar nebula. The solar system's $^{60}$Fe must have been synthesized in a supernova (Mostefaoui, Lugmair \\& Hoppe 2005; Tachibana et al. 2006) and then injected either into the presolar cloud (Vanhala \\& Boss 2002) or onto the surface of the solar nebula (e.g., Boss 2007), processes that both lead to initially strongly heterogeneous distributions. The initial degree of homogeneity of the SLRI is crucial if $^{26}$Al/$^{27}$Al ratios are to be used as precise chronometers for the early solar system (Bizzarro et al. 2004; Halliday 2004; Krot et al. 2005; Thrane et al. 2006).  The degree of uniformity is equally murky even for stable isotopes. For samarium and neodymium isotopes, both nebular homogeneity and heterogeneity have been claimed, depending on the objects under consideration (Andreasen \\& Sharma 2007), e.g., ordinary versus carbonaceous chondrites. However, osmium isotopes seem to require a high degree of homogeneity in both ordinary and carbonaceous chondrites (Yokoyama et al. 2007). The three stable isotopes of oxygen, on the other hand, show clear evidence for heterogeneity in primitive meteorites (Clayton 1993). This heterogeneity is perhaps best explained by self-shielding of molecular CO gas from UV photodissociation at the surface of the solar nebula (Clayton 2002; Lyons \\& Young 2005; however, see Ali \\& Nuth 2007 for an opposing view) or in the presolar molecular cloud (Yurimoto \\& Kuramoto 2004). Oxygen isotopic anomalies formed at the outer surfaces of the solar nebula by definition imply a spatially heterogeneous initial distribution, perhaps capable of explaining the entire range of stable oxygen isotope ratios through mixing between $^{16}$O-rich and $^{16}$O-poor reservoirs (e.g., Sakamoto et al., 2007).  We present here a new set of three dimensional hydrodynamical models of mixing and transport in a protoplanetary disk that demonstrate how the required large-scale transport and mixing may have occurred in the solar nebula. These models are similar to those of Boss (2004, 2007) except that in the new models, the disk is assumed to extend from 1 AU to 10 AU, instead of from 4 AU to 20 AU, in order to better represent the regions of the nebula of most interest for high temperature thermal processing and meteorite formation. ",
        "conclusions": "The discovery of refractory minerals in Comet Wild 2 by the Stardust Mission (Brownlee et al. 2006) is consistent with the rapid (less than 1000 yr), large-scale (10s of AU) outward transport of small dust grains in a marginally gravitationally unstable disk. Such large-scale transport also seems to be required to explain observations of thermally processed, crystalline silicates in comets and protoplanetary disks (e.g., Nuth \\& Johnson 2006; van Boekel et al. 2005; Mer\\'in et al. 2007). Marginally gravitationally unstable disks have the additional feature of providing a robust source of shock fronts capable of thermally processing solids into chondrules (Desch \\& Connolly 2002; Boss \\& Durisen 2005). Furthermore, the degree of spatial heterogeneity in the solar nebula required for the use of the $^{26}$Al/$^{27}$Al chronometer and for the survival of the oxygen isotope anomalies is consistent with the results of these models of mixing and transport of initially highly spatially heterogeneous tracers in a marginally gravitationally unstable disk (Boss 2007). Such a disk appears to be required for the formation of Jupiter by either the core accretion (Inaba et al. 2003) or disk instability (Boss et al. 2002) mechanisms. Marginally gravitationally unstable disk models provide a physical mechanism for self-consistent calculations of the mixing and transport of dust grains across the planet-forming region of the solar nebula.  Mass accretion rates of the disk onto the central protostar are highly variable in marginally gravitationally unstable models, ranging from rates of $\\sim 10^{-6} M_\\odot$/yr to $\\sim 10^{-5} M_\\odot$/yr, much higher than the typical stellar mass accretion rate of $\\sim 10^{-8} M_\\odot$/yr for a classical T Tauri star, but comparable to the rates inferred for T Tauri stars during their periodic FU Orionis outbursts (Boss 1996). Recent observations suggest that most T Tauri disks have lifetimes of no more $\\sim$ 1 Myr (Cieza et al. 2007). Combined with a disk mass of $\\sim 0.1 M_\\odot$, this suggests a time-averaged mass accretion rate for solar-mass T Tauri stars of $\\sim 10^{-7} M_\\odot$/yr. A disk would then need to experience phases of gravitational instability lasting roughly 1\\% to 10\\% of its lifetime in order to accomplish the mass transport implied by these estimates. Phases of gravitational instability are thus expected to be transient phases, but robust and frequent enough to achieve the disk mixing and transport processing described in this paper.  I thank Sasha Krot and a second referee for their improvements to the paper, and Sandy Keiser for her computer wizardry. This research was supported in part by the NASA Planetary Geology and Geophysics Program under grant NNG05GH30G and by the NASA Origins of Solar Systems Program under grant NNG05GI10G, and is contributed in part to the NASA Astrobiology Institute under grant NCC2-1056. Calculations were performed on the Carnegie Alpha Cluster, which was supported in part by NSF MRI grant AST-9976645."
    },
    "0801/0801.1849_arXiv.txt": {
        "abstract": "We investigate the effect of metallicity calibrations,  AGN classification, and aperture covering fraction on the local mass-metallicity relation using 27,730 star-forming galaxies from the Sloan Digital Sky Survey (SDSS) Data Release 4.    We analyse the SDSS mass-metallicity relation with 10 metallicity calibrations, including theoretical and empirical methods.   We show that the choice of metallicity calibration has a significant effect on the shape and y-intercept(\\OH) of the mass-metallicity relation.  The absolute metallicity scale (y-intercept) varies up to $\\Delta[\\log({\\rm O/H})]= 0.7$~dex, depending on the calibration used, and the change in shape is substantial.  These results indicate that it is critical to use the same metallicity calibration when comparing different luminosity-metallicity or mass-metallicity relations.    We present new metallicity conversions that allow metallicities that have been derived using different strong-line calibrations to be converted to the same base calibration.  These conversions facilitate comparisons between different samples, particularly comparisons between galaxies at different redshifts for which different suites of emission-lines are available.   Our new conversions successfully remove the large $0.7$~dex discrepancies between the metallicity calibrations, and we reach agreement in the mass-metallicity relation to within $0.03$~dex on average.      We investigate the effect of AGN classification and aperture covering fraction on the mass-metallicity relation.  We find that different AGN classification methods have negligible effect on the SDSS MZ-relation.  We compare the SDSS mass-metallicity relation with nuclear and global relations from the Nearby Field Galaxy Survey (NFGS).   The turn over of the mass-metallicity relation at ${\\rm M}_{*}\\sim 10^{10}$~\\Msun\\ depends on aperture covering fraction.  We find that a lower redshift limit of $z<0.04$ is insufficient for avoiding aperture effects in fiber spectra of the highest stellar mass (${\\rm M}_{*} > 10^{10}$~\\Msun) galaxies. ",
        "introduction": "The relationship between metallicity and stellar mass provides crucial insight into galaxy formation and evolution.   Theory predicts that as time progresses, the mean stellar metallicity of galaxies increases with age as galaxies undergo chemical enrichment, while  the stellar mass of a galaxy will increase with time as galaxies are built through merging processes \\citep[e.g.,][and references therein]{Pei95,Somerville99,Somerville00,Nagamine01,Calura04}.    A correlation between mass and metallicity arises if low mass galaxies have larger gas fractions than higher mass galaxies, as is observed in local galaxies \\citep{McGaugh97,Bell00,Boselli01}.  The detailed relationship between metallicity and mass may depend critically on galactic-scale outflows driven by supernovae and stellar winds  \\citep[see e.g.,][for a review]{Garnett02,Pettini02}.  Thus, robust measurements of the mass-metallicity (MZ) relation may provide important clues into the impact of galactic-scale winds on the chemical history of galaxies.  The MZ relation was first observed in irregular and blue compact galaxies \\citep{Lequeux79,Kinman81}.  In subsequent work, luminosity was often used as a surrogate for mass because obtaining reliable mass estimates for galaxies was non-trivial.   \\citet{Rubin84} provided the first evidence that metallicity is correlated with luminosity in disk galaxies.  Further investigations solidified the correlation between luminosity and metallicity in nearby disk galaxies \\citep{Bothun84,Wyse85, Skillman89,Vila92, Zaritsky94,Garnett02}.    However, optical luminosity may not be a reliable surrogate for the stellar mass of a galaxy because optical luminosities are sensitive to the level of current star formation and are extinguished by dust.  Near infrared luminosities can be influenced by the age of the stellar population of a galaxy.   Fortunately, reliable stellar mass estimates are now possible, thanks to new state-of-the-art stellar evolutionary synthesis models \\citep[e.g.,][]{Silva98,Leitherer99,Fioc99,Bruzual03}.  Key insight into the mass-metallicity relation has recently been obtained with large spectroscopic surveys such as the Sloan Digital Sky Survey (SDSS) and the  2 degree Field Galaxy Redshift Survey (2dFGRS) \\citep[e.g.,][]{Baldry02,Schulte03,Lamareille04,Tremonti04,Gallazzi05}.  Using the SDSS stellar masses, \\citet[][; hereafter T04]{Tremonti04} characterized the local MZ relation for $\\sim 53,000$ local galaxies.  The MZ relation is steep for masses  $\\lesssim 10^{10.5}$~\\Msun\\ and flattens at higher masses.  T04 use chemical evolution models to interpret this flattening in terms of efficient galactic scale winds that remove metals from low mass galaxies ($M \\lesssim 10^{10.5}$~\\Msun).   Hierarchical galaxy formation models that include chemical evolution and feedback processes can reproduce the observed MZ relation \\citep{DeLucia04,DeRossi06,Finlator07}.  However, these models rely on free parameters, such as feedback efficiency, that are relatively unconstrained by observations.   Alternative scenarios proposed to explain the MZ relation include low star formation efficiencies in low-mass galaxies caused by supernova feedback \\citep{Brooks07}, and a variable integrated stellar initial mass function \\citep{Koppen07}.  The advent of large 8-10~m telescopes and efficient multi-object spectrographs enables the luminosity-metallicity (LZ) and, in some cases, the mass-metallicity relation to be characterized to high redshifts \\citep{Kobulnicky99a,Carollo01,Pettini01,Lilly03,Kobulnicky03,Kobulnicky04,Shapley04,Maier04,Liang04,Maier05,Hoyos05,Savaglio05,Mouhcine06,Erb06,Maier06,Liang06}.   Evolution in the LZ and MZ relations are now predicted by semi-analytic models of galaxy formation within the $\\Lambda$ cold dark matter framework that include chemical hydrodynamic simulations   \\citep{DeLucia04,Tissera05,DeRossi06,Dave07}.   Therefore reliable observational estimates of the LZ and MZ relations may provide important constraints on galaxy evolution theory.  Reliable MZ relations require a robust metallicity calibration.  Common metallicity calibrations are based on metallicity-sensitive optical emission-line ratios.  These calibrations include theoretical methods based on photoionization models \\citep[see e.g.,][for a review]{Kewley02a}, empirical methods based on measurements of the electron-temperature of the gas,  \\citep[e.g.,][]{Pilyugin01,Pettini04}, or a combination of the two \\citep[e.g.,][]{Denicolo02}.  Comparisons among the metallicities estimated using these methods reveal large discrepancies \\citep[e.g.,][]{Pilyugin01,Bresolin04, Garnett04}.  These discrepancies usually manifest as a systematic offset in metallicity estimates, with high values estimated by theoretical calibrations and lower metallicities estimated by electron-temperature metallicities.  Such offsets are found to be as large as 0.6~dex in log(O/H) units \\citep{Liang06,Yin07} and may significantly affect the shape and zero-point of the mass-metallicity or luminosity-metallicity relations.  Initial investigations into the extent of these discrepancies have recently been made by \\citet{Liang06,Yin07,Yin07b}, and \\citet{Nagao06}.    \\citet{Liang06} applied the metallicity calibrations from four authors to $\\sim40,000$ galaxies from the SDSS.  They showed that calibrations based on electron temperature metallicities produce discrepant metallicities when compared with calibrations based on photoionization models.   \\citet{Yin07} compared the theoretical metallicities derived by T04 with \\Te-based metallicities from \\citet{Pilyugin01} and \\citet{Pilyugin05}. They found a discrepancy of $\\Delta[\\log({\\rm O/H})]=0.2$~dex between the two \\Te-based metallicities, and a larger discrepancy of $\\Delta[\\log({\\rm O/H})]=0.6$~dex between the \\Te-based methods and the theoretical method.     Similar results were obtained by  \\citet{Yin07b} who extended the \\citet{Liang06} work to low metallicities typical of low mass galaxies.  The cause of the metallicity calibration discrepancies remains unclear.   The discrepancy has been attributed to either an unknown problem with the photoionization models \\citep{Kennicutt03}, or to temperature gradients or fluctuations that may cause metallicities based on the electron temperature method to underestimate the true metallicities \\citep{Stasinska02,Stasinska05,Bresolin06}.  Until this discrepancy is resolved, the absolute metallicity scale is uncertain.  The metallicity discrepancy issue highlights the need for an in-depth study into the effect of metallicity calibration discrepancies and other effects on the MZ relation.   In this paper, we investigate the robustness of the local MZ relation for star-forming galaxies.  We focus on three important factors that may influence the shape, y-intercept, and scatter of the mass-metallicity relation:  choice of metallicity calibration, aperture covering fraction, and AGN removal method.   Our sample selection is described in Section~\\ref{Sample}.  We describe the stellar mass estimates and metallicities in Sections~\\ref{stellar_masses} and ~\\ref{Metallicity}, respectively. We compare the mass-metallicity relation derived using 10 different, popular metallicity calibrations in Section~\\ref{calibrations}.    We find larger discrepancies between the MZ relations derived with different metallicity calibrations than previously found.  We calibrate the discrepancies between the different calibrations using robust fits, and we provide conversion relations for removing these discrepancies in Section~\\ref{Conversions}. We show that our new conversions  successfully remove the large metallicity discrepancies in the MZ relation.  We investigate the effect of different schemes for AGN removal in Section~\\ref{AGN}, and we determine the effect of fiber covering fraction in Section~\\ref{aperture}.  We discuss the impact of our results on the mass-metallicity and luminosity-metallicity relations in Section~\\ref{Discussion}.  Our conclusions are given in Section~\\ref{Conclusions}.   In the Appendix, we provide detailed descriptions of the metallicity calibrations used in this study and worked examples of the application of our conversions.   Throughout this paper, we adopt the flat $\\Lambda$-dominated cosmology as measured by the WMAP experiment \\citep[$h=0.72$, $\\Omega_{m}=0.29$;][]{Spergel03}). ",
        "conclusions": "\\label{Conclusions}  We present a detailed investigation into the mass-metallicity relation for 27,730 star-forming galaxies from the Sloan Digital Sky Survey.  We apply 10 different metallicity calibrations to our SDSS sample, including theoretical photoionization calibrations and empirical \\Te\\ method calibrations.   We investigate the effect of these  metallicity calibrations on the shape and y-intercept of the mass-metallicity relation.   Using 30,000 galaxy sets sampled randomly from our SDSS sample, we investigate how well the 9 strong-line calibrations maintain relative metallicities.   We find that:  \\begin{itemize} \\item  The choice of metallicity calibrations has the strongest effect on the MZ relation.  The choice of metallicity calibration can change the y-intercept of the MZ relation by up to 0.7~dex.   Until this metallicity discrepancy problem is resolved, absolute metallicities should be used with extreme caution. \\item There is considerable variation in shape and y-intercept of the MZ relations derived using the empirical methods.  We attribute this variation to the different \\HII\\ region samples used to derive the empirical calibrations. \\item The relative difference in metallicities is maintained to an accuracy of \\OH$\\sim 0.02 - 0.1$~dex for 9/10 calibrations, and to within \\OH$\\sim 0.15$~dex for all 9 strong-line calibrations.   For relative metallicity studies where the difference between targets or between samples is $\\lesssim 0.15$~dex, we recommend the use of at least two different calibrations to check that any result is not caused by the metallicity calibrations. \\end{itemize}  We use robust fits to the observed relationship between each metallicity calibration to derive new conversion relations for converting metallicities calculated using one calibration into metallicities of another, \"base\" calibration.   We show that these conversion relations successfully remove the strong discrepancies observed in the MZ relation between the different calibrations.   Agreement is reached to within 0.03~dex on average.  We investigate the effect of AGN classification scheme and fixed-size aperture on the MZ relation.  \\begin{itemize} \\item AGN classification methods have a negligible effect on metallicities derived using \\NIIOII\\ or \\NIIHa.  AGN classification can affect metallicities derived with \\R23\\ by $\\lesssim 0.15$~dex.   For the SDSS sample, AGN classification methods make negligible difference in the shape or y-intercept of the MZ relation.  For samples containing a larger fraction of starburst-AGN composite galaxies, or samples where AGN removal is not possible, we recommend the use of \\NIIOII\\ or \\NIIHa\\ metallicity diagnostics. \\item The median g'-band aperture covering fraction of our SDSS sample is 34.2\\%.  This covering fraction is insufficient for metallicities to represent global values at high masses (M$>10^{10}$\\Msun).   The Nearby Field Galaxy global MZ  is $0.1 - 0.15$~dex lower than the SDSS MZ relation at M$>10^{10}$\\Msun.  Therefore, the metallicity at which the SDSS MZ relation turns over is dependent on both the choice of metallicity calibration, and on the aperture size. \\end{itemize}  We recommend that metallicities be converted into either the \\citet{Pettini04} method or the \\citet{Kewley02a} method to minimize any residual discrepancies, and to maintain accurate relative metallicities compared to other calibrations.  These two diagnostics have the added benefit that at high metallicities, the Kewley \\& Dopita \\NIIOII\\ and Pettini \\& Pagel \\NIIHa\\ calibrations are relatively independent of the method used to remove AGN.  Future work into tailored photoionization models, high S/N spectroscopy of luminous stars in the Milky Way and nearby galaxies, metal recombination lines, IR fine structure lines, and temperature fluctuation studies may help resolve the metallicity discrepancy problem in the near future.  Until then, only relative metallicity comparisons are reliable."
    },
    "0801/0801.3557_arXiv.txt": {
        "abstract": "We present a catalogue of overdensities in the GOODS-South field. We find various high density peaks that are embedded in structures diffused on the entire field, up to $z\\sim$2.5. The slope of their colour-magnitude relation does not show significative evolution with $z$. We find evidence that galaxies forming these structures are more massive than galaxies located in low density regions. We also analyse the variation of galaxy properties with the associated environmental density and we find that the segregation of red galaxies with density is stronger at low redshift and at high luminosities while it gets much weaker for increasing $z$. ",
        "introduction": "The GOODS-MUSIC catalogue \\cite{grazian} covers 143 $arcmin^{2}$ over the Chandra Deep Field South. It contains photometry in 14 bands for 14847 extra-galactic objects selected with $Z<26.18$\u00a0 or with $Ks<23.8$.  About 10\\% of the objects have been observed spectroscopically. Where spectroscopy was not available we have used photometric redshifts obtained from our photometric redshift code (e.g. \\cite{fontana00}) that employs a standard \uf063$\\chi^{2}$ minimization over a large set of templates obtained from synthetic spectral models. We present the application to the GOODS-MUSIC catalogue in the interval $0.4<z<2.4$ of a simple  method to evaluate  galaxy volume density through the use of photo-z \\cite{trevese07}. Here we outline its relevant steps:   - We divide the survey volume in cells whose extension in different directions depends on the relevant positional accuracy (the angular positions being more accurate than distances).  - For each cell in space we then count neighboring objects at increasing  distance, until a number n of objects is reached. We then assign to the cell a comoving density  $\\rho= n/V_{n}$.  - We take into account the brightening of limiting absolute magnitude with increasing redshift for a given apparent magnitude limit, assigning a weight to each detected galaxy at redshift $z$.  %  - Galaxy overdensities of any shape (groups, walls, filaments etc.), are defined as  connected 3-dimensional regions with density exceeding a fixed threshold. Cells are grouped together according to a FoF technique.   Throughout the paper we adopt $\\Omega_{\\Lambda}=0.7, ~\u00a0\\Omega_M \u00a0=0.3,\u00a0~  H_0=70~ km ~ s^{-1} ~\u00a0Mpc^{-1}$.  \\begin{table} \\caption{Structures in the GOODS-South Field} \\label{tab:structures} \\begin{tabular}{rcl} Redshift & N Obj. & References \\\\ \\hline 0.65 - 0.75 &  123 &  Gilli et al. 2003, Adami et al. 2005, Trevese et al. 2007   \\\\ 1.0  - 1.1 & 108 &  Adami et al. 2005, Trevese et al. 2007, Di\\`{a}z Sa\\'{n}chez et al. 2007   \\\\ 1.6& 45  & Castellano et al. 2007 \\\\ 2.2 - 2.4   & 113 &   \\\\ \\hline \\end{tabular} \\end{table} ",
        "conclusions": "We have estimated the galaxy volume density using photometric redshifts in the GOODS-MUSIC catalogue. We selected high-density peaks up to $z$=2.5. Among them we find a structure at redshift 1.6 with properties of a forming cluster of galaxies. A population of red galaxies is present in all of them. The slope of the colour-magnitude relation of this red population is comparable to the slope measured for lower $z$ structures. We have also  analyzed the variation with density of the distribution of galaxy colours, which is related with the distribution of galaxy types. We find that the red galaxy segregation at high densities is stronger at low redshift and at high luminosities but it gets weaker for increasing $z$. At the highest $z$ probed, a blue and star forming population prevails at every density and luminosity value. These results show that the trend previously obtained using a sample of spectroscopic data \\cite{cucciati06} extends to higher redshifts. We have also found that massive galaxies prevails, at every redshift, in dense environments. With the help of photometric redshifts it will be possible to detect overdensities and observe the onset and growth of the bimodality of galaxy types distribution, at the highest reachable redshifts. This will allow to constrain models for galaxy evolution in clustered environments."
    },
    "0801/0801.3282_arXiv.txt": {
        "abstract": "We compare the results of the mark correlation analysis of galaxies in a sample from the Sloan Digital Sky Survey and from two galaxy catalogs obtained by semi-analytical galaxy formation models implemented on the Millennium Simulation. We use the MOPED method to retrieve the star formation history of observed galaxies and use star formation parameters  as weights to the mark correlations. We find an excellent match between models and observations when the mark correlations use stellar mass and luminosity as weights. The most remarkable result is related to the mark correlations associated to the evolution of mass assembly through star formation in galaxies, where we find that semi-analytical models are able to reproduce the main trends seen in the observational data. In addition, we find a good agreement between the redshift evolution of the mean total mass formed by star formation predicted by the models and that measured by MOPED. Our results show that close galaxy pairs today formed more stellar mass $\\sim 10$~Gyr ago than the average, while more recently this trend is the opposite, with close pairs showing low levels of star formation activity. We also show a strong correlation in simulations between the shape and time evolution of the star formation marks and the number of major mergers experienced by galaxies, which drive the environmental dependence in galaxy formation by regulating the star formation process. ",
        "introduction": "Tools to extract physical information from galaxy spectra have evolved spectacularly in the past five years. It is now possible to extract accurately the past star formation (SF) and metallicity histories of galaxies from their integrated stellar light \\citep[e.g.][]{heavens04,cid05,panter07,cid07}. Further, with the advent of mega-spectroscopic surveys, it is now possible to have samples large enough  to study the physical parameters that drive the formation and evolution of galaxies as a function of galaxy properties. One obvious parameter is the galaxy mass, and it seems now firmly established that mass has an influence on the star formation rates (SFRs) of galaxies \\citep[e.g.][]{cowie96}. However, and open question is if other parameters besides mass influence the shape and formation of galaxies \\citep[e.g.][]{gao05,angulo08}.  In a recent study, \\citet{sheth06} used mark correlations of the MOPED SF and metallicities of SDSS galaxies to determine how these quantities were correlated. They found that close pairs ( $ < 1$~Mpc) at $z \\sim 0.1$ dominated the SF activity in the past relative to the mean, while today the SF is dominated by pairs that are separated by more than $10$~Mpc. Similar trends were found for stellar metallicity. While this result fits well within the framework of galaxy ``downsizing'' it is unclear how current models of galaxy formation are able, or not, to reproduce it. In fact, a halo occupation distribution (HOD) study shows that the trends found in \\citet{sheth06} cannot be reproduced by changing the free parameters in the HOD model, thus pointing toward a need for some environmental dependence, which the current HOD paradigm does not incorporate. Thus, it is useful to look into numerical simulations how environment can influence galaxy evolution.  In this Letter, we explore how semi-analytical models (SAMs) of galaxy formation perform at reproducing the mark correlations computed in \\citet{sheth06}. We use the Millennium Simulation and two semi-analytical recipes implemented on it to follow galaxy evolution. We find that models are successful at reproducing the trends observed in the SDSS mark study. The models therefore do well at reproducing the environmental trends of SF suggesting that mass is not the only driver of galaxy formation and evolution. ",
        "conclusions": "In this Letter, we have investigated how the SF in galaxies depends on the scale by comparing the mark correlation analysis of observed and model galaxies. Our results indicate that current semi-analytical galaxy formation models, especially the model by \\cite{bower06} (Durham model), implemented on the large Millennium Simulation are able to reproduce the trends seen in the mark correlations obtained for SDSS galaxies, when the marks are derived from the SFH of galaxies (as measured by the MOPED algorithm). Galaxies in close pairs seen nowadays formed stellar mass at a higher rate than the average population at $z\\sim2$, and at lower redshifts they show low levels of SF activity compared to the average. Using the predicted mark correlations from the models, we find that for galaxies in $z=0$ close pairs the increase in their SFR at high-redshifts is intimately related to an excess in the number of major merger events they experienced in that epoch. From Fig.~\\ref{fig:3} it seems that the measured MOPED SF marks follow closely the marks of the numbers of mergers and seems to suggest that SF follows merger activity. It follows in this scenario that the observed downsizing in SF properties of galaxies was driven by the environment primarily through mergers, which accelerated the evolution of massive galaxies seen today in close pairs (Fig.~\\ref{fig:1}). The overall context of these results gives support to a natural path for galaxy evolution which proceeded via a nurture way, with denser environments playing a fundamental role in defining galaxy properties (see detailed discussion in \\citealt{mateus07} and \\citealt{bundy06}, and references therein).  \\begin{figure} \\plotone{f4.eps} \\caption{(a) Comparison between the logarithm of the mean total stellar mass (SFM; in solar mass) formed in the same redshift bins shown in Fig.~\\ref{fig:3} for the models and for the SDSS data. (b) Distribution of the recent SFM for model galaxies located in high-density regions.} \\label{fig:4} \\end{figure}  The correlation properties of SF cannot be reproduced by models that only use the dark halo mass of the galaxy as the only parameter that drives the SF. This result is also related to the evidence of an environmental dependence of halo formation times found by \\citet{harker06} and \\citet{maulbetsch07}. Actually, current implementations of the HOD paradigm do not include any environmental or spatial correlation. They can therefore not possibly account for the SF effect, which is clearly shown here to be an environmental effect.  It is therefore surprising that current SAMs (especially the Durham one) can reproduce the SF marks so remarkably well, given the complex processes that involve gas-baryon physics. Further, since we had no access to the simulations to fine tune their parameters, the obtained agreement is even more remarkable. The strong correlation between the shape of the SF and the number of mergers mark correlations, which is our main finding,  points toward a possible explanation for regulating the SF activity, which is something that galaxy formation models have inbuilt. However, if this was only the case the Munich models would be similar to the Durham ones. Their main difference resides in the way the merger process is translated to a SF recipe. In order to investigate this issue, in Fig.~\\ref{fig:4}a we show the evolution of the mean SFM  in the same redshift bins of Fig.~\\ref{fig:3}, where the scale dependence of the fluctuations relative to these mean values was previously shown. Galaxies in the Munich model tend to form more stars at recent times than the Durham model but such difference is not too significative to produce the positive correlation shown in Fig.~\\ref{fig:3}. Overall, the model galaxies form less stars at $z > 0.5$ compared to SDSS results and show an excess of SF in the first redshift bin. In Fig.~\\ref{fig:4}b we compare the distributions of the recent SFM predicted by the models but now restricting to galaxies located in dense regions. These galaxies have local galaxy density values above the 25th percentile of the distribution of this quantity estimated through a  $k$-nearest neighbor approach (with $k=10$). Galaxies from the Munich model which inhabit dense environments formed 3.3 times more stars in the first time bin shown in Fig.~\\ref{fig:3} than their counterparts from the Durham model. This higher recent SF activity predicted by the Munich model for galaxies in high-density regions could be the main factor responsible by the positive correlation at small scales shown in Fig.~\\ref{fig:3}. This conclusion is reinforced when we look the SFM distributions for the other time bins, which tend to be similar at high-z, where the mark correlations have almost the same values. Additionally, when we redo the mark analysis for the models considering the same SFM distributions for galaxies in high-density regions, we obtain a lower correlation for the Munich model in the first time bin, confirming that the recent SF activity of galaxies in dense environments predicted by the models is a key factor which defines the differences in their mark correlations. This is also related to the way that models treat energy feedback from stars and AGN and how gas is made available to galaxies to cool and form stars \\citep[e.g][]{BDN07,DT08,DB08}. Clearly, without further exploration of parameter space in the SAMs it is difficult to assert which parameter is the culprit of the good fit shown in Fig.~\\ref{fig:3}. Current SAMs have no specific environmental dependence beyond a scale of $1$~Mpc besides that of the merger process of dark matter and therefore their success must be linked to the way the merger process is then translated into gas forming stars."
    },
    "0801/0801.4377_arXiv.txt": {
        "abstract": "We studied the geometry of the inner (AU-scale) circumstellar environment around the Herbig~Be star MWC\\,147. Combining, for the first time, near- (NIR, $K$~band) and mid-infrared (MIR, $N$~band) spectro-interferometry on a Herbig star, our VLTI/MIDI and AMBER data constrain not only the geometry of the brightness distribution, but also the radial temperature distribution in the disk. For our detailed modeling of the interferometric data and the spectral energy distribution (SED), we employ 2-D radiation transfer simulations, showing that passive irradiated Keplerian dust disks can easily fit the SED, but predict much lower visibilities than observed. Models of a Keplerian disk with emission from an optically thick inner gaseous accretion disk (inside the dust sublimation zone), however, yield a good fit of the SED and simultaneously reproduce the observed 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 strong contributions from the outer dust disk. ",
        "introduction": "\\label{sec:1}  Herbig Ae/Be (HAeBe) stars are intermediate-mass objects which are still accreting material, probably via a circumstellar disk composed of gas and dust. Understanding the structure of these disks and the processes through which they interact with the central star is critical for our understanding of the formation process of stars. Since, until recently, the spatial scales of the inner circumstellar environment (a few AU) were not accessible to infrared imaging observations, conclusions drawn on the 3-D geometry of the circumstellar material were mostly based on the modeling of the SED (e.g.\\ Hillenbrand et al.\\ \\cite{hil92}). However, these fits are known to be ambiguous (e.g.\\ Men'shchikov \\& Henning \\cite{men97}) and have to be complemented with spatial information, as provided by infrared interferometry.  The first systematic studies of HAeBes using the technique of infrared long-baseline interferometry have revealed that for most HAeBes the NIR size correlates with the stellar luminosity $L$ following a simple $R \\propto L^{1/2}$ law.  This suggests that the NIR continuum  emission mainly traces hot dust at the inner sublimation radius. However, for more luminous Herbig Be stars, the NIR-emitting structure is more compact than predicted by the size-luminosity relation (Monnier et al.\\ \\cite{mon02}).  To further investigate the origin of the ``undersized'' Herbig~Be star disks, we observed the Herbig~Be star MWC\\,147 with the VLTI. First infrared interferometric observations of this star by Akeson et al.\\ \\cite{ake00} indicated that the NIR-emitting region is surprisingly compact ($\\sim0.7$~AU, assuming a uniform disk profile). In the course of three ESO open time programmes and using the 8.2~m UT telescopes, we obtained seven VLTI/MIDI measurements and one VLTI/AMBER measurement. The MIDI observations cover baseline lengths between 39 and 102~m. From the AMBER data, one wavelength-dependent visibility, corresponding to a 101~m baseline, could be extracted. For our modeling of MWC\\,147, we adopt the stellar parameters by {Hern{\\'a}ndez} et al.\\ \\cite{her04}, namely a spectral type of B6, a distance of 800~pc, a bolometric luminosity of 1550~$L_{\\odot}$, a mass of 6.6~$M_{\\odot}$, and a stellar radius of 6.63~$R_{\\odot}$. ",
        "conclusions": "\\label{sec:4}  Our VLTI interferometric observations of MWC\\,147 constrain, for the first time, the inner circumstellar environment around a Herbig~Be star over the wavelength range from 2 to $13~\\mu$m. We find evidence that the NIR emission of MWC\\,147 is dominated by the emission from optically-thick gas located inside the dust sublimation radius, while the MIR also contains contributions from the outer, irradiated dust disk.  Our study demonstrates the power of infrared spectro-interferometry to probe the inner structure of the disks around young stars and to disentangle multiple emission components. Future investigations on YSO accretion disks will benefit substantially from the proposed 2nd~generation VLTI instruments, such as MATISSE, increasing not only the number of recorded baselines, but also expanding the spectral coverage to the $L$ and $M$ bands."
    },
    "0801/0801.1764_arXiv.txt": {
        "abstract": "We analyze the bending of light by galaxies or clusters of galaxies in the presence of the $\\La$ term. Going over to the Friedmann-Robertson-Walker (FRW) coordinates, used in fact for the description of actual observations, we demonstrate that the cosmological constant does not influence practically the lensing effect. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.2767_arXiv.txt": {
        "abstract": "{The spectra of the planetary nebula \\object{NGC\\,2392} is reanalysed using spectral measurements made in the mid-infrared with the {\\em Spitzer Space Telescope}. The aim is to determine the chemical composition of this object.  We also make use of {\\em IUE} and ground based spectra.  Abundances determined from the mid-infrared lines, which are insensitive to electron temperature, are used as the basis for the determination of the composition, which are found to differ somewhat from earlier results. The abundances found, especially the low value of helium and oxygen, indicate that the central star was originally of rather low mass. Abundances of phosphorus, iron, silicon and chlorine have been determined for the first time in this nebula. The variation of electron temperature in this nebula is very clear reaching quite high values close to the center.  The temperature of the central star is discussed in the light of the high observed stages of ionization. The nebular information indicates the spectrum of the star deviates considerably from a blackbody.} ",
        "introduction": "\\object{NGC\\,2392} (PN~G197.8+17.3) is a bright planetary nebula with a rather high radial velocity, located in the direction of the galactic anticenter. A photograph of the nebula is shown in Fig.\\,1. The nebula has a bright inner ellipsoidal, almost round, shell of about 18\\arcsec~diameter. This is surrounded by an almost spherical region of considerably lower emission with a diameter of about 40\\arcsec. Both the inner and outer regions contain some structure.  Analysis of the kinematics of the nebula (O'Dell \\& Ball, \\cite{odell}) indicates that the inner region is expanding quite rapidly with respect to the outer region which is also slowly expanding.  The nebula is located 17 degrees above the galactic plane and has only a small extinction. Because of its brightness it is considered a nearby nebula: most of the uncertain distance estimates vary between 1 and 2 kpc.  The nebula has a bright central star (V=10.53) which has been studied by several authors. Pauldrach et al. (\\cite{paul}) have compared the stellar spectrum with a model atmosphere and conclude that the star has an effective temperature T$_{eff}$=40\\,000 K. While this is slightly higher than the hydrogen Zanstra temperature of about 36\\,000 K, it is considerably less than the ionized helium Zanstra temperature which is close to 70\\,000 K. Pauldrach et al. (\\cite{paul}) have noted this discrepancy and admit that they are unable to explain the additional He-ionizing photons which are needed to explain the high T$_z$(HeII). Tinkler \\& Lamers (\\cite{tinkler}) prefer assigning a much higher effective temperature to the star based on the high ionized helium Zanstra temperature and the high energy balance temperature: T$_{eff}$=73\\,000 K. It is clear that the excitation in the nebula is much higher than can be accounted for by a temperature of T$_{eff}$=40\\,000 K. This was already noted by Natta et al. (\\cite{natta}) who showed that higher temperature ionizing radiation was necessary for all ions found in the nebula: the higher the ionization potential of the ion studied the higher the ionizing radiation temperature must be. Why this result is not found in the analysis of the stellar spectrum is not understood.  The purpose of this paper is to study the element abundances in this nebula with the help of mid-infrared spectra, in the hope that the chemical abundances will shed some light on the evolution of this nebula-central star combination. Abundances have been studied earlier by Barker (\\cite{barker}) and by Henry et al. (\\cite{henry}).  Both of these groups use optical nebular spectra (taken by themselves) and ultraviolet {\\em IUE} spectra. Barker (\\cite{barker}) has measurements taken at five different positions in the nebula. He uses low dispersion {\\em IUE} measurements taken with the small aperture (3\\arcsec~diameter).  Henry et al. (\\cite{henry}) also take spectra at five different positions in the nebula but they use low dispersion {\\em IUE} measurements taken with the large aperture (10\\arcsec x 23\\arcsec). No abundance changes as function of the position in the nebula are found by either group. Both groups agree on the abundances of N, O and Ne to within 20\\%. However their carbon abundances differ by more than a factor of 6 and the helium abundances differ by almost 30\\%. Henry et al. (\\cite{henry}) do not discuss the abundance of other elements, while Barker also determines sulfur and argon.  We have measured the spectrum of NGC\\,2392 in the mid-infrared with the IRS spectrograph of the {\\em Spitzer Space Telescope} (Werner et al.\\,\\cite{werner}). The use of the mid-infrared spectrum permits a more accurate determination of the abundances.  The reasons for this have been discussed in earlier studies (e.g. see Pottasch \\& Beintema \\cite{pott1}; Pottasch et al.  \\cite{pott2}, \\cite{pott4}; Bernard Salas et al. \\cite{bernard}), and can be summarized as follows.  The most important advantage is that the infrared lines originate from very low energy levels and thus give an abundance which is not sensitive to the temperature in the nebula, nor to possible temperature fluctuations. Furthermore, when a line originating from a high-lying energy level in the same ion is observed, it is possible to determine an average electron temperature at which the lines in that particular ion are formed. When this average temperature for many ions is determined, it is possible to make a plot of electron temperature against ionization potential, which can be used to determine the average electron temperature (or range of temperatures) for which ions of a particular ionization potential are observed. As will be shown, this is particularly important in NGC\\,2392 for which a large gradient in electron temperature is found. This has not been taken into account in the earlier abundance studies of this nebula.  Use of the infrared spectra have further advantages. One of them is that the number of observed ions used in the abundance analysis is approximately doubled, which removes the need for using large Ionization Correction Factors (ICFs), thus substantially lowering the uncertainty in the abundance. A further advantage is that the extinction in the infrared is almost negligible, eliminating the need to include sometimes large correction factors. This is not very important in this nebula. The number of elements which can be measured increases when including the infrared. In NGC\\,2392 phosphorus, iron and silicon can be measured in the mid-infrared.  This nebula has not been measured by {\\em ISO}, because this part of the sky was not available to this satellite.  This paper is structured as follows. First the Spitzer spectrum of NGC\\,2392 is presented and discussed (in Sect.\\,2). Then the intrinsic H$\\beta$ flux is determined using both the measurements of the infrared hydrogen lines and the radio continuum flux density (Sect.\\,3). The visible spectrum of the nebula is presented in Sect. 4 together with a new reduction of the ultraviolet ({\\em IUE}) spectrum. This is followed by a discussion of the nebular electron temperature and density and the chemical composition NGC\\,2392 (Sect.\\,4). A comparison of the resultant abundances with those in the literature is given in Sect.\\,5. In Sect.\\,6 the central star is discussed especially in relation to the nebular spectrum.  In Sect.\\,7 a general discussion and concluding remarks are given. ",
        "conclusions": "The nebular abundances of eleven elements have been determined. It has been shown that those elements which have not been created in stellar evolution (O, Ne, S, Ar, Cl and possibly Si) have abundances of about a factor two lower than the solar abundance. This is in line with the expectation of a gradient in the stellar abundances as a function of distance from the galactic center.  The helium abundance is lower than solar; this indicates that no helium has been formed during the evolution of the central star. Carbon has apparently been formed in a third dredge-up as it is more abundant than oxygen. Nitrogen has also been formed, probably in the first dredge-up. The abundances are consistent with the central star having evolved from an object having about 1.7M\\smallsun. The low iron abundance is also found in almost all PNe studied and seems to be caused by its being tied up in dust.  The electron temperature is found to vary strongly with the ionization potential of the ion used for its determination. This indicates a strong gradient in the nebula, higher than in any PNe yet studied. The reason for this high gradient is not clear. It could be related to the fact that diffuse X-ray emission has been detected within its inner shell (Guerrero et al. \\cite{guer}), which in turn may be caused by the existence of the exceptionally large expansion velocity of the inner shell.  Finally, the presence of high stages of ionization is confirmed. Both \\ion{O}{iv} and \\ion{Ne}{v} are clearly present. This seems to indicate that the stellar temperature is considerably higher that the value of 40,000 K found from a study of the central star. The suggestion that an additional star is present is discussed, but no solution to this problem can be given."
    },
    "0801/0801.2598_arXiv.txt": {
        "abstract": "{} {Both the black hole mass and the X-ray luminosity of AGNs have been found to be anti-correlated with the normalized excess variance ($\\sigma _{\\rm rms}^2 $) of the X-ray light curves. We investigate which correlation with $\\sigma _{\\rm rms}^2 $ is the intrinsic one.} {We divide a full sample of 33 AGNs (O' Neill et al. 2005) into two sub-samples. The black hole masses of 17 objects in sub-sample 1 were determined by the reverberation mapping or the stellar velocity dispersion. The black hole masses of the remaining 16 objects were estimated from the relationship between broad line region radius and optical luminosity (sub-sample 2). Then partial correlation analysis, ordinary least squares regression and K-S tests are performed on the full sample and the sub-samples, respectively.} {We find that $\\sigma _{\\rm rms}^2 $ seems to be intrinsically correlated with the black hole mass in the full sample. However, this seems to be caused by including the sub-sample 2 into the analysis, which introduces an extra correlation between the black hole mass and the luminosity and strengthens any correlation with the black hole mass artificially. Therefore, the results from the full sample may be misleading. The results from the sub-sample 1 show that the correlation between $\\sigma _{\\rm rms}^2 $ and the X-ray luminosity may be the intrinsic one and therefore the anti-correlation between $\\sigma _{\\rm rms}^2 $ and the black hole mass is doubtful.} {} ",
        "introduction": "X-ray emission from active galactic nuclei (AGNs) exhibits variability on time scales from minutes to days. This indicates that X-rays are likely to be emitted from the inner most regions of AGNs and the variability may be related to the important properties of the central engine. Lawrence \\& Papadakis (1993) utilized long term \\emph{EXOSAT} observations to investigate the power density spectra of 12 AGNs. They found the power density spectra could be described as a power law $P \\propto \\nu ^{ - \\alpha } $ with a mean index $\\alpha  = 1.55$, and the amplitude was anti-correlated with the X-ray luminosity. Detailed studies of the power density spectra have been performed using \\emph{RXTE} and \\emph{XMM-Newton} data. The universal relation between the black hole mass and the \"break time\" in the power density spectra was found both in stellar mass and supermassive black holes (e. g. Uttely \\& McHardy 2005). A tighter relation was discovered when the bolometric luminosity was involved, i.e. $T_B  \\propto M^2 /L $ (McHardy et al. 2006). However, due to the limited observation data, the accurate power density spectra are only available for a small number of AGNs. As an alternative, the normalized excess variance ($\\sigma _{\\rm rms}^2 $) can be easily calculated, and it was found that $\\sigma _{\\rm rms}^2 $ is anti-correlated with X-ray luminosity (e.g. Almaini et al. 2000).   As a result of the progress in determining the black hole masses in AGNs, the relation between the variability and black hole mass has also been investigated. Lu \\& Yu (2001) found the anti-correlation between the black hole mass and $\\sigma _{\\rm rms}^2 $, and suggested this correlation was an intrinsic one, rather than the apparent anti-correlation between the X-ray luminosity and $\\sigma _{\\rm rms}^2 $. Several authors also addressed this problem following Lu \\& Yu's work. (e.g. Bian \\& Zhao 2003; Papadakis 2004; O' Neill et al. 2005). These authors confirmed the anti-correlation between the black hole mass and $\\sigma _{\\rm rms}^2 $, and some models have been subsequently constructed to explain this correlation.  In this paper we revisit this problem, by studying the sample of O' Neill et al. (2005), which includes 33 AGNs and uses nearly the same time scale for all these objects. The black hole masses of 17 objects in this sample were determined by the reverberation mapping or the stellar velocity dispersion (we denote these objects as sub-sample 1 in the following). The black hole masses of the remaining 16 objects were estimated from the relationship between broad line region radius and optical luminosity (we denote these objects as sub-sample 2 in the following). We find that the optical luminosity (which is used in determining the black hole mass for sub-sample 2) has obvious correlation with the X-ray luminosity in 2-10 keV band, as shown in Figure 1 (a). The values of the correlation coefficients between optical luminosity and X-ray luminosity for sub-sample 1 and sub-sample 2 are 0.907 and 0.910, respectively. Even if the datum of NGC 4395 is excluded from sub-sample 2, the value of the correlation coefficient is still 0.819. It is well known that there is a strong correlation between the optical luminosity and the black hole mass (Kaspi et al. 2000). Thus if $\\sigma _{\\rm rms}^2 $ is intrinsically correlated with one of them, an artificial correlation with the other will appear. However, there will be an additional artificial correlation for sub-sample 2 due to the utilization of the optical luminosity in determining the black hole mass, whereas the black hole mass of sub-sample 1 is independently obtained by the reverberation mapping or the stellar velocity dispersion. Therefore, the result from the full sample (especially for that from the sub-sample 2) may be misleading. Furthermore, due to the less reliable black hole mass of sub-sample 2, some unclear systematic biases may be introduced into the analysis. To find out which correlation is intrinsic, we perform the partial correlation analysis to sub-sample 1 and sub-sample 2 separately in \\S2.1. The ordinary least squares regression results are shown in \\S2.2 as another approach to this problem. K-S tests are performed in \\S2.3 to test whether sub-sample 1 and sub-sample 2 are drawn from the same parent population. In \\S3 we discuss our results and make conclusions. ",
        "conclusions": "In \\S2, we have performed the partial correlation analysis and the regression on the sample and found that the apparent intrinsic correlation between $\\sigma _{\\rm rms}^2 $ and the black hole mass is likely to be caused by including the sub-sample 2 into the analysis. Because the black hole masses of AGNs in sub-sample 2 were estimated from their optical luminosity which in turn is positively correlated with their X-ray luminosity, an extra correlation between the black hole mass and X-ray luminosity will be introduced by the sub-sample 2. If the X-ray luminosity is the primary quantity, then this will artificially strengthen any correlation with black hole mass. We therefore should exclude them when investigating the intrinsic correlation with $\\sigma _{\\rm rms}^2 $. According to the results from the sub-sample 1, we conclude that the correlation between $\\sigma _{\\rm rms}^2 $ and the X-ray luminosity may be the intrinsic one, whereas the apparent correlation between $\\sigma _{\\rm rms}^2 $ and the black hole mass is doubtful. Our K-S tests also suggest that sub-samples 1 and 2 are not likely drawn from the same parent population.  As discussed in Lu \\& Yu (2001), several mechanisms may be responsible for the correlation between $\\sigma _{\\rm rms}^2 $ and the X-ray luminosity, such as the hot-spot model, the obscurative variability and so on. After the apparent correlation between $\\sigma _{\\rm rms}^2 $ and the black hole mass was discovered, some models accounting for this correlation were proposed (e.g. O' Neill et al. 2005, Pessah 2007). However, it needs to be verified whether the correlation is intrinsic. Although the black hole masses of about three dozen AGNs have been determined by the reverberation mapping method, the size of our sample is still limited due to the lack of long enough and high quality observation data of these objects. More conclusive results could be obtained when more and higher quality data become available."
    },
    "0801/0801.0554_arXiv.txt": {
        "abstract": "Microwave background temperature and polarization observations are a powerful way to constrain cosmological parameters if the likelihood function can be calculated accurately. The temperature and polarization fields are correlated, partial sky coverage correlates power spectrum estimators at different $l$, and the likelihood function for a theory spectrum given a set of observed estimators is non-Gaussian. An accurate analysis must model all these properties.  Most existing likelihood approximations are good enough for a temperature-only analysis, however they cannot reliably handle a temperature-polarization correlations. We give a new general approximation applicable for correlated Gaussian fields observed on part of the sky. The approximation models the non-Gaussian form exactly in the ideal full-sky limit and is fast to evaluate using a pre-computed covariance matrix and set of power spectrum estimators. We show with simulations that it is good enough to obtain correct results at $l\\agt 30$ where an exact calculation becomes impossible. We also show that some Gaussian approximations give reliable parameter constraints even though they do not capture the shape of the likelihood function at each $l$ accurately. Finally we test the approximations on simulations with realistically anisotropic noise and asymmetric foreground mask. ",
        "introduction": "The Cosmic Microwave Background (CMB) anisotropies are a powerful cosmological probe as they depend simply on the primordial inhomogeneities, content and  geometry of the Universe.  If the perturbations are Gaussian, the full-sky power spectra of the CMB anisotropies and their polarization contain all of the cosmological information.  Information in the polarization power spectra can help to break degeneracies that are present if only temperature information is used, and also helps to reduce cosmic variance uncertainty. Parameter constraints can therefore be significantly improved by using polarization information even if the data is significantly noisier than the temperature.   An accurate joint likelihood analysis of the CMB temperature and polarization data is crucial to estimate cosmological parameters reliably. In principle this is straightforward at linear order if the primordial perturbations are Gaussian as the distribution can be calculated exactly. However calculating the likelihood exactly from partial sky data with anisotropic noise is computationally prohibitive except at low $l$ because large matrices need to be inverted.  Most analyses therefore rely on approximations to the likelihood function at high $l$, using only the information in a set of estimators for the power spectra and a covariance estimated (or calibrated) from simulations. An alternative approach not considered further here would be to use the Gibbs sampling approach of Ref.~\\cite{Wandelt:2003uk,Larson:2006ds}, though this has serious convergence problems of its own~\\cite{Chu:2004zp}.  If the likelihood of the theory power spectrum $C_l$ as a function of the measured estimators $\\Ch_l$ were Gaussian the likelihood could be calculated straightforwardly from the measured $\\Ch_l$. However the distribution is non-Gaussian because for a given temperature power spectrum $C_l$, the $\\Ch_l$, a sum of squares of Gaussian harmonic coefficients, have a (reduced) $\\chi$-squared distribution. At large $l$ the distribution does tend to Gaussian by the central limit theorem; for example the mean and maximum likelihood values of $\\Ch_l$ converge as $1/l$. However the precision with which we can hope to measure the cosmological parameters also improves at $1/l_{\\text{max}}$, so the relative bias due to the non-Gaussianity is potentially independent of $l$. On all scales the distribution must be modeled carefully to get unbiased cosmological parameter constraints.  The importance of the non-Gaussianity of the temperature likelihood function at low $l$ is well known, and there are several well established likelihood approximations to model it~\\cite{Bond:1998qg,Verde:2003ey,Smith:2005ue,Percival:2006ss}. Current polarization data only contributes interesting information at low $l$ where an exact likelihood can be used~\\cite{Gorski94,Slosar:2004fr,Page:2006hz}, however in the future the small-scale polarization signal will be less noise dominated and contain useful information. On small scales the likelihood function cannot be computed exactly in reasonable time, and the likelihood function is significantly more complicated than for the temperature because the temperature and polarization fields are correlated. The only existing attempt to model the polarized likelihood function at high $l$, Ref.~\\cite{Percival:2006ss}, relies on variable transformations that are not guaranteed to be well defined, and is untested in practice. We give a new general well-defined likelihood approximation that can be used with partial-sky Gaussian polarized CMB data, or any other set of correlated Gaussian fields observed on part of the sky. It is exact in the full-sky limit, and can easily be calibrated from simulations. We also discuss under what circumstances a Gaussian likelihood approximation is reliable.  The layout of the paper is as follows.  In section~\\ref{fullsky_like} we present a brief overview of the exact full-sky likelihood function for isotropic noise, and discuss the accuracy required in general for unbiased parameter estimation. We start section~\\ref{like_approx} with a review of various temperature likelihood approximations available in the literature, discuss the accuracy of the various Gaussian approximations, and move on to derive a new general likelihood approximation (Eq.~\\eqref{new_approx}) that is exact in the full-sky limit.  In section~\\ref{tests} we test the approximations by comparing with the exact likelihood function for azimuthal sky cuts and consistency with the binned likelihood. Finally in section~\\ref{anisotropic_parameters} we check the approximations with realistically anisotropic noise and demonstrate consistent parameter estimation from simple Planck-like simulations. Some mathematical and analysis details are described in the appendices: appendix~\\ref{matrices} gives identities relating expressions with symmetric matrices to expressions with a vector of components; appendix~\\ref{goodness} calculates the non-Gaussian correction to the full-sky effective chi-squared; appendix~\\ref{multimap} gives results for the likelihood function when using cross-power spectrum estimators from different maps; appendix~\\ref{cutsky} reviews the basic Pseudo-$C_l$ estimator and exact likelihood formalism and appendix~\\ref{Planck} slightly generalizes previous hybrid Pseudo-$C_l$ estimators for anisotropic noise and gives details of our Planck-like test simulations.  We assume Gaussianity and statistical isotropy of the fields, and focus on the idealized case of pure CMB observations without the complications of foregrounds, point sources, non-linear effects, anisotropic beams, and other observational artefacts. Generalizing our work to more realistic situations will be crucial for application to real data. If the fluctuations turn out to be significantly non-Gaussian or anisotropic a more complicated analysis may also be required. ",
        "conclusions": "In this paper we have attempted to find solutions to the problems facing the likelihood analysis of the CMB temperature and polarization estimators on small scales.  With realistic data we need to be able to calculate the likelihood accurately from partial sky observations. Previous attempts have established some excellent approximations to model the non-Gaussianity of the temperature likelihood function.  However, no good general approximation has been derived to model the polarized likelihood.  At large $l$ computing the likelihood function exactly is computationally prohibitive and the correlation between the temperature and polarization fields makes it more complicated than for the temperature field only.  We gave a new general approximation that can account for this correlation and is exact on the full-sky.  This new approximation is fast to evaluate as it involves a pre-computed covariance independent of $\\mC_l$, and appears to be more than adequate to obtain robust parameter constraints from clean small-scale CMB temperature and polarization data.  In summary, our conclusions regarding the modelling of the likelihood function of power spectrum estimators are:   \\begin{itemize} \\item In the case of binned power spectra, the number of modes per bin ($n_m/n_b$) must be much larger than the number of bins ($n_b$) for non-Gaussian corrections to the likelihood function to be unimportant in all cases; i.e. $n_b \\ll \\sqrt{n_m}$ is required to ensure that parameter bias is much smaller than the error bar.  \\item A Gaussian approximation with fixed fiducial-model covariance gives unbiased results for smooth power spectra at high $l$, but error bars have some dependence on the choice of the fiducial model.  Goodness-of-fit estimators $\\chieff$ can be misleading even for small differences between the fiducial and true model.  \\item A Gaussian approximation with covariance that varies with parameters can give reliable results at high $l$ for smooth spectra, but only if the determinant-term is consistently included; the quadratic approximation without determinant, $\\cll_Q$, is biased in general.  \\item The new likelihood approximation presented in Section~\\ref{sec:newlike} appears to work well for power spectrum estimators with correlated fields and can give nearly optimal results when applied to good power spectrum estimators. It is fast to evaluate as it relies on a pre-computed fiducial-covariance matrix, but is insensitive to small errors in the fiducial model. We recommend it for future work.  \\item Most likelihood approximations with binned estimators ($n_b \\ll \\sqrt{n_m}$) can produce consistent results by the central limit theorem; for smooth power spectra consistency of parameter constraints with those from binned power spectra is a good check. \\end{itemize}  Since the new likelihood approximation is based on estimators and a covariance matrix, it is likely to generalize well to more realistic data where additional uncertainties, non-Gaussianities and correlations can be accounted for via changes to the estimator covariance. It is also likely to produce good results down to low $l$ if positive-definite estimators are used, though this has not been the focus of this paper. Complications such as correlated noise may be well encapsulated in the covariance of a set of maximum-likelihood (or similar) estimators, giving a fast alternative to much slower brute-force likelihood calculations. If the approximation is nearly correct, importance sampling techniques could be used to correct the results with a much smaller number of high-accuracy calculations.    We have not touched at all on the complications of foreground modelling, point sources, non-linear and non-Gaussian anisotropies (e.g. due to SZ), beam uncertainties, or a plethora of other real-world complications. Extending our work to account for these will be crucial for the correct interpretation of future data."
    },
    "0801/0801.2417_arXiv.txt": {
        "abstract": "One of the challenges of the \\textit{Swift} era has been accurately determining \\Ep\\ for the prompt GRB emission.  RHESSI, which is sensitive from 30 keV to 17 MeV, can extend spectral coverage above the \\textit{Swift}-BAT bandpass. Using the public \\textit{Swift} data, we present results of joint spectral fits for 26 bursts co-observed by RHESSI and \\textit{Swift}-BAT through May 2007. We compare these fits to estimates of \\Ep\\ which rely on BAT data alone.  A Bayesian \\Ep\\ estimator gives better correspondence with our measured results than an estimator relying on correlations with the \\textit{Swift} power law indices. ",
        "introduction": "The Reuven Ramaty High-Energy Solar Spectroscopic Imager (RHESSI) \\citep{lin02} is a dedicated solar observatory. RHESSI's nine germanium detectors are sensitive from 30 keV to 17 MeV, with excellent resolution in energy (1-5 keV) and time (1 binary $\\mu$s) \\citep{smit02}.  Each of the nine coaxial detectors is electronically segmented into front and rear segments. Because the detectors are unshielded, RHESSI observes emission from astrophysical sources like GRBs with a $\\sim$2$\\pi$ field of view.  We perform Monte Carlo simulations using the MGEANT package \\citep{stur00} to determine RHESSI's response to off-axis sources like GRBs.  RHESSI's response varies with off axis angle, so we create responses every 15 degrees.  For each response, we simulate monoenergetic photons in 192 logarithmic energy bins ranging from 30 keV -- 30 MeV.  Since RHESSI's per-detector response also varies during the spacecraft's four second spin period, we bin the annular response in six azimuthal bins and weight these bins by the total burst lightcurve to create the final response.  The RHESSI data are extracted in SSW-IDL.  We fit and subtract a time-varying background, allowing for possible periodic modulation with the spin period.  We perform spectral fitting in ISIS \\citep{citeisis}, a forward-fitting package analogous to XSPEC\\footnote{http://heasarc.gsfc.nasa.gov/docs/xanadu/xspec/}, which is extensively programmable and allows computation of rigorous fluence error estimates via exploration of the parameter space.  Energetic charged particles over time have caused radiation damage to RHESSI's germanium detectors.  This radiation damage causes broadened spectral lines due to hole trapping and a general loss of active volume. In this work, we restrict our analysis to those detectors which do not exhibit signs of radiation damage. At the time of \\textit{Swift}'s launch, six RHESSI segments were usable for spectroscopy; by May 2007 only two segments remained undamaged. In November 2007, RHESSI underwent an annealing procedure to reverse the effects of the radiation damage.  The anneal restored some lost sensitivity, but analysis of future bursts will require more sophisticated modeling of the remaining effects of radation damage on RHESSI's spectral response. ",
        "conclusions": ""
    },
    "0801/0801.0909_arXiv.txt": {
        "abstract": "This study reports pulse variation analysis results for the forth discovered accretion-powered millisecond pulsar XTE J1807-294 during its 2003 outburst observed by {\\it Rossi X-ray Timing Explorer}.  The pulsation is significantly detected only in the first $\\sim$90d out of $\\sim$150d observations.  The pulse phase variation is too complex to be described as an orbital motion plus a simple polynomial model.  The precise orbital parameters with $P_{orb}=40.073601(8)$ min and ${\\it a_x}\\sin {\\it i}=4.823(5)$ lt-ms were obtained after applying the trend removal to the daily observed 150s segments pulse phases folded with a constant spin frequency without Keplerian orbit included.  The binary barycenter corrected pulse phases show smooth evolution and clear negative phase shifts coincident with the flares seen on the light curve and the enhancements of fractional pulse amplitude.  The non-flare pulse phases for the first $\\sim$60d data are well described as a fourth order polynomial implying that the neutron star was spun-up during the first $\\sim$60d  with a rate $\\dot \\nu=(1.7\\pm0.3) \\times 10^{-13}$ Hz/s at the beginning of the outburst.  Significant soft phase lags up to $\\sim$500 $\\mu s$ ($\\sim$10\\% cycle) between 2 to 20 keV were detected for the nonflare pulse phases. We conclude that the anomalous phase shifts are unlikely due to the accretion torque but could result from the ``hot spot'' moving on the surface of neutron star. ",
        "introduction": "\\label{intro} Low Mass X-ray Binaries (LMXBs) are considered to be the progenitor of the millisecond pulsars detected in the radio band (see \\citealt{bha91} for an extensive review).  It is widely believed that the weakly magnetized, slowly rotating neutron star is gradually spun-up through the transfer of angular momentum carried by the matter from an accretion disk \\citep{alp82,rad82}.  Millisecond-time-scale variations have been observed in many LMXB systems, including kilohertz Quasi-Periodic Oscillation \\citep[kHz QPO, see][]{van05} and burst oscillation (see \\citealt{str05} for a review); however, the attempts to measure coherent millisecond pulsation from an LMXB, which provides a direct evolutionary link between radio millisecond pulsar and LMXB, were unsuccessful until the discovery of the first accretion-powered millisecond pulsar SAX 1808.4-3658 with a spin period $P_{spin} \\simeq 2.5$ ms \\citep{wij98, cha98}, in 1998.  To date, coherent millisecond pulsations have been detected in ten LMXBs with spin periods ranging from 1.67 ms to 5.5 ms \\citep{wij07,mar07,gav07,cas07,alt07}.   The fourth known accretion-powered millisecond pulsar, XTE J1807-294, was discovered on February 21, 2003 while the Proportional Counter Array (PCA) onboard the {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}) monitored the Galactic-center region \\citep{mar03a}. Its preliminary spin frequency of 190.625 Hz \\citep{mar03a} and orbital period of $40.0741 \\pm 0.0005$ min \\citep{mar03b} were immediately revealed right after its discovery.  Follow-up observations were made with many X-ray observatories. \\citet{mar03b} determined the source location, R.A.=$18^h06^m59^s.8$ DEC=$-29^{\\fdg}24^{\\arcmin}20^{\\arcsec}$ (equinox 2000, $\\sim1^{\\arcsec}$ error) by using  {\\it Chandra} observation. \\citet{cam03}, \\citet{kir04} and \\citet{fal05a} reported more precise orbital parameters and pulsation period using {\\it XMM-Newton} observations. However, \\citet{mar04} found wide swings in the apparent spin change rate vs. accretion rate, which were interpreted as the hot spot moving on the neutron star rather than from accretion torque.  The wide-band (0.5 to 200 keV) spectrum obtained by the combination of simultaneous observations from {\\it XMM-Newton}, {\\it RXTE} and {\\it INTEGRAL} is well described by an absorbed disk blackbody plus thermal Comptonization model \\citep{fal05a}.  {\\it RXTE} has processed extensive monitoring of XTE J1807-294 since February 27, 2003 with a time span of more than 150d until the source went into its quiescent state.  The light curve exhibited an exponential-like decay with a time-scale of $\\sim 19$d at the beginning of the outburst followed by a much slower decay with a time constant of $\\sim 122$d \\citep{fal05a}.  Several broad soft flares lasting for hours to days were seen during the flux decay.  The pulsed fraction increased and the pulse profile became more sinusoidal during the flares \\citep{zha06}.  Twin kHz QPOs were also detected in XTE J1807-294 by {\\it RXTE} \\citep{lin05}.  Owing to their fast rotation, the related parameters of the accretion-powered millisecond pulsars, such as orbital and spin parameters, can be precisely measured during their outbursts lasting for tens to hundreds of days. Unfortunately, the orbital and spin parameters derived by \\citet{cam03}, \\citet{kir04} and \\citet{fal05a} for XTE J1807-294 form {\\it XMM-Newton} observations were based on the presumption that the orbital period is fixed to the preliminary value proposed by \\citet{mar03b} because the time base line is too short to constrain the orbital period.  On the other hand, the $\\sim$150d's {\\it RXTE} observations allow further binary parameter refinement, that is essential for further investigation.  In this paper, we report the pulsation analyses of XTE J1807-294 using archived {\\it RXTE} data (section~\\ref{dar}).  Due to the complexity of pulse phase variation, a trend removal technique is employed to reveal precise orbital parameters for the system.  Pulse phases (binary barycenter-corrected) are found to exhibit anomalous negative shifts during the soft flares in all energy bands between 2-20 keV.  The evolution of the non-flare pulse phases shows that the neutron star was spun-up during the outburst. Significant $\\sim500$ $\\mu s$ soft lags for the non-flare pulses is detected at energies from 2 to 20 keV.  Analysis results, including possible implications for coincident soft flares and anomalous negative phase shifts are discussed in section \\ref{dis}. ",
        "conclusions": "\\label{con}  \\paragraph{1.} Due to complicated intrinsic phase variations, to obtain the precise orbital parameter, we folded our data sets with a constant frequency and used the trend removal technique to remove the local linear pulse phase trend from the observed pulse phase variation for daily observations. We obtained more precise orbital parameters consistent with the ones previously reported (except for $T_{\\pi /2}$).  \\paragraph{2.} The binary barycenter-corrected pulse phases change smoothly except for several anomalous negative phase shifts coincident with flux flares and pulse amplitude enhancements.  The pulse fractional amplitude variation indicates that there were six flares during the first $\\sim$90d of the XTE J1807-294 in its 2003 outburst.  The phase changes of non-flare pulse imply that the neutron star was spun-up during the first 60d of the outburst under the assumption that the hot spot was fixed on the surface of the neutron star during the non-flare states.  \\paragraph{3.} The analysis results of the energy-dependent pulse arrival time indicates that hard pulses lead soft pulses up to $\\sim 500$ $\\mu s$ from 2 to 10 keV, and the lead saturates from 10 to 20 keV, which agrees with the scenario proposed by \\citet{gie02} and \\citet{pou03}.  \\paragraph{4.} Our discovery of the coincidence of flare and phase shift, as well as its correlation with pulse amplitude and KHz QPO~\\citep{zha06} in XTE J1807-294, implies that probably due to weak magnetic confinement on the accreting material, the hot spot position on the accreting millisecond pulsars is sensitive to the accretion flow."
    },
    "0801/0801.2918_arXiv.txt": {
        "abstract": "{Rigorous modeling of the   \\ion{Ba}{ii} ${\\lambda}4554$ formation  is potentially interesting  since  this strongly polarized line  forms  in the solar chromosphere where the magnetic field  is rather poorly known. } { To investigate the role of    isotropic collisions with neutral hydrogen in the   formation of the polarized   \\ion{Ba}{ii} ${\\lambda}4554$  line and, thus, in the determination of the  magnetic field. } { Multipole relaxation and transfer rates of the    $d$ and  $p$-states of  \\ion{Ba}{ii}  by  isotropic collisions with neutral hydrogen are calculated. We consider a plane parallel layer of  \\ion{Ba}{ii} situated at the low chromosphere and anisotropically illuminated  from below which produces   linear polarization in the    ${\\lambda}4554$ line   by scattering processes. To compute that polarization, we   solve the statistical equilibrium equations for   \\ion{Ba}{ii}  levels including collisions, radiation and magnetic field effects. } {Variation laws of the relaxation and transfer rates  with  hydrogen number density $n_{\\textrm {\\scriptsize H}}$  and temperature are deduced.    The   polarization  of the ${\\lambda}4554$  line is  clearly affected  due to isotropic collisions with neutral hydrogen  although the collisional depolarization  of its upper level   $^2P_{3/2}$    is negligible. This is because the alignment  of the  metastable  levels  $^2D_{3/2}$ and $^2D_{5/2}$ of the \\ion{Ba}{ii} are vulnerable to collisions. At   the height of formation of the ${\\lambda}4554$  line where  $n_{\\textrm {\\scriptsize H}} \\sim 2 \\times  10^{14}$ cm$^{-3}$, we find that the neglecting of the   collisions  induces   inaccuracy of $\\sim$ 25\\% on the calculation of the polarization and   $\\sim$ 35  \\%   inaccuracy on     microturbulent   magnetic field determination. } {The polarization of the  ${\\lambda}4554$  line decreases  due to collisions  with hydrogen atoms.  In addition, during scattering processes collisions  could change the frequency of the \\ion{Ba}{ii} photons.  To quantitatively study this line, one should confront   the problem of development of general theory treating partial redistribution of frequencies and including transfer and relaxation  rates     by collisions for a multilevel atom with hyperfine  structure. } \\keywords {Scattering -- Polarization -- Atomic processes --   Sun: chromosphere -- Line: formation} ",
        "introduction": "Depolarization due to the Hanle effect can be  used  to retrieve information on solar magnetic fields  (e.g. Trujillo Bueno    2001). This magnetic depolarization is   usually  mixed with a collisional depolarization. To quantitatively  use the Hanle effect as a technique of investigation of solar magnetic fields,  one needs to know  the collisional rates.  In addition, one    should   include properly both the collisional  depolarizing and  {\\it transfer}  rates in the line formation modeling. For instance,   Derouich et al. (2007)   found that the effect of the collisions, in   typical solar conditions of temperature and hydrogen density,  is particularly important for the very important  Ti {\\sc i}  ${\\lambda}4536$ line--  a collisional depolarization of   more than 25 \\% is obtained. In fact, at a first glance, because the inverse lifetime of the  upper level is larger than  the value of the elastic depolarizing rate  $D^2$, one might think (wrongly) that the effect of the collisions is negligible for the Ti {\\sc i}  ${\\lambda}4536$ line.   Collisional  depolarization  of this line  is mainly due to  {\\it transfer rates}   which are generally neglected (see Eq. 14 of Derouich et al.  2007).  Since their ionization potential is rather low, most Barium atoms are ionized throughout the low chromosphere  (Tandberg-Hanssen \\& Smythe 1970). The core of the Ba {\\sc ii} ${\\lambda}4554$ line is formed at about  800 km above the photosphere (e.g.  Uitenbroek \\& Bruls  1992). The wings are formed in deeper layers of the photosphere. Recently, a renewed interest in the measurement and the physical interpretation  of the linear polarization  of the   Ba {\\sc ii} ${\\lambda}4554$ line have emerged (e.g. Malherbe et al., 2007; Belluzzi et al. 2007; Lopez Ariste et al. 2008).  It has been shown that during the formation   of this line the hyperfine structure and the partial redistribution of the frequencies effects are  very important. In this work,  we   investigate fully the   importance of isotropic collisions with neutral hydrogen in its   modeling (Sects. 2 and 3).   Sect. 4 is dedicated to the calculation of the impact of    neglecting collisional effects on the magnetic field determination. Our concluding remarks are given in Sect. 5. ",
        "conclusions": "It is well known that collisions with neutral hydrogen are   essential   to model    the  processes governing the formation of  many polarized lines in the solar photosphere   (D$_1$ and D$_2$  lines of \\ion{Na}{i}; Ti {\\sc i}  ${\\lambda}4536$;   \\ion{Sr}{i}  ${\\lambda}4607$; etc). We show in this work that, although  the neutral hydrogen is  about ten times less abundant  in the low chromosphere  than in the photosphere, the polarization  of the chromospheric  ${\\lambda}4554$ line  decreases significantly  due to collisions with hydrogen atoms.  Conclusions about the role of the collisions should not be  only based on a simple comparison of the inverse lifetime of the  upper level of the transition  and  the $D^2$  coefficients. Although this can give useful indications,   realistic conclusion   should be inferred only from a full introduction of the depolarization and   collisional transfer         rates  in the SEE for multi-level models.   Lines like    \\ion{Ba}{ii}  ${\\lambda}4554$ having  upper $np$-level   energy larger than the energy of the ($n$-1)$d$-level cannot be treated with a simplified model approximation ($n$ is the principal quantum number). For instance,  to model quantitatively  the K \\ion{Sr}{ii}  ${\\lambda}4078$ line  one should take into account the alignment of the $d$-sates and the   infrared triplet \\ion{Sr}{ii} ${\\lambda}10036$, \\ion{Sr}{ii} ${\\lambda}10327$, and \\ion{Sr}{ii} ${\\lambda}10914$. We notice that the K \\ion{Sr}{ii}  ${\\lambda}4078$ line  was examined by Bianda et al. (1998),  but  they used the traditional Van der Waals approach to calculate collisional rates and neglected the collisional depolarization of the long lived $d$-level. Interestingly, although    the  \\ion{Mg}{ii}  is  in the same isoelectronic sequence as  \\ion{Ca}{ii}, \\ion{Sr}{ii}, and \\ion{Ba}{ii}, the energy of the upper $p$-level of the K  \\ion{Mg}{ii}  ${\\lambda}2796$ line  is smaller than the energy of the $d$-level. This is fortunate because  the modeling of this line could be safely performed using a simplified atomic model neglecting the  role of the $d$-level.  Furthermore, the  K  \\ion{Mg}{ii}  ${\\lambda}2796$ line is expected to be strongly polarized  since the curve showing the anisotropy factor as a function of ${\\lambda}$ reaches its maximum around ${\\lambda} \\sim 2800$ $\\AA$  (see Fig. 2   of Manso Sainz \\& Landi Degl'Innocenti  2002). In addition, in  given physical conditions, collisional effects  are clearly smaller for \\ion{Mg}{ii}  than for \\ion{Sr}{ii} and \\ion{Ba}{ii}. Observations of this line are  difficult  from the ground-based telescopes but it could be observed with the help of high sensitivity modern polarimeters attached to space missions.   To point out trends in the depth dependence of the magnetic field, Derouich et al. (2006) interpreted {\\it spatially-resolved} observations of the  photospheric \\ion{Sr}{i}  ${\\lambda}4607$ line. In order to extend their diagnostic    over larger parts of the solar atmosphere,   quantitative interpretation  of  the \\ion{Ba}{ii}  ${\\lambda}4554$ line,  also observed with spatial resolution, would be highly interesting.  To quantitatively study this line, one has to account for  partial frequency redistribution   effects (e.g. Uitenbroek  \\& Bruls 1992 and Rutten  \\& Milkey 1979). Here also collisions should play an important role since, besides their   effects on the atomic polarization, they  could change the frequency of the \\ion{Ba}{ii} photons.  It remains a challenge to develop a general theory for partial frequency redistribution of polarized radiation in the presence of arbitrary magnetic fields and  including  the effects  of   collisions in a multilevel picture  with/without hyperfine structure."
    },
    "0801/0801.0630_arXiv.txt": {
        "abstract": "\\PRE{\\vspace*{0.3in}}  The origin of both the diffuse high-latitude MeV gamma-ray emission and the 511 keV line flux from the Galactic bulge are uncertain. Previous studies have invoked dark matter physics to independently explain these observations, though as yet none has been able to explain both of these emissions within the well-motivated framework of Weakly-Interacting Massive Particles (WIMPs). Here we use an unstable WIMP dark matter model to show that it is in fact possible to simultaneously reconcile both of these observations, and in the process show a remarkable coincidence: decaying dark matter with MeV mass splittings can explain both observations if positrons and photons are produced with similar branching fractions. We illustrate this idea with an unstable branon, which is a standard WIMP dark matter candidate appearing in brane world models with large extra dimensions. We show that because branons decay via three-body final states, they are additionally unconstrained by searches for Galactic MeV gamma-ray lines. As a result, such unstable long-lifetime dark matter particles provide novel and distinct signatures that can be tested by future observations of MeV gamma-rays. ",
        "introduction": "The existence of dark matter is well-established, yet its identity remains elusive. Standard dark matter candidates include Weakly-Interacting Massive Particles (WIMPs), which have mass $\\sim 0.1-1~\\tev$, and arise from independent attempts in particle physics to understand the mechanism of electroweak symmetry breaking. Well-studied signatures of WIMPs include elastic scattering off nucleons in underground laboratories, missing energy signals at colliders, and particle production via self-annihilation in Galactic and extragalactic sources (e.g.~\\cite{Jungman:1995df,Baer:1997ai,Bergstrom:2000pn,Cheng:2002ej,Bertone:2004pz}).  These standard signatures, however, do not provide the only means to test the properties of dark matter. As an additional example, one may also consider invoking dark matter physics to explain anomalous signatures in observed photon or cosmic ray emission spectra, even at energies seemingly far removed from those associated with weak-scale physics. In fact, in recent years, dark matter models have been constructed to explain astrophysical particle production over a variety of energy scales:  these include ultra-high energy cosmic rays~\\cite{Berezinsky:1997hy}, 511 keV line emission from the Galactic bulge ~\\cite{Boehm:2003bt,Picciotto:2004rp,Hooper:2004qf,Ferrer:2005xva,Kasuya:2006kj,Finkbeiner:2007kk,Pospelov:2007xh}, or the diffuse MeV gamma-ray background~\\cite{Kribs:1996ac,Ahn:2005ck}. Though of course all of these anomalies will not be due to exotic dark matter physics, the observed emissions can provide strong constraints on well-motivated dark matter models, and further they may help identify the interesting regions of parameter space for a given dark matter model.  In this paper we focus on two of the aforementioned anomalies: the diffuse MeV gamma-ray background and the 511 keV line flux from the Galactic bulge. As we discuss below, known ``astrophysical\" sources have both spectral shapes and rates that are unable to account for these observed emissions. The lack of well-motivated sources, as well as the similarities of energy scales, has given rise to speculation that both of these anomalies can be explained within the context of a single dark matter model~\\cite{Lawson:2007kp}, albiet at the cost of abandoning the well-motivated WIMP framework. Remaining within the confines of WIMP models, these emissions have been separately reconciled by invoking unstable WIMPs with nearly degenerate, $\\sim$ MeV scale mass splittings: the 511 keV line flux has been studied in Refs.~\\cite{Finkbeiner:2007kk,Pospelov:2007xh}, and the diffuse MeV photon background has been studied in Refs.~\\cite{Cembranos:2007fj,Cembranos:2006gt}. The goal of this paper is to show that, remarkably, {\\em both} the 511 keV line emission and the diffuse MeV gamma-ray spectrum can be explained in the context of a decaying WIMP model, characterized by a lifetime of $10^{20}$ s and near equal branching fraction to photons and electrons.  Generically, we focus on a scenario in which the WIMP mass spectrum is highly degenerate, characterized by $\\sim~\\mev$ mass splittings between the next-to-lightest particle (NLP) and lightest particle (LP).  Both WIMPs freeze out under the standard conditions in the early Universe, and one of them is unstable with a lifetime in excess of the Hubble time. The NLP decays to three-body final states such as NLP $\\rightarrow$ LP $+\\gamma +\\gamma$, NLP $\\rightarrow$ LP $+ e^+ + e^-$, NLP $\\rightarrow$ LP $+\\bar{\\nu} +\\nu$. Two-body decays are assumed to be either highly suppressed or forbidden. Phenomenological consequences of two-body decays with similar lifetimes, and implications for the diffuse MeV photon background, were introduced in Ref.~\\cite{Cembranos:2007fj}, and further discussed in Ref.~\\cite{Yuksel:2007dr}. Additionally, prospects for detecting long-lifetime decaying dark matter with the GeV gamma-ray background were considered in Ref.~\\cite{Bertone:2007aw}.  As a specific implementation of the above idea, we consider the brane world scenario (BWS), which has become one of the most popular extensions to the Standard Model (SM). In the BWS, particles are confined to live on a three-dimensional brane embedded in a higher dimensional ($D=4+N$) space-time, while the gravitational interaction has access to the entire bulk space. The fundamental scale of gravity in $D$ dimensions, $M_D$, can be lower than the Planck scale, $M_P$. In the original proposal~\\cite{ArkaniHamed:1998rs,Antoniadis:1998ig}, the main aim was to address the hierarchy problem, and for that reason the value of $M_D$ was taken to be around the electroweak scale. However, brane cosmology models have also been proposed in which $M_D$ is much larger than the TeV scale~\\cite{Langlois:2002bb,Multamaki:2003zc}. In this paper, we consider a general BWS with arbitrary fundamental scale $M_D$; since we neglect gravitational effects, our results do not depend on this scale once sufficiently high.  In general, the existence of extra dimensions is responsible for the appearence of new fields on the brane. On one hand, we have the tower of Kaluza-Klein (KK) modes of fields propagating in the bulk space, i.e. the gravitons. On the other, since the brane has a finite tension, $f^4$, its fluctuations will be parametrized by $\\pi^\\alpha$ fields called branons. When traslational invariance in the bulk space is an exact symmetry, these fields can be understood as the massless Goldstone bosons arising from the spontaneous breaking of that symmetry induced by the presence of the brane~\\cite{Sundrum:1998ns,Dobado:2000gr}. However, in the most general case, translational invariance will be explicitly broken and therefore we expect branons to be massive fields. When branons are properly taken into account, the coupling of the SM particles to any bulk field is exponentially suppressed by a factor $\\exp[-M^2_{KK}M_D^2/(8 \\pi^2f^4)]$, where $M_{KK}$ is the mass of the corresponding KK mode~ \\cite{Bando:1999di,Cembranos:2005jc}. As a consequence, if $f\\ll M_D$, the KK modes decouple from the SM particles. Therefore, for flexible enough branes, the only relevant degrees of freedom at low energies in the BWS are the SM particles and branons.  The potential signatures of branons at colliders have been considered in Refs.~\\cite{Creminelli:2000gh,Alcaraz:2002iu,Cembranos:2003aw}, and astrophysical and cosmological implications have been studied in Refs.~\\cite{Kugo:1999mf,Cembranos:2003aw}. Moreover in Ref.~\\cite{Cembranos:2003mr} the possibility that massive branons could account for the observed dark matter of the Universe was studied in detail (see also~\\cite{Cembranos:2003aw}). Here, for the first time, we consider branon phenomenology with MeV mass splittings and lifetimes in excess of the age of the Universe. Due to their universal coupling to the SM, decays of unstable branons produce electrons and photons at the same rate in decays, leading to the aforementioned consequences for diffuse MeV gamma-rays and the 511 keV line flux.  This paper is organized as follows. In section~\\ref{sec:gammarayobservations}, we review the observations of the cosmic and Galactic gamma-ray backgrounds and the 511 keV signal. Section~\\ref{sec:branons} we summarize the salient features of the brane world model, and in section~\\ref{sec:signals} we present the resulting astrophysical signatures. Finally, in sections~\\ref{sec:discussion} and~\\ref{sec:conclusions} we discuss other possible signatures and recap our main conclusions. ",
        "conclusions": "\\label{sec:conclusions}  In this paper, we have analyzed the general possibility that decaying dark matter may be the source of both the 511 keV gamma-ray line from the Galactic center and the diffuse high latitude photon background at energies between 1 and 5 MeV. We have introduced a new scenario where a dark matter particle is nearly degenerate with a lighter daughter, with mass splitting $\\sim$ MeV and lifetime $\\sim 10^3$ times the present age of the Universe. The decays of the dark matter are into three body final states: a Standard Model particle anti-particle pair and a daughter dark matter particle. Remarkably, we find that the above mass splitting and lifetime are exactly that which is required to explain both of the above gamma-ray emissions, provided there is similar branching ratios to electrons and photons.  We have illustrated this idea with a concrete model of brane-world dark matter, and we have shown that a standard WIMP motivated within the brane-world scenario, the branon, can explain both observations with natural values required to obtain the correct dark matter relic abundance. Using the effective low-energy Lagrangian for a model with two branons interacting with the Standard Model fields, we have computed the decay rates of the heaviest branon into the lighter one and pairs of Standard Model particles, in particular focusing on electron-positron pairs and photons. We find that the branon model not only produces positrons and gamma-rays at the required rates, it also provides unique and novel phenomenology that will be tested by future MeV gamma-ray observatories."
    },
    "0801/0801.1770_arXiv.txt": {
        "abstract": "Spectroscopic observations of the 2006 outburst of RS Oph at both infrared (IR) and X-ray wavelengths have shown that the blast wave has decelerated at a higher rate than predicted by the standard test-particle adiabatic shock-wave model. The observed blast-wave evolution can be explained, however, by the diffusive shock acceleration of particles at the forward shock and the subsequent escape of the highest energy ions from the acceleration region. Nonlinear particle acceleration can also account for the difference of shock velocities deduced from the IR and X-ray data. We discuss the evolution of the nova remnant in the light of efficient particle acceleration at the blast wave. ",
        "introduction": "X-ray observations of the latest outburst of RS Oph \\citep{sok06,bod06} have allowed to clearly identify the forward shock wave expanding into the red giant (RG) wind and to estimate the time evolution of its velocity, $v_s$, through the well-known relation for a test-particle strong shock: \\begin{equation} v_s = \\bigg({16 \\over 3} {k T_s \\over \\mu m_H}\\bigg)^{0.5}~, \\end{equation} where $k$ is the Boltzmann constant, $T_s$ is the measured postshock temperature and $\\mu m_H$ is the mean particle mass. The X-ray emission has revealed that after an ejecta-dominated, free expansion stage lasting $\\sim$6~days, the remnant rapidly evolved to display behavior characteristic of a shock experiencing significant radiative cooling. The lack or the very short duration of an adiabatic, Sedov-Taylor phase differs from the remnant evolution model developed after the 1985 outburst \\citep[see][and references therein]{obr92}.  The time-dependence of shock velocity has also been measured by IR spectroscopy \\citep{das06,eva07}. Although the general behavior of the shock evolution was found to be consistent with that deduced from the X-ray emission, the shock velocities determined from the IR data are significantly greater than those obtained using equation~(1) together with the X-ray measurements of $T_s$ (see  Fig.~1a).  We have shown that production of nonthermal particles by diffusive shock acceleration at the blast wave can account for these observations \\citep{tat07}. We have proposed that the difference of shock velocities deduced from the IR and X-ray data is due to the use of equation~(1), which is known to underestimate $v_s$ when particle acceleration is efficient \\citep{dec00,ell07}. The observed early cooling of the blast wave is explained by the escape of high-energy particles from the acceleration region, which we found to be much more effective to cool the shock than radiative losses. Thus, nonlinear diffusive shock acceleration has important implications for the evolution of the nova remnant. ",
        "conclusions": ""
    },
    "0801/0801.1293_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:1}  In ESO period 65 (April-September 2000) the large programme 165.N-0276, led by Roger Cayrel, began making use of UVES  at the Kueyen VLT telescope. Known within the Team and outside as ``First Stars'', it was aimed at obtaining high resolution, high signal-to-noise ratio spectra in the range 320 nm -- 1000 nm for a large sample of extremely metal-poor (EMP) stars identified from the HK objective prism survey \\cite{beers85,beers92}. The goal was to use these spectra to determine accurate atmospheric parameters and chemical composition of these stars which are among the oldest objects amenable to our detailed study. Although these stars are not the first generation of stars they must be very close descendants of the first generation. One may hope to gain insight on the nature of the progenitors from detailed information on the descendants.  The extremely metal-poor stars are very rare objects and finding them in large numbers requires specially designed surveys. All of the proponents of the large programme had been actively working on the medium-resolution follow-up of the HK survey (results still to be published), from either ESO La Silla, Kitt Peak or Roque de los Muchachos. Such a follow-up is mandatory in order to obtain a good list of candidates on which one can invest the time of an 8 m telescope.  The programme was allocated a total of 39 nights between periods 65 and 68, these were split into 8 observational runs of unequal length. The observations were carried out in visitor mode because UVES was used in non-standard settings. The settings selected were Dic1 396+573 and Dic1 396+850, typically with a $1''$ slit for a resolution $R\\sim 43000$. These settings were preferred over the standard Dic1 390+580 and Dic1 390+860, because you gain the Ba\\,{\\sc ii} 455.4 nm line in the blue, Zn\\,{\\sc i} 471 nm in the red and Li\\,{\\sc i} 670.8 nm in the reddest setting.  The main results of the large programme are published on a series of papers ``First Stars'' on A\\&A, so far 10 papers have been published, one is in press, a few more in preparation. In addition a number of papers not in the ``series'' have been published, up to know there are 19 refereed papers published, which make use of the data acquired in the course of this large programme.  In this contribution we highlight the main results of the large programme. ",
        "conclusions": ""
    },
    "0801/0801.4649_arXiv.txt": {
        "abstract": "Almost all properties of a photodissociation region (PDR) depend on its metallicity. The heating and cooling efficiencies that determine the temperature of the gas and dust, the dust composition,  as well as the elemental abundances that influence the chemical structure of the PDR are just three examples that demonstrate the importance of metallicity effects in PDRs. PDRs are often associated with sites of star formation. If we want to understand the star formation history of our own Galaxy and of distant low-metallicity objects we need to understanding how metallicity acts on PDR physics and chemistry. ",
        "introduction": "\\subsection{PDRs and Star Formation} It is common to define PDRs as regions where stellar far ultraviolett (FUV: 6~eV$ \\le h\\nu \\le$ 13.6~eV) radiation dominates the physical and chemical properties of the local interstellar medium (ISM). FUV radiation is predominantely produced in massive stars, which do not live long enough to exit their parental cloud of gas and dust within their lifetime. PDRs are the products of the strong mutual feedback between young massive stars and their parental clouds and hence closely related to the process of star formation (SF). Metallicity ($Z$) is thought to affect the fragmentation properties of gravitationally unstable gas, especially in the early universe (Bromm {\\em et al.\\/} \\cite{bromm01}; Schneider {\\em et al.\\/} \\cite{schneider02}). This should also affect the initial mass function (IMF). However, Jappsen {\\em et al.\\/} \\cite{jappsen07} have shown, that the cooling of hot, very low-$Z$ gas does not seem not to depend on metallicity, since it is dominated by H$_2$ cooling. Hence, the cooling depends on the amount of H$_2$ available for cooling, which might still be a function of metallicity due to the fact that molecular hydrogen is formed on metallic surfaces, i.e. dust. \\subsection{Metallicity} \\begin{figure} \\centering \\includegraphics[width=0.8\\columnwidth]{roellig_fig01.eps} \\caption{Overview of metallicity influence on different PDR aspects.}\\label{influence} \\end{figure} Figure~\\ref{influence} lists the most prominent metallicity effects. An altered metallicity affects the physics (primarily energy balance and radiative transfer) and the chemistry of PDRs. The most important, direct effects are altered elemental abundances leading to a different chemical structure and different total column densities of important species like C$^+$, C, and CO as well as an altered dust-to-gas ratio $D/G$. A lower $D/G$ leads to a weaker attenuation and FUV radiation penetrates deeper into the cloud, destroying fragile molecule up to larger depths. Additionally,  gas heating by FUV photons can act deeper in the cloud leading to a different temperature structure compared to solar metallicity PDRs. On the other hand, the probability for a photon to escape the cloud is increased, improving the cooling efficiency at any particular depth while the smaller elemental abundance reduces the cooling rate. Observations of low metallicity environments also confirm that the dust composition changes at lower $Z$. The PAH (polycyclic aromatic hydrocarbons) fraction drops significantly (Draine {\\em et al.\\/} \\cite{draine07}, Madden {\\em et al.\\/} \\cite{madden06}). PAHs dominantly contribute to the total photoelectric heating (PE) efficiency, thus, at lower values of $Z$ the PE efficiency is significantly diminished. Madden {\\em et al.\\/} \\cite{madden06} also showed, that the spectral shape of the FUV radiation is different in low $Z$ dwarf galaxies. Reduced metallicities lead,  for a given stellar type, to hotter main sequence temperatures and to harder, i.e. more energetic, spectral SEDs. This may affect the dust composition via evaporation and photodissociation processes and leads to a different PE efficiency. Since all these effects are non-linear and mutually coupled it is very difficult to give an analytical solution and one has to rely on numerical computations. Sometimes the results are completely counter-intuitive like displayed in Figure~\\ref{invertedT}. The heating and cooling efficiencies at the cloud surface depend on parameters like density $n$, metallicity and FUV intensity $\\chi$. With increasing FUV intensities and densities above $\\sim10^4$~cm$^{-3}$, the dominant heating contribution shifts from H$_2$-deexcitation to  photoelectric (PE) heating (c.f. R\\\"ollig  {\\em et al.\\/} \\cite{roellig06}). Since $\\Gamma_\\mathrm{PE}\\sim Z^2$ is much less effective for $Z<1$ compared to the cooling terms, which scale linearly with $Z$, the gas temperature drops if $\\chi$ increases, i.e. if we move closer to the FUV source! Vice versa this means, a hypothetical cloud gets hotter the further away it is placed from a strong FUV source. Even though this might be an academic example due to its particular assumptions, it demonstrates the difficulty in predicting metallicity effects. \\begin{figure} \\centering \\includegraphics[width=0.6\\columnwidth]{roellig_fig02.eps} \\caption{Surface gas temperature of a high density PDR vs. metallicity for different FUV intensities. At low $Z$ an inversion occurs, i.e. the surface temperature drops for an intenser FUV illumination.}\\label{invertedT} \\end{figure} \\subsection{The KOSMA-$\\tau$ PDR Code} The model results presented here are computed with the KOSMA-$\\tau$ PDR code. (Stoerzer  {\\em et al.\\/} \\cite{stoerzer96}, R\\\"ollig  {\\em et al.\\/} \\cite{roellig06}). This code features spherical geometry and an isotropical FUV irradiation. It calculates the chemical and physical cloud structure as function of radius as well as the total cloud emission. The metallicity is a free model parameter and changes the elemental abundances and the gas-to-dust ratio. The chemical network in KOSMA-$\\tau$ is fully modular and individual species can be easily added to or removed from the network. Recently, the code has been expanded to calculate the dust temperatures and continuum emission for any given dust distribution (R\\\"ollig \\& Szczerba \\cite{roellig08}).  The model setup can be easily expanded to clump ensembles as presented by Cubick {\\em et al.\\/} (this issue) to calculate the PDR emission of a large, complex superposition of differently sized clouds. R\\\"ollig  {\\em et al.\\/} \\cite{roellig07} gives an overview over PDR modelling as well as a comparison between a number of different model codes. The KOSMA-$\\tau$ results are also available online ({\\small \\tt http://www.astro.uni-koeln.de/workgroups/theo\\_astronomy/pdr/}). ",
        "conclusions": ""
    },
    "0801/0801.1546_arXiv.txt": {
        "abstract": "\\noindent There is an ongoing debate in the literature as to whether the effects of averaging out inhomogeneities (``backreaction'') in Cosmology can be large enough to account for the acceleration of the scale factor in the FLRW models. In particular, some simple models of structure formation studied in the literature seem to indicate that this is indeed possible, and it has also been suggested that the perturbed FLRW framework is no longer a good approximation during structure formation, when the density contrast becomes nonlinear. In this work we attempt to clarify the situation to some extent, using a fully relativistic model of pressureless spherical collapse. We find that whereas averaging during structure formation can lead to acceleration via a selective choice of averaging  domains, the acceleration is not present when more generic domains are used for averaging. Further, we show that for most of the duration of the collapse, matter velocities remain small, and the perturbed FLRW form of the metric can be explicitly recovered, in the structure formation phase. We also discuss the fact that the magnitude of the average effects of inhomogeneities depends on the scale of averaging, and while it may not be completely negligible on intermediate scales, it is expected to remain small when averaging on suitably large scales. ",
        "introduction": "Understanding the origin of Dark Energy is undoubtedly one of the most important problems in cosmology today. Recently, several workers in the field (including the present authors) \\cite{ras,avg,kolb-notari,wilt}, have supported the hypothesis that the acceleration of the scale factor in the standard Friedmann-Lema\\^itre-Robertson-Walker (FLRW) cosmologies, might be attributed to the effects of averaging over inhomogeneities in the real Universe. It is known that explicit averaging of the Einstein equations leads to non-trivial corrections \\cite{ellis,buch,zala}. These corrections or ``backreaction'', it is argued, can mimic a Dark Energy term in the Cosmological equations. Specifically, some simple models have been presented which demonstrate how such an effect may arise \\cite{ras}.  On the other hand, it has also been argued that as long as the standard perturbed FLRW picture for the metric of the Universe is a good approximation, any effects of averaging over the inhomogeneous perturbations must remain small (see, e.g. \\Cites{wald,wett,behr,flan}). The counter given to this argument is that the perturbed FLRW framework is  expected to break down around the time of structure formation \\cite{ras}. That $N$-body simulations do not demonstrate any such breakdown is attributed to the fact that these simulations always work in the Newtonian limit of General Relativity (GR), usually employing periodic boundary conditions, and that the backreaction due to averaging under such conditions can be shown to vanish \\cite{buch2}. Further, since the evolution of the background scale factor is fixed in the simulations, backreaction effects would not be visible even if present.  This situation clearly demands some clarification. One expects that as long as gravitational fields are weak, the perturbed FLRW framework should work well. Is the weak field approximation then actually breaking down during structure formation? If so, then clearly a radical rethink of the approach to Cosmology is in order. If not, on the other hand, then one needs to understand how simple models such as the one presented by R\\\"as\\\"anen \\cite{ras} can achieve an acceleration upon averaging during structure formation.  In this paper we will attempt to address these issues. We employ a simple but fully relativistic model of spherical collapse, \\emph{with appropriate initial conditions imposed}. We will show that the effect observed by R\\\"as\\\"anen can actually be attributed to a very selective choice of averaging domain, and that in fact the effect is not present when a more generic averaging domain is chosen. However, we do find that the effect of averaging inhomogeneities may not be completely negligible. In particular we see $\\sim10\\%$ deviations of the effective deceleration parameter to be defined below, from the expected FLRW value, on the scale over which we define our averaged quantities. This is in line with the findings of Li and Schwarz \\cite{schwarz,schwarz2} in the context of averaging of perturbative inhomogeneities. To support these results, we will show that at any stage of the collapse, \\emph{if matter velocities are small}, then the weak field conditions hold and one can explicitly recover the perturbed FLRW form of the metric. We will see that matter velocities do, in fact, remain small for most of the duration of the collapse. We will also present an argument explaining the origin of the $\\sim10\\%$ effect from perturbative metric inhomogeneities on the scales at hand, and argue that in the real universe, the averaging scale is expected to be large enough for these effects to be much smaller.  We have organised the paper as follows. In section 2 we will set up our model for the spherical collapse, and compare its results with those obtained by R\\\"as\\\"anen. In section 3 we will show how the perturbed FLRW metric is recovered, and we will conclude in section 4 with a discussion of further tests of the backreaction argument. ",
        "conclusions": "In this paper we have addressed the question of whether or not the assumptions involved in using a perturbed FLRW framework to describe the \\emph{metric} of the present Universe, break down during epochs of structure formation, and whether averaging over inhomogeneities in such a context can yield effects large enough to lead to an accelerating scale factor. We have demonstrated two things. First, it is very essential to impose initial conditions properly matched across any boundaries present in the model. Neglecting to do so (for example by ignoring certain regions such as Region 2 of our model) can lead to effects such as an accelerating effective scale factor $a_{mod}$. Also, as we saw in \\fig{fig2d}, it \\emph{is} in principle possible to have very large deviations from FLRW-like conditions after averaging, but this is at the cost of initial conditions which are unrealistic when confronted with the CMB data.  Secondly, we have shown that, provided peculiar velocities of matter are small (which is a reasonable assumption on observational grounds), one can explicitly find coordinates in which the metric of the Universe is of the perturbed FLRW form \\eqref{3eq1}. This by itself is not a surprising result, since one can always do this locally (our Region 1 for example is a closed FLRW sub-Universe by itself, without any coordinate transformations). What is important is the fact that the background FLRW used for comparison is taken to be uniquely and \\emph{globally} defined, so that \\emph{every} local region involved in structure formation can be compared to this background in an unambiguous fashion. The existence of this global FLRW model is not simply an assumption made at the present epoch, it is justified by the fact that density fluctuations in the \\emph{past} were very small.  An issue worth highlighting again is that the time coordinates involved in the transformation from the synchronous coordinates to the perturbed FLRW form, are linked by an infinitesimal transformation (after controlling some degrees of freedom by setting arbitrary functions of time to zero), as expected with weak gravitational fields. While our model is perhaps not the most realistic depiction of observed voids, we expect this feature to hold in more realistic models also. At first glance this may appear to differ from Wiltshire's claim \\cite{wilt} that the clocks of observers in bound regions face a significant calibration when compared with those inside large voids. However the issue is rather subtle, since Wiltshire's arguments involve a different choice of time slicing than the standard constant $t$ hypersurfaces used here, and his model shows cumulative effects over large times. It is possible that these cumulative effects are related to the relative difference between the metric functions which we discussed in Section 3, but a straightforward comparison is not possible at this stage. In this context, it becomes essential to understand the origins of those differences, and moreover the fact that the magnitude of the difference is sensitive to boundary conditions.  A few remarks concerning further tests of the backreaction argument, with more realistic models of structure formation which should account for pressure. We have not modelled voids \\emph{surrounded} by overdensities, since such models are plagued by shell-crossings in the absence of pressure. (Our current model also faces this difficulty beyond region 3.) But qualitatively, we expect our results to hold even in models where a suitable pressure term takes care of shell crossings and leads to stable structures surrounding expanding voids, since we expect peculiar velocities to remain small in this situation as well. As long as the sizes of individual voids are not an appreciable fraction of the Hubble scale, one expects that here as well, it will be possible to find coordinate transformations like the one we have presented. We will return to this problem in future work.  Finally, we comment on an issue which deserves careful consideration. Recall that in the Buchert scheme of averaging perturbative inhomogeneities, which has been employed by several authors \\cite{kolb-notari,behr,schwarz}, there are \\emph{two} scale factors -- the background scale factor $a(t)$ which satisfies the usual FLRW equations, and the effective scale factor $a_D(t)$ which satisfies the Buchert equations \\cite{buch}. It is our belief that this situation has an inherent ambiguity -- which of the two scale factors is relevant for observations? In our opinion, a more consistent way of proceeding would be to employ the fully covariant averaging scheme due to Zalaletdinov \\cite{zala}, developed further in the cosmological context by \\cite{pasi}. In this scheme, there would be only \\emph{one} scale factor, unambiguously associated with a background metric, which  would satisfy the \\emph{corrected} Einstein equations. The fact that backreaction is negligible in the linear regime should allow the problem to be well-posed at say the last scattering epoch. While we have estimated the effects of averaging inhomogeneities to be negligible on large enough scales in Buchert's scheme, this was done in a perturbation theory around the \\emph{fixed} Einstein-deSitter background. It is not clear that this result will continue to hold even at late times when the background itself is ``evolving'', namely, being corrected by the (small but cumulative) effects of inhomogeneities. Note that \\emph{if} these effects build up at late times, then in this scheme with only one scale factor, the evolution of the perturbations will also be modified, which may lead to some interesting effects. We will investigate these issues in the near future.   In conclusion, we wish to suggest that it has not yet been conclusively established that cosmological backreaction becomes large enough during the structure formation phase to cause an acceleration of the scale factor. On the contrary, the calculation presented in the present paper suggests that the backreaction remains small during this phase, and that the intuitive picture concerning weak gravitational fields is actually realized in the form of a perturbed FLRW metric. Self-consistently accounting for the effects that the backreaction would implicitly have on the evolution of the perturbations (via the scale factor), may lead to interesting effects which remain to be explored.    {\\bf Acknowledgments :} It is a pleasure to thank Karel van Acoleyen and T. Padmanabhan for several useful discussions. We are also grateful to Thomas Buchert, Syksy R\\\"as\\\"anen, Dominik Schwarz and David Wiltshire for their critical comments on an earlier draft."
    },
    "0801/0801.4539_arXiv.txt": {
        "abstract": "We report the discovery of a new open cluster (OC) in the Galaxy at $\\ell=167.0^\\circ$ and $b=-1.0^\\circ$. Its field includes the planetary nebula (PN) PK\\,167-0.1. We study the possible associations of the PN/OC pairs NGC\\,2818/NGC\\,2818A, NGC\\,2438/M\\,46 (NGC\\,2437), PK\\,6+2.5/NGC\\,6469, as well as of the PN PK\\,167-0.1 with New\\,Cluster\\,1. The analyses are based on near-infrared colour-magnitude diagrams (CMDs) and stellar radial density profiles (RDPs). NGC\\,6469 is located in a heavily contaminated bulge field.  The CMD morphology, especially for the latter two cases, is defined with a field star decontamination algorithm applied to the 2MASS \\jj, \\hh, and \\ks\\ photometry. Field decontamination for the OCs NGC\\,2818A and M\\,46 produced better defined CMDs and more accurate cluster parameters than in the literature. Those pieces of evidence point to M\\,46 as physically associated with the PN NGC\\,2438. The same occurs for the OC NGC\\,2818A and the PN NGC\\,2818, however previous radial velocity arguments indicate that they are not associated. The OC NGC\\,6469 does not appear to be associated with the PN PK\\,6+2.5, which probably belongs to the bulge. Finally, the distance of the OC New\\,Cluster\\,1 is consistent with a physical association with the PN PK\\,167-0.1. ",
        "introduction": "\\label{Intro}  Planetary nebulae (PNe) are late stages in stellar evolution, occurring for stars with mass in the range 1 --- 6.5\\,\\ms (\\citealt{Weidemann00}), possibly with the upper limit at $\\sim8\\,\\ms$. Most of the difficulty in understanding such stages that lead to PNe is associated with the determination of the PN distance and its progenitor mass. {\\bf Probably, the largest source of error is the distance which, in most cases, is known with a $\\sim50\\%$ uncertainty (e.g. \\citealt{Zhang95}), or more. However, both problems can be minimised if the PN is physically associated with a star cluster. Star cluster distances, in general, can be determined with a high precision and, in the case of physical association, the PN progenitor mass can be assumed to be $\\approx10\\%$ larger than that of the turnoff (TO), stars referred to by \\citet{Meynet93} as red-turnoff objects. Exceptions are the cases where mass transfer from a close binary is involved. In fact, recent radial velocity measurements, although restricted to relatively small samples, suggest than more than 90\\% of the central stars in PNe might have companions (e.g. \\citealt{deM06}).}  Since the early observation of the PN NGC\\,2438, which is surrounded by stars of the open cluster (OC) M\\,46 (NGC\\,2437), photometric and spectroscopic methods have been developed to test physical associations between such objects. \\citet{Ziz75} published a list of 10 OCs with nearby projected PNe. \\citet{MTL07} provided a list of 13 possible physical associations and an additional list of 17 angular coincidences ($\\Delta\\,R<15\\arcmin$). They also present physical criteria for membership based on similarities in velocity, reddening, and the ratio of estimated distances.  \\citet{MTL07} reviewed theoretical and observational aspects of stellar evolution to impose limiting parameters for PNe to be OC members. Many coincidences can be discarded if the PN is a bulge member (\\citealt{Ziz75}), as indicated by position along the disk and radial velocities. OCs younger than 28\\,Myr produce type II SN. The short-lived PN phase varies from $10^3$ to $10^5$ years, according to the progenitor mass (\\citealt{KA2000}, \\citealt{SB96}), and does not favour associations. PNe observed in globular clusters (e.g. \\citealt{KA2000}) imply that their turnoffs at $\\la1\\,\\ms$ are below the lower mass limit for PN progenitors. This invokes mass transfer in binary systems to produce PN in globular clusters. \\citet{deM06}, \\citet{Soker2006}, and \\citet{Zijlstra07} presented a scenario where most PNe arise from binaries. Consequently, care is necessary because the cluster age may not correspond to the PN mass progenitor.  Among the list of 30 PN/OC coincidences provided by \\citet{MTL07}, we selected the PN/OC pairs NGC\\,2818/NGC\\,2818A and NGC\\,2438/M\\,46 (NGC\\,2437). We also include PK\\,6+2.5/NGC\\,6469 and PK\\,167-0.1, which is projected in the field of an as-yet unknown star cluster. We name this low-contrast cluster New\\,Cluster\\,1. In the present paper we derive accurate distances to these OCs, which in turn, can be used to investigate the possible associations.  \\begin{figure*} \\begin{minipage}[b]{0.50\\linewidth} \\includegraphics[width=\\textwidth]{fig1a.ps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.50\\linewidth} \\includegraphics[width=\\textwidth]{fig1b.ps} \\end{minipage}\\hfill \\caption[]{Left panel: $10\\arcmin\\times10\\arcmin$ XDSS R image of the OC NGC\\,2818A and the PN NGC\\,2818. Right panel: $20\\arcmin\\times20\\arcmin$ XDSS R image of the OC M\\,46 and the PN NGC\\,2438. Images centred on the 2MASS coordinates (cols.~5 and 6 of Table~\\ref{tab1}). North to the top and East to the left.} \\label{fig1} \\end{figure*}   \\begin{figure*} \\begin{minipage}[b]{0.50\\linewidth} \\includegraphics[width=\\textwidth]{fig2a.ps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.50\\linewidth} \\includegraphics[width=\\textwidth]{fig2b.ps} \\end{minipage}\\hfill \\caption[]{Left panel: $15\\arcmin\\times15\\arcmin$ DSS B image of the OC NGC\\,6469 and the PN PK\\,6+2.5. The nearly stellar PN is located close to the lower-right corner (see Equatorial coordinates in Table~\\ref{tab2}). Right panel: $10\\arcmin\\times10\\arcmin$ XDSS R image of the OC New\\,Cluster\\,1 and the PN PK\\,167-0.1.} \\label{fig2} \\end{figure*}  \\begin{table*} \\caption[]{General data on the clusters} \\label{tab1} \\tiny \\renewcommand{\\tabcolsep}{0.7mm} \\renewcommand{\\arraystretch}{1.5} \\begin{tabular}{lcccccccccccccc} \\hline\\hline &\\multicolumn{2}{c}{WEBDA}&&\\multicolumn{10}{c}{Derived from 2MASS}\\\\ \\cline{2-3}\\cline{5-13} Cluster&$\\alpha(2000)$&$\\delta(2000)$&&\\rx&$\\alpha(2000)$&$\\delta(2000)$&$\\ell$&$b$&Age&$\\ebv$&\\ds&\\dgc&Alternative names\\\\ &(hms)&($^\\circ\\,\\arcmin\\,\\arcsec$)&&(\\arcmin)&(hms)&($^\\circ\\,\\arcmin\\,\\arcsec$)&($^\\circ$)&($^\\circ$)&(Myr)&&(kpc)&(kpc)&&\\\\ (1)&(2)&(3)&&(4)&(5)&(6)&(7)&(8)&(9)&(10)&(11)&(12)&(13)\\\\ \\hline NGC\\,2818A&09:16:01&$-$36:37:30&&80&09:16:07.7&$-$36:38:09.6&262.0&$+8.6$&$1000\\pm100$&$0.10\\pm0.03$&$2.8\\pm0.1$ &$8.1\\pm0.2$ &ESO\\,372$-$13,Hen\\,2$-$23,PK\\,261$+$08\\\\ M\\,46&07:41:46&$-$14:48:36&&80&07:41:41.1&$-$14:51:45.0&231.9&$+4.0$&$250\\pm50$&$0.10\\pm0.03$ &$1.5\\pm0.1$&$8.3\\pm0.2$  &NGC\\,2437,Mel$-$75,Cr\\,159,OCl$-$601\\\\ NGC\\,6469&17:53:13&$-$22:19:11&&45&17:53:03.4&$-$22:18:14.4&6.55&$+1.97$&$250\\pm50$&$0.58\\pm0.06$ &$1.1\\pm0.1$&$6.1\\pm0.2$   &Mel$-$182,Cr\\,353,OCl$-$21,ESO\\,589\\,SC\\,18 \\\\ New\\,Cluster\\,1&---&---&&60&05:06:20&$+$39:09:50.0&167.0&$-1.0$&$1000\\pm100$&$0.29\\pm0.03$ &$1.7\\pm0.1$&$8.9\\pm0.2$  & \\\\ \\hline \\end{tabular} \\begin{list}{Table Notes.} \\item Cols.~2 and 3: coordinates from WEBDA. Col.~4: extraction radius. Col.~10: reddening in the cluster's central region (Sect.~\\ref{age}). Col.~12: \\dgc\\ for $\\rs=7.2$\\,kpc (\\citealt{GCProp}). \\end{list} \\end{table*}  Besides the investigation on possible PN/OC associations, the present work also serves to provide reliable astrophysical parameters of scarcely studied OCs, as well as to characterise a new one. Such data are important both to studies of the disc structure and to constrain theories of molecular cloud fragmentation, star formation, as well as stellar and dynamical evolutions. The present work employs near-IR \\jj, \\hh, and \\ks\\ photometry obtained from the 2MASS\\footnote{The Two Micron All Sky Survey, All Sky data release (\\citealt{2mass1997}), available at {\\em http://www.ipac.caltech.edu/2mass/releases/allsky/}} Point Source Catalogue (PSC). The spatial and photometric uniformity of 2MASS, which allow extraction of large surrounding fields that provide high star-count statistics, have been important to derive cluster parameters and probe the nature of stellar overdensities (e.g. \\citealt{ProbFSR}).  To this purpose we have developed quantitative tools to disentangle cluster and field stars in CMDs, in particular two different kinds of filters. Basically we apply {\\em (i)} field-star decontamination to uncover cluster evolutionary sequences from the field, which is important to derive reddening, age, and distance from the Sun, and {\\em (ii)} colour-magnitude filters, which proved to be essential for building intrinsic stellar radial density profiles (RDPs), as well as luminosity and mass functions. In particular, field-star decontamination constrains more the age and distance, especially for low-latitude OCs (\\citealt{DiskProp}). These tools were applied to OCs and embedded clusters to enhance main-sequence (MS) and/or pre-MS sequences with respect to the field (\\citealt{M52N3960}; \\citealt{N4755}; \\citealt{N6611}; \\citealt{5LowContr}; \\citealt{M11}; \\citealt{DetAnalOCs}). They were useful in the analysis of faint and/or distant OCs (\\citealt{CygOB2}; \\citealt{3OpticalCl}; \\citealt{5LowContr}; \\citealt{FaintOCs}). In addition, more constrained structural parameters, such as core  and limiting radii (\\rc\\ and \\rl, respectively), and mass function (MF) slopes, have been derived from colour-magnitude-filtered photometry (e.g. \\citealt{DetAnalOCs}).  This paper is organised as follows. Sect.~\\ref{Target_OCs} contains basic properties and reviews literature data (when available) on the possible PN/OC associations. In Sect.~\\ref{2mass} we present the 2MASS photometry, build CMDs, discuss the field-star decontamination, and derive cluster fundamental parameters. Sect.~\\ref{struc} describes cluster structure by means of stellar RDPs and mass density profiles (MDPs). In Sect.~\\ref{MF} MFs are built and cluster masses are estimated. In Sect.~\\ref{CWODS} aspects related to the structure of the present objects are discussed. In Sect.~\\ref{Discus} we discuss membership considerations. Concluding remarks are given in Sect.~\\ref{Conclu}. ",
        "conclusions": "\\label{Conclu}  In the present study an open cluster was discovered (New\\,Cluster\\,1), which is located at $\\rm\\alpha(J2000)=05^h06^m20^s$ and $\\delta(J2000)=+39^\\circ09\\arcmin50\\arcsec$. {\\bf Projected on its field, with an angular separation of $4.9\\arcmin$ with respect to the custer centre, is the PN PK\\,167-0.1.} Field-star decontaminated 2MASS photometry was employed in the analysis of 4 pairs of PN/OC candidates to physical association. We derived accurate ages, distances from the Sun, reddening values, and the mass at the turnoff for the clusters. Cluster core and limiting radii were determined by means of radial density profiles fitted with a King-like function.  {\\bf The values of reddening and distance from the Sun derived in this work suggest that the pairs PN/OC NGC\\,2818/NGC\\,2818A and NGC2438/M46 are compatible with physical associations. However, the available radial velocity data for NGC\\,2818/NGC\\,2818A favour a projection effect.} The pair PK\\,6+2.5/NGC\\,6469 appears to be physically unrelated, since the radial velocity of PK\\,6+2.5  suggests a bulge membership. The PN PK\\,167-0.1 is projected close to the OC New\\,Cluster\\,1, and the similar distances are consistent with physical association. In the cases of physical association, the turnoff masses are compatible with the occurrence of PNe. {\\bf It would be important to determine radial velocities for members of the new cluster, to compare them with that of PK\\,167-0.1, in order to further test the physical association possibility. Besides, to establish the non-association of PK\\,6+2.5 with NGC\\,6469, radial veocities of member stars are necessary.}  From the PN/OC angular positions provided by \\citet{MTL07}, we estimate that the fraction of PNe projected against the central region of clusters, with respect to those in the halo, is significantly higher than that expected from spatially unrelated distributions of PNe and OCs. Besides, this fraction agrees with that expected of clusters characterised by a King-like radial stellar distribution. {\\bf In addition, we analysed the distribution of the number of spatial coincidences per cluster area built with the 32 cases dealt with in this paper. This curve was compared to that expected from a spatially-unrelated distribution of PN/OC separations and that resulting from inversion of the sign of the $b$ PN coordinates. The observed distribution presents a conspicuous excess, for PN/OC separations smaller than the cluster limiting radii, over both simulations, which might indicate physical relation. However, we note that these results do not take into account the rather uncertain PN distances, and radial velocities. Besides, they are based on a relatively small number of spatial coincidences. In this context, it is essential to analyse in more detail (i.e. with accurate fundamental and structural parameters) the whole PN/OC sample of \\citet{MTL07} in order to investigate whether the above excess of relatively close spatial coincidences can be interpreted as physical association. }"
    },
    "0801/0801.2015_arXiv.txt": {
        "abstract": "SN 2006jc is a peculiar supernova (SN), in which the formation of dust has been confirmed at an early epoch of $\\sim$50 days after the explosion. We investigate the possibility of such an earlier formation of dust grains in the expanding ejecta of SN 2006jc, applying the Type Ib SN model that is developed to reproduce the observed light curve. We find that the rapid decrease of the gas temperature in SN 2006jc enables the condensation of C grains in the C-rich layer at 40--60 days after the explosion, which is followed by the condensation of silicate and oxide grains until $\\sim$200 days. The average radius of each grain species is confined to be less than 0.01 $\\micron$ due to the low gas density at the condensation time. The calculated total dust mass reaches $\\simeq$1.5 $M_{\\odot}$, of which C dust shares $0.7$ $M_{\\odot}$. On the other hand, based on the calculated dust temperature, we show that the dust species and mass evaluated to reproduce the spectral energy distribution observed by {\\it AKARI} and {\\it MAGNUM} at day 200 are different from those obtained by the dust formation calculations; the dust species contributing to the observed flux are hot C and FeS grains with masses of $5.6 \\times 10^{-4}$ $M_{\\odot}$ and $2.0 \\times 10^{-3}$ $M_{\\odot}$, respectively, though we cannot defy the presence of a large amount of cold dust such as silicate and oxide grains up to $0.5$ $M_{\\odot}$. One of the physical processes responsible for the difference between calculated and evaluated masses of C and FeS grains could be considered to be the destruction of small-sized clusters by energetic photons and electrons prevailing within the ejecta at the earlier epoch. ",
        "introduction": "Supernova (SN) 2006jc is a peculiar Type Ib SN (SN Ib) in which the formation of dust has been observationally confirmed at a very early time of $\\sim$50 days after the peak brightness (Smith et al. 2008), which is more than a few hundred days earlier than those for the dust-forming SNe observed so far. Evidence for dust formation in SN 2006jc comes from an increase of the red to near-infrared (NIR) continuum (Di Carlo et al. 2007) and simultaneous emergence of blueshifted narrow He I emission lines (Smith et al. 2008). Smith et al. (2008) have proposed that the detected dust does not condense in the ejecta but forms in the dense postshock gas swept up by the forward and reverse shocks. Although X-ray observations of {\\it SWIFT} and {\\it Chandra} (Immler et al. 2008) suggested the presence of the dense circumstellar (CS) shell ejected by the outburst similar to those seen in luminous blue variables (LBVs) two years prior to the explosion (Nakano et al. 2006; Pastorello et al. 2007), the postshock CS gas density estimated from the X-ray light curve (Tominaga et al. 2007) is not high enough for dust grains to condense. Thus, it is reasonable to consider that the appearance of dust toward SN 2006jc should be the outcome of ongoing dust formation in the expanding ejecta of the SN.  Several pieces of evidence for dust formation in the SN ejecta have been reported for Type II SNe (SNe II): SN 1987A (Lucy et al. 1989; Whitelock et al. 1989; Meikle et al. 1993; Wooden et al. 1993; Colgan et al. 1994), SN 1998S (Gerardy et al. 2000; Pozzo et al. 2004), SN 1999em (Elmhamdi et al. 2003), and SN 2003gd (Sugerman et al. 2006; Meikle et al. 2007). For these SNe II except for SN 1998S, the onset of dust formation is estimated to be later than 400 days after explosions. Type IIn SN 1998S with the relatively low-mass hydrogen envelope (e.g., Liu et al. 2000; Fransson et al. 2005) exhibits signatures of dust condensation around 230 days (Gerardy 2000). In addition, SN 1990I is the first SN Ib signifying the ongoing dust formation, in which dust formation is observed at $\\sim$ 230 days (Elmhamdi et al. 2004). On the other hand, the signature of dust formation in SN Ic has not been recorded so far. Although theoretical studies have shown feasibility of dust formation in SNe II (Kozasa et al. 1989, 1991; Todini \\& Ferrara 2001; Nozawa et al. 2003), the possibility of dust formation in SNe Ib/c has never been explored to date. Hence, it is an important subject to be pursued how the formation process of dust in the ejecta depends on the type of SNe.  Since the progenitor stars of SNe Ib/c have lost most of the hydrogen/helium envelopes before the explosion, their ejected masses are smaller, and the expansion velocities are significantly higher than SNe II. This leads to a lower density of gas in the ejecta, and the gas temperature drops down more quickly than those in typical SNe II. This allows us to expect earlier dust formation in SNe Ib/c than SNe II. On the other hand, the lower gas density may even result in unsuccessful dust formation. Furthermore, in the early epoch, energetic photons and electrons generated by the Compton degradation of $\\gamma$-rays from the decay of radioactive elements such as $^{56}$Ni and $^{56}$Co prevail throughout the ejecta, and may affect the formation process of dust grains. SN 2006jc is an ideal laboratory for examining the process of dust formation in SNe Ib because a copious amount of observational data enable us to compare with theoretical models. In this paper, we investigate, for the first time, the possibility of dust formation in the expanding ejecta of SN Ib, based on the SN model that can well reproduce the light curve of SN 2006jc (Tominaga et al. 2007).  The paper is organized as follows. In \\S~2 we summarize the observational evidence for dust formation in SN 2006jc and its interpretation, which are compared with the calculations of dust formation in the later sections. In \\S~3 we investigate the dust formation in the ejecta and show that it is possible for dust grains to condense in the ejecta of SN 2006jc at very early times of about 50 days after explosion. In \\S~4 we calculate the temperature of possible condensates obtained in \\S~3 and evaluate the amount of dust grains contributing to the spectral energy distribution observed with {\\it AKARI} and {\\it MAGNUM} at day 200. Formation processes of dust grains in SN 2006jc are discussed in \\S~5, focusing on the interpretation of the difference between dust mass obtained by the dust formation calculation and that evaluated from the spectral fitting. The summary is presented in \\S~6. In this paper, we assume that the explosion of SN 2006jc has occurred at 15 days before its discovery. ",
        "conclusions": "\\subsection{Mass of dust formed in the ejecta}  As shown in \\S~4.2, the mass of C grains (0.7 $M_\\odot$) obtained by the dust formation calculation is three orders of magnitude larger than that ($5.6 \\times 10^{-4}$ $M_\\odot$) estimated from the spectral fitting to the observations. One of the reasons for this difference is considered to be the clumpy structure in the ejecta. In order to evaluate the dust mass necessary for reproducing the observed spectrum of SN 2003gd, Sugerman et al. (2006) performed the radiative transfer calculation taking account of the clumpy structure and found an increase in the estimated mass of dust in the ejecta by about one order of magnitude, compared with the mass derived for the uniform distribution of dust. Nevertheless, the consideration of dust clumping does not seem to be sufficient to overcome the large gap. Thus, the difference in dust mass should stem from the process of dust formation in the ejecta.  One of the processes influencing on dust formation is the non-local thermal equilibrium (non-LTE) effect. As seen from Figures 1 and 3, the dust temperature is not the same as the gas temperature; at 200 days the temperature of C grains (500--600 K) and silicate and oxide grains (100--200 K) is lower than the gas temperature with 700--1000 K. In the ejecta with strong radiation field, the temperature of small clusters is expected to be higher than the gas temperature depending on the optical property. In this case, the non-LTE effect retards or hinders the formation of dust grains. Assuming that the optical properties of small carbon clusters are the same as those of small C grains, and applying a formula of nucleation rate taking into account the non-LTE effect by Kozasa et al (1996), we realized that the non-LTE effect does not significantly affect the formation of C grains; the non-LTE effect only retards the formation of C grains by less than ten days and does not affect the mass of C grains because the temperature of small carbon clusters is only 50 K higher than that of gas. Meanwhile, the non-LTE effect delays the formation of Si grains from $\\sim$200 days to $\\sim$240 days. This is the reason why we exclude Si grains in reproducing the observed SED at day 200 in \\S~4.2.  Another process affecting the formation of dust in the ejecta is the energy deposition on small-sized clusters through the latent heat deposition at the condensation as well as the collisions with high energy photons and electrons. In the ejecta rich in condensible elements, the destruction of small-sized clusters by the deposition of latent heat could be significant unless the deposited energy is released by collisional and/or radiative processes efficiently. Also, at an early epoch after the explosion, energetic photons and electrons generated from the decays of radioactive elements $^{56}$Ni and $^{56}$Co prevail abundantly throughout the ejecta. Since the nucleation rate is determined by the balance between growth and destruction rate of small-sized clusters, the energy deposition makes small-sized clusters unstable against the growth and as a result reduce the nucleation rate.  Although the nucleation rate can be reduced by increasing the critical cluster size as demonstrated by Bianchi \\& Schneider (2007), no reason is there to cut the number of seed nuclei by increasing the critical size arbitrarily, apart from the uncertainty inherent in the standard nucleation theory arising from the evaluation of chemical potential of small-sized clusters. Thus, we examine the effect of energy deposition for dust formation at early times observed in SN 2006jc by reducing the sticking probability $\\alpha_s$ for simplicity and comparing with the mass of dust evaluated in \\S~4.2. If we take $\\alpha_{\\rm s} \\sim 3 \\times 10^{-3}$, we can obtain the mass of C grains of $\\sim$$5 \\times 10^{-4}$ $M_\\odot$, while $\\alpha_{\\rm s}$ is $\\sim$$0.3$ to produce FeS grains of $2 \\times 10^{-3}$ $M_\\odot$. In this case, the condensation times of C and FeS grains is 40--67 days and 190--200 days, respectively. While the destruction rate of small-sized clusters depends on the properties of materials forming dust grains, the difference in the sticking probability may be caused by the collisions with energetic photons and electrons, reflecting the difference in the condensation time of C and FeS grains; energetic photons and electrons are more abundant at day 50 than at day 200. The details of these microscopic processes are beyond the scope of this paper and will be investigated in the future study on the dust formation in SNe.  \\subsection{Early condensation of dust in SNe Ib}  In \\S~3.2, we show that the early formation of dust around 50 days is possible in SN 2006jc whose progenitor has lost most of He envelope (Foley et al. 2007; Tominaga et al. 2007). However, other dust-forming SNe Ib had not show such an early formation of dust; SN 1990I displayed the diagnostics of dust formation around 230 days (Elmhamdi et al. 2004), though the condensation time is relatively earlier than that of SN II. A peculiar Type Ib SN 2005bf shows a rapid decline of the light curve from 50 days after the explosion, but did not exhibit the NIR brightening in contrast to the case of SN 2006jc (Tominaga et al. 2005). Maeda et al. (2007) have shown that the quickly faded light curve of SN 2005bf can be explained by the newly born magnetized neutron star at the center. On the other hand, they have demonstrated that the blueshifted profiles of O I and Ca II observed at $\\sim$270 days can be explained by the obscuration of dust in the ejecta. If the blueshift of emission lines in SN 2005bf is attributed to the newly formed dust grains, the time of dust formation is considered to be later than 100 days and not to be so early as in SN 2006jc, although the epoch of dust formation in SN 2005bf was not specified by the observations.  The results of calculations revealed that the early formation of dust in SN 2006jc can be realized by the rapid decline of gas temperature, reflecting the large ratio of the explosion energy ($E_{51}= $10) to the ejected mass ($M_{\\rm ej}$ = 4.9 $M_\\odot$); the ratio $E_{51}/(M_{\\rm ej}/M_\\odot) \\simeq 2$ is much greater than 0.125--0.25 for SN 2005bf (Tominaga et al. 2005) and $\\simeq$0.3 for SN 1990I (Elmhamdi et al. 2004). Unfortunately, no literature has reported the explosion energy of Type IIn SN 1998S, which had a time of dust condensation of $\\sim$230 days (Gerardy 2000). Using the ejecta mass estimated from the He core of 4--6 $M_\\odot$ (Fassia et al. 2001; Pooley et al. 2002) and hydrogen envelope of 0.1--2 $M_\\odot$ (Fassia et al. 2000; Anupama et al. 2001) and taking $E_{51}=1$ as a typical explosion energy, we can, however, deduce the ratio to be $E_{51}/(M_{\\rm ej}/M_\\odot) =$ 0.125--0.25 for SN 1998S, which is similar to or slightly less than those for SN 2005bf and SN 1990I. On the other hand, the ratio for the dust forming SNe II is of order 0.1: 0.1--0.2 for SN 1987A (Shigeyama et al. 1988; Woosley 1988), 0.05--0.1 for SN 1999em (Elmhamdi et al. 2003), and $\\simeq$0.12 for SN 2003gd (Hendry et al. 2005). Dust formation has not been yet observed on SNe Ic. However, the early dust formation in SNe Ic is also expected because SNe Ic may well have $M_{\\rm ej}$ similar to SNe Ib.  In conclusion, we can speculate that as the envelope of star at explosion is smaller and the explosion energy is higher, the condensation time of dust tends to be earlier. In order to prove these hypothesises and to investigate formation process of dust in SNe, we need further observational and theoretical studies on dust formation in various types of SNe. In particular, the UV to MIR observations at early times for SNe Ib/c are promising to clarify the formation process of dust because its condensation time is much earlier than that for SNe II, as is shown above."
    },
    "0801/0801.0226_arXiv.txt": {
        "abstract": "We study the Friedmann equation for the warped codimension-two braneworld background which most closely resembles the Randall-Sundrum model.  Extra matter on the (Planck) 4-brane, with equation of state $p_\\theta=(\\alpha-1)\\rho$ for the azimuthal pressure, is required to satisfy the junction conditions.  For $1 < \\alpha < 5$, we show that there are two static solutions to the Einstein equations for given values of the brane stress-energies. Close to the static solutions, the relation between Hubble expansion rate $H$ and brane tension reproduces the standard 4D result for small $H$, but exhibits unusual deviations when $H$ is of order the AdS curvature scale.  The two static branches for $1 <\\alpha < 5$ are shown to come together smoothly at a maximum value of $H$; however the radion is shown to be unstable in the branch with higher $H$. This remains true even with a mechanism for stabilization of the radion, {\\it i.e.,} the Goldberger-Wise (GW) mechanism, since large enough $H$ overcomes the force of stabilization. Even in the unstabilized case, cosmological constraints on the time and spatial variation of Newton's constant are typically satisfied; only fifth force constraints require the stabilization.  For $\\alpha > 5$ the model is intrinsically stable, without the need for a GW field, and in this case we show that inflationary predictions can be modified by the nonstandard Friedmann equation; in particular it is possible to get an upper limit on the spectral index, large deviations from the consistency condition between the tensor spectrum and ratio $r$, and large running of the spectral index even though the slow roll parameters remain small. ",
        "introduction": "The cosmology of braneworld scenarios has been widely studied, particularly in the simplest case of codimension-one branes, involving only a single extra dimension (see \\cite{Langlois,Maartens,Brax,Kanno} for reviews). It was initially noticed that the Friedmann equation was strongly modified from its usual dependence $H^2\\sim \\rho$ to the form $H\\sim\\rho$ \\cite{BDL}, in contradiction to big bang nucleosynthesis and other cosmological tests.  Before the realization that such pathological behavior was linked to the failure to stabilize the extra dimension, it was discovered that in the Randall-Sundrum (RS) model \\cite{RSI,RSII} with a warped extra dimension due to a negative bulk cosmological constant, the Friedmann equation took a more interesting form, \\begin{eqnarray} H^2=\\frac{8\\pi G}{3}\\rho\\left(1+\\frac{\\rho}{T}\\right), \\label{mfe} \\end{eqnarray} where $T$ is the brane tension~\\cite{CGKT,Cline,Binetruy}. Thus the exotic linear dependence $H\\sim\\rho$ could arise as a high-energy correction to the normal low-energy dependence.  The behavior (\\ref{mfe}) is only valid for the single-brane version of RS; in the case of two branes, it is necessary to stabilize the interbrane separation \\cite{CGRT}, which leads to a more complicated dependence for the high-energy corrections \\cite{CV}.  Other interesting forms are possible for the Friedmann equation in codimension-one brane models, as in the DGP model~\\cite{Dvali}, whose Lagrangian includes an extra Einstein-Hilbert term localized on the brane, and leads to a quadratic equation determining $H$, \\begin{eqnarray} H^2=\\frac{8\\pi G}{3}\\rho\\mp\\frac{2M_5^3}{M_4^2}H, \\label{5dfei} \\end{eqnarray} in terms of the 5D Planck mass $M_5$ and 4D Planck mass $M_4$. More intricate possibilities also exist, for example through the addition of a Gauss-Bonnet term to the DGP model \\cite{Brown:2005ug}.  There has also been considerable interest in 6D, codimension-two braneworld models, in part motivated by the suggestion of a self-tuning mechanism for the cosmological constant which was argued to be a special feature of codimension-two branes \\cite{Carroll,Burgess:2001bn,Leblond:2001xr,CDGV}, although this claim is controversial \\cite{Vinet,Papa}.  Regardless of the cosmological constant problem, it is still interesting to consider the predictions of codimension-two branes for the expansion of the early universe, since deviations from the general-relativistic prediction can help constrain the models, as well as lead to new possibilities for inflation \\cite{Maartens:1999hf,CLL}.  In the present work we revisit the problem of modifications to the Friedmann equation, focusing on a warped model which is the natural extension of the RS model to six spacetime dimensions \\cite{Send}-\\cite{regularized}. Part of our motivation is a claim \\cite{CDGV} that the standard Friedmann equation is not recovered even when the extra dimensions are stabilized.  We will show that this claim was erroneous, and that although the unstabilized version of the model indeed exhibits unusual expansion behavior at high energies, the expected results of general relativity (GR) are recovered at low energy in the case where the radion is stabilized through a Goldberger-Wise (GW) \\cite{GW} mechanism, and even in the {\\it unstabilized} model.  The detailed form of the modified Friedmann equation depends on the equation of state $\\alpha$ of extra matter which is placed on the 4-brane that plays the role of the UV brane in our model.  We find that for most values of $\\alpha$, there is an exotic branch of the Friedmann equation which has the Hubble rate $H$ increasing as the brane tension {\\it decreases}.  Through a fluctuation analysis, we show that the radion is unstable on this branch, even in the version of the model which includes the GW mechanism.  On the conventional branch of the Friedmann relation, we obtain interesting deviations from the predictions of GR near (but still below) the point where the radion instability begins.  We show that it is possible to obtain potentially observable deviations from the consistency condition which relates the tensor spectral index to the tensor-to-scalar ratio, in a model of chaotic inflation on the brane. We note that there is a suppression of the spectral index in the braneworld inflation model relative to standard chaotic inflation. Furthermore it is  possible (though it requires more fine tuning) to get large running of the spectral index.  The organization of this paper is as follows. In section \\ref{model}, the basic setup is presented. We derive the bulk equation of  motion and the jump conditions for the codimension two background with the 3-brane undergoing de Sitter expansion. We obtain an exact solution in this background. In section \\ref{Feq}, we study the Friedmann equation in the model in the absence of radion stabilization.  The Friedmann equation is obtained using an analytical approximation valid at low $H$, as well as numerically for arbitrary $H$. We find that standard cosmological expansion is recovered at low $H$, but interesting deviations occur when $H\\sim \\ell$, where $\\ell$ is the curvature scale of the AdS background. In section \\ref{FEGW} we repeat the analysis in the more realistic version of the model which includes stabilization of the bulk using the Goldberger-Wise mechanism.  To check the effectiveness of the stabilization as a function of the Hubble rate,  we analyze the stability of the model, with and without the GW field, in section \\ref{STB}. To explore the physical consequences of the modified Friedmann equation, we study chaotic inflation on the brane in section \\ref{IFL}. Section \\ref{CON} gives our conclusions, including a summary of the most important results of the paper. In appendix \\ref{DD}, details of the derivation of the perturbative approximation to the modified Friedmann equation are supplied.  Appendix \\ref{sec:symbols} gives an index of the many symbols used in this paper as an aid to the bewildered. ",
        "conclusions": "\\label{CON}  In this paper we have studied the warped codimension-two braneworld model which most closely resembles the original codimension-one warped scenario of Randall and Sundrum, with the goal of understanding how cosmological expansion (the Friedmann equation) is altered due to the extra dimensions.  The bulk is approximately AdS$_6$ with a 3-brane at the bottom of the throat and a 4-brane at the top; in addition to tension, the 4-brane has an extra nontensional source whose pressure in the angular direction varies with the brane circumference $\\sim L_4$ as $L_4^{-\\alpha}$.  This extra source is needed in order to find solutions with localized gravity when no stabilization mechanism is included, and it is known \\cite{Burgess:2001bn} that the radion is stable when $\\alpha> 5$.  It generally  plays a less important role when stabilization is induced, for example, using the Goldberger-Wise (GW) mechanism.  We determined the modified Friedmann equation both analytically, in a perturbative expansion in powers of the energy density on the brane, and numerically.  The perturbative result, eq.\\ (\\ref{FRW35}), has the curious feature that $H^2 \\sim \\rho + O(\\rho^3)$, {\\it i.e.,} the $O(\\rho^2)$ correction vanishes.  For larger values of $\\rho$, the numerical solution is necessary, as described in section \\ref{NFE}.  We checked that the two kinds of solutions agree with each other in the small-$\\rho$ region where both are valid.  In our study of the modified Friedmann equation for this model, we have corrected a mistaken claim of ref.\\ \\cite{CDGV}, which did not recover the standard GR result at low energies.   In fact we find standard cosmology at low energy even for the unstabilized models. At high energies, we find exotic features, in which the Hubble rate $H(\\rho)$ turns over and joins with another branch of the function at some critical value of $\\rho$ and $H$,  making $H(\\rho)$ double-valued---see figures \\ref{alpha} and \\ref{fws}.\\footnote{The same phenomenon occurs in the similar model of ref.\\ \\cite{Send}, but it is not commented upon there.}  This behavior even persists in the model with GW stabilization, but we have found that the exotic branch is nevertheless unphysical, because the model becomes destabilized by the Hubble expansion  at the turning point.  We discovered this by doing a small fluctuation analysis around the de Sitter solutions in section \\ref{STB}, in which the eigenvalue problem for the radion mass $m_r^2$ was solved both using analytical approximations and exact numerical methods.  This result  is reassuring, since otherwise we would be left with the puzzle of why the cosmological background is not uniquely determined by the sources of stress energy in the Einstein equation.  Nevertheless, we have found that interesting deviations from the standard Friedmann equation occur near the threshold for radion instability: $dH^2/d\\rho$ diverges at this point, even though $H$ itself is finite.  In fact for the specific example we considered, $H$ only differs from its standard value by 30\\% at this point, but the divergence in $dH^2/d\\rho$ can have observable effects if inflation happened to start near the maximum value of the potential allowed by radion stability.  We demonstrated such effects in the simplest model of chaotic inflation on the 3-brane. Figure \\ref{index} shows that the spectral index is lower than in standard chaotic inflation; for instance with $N_e=55$ e-foldings, $n_s=0.964$ for standard chaotic inflation, whereas it is $0.950$ for the braneworld model with $N_m = 65$ as the maximum number of e-foldings (hence the maximum value of the potential is $65/55$ times its value at horizon crossing).  Thus without much fine tuning, the braneworld model can have a potentially measurable impact on the spectral index.  We then explored the breaking of the  consistency condition between the tensor-to-scalar ratio $r_t$  and the tensor spectral index $n_t$, which may be observable if $r_t$ is close to its experimental upper limit.  Figure \\ref{ntr2} shows that $8n_t/r_t$ differs by 50\\% from the standard prediction of $-1$ even when $V$ is only within 20\\% of its maximum value, so again interesting signatures can occur without excessive fine tuning of the model parameters.  It would be interesting to try to extend our results to a string theoretic realization of inflation in a warped throat, namely the Klebanov-Strassler background \\cite{KS} with stabilization by fluxes \\cite{GKP}.  Although we have used GW rather than flux stabilization in our model, ref.\\ \\cite{Brummer:2005sh} has argued that an effective field theory description of flux stabilization might coincide with the GW picture.  So far, efforts in string theoretic inflation with warped throats have focused on brane-antibrane \\cite{KKLMMT} or DBI \\cite{DBI} inflation.  However the effects we have discussed should also be applicable in string-based constructions with conventional inflation confined to a brane.  The phenomenon of radion destabilization at sufficiently high Hubble rate should also occur in string compactifications, since the decompactified background is always a solution, and the radion, being conformally coupled, gets a mass correction of order $H^2$ in the early universe.  The $H^2\\phi^2$ term in the radion effective action thus drives $\\phi$ away from its nontrivial minimum toward the decompactified $\\phi=0$ vacuum for sufficiently large $H$, and it is likely that deviations from the usual Friedmann equation will be large near this point.  Therefore it is possible that stringy effects due to the extra dimensions may be manifested during inflation even if the inflaton is confined to the standard model brane."
    },
    "0801/0801.2962_arXiv.txt": {
        "abstract": "Due to quantum fluctuations, spacetime is foamy on small scales.  The degree of foaminess is found to be consistent with holography, a principle prefigured in the physics of black hole entropy.  It has bearing on the ultimate accuracies of clocks and measurements and the physics of quantum computation.  Consistent with existing archived data on active galactic nuclei from the Hubble Space Telescope, the application of the holographic spacetime foam model to cosmology requires the existence of dark energy which, we argue, is composed of an enormous number of inert ``particles\" of extremely long wavelength.  We suggest that these ``particles\" obey infinite statistics in which all representations of the particle permutation group can occur, and that the nonlocality present in systems obeying infinite statistics may be related to the nonlocality present in holographic theories.  We also propose to detect spacetime foam by looking for halos in the images of distant quasars, and argue that it does not modify the GZK cutoff in the ultra-high energy cosmic ray spectrum and its contributions to time-of-flight differences of high energy gamma rays from distant GRB are too small to be detectable.\\\\   {\\it Keywords}: spacetime foam, quantum foam, holography, dark energy, infinite statistics, nonlocality\\\\ ",
        "introduction": "Following John Wheeler, many physicists (including the author) believe that space is composed of an ever-changing arrangement of bubbles called spacetime foam, also known as quantum foam.  To understand the terminology, let us consider the following simplified analogy which Wheeler gave in a conference on gravity at University of North Carolina in 1957.  Imagine yourself flying an airplane over an ocean.  At high altitude the ocean appears smooth.  But as you descend, it begins to show roughness.  Close enough to the ocean surface, you see bubbles and foam.  Analogously, spacetime appears smooth on a large scale, but on sufficiently small scales, it will appear rough and foamy, hence the term ``spacetime foam'' \\cite{Wheeler}. Many physicists believe the foaminess is due to quantum fluctuations of spacetime, hence the alternative term ``quantum foam.''  This reveiw article is devoted to a discussion of spacetime foam models, or rather, a specific one of them for which the concept of entropy plays a crucial role \\cite{wigner,Karol, found,PRL,llo04}. The only ingredients we use in the whole discussion are quantum mechanics and general relativity, the two pillars of modern physics.  We hope that the results are very general and have wider validity than most of the candidates of quantum gravity theory. Assuming that there is a unity of physics connecting the Planck scale to the cosmic scale, we also apply quantum foam physics to cosmology \\cite{CNvD,Arzano,plb}.  In the next section, by a process of mapping the geometry of a region of spacetime, we show how foamy spacetime is.  It turns out that the degree of foaminess of spacetime determines the maximum amount of entropy a spatial region can hold.  This maximum amount of information can be shown to be consistent with that encoded in the holographic principle which has its origin in black hole physics.  Appropriately this spacetime foam model has come to be known as the holographic model.  In section 3, we propose to probe spacetime foam by looking for halos in the images of distant quasars or bright active galactic nuclei (AGN). There we show that the archive from the Hubble Space Telescople (HST) can already be used to yield useful bounds.  Among the spacetime foam models, the holographic foam model is unique in its correspondence with the case of maximum energy density that a spatial region can hold without collapsing into a black hole.  Applied to cosmology it ``predicts\" a critical energy density as observed in recent years.  This potential connection between the extremely large and the extremely small is explored in section 4; there we show that the archive on quasars and AGN from HST can be used, in conjuction with the physics of quantum computation, to ``prove\" the existence of dark energy, {\\it independent} of the various cosmological observations of recent years.  Indeed, from our perspective, dark energy is arguably a cosmological manifestation of quantum foam!  In section 5, we examine how and why dark energy is so different from ordinary energy/matter.  There we show that the holographic model of spacetime foam naturally leads us to speculate that dark energy consists of an incredibly large number of extremely-long-wavelength ``particles\". Then we exploit the positivity of the entropy for these particles to show that they obey not the ordinary bose or fermi statistics, but the exotic infinite statistics (also known as quantum Boltzmann statistics).  We start our discussion with holography and end with infinite statistics.  We conclude section 5 by speculating that the nonlocality known to be present in both of them may be related to each other.  We give a summary in section 6.  For completeness, we discuss various other relevant topics of spacetime foam physics in six short appendices.  In Appendix A, using a gedanken experiment to measure distances we rederive the holographic foam model.  In Appendix B, we apply the results in Appendix A to set bounds on the accuracies of clocks and limits on computations, and derive the black hole entropy and lifetime. We derive the holographic principle and the Margolus-Levitin theorem (which is liberally used in the text) in Appendices C and D respectively.  Appendix E is devoted to a discussion of the uncertainties in energy-momentum measurements consistent with spacetime foam. In Appendix F we apply the results obtained in Appendix E to discuss high-energy $\\gamma$ rays from distant gamma ray bursts and ultra-high energy cosmic ray events.  On notations, the subscript ``P'' denotes Planck units; thus $l_P \\equiv (\\hbar G /c^3)^{1/2} \\sim 10^{-33}$ cm is the Planck length etc.  On units, $k_B$ (the Boltzmann constant) and sometimes $\\hbar$ and c are put equal to 1 for simplicity.\\\\ ",
        "conclusions": "A short summary of this review article is given by the accompanying flow chart Figure~\\ref{figure2}.  \\begin{figure} \\[ \\xymatrix{ & *\\txt{QM $+$\\\\ bh physics} \\ar[d] \\ar[dr]^(.6){\\txt{bh entropy}}  & \\\\ \\Omega \\cong 1 \\ar@{<->}[r] \\ar[d]_{\\txt{AGN data}} & *\\txt{holographic\\\\ st foam} \\ar@{<->}[r] \\ar[d] & *\\txt{holography} \\ar@{--}[dl]^(.4){\\txt{nonlocality}} \\\\ \\txt{dark energy} \\ar[r] & \\infty~\\textrm{statistics} & } \\] \\caption{Connecting the different ideas.} \\label{figure2} \\end{figure}   Starting from quantum mechanics and some rudimentary black hole physics (viz, there is an upper bound of matter/energy that can be put in a spatial region without the region collapsing into a black hole), we derive the holographic model of spacetime foam. From the simple observation that the cosmic energy density is critical (i.e., the density parameter of the universe $\\Omega \\cong 1$), aided by some archived data on quasar or AGN from the Hubble Space Telescope, we are led to conclude that dark energy exists. This observation is a solid piece of evidence in favor of the holographic foam model which, after all, ``predicts'' that the cosmic energy density has the critical value. One is then invited to apply the holographic model to cosmology, leading to the logical speculation that the constituents of dark energy obey infinite statistics.  We have now come full circle, starting with holography and ending with infinite statistics, with both sharing one common property: nonlocality.  In the whole discussion, entropy plays a leading role, first in motivating holography, and secondly in determining, via its positivity property, that the ``particles'' constituting dark energy obey infinite statistics.  We conclude with some observation and a partial list of open questions.  (1) We have considered a perfect gas consisting of ``particles\" of extremely long wavelength, obeying Boltzmann statistics in the Universe at temperature $T$.  But we have seen that these ``particles\" have had interactions only of order one time on the average during the entire cosmic history. A question can be raised as to whether such an inert gas can come to thermal equilibrium at any well defined temperature.  We do not have a good answer; but the fact that all these ``particles\", though extremely inert, have a wavelength comparable to the observable Hubble radius may mean that they overlap significantly, and accordingly can perhaps share a common temperature.  (2) These ``particles'' provide a (more or less) spatially uniform energy density, like a time-dependent cosmological constant \\cite{ChenWu}.  But in a way, this type of models is preferrable to the cosmological constant because it may be easier to understand a zero cosmological constant (perhaps due to a certain not-yet-understood symmetry or initial condition \\cite{unimod}) than an exceedingly small (but non-zero) cosmological constant.  (3) Recall that earlier cosmic epochs are associated with $\\rho \\propto a^{-4}$ (radiation-domination) and (followed by) $\\rho \\propto a^{-3}$ (matter-domination).  If the holographic foam cosmology is correct, these epochs are succeeded by the dark-energy-dominated era with $\\rho \\propto a^{-2}$.  We note that the successes of the conventional big bang cosmology, such as the primordial nucleosynthesis, are not affected by this form of dark energy.  (4) An obvious question concerns the sign of the pressure for the gas of dark energy ``particles''. For cosmic energy density $\\rho \\sim H^2/G$, the equation of state is given by $w = p/\\rho \\sim -1/3$, not negative enough to give an accelerating expansion.  However, it has been pointed out \\cite{negativep} that a transition from an earlier decelerating to a recent and present accelerating cosmic expansion can arise as a pure interaction phenomenon if pressureless dark matter is coupled to holographic dark energy \\cite{holocos}. As a bonus, within the framework of such cosmological models, we can now understand why, in addition to dark energy, dark matter has to exist. On the other hand, this scenario will become less natural if the equation of state $w$ turns out to be very close to $-1$.  (5) Critical cosmic energy density is the hallmark of the inflationary paradigm \\cite{zizzi}.  Can the holographic foam cosmology (which requires $\\Omega \\cong 1$) supplement inflation in solving cosmological problems and in providing the necessary primordial perturbations to yield the observed astronomical structures?  What are the phenomenological consequences of HFC?  (6) While it appears quite reasonable that holography (and quantum gravity in general) and infinite statistics are compatible, exactly how they are related is not yet clear.  In particular, how is the nonlocality in holography related to the nonlocality present in theories with infinite statistics?  (We venture a guess: Entropy is the common link, so it may hold the key in understanding the relation.)  Before last century, spacetime was regarded as nothing more than a passive and static arena in which events took place. Early last century, Einstein's general relativity changed that viewpoint and promoted spacetime to an active and dynamical entity.  Quantum mechanics blossomed around that time.  But the challenge to understand the quantum nature of spacetime was not taken up seriously until quite a bit later. Now many physicists believe that spacetime, like all matter and energy, undergoes quantum fluctuations.  These quantum fluctuations make spacetime foamy on small spacetime scales. In this article, we have used a global positioning-like system to measure the geometry of spacetime, to show that spacetime fluctuations, on the average, scale as the cube root of distances.  This result is very general, depending only on quantum mechanics and limited black hole physics (viz., the size of a black hole scales linearly with its mass). The cube root dependence is strange, but it has been shown to be consistent with the holographic principle and (as shown in Appendix B) semi-classical black hole physics in general.  Furthermore, applied to cosmology, it successfully ``predicts\" the existence of dark energy and that the cosmic energy density is critical.  To the author at least, the cube-root result for spacetime foam is as beautiful as it is strange --- and, when the Very Large Telescope Interferometer reaches its design performance, it may even be proven to be true.\\\\  \\bigskip"
    },
    "0801/0801.0683_arXiv.txt": {
        "abstract": "{Several approaches to estimate frequency, phase and amplitude errors in time series analyses were reported in the literature, but they are either time consuming to compute, grossly overestimating the error, or are based on empirically determined criteria.} {A simple, but realistic estimate of the frequency uncertainty in time series analyses.} {Synthetic data sets with mono- and multi--periodic harmonic signals and with randomly distributed amplitude, frequency and phase were generated and white noise added. We tried to recover the input parameters with classical Fourier techniques and investigated the error as a function of the relative level of noise,  signal and frequency difference.} {We present simple formulas for the upper limit of the amplitude, frequency and phase uncertainties in time--series analyses. We also demonstrate the possibility to detect frequencies which are separated by less than the classical frequency resolution and that the realistic frequency error is at least 4 times smaller than the classical frequency resolution.} {} ",
        "introduction": " ",
        "conclusions": "Based on extensive simulations, we have shown that there is an upper limit for the amplitude and frequency error in time series data analyses. Compared to the statistically {\\em expected} value for the uncertainties given by \\cite{Montgomery}, our {\\em upper limits} cover the possible error due to white noise and leaves even room for additional error sources like atmospheric scintillation. A major advantage of calculating amplitude, frequency and phase errors in terms of spectral significance rather than signal--to--noise ratio is that the time--domain noise need not be Gaussian. As pointed out by \\cite{Reegen}, the spectral significance does not depend on the probability distribution associated to the noise, and the only precondition is uncorrelatedness of consecutive data points. There has to be mentioned that amplitude, frequency and phase errors derived from spectral significances are only comparable to errors derived from SNR if the time series is well sampled (e.g. continuos space observations). Contrary to spectral significance based errors, SNR based error estimations (time--domain as well as frequency--domain) do not take into account the data sampling and can yield in a crude underestimation of the errors for ``bad\" sampling like it is more or less always the case for single--site ground based observations.  We have shown that the phase error defined by \\cite{Montgomery} is consistent with our simulations.  Furthermore we have shown that the determination of frequency pairs closer than the Rayleigh frequency resolution is possible and that the resulting frequency error is still 4 times smaller than the Rayleigh frequency resolution. However, our simulation does not say anything about the reliability of close frequency pairs in general. It tells us about the frequency uncertainty of a peak if, after prewhitening this peak, a second significant peak is present. It tells us that peaks do not influence each other's frequency determination if they are separated in frequency by 3 times the Rayleigh frequency resolution. For closer peaks the frequency uncertainty is at least 4 times below the Rayleigh resolution even for peaks within the Rayleigh resolution."
    },
    "0801/0801.1685_arXiv.txt": {
        "abstract": "We report the discovery of WASP-5b, a Jupiter-mass planet orbiting a 12th-mag G-type star in the Southern hemisphere.  The 1.6-d orbital period places WASP-5b in the class of very-hot Jupiters and leads to a predicted equilibrium temperature of 1750 K. WASP-5b is the densest of any known Jovian-mass planet, being a factor seven denser than TrES-4, which is subject to similar stellar insolation, and a factor three denser than WASP-4b, which has a similar orbital period. We present transit photometry and radial-velocity measurements of WASP-5 (= USNO-B1\\,0487-0799749), from which we derive the mass, radius and density of the planet: M$_{\\rm P}$ = 1.58 $^{+0.13}_{-0.08}$ M$_{\\rm J}$, R$_{\\rm P}$ = 1.090 $^{+0.094}_{-0.058}$ R$_{\\rm J}$ and $\\rho_{\\rm P}$ = 1.22 $^{+ 0.19}_{- 0.24}$ $\\rho_{\\rm J}$. The orbital period is P = 1.6284296 $^{+0.0000048}_{-0.0000037}$ d and the mid-transit epoch is T$_{\\rm C}$ (HJD) = 2454375.62466 $^{+0.00026}_{-0.00025}$. ",
        "introduction": "Several wide-angle, ground-based surveys for transiting extrasolar planets have now been successful, including HAT \\citep{2002PASP..114..974B}, TrES \\citep{2006ApJ...651L..61O}, WASP \\citep{2006PASP..118.1407P} and XO \\citep{2005PASP..117..783M}.  In the Southern hemisphere relatively faint ($V$ = 16--17) transiting systems have been found by OGLE \\citep{1992AcA....42..253U,2007arXiv0710.5278P} and by \\citet{2007arXiv0711.1746W}. The WASP-South survey is the first to find much brighter transiting systems in the South \\citep{WASP-4b}. We report here the second planet found by the WASP-South and CORALIE collaboration, a dense, very-hot Jupiter transiting a $V$ = 12.3 star with an orbital period of 1.6 d.  Transit searches are most sensitive to large, Jupiter-sized planets in very short orbits. WASP-5b, like the recently announced WASP-3b and WASP-4b, has an ultra-short period below 2 d \\citep{WASP-3b,WASP-4b}.  We discuss the characteristics of these ultra-close, highly irradiated planets. ",
        "conclusions": "The discovery of WASP-3b \\citep{WASP-3b}, WASP-4b \\citep{WASP-4b} and WASP-5b, all transiting, Jupiter-mass planets with orbital periods of less than 2 d, increases the number of such systems known from four to seven\\footnote{http://www.inscience.ch/transits/}. It is notable that they show a large disparity in density (Fig.~\\ref{dens-per}). WASP-5b is the densest of these systems, whereas WASP-4b is the least dense, with a density a third that of WASP-5b. This illustrates that planets with a similar orbital period and radiation environment must still have different compositions or past histories.  For example, WASP-5b will receive the same insolation and thus has a similar estimated equilibrium temperature (1753 K) to TrES-4 (1720 K), which has a longer orbit about a more luminous star \\citep{2007ApJ...667L.195M}. Despite this, WASP-5b is denser by a factor of 7 (1.22 $\\rho_{J}$ compared with 0.167 $\\rho_{J}$ for TrES-4). \\citet{2007ApJ...661..502B} postulate that the bloating of the under-dense planets could be due to an enhanced atmospheric opacity, which could be absent in WASP-5b.  Although it is denser than other planets with a similar orbital period, WASP-5b is still within the expected theoretical range \\citep*{2007ApJ...659.1661F}; depending on the age of the system, a core mass of $\\sim$\\,50--100 M$_{\\varoplus}$ is expected, with a younger planet requiring a more massive core. However, if the bloating of some planets were due to an extra energy source that is common to all hot Jupiters, then the core of a dense planet such as WASP-5b would have to be more massive to compensate \\citep{2006ApJ...642..495F,2006A&A...453L..21G}. If the core mass is toward the higher end of the estimate and the metallicity is no more than solar then there could be a discrepancy with the relationship between core mass and stellar metallicity postulated by \\citealt{2007ApJ...661..502B}. High resolution spectra with better S/N are required to reduce the uncertainties in the metallicity and age.  WASP-5b ($\\log g_{\\rm *}$ = 3.48) has the highest surface gravity of planets with orbital periods less than 5 d.  Thus it could be a good system for studying evaporation as a function of surface gravity. In comparison HD\\,209458b is heavily bloated ($\\log g_{\\rm *}$ = 2.97) and appears to be undergoing evaporation \\citep{2003Natur.422..143V,2007ApJ...671L..61B,2008arXiv0802.0587V}. Given its high equilibrium temperature, WASP-5b is also expected to have a prominent secondary transit, which should be detectable with $\\itl Spitzer$ (e.g. \\citealt{2007Natur.447..691H}). Further, with the moderately bright host star having a moderate rotation rate of $v \\sin i = 3.4$ km~s$^{-1}$, the Rossiter--McLaughlin effect (40 $\\pm$ 8 m~s$^{-1}$) should be detectable \\citep{2007ApJ...655..550G},  allowing us to check the alignment between the planetary orbit and the stellar spin axis.  \\begin{figure} \\includegraphics[width=60mm, angle=270]{dens-per2.ps} % \\caption{Planet density versus orbital period for planets with periods less than 5 d. The dense, hot Neptune, GJ\\,436b, and under-dense TrES-4 are labelled; the five WASP planets are labelled numerically. The error bars on the densities were calculated from the uncertainties on the masses and radii, which were assumed to be uncorrelated. Data taken from http://www.inscience.ch/transits/} \\label{dens-per} \\end{figure}"
    },
    "0801/0801.1199_arXiv.txt": {
        "abstract": "We study the dynamics of the closed scalar field FRW cosmological models in the framework of the so called {\\it Unified Dark Matter} (UDM) scenario. Performing a theoretical as well as a numerical analysis we find that there is a strong indication of chaos in agreement with previous studies. We find that a positive value of the spatial curvature is essential for the appearance of chaoticity, though the Lyapunov number seems to be independent of the curvature value. Models that are close to flat ($k \\rightarrow 0^{+}$) exhibit a chaotic behavior after a long time while pure flat models do not exhibit any chaos. Moreover, we find that some of the semiflat models in the UDM scenario exhibit similar dynamical behavior with the $\\Lambda$ cosmology despite their chaoticity.  Finally, we compare the measured evolution of the Hubble parameter derived from the differential ages of passively evolving galaxies with that expected in the semiflat unified scalar field cosmology. Based on a specific set of initial conditions we find that the UDM scalar field model matches well the observational data. ",
        "introduction": "There is by now convincing evidence that the available high quality cosmological data (Type Ia supernovae, CMB, etc.) are well fitted by an emerging cosmological model, which contains cold dark matter to explain clustering and an extra component with negative pressure, the vacuum energy (or in a more general setting the ``dark energy''), to explain the observed accelerated cosmic expansion (Refs. \\cite{Riess07,Spergel07} and references therein). Due to the absence of a physically well-motivated fundamental theory, there have been many theoretical speculations regarding the nature of the above exotic dark energy. Indeed, many authors claim that a real scalar field which rolls down the potential $U(\\phi)$ \\cite{Ozer87,Peebles88,Weinberg89,Turner97,Caldwell98,Varum00,Padm03,Peebles03} could resemble the dark energy.  On the other hand, the global dynamics, in models with dark energy, become more complicated than in matter dominated models. Within this framework, a serious issue is the existence (or nonexistence) of chaos in the scalar field cosmology. This is very important because chaotic fields can provide possible solutions to the cosmological coincidence problem (the question why $\\Omega_m$ and $\\Omega_\\Lambda$ are of the same order at the present time) as well as to the problem of uniqueness of vacua. The first paper on chaos in a closed Friedman Robertson Walker geometry coupled to a scalar field (FRW with a scalar field potential $U(\\phi)\\propto \\phi^{2}$) was written by Page \\cite{Page84}. This author, suggested that the dynamical behavior of this non-integrable model is similar to the chaos appearing in ergodic systems. Since then, the problem of chaos in FRW scalar field cosmologies has been investigated thoroughly in several papers \\cite{Calzetta93,Cornish96,Kamen97,Peebles03,Joras03,Kanekar01,Beck04,Heinz05,Szydlow05,Toporensky06,Szydlow07}. The dynamical properties of such cosmological models are met in dynamical systems exhibiting chaotic scattering. However, in this kind of studies there is a crucial question which has to be addressed: {\\it how can we estimate the amount of chaos?} It has been shown that in chaotic scattering problems classical tools, like the Lyapunov characteristic number, are unable to provide a clear answer whether an orbit is chaotic or not \\cite{Contop99}.  The aim of this work is to investigate the dynamical properties of a family of cosmological models based on a scalar field potential which appears to act both as a dark matter and dark energy \\cite{Varum00,Bertacca06}. In order to measure the amount of chaos we follow the notation of Ref.\\cite{Contop99} that was able to give useful answers to their scattering problem. These authors used an index of chaoticity which is capable to detect stochastic areas in scattering problems. To this end, we attempt to compare our theoretical predictions with observational data by utilizing the measured evolution of the Hubble parameter \\cite{Simon05}.  The structure of the paper is as follows. The basic theoretical elements of the problem are presented in section 2. In particular, in the first subsection we study the dynamical system analytically and in the second subsection we consider various chaoticity indices. Section 3 outlines the numerical results of the present study together with a thorough discussion, and finally we draw our conclusions in section 4. ",
        "conclusions": "In this work we investigate the dynamics of the closed scalar field FRW cosmological models in the framework of the so called {\\it Unified Dark Matter} scenario. We find that for the closed geometry there is a strong indication of chaos in agreement with previous studies and that the $LCN$ of these models is independent from the curvature. In particular, we find that there are semiflat cosmological models with specific initial conditions in which there is a clear indication for a chaotic behavior and at the same time the corresponding dynamics is almost the same as in the $\\Lambda$ cosmology. We verify this by combaning the measured evolution of the Hubble parameter with that expected in the scalar field cosmology (with specific initial conditions) and we find a very good agreement. If that is the case, then the chaotic fields may provide an alternative theory for the solution of the cosmological coincidence problem."
    },
    "0801/0801.4891_arXiv.txt": {
        "abstract": "#1{\\vskip 7mm \\begin{center}{\\large Abstract}\\par \\smallskip \\begin{minipage}[c]{12cm} \\small #1 \\end{minipage} \\end{center} } \\def\\title#1{\\begin{center}{\\Large\\bf #1}\\end{center}} \\def\\author#1{\\vskip 5mm \\begin{center}{#1}\\end{center}} \\def\\address#1{\\begin{center}{\\it #1}\\end{center}} \\makeatletter \\def\\vereq#1#2{\\lower3pt\\vbox{\\baselineskip1.5pt \\lineskip1.5pt \\ialign{$\\m@th#1\\hfill##\\hfil$\\crcr#2\\crcr\\sim\\crcr}}} \\makeatother   \\usepackage{graphicx} \\usepackage{amssymb} \\usepackage[mathscr]{eucal}   \\newcommand{\\gs}{g_{\\mathrm{s}}} \\newcommand{\\Nflux}{\\mathcal{N}} \\newcommand{\\phistrg}{\\phi_{\\mathrm{strg}}} \\newcommand{\\phiuv}{\\phi_{_{\\mathrm{UV}}}} \\newcommand{\\alphas}{\\alpha'} \\newcommand{\\mpl}{m_{\\mathrm{Pl}}} \\newcommand{\\phiini}{\\phi_{\\mathrm{in}}} \\newcommand{\\phiend}{\\phi_{\\mathrm{end}}} \\newcommand{\\phieps}{\\phi_{\\epsilon}} \\newcommand{\\phiDBI}{\\phi_{_{\\mathrm{DBI}}}}     \\begin{document}  \\title{% Constraints on brane inflation from WMAP3 }  \\author{Larissa Lorenz\\footnote{lorenz@iap.fr}}% \\address{$^{1}$Institut d'Astrophysique de Paris, UMR 7095-CNRS, Universit\\'e Pierre et Marie Curie,\\\\98bis boulevard Arago, 75014 Paris, France}    \\abstract{ Considerable and ongoing effort is made to identify promising scalar field candidates in string theory to drive a cosmological period of inflation. At stake is the possibility that fundamental string parameters could be encoded in observables such as the CMB perturbation spectrum. In this contribution, we hold a concrete model of string inflation (KKLMMT) up against WMAP3 and discuss the constraints obtained. } ",
        "introduction": "In recent years, the hope of embedding cosmological inflation into superstring theory has been put on more solid grounds. While they remain challenging, issues such as moduli stabilization are better understood, and scenarios for both open and closed string mode inflatons have been constructed. With its tight relation to observables of current and future CMB experiments, inflation could provide the decisive missing link between string theory and observation. We investigate if the WMAP3 data provides constraints on the parameters of one particular (open string) scenario, known as the KKLMMT model of brane inflation \\cite{kklmmt}. To this end, we identify its cosmological parameters and how they relate to the underlying string geometry, followed by a comparison to the WMAP3 data using numerical integration of the perturbations and MCMC methods \\cite{ourpaper} (see also \\cite{tye}). ",
        "conclusions": "A general result of this work is that, \\emph{in principle}, it seems possible to constrain stringy parameters from cosmology. However, the accuracy of present data does not suffice to break the degeneracies. Moreover, one must not underestimate the strong \\emph{theoretical} prior that comes with any attempt at cosmological model building in string theory, since the testable inflationary quantities derive from fundamental (e.g. geometric) choices for the background.  In \\cite{ourpaper}, we presented the first complete MCMC analysis of  the pure KKLMMT model, considering $\\gs$ and $\\alpha'$ as free parameters. We also suggest how to systematically scan the parameter space for arbitrary $\\gs,\\alpha'$. The data favour those cases where inflation occurs in the usual slow-roll way, i.e. inflation ends at $\\phieps$ and not earlier at brane annihilation. This is because $\\phiend=\\phistrg>\\phieps$ would push $n_{\\mathrm{s}}\\rightarrow 1$, while preserving a low level of gravitational waves (see \\cite{ourpaper}). A weak limit on $v$, i.e. on a parameter of the 6d compactification, is also obtained: $\\log v > -10$ at $95\\%$ CL.  The choice of the pure KKLMMT scenario \\cite{kklmmt} comes with a considerable caveat: All moduli are considered stabilized, and various additional contributions to $V(\\phi)$ are assumed to combine in such a way as to only leave the Coulomb term (\\ref{potential}). Recently, these contributions became calculable \\cite{corrections}, and their general cancellation is unlikely. In practice, the full potential should be of the form $V(\\phi)=V_{D\\bar{D}}(\\phi)+m^{2}\\phi^{2}+\\dots$, leading to a completely different inflaton evolution and notably rendering the DBI phase important. The next step would be to include these terms in our analysis, at the expense of introducing additional parameters."
    },
    "0801/0801.4319_arXiv.txt": {
        "abstract": "We describe and discuss the spectral and temporal characteristics of the prompt emission and X-ray afterglow emission of X-ray flashes (XRFs) and X-ray-rich gamma-ray bursts (XRRs) detected and observed by {\\it Swift} between December 2004 and September 2006.  We compare these characteristics to a sample of conventional classical gamma-ray bursts (C-GRBs) observed during the same period.  We confirm the correlation between $\\eop$ and fluence noted by others and find further evidence that XRFs, XRRs and C-GRBs form a continuum.  We also confirm that our known redshift sample is consistent with the correlation between the peak energy in the GRB rest frame ($\\esp$) and the isotropic radiated energy ($\\eiso$), so called the $\\esp$-$\\eiso$ relation.  The spectral properties of X-ray afterglows of XRFs and C-GRBs are similar, but the temporal properties of XRFs and C-GRBs are quite different.  We found that the light curves of C-GRB afterglows show a break to steeper indices (shallow-to-steep break) at much earlier times than do XRF afterglows.  Moreover, the overall luminosity of XRF X-ray afterglows is systematically smaller by a factor of two or more compared to that of C-GRBs.  These distinct differences between the X-ray afterglows of XRFs and C-GRBs may be the key to understanding not only the mysterious shallow-to-steep break in X-ray afterglow light curves, but also the unique nature of XRFs. ",
        "introduction": "Despite the rich gamma-ray burst (GRB) sample provided by BATSE \\citep[e.g.,][]{paciesas1999,kaneko2006}, Beppo{\\it SAX} \\citep[e.g.,][]{frontera2004}, $Konus$-$Wind$ \\citep[e.g.,][]{ulanov2004}, and $HETE$-2 \\citep[e.g.,][]{barraud2003,sakamoto}, the emission properties of GRBs are still far from being well-understood.  In recent years, however, another phenomenon that resembles GRBs in almost every way, except that the flux comes mostly from X rays instead of $\\gamma$ rays, has been discovered and studied.  This new class of bursts has been dubbed ``X-ray flashes'' (XRFs; \\citet{heise2,barraud2003,sakamoto}), and there is strong evidence to suggest that ``classical'' GRBs (hereafter C-GRBs) and XRFs are closely-related phenomena. Understanding what physical processes lead to their differences could yield important insights into their nature and origin.  \\citet{strohmayer} identified 22 bursts observed by $Ginga$ that occurred between March of 1987 and October of 1991, and for which the spectra could be reliably analyzed. About 36\\% of GRBs observed by $Ginga$ had very soft spectra. They noted that these bursts resembled BATSE long GRBs in duration and general spectral shape, but the peak energies of the $\\nu$F$_{\\nu}$ spectrum, $\\eop$, extended to lower values than those of the BATSE bursts \\citep{preece2000,kaneko2006}. \\citet{heise2} reported that among the sources imaged by the Wide Field Cameras (WFCs) on board {\\it Beppo}SAX was a class of fast X-ray transients with durations less than 1000 s that were not ``triggered'' (that is, detected) by the Gamma Ray Burst Monitor (GRBM).  This became their working definition of XRFs. \\citet{kippen} searched for C-GRBs and XRFs which were observed simultaneously by WFC and BATSE.  They found 36 C-GRBs and 17 XRFs in a 3.8-year period.  Joint WFC and BATSE spectral analysis was performed for the sample, and they found that XRFs have a significantly lower $\\eop$ compared with C-GRBs.  They also found that there is no systematic difference between XRFs and C-GRBs in their low-energy photon indices, high-energy energy photon indices, or durations. The systematic spectral analysis of a sample of 45 {\\it HETE-2} GRBs confirmed these spectral and temporal characteristics of XRFs. It is worth noting that nine out of sixteen XRF samples of {\\it HETE-2} have $\\eop$ $<$ 20 keV \\citep{barraud2003,sakamoto}.  Although the XRF prompt emission properties have been studied, until the launch of {\\it Swift}~\\citep{gehrels}, only a handful of X-ray afterglows associated with XRFs were reported. \\citet{alessio2005} studied the prompt and afterglow emission of XRFs and X-ray-rich GRBs (XRRs) observed by {\\it Beppo}SAX and {\\it HETE-2}. They found that the XRF and XRR afterglow light curves seem to be similar to those of C-GRBs, including the break feature in the light curves. They also investigated the off-axis viewing scenarios of XRFs for the top-hat shaped jet \\citep{yamazaki2002,yamazaki2004}, the universal power-law shaped jet \\citep{rossi2002,zhang2002,lamb2005}, and the Gaussian jet \\citep{zhang2004}, and concluded that these models might be consistent with the data. Their sample, however, only contains 9 XRFs/XRRs with measured X-ray afterglows.  Furthermore, the data points in the X-ray light curves were not well sampled, so that there are large uncertainties in the decay indices and the overall structures of the light curve in most cases.  Moreover, since the X-ray afterglow observations began $>$ 10$^{4}$ seconds after the trigger, their sample is able to say little about the early afterglow properties, which contain rich information that can constrain jet models for XRFs. Other XRF theoretical models are the inhomogeneous jet model \\citep{toma2005}, the internal shock emission from high bulk Lorentz factor shells \\citep{mochkovitch2003,barraud2005}, the external shock emission from low bulk Lorentz factor shells \\citep{dermer1999,dermer2003}, and the X-ray emission from the hot cocoon of the GRB jet if viewed from off-axis \\citep{meszaros2002,woosley2003}.  Because of the sophisticated on-board localization capability of the {\\it Swift} Burst Alert Telescope (BAT; \\citet{barthelmy}) and the fast spacecraft pointing of {\\it Swift}, more than 90\\% of {\\it Swift} GRBs have an X-ray afterglow observation from the {\\it Swift} X-Ray Telescope (XRT; \\citet{burrows}) within a few hundred seconds after the trigger. Due to the fact that BAT is sensitive to relatively low energies (15-150 keV) and also a large effective area ($\\sim$ 1000 cm$^{2}$ at 20 keV for a source on-axis), BAT is detecting also XRFs and XRRs.  However, because of the BAT's lack of response below 15 keV, it is very challenging to detect XRFs with $\\eop$ of a few keV which dominated the XRF samples of the {\\it Beppo}SAX and {\\it HETE-2} \\citep[e.g.,][]{kippen,sakamoto}.  Nonetheless, {\\it Swift} has an unique capability for studying the detailed X-ray afterglow properties just after the burst for XRFs and XRRs with $\\eop$ $\\gtrsim$ 20 keV for the first time.  The systematic study of the X-ray emissions of GRBs observed by XRT reveals a very complex power-law decay behavior consisting typically of an initial very steep decay (t$^{\\alpha}$ with $-10 \\lesssim \\alpha_{1} \\lesssim -2$) \\citep[e.g.,][]{obrien2006,sakamoto2007}, followed by a shallow decay ($-1 \\lesssim \\alpha_{2} \\lesssim 0$), followed by a steeper decay ($-2 \\lesssim \\alpha_{3} \\lesssim -1$) \\citep[e.g.,][]{nousek2006, obrien2006,willingale2007}, sometimes followed by a much steeper decay ($\\alpha_{4} \\lesssim -2$) \\citep[e.g.,][]{willingale2007} and, in some cases (about 50\\%), overlaid X-ray flares \\citep[e.g.][]{burrows2005,chincarini2007,kocevski2007}.  Although there is increasing evidence that the initial very steep decay component $\\alpha_{1}$ is a tail of the GRB prompt emission \\citep[e.g.,][]{liang2006,sakamoto2007}, the origin of the phase from a shallow $\\alpha_{2}$ to a steeper decay $\\alpha_{3}$ (hereafter a shallow-to-steep decay) is still a mystery.  Moreover, not all GRBs have a shallow-to-steep decay phase in their X-ray afterglow light curves.  Thus, it is very important to investigate the X-ray afterglow light curves of bursts along with their prompt emission properties to find a difference in their characteristics between C-GRBs and XRFs.  In this paper, we report the systematic study of the prompt and afterglow emission of 10 XRFs and 17 XRRs observed by {\\it Swift} from December 2004 through September 2006. Although the data from {\\it Swift} BAT is the primary dataset for investigation of the prompt emission properties, we also use information from $Konus$-$Wind$ and {\\it HETE-2} as reported on the Gamma-ray burst Coordinate Network\\footnote{http://gcn.gsfc.nasa.gov/gcn\\_main.html} or in the literature, when available, to obtain better constraints on $\\eop$.  We focus on X-ray afterglow properties observed by {\\it Swift} XRT in this study.  In \\S 2, we discuss our classification of GRBs, the analysis methods of the BAT and the XRT data, and the selection criteria of our sample. In \\S 3 and \\S4, we show the results of the prompt emission and the X-ray afterglow analysis, respectively.  We found distinct differences between XRFs and C-GRBs in the shape and in the overall luminosity of X-ray afterglows. We discuss the implications of our results in \\S 5.  Our conclusions are summarized in \\S 6. We used the cosmological parameters of $\\Omega_{\\rm m}$ = 0.3, $\\Omega_{\\rm \\Lambda}$ = 0.7, and $H_{0}$ = 70 km s$^{-1}$ Mpc$^{-1}$. The quoted errors are at the 90\\% confidence level unless we state otherwise. ",
        "conclusions": "We have seen that the XRFs observed by {\\it Swift} form a continuum with the C-GRBs observed by {\\it Swift} and by other missions, having systematically lower fluences and lower $\\eop$ than C-GRBs.  We have noted that the X-ray light curves of XRFs tend to follow a different ``template'' than those of C-GRBs.  The light curves of the C-GRB afterglows show a break to steeper indices (shallow-to-steep decay) at earlier times, whereas XRF afterglows show no such break. This break is evident in the X-ray but not in the optical light curve.  Moreover, the overall luminosity of XRF X-ray afterglows is smaller by a factor of two or more compared to that of C-GRBs. These distinct differences in the X-ray afterglow between XRFs and C-GRBs are keys to understanding not only the shallow-to-steep decay phase in the X-ray afterglow but also the nature of XRFs in an unified picture.  We have discussed the geometrical jet models based on the trend which we found that the shallow-to-steep break in the X-ray afterglow preferentially is seen in the C-GRB sample. We concluded that none of the jet models can explain the behavior of a shallow-to-steep decay phase observed only in the X-ray afterglow.  We also emphasize the importance of having simultaneous X-ray and optical afterglow observations along with the characteristics of the prompt emission such as $\\eop$ to constrain the various geometrical jet models."
    },
    "0801/0801.1016_arXiv.txt": {
        "abstract": "In this paper we review the possible radiation mechanisms for the observed non-thermal emission in clusters of galaxies,  with a primary focus on the radio and hard X-ray emission. We show that the difficulty with the non-thermal, non-relativisitic Bremsstrahlung model for the hard X-ray emission, first pointed out by \\citet{Petrosian2001} using a cold target approximation, is somewhat alleviated when one treats the problem more exactly by including the fact that the background plasma particle energies are on average a factor of 10 below the energy of the non-thermal particles. This increases the lifetime of the non-thermal particles, and as a result decreases the extreme energy requirement, but at most by a factor of three. We then review the synchrotron and so-called inverse Compton emission by relativistic electrons, which when compared with observations can constrain the value of the magnetic field and energy of relativistic electrons. This model requires a low value of the magnetic field which is far from the equipartition value. We briefly review the possibilities of gamma-ray emission and prospects for {\\sl GLAST} observations. We also present a toy model of the non-thermal electron spectra that are produced by the acceleration mechanisms discussed in an accompanying paper. ",
        "introduction": "\\label{intro}  The observed signatures of the non-thermal (NT) activity in the intra-cluster medium (ICM) were described in details  by \\citealt{Rephaeli2008} - Chapter 5, this volume and  \\citealt{Ferrari2008} - Chapter 6, this volume.  Here we give a brief summary. The first and least controversial radiative signature comes from radio observations. Synchrotron emission by a population of relativistic electrons is the only possible model for the production of this radiation. In the case of the Coma cluster, the radio spectrum may be represented by a broken power law \\citep{Rephaeli1979}, or a power law with a rapid steepening \\citep{Thierbach2003}, or with an exponential cutoff \\citep{Schlickeiser1987}, implying the presence of electrons with similar spectra. Unfortunately, from radio observations alone one cannot determine the energy of the electrons or the strength of the magnetic field. Additional observations or assumptions are required. Equipartition or minimum total (relativistic particles plus fields) energy arguments imply a population of relativistic electrons with Lorentz factor $\\gamma\\sim10^4$ and magnetic field strength of $B\\sim\\mu$G, in rough agreement with the Faraday rotation measurements \\citep[e.g.][]{Kim1990}. \\citet{Rephaeli1979} and \\citet{Schlickeiser1987} also pointed out that the electrons responsible for the radio emission, should also produce a spectrum of hard X-ray (HXR) photons  (similar to that observed in the radio band),  via inverse Compton (IC) scattering of the Cosmic Microwave Background (CMB) photons. This emission  is estimated to be the dominant emission component around 50 keV. Detection of HXR radiation could break the degeneracy and allow determination of the magnetic field and the energy of the radiating electrons. In fact, because the energy density of the CMB radiation (temperature $T_0$) $u_{\\rm CMB} = 4\\times 10^{-13}(T_0/2.8\\,{\\rm K})^4$~erg\\,cm$^{-3}$ is larger than the magnetic energy density $u_B = 3\\times 10^{-14}(B/\\mu{\\rm G})^2$~erg\\,cm$^{-3}$, one expects a higher flux of HXR than radio radiation.  As already described in the above mentioned papers by Rephaeli et al. and Ferrari et al., recently there has been growing evidence for this and other signatures of the NT activity. Excess HXR and extreme ultraviolet (EUV) radiation are observed at the high and low ends of the usual soft X-ray (SXR) thermal Bremsstrahlung (TB) radiation. Fig.~1 shows all the flux $\\nu F (\\nu)$ (or equivalently the energy density $\\nu u(\\nu) = 4\\nu F (\\nu)/c)$ of the above mentioned and other radiation for the Coma cluster. However, for the excess radiation not only the exact mechanisms are controversial but even their NT nature is questioned. The observed spectra of the excess radiation often can be fit by thermal spectra with  higher and lower temperatures than that needed for the SXR observations with almost the same confidence as with a NT power law. The most natural NT process for these excesses (specially for HXRs) is the IC scattering of the CMB photons. However, the relatively high observed HXR fluxes require a large number of relativistic electrons, and consequently a relatively low magnetic field for a given observed radio flux. For Coma, this requires the (volume averaged) magnetic field to be ${\\bar B}\\sim 0.1-0.3$~$\\mu$G, while equipartition gives ${\\bar B}\\sim 0.4$~$\\mu$G and Faraday rotation measurements give the (average line-of-sight) field of ${\\bar B}\\sim 3$~$\\mu$G \\citep{Giovannini1993,Kim1990,Clarke2001,Clarke2003}. In general the Faraday rotation measurements of most clusters give $B > 1$~$\\mu$G; see e.g. \\citet{Govoni2003}. Because of this apparent difficulty, various authors (see, e.g. \\citealt{Ensslin1999,Blasi2000}) suggested that the HXR radiation is due to non-thermal Bremsstrahlung by a second population of NT electrons with a power law distribution in the 10 to 100 keV range. In what follows we examine the merits and shortcomings of the mechanisms proposed to interpret these observations. We first consider the EUV observations briefly and then address the thermal and NT (IC and non-thermal Bremsstrahlung) radiation model for the HXR observations. ",
        "conclusions": "In this paper we address the processes that produce radio, EUV and HXR radiation in the ICM. First we consider NT Bremsstrahlung as the source of HXRs which, as shown earlier (P01), faces the difficulty of its low yield compared to Coulomb losses. We describe results from a more detailed analysis of the lifetimes of NT electron tails (or bumps) in a hot ICM than that presented in P01, where cold target loss rates were used. We find that the lifetimes of NT tails is increased by a factor of $<3$ so that the above difficulty becomes less severe but production of HXRs via NT Bremsstrahlung remains problematical. Next we discuss the expected radiative signature of relativistic electrons and show that radio and HXR observations can be explained by synchrotron and IC scattering of CMB photons. But now one requires a low magnetic field which is far from equipartition with the electrons. Finally we give a rough estimate of the gamma-ray signature of the relativistic electrons and point out several possible scenarios in which the gamma-ray fluxes might exceed the {\\sl GLAST} threshold. Based on these results we present an average spectrum of electrons that is required and possible extensions of it into the low energy regime. Production of these spectra is discussed by \\citealt{Petrosian2008} - Chapter 11, this volume."
    },
    "0801/0801.3822_arXiv.txt": {
        "abstract": "We present an empirical method for estimating the underlying redshift distribution $N(z)$ of galaxy photometric samples from photometric observables. The method does not rely on photometric redshift (photo-z) estimates for individual galaxies, which typically suffer from biases. Instead, it assigns weights to galaxies in a spectroscopic subsample such that the weighted distributions of photometric observables (e.g., multi-band magnitudes) match the corresponding distributions for the photometric sample. The weights are estimated using a nearest-neighbor technique that ensures stability in sparsely populated regions of color-magnitude space. The derived weights are then summed in redshift bins to create the redshift distribution. We apply this weighting technique to data from the Sloan Digital Sky Survey as well as to mock catalogs for the Dark Energy Survey, and compare the results to those from the estimation of photo-z's derived by a neural network algorithm. We find that the weighting method accurately recovers the underlying redshift distribution, typically better than the photo-z reconstruction, provided the spectroscopic subsample spans the range of photometric observables covered by the photometric sample. ",
        "introduction": "\\label{sec:int}   On-going, wide-field surveys are delivering photometric galaxy samples of unprecedented scale. Optical and near-infrared surveys planned for the next decade will increase the sizes of such samples by an order of magnitude. Much of the utility of these samples for astronomical and cosmological studies rests on knowledge of the redshift distributions of the galaxies they contain. For example, surveys aimed at probing dark energy via clusters, weak lensing, and baryon acoustic oscillations (BAO) will rely on the ability to coarsely bin galaxies by redshift, enabling approximate distance-redshift measurements as well as study of the growth of density perturbations. The power of these surveys to constrain cosmological parameters will be limited in part by the accuracy with which the galaxy redshift distributions can be determined \\citep{HutKimKraBro04,HutTakBerJai06,ZhaKno06,Zha06,MaHuHut06,LimHu07}.  Photometric redshifts (photo-z's, denoted $z_{\\rm phot}$ below) -- approximate estimates of galaxy redshifts based on their broad-band photometric observables, e.g., magnitudes or colors -- offer one technique for approaching this problem. Photo-z's have the advantage that they provide redshift estimates for each galaxy in a photometric catalog; such information is useful for certain studies \\citep{Manetal07}. However, in many applications we do not need such galaxy-by-galaxy information -- instead, we only require an estimate of the redshift distribution of a sample of galaxies selected by some set of photometric observables. For example, cosmic shear weak lensing or angular BAO measurements rely on relatively coarse binning of galaxies in redshift, and it suffices to have an accurate estimate of the redshift distribution $N(z)$ for galaxies satisfying certain color or magnitude selection criteria \\citep{she04,Sheetal07a,jain07}. Photo-z estimators are {\\it not} typically designed to provide unbiased estimates of the redshift distribution: $N(z_{\\rm phot})$ is biased by photo-z errors.   Although deconvolution \\citep{pad05} or other techniques \\citep{Sheth07} can be used to obtain improved estimates of the redshift distribution from photo-z measurements, this problem motivates the development of a method optimized to directly estimate the underlying redshift distribution $N(z)$ for a photometric sample. In addition to its direct utility, a precise, unbiased estimate of the redshift distribution is useful even for probes that do require individual galaxy redshifts, since it provides a template for characterizing photo-z errors.  In this paper we present an empirical technique to estimate $N(z)$ for a photometric galaxy sample that is based upon matching the distributions of photometric observables of a spectroscopic subsample to those of the photometric sample. The method assigns weights to galaxies in the spectroscopic subsample (hereafter denoted the training set, in analogy with machine-learning methods of photo-z estimation), so that the weighted distributions of observables for these galaxies match those of the photometric sample. The weight for each training-set galaxy is computed by comparing the local ``density'' of training-set galaxies in the multi-dimensional space of photometric observables to the density of the photometric sample in the same region. We estimate the densities using a nearest neighbor approach that ensures the density estimate is both local and stable in sparsely occupied regions of the space. The use of the nearest neighbors ensures optimal binning of the data, which minimizes the requisite size of the spectroscopic sample. After the training-set galaxy weights are derived, we sum them in redshift bins to estimate the redshift distribution.  As we will show, this method provides a precise and nearly unbiased estimate of the underlying redshift distribution for a photometric sample and does not require photo-z estimates for individual galaxies. Moreover, the spectroscopic training set does {\\it not} have to be representative of the photometric sample, in its distributions of magnitudes, colors, or redshift, for the method to work. We only require that the spectroscopic training set {\\it cover}, even sparsely, the range of photometric observables spanned by the photometric sample. The method can be applied to different combinations of photometric observables that correlate with redshift -- in this paper, we confine our analysis to magnitudes and colors. In a companion paper (Cunha et al., in preparation), we compare this weighting technique to the deconvolution method of \\cite{pad05} and show that the weights can be used to naturally regularize and improve the deconvolution.   The paper is organized as follows. In \\S~\\ref{sec:cat} we present the simulated and real galaxy catalogs used to test the method. In \\S~\\ref{sec:wei} we describe the algorithm for the calculation of the weights of training-set galaxies using a nearest neighbor method. In \\S~\\ref{sec:qua} we define simple statistics to assess the quality of the reconstructed distributions. In \\S~\\ref{sec:res} we present $N(z)$ estimates derived from the weighting method and compare with results using photo-z's for individual galaxies derived from a neural-network algorithm. We discuss the results and present our conclusions and perspectives in \\S~\\ref{sec:dis}. In Appendix~\\ref{app:neu}, we provide a brief description of the neural network photo-z algorithm that we use for comparison with the weighting method. ",
        "conclusions": "\\label{sec:dis}  We have presented a new technique to estimate the underlying redshift distribution of photometric galaxy samples. The method relies on a spectroscopic training set and re-weighting of the training-set galaxies to match the distribution of photometric observables of the photometric sample. The weights are estimated using a flexible nearest-neighbor approach in color-magnitude space and the redshift distribution is estimated by summing the galaxy weights in redshift bins. Tests on mock catalogs and on existing data sets show that this procedure yields an accurate estimate of the redshift distribution and that it performs significantly better than simply binning the photo-z estimates of individual galaxies in the photometric sample. The weighting method also appears to be robust, in the sense that the spectroscopic sample can have very different distributions of photometric observables and redshift from the photometric sample. The main requirement is that the training set should span the range of photometric observables found in the photometric sample.  The key assumption underlying the technique is that two samples (e.g., the spectroscopic and photometric) with the same distribution of photometric observables will have very similar redshift distributions. This assumption holds if the selection criteria used to define the two samples differ only in the space of photometric observables. Several effects can cause this condition to be violated: statistical errors, LSS and spectroscopic failures.  Statistical errors are the simplest to quantify and are significant in regions of magnitude space where the training set is sparse, typically at fainter magnitudes. LSS can be significant if certain regions of the space of photometric observables are only represented in the training set by a survey that covers small solid angle, in which one or a few large structures dominate. We showed that, even in such cases, the weighting method works quite well (\\S~\\ref{subsec:SDSSres}).  Spectroscopic failures (i.e., targeted objects for which a redshift could not be obtained) in the training set can have a similar effect if the failures happen systematically, for instance in a particular galaxy spectral type. If the effects of spectroscopic failures are prevalent in regions of magnitude space where the redshift distribution is broad or multiply peaked, they can potentially cause systematic errors in the recovery of the redshift distribution. However we also showed that the weighting method performs well even when we take an arbitrary type distribution in the training set (\\S~\\ref{subsec:DESres}).  The weighting method requires a training set with a size (density) such that the inter-(training-set)-galaxy separation in the space of photometric observables is comparable to the characteristic (curvature) scale of the redshift/photometric-observables manifold or the scale defined by the typical photometric errors - whichever is smaller. That condition ensures that on average at least one neighbor to the galaxy is meaningful; in practice it would be safer to have the density a few times larger than this minimum density. The use of mock catalogs can shed light on the optimal parameters to employ on the weighting method, such as the number of neighbors $N_{\\rm nei}$ (possibly varying according to the local density), the minimum training set size, the need or not for renormalization, etc. Since these simulations are necessary for other typical calibration reasons, their need does not put any strong restrictions to the application of the weighting method. For instance, simulations and calibration samples are necessary to calibrate the photo-z errors.  This weighting technique has been used to estimate the redshift distribution of the SDSS DR6 photometric sample~\\citep{OyaLimCunLinFriShe08} and to help assess the quality of the photo-z's computed for that sample. It has also been used in conjunction with photo-z's in the measurement of the SDSS cluster-mass cross-correlation function via weak lensing~\\citep{Sheetal07a}, allowing for the inversion of cluster mass profiles~\\citep{Johetal07} and estimation of cluster mass-richness relations~\\citep{Johetal07} and mass-to-light ratios~\\citep{Sheetal07b}. Finally, this weighting scheme has recently been employed in the study of galaxy-galaxy weak lensing calibration bias~\\citep{Manetal07}, where it was shown to yield much smaller biases than those arising from photo-z estimates. For future photometric surveys, the weighting method can complement and provide cross-checks on photo-z estimates and help control photo-z errors."
    },
    "0801/0801.2108_arXiv.txt": {
        "abstract": "We investigate the orbital evolution of planetesimals in a self-gravitating circumstellar disc in the size regime ($\\sim 1-5000$ km) where the planetesimals behave approximately as test particles in the disc's non-axisymmetric potential. We find that the particles respond to the stochastic, regenerative spiral features in the disc by executing large random excursions (up to a factor of two in radius in $\\sim 1000$ years), although typical random orbital velocities are of order one tenth of the Keplerian speed. The limited time frame and small number of planetesimals modeled does not permit us to discern any {\\it net} direction of planetesimal migration. Our chief conclusion is that the high eccentricities  ($\\sim 0.1$) induced by interaction with spiral features in the disc is likely to be highly unfavourable to the collisional growth of planetesimals in this size range  while the disc is in the self-gravitating regime. Thus {\\it if}, as recently argued by Rice et al 2004, 2006, the production of planetesimals gets under way when the disc is in the self-gravitating regime (either at smaller planetesimal size scales, where gas drag is important, or via gravitational fragmentation of the solid component), then the planetesimals thus produced would not be able to grow collisionally until the disc ceased to be self-gravitating. It is unclear, however, given the large amplitude excursions undergone by planetesimals in the self-gravitating disc, whether they would be retained in the disc throughout this period, or whether they would instead be lost to the central star. ",
        "introduction": "It is currently unclear whether the process of planet formation begins during the earliest phases of star formation (i.e. in circumstellar discs that are strongly self-gravitating). The high disc mass at this stage (with gas mass $\\sim 10 \\%$ of the central star's mass (e.g. Eisner and Carpenter 2006) is favourable to this hypothesis and provides the only opportunity for gas giants to form via (gas phase) gravitational instability \\citep{boss00}. A more subtle effect - again strongly favoured during the self-gravitating case - arises through the interplay between pressure gradients in spiral arms and gas drag on solid particles and results in the accumulation of dust in spiral arms \\citep{rice04}. This dust may then be able to accumulate further through the action of collisions and/or self-gravity in the dust phase \\citep{rice06}. \\footnote{Note that such rapid accumulation, which might result in the rapid formation of large size planetesimals, is required in order for the growing planetesimals to escape the dangerous meter-size range, where they are subject to very fast migration down into the hottest inner disc \\citep{weiden77}.}  Once the disc is no longer self-gravitating in the dominant (gas) component, it is evidently impossible to form planets through either gas phase Jeans instability or dust accumulation in spiral arms and thus the most successful models invoke the collisional growth of dust \\citep{pollack96} followed possibly by the accretion of a gaseous envelope. Despite the lower surface densities during later phases (estimates of gas disc mass from sub-millimetre dust measurement suggest typical masses that are around an order of magnitude lower than in the self-gravitating case; \\citealt{andrews05}), the great advantage for planet formation is simply the longer duration of this later phase. Discs with sufficient gas and dust to form planets typically survive for $10^6-10^7$ years \\citep{haisch01,armitage03,wyatt05}, whereas the existence of dusty debris discs  around older stars suggests that (potentially planet building) collisions between rocky planetesimals proceeds for $10^8$ years or more \\citep{bromley04}. These numbers should be contrasted with the brief window of opportunity ($\\sim 10^5$ years) in the self-gravitating phase.  In summary, then, it is unclear without detailed calculation which of these regimes is favoured in the trade-off between lifetime and disc mass.  The viability of planet formation in the self-gravitating regime however requires that a number of conditions are satisfied beyond the initial phases of accumulation outlined above. In particular, we need to understand the dynamical evolution of the protoplanetary components (whether at the scale of dust aggregates, planetesimals or proto-giant planets) with a view to answering the following questions: (i) does orbital migration result in such bodies spiralling into the central star, or can they be retained in the disc at least over the initial self-gravitating phase, and (ii) (in the case of dust or planetesimals) are the kinematics of such bodies conducive to further collisional growth?  As a first step in answering this question we here undertake pilot simulations of a small number of test particles in a self-gravitating disc, with a view to both considering the issue of their orbital migration and  to establishing the velocity dispersion that is set up through interaction with spiral structure in the disc. It should be stressed that the calculation differs in several important  ways from most calculations of orbital migration and planetesimal kinematics  reported in the literature. Firstly, conventional planetary migration, whether classified as Type I or Type II, is the consequence of the gravitational torque exerted on the planet (or planetesimal)  by non-axisymmetric structure {\\it induced by the planet/planetesimals}. Since this torque contribution can have significant contributions at the size scale of the Hill radius (or below: \\citealt{bate03}), it is evidently important that this scale is well resolved in numerical calculations, which makes the modeling of the migration of low mass objects particularly challenging, especially with codes such as Smoothed Particles Hydrodynamics (SPH) \\citep{valborro06}. In the present case, however, we are interested in  a regime where the disc has pronounced non-axisymmetric structure {\\it in the absence of the planet/planetesimal} (due to the presence of self-gravitating spiral modes in the disc) and we are therefore not obliged to resolve the Hill radius in order to capture the dominant torque contribution.  Evidently, as the mass of the planetesimal is increased, the contribution to the torque which it experiences from self-induced spiral structure becomes significant and we will assess this {\\it a posteriori} in order to set an upper mass scale on the applicability of our calculations. Secondly, in conventional calculations of the kinematics of planetesimal swarms \\citep{kokubo2000,thommes03}, the velocity dispersion of planetesimals is set by an equilibrium between what is sometimes termed 'viscous stirring' (i.e. the transfer of kinetic energy from larger bodies to the planetesimal swarm via two-body scattering) and eccentricity damping by gas drag. In the present case, however, the velocity dispersion of our  particles is set by the level to which eccentricity is pumped through interaction with the fluctuating potential of the self-gravitating disc. In this pilot calculation we omit gas drag, which places a lower limit to  the size scale of objects to which the calculation is applicable; again, we set this lower limit {\\it a posteriori}.  Although we have stressed the qualitative differences between this calculation and that reported in the large body of literature on planetary migration and planetesimal kinematics, we point out that it bears close comparison with the calculation of \\citep{nelson05}, which studies the response of planetesimals to the fluctuating non-axisymmetric structure in a disc that is subject to the magneto-rotational (MRI) instability. In this study it was found that the planetesimals underwent stochastic migration, but that it was not possible to discern the mean magnitude (or even sign) of the migration rate over the limited timeframe of the simulation (a few hundred orbits). In our experiment also, conducted over a similar time frame to Nelson (2005), we are unable to discern a net direction of migration, but we find (perhaps unsurprisingly given the fact that our disc is about an order of magnitude larger in mass) that the amplitude of the stochastic variations in the orbital radius is much larger, with several planetesimals undergoing radial excursions of a factor two or more in both directions. Evidently, this behaviour has important implications for the ability of planetesimals to grow by collisions during the self-gravitating phase.  The structure of the paper is as follows. In Section 2 we describe the modeling of the self-gravitating disc, which is maintained in a 'self-regulated' (i.e. non fragmenting) state through the imposition of an {\\it ad hoc} cooling law for which the cooling time is a suitably large multiple of the local dynamical time. In Section 3  we analyse the orbital response of $12$ test particles placed in a disc of mass  $0.1$ and $0.5 M_\\odot$. We also analyse the location of the dominant torque contribution on the planetesimals and assess the range of planetesimal mass and size scales for which we expect 'test particle' behaviour and for which gas drag can be ignored. In section 4 we discuss the implications of our numerical findings for the migration and collisional growth of planetesimals in self-gravitating discs. Section 5 contains our chief conclusions. ",
        "conclusions": "We can summarise our numerical findings as follows. Planetesimals in the size/mass range to which our  calculations apply (see equation \\ref{e:valid}) are subject to large scale fluctuations in orbital radius due to interaction with regenerative spiral structure in the disc. We can discern no net sign to the mean migration rate over the timescale of the simulation and are thus unable to answer one of the goals of the present study, i.e. whether we expect planetesimals to migrate into the star during the self-gravitating phase of the disc. Since the disc's lifetime in the self-gravitating phase exceeds the duration of the present experiment by about a factor $30$ it will be computationally challenging to answer this question, especially since (given the stochastic nature of the orbital evolution) it will be necessary to assess the fraction of planetesimals retained in the disc through following  the evolution of a large ensemble of particles.  We can however make some comments about how the dynamical evolution of our particles is likely to affect their collisonal growth during the self-gravitating phase.  As noted by Nelson (2005) in the context of the evolution of planetesimals in a disc subject to MRI, the fluctuations in orbital radius and eccentricity could in principle either aid or hinder planetesimal growth. On the positive side, orbital migration provides a mechanism for avoiding the self-limitation of planetesimal growth through the clearing out of a proto-planet's ``feeding zone\". In conventional models, this imposes an ``isolation mass\" limit which makes it difficult to grow beyond a  terrestrial mass scale in the inner regions of the disc \\citep{ida04}. Although this limitation is by-passed in our self-gravitating models, since the feeding zone would be continuously replenished by planetesimals undergoing an orbital random walk, this factor is not very significant since in any case the high surface density during the self-gravitating phase means that isolation masses are high (higher, that is, than the upper limit for the applicability of our calculations). Another potentially positive aspect of evolution in a self-gravitating disc is the possibility of concentrating planetesimals in high density spiral features, where they might be more prone to undergo collisional growth. In the case of the more massive -- $0.5 M_\\odot$ -- disc (see Section 4), there is some evidence for the dominance at times of spiral modes which persist over a number of dynamical times, allowing particles to settle into orbits which conserve a Jacobi integral in the co-rotating frame. Some of these orbits involve a preference for particles to spend time in the spiral arms, an impression strengthened by inspection of animations which show particles which co-rotate for a while in the vicinity of an  arm and which then respond to changes in the spiral structure by passing  rapidly to another spiral feature. This tendency will not lead to a significantly enhanced collision rate, however, unless there is a strong enhancement in the {\\it solid to gas ratio} in the spiral arms: since the amplitude of the gaseous arms is around a factor two at most, then even a particle that spent {\\it all} its time in the arms would only, at fixed solid to dust ratio, experience an enhancement in the local planetesimal density of a factor two at most. Since we have only studied the evolution of ten planetesimals, we cannot, on the basis of the present study, comment on the degree to which planetesimals are differentially concentrated in the arms. Nevertheless, the results of \\citet{rice04} suggest that the differential accumulation of solids in the arms is weak at  size scales where gas drag is unimportant.  We now turn to the negative implication of stochastic migration noted by Nelson (2005), i.e. the growth of planetesimal velocity dispersion. This hinders planetesimal growth in two respects: higher collisional velocities are associated with disruption of parent bodies instead of collisional growth (e.g. Leinhardt et al 2007) and also suppress the role of  gravitational focusing in effecting physical collisions. Quantifying the second effect, the timescale for collisional growth of a planetesimal of mass $M_{\\rm p}$, density $\\rho$ and  radius $R_{\\rm p}$ is given by: \\begin{equation} t_{\\rm grow} =  \\frac{ \\rho^{2/3} M_{\\rm p}^{1/3}}{\\Sigma_p \\Omega (1+4GM_{\\rm p}/(R_{\\rm p} \\sigma^2))} \\end{equation} where $\\Sigma_{\\rm p}$ is the surface density of planetesimals, $\\sigma$ is the planetesimal velocity dispersion and $\\Omega$ is the local Keplerian orbital frequency. For a planetesimal of density $\\sim 3$ g/cm$^3$  located at $10$ A.U. in a disc with surface density of planetesimals  equal to 0.4 g/cm$^ {2}$, the growth timescale via two body collisions is: \\begin{equation} t_{\\rm grow} = 5\\times 10^9 \\frac{\\tilde M^{1/3}}{1+\\tilde M^{2/3}\\tilde\\sigma^{-2}}~~\\mbox{yrs} \\end{equation} where $\\tilde M=(M_p/6 ~10^{-4}M_{\\oplus})$ and $\\tilde \\sigma=(\\sigma/1 \\mbox{km s}^{-1})$. Evidently this timescale is far too long to be of interest unless the second term on the denominator (due to enhancement of the collisional cross section by gravitational focusing) is large. {\\it If} we simply set $\\sigma$ to be the product of the orbital eccentricity and local Keplerian velocity (an assumption we revisit below),  it follows that at a given location in the disc and for a given mass scale of planetesimal, then the growth timescale in the gravitationally focused regime scales as the {\\it square of the eccentricity} divided by the surface density normalisation. In other words, for a planetesimal at a given mass scale to be able to grow significantly over a disc phase of duration $t_{\\rm life}$ we require \\begin{equation} \\frac  {t_{\\rm grow}} {t_{\\rm life}} < 1 \\end{equation} which translates into a lower limit on the product $t_{\\rm life} \\Sigma e^{-2}$ (where $\\Sigma$ is the surface density normalisation of the disc).  We can now return to the issue raised in the Introduction of whether or not planetesimal growth is favoured in the high density but short lived phase during which the disc is self-gravitating. Assuming for simplicity that the self-gravitating phase lasts $\\sim 10 \\%$ of the subsequent lifetime of the primordial gas/dust disc but that its surface density (mass) is typically ten times greater, we see that these two factors roughly cancel. Thus whether the growth of planetesimals is favoured during the self-gravitating phase actually boils down to the relative values of the eccentricity in the two phases.  Our simulations have shown that this factor is decisive in disfavouring the self-gravitating regime. In conventional models for the dynamical evolution of planetesimal swarms (e.g. \\citealt{kokubo2000}) the equilibrium eccentricity (attained through the balance between stirring by larger bodies and damping by gas drag) is very low (of order $0.01$). We have however seen that in the self-gravitating phase, the typical eccentricity is $0.1$ or more. Since the growth time scales as the square of the eccentricity, this means that the self-gravitating phase is highly unfavourable (compared with the later more quiescent phase) for the collisional growth of planetesimals. We note that the above estimates for conventional proto-planetary discs assume a laminar disc structure;  eccentricity driving by interaction with turbulent structures in a disc exhibiting MRI would further hinder collisional growth, whether or not the disc is self-gravitating.   There however remain two possible qualifiers of the above conclusion. Firstly,  it is only valid to set the local velocity dispersion equal to the product of the eccentricity and local Keplerian velocity in the case that the kinematics is well described as a set of elliptical orbits which are elongated in a random direction. Such a description does not of course allow for the possibility of {\\it local velocity coherence} in a planetesimal swarm in which velocities are significantly non-circular. In this sense, therefore, the above estimates are pessimistic. Further calculations, involving a large swarm of planetesimals from which one could derive the {\\it local} velocity dispersion, are required in order to settle this issue. We however note that our difficulty in identifying coherent orbital families, interacting with a long lived spiral mode, probably implies that this effect is unlikely to improve the prospects for particle growth significantly. The second possible qualifier is that, as noted in Section 2, our simulations employ a cooling timescale which is close to the critical one for disc fragmentation and therefore represent conditions where spiral modes are of relatively high amplitude. In the inner regions of protoplanetary discs, however, cooling timescales are longer than employed here and thus one would expect lower amplitude spiral structures, with a correspondingly lower amplitude of stochastic driving of planetesimal orbits. This possibility needs to be quantified through further numerical investigation.  If one leaves aside these possibilities, then we have shown that the self-gravitating phase is unfavourable for collisional growth in the regime of planetesimal scale studied. We have not ruled out the possibility of collisional growth at smaller scales (where gas drag can effect a significant enhancement of the local solid to gas ratio in spiral features; \\citealt{rice04}) nor that (where such enhancement exists) it may not be possible to produce much larger bodies through the gravitational fragmentation of the solid component (e.g., \\citealt{rice06}). We nevertheless consider it likely that whatever the size scale attained in this way, there will be no further collisional growth of planetesimals while the disc is self-gravitating.  This obviously limits the extent to which planet formation can get under way in the self-gravitating phase. Either a massive planet is produced via gas phase Jeans instability \\citep{boss00}, thus obviating the need for collisional growth, or else, if self-gravity aids the initial accumulation of solid bodies \\citep{rice04,rice06},  then it is unlikely that such bodies can undergo further collisional growth once they become large enough for gas drag to be unimportant. In this latter case, further growth of such bodies could only resume once the disc had evolved to the point where its self-gravity (and the associated pumping up of the planetesimal velocity dispersion) has become unimportant. The vital unanswered question, therefore, is whether orbital migration would prevent the planetesimals from remaining in the disc throughout the self-gravitating phase. This question can only be answered through further, computationally challenging, calculations."
    },
    "0801/0801.2278_arXiv.txt": {
        "abstract": "We present photometric and spectroscopic data of the peculiar SN~2005la, an object which shows an optical light curve with some luminosity fluctuations and spectra with comparably strong narrow hydrogen and helium lines, probably of circumstellar nature. The increasing full-width-half-maximum velocity of these lines is indicative of an acceleration of the circumstellar material. SN~2005la exhibits hybrid properties, sharing some similarities with both type IIn supernovae and 2006jc-like (type Ibn) events. We propose that the progenitor of SN~2005la was a very young Wolf-Rayet (WN-type) star which experimented mass ejection episodes shortly before core collapse. ",
        "introduction": "\\label{intro}   The interest of the astronomical community has been recently caught by a few very luminous supernovae (SNe) that have exploded within dense, massive and pre-existent circumstellar media. Some of them have been spectroscopically classified as type IIn supernovae \\citep[SNe IIn,][]{sch90} because of the presence of prominent narrow (about 1000 km s$^{-1}$) H Balmer emission lines (e.g. SN 2006gy, \\citet{smi06,ofek07}; SN 2007bw, \\citet{NSF07}). For few other events (e.g. SN 1999cq, \\citet{mat00}; SN 2006jc, \\citet{fol07,pasto07,pasto07b}), whose spectra show relatively narrow He I emission lines ($\\sim$ 2000 km s$^{-1}$), a different nomenclature was coined: type Ibn SNe \\citep{pasto07,pasto07b}. In the spectra of these objects, the narrow H lines are much weaker than the He I features, or completely missing.  In addition, a number of regular type IIn SNe have developed at some stages clear evidence of narrow (from a few $\\times$100 km s$^{-1}$ to $\\sim$2000 km s$^{-1}$) He I lines although, contrary to SNe Ibn, they are always less prominent than H$\\alpha$. This group includes SN 1978K \\citep{chu95}, SN 1987B \\citep{sch96}, SN 1988Z \\citep{sta91,tur93,are99}, SN 1995G \\citep{pasto02}, SN 1995N \\citep{fran02,pasto05,zam05}, SN 1996L \\citep{ben99}, SN 1997ab \\citep{heg97}, SN 1997eg \\citep{sal02}, SN 1998S \\citep{leo00,ger00,liu00,fas01,anu01}. In all these SN events, the narrow lines (both of H and He I) are thought to originate in a slowly-moving circumstellar medium (CSM) which is photoionized either by early UV flux or interaction with the SN ejecta.  A number of authors \\citep{kot06,gal07,smi06} suggest that a significant fraction of SNe IIn might be produced by the explosion of massive Luminous Blue Variable (LBV) stars, even if the nature of at least some of type IIn events is still debated \\citep[e.g. SN~2002ic and similar events, see][and references therein]{ham03,ben06,pri07a,tru08}. As an alternative view, \\cite{woo07} propose a pair-instability pulsation scenario as a trigger of major mass ejection episodes. The ejecta collide with the material expelled by the star in previous eruptions converting large amounts of kinetic energy into radiation and generating a Type IIn SN.  Type Ibn SNe seem to be connected with WR progenitors whose He envelope is almost completely stripped away, and that explode while they are still embedded within the He-rich material \\citep{pasto07,pasto07b,fol07,smi07,tom07,imm07}. \\citet{pasto07} showed that SN 2006jc was preceded by a giant outburst 2 years before core-collapse. Although the SN itself suggests a WC star progenitor, the outburst had similar characteristics to an LBV eruption.  In \\citet[][paper 1 of this series]{pasto07b} the common properties of objects similar to SN~2006jc (SNe Ibn) are discussed. In paper 3 \\citep{seppo07} we investigate the origin of the strong infrared excess observed in SN 2006jc. In this paper (paper 2) we describe the observed characteristics of the unusual SN~2005la, an object whose spectra show simultaneously narrow H and He I features, with the most prominent He I lines having comparable strength (at least at early stages) as H$\\alpha$, and showing unusual evolution with time. We propose that SN~2005la is a transitional object lying somewhere between the newly defined SNe Ibn and more classical type IIn events. ",
        "conclusions": ""
    },
    "0801/0801.3081_arXiv.txt": {
        "abstract": "We present high resolution (R~20000) spectra in the blue and the far red of cicumnuclear star-forming regions (CNSFRs) in three early type spirals (NGC\\,3351, NGC\\,2903 and NGC\\,3310) which have allowed the study of the kinematics of stars and ionized gas in these structures and, for the first time, the derivation of their dynamical masses for the first two. In some cases these regions, about 100 to 150\\,pc in size, are seen to be composed of several individual star clusters with sizes between 1.5 and 4.9\\,pc estimated from Hubble Space Telescope (HST) images. The stellar dispersions have been obtained from the Calcium triplet (CaT) lines at $\\lambda\\lambda$\\,8494,8542,8662\\,\\AA, while the gas velocity dispersions have been measured by Gaussian fits to the H$\\beta$ and [OIII]\\,$\\lambda\\lambda$\\,5007\\,\\AA lines on the high dispersion spectra. Values of the stellar velocity dispersions are between 30 and 68\\,km/s. We apply the virial theorem to estimate dynamical masses of the clusters, assuming that systems are gravitationally bounded and spherically symmetric, and using previously measured sizes. The measured values of the stellar velocity dispersions yield dynamical masses of the order of 10$^7$ to 10$^8$ solar masses for the whole CNSFRs. Stellar and gas velocity dispersions are found to differ by about 20 to 30\\,km/s with the H$\\beta$ emission lines being narrower than both the stellar lines and the [OIII]\\,$\\lambda\\lambda$\\,5007\\,\\AA lines. The twice ionized oxygen, on the other hand, shows velocity dispersions comparable to those shown by stars, in some cases, even larger. We have found indications of the presence of two different kinematical components in the ionized gas of the regions. We have mapped the velocity field in the central kiloparsec of the spiral galaxies NGC\\,3351 and NGC\\,2903. For the first object the radial velocity curve shows deviations from circular motions for the ionized hydrogen consistent with its infall towards the central regions of the galaxy at a velocity of about 25\\,km/s. For NGC\\,3310 we present preliminary results of the velocity dispersions for one of the two observed slit position angles, two CNSFRs and the nucleus. ",
        "introduction": "The high-resolution blue and far red spectra were acquired as part of an observing run of three nights in 2000. They were simultaneously obtained with ISIS, using the H2400B and R1200R gratings, respectively, on the 4.2m William Herschell Telescope. In the blue arm the spectral range is 4780-5200\\,\\AA\\ and in the red arm it is 8360-8760\\,\\AA. The width of the slit is 1'' providing a resolution of ~0.21\\,\\AA/pix and ~0.39\\,\\AA/pix, respectively. The spatial resolutions are 0.4 and 0.39 arcsec/pix for the blue and red arms, respectively.  Three different slit position angles were chosen for NGC\\,3351 to observe 5 CNSFRs and the nucleus \\citep[see][]{hagele07}. In the case of NGC\\,2903 two different slit position angles allow us to study  4 CNSFRs and the nucleus \\citep[see][]{hagele08}. For NGC\\,3310 we have observed in two slit positions but we only present preliminary results for one of them. The images were processed and analysed using  IRAF routines in the usual manner.  In addition to the galaxy images, observations of 11 velocity template stars were made to provide good stellar reference frames. They are late type giant and supergiant stars which have strong CaT features \\citep[see][]{diaz89}.  The different methods used to derivate the quantities presented here are explained in detail in \\cite{hagele07}. ",
        "conclusions": "Stellar velocity dispersions are between 30 and 68\\,km/s, about 25\\,km/s larger than those measured for the gas. However, the best Gaussian fits involved two different components for the gas: a \"broad component\" with a velocity dispersion  similar to that measured for the stars, and a \"narrow component\" with a dispersion lower than the stellar one by about 30\\,km/s. This last component seems to have a relatively constant value for all the CNSFRs in each galaxy, with estimated values close to 25\\,km/s for the two gas emission lines. The velocities of the two components of the fits in the CNSFRs of NGC\\,3351 are the same within the observational errors, but in the cases of NGC\\,2903 and NGC\\,3310 we find a shift between the narrow and the broad component that vary between -10 and 35\\,km/s in radial velocity.  The dynamical masses estimated from the stellar velocity dispersion using the virial theorem for the CNSFRs of NGC\\,3351 are in the range between 4.9\\,$\\times\\,10^6$ and 4.8\\,$\\times\\,10^7$\\,M$_\\odot$, and is 3.5\\,$\\times\\,10^7$ for its nuclear region inside the inner 11.3\\,pc \\citep{hagele07}. In the case of NGC\\,2903 the masses are in the range between 6.4\\,$\\times\\,10^7$ and 2.1\\,$\\times\\,10^8$\\,M$_\\odot$ for the CNSFRs and is 2.2\\,$\\times\\,10^7$ for its nuclear region inside the inner 3.8\\,pc \\citep{hagele08}. Masses derived from the H$\\beta$ velocity dispersions under the assumption of a single component for the gas would have been underestimated by factors between approximately 2 to 4.  Masses of the ionising stellar clusters of the CNSFRs have been derived from their H$\\alpha$ luminosities \\citep{planesas97} under the assumption that the regions are ionisation bound and without taking into account any photon absorption by dust. Their values for NGC\\,3351 are between 8.0\\,$\\times\\,10^5$ and 2.5\\,$\\times\\,10^6$ M$_\\odot$ for the starforming regions, and is 6.0\\,$\\times\\,10^5$  for the nucleus \\citep{hagele07}. And for the regions of NGC\\,2903 they are between 3.3\\,$\\times\\,10^6$ and 4.9\\,$\\times\\,10^6$ M$_\\odot$, and is 2.1\\,$\\times\\,10^5$ for its nucleus \\citep{hagele08}. These values are comparable to that derived by \\cite{gonzalez-delgado95} for the circumnuclear region A in NGC\\,7714. Therefore, the ratio of the ionising stellar population to the total dynamical mass is between 0.01 and 0.16.  Derived masses for the ionised gas vary between 7.0\\,$\\times\\,10^3$ and 8.7\\,$\\times\\,10^4$ M$_\\odot$ for the CNSFRs of NGC\\,3351 \\citep{hagele07}, and is 2\\,$\\times\\,10^3$ M$_\\odot$ for the nucleus, and between 6.1\\,$\\times\\,10^4$ and 1.3\\,$\\times\\,10^5$ M$_\\odot$ for the regions and is 2\\,$\\times\\,10^3$ M$_\\odot$ for the nucleus of NGC\\,2903 \\citep{hagele08}. These values are also comparable to that derived by \\cite{gonzalez-delgado95}.  For NGC\\,3351, the rotation velocities derived for both stars and gas are in reasonable agreement, although in some cases the gas shows a velocity slightly different from that of the stars \\citep{hagele07}. The rotation curve corresponding to the position going through the centre of the galaxy shows maximum and minimum values at the position of the circumnuclear ring, as observed in other galaxies with CNSFRs \\citep[][and reference therein]{diaz99}. The differences in velocity between gas and stars can be interpreted as motions of the ionised hydrogen deviating from rotation and consistent with a radial infall to the central regions of the galaxy. Our results are consistent with those found by \\cite{rubin75} and would yield an infall velocity of about 25 km/s.  On the other hand, the observed stellar and [OIII] velocities of NGC\\,2903 are in good agreement, while the H$\\beta$ measurements show shifts similar to those find between the narrow and the broad components \\citep{hagele08}. This different behavior can be due to that the position of the single Gaussian fits are dominated by the broad component in the case of the [OIII] emission line while in the case of the H$\\beta$ are dominated by the narrow one. Again, the rotation curve corresponding to the position going through the nucleus shows maximum and minimum values at the positions of the circumnuclear regions."
    },
    "0801/0801.3562_arXiv.txt": {
        "abstract": "We describe simulations of the response of a gaseous disc to an active spiral potential. The potential is derived from an N-body calculation and leads to a multi-armed time-evolving pattern. The gas forms long spiral arms typical of grand design galaxies, although the spiral pattern is asymmetric. The primary difference from a  grand-design spiral galaxy, which has a consistent 2/4-armed pattern, is that instead of passing through the spiral arms, gas generally falls into a developing potential minimum and is released only when the local minimum dissolves. In this case, the densest gas is coincident with the spiral potential, rather than offset as in the grand-design spirals.  We would therefore expect no offset between the spiral shock and star formation, and no obvious co-rotation radius. Spurs which occur in grand-design spirals when large clumps are sheared off leaving the spiral arms, are rare in the active, time-evolving spiral reported here. Instead, large branches are formed from spiral arms when the underlying spiral potential is dissolving due to the N-body dynamics. We find that the molecular cloud mass spectrum for the active potential is similar to that for clouds in grand design calculations, depending primarily on the ambient pressure rather than the nature of the potential. The largest molecular clouds occur when spiral arms collide, rather than by agglomeration within a spiral arm. ",
        "introduction": "A key objective to understanding star formation is to explain how molecular clouds and stars form in different types of galaxies. The presence of spiral shocks in grand design galaxies explains the predominance of young stars in spiral arms \\citep{Roberts1969,Woodward1976,Cepa1990, Knapen1992,Knapen1996}, and can also account for the observed velocity dispersions in molecular clouds (\\citealt{Zhang2001,Bonnell2006}; \\citealt*{KKO2006}). Otherwise, star formation is thought to occur through triggering by supernovae \\citep{Mueller1976,Gerola1978}, spontaneously through gravitational, thermal or magnetic instabilities, or by turbulent compression (as reviewed in \\citealt{Elmegreen1996}).  From optical images, spiral galaxies can be classified as i) grand design, usually consisting of 2 symmetrical spiral arms, ii) multi-armed, with several asymmetric spiral arms or iii) flocculent, with multiple shorter arms. Using the classification scheme of \\citet{Elmegreen1987}, grand design galaxies are inclusive of multi-armed spiral galaxies, but we retain this terminology to distinguish between the two types. Classical grand design structure is thought to originate from perturbations due to bars or companion galaxies which induce a density wave in the underlying stellar disk. Flocculent (and multi-armed) structure develops from local gravitational instabilities, which are sheared into short, transient spiral arms. Most flocculent galaxies are also found to have a weak density wave in the older stellar population visible in the K band \\citep{Thornley1997,Grosbol1998}.  Several hydrodynamical simulations have investigated the response of the ISM to spiral density waves in grand design galaxies, by assuming a rigidly rotating spiral potential, usually with an isothermal medium \\citep*{Patsis1997,Chak2003,Kim2002,Dobbs2006}. The spiral pattern in these simulations is assumed to be long-lasting, of at least several rotation periods. The formation of dense structures along the spiral arms are associated with molecular clouds \\citep*{Kim2003,DBP2006,Dobbs2007} and the shearing of these features produces spurs perpendicular to the arms \\citep*{Kim2002,Dobbs2006,Shetty2006}. For the purely hydrodynamic simulations \\citep{Wada2004,DBP2006}, the development of molecular clouds and spurs is controlled by the dynamics of the shock. Similar features are formed through gravitational instabilities \\citep{Kim2002,Shetty2006}.  The generation of turbulence has also been examined in simulations of galactic disks. \\citet{Wada2002} show that gravitational instabilities combined with galactic rotation can drive turbulence, producing flocculent spiral structure. Supernovae are expected to be a major progenitor of turbulence \\citep{deAvillez2005,Joung2006} in the ISM, whilst \\citet{Piontek2005} suggest that MRI driven turbulence may be important in the outer parts of discs where supernovae are less frequent. An alternative possibility which has emerged recently is that spiral shocks may induce turbulence \\citep{Zhang2001,Bonnell2006,KKO2006,Dobbs2007}. In particular, \\citet{Bonnell2006} show the passage of gas through a spiral arm and the formation of molecular clouds which exhibit a $\\sigma \\propto r^{0.5}$ velocity dispersion law.  Realistic galactic potentials are not rigid, but time dependent. Even in grand design galaxies, the spiral pattern and pitch angle of the spiral arms may change within a few rotation periods \\citep*{Merri2006}. For galaxies where the spiral structure arises through more localised gravitational instabilities, the number and shape of spiral arms are expected to be continually evolving. Numerous N-body simulations have investigated the spiral structure arising from such instabilities in a stellar disc. If the mass of the halo is sufficient, the bar instability is prevented and the resulting structure then depends on the Toomre $Q$ parameter. Generally gravitational instabilities produce flocculent and multi-armed spirals \\citep*{Bottema2003,Li2005b}, unless an increase of $Q$ in the centre of the disk is induced to suppress $m>2$ modes \\citep*{Thomasson1990}. With $Q \\sim 1-2$, multi-armed spirals with modes up to $m=6$ are formed \\citep{Sellwood1984}. Lowering the disk mass or raising the velocity dispersion increases $Q$ and produces a flocculent galaxy \\citep{Elmegreen1993}. In all these cases, the gaseous component of the disc develops a much more pronounced spiral structure than present in the stellar component. Hydrodynamical simulations show that the gas becomes dense enough for star formation to occur in the spiral arms, providing the Toomre criterion \\citep{Toomre1964} is satisfied ($Q \\lesssim 1$ for gas \\citep{Tasker2006} or gas and stars \\citep{Li2005b}).  In this paper, we describe hydrodynamical simulations of multi-armed spiral galaxies, which use an active galactic potential. This potential is taken from an N-body simulation described in \\citet{Sellwood1984}. The potential has also recently been used for 2D gas simulations \\citep{Clarke2006} with the grid code CMHOG, which highlight the dependence of the star formation rate on the changing spiral structure. Whilst the gas dynamics in the grand design simulations are driven by a density wave, the potential from the N-body simulation arises through transient gravitational instabilities in the disc stars. We describe the galactic structure, gas dynamics and the properties of molecular clouds which arise with the active potential, and compare with previous simulations where a rigid potential was adopted \\citep{Dobbs2007}. ",
        "conclusions": "We have extended previous analysis of the structure and dynamics of the ISM in spiral galaxies by applying a time-dependent spiral potential derived from an N-body calculation. The gas forms long spiral arms, exhibiting a multi-armed structure where the underlying stellar component is determined by gravitational instabilities rather than a spiral density wave. The main difference compared to the grand design case is that the development, collision and dissipation of spiral arms in the underlying potential is  the main driver for generating high density structure in the ISM. The timescale for the evolution of structure is thus linked to the lifetime of the local minimum.  Gas accumulates in the potential minima as they develop, thus the densest gas is coincident with the stellar spiral arms. Consequently no offset would be expected between the dust lanes and star formation in such a galaxy. Spurs associated with the shearing of GMCs as they leave the spiral arm, are largely absent with the active potential. This is because gas only leaves the spiral arms once the potential minimum dissipates. The gas then retains the large scale spiral arm structure, becoming slowly wound up until it fully dissipates or encounters another spiral arm. Instead of spurs, much larger bridges, which are remnants of former spiral arms, extend between the main spiral arms.  We further identified molecular clouds from these simulations, finding that the largest GMCs occur where spiral arms are colliding. We found the surprising result that the mass spectra are similar to previous results where the potential corresponded to a density wave. This suggests that the external pressure applied to the cold gas determines the sizes and masses of the clouds, with a steeper index when the cold clumps are pressure confined.  A natural extension of this work would be to model both the stellar and gaseous components of the disc, to include the gravitational instabilities of the stars and the response of the gas simultaneously \\citep{Li2005b}. Furthermore incorporating stellar feedback is required to treat the stellar and gaseous components self consistently. We have also neglected heating and cooling, magnetic fields and self gravity, which we are in the process of analysing with respect to grand design galaxies. We leave the consideration of these processes for more general stellar potentials to future work."
    },
    "0801/0801.3879_arXiv.txt": {
        "abstract": "The behavior of spin-half particles is discussed in the $3 + 1$-dimensional constant curvature black hole (CCBH) spacetime. We use Schwarzschild-like coordinates, valid outside the black hole event horizon. The constant time surfaces corresponding to the time-like Killing vector are degenerate at the black hole event horizon and also along an axis. We write down the Dirac equation in this spacetime using Newman-Penrose formalism which is not easily separable unlike that in the Kerr metric. However, with a particular choice of basis system the equation is separable and we obtain the solutions. We discuss the structural difference in the Dirac equation in the CCBH spacetime with that in the Kerr geometry, due to difference in the corresponding spacetime metric, resulting complexity arised in separation in the earlier case. ",
        "introduction": "A new type of black hole solution has been found by Ba{\\~n}ados \\cite{B1}. This solution results from an identification of space-time points in the anti-de Sitter space and represents a higher dimensional generalization of the $2+1$ dimensional BTZ black hole \\cite{B2}. This is a constant curvature black hole (CCBH) with a negative cosmological constant. The corresponding spacetime geometry is $R^3 \\times S^1$. Unlike the BTZ black hole, the $3+1$ dimensional black hole does not have rotating solution and is dynamic. However Schwarzschild-like coordinates, valid outside the horizon, can be found which covers a sub-manifold of the space time.  The scalar field solution in the $3+1$ dimensional CCBH was discussed in \\cite{K1} while studying their thermodynamic behavior. On the other hand, the spinor field solution was studied extensively in black hole spacetimes over last three decades by various authors including one of the present ones (e.g. \\cite{teu,chandra76,page76,chandra83,guv90,mc99,mc,md,m00}). These authors discussed in detail the separability of the Dirac equation in black hole spacetimes and solutions. In the present article we study the behavior of a spin-half field in the $3+1$ dimensional CCBH spacetime. We use Newman-Penrose formalism to establish the Dirac equation. It is found that in this background the Dirac equation is not easily separable, unlike the situation in the Kerr geometry, unless we change the basis appropriately. This is, as we discuss in the following sections in detail, due to very nature of the background geometry which is, by definition, quite different from conventional black hole spacetimes.  The paper is organized as follows. In the next section we recall the $3+1$ CCBH metric and describe its properties briefly. In \\S III we calculate the basis vectors of null tetrad in this spacetime and write down the Dirac equation following the Newman-Penrose formalism. We show that the set of equations is separable into radial and angular part by defining new spinors obtained with a suitable combination of original spinors. Subsequently we discuss the solutions in \\S IV. The section V is devoted in obtaining the effective potentials of spinors in the CCBH spacetime and comparing with that in the Kerr geometry. We show that due to difference in structure of spacetime metric the Dirac equation in the CCBH spacetime appears different than that in the Kerr geometry hindering its separation in regular basis where particles are expressible as {\\it pure} up and/or down spinors. Finally, we summarize our results in \\S VI. ",
        "conclusions": "We have considered the spin-half fermion field in the $3+1$-dimensional CCBH background. The Dirac equation is obtained by using the Newman-Penrose formalism. We show that the Dirac equation is separable only with the change of basis. The new spinors are linear combination of original spinors giving rise to the Dirac equation separable into radial and angular parts. We obtain the corresponding solutions. We discuss the difference in the Dirac equation in the CCBH spacetime with that in the Kerr geometry recalling the geometrical conditions to be satisfied to separate them out \\cite{cm}. We show that the structure of equations in the earlier case is different than the later one (due to difference in structure of the spacetimes) leading to a different separation and solution procedure. Now one can consider the Dirac equation solution in the CCBH spacetime to study the thermodynamical properties of spinor field as was done earlier for scalar field \\cite{K1}.  Note that the coordinate system used in the text does not cover the entire manifold outside the horizon completely. There are other foliations which cover the outside horizon region completely \\cite{HP}. However in this foliation the metric becomes explicitly time dependent."
    },
    "0801/0801.1421_arXiv.txt": {
        "abstract": "{In this short note I discuss the hypothesis that bursting activity of magnetars evolves in time analogously to the glitching activity of normal radio pulsars (i.e. sources are more active at smaller ages), and that the increase of the burst rate follows one of the laws established for glitching radio pulsars. If the activity of soft gamma repeaters decreases in time in the way similar to the evolution of core-quake glitches ($\\propto t^{5/2}$), then it is more probable to find the youngest soft gamma repeaters, but the energy of giant flares from these sources should be smaller than observed $10^{44}$~--$10^{46}$~ergs as the total energy stored in a magnetar's magnetic field is not enough to support thousands of bursts similar to the prototype 5 March 1979 flare.} ",
        "introduction": "Soft gamma repeaters (SGRs) are known to demonstrate three main types of activity: weak, intermediate, and giant flares (see a recent review in Woods \\& Thompson 2006). \\footnote{Here often I do not discuss separately, in the context of the proposed hypothesis, hyperflares (HFs) like the one from SGR 1806-20 observed on 2004 Dec 27. Such events are just included into the class of giant flares.} More energetic events are less frequent than weak bursts. The distribution of SGRs burst energy has a power-law shape, $dN/dE \\propto E^{-\\alpha}, \\alpha\\sim 5/3$. This distribution is similar to the Gutenberg-Richter law established for terrestrial quakes (Cheng et al. 1996; G\\\"o\\u{g}\\\"u\\c{s} et al. 1999). A lot is known about properties of bursts of different kinds, but the evolution of the rate of bursts during lifetime of SGRs in not known due to poor statistics and relatively short period of exploration of these sources.  Here I discuss a possible link between SGRs bursts and period glitches, especially glitches driven by starquakes. Such connection was suggested long ago. Also the possibility that starquakes can power cosmic gamma-ray bursts was discussed by Blaes et al. (1989) (see references to earlier papers therein). However, to the author's knowledge, the possible similarity between the evolution of the bursts rate during a SGR lifetime  and the evolution of the rate of glitches derived from radio pulsar observations has never been discussed in literature. The reason lies in the accepted difference between radio pulsar and magnetar glitches. In the case of radio pulsars,  two types of glitches are distinguished: Vela-like and Crab-like. They are thought to be due to either starquakes or superfluid vortex unpinning. Glitches in magnetars are believed to be connected with strong magnetic fields of these stars. Nevertheless, I think that a discussion of possible similarities between  the evolution of glitch rates in different kinds of neutron stars (NSs) is worth while, especially in connection with giant flares.  During $\\sim 25$~--~$30$ years of observation four strong flares have been detected from four known SGRs (Mazets et al. 2004)\\footnote{Often only two bursts -- 05 March 1979 from SGR 0526-66 and 27 August 1998 from SGR 1900+14 -- are called giant bursts. Many authors do not include the flare from SGR 1627-41 on 18 June 1998 into the list of giant flares, as for this event a pulsating ``tail'' was not observed. The flare of SGR 1806-20 is also sometimes classified as an event of a different type -- a {\\it hyperflare} - HF.}. Despite their smaller rate, due to large energy release these bursts dominate among other types of bursting activities of SGRs. The rate of such flares was usually assumed to be $\\sim 0.02-0.01$ per year per source (Woods \\& Thompson 2006) based on just two giant flares (GFs) (5 March 1979 and 28 August 1998). However, with inclusion of the 27 December 2004 flare and with inclusion of the 18 June 1998 burst of SGR 1627-41 the rate becomes higher. Still, below I shall use a conservative estimate $0.02$ GFs per year per source. Data on possible extragalactic SGR flares (Ofek 2007) is only marginally compatible with this rate, and I briefly discuss this problem below.  Previously several authors (Nakar et al. 2006; Popov \\& Stern 2006; Lazzati, Ghirlanda \\& Ghisellini 2005; Ofek 2007) addressed the problem of extragalactic SGR flares. They used different archive data (BATSE, IPN) to search for examples of extragalactic GFs and HFs, and to put limits on the fraction of SGR flares among short GRBs (see Nakar 2007 for a review on short GRBs and for some discussion of the fraction of SGR flares among them). Estimates show that GFs can be observed by most of the detectors (from BATSE to SWIFT) up to distances of few Mpc (Palmer et al. 2005). I.e., even SWIFT cannot observe GFs from most of the Virgo cluster -- only HFs can be detected from such distance (Hurley et al. 2005). \\footnote{Details about SWIFT sensitivity can be found in (Band 2006). SWIFT appears to be more sensitive than BATSE to long bursts, but not for short hard events. The advantage of SWIFT is in better localization of short events. BAT onboard SWIFT is less sensitive than BATSE to energies $>100$ keV. HFs are relatively hard events: spectra of the HF 27 December 2004 can be described as a blackbody with T~$\\sim$~200 keV. So, in general, SWIFT is not much better than BATSE for detection of bursts like the only known HF. Other experiments, like BeppoSAX and HETE-II, also are not more sensitive than BATSE to short hard bursts. However, as noted by Band (2006), an analysis of short-duration triggers by SWIFT which do not result in successful image triggers can result in discovery of more short flares. Up to the limiting distance for SWIFT detection the extrapolation of the galactic rate of GFs gives an estimate of few  events in a year, so SWIFT with its sky coverage of about 1.4 steradians can detect approximately one GF per year. HFs can be detected from larger distances: 30-40 Mpc by BATSE, and up to 70 Mpc -- by SWIFT (Hurley et al. 2005). Scaling the galactic rate of HFs -- one in 30 years -- Hurley et al. (2005) gives the rate about 50 per year for SWIFT. However, any rescaling of the galactic rate of HFs is extremely uncertain as only one event has been observed. Results presented in (Popov \\& Stern 2006; Lazzati et al. 2005; Ofek 2007) show that the rate of extragalactic HFs is much lower than that  expected from a naive  scaling (based on starformation rate or supernova rate, etc.) given one observed event from SGR 1806-20 in our Galaxy. Upper limits from non-detection of HFs from the Virgo cluster by BATSE provide an estimate $\\sim 10^{-3}$ yr$^{-1}$ per Milky Way-like galaxy (Popov \\& Stern 2006). This number after rescaling again gives an estimate of about one detection per year for SWIFT. Still, larger statistics of HFs detection is necessary. Recent observations (Frederiks et al. 2007; Ofek et al. 2006) have already provided a candidate for a HF in a close-by galaxy with high starformation rate (such stellar systems are the most prominent hosts of GFs and HFs). Hopefully, more events can be detected by SWIFT during its lifetime.}  Usualy authors apply the value of GFs rate $\\sim 0.02-0.01$ per year per source to the whole life period of SGRs, i.e. the rate of GFs is assumed to be constant in time. In this short note I discuss an alternative assumption. NS activity of different kinds usually decreases with time, and the GF rate should not be an exception. Possibly the closest analogue is the evolution of glitch (especially, starquake) rate. ",
        "conclusions": "I propose and discuss the hypothesis that the bursting activity of SGRs decreases with time similar to the glitching activity of normal radio pulsars. This assumption leads to several interesting consequences. Namely,  \\begin{itemize} \\item For the same burst energy the probability to detect a SGR at least does not diminish with decreasing  SGR age. Even it can be more probable to detect young SGRs due to their more frequent bursts. For example, the galactic (in general, Local group) population of SGRs should already include all the youngest objects of this type. \\item If the bursting activity of magnetars rapidly decreases with time, but a typical GF energy is constant (or decreases too), then we face an energy crisis: there is not enough energy in the magnetic field to support thousands of flares similar to, say, the 5 March 1979 burst. \\item Young SGRs should burst more often, but their giant flares should be less powerful than the known ones. \\end{itemize}  Absence of many identified extragalactic giant and hyper flares by itself can rule out joint decrease of the flare rate and flare energy with time."
    },
    "0801/0801.4803_arXiv.txt": {
        "abstract": "{S147 is a large faint shell-type supernova remnant (SNR) known for its remarkable spectral break at cm-wavelength, which is an important physical property to characterize SNR evolution. However, the spectral break is based on radio observations with limited precision.} {New sensitive observations at high frequencies are required for a detailed study of the spectral properties and the magnetic field structure of S147.} {We conducted new radio continuum and polarization observations of S147 at $\\lambda$11\\ cm and at $\\lambda$6\\ cm with the Effelsberg 100-m telescope and the Urumqi 25-m telescope, respectively. We combined these new data with published lower-frequency data from the Effelsberg 100-m telescope, and with very high-frequency data from WMAP to investigate the spectral turnover and polarization properties of S147. } {S147 consists of numerous filaments embedded in diffuse emission. We found that the integrated flux densities of S147 are $34.8\\pm 4.0~$Jy at $\\lambda$11\\ cm and $15.4\\pm 3.0~$Jy at $\\lambda$6\\ cm. These new measurements confirm the known spectral turnover at $\\sim$1.5\\ GHz, which can be entirely attributed to the diffuse emission component. The spectral index above the turnover is $\\alpha = -1.35 \\pm 0.20~(\\rm S\\sim\\nu^{\\alpha})$.  The filamentary emission component has a constant spectral index over the entire wavelength range up to 40.7~GHz of $\\alpha = -0.35 \\pm 0.15$. The weak polarized emission of S147 is at the same level as the ambient diffuse Galactic polarization. The rotation measure of the eastern filamentary shell is about $-70$\\ rad m$^{-2}$.  } {The filamentary and diffuse emission components of S147 have different physical properties, which make S147 outstanding among shell type SNRs. We attribute the weak polarization of S147 at $\\lambda$11\\ cm and at $\\lambda$6\\ cm to a section of the S147 shell showing a tangetial magnetic field direction. } ",
        "introduction": "The extended object Shajn 147 (S147) has long, beautiful delicate filaments visible in the optical bands, and is located in the anti-center of the Galaxy. It was first classified as a possible shell-type supernova remnant (SNR) by \\citet{m+58} and \\citet{v+60}. The distance to S147 was estimated to be around 1~kpc \\citep{cc76,k+80}, using the so called $\\Sigma$-D relation \\citep{milne+79}. The interstellar reddening toward the SNR \\citep{f+85} suggests a smaller distance of 0.8~kpc. This is supported by absorption lines of the B1e star HD36665 at 880~pc originating from high velocity gas of S147 \\citep{p+81,sw+04}. The expansion velocity of the shell is about 80 - 120 km s$^{-1}$ \\citep{p+81,ka+79}. The age of S147 was estimated to be $\\sim 2 \\times 10^5$~yr through the radiative blast wave's Sedov-Taylor solution \\citep{sf+80,k+80}. The PSR J0538+2817 has been discovered within the boundary of S147 \\citep{a+96}, and is probably physically associated with the SNR \\citep{a+96,n+07}. The kinematic age of the pulsar, if originated from the SNR center, is about $\\sim 4\\times 10^4$~yr. The parallax of the pulsar suggests a distance of $1.47_{-0.27}^{+0.42}$~kpc \\citep{n+07}.  \\citet{d+74} summarized early radio maps of S147 between 178~MHz and 1420~MHz. She obtained a flux density spectral index of $\\alpha= -0.67\\pm 0.37$, revealing S147 as a SNR. Later, radio observations of S147 at 1648\\ MHz and 2700\\ MHz were made by \\citet{k+80}, and the southern part of the SNR was observed at 4995\\ MHz by \\citet{sf+80}.  These high-frequency observations suggested a spectral turnover near 1.5\\ GHz with a flat spectrum at lower frequencies and a steep spectrum at higher frequencies. The turnover needs to be confirmed due to large uncertainties in the low-frequency measurements and the fact that the observation at 4995\\ MHz covered only  the southern part of the remnant. Later sensitive measurements, at 1425\\ MHz and 2695\\ MHz, by \\citet{fr+86} could not improve the integrated radio spectrum, but revealed spectral index variations across S147. Regions of bright filaments show a flux density spectral index of about $\\alpha = -0.5$, while regions in between show a spectral index near $\\alpha = -1.0$. The observation at 863\\ MHz by \\citet{rzf+03} provided a more precise flux density at a low frequency, but did not improve on the uncertainty in the spectrum. The 2695\\ MHz radio map made by \\citet{fr+86} for the first time showed linearly polarized emission associated with some filamentary structure. Detailed information from this polarization observation cannot be derived due to an insufficient signal-to-noise ratio.  One goal of this work is to use the new sensitive $\\lambda$6\\ cm receiver at the Urumqi 25-m radio telescope to obtain a complete radio map at this band for a detailed study of the spectral turnover. Here we present results of total power and linear polarization maps at 4800\\ MHz ($\\lambda$6\\ cm). We also present new observations at 2639\\ MHz ($\\lambda$11\\ cm) with the Effelsberg 100-m radio telescope and unpublished polarization data at 1420\\ MHz ($\\lambda$21\\ cm)  obtained from the ``Effelsberg Medium Galactic Latitude Survey\" \\citep{rf+04}. The new observations and the data processing are described in Sect.~2. The investigation of the integrated radio spectrum and possible spectral variations across the remnant are discussed in Sect.~3. The radio polarization and the rotation measure toward S147 are analyzed in Sect.~4. The results are summarized in Sect.~5. ",
        "conclusions": "We present new sensitive maps of S147 for the total  and polarized intensity at $\\lambda$6\\ cm and $\\lambda$11\\ cm. The $\\lambda$6\\ cm maps are the first complete ones obtained so far. Our new measurements confirm the turnover of the integrated spectrum of S147 at a frequency of about 1.5~GHz.  The measured structures have been decomposed into filamentary structures (smaller than $18\\arcmin$) and large-scale diffuse emission component. The T-T plots between 2639\\ MHz and 4800\\ MHz for these two emission components prove that the turnover is attributed entirely to the diffuse emission component, which has a spectral index of $\\alpha$ $\\sim$-1.35. The filamentary emission component has a spectral index of $\\alpha$ $\\sim$-0.35 and can be traced up to 40.7\\ GHz on WMAP maps.  We discussed several possible mechanisms to explain the turnover of the diffuse emission component. We can exclude the scenarios of two populations of electrons as well as the diffuse shock accelerated of electrons. However, a compressed magnetic field is a possible explanation. The SNR shock wave compresses the local magnetic field and shifts the turnover in the Galactic radio spectrum from about 400\\ MHz to about 1.5\\ GHz.  High synchrotron losses, during the early phase of the SNR, would cause a bend at a rather high frequency, but subsequent expansion of S147 would shift the frequency toward about 1.5\\ GHz as well.  S147 is a faint source located on the near side of the Perseus arm. Polarized emission from S147 is confused  with large-scale Galactic emission, except for the bright filamentary regions. The RM in the eastern filamentary region is about $-70$\\ rad m$^{-2}$. The magnetic field in the filamentary region has a strength of typically $\\sim~30\\mu$G."
    },
    "0801/0801.2202_arXiv.txt": {
        "abstract": "Our understanding of coronal phenomena, such as coronal plasma thermodynamics, faces a major handicap caused by missing coronal magnetic field measurements. Several lines in the UV wavelength range present suitable sensitivity to determine the coronal magnetic field via the Hanle effect. The latter is a largely unexplored diagnostic of coronal magnetic fields with a very high potential. Here we study the magnitude of the Hanle-effect signal to be expected outside the solar limb due to the Hanle effect in polarized radiation from the H~{\\sc{i}} Ly$\\alpha$ and $\\beta$ lines, which are among the brightest lines in the off-limb coronal FUV spectrum. For this purpose we use a magnetic field structure obtained by extrapolating the magnetic field starting from photospheric magnetograms. The diagnostic potential of these lines for determining the coronal magnetic field, as well as their limitations are studied. We show that these lines, in particular H~{\\sc{i}} Ly$\\beta$, are useful for such measurements. ",
        "introduction": "The magnetic field is the main driver of all physical phenomena in the solar corona. The field lines of force thread the solar atmosphere coupling its different layers to the solar interior. They transport energy to the corona, where it is dissipated and heats the plasma to several MK as well as accelerating particles to several hundreds of km~s$^{-1}$ within a few solar radii above the surface. The magnetic field also affects the interplanetary medium through the transfer of angular momentum from the Sun through the solar wind and transient events. The latter also have important consequences on planets such as the Earth. In spite of its key role, the magnetic field remains an almost unknown parameter for any study involving the solar upper atmosphere.  Measuring the magnetic field vector is a routine exercise in the photosphere and to a smaller degree in the chromosphere. Magnetic fields in these two layers are strong enough to yield an easily measurable Zeeman signal (and also Hanle effect for numerous spectral lines). Inversion techniques of the observed Stokes parameters are sufficiently sophisticated from both, a physical and computational point of view. However, the coronal magnetic field is rather weak, so that the Zeeman splitting in (generally weak) coronal lines is too small to be accurately measured except above active regions with relatively strong field for few spectral lines with exceptionally large Land\\'e factors (Lin et al. 2004). Significant Zeeman splitting is usually accompanied by the so-called ``strong regime of the Hanle effect'' where the direction of the magnetic field projected on the plane of the sky can be obtained from the linear polarization signal of spectral lines such as the Fe~{\\sc{xiv}} 530.3~nm green line. However, the coronal plasma is optically thin and direct coronal magnetic field measurements have a line-of-sight integrated character. To obtain the 3D structure of the magnetic field, vector tomographic methods are required (see Kramar et al. 2006 and Kramar \\& Inhester 2007).  Techniques based on different physical mechanisms, such as radio gyroresonant and bremsstrahlung emission, are being developed. However, they provide field strengths in different layers (depending on the frequency-opacity relation) without height information on where the signal comes from (see Lee et al. 1997). They also remain restricted to areas above active regions. These measurements are, however, important to constrain coronal magnetic fields obtained though the extrapolation of photospheric and chromospheric field measurements. Extrapolation techniques of magnetic fields have undergone a rapid development and could benefit enormously from high quality and spatially extended measurements of magnetic fields in the photosphere and chromosphere and probably the lower corona (Solanki et al. 2003).  Theoretically, direct determination of the magnetic field in the solar corona could be achieved through linear polarization of spectral lines with suitable sensitivity to the Hanle effect. Raouafi et al. (1999a,b) successfully measured the linear polarization of the O~{\\sc{vi}} 103.2~nm line and interpreted it in terms of the Hanle effect signature of the coronal magnetic field (see also Raouafi et al. 2002a,b).  The aim of the present paper is to study the possibility of measuring the coronal magnetic field vector through the Hanle effect using coronal lines with both, relatively adequate sensitivity to the field and high emissivity. The study aims to be relatively realistic in the sense that the magnetic field used is an extrapolation of photospheric measurements. The results presented here are in a preliminary stage and the different inputs to the method have to be improved in the future in order to be used to underpin future missions to measure the coronal magnetic field directly. ",
        "conclusions": "Although the model used to compute the polarization of the hydrogen lines is very simple and can be significantly improved in numerous aspects, significant and practically measurable polarization parameters (rates and direction angles of the plane of polarization) due to the Hanle effect resulting from the coronal magnetic field are obtained. The use of coronal magnetic field data obtained from the potential field extrapolation of photospheric magnetograms is a step forward in realistically modeling solar coronal magnetic phenomena. Future missions will benefit from such simulations of the measurement of the coronal magnetic field.  Hanle effect measurements in the UV are a promising way of measuring the Sun's coronal magnetic field, although due to the lack of appropriate instruments the use of the technique has so far been limited (e.g. Raouafi et al. 1999a,b). Ly$\\alpha$ and $\\beta$ are promising spectral lines for coronal Hanle effect measurements. Their Hanle sensitivity to different field strengths and their different formation processes make them complementary to each other. The polarization of several other spectral lines (e.g., O~{\\sc{vi}} 1032~{\\AA}) have to be explored in order to widen the choice range for any attempt to measure the magnetic field in the solar corona."
    },
    "0801/0801.0031_arXiv.txt": {
        "abstract": "We have observed the supernova remnant MSH\\,15$-$52 (G320.4$-$1.2), which contains the gamma-ray pulsar PSR B1509$-$58, using the CANGAROO-III imaging atmospheric Cherenkov telescope array from April to June in 2006. We detected gamma rays above 810\\,GeV at the 7 sigma level during a total effective exposure of 48.4~hours. We obtained a differential gamma-ray flux at 2.35\\,TeV of $(7.9\\pm 1.5_{\\rm stat}\\pm 1.7_{\\rm sys}) \\times 10^{-13}$\\,cm$^{-2}$\\,s$^{-1}$\\,TeV$^{-1}$ with a photon index of $2.21 \\pm 0.39_{\\rm stat} \\pm 0.40_{\\rm sys}$, which is compatible with that of the H.E.S.S.\\ observation in 2004. The morphology shows extended emission compared to our Point Spread Function. We consider the plausible origin of the high energy emission based on a multi-wavelength spectral analysis and energetics arguments. ",
        "introduction": "Recent imaging atmospheric Cherenkov telescopes (IACTs) have achieved remarkably high sensitivity in the very high energy gamma-ray band. This is well illustrated by the Galactic plane survey carried out by the H.E.S.S.\\ collaboration \\citep{aha06,hop07}. Many of the discovered TeV sources are associated with pulsar wind nebulae (PWNe) \\citep[e.g.,][]{gae06}, which are now established as the most populous category among Galactic TeV sources: 18 PWNe have been found so far \\citep{hin07} and a portion of 21 unidentified galactic TeV sources could be PWNe as well \\citep{aha07b,fun07b}.  The Crab nebula, the prototype PWN, has been observed in every accessible wave band. The nebula contains the Crab pulsar, which has the highest spindown energy loss of known gamma-ray pulsars \\citep{tho03}. However, the radiation mechanism remains controversial for the Crab nebula and for other PWNe: it has been argued that the Crab nebula's broadband spectral energy distribution (SED) can be well explained by synchrotron emission in an average magnetic field of $\\sim 0.1-0.3$\\,mG for radio to soft gamma-ray bands, and by inverse Compton (IC) scattering of synchrotron, IR, millimeter and cosmic microwave background (CMB) photons for the hard gamma-ray band up to 100\\,TeV \\citep{ahat96,ato96,aha97,aha98,aha04}. On the other hand, it was claimed that gamma rays from the decay of neutral pions, produced by hadrons, contribute significantly at $\\sim 10$\\,TeV for the Crab nebula \\citep{bed03}. Consideration has also been given to whether the energy source of TeV gamma-ray emission from an SNR containing a pulsar (composite SNR) is the pulsar's spindown energy or related to the supernova explosion \\citep[e.g.][]{fun07b}, and the efficiency of energy conversion to the particle acceleration for such models.  PSR B1509$-$58 has the third highest spindown energy loss after the Crab pulsar \\citep{tho03} and PSR J1833-1034 \\citep{cam06} in the Galaxy, and its nebula has also been well studied across the electromagnetic spectrum. TeV gamma-ray observations, combined with those at other energy bands, provide additional information which may lead to solutions for the above problems and a unified comprehension of pulsar and nebula systems. PSR B1509$-$58 was detected in the radio supernova remnant MSH\\,15$-$52 (G320.4$-$1.2) \\citep{cas81}, initially as a 150-ms X-ray pulsar with the {\\it Einstein} satellite \\citep{sew82}, which was confirmed by later X-ray/soft gamma-ray observations \\citep{kaw93,mat94,gun94,sai97,mar97,cus01,for06}. Subsequently its pulse period was detected at radio frequencies \\citep{man82}, and at soft gamma-ray energies with COMPTEL \\citep{kui99}. The pulsar was detected above 30\\,MeV by EGRET at the 4.4$\\sigma$ level with some suggestive, if not statistically compelling, evidence of modulation at the pulsar period \\citep{kui99}. Although optical \\citep{she98} and near-IR \\citep{kap06} searches found possible pulsar counterparts, the pulse period was not detected. As one of the most energetic young pulsars, PSR B1509$-$58, has been particularly well studied at radio wavelengths. A detailed timing analysis yielded a period derivative of $\\dot{P} = 1.5\\times 10^{-12}$, and a high spin-down luminosity of $\\dot{E} = 1.8\\times 10^{37}I_{45}$\\,ergs\\,s$^{-1}$, where $I_{45}$ is the moment of inertia in units of $10^{45}$g cm$^{-2}$. A braking index of $n = 2.84$ was measured, corresponding to an age, assuming an initial period of $P_0 \\ll P$, of $\\tau = (P/(n-1)\\dot{P})[1-(P_0/P)^{n-1}] \\sim(P/(n-1)\\dot{P}) \\sim 1700$\\,yr. A large dipole surface magnetic field of $B=1.5\\times10^{13}$\\,G was also inferred \\citep{kas94,liv05}. However, in contrast to other young pulsars, no pulsar glitch has been observed to date. Initially, the age of the SNR was estimated to be $6-20$\\,kyr \\citep{sew83} or  $\\sim 10$\\,kyr \\citep{van84} prompting debate about the disagreement with the pulsar's age. \\citet{bla88} and \\citet{gva01}, for instance, suggested the pulsar age was actually $\\ge 20$\\,kyr. The latter was based on the assumption that pulsar's braking torque was enhanced by the interaction between the pulsar's magnetosphere and circumstellar dense clumps.  The radio morphology of MSH\\,15$-$52 consists of southeast and northwest shells. The latter, $\\sim 10'$ from the pulsar, spatially coincides with the H$\\alpha$ nebula RCW89 \\citep{rod60}, and \\citet{gae99} concluded that MSH\\,15$-$52, PSR B1509$-$58 and RCW89 were associated systems. The distance, derived from an H\\,I absorption measurement, is $5.2 \\pm 1.4$\\,kpc, consistent with the value of $5.9 \\pm 0.6$\\,kpc determined from the pulsar dispersion measure \\citep{tay93}. We adopt $d=5.2$\\,kpc throughout this paper.  Symmetric jets, similar to the Crab pulsar \\citep{bri85} and the Vela pulsar \\citep{hel01}, were observed by {\\it ASCA} \\citep{tam96}, {\\it ROSAT} \\citep{tru96} and {\\it Chandra} \\citep{gae02}. The jet directed towards the northwest was observed to terminate at RCW89. Precise {\\it Chandra} observations revealed the sequential heating of the knots in RCW89 by the pulsar jet \\citep{yat05}. The high resolution {\\it Chandra} image also showed the arc structure where the pulsar wind may be terminated, and the diffuse pulsar wind nebula \\citep{gae02}, which emitted non-thermal synchrotron radiation. However the PWN has not been observed at other wavelengths such as IR or optical. Although IRAS found an infrared source, {\\it IRAS} 15099$-$5856, spatially coincident with PSR B1509$-$58, the IR emission was thermal and therefore not related to the PWN \\citep{are91}. A faint radio structure was detected \\citep{gae99} but its flux density was obtained only within a large error \\citep{gae02}. A larger extent than that of X-ray was expected due to the difference of the cooling lifetime so that it was thought to be partially hidden by the bright RCW89.   The {\\it Ginga} LAC discovered single power-law emission up to 20\\,keV with a photon index of $\\sim2$, indicating synchrotron emission and the existence of accelerated electrons \\citep{asa90}. The synchrotron nebula spectrum was also measured by {\\it EINSTEIN} \\citep{sew83}, {\\it EXOSAT} \\citep{tru90} and {\\it RXTE} \\citep{mar97}. {\\it BeppoSAX} detected nonthermal emission from 1\\,keV up to 200\\,keV with a photon index of $\\Gamma = 2.08\\pm 0.01$ \\citep{min01}, while recent observations with {\\it INTEGRAL} IBIS found a possible ($2.9\\sigma$) spectral cut off at $\\sim 160$\\,keV \\citep{for06}. Since high energy electrons were demonstrated to exist, very high energy gamma-ray emission was predicted via IC scattering with CMB photons \\citep{du95,har96}. CANGAROO-I also suggested a possible VHE gamma-ray detection of $\\sim 10 \\%$ of the Crab flux above 1.9\\,TeV, assuming a spectral photon index of 2.5 \\citep{c1}. H.E.S.S.\\ subsequently reported extended VHE gamma-ray emission along with the pulsar jet. Their TeV morphology showed a good coincidence with X-ray images, indicating that the TeV gamma-rays originate from the inverse Compton scattering of relativistic electrons. It was pointed out that IC scattering of the CMB could not account for most of the TeV gamma-ray flux with an assumed magnetic field of 17~$\\mu$G, which indicated the contribution of IR photons as seed photons for the IC process \\citep{aha05,khe05}. The necessity for target photons in addition to the CMB has recently been suggested for other Galactic sources also \\citep{khe05,hin07a}. Here we report the results of TeV gamma-ray observations with the CANGAROO-III telescopes and consider the origin of the TeV gamma-ray emission based on a discussion of the energetics. ",
        "conclusions": "CANGAROO-III observed the SNR MSH\\,15-52 containing PSR B1509$-$58 for 48.4 hours in 2006 and detected VHE gamma-ray emission at the 7$\\sigma$ level.  The obtained differential flux and the intrinsic extent of TeV gamma-ray emission are consistent with the previous H.E.S.S.\\ result. Studies of the multi-wavelength spectra was performed, based on both hadronic and leptonic models. In the leptonic scenario, an IR photon field and a cooled broken power-law spectrum of electrons are necessary to reproduce the TeV gamma-ray emission. From the point of view of the energetics, if we do not take into account the expansion loss, a typical supernova could provide sufficient energy for electrons to reproduce the SED, while hardly to protons. The morphology of the TeV gamma-ray emission, however, does not support the supernova explosion as the global energy source. Electrons can also be accelerated enough to reproduce the SED when $\\ge 4$\\% ($\\tau _0=30$\\,yr is assumed) of the rotational energy is supplied to the kinetic energy.  If the pulsar was older than its characteristic age of 1700\\,yr, e.g., 20\\,kyr, the required acceleration efficiency would be reduced.   Filling in the gaps in the SED is very important for the discussion above. IR/optical measurements of the PWN emission are crucial to determine the synchrotron spectrum and the electron scenario. Additionally, the determination of the spectral break in the synchrotron or IC component would help to resolve the long-standing question of the age of this complex system. We await the results from all sky survey of the recently launched IR satellite {\\it Akari} \\citep{mur07}. GLAST and the next generation of the ground-based IACTs such as CTA \\citep{cta} and AGIS \\citep{agis} are expected to determine the IC spectra. In addition, X-ray observations of the whole PWN, excluding PSR B1509$-$58, are required in order to accurately estimate the flux of synchrotron nebula emission."
    },
    "0801/0801.2172_arXiv.txt": {
        "abstract": "We present a {\\it Spitzer Space Telescope} imaging survey of the most massive Galactic globular cluster, $\\omega$\\,Centauri, and investigate stellar mass loss at low metallicity and the intracluster medium (ICM). The survey covers approximately 3.2$\\times$ the cluster half-mass radius at 3.6, 4.5, 5.8, 8, and 24~\\micron{}, resulting in a catalog of over 40,000 point-sources in the cluster. Approximately 140 cluster members ranging 1.5 dex in metallicity show a red excess at 24~\\micron{}, indicative of circumstellar dust. If all of the dusty sources are experiencing mass loss, the cumulative rate of loss is estimated at 2.9 -- 4.2 $\\times$ 10$^{-7}~M_\\odot$ yr$^{-1}$, 63\\% -- 66\\% of which is supplied by three asymptotic giant branch stars at the tip of the Red Giant Branch (RGB). There is little evidence for strong mass loss lower on the RGB.  If this material had remained in the cluster center, its dust component ($\\gtrsim$ 1 $\\times$ 10$^{-4}~M_\\odot$) would be detectable in our 24 and 70~\\micron{} images. While no dust cloud located at the center of $\\omega$\\,Cen is apparent, we do see four regions of very faint, diffuse emission beyond two half-mass radii at 24~\\micron{}.  It is unclear whether these dust clouds are foreground emission or are associated with $\\omega$\\,Cen. In the latter case, these clouds may be the ICM in the process of escaping from the cluster. ",
        "introduction": "} \\label{sec:intro}  Despite their old age and low metal content, Galactic globular clusters (GCs) are now known to be dust factories \\citep{ramdani01, origlia02, evans03, boyer06, jaccoea06a, lebzelter06, ita07}. This has important implications for our understanding of mass loss from low-mass and metal-poor stars and for the replenishment of the interstellar medium (ISM) within old, metal-poor stellar systems, such as GCs, dwarf spheroidal (dSph) galaxies, and galactic halos.  All stars more massive than $\\simeq~0.8~M_\\odot$ lose a significant fraction of their mass before they deflagrate or leave a compact remnant \\citep[see reviews by][]{iben83,chiosi86}. In the final stages of evolution, most stars enter higher stages of nuclear burning, and the products of this nucleosynthesis are transported to the stellar surface through convection or rotation. The chemically enriched mass that is recycled (with some delay) into the ISM has a significant impact on any further star formation, and drives the chemical evolution within galaxies. In particular, the enrichment of the ISM by dust is of great importance, as grains play important roles in many ISM processes, including the formation of molecular hydrogen \\citep{gould63} and planet formation. The winds of low and intermediate-mass stars on the upper red giant and asymptotic giant branches (RGB and AGB) and massive red supergiants are prolific dust producers \\citep{gehrz89}. The dust is an integral part of the driving mechanism of the most massive winds \\citep[e.g.,][]{gehrz71, jura85, gail87}, and our understanding of red giant mass loss relies heavily on the details of the dust condensation process \\citep[e.g.,][]{salpeter74,gauger90,hofner97} --- which is poorly understood in particular at low metallicity \\citep[for a review, see][]{jacco06}.  GCs are exquisite laboratories in which to study stellar winds because they represent the closest match to a single stellar population at metallicities that span over two orders of magnitude down to $\\lesssim$ 1\\% solar \\citep[cf.][]{gratton04}. Mass loss remains important throughout the history of a GC \\citep{jaccomc07}; even in old GCs, stars lose $\\sim$30\\% of their mass, most or all of which is removed from the cluster via a host of mechanisms including ram-pressure stripping during Galactic plane crossings. RGB mass loss may be responsible for the blue horizontal branches observed in metal-poor GCs \\citep{rood73}, and evidence for mass loss is seen in the line profiles of chromospherically active or pulsating red giants in GCs \\citep{dupree84,cacciari04,mcdonald07,meszaros08}.  There is mounting evidence for mass loss from GC red giants in the form of dusty winds \\citep{ramdani01, origlia02}. The first spectroscopic identifications of silicate dust are in the tip-AGB star V1 in the massive GC 47\\,Tucanae \\citep{jaccoea06a} and subsequent identification of minerals in other AGB stars in 47\\,Tuc with {\\it Spitzer} IRS observations \\citep{lebzelter06}. With [Fe/H]~$=-0.7$, 47\\,Tuc is a metal-rich GC, providing an environment similar to that found, for instance, in the Small Magellanic Cloud. To probe uncharted terrain, it is important to study dust in GCs at lower metallicity. The surprising detection of both interstellar and circumstellar dust and gas in the metal-poor GC M\\,15 \\citep{evans03,boyer06,jaccoea06b} demonstrates that dust forms even in extremely metal-deficient environments.  \\subsection{{\\it $\\omega$\\,Centauri}} \\label{sec:ocen}  The most massive Galactic GC, $\\omega$\\,Centauri, contains enough stars in the dusty wind phase for an empirical reconstruction of the evolution of mass loss. Its gravitational well \\citep[$v_{\\rm esc}~=~44$ km s$^{-1}$;][]{gnedin02} could retain interstellar matter if winds are slow. $\\omega$\\,Cen is located approximately 15\\degr{} out of the Galactic plane at a distance of $\\approx$5 kpc \\citep{lub02,vandeven06,reddening}, so it is possible to probe the stellar population to the cluster core.  The bulk of the stars in $\\omega$\\,Cen have [Fe/H] $=-1.7$ \\citep{norris96,smith00}, a metallicity much lower than the environments in which circumstellar dust has traditionally been studied. Complications arise because $\\omega$\\,Cen also harbors a small fraction of stars up to an order of magnitude more metal-rich \\citep{lee99, pancino00, pancino02} and an intermediate population of stars with an abnormally high helium content \\citep{norris96, norris04}. Red giants in $\\omega$\\,Cen display a range in surface abundances, from oxygen-rich, M-type stars with titanium oxides, stars enhanced in CH or in cyanide (CN), to genuine carbon stars with molecular carbon \\citep[C$_2$;][]{jacco07}.  These disparate populations cause difficulties in stellar evolution studies of $\\omega$\\,Cen, but since stars from the metal-rich and metal-poor populations can be distinguished by means of optical spectroscopy, $\\omega$\\,Cen presents us with a unique opportunity to study the dust and mass loss within a single environment as a function of atmospheric chemistry.  Consensus has yet to be reached regarding the formation history of the different subpopulations \\citep{sollima05, stanford06,villanova07}. The spread and peculiarities in elemental abundances in $\\omega$\\,Cen have been interpreted as evidence that this GC may be the remnant nucleus of a tidally disrupted dSph \\citep{zinnecker88, freeman93}. If so, $\\omega$\\,Cen may offer us insight into the mass loss and chemical enrichment in the much more distant nucleated dSphs, and perhaps dSphs in general.  Here, we present NASA {\\it Spitzer Space Telescope} \\citep{werner04, gehrz07} images of $\\omega$\\,Cen and a catalog including the magnitudes of over 40,000 point-sources from 3.6 to 70~\\micron{}. The 24~\\micron{} mosaic, which enables reliable detection of individual circumstellar dust shells near the peak wavelength of their spectral energy distributions (SEDs), provides the most dramatic improvement over previous data.  We make use of proper motion \\citep{vanleeuwen00} and radial velocity \\citep{jacco07} measurements to separate cluster members from the substantial foreground stellar and background galaxy populations. Optical spectroscopic information permits us to compare the contributions from stars of different composition and evolutionary status to their infrared (IR) emission and dust production. Our observations are discussed in \\S\\,\\ref{sec:obs}, photometry and the $\\omega$\\,Cen point-source catalog are described in \\S\\,\\ref{sec:phot}.  The color-magnitude diagrams (CMDs), stellar populations, mass loss, and the intracluster medium (ICM) dust are discussed in \\S\\,\\ref{sec:disc}, and our conclusions are presented in \\S\\,\\ref{sec:concl}. ",
        "conclusions": "} \\label{sec:concl}  Our {\\it Spitzer} multi-wavelength and dual-epoch study of the most massive Galactic globular cluster, $\\omega$\\,Cen, and the resulting point-source catalog provide the most complete mid- to far-IR atlas of any GC to date.  Despite $\\omega$\\,Cen's rather unusual stellar population, the {\\it Spitzer} CMDs show little structure beyond the division between the RGB stars belonging to the cluster and background galaxies. The HB stars and RGB-a stars are not separated from the RGB, except for a slightly blue [3.6] -- [4.5] color for the RGB-a.  Nine sources are detected at 70~\\micron{}, one of which is a resolved spiral galaxy, and two of which are spatially coincident with optical sources that are known cluster members. If emission at 5.8 and 8~\\micron{} in the SEDs of the eight point sources is PAH emission, then these sources are background galaxies.  By cross-referencing with a catalog of optical spectra \\citep{jacco07}, we confirm that super-Ba-rich stars delineate the tip of the AGB, likely due to a third dredge-up, and find that three M-type stars not only have strong red excess, but are brighter than the 24~\\micron{} tRGB. Attributing the excess IR luminosity to dust suggests that the three bright M-type stars dominate the cluster's cumulative mass-loss rate (gas $+$ dust) of 2.9 -- 4.2 $\\times$ 10$^{-7}$ $M_\\odot$ yr$^{-1}$. If the cluster mass-loss rate has remained constant since the last Galactic plane crossing, then the cluster has lost at least 1 -- 2 $M_\\odot$ of dust-traced material over the last 3.4 $\\times$ 10$^{6}$ years.  The dusty stars in our sample range over $\\approx$1.5 dex in metallicity, suggesting that dust production is not inhibited even in stars with very low metal contents. In addition, our results show that, in $\\omega$\\,Cen, significant dusty mass loss occurs near the very tip of the RGB or AGB, concentrated in only a few individual stars, with little evidence for such mass loss lower on the RGB.  While our images are sensitive enough to detect the amount of dust expected in $\\omega$\\,Cen based on the current mass-loss rate, no obvious ICM dust clouds are apparent in the cluster center at 24 or 70~\\micron{}.  We do find several regions of faint extended 24~\\micron{} emission on the outskirts of the cluster, but it is unclear whether these clouds are associated with the ICM or are in the foreground. If the dust features are associated with the cluster, they are well outside the cluster half-mass radius, and could therefore be ICM material that is in the process of escaping the cluster."
    },
    "0801/0801.4562_arXiv.txt": {
        "abstract": "An arbitrary initial magnetic field in an A star evolves into a stable equilibrium. Simulations are presented of the formation of non-axisymmetric equilibria consisting of twisted flux tubes meandering under the surface of the star, and analytic arguments are given relating the stability and form of these equilibria. These results may help to explain observations of Ap stars with very non-dipolar fields. This work is also applicable to other essentially non-convective stars such as white dwarfs and neutron stars. ",
        "introduction": "Convective stars of various types tend to have small-scale, time-varying magnetic fields, which is interpreted as dynamo action driven by differential rotation and convection/buoyancy instabilities. In contrast, intermediate-mass main-sequence stars ($1.5M_\\odot<M<6M_\\odot$), which have a small convective core and a radiative envelope, have either no detectable field or a large-scale, steady magnetic fields -- the Ap stars (see e.g. Auri\\`ere et al. 2008). Many of these stars have roughly dipolar fields but many have rather more complex fields, such as $\\tau$ Sco (Donati et al. 2006).  Given the lack of necessary dynamo ingredients, we infer that the magnetic fields are `fossil' remnants in some stable equilibrium, formed from the magnetic field left over from the convective/accretive protostellar phase. On the theoretical side, the emphasis has been on finding stable magnetic equilibria. Analytically one can produce an equilibrium configuration simply by making sure all forces are balanced, and then test its stability using an energy method (Bernstein et al. 1958). Unfortunately, it has been easier to demonstrate the {\\it instability} of various equilibria than to find {\\it stable} configurations. For instance, both axisymmetric poloidal fields and axisymmetric toroidal fields have been shown to be unstable (Wright 1973; Tayler 1973; Braithwaite 2006). More recently, it has become possible to find stable equilibria using numerical methods. Braithwaite \\& Nordlund (2006, see also Braithwaite \\& Spruit 2004) found such a configuration by evolving an arbitrary initial magnetic field in time and watching it find its way into a stable equilibrium. This equilibrium is roughly axisymmetric, consisting of both toroidal and poloidal components (logical, since both are unstable on their own) in a twisted-torus configuration. From the outside, the field looks approximately dipolar, as do the fields of many Ap stars. Here, I demonstrate the existence of more complex, non-dipolar configurations. I first describe simulations of magnetic equilibria evolving from arbitrary initial conditions, then use simple analytic methods to explore the nature and stability of these equilibria. ",
        "conclusions": "Once an equilibrium has formed, it continues to evolve as a result of magnetic and thermal diffusivities ($\\eta$ and $\\kappa$ respectively). Magnetic diffusion (a result of finite conductivity) causes the field to evolve on a timescale $\\mathcal{L}^2/\\eta$ where $\\mathcal{L}$ is the relevant length scale. For an A star, this is around $10^{10}$yr if $\\mathcal{L}=R$. Thermal diffusion can also cause field evolution. To be in pressure and density equilibrium with its non-magnetised surroundings, a flux tube must be at a lower temperature than its surroundings. Heat diffuses into the tube, causing it to rise on a timescale of order $\\mathcal{L}^2\\beta/\\kappa$, where $\\beta$ is the ratio of thermal to magnetic pressure. These effects are discussed in Braithwaite 2008 (in prep.); it is likely that both are too slow to be important in A stars, except perhaps in stars with very complex, small-scale equilibria (i.e. with smaller $\\mathcal{L}$). Whether there is any observational evidence for field evolution in A stars is controversial.  That we observe both axisymmetric and non-axisymmetric fields in Ap stars implies that initial conditions vary, and that probably the central concentration of the initial field varies. The reason for this is unclear. However, it is a common phenomenon that Ap stars share with white dwarfs and neutron stars (as is the enormous range in field strengths), while the apparent cutoff in field strengths below $300$G (Auri\\`ere et al. 2008) is unique to A stars, so the two problems are probably unconnected.  In summary, I have performed numerical simulations of the formation of stable magnetic equilibria in a non-convective star, starting from turbulent initial conditions. A sufficiently centrally-concentrated initial field evolves into an approximately axisymmetric equilibrium; a more spread-out initial field evolves into a more complex, non-axisymmetric equilibrium consisting of twisted flux tubes meandering under the surface of the star. In this case, the adjustment to equilibrium consists in the stretching or contracting of these flux tubes until an energy minimum is reached. Still needed is a more quantitative understanding of the effect of initial conditions on the resulting equilibrium, and more work is also required on the quasi-static diffusive evolution of these equilibria. The results are applicable to all predominantly non-convective stars: main sequence above $1.5M_\\odot$, white dwarfs and neutron stars."
    },
    "0801/0801.4424_arXiv.txt": {
        "abstract": "{Astrometric observations of resolved binaries provide estimates of orbital periods and will eventually lead to measurement of dynamical masses. Only a few very low mass star and brown dwarf masses have been measured to date, and the mass-luminosity relation still needs to be calibrated.} {We have monitored 14 very low mass multiple systems for several years to confirm their multiplicity and, for those with a short period, derive accurate orbital parameters and dynamical mass estimates.} {We have used high spatial resolution images obtained at the Paranal, Lick and HST observatories to obtain astrometric and photometric measurements of the multiple systems at several epochs. The targets have periods ranging from 5 to 200 years, and spectral types in the  range M7.5 -- T5.5. } {All of our 14 multiple systems are confirmed as common proper motion pairs. One system  (2MASSW~J0920122+351742) is not resolved in our new images, probably because the discovery images were taken near maximum elongation. Six systems have periods short enough to allow dynamical mass measurements within the next 15 to 20~years. We estimate that only 8\\% of the ultracool dwarfs in the solar neighborhood are binaries with separations large enough to be resolved, and yet periods short enough to derive astrometric orbital fits over a reasonable time frame with current instrumentation. A survey that doubles the number of ultracool dwarfs observed with high angular resolution is called for to discover enough binaries for a first attempt to derive the mass-luminosity relationship for very low-mass stars and brown dwarfs.} {} ",
        "introduction": "Over the last few years, intensive computational and observational efforts have been made to improve our understanding of the formation processes and evolution of brown dwarfs (BDs) and very low mass (VLM) stars. The determination of their Initial Mass Function (IMF) is a crucial step in this direction. Translating an observed luminosity function into an IMF requires an accurate determination of their mass-luminosity relationship at different ages, which up to now relies primarily on theoretical mass-luminosity  relationships. Although the empirical constraints on these relationships for VLM stars have considerably improved within the past years \\citep[see e.g][]{2004ApJ...604..741H,2000A&A...364..217D,2000A&A...364..665S} only a few observational constraints are currently available and large uncertainties remain \\citep[][]{2001A&A...367..183L, 2004A&A...423..341B, 2004A&A...428..205B, 2004ApJ...615..958Z,2005Natur.433..286C,2006Natur.440..311S,2008arXiv0801.1525I}.  The degeneracy in the mass-luminosity relation for ultra-cool dwarfs (UCDs) makes it difficult to accurately estimate their physical properties. Dynamical masses, which are not model-dependent, are a unique way to calibrate this relation.  The components of a multiple system are expected to be coeval, removing part of the above mentioned degeneracy. Although the ages of the targets studied in this work are not well constrained, it will be possible, once their dynamical masses are known, to take advantage of their coevality to test the evolutionary models. By adjusting the theoretical isochrones empirically to fit both the observed total masses and the individual luminosities of the multiple systems, it will be possible to directly check the consistency of the models with the observations. The corresponding predictions on the age can then be compared to other indicators such as the activity, the rotation, and the presence and strength of particular spectral features (such as Li, H$\\alpha$), but also to more recent techniques based on spectral analysis of gravity sensitive features as described by \\citet{2004ApJ...609..885M}, \\citet{2004ApJ...600.1020M}, \\citet{2003ApJ...593L.113M}, \\citet{2004ApJ...615..958Z} and \\citet{2006ApJ...639.1095B}. Finally, by studying the physical characteristics of objects with known dynamical masses, it will be possible to provide crucial information for our understanding of their physical properties, such as their interior structure, the formation of dust, the settling and depletion of refractory elements, and the underlying opacities. An accurate determination of the mass of an object based on dynamical masses in binary systems therefore provides not only a reality check for the theory but also a cornerstone in the understanding of the mass distribution of brown dwarfs.  In this work, we present a time-series of high angular resolution observations aimed at monitoring binary ultra-cool dwarfs. These observations confirm the common proper motion of the binary candidates and represent a first step towards the derivation of orbital parameters and dynamical masses. Most objects presented here were monitored over timescales too short in comparison with their periods, allowing us to estimate rough orbital periods, but preventing us from obtaining detailed orbital fits. ",
        "conclusions": "} We present astrometric and photometric results of follow-up observations of 14 UCD binaries. Only half of them are rotating fast enough to provide accurate dynamical masses within the next 15-20 years. The HST, but also the recently commissioned Laser Guide Stars for Adaptive Optics on 8~m class telescopes should allow to discover and follow more UCD binaries, usually too faint and too red even for the IR-WFS of NACO. Some targets not included in the present sample are already part of other on-going programs, and more follow-up observations are likely to be published in the coming months/years. We are currently closely monitoring three additional targets for which dynamical masses will be derived within one year (Bouy et al., in prep.). Another two \\citep[$\\epsilon-$Indi~Bab and GJ~1001BC, respectively ][]{2004A&A...413.1029M, 2007IAUS..240..329G} are the targets of additional monitoring programs. The total number of \"short\" period VLM  multiple systems (short meaning periods allowing dynamical mass measurements within 15--20~yr) roughly adds up to a dozen of objects, which has been extracted from original samples of UCDs made of $\\approx$140 objects \\citep{2003AJ....126.1526B,2003ApJ...587..407C,2003AJ....125.3302G}, i.e. the frequency of short-period resolved binaries is about 8\\%. If we consider that about 20 binaries (40 masses) are required in order to start calibrating the mass-luminosity relationship, the current study shows that we would need to observe a total of roughly 140/12$\\times$20=250 UCDs at high spatial resolution. This estimate means that another survey of about 140 more UCDs is needed to discover enough binaries that can yield dynamical masses in the near future for a calibration of the mass-luminosity relationship. Even more dynamical masses will be required to extend the study of UCD physical properties to additional parameters, such as age, gravity, and  metallicity. The study of UCDs would therefore greatly benefit from new high spatial resolution surveys dedicated to searching for new multiple systems, and from complementary monitoring programs targeting the shortest period binaries."
    },
    "0801/0801.1806_arXiv.txt": {
        "abstract": "Red supergiants (RSGs) are an evolved stage in the life of intermediate massive stars ($<25M_\\odot$). For many years, their location in the H-R diagram was at variance with the evolutionary models. Using the MARCS stellar atmospheres, we have determined new effective temperatures and bolometric luminosities for RSGs in the Milky Way, LMC, and SMC, and our work has resulted in much better agreement with the evolutionary models. We have also found evidence of significant visual extinction due to circumstellar dust. Although in the Milky Way the RSGs contribute only a small fraction ($<1$\\%) of the dust to the interstellar medium (ISM), in starburst galaxies or galaxies at large look-back times, we expect that RSGs may be the main dust source. We are in the process of extending this work now to RSGs of higher and lower metallicities using the galaxies M31 and WLM. ",
        "introduction": "Those of us here who have worked on massive stars for a while are probably all attracted by stellar physics at the extremes.  For O-type stars, we are dealing with stars that are as massive and luminous as stars come, and as hot as main-sequence stars can get.  To properly assess their physically properties through spectroscopic analysis has required not only the introduction of non-LTE atmosphere models (\\cite[Mihalas \\& Auer 1970]{MihAue70}, \\cite[Auer \\& Mihalas 1972]{AueMih72}) but an additional thirty years of developments, such as the inclusion of mass loss (\\cite[Abbott \\& Hummer 1972]{AbbHum72}, \\cite[Kudritzki 1976]{Kud76}), the inclusion of hydrodynamics of the stellar wind (\\cite[Gabler, Gabler, Kudritzki, et al.\\ 1989]{Gab89}; \\cite[Kudritzki \\& Hummer 1990]{KudHum90}; \\cite [Puls, Kudritzki, Herrero, et al.\\ 1996]{Puls96}), and, most recently, the full inclusion of line blanketing (\\cite [Hillier \\& Miller 1989]{HilMil89}; \\cite[Hillier, Lanz, Heap, et al.\\ 2003]{Hil03}; \\cite[Herrero, Puls, \\& Najarro 2002]{Her02}). (For a recent summary, see \\cite[Massey, Bresolin, Kudritzki, et al.\\ 2004]{OPapI} and \\cite[Massey, Puls, Pauldrach, et al.\\ 2005]{OPapII}).  The study of LBVs and WRs is equally exciting, stars where radiation pressure dances with gravity to see who will lead (\\cite[Lamers 1997]{Lam97}, \\cite[Smith \\& Owocki 2006]{SmiOwo06}), and where high mass loss rates are continuous rather than episodic due to high metal content in the stellar atmosphere (\\cite[Crowther 2007 and references therein]{Cro07}).  However, largely ignored until now are the red supergiants (RSGs).  The physical conditions in these stars are, in their own way, equally extreme.  They have the largest physical sizes of any stars (up to 1500$\\times$ the radius of the sun; see \\cite [Levesque, Massey, Olsen, et al.\\ 2005, here after Paper I] {RSGPapI}).  This large physical size invalidates the usual assumptions of plane parallel geometry.  The velocities of the convective layers in these stars' atmospheres are supersonic, giving rise to shocks (\\cite [Freytag, Steffen, \\& Dorch 2002]{Fre02}), and making the stars' photospheres very asymmetric and invalidating mixing-length assumptions. Their extremely cool temperatures (3400 - 4300~K) lead to the the presence of molecules in their atmospheres, requiring the inclusion of extensive molecular opacity sources in any realistic model atmosphere.  From an observational point of view, the large (negative) bolometric corrections and their sensitivity to the adopted temperature complicate the transformation from the observed color-magnitude diagram to the physical H-R diagram in much the same way as it does for the O-type stars.  Recent advances in stellar atmosphere models for cool stars (e.g., \\cite[Plez 2003]{Plez03}) have allowed the first reasonable determination of the physical properties of these stars, in much the same way that the  non-LTE H and He models of \\cite{AueMih72} allowed the first reasonable determiation of the physical properties of O-type stars by \\cite{Con73}.  And while we may still have a way to go, we believe our answers will hold up as well as those that \\cite{Con73} have, which is really pretty well (see \\cite[Massey, Puls, Pauldrach, et al.\\ 2005] {OPapII}).  I find it personally interesting that there has been a real aversion to looking at what happens to a massive star as it heads over to the far right side of the H-R diagram. I think this is cultural---for many years much of the ``massive star community\" really thought of itself as the ``hot star community\", with the exception of a few workers, most notably our good colleague Roberta Humphreys, whose early work on supergiants in the Milky Way and other Local Group galaxies (such as \\cite[Humphreys 1978] {Hump78}, \\cite[1979a]{Hump79a}, \\cite[1979b]{Hump79b}, \\cite[1980a]{Hump80a}, \\cite[1980b]{Hump80b}, \\cite[1980c]{Hump80c}, \\cite[Humphreys \\& Davidson 1979] {HumpDavid79}, \\cite[Humphreys \\& Sandage 1980] {HumpSand80}) certainly spurned my own interest in the field, and whose presence at these  symposia always reminds us that there's more to the life of a massive star than the O and WR stages.  Let us briefly review what the evolutionary tracks predict RSGs come from. In Fig.~\\ref{fig:HRDs} we show the tracks covering a range of a factor of 10 in metallicity, from z=0.004 (SMC-like) to z=0.020 (solar) to z=0.040 (M31-like). I have drawn in a vertical line at an effective temperature of 4300~K, which roughly corresponds to that of a K0~I. Stars to the right of that line we are calling RSGs.   At solar metallicities we expect that stars with initial masses $\\leq 25M_\\odot$ will become RSGs.  At lower metallicities (SMC-like) the upper mass limit for RSGs is probably a bit higher---maybe $30M_\\odot$?---it's hard to tell because of the quantization of the tracks. The upper mass tracks go much further to the right at this low metallicity, but stop short of the RSG dividing lines---these 30-60$M_\\odot$ stars become F- and G-type supergiants, but not K or M.  At higher metallicity (M31) the limit is definitely lower, around $20M_\\odot$.   In the case of solar metallicity the $25M_\\odot$ track turns back to the blue, and in fact such a star should become a WR after the RSG phase.  One more thing of note is that the tracks don't extend very far to the right of the K0~I (vertical) line in the SMC---the RSGs in the SMC shouldn't be very late, mostly K through M0, say.  At higher metallicities they extend further to cooler effective temperatures.  This is consistent with the change in the average RSG type observed in the SMC, LMC, and Milky Way (\\cite[Elias, Frogel, \\& Humphreys 1985] {EFH85}, \\cite[Massey \\& Olsen 2003]{MO03}).  \\begin{figure} \\begin{center} \\includegraphics[width=2.65in]{masseyf1a.eps} \\end{center} \\includegraphics[width=2.65in]{masseyf1b.eps} \\includegraphics[width=2.65in]{masseyf1c.eps} \\caption{\\label{fig:HRDs}Effects of metallicity on the evolutionary tracks in the RSG region. The tracks for z=0.020 (solar) are from \\cite{MM03}, for z=0.004 (SMC-like) are from \\cite{MM00}, and for z=0.040 (M31-like) are from \\cite{MM05} and \\cite{MM94}. Solid curves denote the tracks with no initial rotation, while the dashed lines correspond to initial rotations of 300 km s$^{-1}$.  The black vertical line marks a temperature of 4300 K, roughly that of a K0~I at both solar and SMC metallicity (Papers I, II).} \\end{figure}  That said, when we began worrying about the issue of the physical properties of RSGs it was because if one placed the ``observed\" location of RSGs on the H-R diagram, they missed the tracks entirely: the alleged effective temperatures and luminosities were cooler and higher than those predicted by the evolutionary tracks. This was first noticed by \\cite{MO03} for the SMC and LMC, but we quickly confirmed that the problem also existed for the Milky Way sample (\\cite[Massey 2003a]{MasARAA}). When I mentioned this issue at the Lanzarote meeting (\\cite [Massey 2003b]{MasIAU212}) Daniel Schaerer came up to me afterwards and said really, this was not a problem, since how far to the right the tracks went were heavily dependent upon such issues as how the mixing length was treated (see, for example, \\cite[Maeder \\& Meynet 1987]{MM07}).  However, this did not explain the issue of the luminosities being too high, and as an observer I was more concerned with what if the ``observations\" were wrong?  Because, of course we don't ``observe\" effective temperatures and bolometric luminosities; instead, we obtain photometry and spectroscopy and use some relationship to convert these to physical properties.  Indeed, further reading convinced me that there could be a serious problems, as much of what we ``knew\" about the effective temperature scale of RSGs were derived from lunar occultations of red {\\it giants} (not supergiants); see discussion in \\cite{MO03}.  What would be involved in determining the effective temperatures of RSGs ``right\"? We really need models that have enough physics in them to correctly reproduce temperature-sensitive spectral features. The participants at this conference (mostly) understand what was involved in getting there for O-type stars.  RSGs present their own challenges, as noted above.  Fortunately, at the time I got intrigued by this problem, sophisticated models that were up to the task were becoming available. A modern version of these MARCS models was described by \\cite{Plez92}, based upon the earlier work of \\cite{Gust75}.   These are static, LTE, opacity-sampled models, and the current version (\\cite[Plez 2003]{Plez03}; \\cite[Gustafsson, Edvardsson, Eriksson, et al.\\ 2003]{Gust03}) includes improved atomic and molecular opacities and sphericity. ",
        "conclusions": ""
    },
    "0801/0801.3803_arXiv.txt": {
        "abstract": "{\\axj\\ is an unclassified transient X-ray source discovered during surveys by \\asca\\ in 1993--1999. The transient behaviour and the short and bright flares of the source have led to the idea that it is part of the recently revealed subclass of supergiant fast X-ray transients.} {A multi-wavelength study in NIR, optical, X-rays, and hard X-rays of \\axj\\ is undertaken in order to determine its nature.} {Public \\integral\\ data and our target of opportunity observation with \\xmm\\ were used to study the high-energy source through timing and spectral analysis. Multi-wavelength observations in optical and NIR with the ESO/NTT telescope were also performed to search for the counterpart.} {\\axj\\ is a new high-mass X-ray binary pulsar with an orbital period of $185.5\\pm1.1$~days (or $185.5/f$ with $f=2,3\\ \\mathrm{or}\\ 4$) and a spin period of $\\sim 66$~s, parameters typical of a Be/X-ray binary. The outbursts last $\\sim12$~d. A spin-down of $\\dot{P}=0.08\\pm0.02\\ \\mathrm{s}\\,\\unit{yr}{-1}$ is also observed, very likely due to the propeller effect. The most accurate X-ray position is R.A.~(2000)~$=\\ra{17}{49}{06.8}$ and Dec.~$=\\dec{-27}{32}{32.5}\\ (\\mathrm{uncertainty}\\ 2\\arcsec)$. The high-energy broad-band spectrum is well-fitted with an absorbed powerlaw and a high-energy cutoff with values $\\nh=20.1_{-1.3}^{+1.5}\\times10^{22}\\ \\unit{cm}{-2}$, $\\Gamma=1.0_{-0.3}^{+0.1}$, and $\\Ec=21_{-3}^{+5}$~keV. The only optical/NIR candidate counterpart within the X-ray error circle has magnitudes of $R=21.9\\pm0.1$, $I=20.92\\pm0.09$, $J=17.42\\pm0.03$, $H=16.71\\pm0.02$, and $K_{\\mathrm{s}}=15.75\\pm0.07$, which points towards a Be star located far away ($> 8.5$~kpc) and highly absorbed ($\\nh\\sim 1.7\\times10^{22}\\ \\unit{cm}{-2}$). The average 22--50~keV luminosity is $0.4-0.9\\times10^{36}\\ \\es$ during the long outbursts and $3\\times10^{36}\\ \\es$ during the bright flare that occurred on MJD~52891 for an assumed distance of 8.5~kpc.} {} ",
        "introduction": "High-mass X-ray binaries (HMXB) are systems composed of a compact object (either a neutron star (NS) or a black hole (BH)) and a massive companion star. The emission in X-rays is explained by the accretion of the stellar wind coming from the companion and captured by the compact  object. There are mainly two classes of HMXB, depending on the nature of the massive primary star, which has important implications for the X-ray observational features. One might read \\citet{Whiteal95} and \\citet{Psaltis06} (and references therein) for reviews of HMXB.  The first class is composed of persistent high-energy binary systems made of a supergiant star (spectral type OB) (later called SGXB) orbiting close to the compact object. The other group is composed of a Be star accompanying the compact object, mostly an NS. A Be star is a fast-rotating B star possessing a dense and slow stellar outflow at the equator (thus creating a disc-like ring of matter surrounding the B star). The group of BeXB is the most numerous in the HMXB family \\citep{Liual00, Liual06}. They possess wide orbits ($P_{\\mathrm{orb}}>15$~days). They display either recurrent outbursts near the periastron passage whose duration is normally linked to the orbital period (most common outbursts), or giant outbursts that can last from weeks to months. In most cases, their X-ray emission is modulated by the strong magnetic field of the NS. Their spectra are similar to those of supergiant X-ray binaries \\citep[see reviews of][]{Bildstenal97,Coe00,Negueruela07}.  Among the new sources, \\integral\\ might have unveiled a new subclass of SGXB. Indeed, several newly discovered sources were identified as galactic X-ray sources with a supergiant companion, and they display transient behaviour \\citep[e.g. IGR~J17544$-$2619,][]{Zand05, Pellizzaal06}. The quiescent level in these systems is of the order of $10^{32-33}$~ergs$\\,\\mathrm{s}^{-1}$ \\citep[e.g. IGR~J08408$-$4503,][]{Gotzal07,Leyderal07}. The X-ray luminosity increases up to $10^{36}$~ergs$\\,\\mathrm{s}^{-1}$ (as observed in the known SGXB) only during short luminous-flare activity. The X-ray luminosity remains at a very low level (if not totally absent) during most of the time, except during the flares. These flaring periods last for a few hours at most, and then the source goes back to an undetectable level of emission \\citep[e.g. XTE~J1739$-$302,][]{Smithal06}. Therefore, they received the name of Supergiant Fast X-ray Transient (SFXT) \\citep{Negueruelaal06}. Except for the transient nature of SFXT, their spectral features are similar to the previously known persistent SGXB.  The unclassified transient X-ray source AX J1749.1-2733 was discovered at a very low-luminosity level during the ASCA galactic centre survey performed between 1993 and 1999 \\citep{Sakanoal02}. Its position is R.A.~(2000)~$=\\ra{17}{49}{09.0}$ and Dec.~$=\\dec{-27}{33}{13.9}$ (50$\\arcsec$ at 90\\% confidence). The source was detected three times over the six observations when the source was located within the instrument field of view. The observed 0.7--10~keV flux varies between $(2-6)\\times 10^{-12}\\ \\ecms$. Each spectrum is fitted with an absorbed power-law model. However, only in one observation are the parameters well-constrained with $\\Gamma=2.1_{-2.6}^{+4.7}$ and $\\nh=25_{-21}^{+57}\\times10^{22}\\ \\unit{cm}{-2}$.  A new flare of this source was observed by INTEGRAL on Sept. 9, 2003, lasting 1.3~days \\citep{Sgueraal06}. The 20--60~keV peak flux was 40~mCrab, and they fitted the spectrum with a combination of a black body and a power-law ($KT=0.7_{-0.1}^{+0.3}\\ \\mathrm{keV}$ and $\\Gamma=2.5\\pm0.2$). Due to the shortness of the flare, they classified this source as a candidate SFXT. This tentative classification has already been pointed out by \\citet{Zand05} and references therein. \\citet{Grebeneval07} discussed one possible mechanism that could trigger such bright and fast flare in this system. If the source is a supergiant wind-fed accreting pulsar, the propeller mechanism would inhibit the accretion when the neutron star's spin period is less than the critical value $P_{\\mathrm{s}}^{*}\\sim3\\,\\mathrm{s}$. However, an increase in the local stellar wind density of the order of $(P_{\\mathrm{s}}^{*}/P_{\\mathrm{s}})^{7/3}$ during a change in the stellar outflow would result in the observed flare. \\citet{Kuulkersal07} also report \\integral\\ detections in the periods Feb.--Apr. in 2005 and 2006 and non detection in 2005 Aug.--Oct.  with an average 20--60~keV flux of $<2$~mCrab in both detections\\footnote{see the public page of the \\integral\\ Galactic Bulge program at {\\tt http://isdc.unige.ch/Science/BULGE} for further information.}. Analysing all \\integral\\ public data and using the data available on the public page of the \\integral\\ Galactic Bulge program$^{1}$, \\citet{Zuritaal07} discovered a likely period of $\\sim185$~days as the source was detected by \\integral\\ at a level of $\\gtrsim 5 \\sigma$ for a few days  during March and September each year between 2003--2007.  Following this announcement, \\swift\\ performed a target of opportunity (ToO) observation of 5~ks on March 30, 2007 \\citep{Romanoal07a}\\footnote{From the same observation, \\citet{Kong07} also report a slightly more accurate position with R.A.~(2000)~$=\\ra{17}{49}{06.8}$ and Dec.~$=\\dec{-27}{32}{31.5}$ (3.8$\\arcsec$ at 90\\% confidence). He also fitted the spectrum with an absorbed power-law with $\\nh=(19\\pm9)\\times10^{22}\\ \\unit{cm}{-2}$, $\\Gamma=2.1\\pm1.2$, and an absorbed 2--10~keV flux of $\\sim 6\\times 10^{-12}\\ \\ecms$.}. They detected a bright source at the position R.A.~(2000)~$=\\ra{17}{49}{06.8}$ and Dec.~$=\\dec{-27}{32}{30.6}$ (6.3$\\arcsec$ at 90\\% confidence), $51\\arcsec$ away from the \\asca\\ position\\footnote{The distance was wrongly given in the ATel as their \\asca\\ position does not correspond to the one published in \\citet{Sakanoal02}.}. Its spectrum could be fitted with an absorbed power-law ($\\nh=23_{-10}^{+14}\\times10^{22}\\ \\unit{cm}{-2}$, $\\Gamma=2.5_{-1.7}^{+2.0}$) and the observed  2--10~keV flux is $\\sim10^{-11}\\ \\ecms$. The \\swift/UVOT telescope did not detect any optical counterpart with a $3\\sigma$ upper limit of $V=20.67$~mag. Within the error circle, \\citet{Romanoal07a} report three 2MASS candidate counterparts whose infrared (IR) magnitudes in the JHK bands suggest a strong optical extinction of 20~mag. Only one candidate is compatible with a supergiant. Thus, the nature of \\axj\\ remains a mystery.  \\xmm\\ also performed a ToO observation of the source during the expected outburst of 2007 March. From this observation, \\citet{Karaseval07,Karaseval08} reported the discovery of a pulsation of $\\sim132$~s. The pulse profile displays a double-peak structure with a pulse fraction of $\\sim30$\\% in the 2--10~keV energy range. They also detected a pulsation during the outburst detected by \\integral\\ on Sept. 8--10, 2003 in the 20--60~keV energy range with a higher pulse fraction of 50\\%.  Here we report multiwavelength observations performed on \\axj\\ with \\integral, \\xmm, and the ESO/NTT telescope. In Sect.~\\ref{secObs}, we first describe the observations and the data analysis of each instrument. Then, we present the results in Sect.~\\ref{secRes}. Finally, we finish with a discussion and the conclusion on the nature of the source in Sect.~\\ref{secDis}. ",
        "conclusions": "\\label{secDis}  The source possesses a long orbital period of $(185.5\\pm1.1)/f$~d where $f=1,2,3,\\ \\mathrm{or}\\ 4$ and a spin period of $\\sim 66$~s. At the likely periastron passage, the source goes through an outburst whose duration is estimated to $\\sim 12$~d (with an upper-limit of $\\lesssim 18$~d). With such orbital period ($f=1$) and outburst duration, the source is bright $<10$\\% of its time. The recurrence of such long outbursts related to the binary motion corresponds to type I outbursts observed in other BeXB \\citep{Coe00, Negueruela07}. Furthermore, combining the observed pulsation and orbital period, \\axj\\ falls well in the BeXB arm of the Corbet diagram \\citep[{\\it i.e.} $P_{\\mathrm{spin}}$ vs $P_{\\mathrm{orb}}$;][]{Corbet84,Corbet86} strengthening its identification as a BeXB. The large orbital period rules out the possibility to associate this source with a low-mass X-ray binary whose companion is a star of the main sequence. Flaring activity in BeXB has never been observed except in EXO~2030$+$375 and recently in SWIFT~J1626.6$-$5156 \\citep[][and references therein]{Reigal08}. For SWIFT~J1626.6$-$5156, the intensity varied by a factor 4 on a timescale of 450~s with sharp rise/decay of the count rate. This behaviour differs from the bright flare of \\axj\\ detected with ISGRI in revolution {\\tt 110}, which shows a smooth rate's rise/decay, a variability factor of 8, and a duration of 0.5~d. This bright flare in revolution {\\tt 110} is likely to be due to a sensitivity bias where only the peak flux of the outburst is significantly detected in individual pointings. However, flaring activity cannot be discarded for \\axj\\ because of the detections in pointings {\\tt 01100038} and {\\tt 01730024}, where the rate varies by a factor 3--4 on a timescale of 1~h in comparison to the average rates during the outbursts. Still, these detections may also be due to a sensitivity bias, because the count rates are similar to the 5.1$\\sigma$ detection threshold.  A spin-down of \\axj\\ is observed between the 2 observations of the bright flare with \\integral\\ and the \\xmm\\ observation. We derived a spin-down of $\\dot{P}=0.08\\pm0.02\\ \\mathrm{s}\\,\\unit{yr}{-1}$ ({\\it i.e.} $\\dot{\\nu}=(-6.2\\pm1.4)\\times 10^{-13}\\ \\mathrm{Hz}\\,\\unit{s}{-1}$). Such spin-down episodes have already been observed in other accreting pulsars, some examples being EXO~2030$+$375 with $\\dot{\\nu}=-3.4\\times 10^{-14}\\ \\mathrm{Hz}\\,\\unit{s}{-1}$ and A0535$+$26 with $\\dot{\\nu}=-2.2\\times 10^{-13}\\ \\mathrm{Hz}\\,\\unit{s}{-1}$ \\citep[][and references therein]{Bildstenal97}. To have accretion onto an NS, the magnetic radius $R_{\\mathrm{mag}}$ must be lower than the co-rotation radius $R_{\\mathrm{cor}}$. Otherwise, the infalling matter is centrifugally expelled by the magnetic field \\citep[the propeller regime,][]{Illarionoval75}. For an NS ($M=1.4\\ \\Ms$, $R=10$~km) with a spin period of 66~s, at the equilibrium point $R_{\\mathrm{mag}}=R_{\\mathrm{cor}}$, one estimates the magnetic field to be $B=6.5\\times10^{12}\\ \\mathrm{G}\\times (\\dot{M}/10^{-10}\\,\\Ms\\,\\unit{yr}{-1})^{1/2}$. The magnetic radius depends on the accreted matter flow as $R_{\\mathrm{mag}}\\propto \\dot{M}^{-2/7}$. Thus, the magnetic radius may exceed the co-rotation radius if the accretion rate becomes low enough, stopping the accretion. The infalling matter, expelled by the propeller effect, can then remove some angular momentum and decrease the spin period of the neutron star. With such a long orbital period, the NS is in fact usually far from the dense accreting zone, where such effects leading to a decrease in the spin period can occur.  The broad X- and hard X-ray spectrum of \\axj\\ is similar to other HMXB pulsars \\citep{Whiteal95}. It is well-fitted with an absorbed power law and a high-energy cutoff with values $\\nh=20.1_{-1.3}^{+1.5}\\times10^{22}\\ \\unit{cm}{-2}$, $\\Gamma=1.0_{-0.3}^{+0.1}$, and $\\Ec=21_{-3}^{+5}$~keV. The large absorption is in good agreement with the previous observation \\citep[\\asca, \\swift;][]{Sakanoal02,Romanoal07a,Kong07}; still, it is strongly constrained with the EPIC/pn data. This large absorption is an order of magnitude higher than the galactic absorption expected in the line of sight of \\axj, $\\nh(\\mathrm{HI})=1.7\\times10^{22}\\ \\unit{cm}{-2}$ \\citep[interpolated from the HI map of][]{Dickeyal90}. Thus, the X-ray source is intrinsically absorbed. The combination of the EPIC/pn and IBIS/ISGRI spectra allowed us to better constrain the continuum getting a harder spectrum with a $\\Gamma$ that decreases by a factor 2 in comparison with former estimates and the necessity of the high-energy cutoff.  \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{cmd_axj1749_new.pdf}} \\caption{Colour-magnitude diagram. The magnitudes here correspond to absolute magnitude. The location of the source \\#5 in the diagram is shown for  the galactic $\\nh$ and several distances (1, 8.5 and 16 kpc). The arrows simulate the colour shift of $J-K\\leq-0.5$ due to a disc-like ring of matter surrounding the star, assuming that it is a Be star.} \\label{ima_cmd} \\end{figure} Only one candidate counterpart is located within the most accurate X-ray position of \\xmm\\ (see Sect.~\\ref{secPos}). Using the relation of \\citet{Predehlal95} that links the X-ray absorption and the optical extinction, $\\nh\\,(\\unit{cm}{-2})/\\Av\\,(\\mathrm{mag})=1.79\\times 10^{21}$, the expected optical extinction $A_{\\mathrm{V}}$ is 9.5~mag assuming the reported galactic $\\nh=1.7\\times10^{22}\\ \\unit{cm}{-2}$. Then, the NIR extinctions are estimated as $A_{\\mathrm{J}}=2.680$,  $A_{\\mathrm{H}}=1.663$ and $A_{\\mathrm{K}}=1.064$ using $A_{\\mathrm{J}}/\\Av=0.282$,  $A_{\\mathrm{H}}/\\Av=0.175$ and $A_{\\mathrm{K}}/\\Av=0.112$ \\citep{Riekeal85}. We derive the absolute magnitude as $M_{\\mathrm{abs}}=m_{\\mathrm{app}}-5\\,log(\\mathrm{distance}/10 \\mathrm{pc})-\\Av$. Assuming a distance of 8.5~kpc, we obtain $M_{\\mathrm{J}}=0.09$, $M_{\\mathrm{H}}=0.39$ and $M_{\\mathrm{K}}=-0.05$\\footnote{The \\Ks\\ magnitude is transformed into the $K$ magnitude using the relation $K-K_{\\mathrm{s}}=-0.005(J-K)$ (SOFI user manual, p.7).}, so the colour is $(J-K)_{0}=0.14$. Thus, in a colour-magnitude diagram (CMD), the candidate companion star of \\axj\\ falls near the zone of the main sequence stars with spectral type between B5 and B8--A0 \\seefig{ima_cmd}. In this case of a BeXB system, the disc-like ring of matter surrounding the B star can modify its colour $J-K$ by a large factor \\citep[e.g. up to $-0.5$~mag in 4U~0115+63][]{Reigal07}. This is illustrated in Fig.~\\ref{ima_cmd} by arrows of maximum length of 0.5 mag that strengthen the scenario of a B star located far away in the Galaxy, thus highly absorbed by the interstellar medium. In another case where the star would be located nearer to us ($\\sim 1$~kpc), thus with a lower extinction, the star would be a main sequence star of type KM. In that case, this star would probably not be the counterpart of \\axj, because the large orbital period completely rules out the possibility of combining such KM star with this X-ray source. The possibility that the star is a red giant is also unlikely as it should combine a location far away in the Galaxy with low extinction. Spectroscopy of the candidate counterpart is needed to disentangle between the 2 possibilities, either a B star located far away in the Galaxy $\\gtrsim8.5$~kpc and strongly absorbed or a normal star located nearer to us. Still, because of the X-ray properties observed for \\axj\\ and the absence of any other candidate, the B star is favoured. As the known BeXB have an earlier spectral type than B3 \\citep{Negueruela98}, this would imply that the source is located at a very long distance of $>8.5$~kpc. Another argument in favour of the BeXB scenario is the reddening-free parameter $Q=(J-H)-1.70(H-K_{\\mathrm{S}})$ \\citep{Negueruelaal07a}. For candidate \\#5, $Q\\sim-0.9\\pm0.1$, thus suggesting a high IR excess. Such excess is not expected in KM stars whose $Q$ values lie between 0.4--0.5; however, a factor $Q<0$ is expected for Be stars. Assuming a distance of 8.5~kpc, the 22--50~keV luminosity is $3\\times10^{36}\\ \\es$ during the bright flare that occurred on MJD~52891 and $0.4-0.9\\times10^{36}\\ \\es$ during the long outbursts, values typical of BeXB. The 0.2--10~keV luminosity is $0.2\\times10^{36}\\ \\es$, a factor 2--4 lower than in \\integral. We point out that, in the case of a BeXB with typical outbursts of the order of $\\sim10^{36}\\ \\es$, and if the system was located nearer to us ($\\lesssim 8$~kpc), \\integral\\ would detect such a source with a 22-50~keV count rate $\\gtrsim 2$~\\cps, which is not the case.  In the case of a BeXB, \\axj\\ shows a high intrinsic absorption of $\\nh=20.1_{-1.3}^{+1.5}\\times10^{22}\\ \\unit{cm}{-2}$ that is quite uncommon in other BeXB. Indeed, most of the BeXB have density columns closer to the galactic absorption with $\\nh\\sim 1-3\\times N_{\\mathrm{H}} ^{\\mathrm{Gal}}$ \\citep[see Fig.~15 of][]{Bodagheeal07}. Only one BeXB exhibits such high absorption: 2S~1845$-$024. The properties of this source are very similar to \\axj\\ with $P_{\\mathrm{orb}}=242.18\\pm0.01$~d, outburst duration of $\\sim 13$~d, $P_{\\mathrm{spin}}=94.9$~s (with episodes of spin-up and spin-down) and $\\nh=25\\pm10\\times10^{22}\\ \\unit{cm}{-2}$ \\citep[][and references therein]{Koyamaal90,Soffittaal98,Fingeral99}. They classified 2S~1845$-$024 as a BeXB (located at 10~kpc\\footnote{This value can be considered as an upper limit as they derive it taking the minimum observed $\\nh$ as interstellar absorption without taking the intrinsic absorption into account.}) from its X-ray properties; however, no counterpart has been identified yet. Thus, \\axj\\ might be the first identified counterpart BeXB with high intrinsic absorption similar to the obscured SGXB revealed by \\integral. Spectral properties of pulsars, either SGXB or BeXB, are very similar \\citep[see e.g.][]{Whiteal95}. As IBIS/ISGRI is not affected by absorption, \\integral\\ allowed a large population of obscured SGXB to be revealed with $\\nh$ values of $\\gtrsim10^{23}\\ \\unit{cm}{-2}$. That such highly-absorbed BeXB have not been observed yet \\citep{Bodagheeal07} is perhaps not so surprising since BeXB are mostly transient sources, and the sources must also be observed during the outbursts by another X-ray emission with high sensitivity at low energies (such as \\xmm\\ or \\chandra) in order to constrain the absorption. Such a population of obscured BeXB might still have to be revealed. Indeed, the only BeXB among the \\integral\\ sources reported in Fig.~15 of \\citet{Bodagheeal07} is the 3rd highly-absorbed BeXB (after \\axj\\ and 2S~1845$-$024).  All the X-ray properties observed in \\axj\\ lead to classifing this object as an X-ray binary, most probably an HMXB with a Be companion star whose compact object is an NS. The only optical/NIR candidate counterpart located inside the best X-ray error circle is compatible with a B star located far inside the galaxy and suffering large extinction. Besides, the lack of a supergiant companion and the duration/smoothness of the outbursts rule out the classification of this source as an SFXT, as proposed in \\citet{Zand05} and \\citet{Sgueraal06}. Instead, \\axj\\ seems to be an obscured BeXB located far away in the Galaxy."
    },
    "0801/0801.4338_arXiv.txt": {
        "abstract": "{  The $BeppoSAX$ archive is currently the largest reservoir of high sensitivity simultaneous soft and hard-X ray data of Seyfert galaxies. From this database all the Seyfert galaxies (105 objects of which 43 are type I and 62 are type II) with redshift lower than 0.1 have been selected and analyzed in a homogeneous way (Dadina 2007). Taking advantage of the broad-band coverage of the $BeppoSAX$ MECS and PDS instruments ($\\sim$2-100 keV), the X-ray data so collected allow to infer the average spectral properties of nearby Seyfert galaxies included in the original sample and, most notably: the photon index ($\\Gamma$$\\sim$1.8), the high-energy cut-off (Ec$\\sim$290 keV), and the relative amount of reflection (R$\\sim$1.0). The data collected have been also used to test some basic assumptions of the unified scheme for the active galactic nuclei. The distributions of the isotropic indicators used here (photon index, relative amount of reflection, high-energy cut-off and narrow FeK$\\alpha$ energy centroid) are similar in type I and type II objects while the absorbing column and the iron line equivalent width significantly differ between the two classes of active galactic nuclei with type II objects displaying larger columns (N$_{H}$$\\sim$3.7$\\times$10$^{22}$ and 6.1$\\times$10$^{23}$ cm$^{-2}$ for type I and II objects respectively) and equivalent width (EW$\\sim$220 and 690 eV for type I and II sources respectively). Confirming previous results, the narrow FeK$\\alpha$ line is consistent, in Seyfert 2, with being produced in the same matter responsible for the observed obscuration. These results, thus, support the basic picture of the unified model. Moreover, the presence of a X-ray Baldwin effect in Seyfert 1 has been here measured  using the 20-100 keV luminosity (EW$\\propto$L(20-100)$^{-0.22\\pm0.05}$). Finally, the possible presence of a correlation between the photon index and the amount of reflection is confirmed thus indicating thermal Comptonization as the most likely origin of the high energy emission for the active galactic nuclei included in the original sample. ",
        "introduction": "The high energy emission from active galactic nuclei (AGN) is thought to come from the innermost  regions of accretting systems that are centered around super-massive black-holes (SMBH). For this reason, X-rays are expected to be tracers of the physical conditions experienced by matter before disappearing into SMBH. Moreover, thanks to their high penetrating power, energetic photons, escaping from the nuclear zones, test the matter located between their source and the observer. Thus, they offer powerful diagnostics to understand the geometry and the physical conditions of the matter surrounding the SMBH.  The broad-band of $BeppoSAX$ offered for the very first time the opportunity to measure with a remarkable sensitivity, the spectral shape of AGN in the 0.1-200 keV range. This potential has been previously exploited to study in details a number of sample selected in different manners (see for example Maiolino et al. 1998; Perola et al. 2002). These studies were fundamental in making important steps forward in the comprehension of the emitting mechanism at work in the production of X-rays (Perola et al. 2002) and to partially unveil the geometry of the cold matter surroundings the central SMBH (Maiolino et al. 1998; Bassani et al. 1999; Risaliti et al. 1999).  The $BeppoSAX$ database full potential, however, was never exploited before. In a previous paper, the entire catalog of the Seyfert galaxies at z$\\leq$0.1 contained in the $BeppoSAX$ archive has been presented (Dadina 2007). This sample was selected starting from the catalog of Seyfert galaxies contained in the V{\\'e}ron-Cetty \\& V{\\'e}ron (2006) sample of AGN and contains 13 radio-loud objects and 7 narrow-line Seyfert 1.  The spectral analysis was performed in the 2-100 keV band whenever possible and the data were fit with a set of template models to obtain a homogeneous dataset. Here the X-ray data so collected are statistically inspected in order to infer what are the average characteristics of the nearby Seyfert galaxies contained in this sample in the 2-100 keV band. Finally, present dataset is used to perform simple tests on the unified scheme (UM) for the AGN (Antonucci 1993) and on the emission mechanism acting in the core of the Seyfert galaxies. More detailed analysis on this latter topic will be presented in another paper (Petrucci, Dadina \\& Landi, submitted) where detailed thermal Comptonization models (Poutanen \\& Svensson 1996, Haardt \\& Maraschi 1993 ) will be used to fit the BeppoSAX data with the main scope to study the dependence of the spectral properties in the ``two phase'' scenario (Haardt 1991, Haardt \\& Maraschi 1993) assuming different geometries of the corona. ",
        "conclusions": "The average properties of Seyfert galaxies in the local Universe (z$\\leq$0.1) as seen by $BeppoSAX$  has been investigated, analyzing the sample of objects presented in Dadina (2007). Multiple observations of single objects were treated independently, i.e. the multiple measurements of parameters were not averaged for statistical purposes. This method has been chosen since, in the framework of the simplest version of UM for AGN (Antonucci, 1993) the AGN are thought to be very similar to each-other and only orientation/absorption effects and the activity-level of the targets could introduce observational differences between different objects. In this scenario, the monitoring of a single source could be reproduced by the observations of many sources in different states and ``vice versa''.   $BeppoSAX$ offered the advantage of a useful X-ray broad energy band. Data studied here fall in  the 2-100 keV band for a majority of objects. This advantage has been used to investigate the properties of the high-energy continuum of Seyfert galaxies. As stated in the previous paper (Dadina, 2007), the basic template was a power-law with a high energy cut-off plus a reflection component (namely PEXRAV model in XSPEC). The results of this analysis can be summarized as follow:  \\begin{itemize}  \\item the average slope of the power-law is 1.84$\\pm$0.03 for the entire sample of objects. Considering the two families of AGN separately, it turns-out that $\\Gamma$=1.89$\\pm$0.03 for type I objects (including Seyfert 1, 1.2 and 1.5) while $\\Gamma$=1.80$\\pm$0.05 for type II objects (including Seyfert 1.8, 1.9 and 2);  \\item the average value of the relative reflected-to-direct normalization parameter R is 1.01$\\pm$0.09 with a slight difference between the two classes of Seyfert galaxies (R=1.23$\\pm$0.11 for type I objects and R=0.87$\\pm$0.14 for the type II ones);  \\item the high energy cut-off was measured to be Ec=287$\\pm$24 (Ec=230$\\pm$22 keV for Seyfert 1 and Ec=376$\\pm$42 keV for Seyferts 2);  \\item as expected and as known from previous works, the absorbing column is very different in the two classes of objects. On average N$_{H}$$\\sim$3.66$\\times$10$^{22}$ cm$^{-2}$  type I and N$_{H}$$\\sim$6.13$\\times$10$^{23}$ cm$^{-2}$ for type II AGN. The high mean value obtained for Seyfert 1 is caused by a selection effect induced by the energy coverage of the MECS+PDS instruments (2-100 keV).   \\item evidence of a X-ray Baldwin effect is found in Seyfert 1 galaxies when the EW of the  FeK$\\alpha$ line is plotted against the 20-100 keV luminosity.  \\item a significant correlation has been found between R and $\\Gamma$.   \\end{itemize}   These results are well in agreement with the basic assumptions of the UM for AGN (Antonucci 1993). In fact, no differences are measured in the observables that are supposed to be isotropic while the absorbing column seems to be the only discriminator between the different types of Seyfert galaxies. This reflects also in the properties of the FeK$\\alpha$ line. No difference is measured in the line centroid (see table 3) between the two classes of Seyfert galaxies. Type II objects, however, display more intense features (EW=222$\\pm$33 eV for Seyfert 1 and EW=693$\\pm$195 eV for Seyfert 2). The  physical origin of the X-ray ``Baldwin effect'' here measured for Seyfert 1 using the 20-100 luminosity is unclear. Both light bending (Miniutti \\& Fabian  2004) and torus models (Page et al. 2004) are consistent with present data even if the strong relation of the FeK$\\alpha$ line EW in type II objects with the absorption column indicates that the most of the narrow line component should be due to the torus. Finally, the measured $\\Gamma$-R relationship is consistent with thermal Comptonization models."
    },
    "0801/0801.4612_arXiv.txt": {
        "abstract": "The results from a $Hubble$ $Space$ $Telescope$ ($HST$) snapshot survey of post-AGB objects are shown. The aim of the survey is to complement existing $HST$ images of PPN and to connect various types of nebulosities with physical and chemical properties of their central stars. Nebulosities are detected in 15 of 33 sources. Images and photometric and geometric measurements are presented. For sources with nebulosities we see a morphological bifurcation into two groups, DUPLEX and SOLE, as previous studies have found. We find further support to the previous results suggesting that this dichotomy is caused by a difference in optical thickness of the dust shell. The remaining 18 sources are classified as stellar post-AGB objects, because our observations indicate a lack of nebulosity. We show that some stellar sources may in fact be DUPLEX or SOLE based on their infrared colors. The cause of the differences among the groups are investigated. We discuss some evidence suggesting that high progenitor-mass AGB stars tend to become DUPLEX post-AGB objects. Intermediate progenitor-mass AGB stars tend to be SOLE post-AGB objects. Most of the stellar sources probably have low mass progenitors and do not seem to develop nebulosities during the post-AGB phase and therefore do not become planetary nebulae. ",
        "introduction": "The post-AGB\\footnote{The term ``proto-planetary nebulae'' (PPN) is also often used to describe objects in transition between AGB and PN. However, one has to remember that low-mass objects that evolve very slowly will not be able to ionize the ejected matter and become planetary nebulae. Hence, the term ``post-AGB'' describes a wider group of evolved objects.} phase is a short period in the evolution of low- and intermediate-mass stars ($0.8 - 8M_{\\odot}$) between the asymptotic giant branch (AGB) and the planetary nebula (PN) phases. The theoretical evolutionary tracks of the central stars \\citep{blo95} predict that this phase lasts $10^2 - 10^5$ yrs, depending on the core mass, yet even shorter kinematical age for nebulosities, of order of $10^3 - 10^4$ yrs \\citep{sah07}. Shorter kinematical life time is obvious from the observational point of view since the entire shell (especially an extended cold shell of post-AGB object) is not necessarily detected at the optical wavelength.  Many significant changes occur during this phase to both the star and the nebula. The superwind (suddenly increased mass loss at the end of the AGB phase) stops, large amplitude pulsations cease and the circumstellar envelope slowly expands and cools revealing the central star. The post-AGB star evolves at constant luminosity on the Hertzsprung-Russell diagram towards higher temperatures. The morphology of the expanding envelope changes, resulting in the fascinating shapes of the nebula itself.  Despite many efforts the post-AGB phase is still poorly understood. Especially the departure from spherically symmetric AGB circumstellar envelopes \\citep[e.g.,][]{olo01} and the formation of diverse, axisymmetric, nebulae \\citep[e.g.,][]{bal02} are of great interest to astronomers. Determining the main cause(s) of the breaking of mass-loss symmetry is an important issue in understanding the late stages of stellar evolution.  Previous morphological studies of post-AGB objects \\citep[e.g.,][hereafter UMB00]{sah99,kwok00,hri00,hri01,ueta00} revealed many asymmetrical nebulae around central stars, where jets and concentric arcs or rings were not unusual. Further, the ground-based mid-IR (8--21 $\\mu$m) observations done by \\citet{mei99} showed two morphological classes of extended axisymmetric nebulae among post-AGB objects: ``core/elliptical'' with unresolved cores and elliptical nebulae and ``toroidal'' with limb-brightened peaks that suggest equatorial density enhancements.  Optical observations revealed more interesting details on structures seen in post-AGB nebulae. UMB00 observed 27 PPN candidates and divided them into two classes: the Star-Obvious Low-level Elongated (SOLE) nebulae which show star-dominated emission with faint extended nebulosity and DUst-Prominent Longitudinally EXtended (DUPLEX) nebulae which are dust dominated with a faint or completely obscured central star.  They also noticed that those two optical classes corresponded to the mid-IR ones, i.e. core/elliptical $\\approx$ DUPLEX and toroidal $\\approx$ SOLE, and that the main difference between SOLE and DUPLEX nebulae was the degree of the equatorial enhancement (i.e., DUPLEXes tend to have a higher equator-to-pole density ratio than SOLEs) and could not be attributed only to the inclination-angle effects. They also suggested that SOLE nebulae were optically thinner than DUPLEX nebulae and that SOLE PPNs would have evolved from low-mass progenitors and DUPLEX from high-mass progenitors. Later model calculations \\citep{mei02} confirmed the physical differences between those two classes of objects. Other high-resolution near-infrared $HST$ observations \\citep[e.g.,][]{ueta05,su03} also supported optical results. Follow-up studies by e.g., \\citet{ueta05} were conducted to further investigate the asymmetry of observed objects. They also found that differences in observed morphologies are caused (mainly) by optical depth and (to a lesser extent) by inclination of the objects. All those observations strengthen the hypothesis of an equatorial mass loss enhancement at the end of the AGB (superwind) phase \\citep{mei99} and that morphological shaping of nebulae occurs before the post-AGB phase. Moreover, the resulting diverse shapes observed in proto-planetary as well as PNs may be caused by some properties of the central stars such as chemical composition, mass and/or metallicity.  In order to find further support to those results we have executed an $HST$ snapshot survey of post-AGB objects by increasing the range of masses and chemical composition (C/O ratio) and by enlarging the sample of studied objects. The goal of the project is to connect the observed diversity of nebular shapes with physical and chemical properties of the central stars. In this paper we present the results for post-AGB objects observed recently with $HST$. Observation and data reduction are summarized in $\\S$2. The results and differences between selected groups of objects are shown in $\\S$3. In $\\S$4 we discuss obtained results with references to physical and chemical properties of stars. The conclusions are included in $\\S$5.  \\begin{table*}[!ht] \\begin{center} {Table 1: Target summary for $HST$ ACS/HRC observation of post-AGB objects} \\small \\begin{tabular}{llrrrrrl} \\tableline\\tableline \\multicolumn{1}{c}{IRAS ID} & \\multicolumn{1}{c}{Other name} & \\multicolumn{1}{c}{Prop. ID$^{1}$} & \\multicolumn{1}{c}{$V$ mag$^{2}$} & \\multicolumn{1}{c}{J mag$^{2}$} & \\multicolumn{1}{c}{T$_{\\rm eff}$$^{3,4}$} & \\multicolumn{1}{c}{Mass$^4$} & \\multicolumn{1}{c}{C/O$^4$} \\\\ \\tableline 01005+7910       & ...         & 10627~~~~ & 10.96~~ & 10.274~ &21000~~    &     0.55~~~~ & ~~C     \\\\ 04395+3601       & AFGL 618    &  9430~~~~ & ...~~~~ & 13.510~ &25000\\tablenotemark{a}~&    ...~~~~~~ &~~C\\tablenotemark{a} \\\\ 06034+1354       & DY Ori      & 10627~~~~ & ...~~~~ &  8.073~ & 5900~~    & $<$ 0.55~~~~ & ~~O~    \\\\ 08143$-$4406     & PM 1$-$39   & 10627~~~~ & ...~~~~ &  9.172~ & 7150~~    & $<$ 0.55~~~~ & ~~C~    \\\\ 09256$-$6324     & IW Car      & 10627~~~~ &  8.33~~ &  5.875~ & 6700~~    & $<$ 0.55~~~~ & ~~C~    \\\\ 11385$-$5517$^5$ & V885 Cen    &  9436~~~~ &  7.04~~ &  5.947~ & 8500~~    &     0.55~~~~ & ~~O~    \\\\ 12067$-$4508     & RU Cen      & 10627~~~~ &  9.13~~ &  7.616~ & 6000~~    & $<$ 0.55~~~~ & ~~O~    \\\\ 12175$-$5338     & V1024 Cen   & 10627~~~~ &  9.40~~ &  8.326~ & 7350~~    &     0.62~~~~ & ~~O~    \\\\ 12222$-$4652     & HD 108015   & 10627~~~~ &  8.01~~ &  6.941~ & 6800~~    & $<$ 0.55~~~~ & ~~O~    \\\\ 12538$-$2611     & LN Hya      & 10627~~~~ &  6.82~~ &  5.251~ & 6000~~    & $<$ 0.55~~~~ & ~~O~    \\\\ 13416$-$6243     & ...         & 10627~~~~ & ...~~~~ & 10.302~ &...~~~~    &    ...~~~~~~ & ~~C\\tablenotemark{b} \\\\ 13428$-$6232$^5$ & ...         &  9463~~~~ & ...~~~~ & 13.176~ &...~~~~    &    ...~~~~~~ & ~~C\\tablenotemark{b} \\\\ 15039$-$4806     & HD 133656   & 10627~~~~ &  7.58~~ &  6.680~ & 8000~~    &     0.55~~~~ & ~~O~    \\\\ 15469$-$5311     & ...         & 10627~~~~ & 10.82~~ &  7.190~ & 7500~~    & $<$ 0.55~~~~ & ~~O~    \\\\ 15553$-$5230$^5$ & ...         & 10627~~~~ & ...~~~~ & 13.380~ &...~~~~    &    ...~~~~~~ & ~~...   \\\\ 16206$-$5956     & LS 3591     & 10627~~~~ & 10.00~~ &  9.002~ & 8500~~    &     0.66~~~~ & ~~O~    \\\\ 17163$-$3907     & Hen 3$-$1379& 10185~~~~ & 12.45~~ &  4.635~ &...~~~~    &    ...~~~~~~ & ~~...   \\\\ 17243$-$4348     & LR Sco      & 10627~~~~ & 10.49~~ &  8.035~ & 6750~~    &     0.94~~~~ & ~~O~    \\\\ 17279$-$1119     & V340 Ser    & 10627~~~~ &  9.78~~ &  7.845~ & 7300~~    & $<$ 0.55~~~~ & ~~C~    \\\\ 17516$-$2525     & ...         &  9436~~~~ & 17.76~~ &  8.695~ &...~~~~    &    ...~~~~~~ & ~~O\\tablenotemark{b} \\\\ 17534+2603       & 89 Her      & 10627~~~~ &  5.51~~ &  4.998~ & 6550~~    &     0.61~~~~ & ~~O~    \\\\ 18135$-$1456     & OH 15.7+0.8 &  9436~~~~ & 16.61~~ & ...~~~~ &...~~~~    &    ...~~~~~~ & ~~O\\tablenotemark{c} \\\\ 19125+0343       & BD+03 3950  & 10627~~~~ & 10.46~~ &  7.903~ & 7750~~    &     0.58~~~~ & ~~O~    \\\\ 19157$-$0247     & BD$-$02 4931& 10627~~~~ & 10.87~~ &  8.872~ & 7750~~    &     0.58~~~~ & ~~O~    \\\\ 19306+1407$^5$   & ...         &  9436~~~~ & ...~~~~ & 11.286~ & B~~~~     &    ...~~~~~~ & ~~O\\tablenotemark{b} \\\\ 19475+3119$^5$   & HD 331319   &  9436~~~~ &  9.60~~ &  7.773~ & 7750~~    &     0.58~~~~ & ~~O~    \\\\ 20000+3239$^5$   & ...         &  9436~~~~ & ...~~~~ &  8.021~ & 5500\\tablenotemark{d}~&    ...~~~~~~ & ~~C\\tablenotemark{b} \\\\ 20117+1634       & R Sge       & 10627~~~~ &  9.31~~ &  7.818~ & 5000~~    &     0.93~~~~ & ~~O~    \\\\ 20547+0247       & U Equ       &  9436~~~~ & ...~~~~ & 11.561~ & G~~~~     &    ...~~~~~~ & ~~O\\tablenotemark{e} \\\\ 22036+5306$^5$   & ...         & 10185~~~~ & ...~~~~ & 11.666~ &...~~~~    &    ...~~~~~~ & ~~O\\tablenotemark{b} \\\\ 22223+4327$^5$   & BD+42 4388  &  9436~~~~ & 10.00~~ &  7.812~ & 6500~~    &     0.55~~~~ & ~~C~    \\\\ 23304+6147$^5$   & PM 2$-$47   &  9436~~~~ & ...~~~~ &  8.501~ & 6750~~    &     0.66~~~~ & ~~C~    \\\\ 23541+7031       & M 2$-$56    &  9436~~~~ & ...~~~~ & 13.857~ & B~~~~     &    ...~~~~~~ & ~~O\\tablenotemark{b} \\\\ \\tableline \\end{tabular} \\end{center} \\small $^1$ 9430 - PI S. Trammell, 9436 \\& 10185 - PI R.Sahai, 10627 - PI M. Meixner.\\\\ $^2$ Magnitudes from GSC2.2 and 2MASS catalogues.\\\\ $^3$ Spectral type is shown if T$_{\\rm eff}$ is not known from the model atmosphere analysis.\\\\ $^4$ Data collected by \\citet{sta06} with few exceptions: $^a$ \\citet{cer01} and references therein, $^b$ C/O ratio established by the inspection of the ISO spectrum, $^c$ \\citet{eng02} and references therein, $^d$ \\citet{volk02}, $^e$ \\citet{geb05} and references therein.\\\\ $^5$ Images of those objects were also published in the newest paper by \\citet{sah07}. \\vspace{0.5cm} \\end{table*} ",
        "conclusions": "We analyzed HST images of 31 post-AGB objects.  We increased the number of observed PPNs and covered a wider variety of PPNs. Following previous results of UMB00 we selected SOLE and DUPLEX sources based on their morphology. Our results are consistent with UMB00, and therefore, strengthen the validity of the general PPN structure suggested by UMB00, in which an intrinsically axisymmetric shell assumes a different morphology mainly due to the varying degree of the equatorial density enhancement which determines the presence of the dust lane in DUPLEXes and the absence of it in SOLEs. We selected a separate group of stellar post-AGB objects without visible nebulosities. We searched for connections between morphology of the nebulosities and the chemical and physical properties of their central stars and did not find any obvious connection. We confirmed results of UMB00 that SOLE objects with thin envelopes are bluer then dusty DUPLEX sources. We also suggested that some of the stellar sources may be objects with the nebulosity around central star, but seen pole-on and/or too faint and distant to be resolved. We also found the differences in galactic distribution of different groups with the SOLEs lying farther from the galactic plane than DUPLEXes, which could suggest that DUPLEX objects may have more massive progenitors. However, we were not able to confirm this suggestion directly because the masses of DUPLEXes and some SOLEs are not known."
    },
    "0801/0801.1215_arXiv.txt": {
        "abstract": "We study clusters in the BCS cluster sample which are observed by {\\it Chandra} and are  more distant than redshift, $z>0.1$.  We select from this subsample the clusters which have both a short central cooling time and a central temperature drop, and also those with a central radio source. Six of the clusters have clear bubbles near the centre.  We calculate the heating by these bubbles and express it as the ratio $r_{\\rm heat}/r_{\\rm cool}=1.34\\pm0.20$.  This result is used to calculate the average size of bubbles expected in all clusters with central radio sources.  In three cases the predicted bubble sizes approximately match the observed radio lobe dimensions.  We combine this cluster sample with the B55 sample studied in earlier work to increase the total sample size and redshift range.  This extended sample contains 71 clusters in the redshift range $0\\leq z \\leq0.4$.  The average distance out to which the bubbles offset the X-ray cooling in the combined sample is at least $r_{\\rm heat}/r_{\\rm cool}=0.92\\pm0.11$.  The distribution of central cooling times for the combined sample shows no clusters with clear bubbles and $t_{\\rm cool}>1.2 \\gyr$.  An investigation of the evolution of cluster parameters within the redshift range of the combined samples does not show any clear variation with redshift. ",
        "introduction": "\\label{sec:intro}  Since the discovery of ``holes'' in the X-ray emission at the centre of the Perseus Cluster \\citep{Bohringer93}, and especially since the launch of \\emph{Chandra} and \\emph{XMM-NEWTON}, many depressions in the central intra-cluster medium (ICM) of low redshift clusters have been found (e.g. Hydra A, \\citealp{McNamara00}; A2052, \\citealp{Blanton01}; A2199, \\citealp{Johnstone02}; Centaurus, \\citealp{Sanders02}).  Recent compilations are given by \\citet{Dunn06f,Rafferty06} and \\citet{Birzan04}.  These holes have been observed to anti-correlate spectacularly with the radio emission from the active galactic nucleus (AGN) at the centres of these clusters.  Their morphology, particularly in the closest clusters, has led to their interpretation as bubbles of relativistic gas blown by the AGN into the thermal ICM. The relativistic gas is less dense than the ICM, so the bubbles detach from the core and rise up buoyantly through the cluster, e.g. Perseus \\citep{Churazov00,Fabian03b}.  The older bubbles tend not to have $\\ghz$ radio emission associated with them and have been termed ``Ghost'' bubbles.  The X-ray emission of the ICM represents an energy loss, so leads the plasma cooling in the core if there were no compensating energy source.  To maintain  pressure support, the gas would flow on to the central galaxy as a ``cooling flow.''  However, with the high spatial and spectral resolution of \\emph{Chandra} and \\emph{XMM-Newton}, the gas temperature is seen to drop by a factor of several in the core but no continuous cooling flow is found (\\citealp{Peterson03}, see \\citealp{Peterson06} for a review).  Some form of gentle, continuous and distributed heating appears to be at work \\citep{Voigt04}.  The interaction of the AGN with the ICM is the likeliest explanation as the heat source for cool-core clusters.  The bubble model for radio sources was proposed by \\citet{Gull73}, and further theoretical results were presented in \\citet{Churazov00, Churazov02}.  Recent studies on the heating ability of AGN within clusters are \\citet{Birzan04, Rafferty06} and \\citet{Dunn06f}.  A majority (71 per cent) of ``cooling core''  clusters harbour radio sources \\citep{Burns90}, and a similar number of clusters which require heating (likely to be a cooling core) harbour clear bubbles \\citep{Dunn05, Dunn06f}.  The action of creating the bubbles at the centre of the cluster by the AGN is a favoured method of injecting energy into the central regions of the cluster and so prevent the ICM from cooling.  This process sets up sound/pressure waves in the ICM, the dissipation of which requires the ICM is to be viscous \\citep{Fabian03a}.  The viscous dissipation of the pressure waves allows the energy from the bubble creation to be dissipated far from the cluster centre where it is required.  In \\citet{Dunn06f} we analysed clusters in the Brightest 55 (B55) sample \\citep{Edge90}.  These are all at comparatively low redshift (only two are at $z>0.1$).  Using the results of \\citet{Peres98} and the literature at the time we selected those clusters with both a short central cooling time and a large central temperature drop, as well as those with a central radio source.  We found that at least 70 per cent of clusters in which high cooling rates are expected (those with short central cooling times and central temperature drops) harbour AGN bubbles.  We also found that 95 per cent of these cool core clusters harbour radio cores.  The total fraction of clusters which harbour a radio source is 53 per cent (29/55).  Where AGN bubbles are present, on average they show that the AGN is injecting enough energy into the central regions of the cluster to offset the X-ray cooling out to the cooling radius ($r_{\\rm cool}$ where $t_{\\rm cool}=3\\gyr$).  To increase the number of clusters studied and also investigate any evolution in the AGN heating with redshift we now turn to the Brightest Cluster Sample \\citep{Ebeling98}, initially performing a similar analysis and then combining the results from both cluster samples.  The sample selection for this work is outlined in Section \\ref{sec:selection} and the data preparation and reduction in Section \\ref{sec:data_prep}.  We then look at the two different cluster subsets, those with bubbles in Section \\ref{sec:bubbles} and those with central radio sources in Section \\ref{sec:radio}.  The results of our investigations into the global differences within the different subsets is presented in Section \\ref{sec:heating}.  The combined cluster sample and the redshift evolution results are presented in Section \\ref{sec:discuss}.  We use $H_0=70\\kmpspmpc$ with $\\Omega_{\\rm m}=0.3$, $\\Omega_\\Lambda=0.7$ throughout this work. ",
        "conclusions": "We have extended our study of the effect of central AGN in clusters to clusters with $z>0.1$, using the BCS.  There are only six clusters which harbour clear bubbles, and on average they overcompensate for the X-ray cooling.  These clusters are combined with an updated analysis of those in the B55 cluster sample.  The average distance, as a fraction of the cooling radius,  out to which the bubbles can offset the X-ray cooling from the combined sample is $r_{\\rm heat}/r_{\\rm cool}=0.92\\pm0.11$ for clusters where this distance can be reliably determined.  Adding in all clusters with bubbles results in $r_{\\rm heat}/r_{\\rm cool}=2.10\\pm0.40$.  Although it appears as if the AGN bubbles can easily counteract the X-ray cooling within the cooling radius, there is a selection effect, especially at the higher redshifts, which means we are missing small bubbles in the most distant clusters.  Using the result from the BCS bubbles, we calculate the expected size of bubbles within the clusters which harbour radio sources.  In three cases (A1068, A1763 and A2204) the predicted bubble sizes and observed radio emission are similar.  In some of the clusters the radio source identified in the NVSS is offset to the cluster centre in VLA observations, and so these may be background AGN rather than associated with the cluster.  A category of clusters with cool cores, radio sources and no clear bubbles exists in both the B55 and BCS samples. The lack of bubbles is explained by low signal-to-noise ratio in the X-ray images of some but not all of these objects.  The distribution of the central cooling time of the clusters shows that there are no clusters with bubbles which have $t_{\\rm cool}>1.2\\gyr$.  Also there are only 10/24 clusters with  $t_{\\rm cool}>4.0\\gyr$ which harbour a central radio source, only 5/47 clusters with  $t_{\\rm  cool}<4.0\\gyr$ have no central radio source.  We investigated the evolution of cluster parameters with redshift and over the range $0.0<z<0.4$ we do not find any significant change.  The cooling time profiles show some variation but this can be explained as being the result of the resolution differences as the distances increase."
    },
    "0801/0801.4606_arXiv.txt": {
        "abstract": "We investigate the modified gravity theories in terms of the effective dark energy models. We compare the cosmic expansion history and the linear growth in different models. We also study the evolution of linear cosmological perturbations in modified theories of gravity assuming the Palatini formalism. We find the stability of the superhorizon metric evolution depends on models. We also study the matter density fluctuation in the general gauge and show the differential equations in super and sub-horizon scales. ",
        "introduction": "Type Ia supernova distance-redshift measurement shows that the present expansion of the universe is accelerating \\cite{SNe}. To explain this phenomena we need to introduce a homogeneous component of energy with a negative pressure, dubbed dark energy \\cite{DE}, or a modification of gravity \\cite{BD,DGP,Buc,Ferraris}.  We can parameterize the effective dark energy component in the modified gravity theories \\cite{Linder}. The evolution equation of Hubble parameter and that of energy density fluctuation can be simply expressed with this parametrization.  The Palatini formalism where the metric and the connections are treated as independent variables and the energy momentum tensor only depends on the metric leads to a different theory from what is obtained from the metric formalism. The Palatini formalism of f(R) gravity results in second order differential equations due to the algebraic relation between the curvature scalar and the trace of the energy momentum tensor. We will use the evolution of linear perturbations in f(R) models in the physical frame to check the stability of the theory. We also investigate the evolution of the matter density fluctuation.  In the next section we introduce the parametrization of modified gravities as an effective dark energy. We review the Palatini f(R) gravity in section 3. In section 4, we investigate the linear perturbation of f(R) gravity. We also derive the stability equation of metric fluctuations in the high curvature limit and show the stability in a specific model. We show the evolutions of the metric fluctuation and the density fluctuation in the superhorizon and the subhorizon scales in section 5. We reach our conclusions in section 6. ",
        "conclusions": "We investigate the different modified gravity models as an effective dark energy models. We can parameterize the effective equation of state of dark energy in terms of modified gravities. We show the evolution of energy density fluctuation compared to that of general relativity.  We have analyzed the stability of metric fluctuations by checking the cosmological evolution of linear perturbations in Palatini f(R) gravity. We have also considered the matter density fluctuation in the Newtonian gauge.  We have shown that the stability of metric fluctuations in the high redshift limit of high curvature is not simply expressed. We need to check each model for the stability. However, we have found that the deviation from the superhorizon metric evolution is null for a specific choice of the nonlinear Einstein-Hilbert action, $f(\\hat{R}) \\sim \\hat{R}^{n}$ and stability of this model is guaranteed.  We have investigated the evolution equations of Newtonian potential and matter density contrast in super and sub-horizon scales. In the specific model, superhorizon evolutions of Newtonian potential and matter density fluctuation are same to those of general relativity. However, subhorizon evolutions show the different behaviors from the general relativity case. This will give us the method to probe the possibility of $f(R)$ theory."
    },
    "0801/0801.1509_arXiv.txt": {
        "abstract": "\\sloppy We report the discovery of WASP-4b, a large transiting gas-giant planet with an orbital period of 1.34 days.  This is the first planet to be discovered by the SuperWASP-South observatory and CORALIE collaboration and the first planet orbiting a star brighter than 16$^{th}$ magnitude to be discovered in the Southern hemisphere.  A simultaneous fit to high-quality lightcurves and precision radial-velocity measurements leads to a planetary mass of 1.22$^{+0.09}_{-0.08}$ M$\\rm_{Jup}$ and a planetary radius of 1.42$^{+0.07}_{-0.04}$ R$\\rm_{Jup}$. The host star is USNO-B1.0 0479-0948995, a G7V star of visual magnitude 12.5. As a result of the short orbital period, the predicted surface temperature of the planet is 1761 K, making it an ideal candidate for detections of the secondary eclipse at infrared wavelengths. ",
        "introduction": "Transiting extrasolar planets offer our best opportunity for measuring fundamental parameters of extrasolar planets, such as radius and mass, allowing us to test our models for planetary formation and evolution. To-date the only transiting planets known in the Southern sky are those discovered by surveys targeting the galactic plane (e.g. \\citealp{ogle, ogle2,lupus}) and find planets around stars with typical visual magnitudes of 16--17.  The SuperWASP-South (WASP-S) instrument looks for transiting planets around stars of visual magnitude 8--13, and will eventually cover the entire southern sky, excluding the Galactic plane.  We report here the discovery of WASP-4b, a transiting hot-Jupiter orbiting a star of magnitude 12.5.  This is the first discovery by WASP-S, made in collaboration with radial-velocity measurements from the CORALIE spectrograph on the 1.2-m Euler telescope. The discovery marks the beginning of a campaign to discover brighter transiting extrasolar planets in the Southern hemisphere, to complement those discovered in the North by projects including HAT \\citep{hat}, TrES \\citep{tres}, XO \\citep{xo} and WASP \\citep{sw}.  Transit searches are most sensitive to large planets with short orbital periods, and have now found nine Jupiter-mass planets in orbits of less than 2 d\\footnote{http://exoplanet.eu}. This has led to suggestions of a class of very-hot Jupiters \\citep{melo}, with highly irradiated atmospheres (e.g. \\citealt{fortney}), to which WASP-4b likely belongs. ",
        "conclusions": "A simultaneous fit of precision photometry and radial velocity measurements result in a planetary radius of 1.42 R$_{\\rm Jup}$ and mass of 1.22 M$_{\\rm Jup}$ for WASP-4b.  This is the second largest transiting planet discovered to date, second only to TrES-4b (\\citealp{tres4}; for plots comparing WASP-4b with other transiting extrasolar planets see \\citealt{wasp3}).  WASP-4b, with an orbital period of 1.3 d, can be compared to other very-short-period planets such as OGLE-TR-56b \\citep{ogle56}, TrES-3b \\citep{tres3} and WASP-3b \\citep{wasp3}.  The low rotation velocity of the host star WASP-4 is comparable with that of TrES-3; in contrast, OGLE-TR-56 and WASP-3 have much higher rotation velocities, which is thought to be due to the young age of these stars.  The stars  WASP-4 and TrES-3 are thought to be older, however the synchronisation time for WASP-4 would be longer still at 8 Gyr (using the  method of \\citealp{marcy}), so their low rotational velocity is consistent with their not having been spun-up by their planets.  A preliminary estimate of the stellar parameters indicates that the parent star is spectral type G7V with solar metallicity. Together with the short orbital period of 1.34 d this results in a blackbody planetary surface temperature, assuming zero albedo and isotropic re-radiation, of 1760 K.  The orbital distance of 0.023 AU places WASP-4b well within the criterion ($a < 0.04$ AU) for the proposed new class of pM planets \\citep{fortney}.  These planets display a temperature inversion due to low-pressure stellar absorption by gaseous TiO and VO.  As a result they exhibit an unusually hot stratosphere which emits strongly in the mid-infrared.  The WASP-4b system has a larger flux ratio than the very-hot-Jupiters WASP-3b and TrES-3b and so is an ideal target for secondary eclipse studies by {\\sl Spitzer}.  The study of more of these systems, including WASP-4b, should help constrain models for atmospheric dynamics and heat distribution."
    },
    "0801/0801.1023_arXiv.txt": {
        "abstract": "Modern cosmological observations allow us to study in great detail the evolution and history of the large scale structure hierarchy. The fundamental problem of accurate constraints on the cosmological parameters, within a given cosmological model, requires precise modelling of the observed structure. In this paper we briefly review the current most effective techniques of large scale structure simulations, emphasising both their advantages and shortcomings. Starting with basics of the direct N-body simulations appropriate to modelling cold dark matter evolution, we then discuss the direct-sum technique {\\sl GRAPE}, particle-mesh ({\\sl PM}) and hybrid methods, combining the {\\sl PM} and the tree algorithms. Simulations of baryonic matter in the Universe often use hydrodynamic codes based on both particle methods that discretise mass, and grid-based methods. We briefly describe Eulerian grid methods, and also some variants of Lagrangian smoothed particle hydrodynamics ({\\sl SPH}) methods. ",
        "introduction": "In the hierarchical picture of structure formation, small objects collapse first and then merge to form larger and larger structures in a complex manner. This formation process reflects on the intricate structure of galaxy clusters, whose properties depend on how the thousands of smaller objects that the cluster accretes are destroyed or survive within the cluster gravitational potential. These merging events are the source of shocks, turbulence and acceleration of relativistic particles in the intracluster medium, which, in turn, lead to a redistribution or amplification of magnetic fields, and to the acceleration of cosmic rays. In order to model these processes realistically, we need to resort to numerical simulations which are capable of resolving and following correctly the highly non-linear dynamics. In this paper, we briefly describe the methods which are commonly used to simulate galaxy clusters within a cosmological context.   Usually, choosing the simulation setup is a compromise between the size of the region that one has to simulate to fairly represent the object(s) of interest, and the resolution needed to resolve the objects at the required level of detail. Typical sizes of the simulated volume are a megaparsec scale for an individual galaxy, tens to hundreds of megaparsecs for a galaxy population, and several hundreds of megaparsecs for a galaxy cluster population. The mass resolution varies from $\\approx 10^5$~M$_\\odot$ up to $\\approx 10^{10}$~M$_\\odot$, depending on the object studied, while, nowadays, one can typically reach the resolution of a few hundred parsec for individual galaxies and above the kiloparsec scale for cosmological boxes. ",
        "conclusions": ""
    },
    "0801/0801.1486_arXiv.txt": {
        "abstract": "In this work we show that the presence of a  vector field on cosmological scales could explain the present phase of accelerated expansion of the universe. The proposed  theory contains no dimensional parameters nor potential terms and does not require unnatural initial conditions in the early universe, thus avoiding the so called cosmic coincidence problem. In addition, it  fits the data from high-redshift supernovae with excellent precision, making definite predictions for cosmological parameters. Upcoming observations will be able to clearly discriminate this model from standard cosmology with cosmological constant. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.2485_arXiv.txt": {
        "abstract": "Previous studies of elemental abundances in Mercury-Manganese (HgMn) stars have occasionally reported the presence of lines of the ionized rare noble gas \\Xeii,\\ especially in a few of the hottest stars with \\mbox{$\\teff\\sim13000$--15000 K}. A new study of this element has been undertaken using observations from Lick Observatory's Hamilton \\'{E}chelle Spectrograph. In this work, the spectrum synthesis program UCLSYN has been used to undertake abundance analysis assuming LTE. We find that in the Smith \\& Dworetsky sample of HgMn stars, Xe is vastly over-abundant in 21 of 22 HgMn stars studied, by factors of 3.1--4.8 dex. There does not appear to be a significant correlation of Xe abundance with \\teff. A comparison sample of normal late B stars shows no sign of \\Xeii\\ lines that could be detected, consistent with the expected weakness of lines at normal abundance. The main reason for the previous lack of widespread detection in HgMn stars is probably due to the strongest lines being at longer wavelengths than the photographic blue. The lines used in this work were $\\lambda$4603.03, $\\lambda$4844.33 and $\\lambda$5292.22. ",
        "introduction": "Among the chemically peculiar (CP) stars of the upper main sequence, the HgMn stars \\mbox{(10500 K $\\leq \\teff \\leq 14500$ K)} constitute some of the objects most intensively studied with a wide range of optical and ultraviolet spectra (summarised by \\citealt{wal04}). The reasons for this include that: many of these stars are extremely slow rotators, making details of their line profiles easily studied for isotopic anomalies, for example; they are not affected by strong magnetic fields, which simplifies theoretical considerations; and there are many of these stars bright enough to study with the highest spectral resolutions available. Most often these stars have been studied individually, but a few studies have considered single elements among a set of stars to determine the overall behaviour of abundances and isotope distributions as a function of physical variables such as \\teff\\ (for example, \\citealt{s93,s94,s96,s97,sd93,wl99,db2000}). Here, we report an analysis of strong unblended lines of the rare noble gas xenon (atomic number 54; cosmic \\mbox{log {\\em A}\\,(Xe) = 2.27} on the scale where \\mbox{log {\\em A}\\,(H) = 12.00}) with a view to determining its abundance in a sample of HgMn stars. As will be shown, Xe is found in nearly all HgMn stars and exhibits abundance excesses of 3 dex or more.  Previous observations of \\Xeii\\ in the spectra of HgMn stars are scanty, although it was originally reported as a possible identification in 3 Cen A by \\citet{bid62} and a probable identification was reported by \\citet*{andersen_etal84} in the He-weak Bp star HR\\,6000 which seems to be closely related to the HgMn stars. More recently, \\Xeii\\ was identified in HR 7361 by \\citet{adel92}; 46 Aql by \\citet*{sada-etal01}; $\\kappa$ Cnc by \\citet{rs88}; 33 Gem by \\citet*{adel-etal96}; 112 Her by \\citet*{ryab-etal96}. This sparseness may be attributable to the past tendency to observe these stars mainly in the blue photographic region, where \\Xeii\\ lines are relatively weak. For example, \\citet{guth85} explored Cowley's Dominion Astrophysical Observatory 2.4\\,\\AA\\,mm$^{-1}$ spectrograms of several HgMn stars but did not report \\Xeii.  In this work, the strong lines of \\Xeii\\ used were $\\lambda$5292.22 and $\\lambda$4844.33, supplemented by the weaker $\\lambda$4603.03 visible in most of the programme stars. ",
        "conclusions": "\\subsection{The xenon enhancement in HgMn stars}  Xenon (\\Xeii) is yet another of the spectra originally omitted from the Revised Multiplet Tables by \\citet{moore45} as `not of astrophysical interest' yet subsequently found in the HgMn stars and other CP types (e.g., \\Gaii,\\ \\Krii,\\ \\Ptii,\\ \\Auii,\\ \\Hgii,\\ \\Tlii\\ and \\Biii\\ have also been identified in stars).  The xenon enhancement in HgMn stars rivals that of other rare elements such as Hg \\citep{s97,wl99} with many of the stars having Xe enhancement factors of \\mbox{40--50\\,000}. Such a large enhancement of abundance implies a strong upwards radiative acceleration on Xe over a wide range of \\teff\\ and \\logg,\\ and it will be interesting to see if this prediction is confirmed by future calculations, and to see whether Xe is expected to be stratified in a thin layer. Radiative diffusion processes must be dredging up xenon from deep inside the stellar envelope to get an enhancement of this size. As can be seen in Figure~\\ref{fig-abunds}, the enhancements found in the cooler HgMn stars clustered around \\mbox{11\\,000 K} are all very similar at about 3.9 dex. At higher \\teff,\\ the enhancements show much more scatter, which is not an artifact of the measurement uncertainties but must be a real effect. The one star with an upper limit only, $\\phi$\\,Her, could itself have a considerable enhancement of Xe but yet be below the limit of detectability in the present work.  We looked at what we feel are the 10 best-observed stars in Table~\\ref{abunds} to see if there were any systematic differences in results from the lines used. The mean difference between \\mbox{log {\\em{A}}} derived from $\\lambda$4844 and $\\lambda$5292 is \\mbox{0.03 $\\pm$ 0.06 dex}, so these two lines give results that agree very well. The mean difference between abundances from $\\lambda$4603 and the average of the other two lines is \\mbox{-0.22 $\\pm$ 0.08 dex}, which seems obviously significant. The most likely explanation is an error in one or more of the adopted oscillator strengths. However, this does not affect any of the essential conclusions in this work.  The rarity of \\Xeii\\ identifications in previous work is explicable by the relative weakness of its lines in the photographic blue region, where many observations were done only below 4600\\AA. It is easy to understand why the lines might be overlooked in some earlier work of this sort -- and in at least one case, a strong line of another element masks the presence of a \\Xeii\\ line, as discussed below.   \\subsection{A comment on strontium abundances}  In the course of an investigation of strontium abundances in HgMn stars \\citep*{ddp08} it was noticed that \\Srii\\ $\\lambda$4215.519 apparently persisted even when \\Srii\\ $\\lambda$4077.714 was completely undetectable. An example of this is seen in the \\Srii\\ results of \\citet{ryab-etal96}, where $\\lambda$4215.519 has a measured equivalent width of 6.5m\\AA\\ in 112\\,Her. We also found a similar strength for this line, but a wavelength slightly shifted.  We were unable to detect \\Srii\\ $\\lambda$4077.714 at all ($\\leq$0.6m\\AA), nor did \\citet{ryab-etal96} report seeing this line.  The solution to this anomaly is our identification of a blend of \\Srii\\ with \\Xeii\\ $\\lambda$4215.60 \\citep{hp87}, which has an estimated strength of 4m\\AA\\ in 112\\,Her from the abundance in Table~\\ref{abunds} and an astrophysical \\loggf\\ of $-1.06$, based on the Lick observations of the ultra-sharp-lined, Sr-deficient, ``mild'' HgMn star 46\\,Aql \\citep{cow80,sd93}. The remaining strength of this feature in 112\\,Her is probably due to weak lines of \\hbox{Mn\\,{\\sc ii}} and \\hbox{Cr\\,{\\sc ii}}. Strontium is at least 1.0 dex below solar abundance in both 46\\,Aql and 112\\,Her."
    },
    "0801/0801.2166_arXiv.txt": {
        "abstract": "We consider a Hartmann layer, stationary flow of a viscose and resistive fluid between two plates with superimposed transverse magnetic field,  in the limit of  gyrotropic  plasma, when viscosity across the field is strongly suppressed. For zero cross-field viscosity, the problem is not well posed, since viscosity then vanishes on the boundaries and in the middle of the layer, where there is no longitudinal field. An additional arbitrarily small isotropic viscosity allows one to find magnetic field and velocity profiles which are independent of this viscosity floor and  different from flows with isotropic viscosity. Velocity sharply rises in a thin boundary layer, which thickness depends both on the Hartmann number and  on the Lundquist number of the flow.  The implication of the work is that,  in simulating ICM dynamics it is imperative to use  numerical schemes which take into account anisotropic viscosity. Although magnetic fields  are dynamically subdominant in the ICM  they do  determine its the dissipative properties,  stability of embedded structures and the transition to turbulence. ",
        "introduction": "Gasdynamical interactions  of  magnetized flows  in the cores of clusters of galaxies play an important role in formation of the observed morphological structures, \\eg\\ through expansion of AGN blown bubble  into ICM medium and resulting plasma heating, and  interaction  of two gas components in merging clusters \\cite[\\eg][]{MarkevitchVikhlinin07}. Majority of theoretical work on these topics has been numerical, mostly using the existing fluid and MHD codes, like ZEUS. Unfortunately, simple hydrodynamic models \"face multiple failures\" \\citep{Reynolds05} for cluster cores. Perhaps, the most evident  example is  the  Raleigh-Taylor and Kelvin-Helmholtz instabilities of AGN blown bubbles, which disrupt the bubbles approximately on one rise time \\citep[\\eg][]{Kaiser05}. On the other hand, artificial fiddling with viscosity (which is usually parametrized with respect to \\cite{Braginsky} value; this procedure is not justified in ICM  \\citep{Schekochihin05,lyutikov07}), shows that \"modest\" changes of shear viscosity lead to qualitatively different results \\citep[\\eg][ increased viscosity make the ICM plasma gel-like and quenches the instability]{Reynolds05,Sijacki}.  One of principal reasons, perhaps,  for the  failure  of simple MHD codes is that they use {\\it isotropic} viscosity, while   ICM  plasma is strongly gyrotropic, in a sense that it is  weakly collisional, $r_L \\ll \\lambda \\leq L$ ($r_L$ is ion Larmor radius, $\\lambda$ is mean free path and $L$ is a typical size of the system),   and weakly magnetized, kinetic pressure $p$ is much larger than magnetic pressure, $\\beta = 8 \\pi p/B^2\\gg 1$ ($B$ is a typical \\Bf). In ICM the  ion  Larmor radius $r_L \\sim 10^{8-9}$ cm is many orders of magnitude smaller than the  mean free path $\\lambda \\sim 10^{22-23} $ cm and the size $L\\sim 10^{24}$ cm, while $\\beta \\sim 100$ \\citep{ct02}.   In a strongly gyrotropic plasma the local  transport properties, primarily viscosity and conductivity,  become anisotropic \\citep{Braginsky}. The effects of anisotropic viscosity and conductivity are expected to change substantially  the results of ICM simulations. As a simple example, note that \\Bf\\ drapes around the contact surface separating the two interacting media  \\citep{lyutikov07}. In a strongly gyrotropic plasma the shear viscosity inside the draping layer, with a  flow along \\Bf\\ lines, then becomes 0. This runs contrary to the idea that high viscosity may provide stabilization  \\cite{Reynolds05}. (Note that draping itself can provide stabilization against KH instability, in a way similar to the effect of a thin oil layer on water waves \\citep{Dursi07}).  Understanding the basic properties of  strongly gyrotropic plasma is imperative for further progress, especially for parametrization of the 'sub-grid' physics in large numerical simulations.  In this Letter we adapt one of the basic solutions of MHD, a Hartmann flow  \\citep[\\eg][]{LLVIII}, to the case of anisotropic viscosity. ",
        "conclusions": "In this letter we considered a basic plasma physics problem: Hartmann flow with anisotropic viscosity. We first argued that in the case when transverse viscosity  is suppressed completely, the problem cannot be formulated in  a physically meaningful way: there should be some small isotropic contribution to viscosity. We derived  \\Bf and velocity profiles, which in the limit of small  isotropic viscosity floor  are independent of its exact value. These profiles are considerably different from the case of isotropic viscosity. The velocity gradients are much more concentrated close to the walls of the channel than in the case of isotropic viscosity.  How important is the structure of a boundary layer for the overall structure of the flow? On the one hand, in a laminar regime at low Reynolds numbers, the structure of the boundary layer is probably not important: in the boundary layer the relative velocity just drops to zero, according to some law,  without affecting the overall structure of the flow. On the other hand, properties of the  boundary layer  determine its stability  and, thus, transition to turbulence \\citep{LLVI}. In ICM plasma the Reynolds numbers are in the range $Re \\sim 10-1000$, while, typically, transition to turbulence occurs at $Re \\sim 100-1000$.  (Transition to turbulence occurs in the boundary layer). Thus, we expect that effects of anisotropic viscosity  are likely to be very important for ICM plasma, especially for determining its transition to turbulence.  Thus, in simulating ICM dynamics it is imperative to use appropriate numerical schemes which take into account anisotropic viscosity (and to a lesser extent, conductivity), like the ones that have been applied to accretion disks \\citep{sharma07}. Though \\Bfs are dynamically subdominant in ICM (plasma beta  parameter is large), they do, in fact,  determine its the dissipative properties,  stability of embedded structures and transition to turbulence."
    },
    "0801/0801.2399_arXiv.txt": {
        "abstract": "Even if nothing but a light Higgs is observed at the LHC, suggesting that the Standard Model is unmodified up to scales far above the weak scale, Higgs physics can yield surprises of fundamental significance for cosmology. As has long been known, the Standard Model vacuum may be metastable for low enough Higgs mass, but a specific value of the decay rate holds special significance: for a very narrow window of parameters, our Universe has not yet decayed but the current inflationary period can not be future eternal. Determining whether we are in this  window requires exquisite but achievable experimental precision, with a measurement of the Higgs mass to 0.1 GeV  at the LHC, the top mass to 60 MeV at a linear collider, as well as an improved determination of $\\alpha_s$ by an order of magnitude on the lattice. If the parameters are observed to lie in this special range, particle physics will establish that the future of our Universe is a global big crunch, without harboring pockets of eternal inflation, strongly suggesting  that eternal inflation is censored by the fundamental theory. This conclusion could be drawn even more sharply if metastability with the appropriate decay rate is found in the MSSM, where the physics governing the instability can be directly probed at the TeV scale. ",
        "introduction": "The most remarkable recent discovery in fundamental physics is that the Universe is accelerating \\cite{SSTC,SCPC,WMAP}. Will this acceleration last forever or eventually terminate? What will be the global picture of space-time, and is this a meaningful question? These are some of the most pressing theoretical issues of our time. The idea of eternal inflation \\cite{Vilenkin:1983xq,Linde:1986fd} appears to be supported by the existence of a landscape  of metastable vacua in the string theory \\cite{Bousso:2000xa,Douglas:2006es}. It provides a plausible mechanism to populate the different vacua making it possible to apply environmental arguments to explain the smallness of the cosmological constant \\cite{Weinberg:1987dv}. However,  the necessity to talk about causally disconnected space-time regions to describe eternal inflation makes it extremely challenging to implement this picture in a full theory of quantum gravity. At the moment we do not even know if there are any sharply-defined observables for such a theory to compute; if any exist, they likely reside at future boundaries, with no obvious connection to our own observations (see \\cite{Witten:2001kn,Strominger:2001pn,Freivogel:2006xu,Susskind:2007pv} for some of attempts to make sense of this picture). This problem and the analogy of the emerging causal structure with that of the evaporating black hole, where a naive semiclassical logic fails by predicting the information loss, naturally suggests that semiclassical physics may also be misleading in this case, and a qualitatively new approach is needed to describe the eternally inflating Universe (see, e.g., \\cite{ArkaniHamed:2007ky} for a recent discussion). Clearly, in this situation any experimental data shedding some light on the nature of the cosmological acceleration is extremely welcome.  At the classical level the choice is very simple. Either we live in a local minimum of energy and the cosmic acceleration will last forever, or the acceleration is driven by the potential energy of the slowly rolling scalar field and may eventually terminate. Things change qualitatively at the quantum level. In the slow roll case the Universe can still be eternally inflating if the acceleration parameter is sufficiently close to being constant~\\cite{ArkaniHamed:2007ky}, \\be \\label{HdH4} {\\dot H_\\Lambda\\over H_\\Lambda^4}\\lesssim M_{Pl}^{-2}\\;, \\ee where $H_\\Lambda$ is the Hubble expansion rate corresponding to the current vacuum energy. Numerically, if the bound (\\ref{HdH4}) does not hold  the current stage of the cosmological acceleration may last as long as $\\sim 10^{130}$ years without being eternal. Current data show that the equation of state of the negative energy component is rather close to that of the vacuum energy, $\\dot{H}_\\Lambda/H_\\Lambda^2\\lesssim{\\rm few}\\cdot 10^{-2}$. The bound (\\ref{HdH4}) implies that the next theoretically meaningful value for this ratio is at the level $H_\\Lambda^2/M_{Pl}^2\\sim 10^{-120}$. This indicates that in the slow roll case the  chances to test  the eternal nature of the cosmological acceleration by directly measuring the time dependence of the acceleration parameter   $\\dot{H}_\\Lambda/H_\\Lambda^2$ are not high in the next $10^{130}$ years or so.  The situation is more interesting if we live in a local energy minimum. In the string landscape our positive energy vacuum is inevitably metastable towards decay into a lower energy state \\cite{Coleman:1980aw}. Given the extreme smallness of the vacuum energy this decay is likely to proceed into a rapidly crunching negative energy Universe. The decay goes through the non-perturbative nucleation of  bubbles of the new vacuum that afterwards expand with the speed of light. Whether or not the current stage of the cosmological acceleration is eternal depends on the rate $\\Gamma$ of the bubble nucleation. If the decay rate is fast enough \\be \\label{Gbound} \\Gamma\\geq \\Gamma_{\\rm cr}={9\\over 4\\pi}H_\\Lambda^4\\;, \\ee bubbles collide and eventually all the Universe suffers a transition into the new vacuum. The current stage of inflation is not eternal in this case. On the other hand, for slower decay rates the expansion of the Universe wins and bubbles never fill the whole inflating volume. Of course, bubble walls propagate with the speed of light so we have no chances of observing bubbles without being immediately killed by a domain wall. However, in principle, we can establish whether the bound (\\ref{Gbound}) holds by performing high precision measurements of the underlying particle physics parameters, so that  it becomes possible to identify a nearby negative energy minimum and to calculate a lower bound for the decay rate with high enough precision. We say `a lower bound' because there are always potential instabilities coming from high energy physics we cannot control, but it suffices to find fast enough infrared calculable instabilities to see that inflation cannot be eternal.  High precision measurements are needed since, given that the Universe did not decay up to the present we know that the decay rate $\\Gamma$ cannot be significantly faster than the lower bound (\\ref{Gbound}). Imposing that the probability ($p$) not to decay up to now is larger than, say, 5\\% ($2\\sigma$) gives $\\Gamma\\leq 52.3\\,\\Gamma_{\\rm cr}$, for generic $p$ the rate will be in the window \\beq \\label{eq:decayreg} \\Gamma_{\\rm cr}\\leq\\Gamma\\leq 17.45\\, \\log{(p^{-1})}\\, \\Gamma_{\\rm cr}\\,. \\eeq Given that $\\Gamma$ depends exponentially on the underlying particle physics parameters, we will see that  even a change in the value of $\\Gamma$ by a factor of 50 corresponds to a tiny change of particle masses and coupling constants.  In this paper we argue that, with a certain amount of luck,   challenging but feasible measurements at the TeV scale may establish the existence of a fast enough instability, such that the bound (\\ref{Gbound}) holds. Our main focus will be a minimal scenario of this kind, which is realized if there is no new physics up to  very high energies of order the GUT scale. As reviewed in section~\\ref{sec:SM}, in this case if the Higgs mass is light enough, the Standard Model (SM) potential develops a negative energy minimum at large values of the Higgs field. We show then that one can verify whether the bound (\\ref{Gbound}) holds for this particular decay channel by performing high precision measurements of the Higgs mass (with uncertainty $\\Delta m_H\\sim 0.2$~GeV), the top mass ($\\Delta m_t\\sim 60$~MeV) and the strong coupling constant ($\\Delta\\alpha_s/\\alpha_s\\sim 10^{-3}$). These numbers are achievable with the future data from the LHC and a linear collider. On the theoretical side one will need to improve the existing calculations of the decay rate by including at least one extra electroweak and two strong loops.  Of course, we will never be able to directly verify that the Standard Model holds up to the GUT scale. However, the window for the SM parameters allowing the decay rate to be fast enough, such that the current cosmological acceleration is not eternal, is extremely narrow. Consequently, if no new physics is found at the LHC and the parameters turn out to be in that window, it is reasonable to take this scenario seriously and to try to find an underlying reason why Nature worked so hard to avoid eternal inflation. We outline reasons why Nature might want to censor eternal inflation  in section~\\ref{sec:no}. We stress that we do not necessarily find these arguments plausible---oNe of the AutHors doesn't believe them---they certainly require what looks like a conspiratorial restriction on particle physics models, (although milder examples of surprising gravitational limitations on effective field theories have been seen in other contexts \\cite{Vafa:2005ui,ArkaniHamed:2006dz,Ooguri:2006in}).  It seems much more reasonable that eternal inflation {\\it does} occur, and that we have to learn how to properly deal with it.  Nonetheless the counter arguments are not unreasonable, and it is interesting that entirely feasible experimental observations could strongly push us to take them seriously.  These arguments {\\it predict} that the vacuum decay rate is within the region of eq.~(\\ref{eq:decayreg}). {\\it A priori}, there is no reason to expect this region to be so narrow. However it  is very narrow in the real world. This is nothing but a reflection of the famous cosmic coincidence problem (i.e. why $T_{\\rm U}\\sim H_\\Lambda^{-1}$). At the moment, the best solution to this problem is provided by the Weinberg's argument~\\cite{Weinberg:1987dv}. However, we would like to stress that the fact that the non-eternal inflation bound predicts a very narrow range in the Standard Model parameters is independent of whether the Weinberg's explanation is the correct one and is just a consequence of experimental data. On a similar note, observing the Standard Model parameters in this tiny window would  neither support nor disfavor Weinberg's argument.  The situation can be even more interesting if new physics is discovered at the LHC, as in this case a negative energy minimum could be found at directly accessible energy scales. In section~\\ref{sec:MSSM} we briefly review one scenario of this kind, which is supersymmetry with large soft-breaking trilinear couplings. We conclude by summarizing our results and outlining future directions. ",
        "conclusions": "To summarize,  forthcoming particle physics data may provide a surprising twist for fundamental physics and cosmology, even in the ``nightmare scenario\" where nothing but a light Higgs is observed at the LHC. As has long been known, our vacuum may be metastable for a light enough Higgs. For a very narrow window of parameters, our Universe has not yet decayed but the current inflationary period can not be future eternal. To support this conclusion one would need to measure the Standard Model parameters with extreme precision, requiring significant but achievable progress in the theoretical calculations and experimental measurements.  As a benchmark value for the desired accuracy  we used $0.2$~GeV for the Higgs mass. Given the relation (\\ref{eq:res}) this corresponds to the precision at the level of 60~MeV for the top mass (as opposed to the present $\\sim 1.8$~GeV) , and at the level of $0.14$\\% for the strong coupling (as opposed to the current precision at the level $\\sim1.7$\\%). With this precision one would be able to conclusively establish that the current stage of the cosmological acceleration is not eternal if the decay rate of our vacuum is fast enough, so that the probability that we did not decay up to now is $5$\\%. Even higher precision is needed if the decay rate is slower. Note that  in this case, even if we are not able to resolve whether the Higgs mass lies above or below the percolation bound, finding  the Higgs within $0.2$~GeV from that bound would be remarkable enough to take seriously the idea that eternal inflation may be impossible. This may be true for the heavier Higgs masses, when the incalculable  gravitational corrections may in principle affect the bound at this level of precision.   For the Higgs mass itself the desired level of accuracy is likely to be achievable  directly at the LHC \\cite{TDRs}. Indeed, for a light  Higgs, a very precise determination of the mass is possible by measuring the location of the maximum of the photon distribution from the $H\\to\\gamma\\gamma$ decays. It is important that we need the mass itself; much higher statistics would be needed to determine the width with the same precision.  The top mass precision will not improve much at the LHC~\\cite{TDRs,Stavropoulos:2005he}. The statistical uncertainties will be reduced quite significantly and will be practically negligible. The problem is due to the large systematic uncertainties from the jet energy scale and final state radiation. The best one can hope for is $\\Delta m_t\\sim1$~GeV. However, the situation improves a lot at a linear collider \\cite{Hoang:2000ib,Penin:2002zv}. Here one can use the shape of the cross-section for $t\\bar{t}$ production near threshold to determine the top mass. This method is analogous to the extraction of the bottom mass from the spectrum of the bottomonium. At the moment the theoretical uncertainty from this method has been reduced to the level $\\sim 80$~MeV and may  improve somewhat in the future.  The best determination of $\\alpha_s$ at the moment (and apparently in the reasonable future as well) is from lattice QCD. It appears reasonable  to expect an order of magnitude improvement ({\\it i.e.}, the required level) within the time-scale of constructing a new linear collider.  Finally, on the theoretical side one will need to improve both the running and the matching conditions for gauge couplings and masses by at least one extra electroweak loop and two strong loops. This also requires to include another loop for the calculation of the bounce action.   It is worth mentioning that there is a theoretically motivated natural level of precision, that one may wish to achieve. Indeed, even if the Universe is eternally inflating,  bubbles of new vacuum may form infinite volume clusters if the decay rate is fast enough. The exact value of $\\eps\\equiv\\eps_p$ when this transition (usually referred to as the ``percolation transition\", contrary the terminology adopted in the current paper) happens is currently unknown and is limited to be in the range (\\ref{truepercolation}). Most likely $\\eps_p$ is quite close to the upper boundary $\\eps_p\\sim 0.24$. It appears natural to try to achieve an accuracy allowing to distinguish this critical value $\\eps_p$ from $\\eps=9/4\\pi$ (corresponding to the transition to the eternal regime). If $\\eps_p=0.24$ this would require the Higgs mass resolution at the level of $0.05$~GeV, but may be somewhat less challenging if $\\eps_p$ is smaller.   There are other scenarios motivating the Higgs mass to be in the proximity of the life-time bound. For instance, a  proposal of Ref.~\\cite{Feldstein:2006ce} is based on specific assumptions on how the Standard Model parameters scan in the landscape. More generally, this may be a natural expectation if one adopts the ``living dangerously\" logic~\\cite{Weinberg:1987dv,ArkaniHamed:2005yv}.  The key observation of the current paper is the existence of a sharp critical value of the Higgs mass,  having a direct physical meaning independently of any assumptions on the landscape statistics and on what one understands by being ``close\" to the life-time bound. A high precision is of crucial importance to take the full advantage of the existence of this sharp number. One of the benefits of having a sharp prediction, is that although we will never be able to directly test that there is no new physics up to the GUT scale, finding the Higgs mass in the correct window would provide a strong case for such a scenario.There are also very well-motivated ideas for physics at the TeV scale, such as low-energy SUSY with large $A$~terms, where we can conclude that our vacuum is metastable. In this case high-precision measurements can in principle again tell us that the instability rate is subcritical for eternal inflation, and in this case the conclusion can be drawn much more sharply since the physics of the instability involves no further theoretical extrapolation and can be studied at the TeV scale."
    },
    "0801/0801.2350_arXiv.txt": {
        "abstract": "{The imaging performance of X-ray optics (expressed in terms of HEW, Half-Energy-Width) can be severely affected by X-ray scattering caused by the surface roughness of the mirrors. The impact of X-ray scattering has an increasing relevance for increasing photon energy, and can be the dominant problem in a hard X-ray telescope like SIMBOL-X. In this work we show how, by means of a novel formalism, we can derive a surface roughness tolerance -- in terms of its power spectrum -- from a specific HEW requirement for the SIMBOL-X optical module. ",
        "introduction": "The SIMBOL-X X-ray telescope (\\cite{Pareschi06}) will be the first space observatory in the soft ($E <$~10~keV) and hard ($E >$~10~keV) X-ray band with imaging capabilities comparable to the presently operated soft X-ray telescopes Newton-XMM (15~arcsec HEW) and Swift-XRT (20~arcsec HEW). The SIMBOL-X angular resolution goal is $\\sim$~15~arcsec HEW at 1~keV and $<$ 20~arcsec HEW at 30~keV.  In order to ensure good imaging capabilities to a hard X-ray grazing-incidence optic like those of SIMBOL-X, some critical points have to be carefully evaluated and controlled: \\begin{enumerate} \\item{{\\it Deformation of mirrors profiles, and mirrors misalignment.} The effect on the HEW can be considered as independent of the photon energy $E$, and may be estimated using the geometrical optics (ray-tracing). The impact of this factor is strongly affected by the manufacturing, handling, positioning process of the mirror shells.} \\item{{\\it Microroughness of reflecting surfaces.} It causes the scattering of reflected X-rays (XRS) (see e.g. \\cite{Church86}). This effect falls in the physical optics domain: it has an increasing relevance for increasing photon energy (decreasing photon wavelength $\\lambda$).} \\end{enumerate} As the SIMBOL-X energy band is extended up to 80~keV and beyond, XRS can be the dominant problem at high energies. Hence, a detailed study has to be carried out to establish microroughness tolerances from the angular resolution mission requirements. In the following sections we show how this can be done. ",
        "conclusions": "The evaluation of the XRS contribution to the HEW in the SIMBOL-X hard X-ray telescope is a very important issue, since at high energies it can be the dominant imaging degradation factor, if the reflecting surfaces are not sufficiently smooth. By means of the formulae mentioned above, a given HEW($E$) requirement -- in the SIMBOL-X energy band -- can be directly translated into important information concerning the surface microroughness tolerance, conveniently expressed in terms of PSD. Conversely, we can use the PSD of the mirror roughness -- measured over a very wide spatial frequency range -- to preview the HEW($E$) function of a single mirror shell, in the range of photon energies and incidence angles of SIMBOL-X. The HEW can be calculated for each mirror shell from the surface PSD, and from that we can easily recover the HEW estimation for the entire optical module.  Calculations performed on the PSD of samples replicated from smooth masters by Nickel electroforming proved that the effort has to be concentrated on the damping of the microroughness at spatial wavelengths shorter than $\\sim$~300~$\\mu$m, in order to improve the angular resolution at high energies and therefore to meet the SIMBOL-X angular resolution requirements in soft and hard X-rays."
    },
    "0801/0801.0739_arXiv.txt": {
        "abstract": "Supermassive black hole binaries (BHBs) produced in galaxy mergers recoil at the time of their coalescence due to the emission of gravitational waves (GWs).  We simulate the response of a thin, two--dimensional disk of collisionless particles, initially on circular orbits around a $10^6~{\\rm M_\\odot}$ BHB, to kicks that are either parallel or perpendicular to the initial orbital plane. Typical kick velocities ($v_{\\rm kick}$) can exceed the sound speed in a circumbinary gas disk. While the inner disk is strongly bound to the recoiling binary, the outer disk is only weakly bound or unbound.  This leads to differential motions in the disturbed disk that increase with radius and can become supersonic at $\\gsim 700$ Schwarzschild radii for $v_{\\rm kick}=500~{\\rm km~s^{-1}}$, implying that shocks form beyond this radius.  We indeed find that kicks in the disk plane lead to immediate strong density enhancements (within weeks) in a tightly wound spiral caustic, propagating outward at the speed $\\sim v_{\\rm kick}$. Concentric density enhancements are also observed for kicks perpendicular to the disk, but are weaker and develop into caustics only after a long delay ($>$one year). Unless both BH spins are low or precisely aligned with the orbital angular momentum, a significant fraction ($\\gsim$ several \\%) of kicks are sufficiently large and well aligned with the orbital plane for strong shocks to be produced. The shocks could result in an afterglow whose characteristic photon energy increases with time, from the UV ($\\sim 10$eV) to the soft X--ray ($\\sim 100$eV) range, between one month and one year after the merger.  This could help identify EM counterparts to GW sources discovered by {\\it LISA}. ",
        "introduction": "The recent break--through in numerical relativity has allowed a direct computation of the linear momentum flux produced during the coalescence of a BH binary (Baker et al. 2006, 2007; Campanelli et al. 2007a,b; Gonz\\'alez et al. 2007a,b; Herrmann et al. 2007a,c; Koppitz et al. 2007).  The resulting final recoil depends on the masses, orbital parameters, and spins of the BHs, and in special configurations (i.e. with spins anti-aligned with each other), it can reach velocities as high as $\\approx 4,000~{\\rm km~s^{-1}}$ (Gonz\\'alez 2007b).  Much of the recent literature focused on the astrophysical implications of high--velocity kicks, which may displace or remove supermassive BHs from galactic centers (e.g. Merritt et al. 2004; Madau \\& Quataert 2004; Loeb 2007), and inhibit the growth of BHs at high redshifts ($z\\gsim 6$), where the escape velocities from low--mass galactic halos is small ($\\lsim 100~{\\rm km~s^{-1}}$; Haiman 2004, Yoo \\& Miralda-Escude 2004, Shapiro 2005; Volonteri \\& Perna 2005).  Another exciting implication of kicks is that they may help produce an electromagnetic (EM) counterpart of gravitational wave sources detected by the future {\\it Laser Interferometric Space Antenna} ({\\it LISA}) satellite. The discovery of such a counterpart would constitute a milestone for fundamental physics and astrophysics (e.g. Kocsis et al. 2007b). If the BHB is surrounded by a circumbinary gas disk, the disk will respond promptly (on the local orbital timescale) to such a kick. If this results in warps or shocks, the disturbed disk could produce a transient EM signature (Milosavljevi\\'c \\& Phinney 2005). The sky localization uncertainty from the {\\it LISA} instrument several weeks prior to merger is typically a few square degrees (Kocsis et al. 2007a; Lang \\& Hughes 2007), and the kick can begin building up at the end of the inspiral phase, before the final coalescence (Schnittman et al. 2007).  This will make it possible, in many cases, to monitor a few square--degree area on the sky prior to, during, as well as immediately following the coalescence of BHs in the mass range $\\sim (10^5$--$10^7)\\,{\\rm M_\\odot}/(1+z)$ at redshifts out to $z\\sim 3$, and search for a prompt transient signature associated with the kick (Kocsis et al. 2007b).  In the context of producing a prompt EM counterpart, we expect that the {\\it direction}, in addition to the magnitude of the kick, will be important. Naively, one expects that a kick within the plane of any circumbinary disk will be more likely to cause density enhancements and ``light up'' the disk than a kick perpendicular to it.  The kick direction can be clearly important on larger scales, as well, and affect phenomena that would occur on long timescales ($\\gg$ years) after the merger. For example, the angle relative to a large--scale galactic disk can determine whether a kicked BHB ends up outside the galaxy or not, and perhaps also whether shocks are produced when a BHB plunges into such a large--scale disk.  While these issues merit investigation, in this {\\em Letter} we focus on the smaller scales and shorter timescales ($\\lsim$ months) that are relevant to prompt {\\it LISA} counterparts.  Schnittman \\& Buonanno (2007; hereafter SB07) used the ``effective one--body approach'' and derived a scaling formula that yields the recoil velocity vector for arbitrary mass ratios and spin vectors. While their results have not been tested for generic spins, they agree well (to within 20-30\\%) with numerical results in those special configurations where they were tested (including numerical calculations for configurations with the spins parallel or at intermediate angles with respect to the orbital plane; see Campanelli et al. 2007a,b, Gonz\\'alez et al. 2007a; Herrmann et al. 2007b, Tichy \\& Marronetti 2007).  In this {\\em Letter}, we investigate the response of a circumbinary disk to the kick, by following the perturbed, Keplerian orbits of collisionless massless test particles around the recoiling BHB.  The purpose of this exercise is to demonstrate that prompt shocks or strong density enhancements are likely to arise when the kick is aligned with the plane of the circumbinary disk, whereas they may be less likely for highly inclined kicks.  We will then use the formula of SB07 for the dependence of the kick speed and direction on the mass ratio and spins of the BHs, to argue that a significant fraction ($\\gsim$ few \\%) of kicks may be both sufficiently large and well aligned with the orbital plane for strong shocks to be produced within a few weeks after coalescence. ",
        "conclusions": "\\label{sec:discuss}  The main result of this {\\it Letter} is that strong density enhancements can form promptly after a supersonic kick in the plane of the circumbinary disk, within a few weeks of the coalescence of a $\\sim 10^6~{\\rm M_\\odot}$ BHB.  Because the disk is cold, and caustics (Fig.~\\ref{fig:spirals}) are formed when particles first cross each other along their orbits, this implies that corresponding shocks could occur in a gas disk.  For hydrodynamical shocks to occur within a finite--pressure gas, the relative motions $v_{\\rm c}$ between the neighboring particles that produce the caustic must exceed the sound speed.  At the outermost radius where the disk is marginally bound to the BHB, one expects $v_{\\rm c}\\sim v_{\\rm kick} \\sim v_{\\rm orbit}$; relative motions will be slower further inside.  The relative speed should roughly correspond to covering the epicyclic amplitude $\\sim (v_{\\rm kick}/v_{\\rm orbit}) r_c$ in the caustic--formation time $t_{\\rm c}\\sim r_c/v_{\\rm kick}$, yielding $v_c\\sim v_{\\rm kick}^2/v_{\\rm orbit}$.  For $v_{\\rm kick}= 500 {\\rm km~s^{-1}}$, this predicts $v_{\\rm c} \\sim 25 {\\rm km~s^{-1}} (r/1000r_{\\rm S})^{1/2}$; we have verified in our simulations that particles cross the caustics with speeds at about $\\sim 30\\%$ above this predicted value.  Compared with the sound speed $c_s\\approx 25 {\\rm km~s^{-1}} (r/1000r_{\\rm S})^{-9/20}$, this suggests that the density waves produced by the kick in the gas beyond $\\sim 700r_{\\rm S}$ will indeed steepen into shocks.  We also found that the inclination of the kick may be important in determining the strength and timing of such shocks -- perpendicular kicks would only produce weaker density enhancements, at least until a delay of about a year.  A non--negligible fraction ($\\gsim$ several \\%) of kicks could, however, be sufficiently large and well aligned with the orbital plane for shocks to be produced within a few weeks after coalescence.  The nature of the emission resulting from the shocks or density enhancements will have to be addressed in future work, by computing the heating rate at the spiral shocks, and modeling the overall disk structure and vertical radiation transport.  However, the disk's luminosity resulting from the kick, $L_{\\rm kick}$ could be a non--negligible fraction of the BHB's Eddington luminosity, and therefore potentially observable (Kocsis et al. 2007b).  If $M_{\\rm shock}=f_{\\rm shock}(M_1+M_2)$ is the mass of the shocked gas that is heated to temperatures corresponding to $v_{\\rm shock}$, and $t_{\\rm shock}$ is the time--scale on which the corresponding thermal energy is converted to photons, then $L_{\\rm kick} \\approx (1/2) M_{\\rm shock} v_{\\rm shock}^2 / t_{\\rm shock}$, and $L_{\\rm kick}/L_{\\rm Edd} \\approx 0.1 (f_{\\rm shock}/10^{-5}) (v_{\\rm shock}/500~{\\rm km~s^{-1}})^2 (t_{\\rm shock}/{\\rm 1 month})^{-1}$.  We may also speculate on the spectral evolution of the ``kick after--glow''. Assuming $M_{\\rm shock} \\propto \\Sigma r dr \\propto r ^{19/10}$ (with $\\Sigma \\propto r^{-3/5}$ and $dr\\propto r^{1/2}$, the epicyclic amplitude), and $v_{\\rm shock} \\approx v_c$, we find $L_{\\rm kick} \\propto r^{24/10}$. This suggests that the luminosity may be dominated by the outermost shocked shells. The spectrum will then peak at the characteristic photon energy corresponding to $k T_{\\rm shock}\\propto v_{\\rm c}^2 \\propto v_{\\rm orbit}^{-2} \\propto r$.  The shocks could therefore result in an afterglow, starting from $r_{\\rm cavity}/v_{\\rm kick}\\sim 30$ days, first peaking in the UV band ($\\sim 10$eV), and then hardening to the soft X--ray ($\\sim 100$eV) range after one year. The detection of such an afterglow would help identify EM counterparts to GW sources discovered by {\\it LISA}.  Our results need to be verified in three--dimensional simulations that resolve the orbits within a finite disk thickness.  A realistic disk model, incorporating gas dynamics, is needed to study the correspondence between collisionless caustics and gaseous shocks. Finally, the SB07 formula needs to be confirmed for generic spin configurations.  Nevertheless, our results do suggest that kicks due to gravitational waves may produce a prompt EM signal.   We thank A. Buonanno, M. Milosavljevi\\'c, K. Menou, and B. Kocsis for useful discussions, and M. Milosavljevi\\'c for suggesting the arguments based on epicycles. This work was supported by the Pol\\'anyi Program of the Hungarian National Office for Research and Technology (NKTH) and by NASA grant NNG04GI88G (to ZH)."
    },
    "0801/0801.3606_arXiv.txt": {
        "abstract": "We study Tully Fisher relations for a sample that combines extremely faint (M$_{\\rm B} < -14.$) galaxies along with bright (i.e. $\\sim L_{*}$) galaxies. Accurate ($\\sim 10\\%$) distances, I band photometry, and B-V colors are known for the majority of the galaxies in our sample. The faint galaxies are drawn from the Faint Irregular Galaxy GMRT survey (FIGGS), and we have HI rotation velocities derived from aperture synthesis observations for all of them. For the faint galaxies, we find that even though the median HI and stellar masses are  comparable, the HI mass correlates significantly better with the circular velocity indicators than the stellar mass. We also find  that the velocity width at the 20\\% level (W$_{20}$) correlates better with mass than the rotation velocity, although the difference is not statistically significant. The faint galaxies lie systematically below the I band TF relation defined by bright galaxies, and also show significantly more intrinsic scatter. This implies that the integrated star formation in these galaxies has been both less efficient and also less regulated than in large galaxies. We estimate the intrinsic scatter of the faint galaxies about the I band TF to be $\\sim 1.6$~mag. We find that while the faint end deviation is greatly reduced in Baryonic Tully-Fisher (BTF) relations, the existence of a break at the faint end of the BTF is subject to systematics such as the assumed stellar mass to light ratio. If we {\\it assume} that there is an intrinsic BTF and try to determine the baryonic mass by searching for prescriptions that lead to the tightest BTF, we find that scaling the HI mass leads to a much more significant tightening than scaling the stellar mass to light ratio. The most significant tightening that we find however, is if we scale the entire baryonic mass of the faint (but not the bright) galaxies. Such a scenario would be consistent with models where dwarf (but not large) galaxies have a large fraction of dark or ``missing'' baryons. In all cases however the minimum in the $\\chi^2$ curve is quite broad and the corresponding parameters are poorly constrained. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.4210_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{intro}  Stars form from the gravitational collapse of dense cloud cores in the molecular interstellar medium of galaxies.  Studying and characterizing the properties of  dense cores is thus of great interest to gain insight into the initial conditions and initial stages of the star formation process.  Our observational understanding of low-mass dense cores has made significant progress in recent years and three broad categories of cores can now be distinguished within nearby molecular clouds, which possibly represent an evolutionary sequence: starless cores, prestellar cores, and ``Class~0'' protostellar cores. Starless cores are possibly transient concentrations of molecular gas and dust without embedded young stellar objects (YSOs), typically observed in tracers such as C$^{18}$O (e.g. Onishi et al. 1998), NH$_{3}$ (e.g. Jijina, Myers, \\& Adams 1999), or dust extinction (e.g. Alves et al. 2007), and which do not show evidence of infall. Prestellar cores are also starless ($M_\\star = 0$) but represent a somewhat denser and more centrally-concentrated population of cores which are self-gravitating, hence unlikely to be transient. They are typically detected in (sub)millimeter dust continuum emission and dense molecular gas tracers such as NH$_3$ or N$_2$H$^+$ (e.g. Ward-Thompson et al. 1994; Benson \\& Myers 1989; Caselli et al. 2002), often seen in absorption at mid- to far-infrared wavelengths (e.g. Bacmann et al. 2000, Alves et al. 2001), and frequently exhibit evidence of infall motions (e.g. Gregersen \\& Evans 2000). Conceptually, all prestellar cores are starless but only a subset of the starless cores evolve into prestellar cores; the rest are presumably ``failed'' cores that eventually disperse and never form stars.  In practice, prestellar cores are characterized by large density contrasts over the local background medium. Specifically, the mean densities of observed prestellar cores exceed the mean densities of their parent clouds by a factor $\\simgt $~5--10, while their mean column densities exceed the background column densities by a factor $\\simgt $~2. For comparison, a critical self-gravitating Bonnor-Ebert isothermal spheroid has a mean density contrast $\\bar{\\rho}_{\\rm BE}/\\rho_{\\rm ext} \\sim 2.4$ (e.g. Lombardi \\& Bertin 2001) and a mean column density contrast $\\bar{\\Sigma}_{\\rm BE}/\\Sigma_{\\rm ext} \\sim 1.5$ over the external medium.  Finally, Class~0 cores/objects are young accreting protostars observed early after point mass formation while most of the mass of the system is still in the form of a dense core/envelope as opposed to a YSO ($M_\\star \\ll M_{\\rm env} $) (Andr\\'e, Ward-Thompson, Barsony 1993). They are believed to result from the gravitational collapse of prestellar cores. Class~0 protostars themselves evolve into Class~I objects with $M_{\\rm env} < M_\\star$ (Lada 1987;  Andr\\'e \\& Montmerle 1994) as the protostellar envelope dissipates through accretion and ejection of circumstellar material.  Class~I objects subsequently evolve into (Class~II and Class~III)  pre-main sequence stars surrounded by a circumstellar disk (optically thick and optically thin in the near-/mid-IR, respectively), but lacking a dense circumstellar envelope ($M_{\\rm env} \\sim 0$).  Improving our understanding of the formation and evolution of dense cores in molecular clouds is crucially important since there is now good evidence that these early stages largely control the origin of the stellar initial mass function (IMF). Indeed, observations indicate that the prestellar core mass function resembles the IMF (see \\S ~\\ref{sec:surveys} below),  suggesting that the effective reservoirs of mass required for the formation of individual stars are already selected at the prestellar core stage. Furthermore, it is at the very end of the prestellar stage that multiple systems are believed to form and during the protostellar stage that a fraction of the prestellar core mass reservoir is accreted by the central protostellar components as a result of the accretion/ejection process.  \\begin{figure} \\centering \\includegraphics[width=8cm,angle=270]{cores_chamonix_fig1.ps} \\caption{SCUBA $850\\:\\mu$m dust continuum map of the NGC~2068 protocluster extracted from the mosaic of Orion~B by Motte et al. (2001). A total of 30 compact prestellar condensations (marked by crosses), with masses between $\\sim 0.4\\, M_\\odot $ and $\\sim 4.5\\, M_\\odot $, are detected in this $\\sim 1$~pc $\\times$ 0.7~pc field.} \\label{fig:ngc2068} \\end{figure} ",
        "conclusions": "\\label{conclusions}  The prestellar core mass function (CMF) appears to be consistent with the stellar IMF between $\\sim 0.1\\, M_\\odot$ and $\\sim 5\\, M_\\odot$, although large uncertainties remain especially at the low- and high-mass ends (cf. \\S ~\\ref{sec:surveys}). Small internal and relative motions are measured for these prestellar cores, implying that they are much less turbulent than their parent cloud and generally do not have time to interact before collapsing to (proto)stars (cf. \\S ~\\ref{velocity}). These results strongly support scenarios according to which the IMF is largely determined at the prestellar stage.  None of the extreme scenarios proposed for the formation of prestellar cores can explain all observations. Pure ambipolar diffusion is too slow to be the main core formation mechanism, for typical levels of cloud ionization. Purely hydrodynamic pictures have trouble accounting for the inefficiency of core formation and the detailed velocity structure of individual cores. A mixed scenario (cf. Nakamura \\& Li 2005, Basu et al. 2008) may be the solution: supersonic MHD turbulence in a molecular cloud close to magnetic criticality generates seed cores, which grow in mass until they become gravitationally unstable and collapse in a magnetically-controlled fashion while decoupling from their turbulent environment. Such a mixed picture has  several advantages. On the one hand, relatively strong magnetic fields prevent global collapse and lead to inefficient core/star formation on large (GMC) scales. On the other hand, the rate of ambipolar diffusion is enhanced by turbulence in shock-compressed regions (e.g. Zweibel 2002, Fatuzzo \\& Adams 2002, Nakamura \\& Li 2005). The difference between the clustered and the isolated or distributed mode of star formation may result from local variations around the critical mass-to-flux ratio and/or from the presence or absence of external triggers.  Cluster-forming clumps such as L1688 in Ophiuchus or NGC2264-C in Monoceros are likely (slightly) magnetically supercritical. There is some evidence that these clumps are in a state of global collapse induced by large-scale external triggers (e.g. Peretto et al. 2006, Nutter et al. 2006). In this case, the local star formation efficiency within each condensation is high ($\\simgt 50\\% $ -- Motte et al. 1998) and a promising core formation mechanism  is purely gravitational, Jeans-like fragmentation of compressed cloud layers (cf. Palous 2007, Peretto et al. 2007), possibly followed by subsequent core growth (e.g. Basu \\& Jones 2004).  Regions of more distributed and less efficient star formation, such as Taurus or the Pipe Nebula, are likely slightly subcritical or nearly critical. The size of individual cores is then larger (e.g. Motte et al. 1998), the local star formation efficiency is lower ($\\sim$~15\\%--30\\% -- Onishi et al. 2002, Alves et al. 2007), and the feedback from protostellar outflows may be more important in limiting accretion and defining stellar masses (e.g. Shu et al. 2004).  To fully understand how stars form and how the IMF comes about, a better knowledge of the presently poorly known initial conditions for molecular cloud formation is required, a subject of growing interest (see, e.g., Hennebelle et al. in this volume). It is also crucial to further investigate the processes by which prestellar cores form, evolve, and eventually collapse and fragment into multiple protostellar systems. With present submillimeter instrumentation, observational studies are limited by small-number statistics and restricted to the nearest regions. The advent of major new facilities in the coming years should yield several breakthroughs in this field. With an angular resolution at 75--300~$\\mu$m comparable to, or better than, the largest ground-based millimeter-wave radiotelescopes, $Herschel$, the Far InfraRed and Submillimeter Telescope to be launched by ESA in 2008 (cf. Pilbratt 2005), will make possible complete surveys for prestellar cores down to the proto-brown dwarf regime in the cloud complexes of the Gould Belt (cf. Andr\\'e \\& Saraceno 2005). High-resolution ($0.01^{\\prime\\prime}-0.1^{\\prime\\prime}$) studies with the `Atacama Large Millimeter Array' (ALMA, becoming partly available in 2011, fully operational in 2013 -- cf. Bachiller 2008) at $\\sim $~450~$\\mu $m~--~3~mm will allow us to probe the kinematics of individual condensations in distant, massive protoclusters. Complementing each other nicely, $Herschel$ and ALMA will tremendously improve our global understanding of the initial stages of star formation in the Galaxy."
    },
    "0801/0801.1329_arXiv.txt": {
        "abstract": "As part of an effort to understand the origin of open clusters, we present a statistical analysis of the currently observed Pleiades. Starting with a photometric catalog of the cluster, we employ a maximum likelihood technique to determine the mass distribution of its members, including single stars and both components of binary systems. We find that the overall binary fraction for unresolved pairs is 68\\%. Extrapolating to include resolved systems, this fraction climbs to about 76\\%, significantly higher than the accepted field-star result. Both figures are sensitive to the cluster age, for which we have used the currently favored value of 125~Myr. The primary and secondary masses within binaries are correlated, in the sense that their ratios are closer to unity than under the hypothesis of random pairing. We map out the spatial variation of the cluster's projected and three-dimensional mass and number densities. Finally, we revisit the issue of mass segregation in the Pleiades. We find unambiguous evidence of segregation, and introduce a new method for quantifying it. ",
        "introduction": "Open clusters, with their dense central concentrations of stars, are relatively easy to identify. Over a thousand systems are known, and the census is thought to be complete out to 2~kpc \\citep{br01,d02}. Because the clusters are no longer buried within interstellar gas and dust, their internal structure and dynamics is also more accessible than for younger groups.  Despite these favorable circumstances, many basic questions remain unanswered. Most fundamentally, how do open clusters form? All observed systems have undergone some degree of dynamical relaxation. Thus, the present-day distribution of stellar mass differs from the one just after disruption of the parent cloud. Recovering this initial configuration will clearly be of value in addressing the formation issue. But such reconstruction presupposes, and indeed requires, that we gauge accurately the stellar content of {\\it present-day} clusters. Even this simpler issue is a non-trivial one, as we show here.  We consider one of the most intensively studied open clusters, the Pleiades. Again, our ultimate goal is to trace its evolution from the earliest, embedded phase to the present epoch. Focusing here on the latter, we ask the following questions: What are the actual masses of the member stars? Do they follow the field-star initial mass function? How many of the members are single, and how many are in binary pairs? Are the primary and secondary masses of binaries correlated? What is the overall density distribution in the cluster? What is the evidence for mass segregation, and how can this phenomenon be quantified?  All of these questions have been addressed previously by others. \\citet{dh04} constructed a global mass function for the Pleiades. Their method was to assign masses based on the observed distribution of $R$-magnitudes. A more accurate assessment should account for the photometric influence of binaries. Several studies directed specifically at binaries have probed selected regions for spectroscopic pairs \\citep{rm98,b97}. However, the overall binary fraction has not been carefully assessed, despite some preliminary attempts \\citep{sj95,b97,m03}.  A fuller investigation of these and the other issues raised requires statistical methods; these should prove generally useful in characterizing stellar populations. Section~2 describes our approach, which employs a regularized maximum likelihood technique \\citep[e.g.][]{c98}. A similar method has been applied to other astronomical problems, including the reconstruction of cloud shapes \\citep{t07} and the investigation of binarity within globular clusters \\citep{rw91}. Our study is the first to apply this versatile tool to young stellar groups. In doing so, we also relax many of the restrictive assumptions adopted by previous researchers. Section~3 presents our derived mass function for the Pleiades, along with our results for binarity. The density distribution, as well our quantification of mass segregation, are the topics of Section~4. Finally, Section~5 summarizes our findings, critically reexamines the binarity issue, and indicates the next steps in this continuing study. ",
        "conclusions": "In this paper, we have applied a versatile statistical tool, the maximum likelihood technique, to assess the distribution of stellar mass and the incidence of binaries in the Pleiades. We began with a near-infrared catalog of cluster members. Our basic assumption was that all cluster members share the same evolutionary age, and that any dispersion in the color-magnitude plane stems from binarity and random photometric errors. We were then able to infer the most probable distribution of masses, both for the cluster as a whole, and as a function of distance from its center. Finally, we introduced a simple method for gauging the degree of mass segregation in the cluster.  One of our surprising results is the relatively high fraction of binaries. We estimate that 68\\% of all systems in the cluster are unresolved binaries; this figure climbs to about 76\\% if resolved pairs are included. These fractions are significantly higher than the accepted field-star result, so we should scrutinize them carefully. Could they stem from an underestimate of the photometric error at faint magnitudes? Since the error in $I$ is greater than $K$, we artificially increased the dispersion $\\sigma_I$. We kept $\\sigma_I$ at 0.15 until \\hbox{$M_I \\,=\\,9.5$}, below which we increased it linearly, reaching \\hbox{$\\sigma_I\\,=\\,0.20$} at \\hbox{$M_I \\,=\\, 12$}. After redoing the maximum likelihood analysis, the global binary fraction $b$ for unresolved pairs is unchanged.  Another potential difficulty is our neglect of the physical thickness of the cluster. We have assigned all members a distance of 133~pc, although there will naturally be some variation. However, this effect is also relatively small. From Section~4.1, the volumetric number density falls off with radius approximately as a King model with core and tidal radii of 2.1 and 19~pc, respectively. Consider the front half of a spherical cluster with such a density distribution. It may readily be shown that the mean distance from the plane of the sky of any cluster member is \\hbox{$d\\,=\\,2.6\\,{\\rm pc}$}. For a cluster at mean distance $D$, the induced magnitude spread is \\hbox{$5\\ {\\rm log}\\,[(D+d)/D]$}, which is 0.04 in our case. Although the actual spread in magnitudes is not Gaussian, we have added this figure in quadrature to both $\\sigma_I$ and $\\sigma_K$, and rerun the analysis. Again, the binarity is unaffected.  The errors due to both photometry and finite cluster thickness induce a symmetric spread in stellar magnitudes. That is, they scatter as many sources below the fiducial isochrone as above it. Thus, they cannot reduce the estimated binarity, which stems from an excess of stars above the isochrone. One systematic error that {\\it would} affect $b$ is an overestimation of the cluster distance. If $D$ were lowered, the absolute magnitudes of all sources would decrease equally, as would the inferred $b$-value. Quantitatively, the distance would have to decrease by about 15~pc to bring the binary fraction down to the field-star result for G-dwarf primaries. An error of this size for the average distance is excluded by current observations, for which the estimated uncertainty is only 1~pc \\citep{s05}.  Since our method relies solely on photometry to assess binarity, we cannot distinguish between physically linked pairs and chance alignments. As mentioned in Section~3.4, the resolution limit of our data is $10^{\\prime\\prime}$, or \\hbox{$\\Delta r_\\circ \\,=\\, 1400~{\\rm AU}$} at the distance of the Pleiades. Consider a star at a radius $r$ from the cluster center. Its average number of neighbors within $\\Delta r_\\circ$ is \\hbox{$\\pi\\,\\Delta r_\\circ^2\\,n_s (r)$},where $n_s (r)$ is the projected surface number density of the cluster. Since each ring of width $dr$ contains \\hbox{$2\\,\\pi\\,n_s\\,r\\,dr$} stars, integration over all members yields the total number of chance alignments: \\begin{equation} N_{\\rm chance} \\,=\\, 2\\,\\pi^2\\,\\Delta r_\\circ^2 \\int_0^R\\!n_s^2 (r)\\,r\\,dr \\,\\,. \\end{equation} where $R$ is the cluster's outer radius, Using $n_s (r)$ from Figure~10, we find\\hbox{$N_{\\rm chance} \\,=\\,2.4$}. Thus, chance alignments have no quantitative impact.  Yet another source of systematic error is the cluster age. We have adopted the lithium-based figure of 125~Myr from \\citet{sta98}. Earlier estimates, using the main-sequence turnoff, yielded a range of answers. For example, \\citet{m93} found 100~Myr. Even this minor reduction affects our results, since it lifts the low-mass end of the isochrone toward higher luminosity. For an age of 100~Myr, our analysis gives \\hbox{$b\\,=\\,0.57\\,\\pm\\,0.02$} and \\hbox{$c\\,=\\,0.28\\,\\pm\\,0.06$}. The binary fraction is augmented to 0.64 when we include resolved pairs. From Figure~3 of \\citet{sta98}, a 100~Myr age corresponds to a lithium edge at \\hbox{$M_I\\,=\\,11.7$}, or \\hbox{$I\\,=\\,17.3$} at the Pleiades distance. Such a result seems incompatible with the lithium data shown in Figure~2 of \\citet{sta98}, but the total number of observations is relatively small.  We conclude that the enhanced binarity is a real effect. What this fact tells us about the origin of open clusters remains to be seen. Our next step in addressing this basic issue is to try and account theoretically for the empirical properties just obtained through our statistical analysis. We will ascertain, using direct numerical simulations, the range of initial states that can relax dynamically to the present-day Pleiades. Such a study will bring us one step closer to understanding the molecular cloud environments that give rise to open clusters."
    },
    "0801/0801.3430_arXiv.txt": {
        "abstract": "{ We present a study of the structural and scaling properties of the gas distributions in the intracluster medium (ICM) of 31 nearby ($z < 0.2$) clusters observed with {\\it XMM-Newton}, which together comprise the Representative {\\it XMM-Newton} Cluster Structure Survey (REXCESS). In contrast to previous studies, this sample is unbiased with respect to X-ray surface brightness and cluster dynamical state, and it fully samples the cluster X-ray luminosity function. The clusters cover a temperature range of 2.0 -- 8.5 keV and possess a variety of morphologies. The sampling strategy allows us to compare clusters with a wide range of central cooling times on an equal footing. We applied a recently developed technique for the deprojection and PSF-deconvolution of X-ray surface brightness profiles to obtain non-parametric gas-density profiles out to distances ranging between $0.8R_{500}$ and $1.5 R_{500}$. We scaled the gas density distributions to allow for the systems' differing masses and redshifts. The central gas densities differ greatly from system to system, with no clear correlation with system temperature. At intermediate radii ($\\sim 0.3 R_{500}$), the scaled density profiles show much less scatter, with a clear dependence on system temperature. We find that the density at this radius scales proportionally to the square root of temperature, consistent with the presence of an entropy excess as suggested in previous literature. However, at larger scaled radii this dependence becomes weaker: clusters with $kT > 3$ keV scale self-similarly, with no temperature dependence of gas-density normalisation. The REXCESS sample allows us to investigate the correlations between cluster properties and dynamical state. We find no evidence of correlations between cluster dynamical state and either the gas density slope in the inner regions or temperature, but do find some evidence of a correlation between dynamical state and outer gas density slope. We also find a weak correlation between dynamical state and both central gas normalisation and inner cooling times, but this is only significant at the 10\\% level. We conclude that, for the X-ray cluster population as a whole, both the central gas properties and the angle-averaged, large-scale gas properties are linked to the cluster dynamical state. We also investigate the central cooling times of the clusters. While the cooling times span a wide range, we find no evidence of a significant bimodality in the distributions of central density, density gradient, or cooling time. Finally, we present the gas mass-temperature relation for the REXCESS sample, finding that $h(z)M_{gas} \\propto T^{1.99\\pm0.11}$, which is consistent with the expectation of self-similar scaling modified by the presence of an entropy excess in the inner regions of the cluster and consistent with earlier work on relaxed cluster samples. We measure a logarithmic intrinsic scatter in this relation of $\\sim 10\\%$, which should be a good measure of the intrinsic scatter in the $M_{gas}$--$T$ relation for the cluster population as a whole.}  \\authorrunning{J.H. Croston et al.} \\titlerunning{REXCESS gas density distributions} ",
        "introduction": "The X-ray emitting gas in galaxy clusters contains the signatures of important evolutionary processes such as mergers, AGN activity, galaxy interactions and tidal stripping. The impact of gas physical processes on the observable X-ray properties of galaxy clusters must be fully understood in order to use galaxy clusters to test the predictions of structure formation models, and to understand the relationship between AGN activity, galaxy evolution and the evolution of large-scale structure. The global scaling relations between cluster observables such as X-ray luminosity/temperature and cluster mass must be well constrained in the local Universe so that cluster evolution to high redshifts can be investigated; this is not possible without a good understanding of the processes that lead to the observed scatter in these relationships.   Since the first evidence that the gas and dark matter content of galaxy clusters does not scale according to the simplest self-similar predictions (e.g. Edge \\& Stewart 1991; Arnaud \\& Evrard 1999), non-gravitational heating processes have been thought to play an important role in determining the X-ray properties of galaxy clusters. Recent results have shown that the dark matter component of galaxy clusters does scale self-similarly, in good agreement with theoretical predictions (e.g. Pointecouteau et al. 2005; Vikhlinin et al. 2006). It has also been shown that the departure of the observed $L_{X}/T_{X}$ relation from theoretical predictions is due to increased entropy in cluster gas (e.g. Ponman et al. 2003; Pratt, Arnaud \\& Pointecouteau 2006). A combination of radiative cooling and feedback processes associated with galaxy formation, either from galaxy winds or AGN activity associated with central supermassive black holes, is the most widely accepted explanation for these results (e.g. Voit 2005). Cluster mergers are also expected to have an important effect on their observed X-ray properties: for example, simulations have shown that cluster merger processes are likely to have an important effect on the $L_{X}/T_{X}$ relation (Rowley et al. 2004), although cooling and heating processes in the inner regions may be the dominant source of scatter (O'Hara et al. 2006).   Radial surface brightness profiles and gas density profiles have been the key means of obtaining information about the structure and scaling properties of the intracluster medium (ICM) since the advent of X-ray imaging of clusters, providing an essential tool to enable study of three-dimensional gas distributions, and, together with temperature, the study of the gas entropy distributions and the total mass profiles via the assumption of hydrostatic equilibrium. {\\it ROSAT} studies (e.g. Neumann \\& Arnaud 1999) suggested that a $\\beta$ model was an adequate parametrization of the X-ray surface brightness distribution outside the core region in galaxy clusters, with clear evidence for self-similarity, but a large dispersion in the central regions. The core regions of clusters are now much better resolved with {\\it XMM-Newton} and {\\it Chandra}, and more complex models are required to fit their surface brightness profiles (e.g. Pratt \\& Arnaud 2002; Pointecouteau et al. 2004; Vikhlinin et al. 2006). There is also some evidence that the slope of the gas density profile may steepen at large radii in some clusters (e.g. Viklinin et al. 1999; Neumann 2005; Vikhlinin et al. 2006). Constraints on the radial distribution of hot gas in the galaxy cluster population are crucial both for accurate estimation of total cluster masses and for studies of gas entropy. Spherically symmetric methods for constraining the gas distributions of galaxy clusters have obvious limitations when applied to samples that are not selected for regularity; however, as a straightforward method that is easy to apply to both observations and simulations, they remain an essential tool for studying cluster properties in three dimensions.   \\begin{table*}[t] \\begin{minipage}[t]{17.5cm} \\caption{Cluster properties. Columns: (1) Cluster ID, (2) redshift, (3) column density in $10^{20}$ cm$^{-2}$, (4) X-ray temperature from a {\\it mekal} fit in the $[0.15 - 0.75] R_{500}$ region in keV, (5) X-ray temperature from a {\\it mekal} fit in the $[0.15-1]R_{500}$ region in keV, (6) abundance obtained from a {\\it mekal} fit in the $[0.15 - 1] R_{500}$ region, in units of solar abundance, (7) $R_{500}$ in kpc, iteratively determined from the $M_{500}-Y_{X}$ relation of Arnaud, Pointecouteau \\& Pratt (2007) as described in the text, (8) logarithm of gas mass to $R_{500}$ determined as described in the text in $M_{\\sun}$, (9) inner slope of the gas density profile determined in the region with $r < 0.05 R_{500}$, (10) outer slope of the gas density profile determined in the region$[0.3 - 0.8]R_{500}$ as defined in the text.} \\label{sample} \\centering \\renewcommand{\\footnoterule}{}  % \\begin{tabular}{llcccccccc} \\hline \\hline Cluster& z & $N_{H}$& $T_{X} (0.75R_{500})$&$T_{X} (R_{500})$ & $Z$ & $R_{500}$ & log($M_{gas}$) & $\\alpha$ ($<0.05R_{500}$) & $\\beta$ ($[0.3-0.8]R_{500}$)\\\\ \\hline RXCJ0003.8$+$0203 & 0.0924 & 3.0 & 3.87$_{-0.10}^{+0.10}$ & 3.64$_{-0.09}^{+0.09}$ & 0.27$_{-0.04}^{+0.04}$ & 876.69 & 13.298$\\pm$0.006 & 0.55$\\pm$0.00 & 0.63$\\pm$0.01 \\\\ RXCJ0006.0$-$3443 & 0.1147 & 1.1 & 5.18$_{-0.20}^{+0.20}$ & 4.60$_{-0.16}^{+0.21}$ & 0.34$_{-0.06}^{+0.06}$ & 1059.31 & 13.642$\\pm$0.010 & 0.48$\\pm$0.00 & 0.50$\\pm$0.01 \\\\ RXCJ0020.7$-$2542 & 0.1410 & 2.1 & 5.55$_{-0.13}^{+0.13}$ & 5.24$_{-0.15}^{+0.15}$ & 0.18$_{-0.04}^{+0.04}$ & 1045.30 & 13.606$\\pm$0.008 & 0.21$\\pm$0.00 & 0.74$\\pm$0.01 \\\\ RXCJ0049.4$-$2931 & 0.1084 & 1.9 & 3.03$_{-0.12}^{+0.12}$ & 2.79$_{-0.11}^{+0.11}$ & 0.26$_{-0.04}^{+0.05}$ & 807.79 & 13.225$\\pm$0.009 & 0.48$\\pm$0.01 & 0.65$\\pm$0.01 \\\\ RXCJ0145.0$-$5300 & 0.1168 & 2.8 & 5.63$_{-0.14}^{+0.14}$ & 5.51$_{-0.16}^{+0.16}$ & 0.30$_{-0.05}^{+0.05}$ & 1089.28 & 13.674$\\pm$0.009 & 0.17$\\pm$0.00 & 0.56$\\pm$0.00 \\\\ RXCJ0211.4$-$4017 & 0.1008 & 1.6 & 2.07$_{-0.05}^{+0.05}$ & 2.02$_{-0.06}^{+0.06}$ & 0.27$_{-0.02}^{+0.02}$ & 685.04 & 12.991$\\pm$0.008 & 0.70$\\pm$0.02 & 0.61$\\pm$0.03 \\\\ RXCJ0225.1$-$2928 & 0.0604 & 1.6 & 2.67$_{-0.13}^{+0.13}$ & 2.61$_{-0.16}^{+0.16}$ & 0.69$_{-0.09}^{+0.11}$ & 693.91 & 12.874$\\pm$0.012 & 0.04$\\pm$0.01 & 0.56$\\pm$0.00 \\\\ RXCJ0345.7$-$4112 & 0.0603 & 1.8 & 2.30$_{-0.06}^{+0.09}$ & 2.15$_{-0.08}^{+0.08}$ & 0.37$_{-0.04}^{+0.05}$ & 688.40 & 12.919$\\pm$0.013 & 1.19$\\pm$0.05 & 0.62$\\pm$0.02 \\\\ RXCJ0547.6$-$3152 & 0.1483 & 2.1 & 6.06$_{-0.14}^{+0.14}$ & 5.68$_{-0.11}^{+0.11}$ & 0.27$_{-0.03}^{+0.03}$ & 1133.74 & 13.768$\\pm$0.005 & 0.26$\\pm$0.01 & 0.62$\\pm$0.01 \\\\ RXCJ0605.4$-$3518 & 0.1392 & 4.5 & 4.91$_{-0.11}^{+0.11}$ & 4.81$_{-0.12}^{+0.12}$ & 0.31$_{-0.04}^{+0.04}$ & 1045.94 & 13.659$\\pm$0.008 & 0.87$\\pm$0.04 & 0.70$\\pm$0.01 \\\\ RXCJ0616.8$-$4748 & 0.1164 & 5.1 & 4.17$_{-0.11}^{+0.11}$ & 4.16$_{-0.12}^{+0.12}$ & 0.31$_{-0.04}^{+0.04}$ & 939.16 & 13.452$\\pm$0.008 & 0.61$\\pm$0.04 & 0.50$\\pm$0.02 \\\\ RXCJ0645.4$-$5413 & 0.1644 & 6.5 & 7.27$_{-0.18}^{+0.18}$ & 6.97$_{-0.19}^{+0.19}$ & 0.22$_{-0.04}^{+0.04}$ & 1279.98 & 13.994$\\pm$0.007 & 0.49$\\pm$0.01 & 0.60$\\pm$0.01 \\\\ RXCJ0821.8$+$0112 & 0.0822 & 4.2 & 2.84$_{-0.10}^{+0.10}$ & 2.44$_{-0.12}^{+0.12}$ & 0.28$_{-0.04}^{+0.05}$ & 755.86 & 13.071$\\pm$0.010 & 0.68$\\pm$0.02 & 0.64$\\pm$0.02 \\\\ RXCJ0958.3$-$1103 & 0.1669 & 5.4 & 6.30$_{-0.44}^{+0.50}$ & 5.85$_{-0.40}^{+0.45}$ & 0.25$_{-0.00}^{+0.00}$ & 1077.39 & 13.648$\\pm$0.016 & 0.64$\\pm$0.02 & 0.81$\\pm$0.01 \\\\ RXCJ1044.5$-$0704 & 0.1342 & 3.6 & 3.57$_{-0.05}^{+0.05}$ & 3.52$_{-0.05}^{+0.05}$ & 0.26$_{-0.02}^{+0.02}$ & 931.85 & 13.518$\\pm$0.008 & 1.11$\\pm$0.06 & 0.69$\\pm$0.01 \\\\ RXCJ1141.4$-$1216 & 0.1195 & 3.2 & 3.54$_{-0.05}^{+0.05}$ & 3.40$_{-0.06}^{+0.06}$ & 0.38$_{-0.03}^{+0.03}$ & 885.24 & 13.385$\\pm$0.012 & 0.98$\\pm$0.05 & 0.62$\\pm$0.01 \\\\ RXCJ1236.7$-$3354 & 0.0796 & 5.5 & 2.73$_{-0.01}^{+0.09}$ & 2.57$_{-0.03}^{+0.11}$ & 0.42$_{-0.04}^{+0.04}$ & 753.50 & 13.078$\\pm$0.009 & 0.41$\\pm$0.01 & 0.61$\\pm$0.01 \\\\ RXCJ1302.8$-$0230 & 0.0847 & 1.7 & 3.44$_{-0.07}^{+0.07}$ & 2.92$_{-0.07}^{+0.09}$ & 0.26$_{-0.03}^{+0.03}$ & 842.12 & 13.247$\\pm$0.013 & 1.04$\\pm$0.06 & 0.50$\\pm$0.02 \\\\ RXCJ1311.4$-$0120 & 0.1832 & 1.8 & 8.44$_{-0.12}^{+0.12}$ & 8.24$_{-0.13}^{+0.13}$ & 0.26$_{-0.02}^{+0.02}$ & 1319.18 & 14.019$\\pm$0.004 & 0.56$\\pm$0.01 & 0.71$\\pm$0.01 \\\\ RXCJ1516.3$+$0005 & 0.1181 & 4.7 & 4.48$_{-0.07}^{+0.07}$ & 4.18$_{-0.08}^{+0.08}$ & 0.25$_{-0.03}^{+0.03}$ & 989.86 & 13.548$\\pm$0.008 & 0.44$\\pm$0.00 & 0.65$\\pm$0.01 \\\\ RXCJ1516.5$-$0056 & 0.1198 & 5.4 & 3.74$_{-0.09}^{+0.10}$ & 3.40$_{-0.08}^{+0.08}$ & 0.30$_{-0.03}^{+0.03}$ & 927.02 & 13.472$\\pm$0.009 & 0.51$\\pm$0.01 & 0.41$\\pm$0.01 \\\\ RXCJ2014.8$-$2430 & 0.1538 & 13.1 & 5.73$_{-0.10}^{+0.10}$ & 5.63$_{-0.11}^{+0.11}$ & 0.27$_{-0.03}^{+0.03}$ & 1155.29 & 13.843$\\pm$0.013 & 0.88$\\pm$0.15 & 0.64$\\pm$0.01 \\\\ RXCJ2023.0$-$2056 & 0.0564 & 5.3 & 2.72$_{-0.09}^{+0.09}$ & 2.46$_{-0.12}^{+0.12}$ & 0.20$_{-0.04}^{+0.04}$ & 739.51 & 13.014$\\pm$0.010 & 0.46$\\pm$0.00 & 0.58$\\pm$0.01 \\\\ RXCJ2048.1$-$1750 & 0.1475 & 4.7 & 5.01$_{-0.11}^{+0.11}$ & 4.59$_{-0.08}^{+0.08}$ & 0.22$_{-0.03}^{+0.03}$ & 1077.96 & 13.730$\\pm$0.008 & 0.01$\\pm$0.00 & 0.51$\\pm$0.01 \\\\ RXCJ2129.8$-$5048 & 0.0796 & 2.2 & 3.88$_{-0.14}^{+0.14}$ & 3.64$_{-0.12}^{+0.16}$ & 0.46$_{-0.06}^{+0.06}$ & 900.60 & 13.350$\\pm$0.008 & 0.37$\\pm$0.00 & 0.46$\\pm$0.01 \\\\ RXCJ2149.1$-$3041 & 0.1184 & 2.3 & 3.50$_{-0.07}^{+0.07}$ & 3.40$_{-0.08}^{+0.08}$ & 0.26$_{-0.03}^{+0.03}$ & 886.56 & 13.393$\\pm$0.010 & 0.75$\\pm$0.06 & 0.56$\\pm$0.01 \\\\ RXCJ2157.4$-$0747 & 0.0579 & 3.6 & 2.76$_{-0.07}^{+0.07}$ & 2.30$_{-0.06}^{+0.10}$ & 0.28$_{-0.03}^{+0.03}$ & 751.45 & 13.047$\\pm$0.008 & 0.42$\\pm$0.00 & 0.37$\\pm$0.01 \\\\ RXCJ2217.7$-$3543 & 0.1486 & 1.1 & 4.65$_{-0.08}^{+0.10}$ & 4.45$_{-0.09}^{+0.09}$ & 0.21$_{-0.03}^{+0.03}$ & 1022.61 & 13.638$\\pm$0.007 & 0.45$\\pm$0.00 & 0.60$\\pm$0.01 \\\\ RXCJ2218.6$-$3853 & 0.1411 & 1.4 & 6.16$_{-0.19}^{+0.19}$ & 5.88$_{-0.15}^{+0.20}$ & 0.34$_{-0.05}^{+0.05}$ & 1130.13 & 13.747$\\pm$0.006 & 0.31$\\pm$0.00 & 0.66$\\pm$0.01 \\\\ RXCJ2234.5$-$3744 & 0.1510 & 1.2 & 7.30$_{-0.12}^{+0.12}$ & 6.95$_{-0.14}^{+0.14}$ & 0.23$_{-0.03}^{+0.03}$ & 1283.21 & 13.984$\\pm$0.007 & 0.13$\\pm$0.01 & 0.71$\\pm$0.01 \\\\ RXCJ2319.6$-$7313 & 0.0984 & 2.9 & 2.52$_{-0.07}^{+0.07}$ & 2.48$_{-0.08}^{+0.08}$ & 0.31$_{-0.04}^{+0.04}$ & 788.73 & 13.238$\\pm$0.010 & 0.81$\\pm$0.07 & 0.51$\\pm$0.02 \\\\ \\hline \\hline \\end{tabular} \\end{minipage} \\end{table*}  In this paper we examine the gas distributions in 31 clusters drawn from the Representative {\\it XMM-Newton} Cluster Structure Survey (REXCESS). Full details of the sample selection function can be found in B\\\"ohringer et al. (2007). To achieve a statistically well-defined sample which fully samples the X-ray luminosity function and is unbiased with respect to dynamical state, a sample of 33 clusters was constructed from the REFLEX catalogue (B\\\"ohringer et al. 2004) based on the following criteria: redshift $z < 0.2$; close to homogeneous coverage of luminosity space; $kT > 2.0$ keV; detectable with {\\it XMM-Newton} to a radius of $\\sim R_{500}$; and distances selected to optimise the extent of the cluster within the {\\it XMM-Newton} field-of-view. In addition, a firm detection of more than 30 photons in the original RASS detection, and a low column density towards the source were also required. Observations of the full sample of 31 clusters plus 2 archive observations have now been completed. The analysis presented here consists of the 31 clusters for which it is reasonable to carry out a 1-dimensional analysis. We exclude the clusters RXCJ2152$-$1942 (Abell 2384B) and RXCJ0956$-$1004 (Abell 901/902) which are highly irregular/multiple cluster systems and hence cannot be analysed in this way.  Throughout this paper we adopt a $\\Lambda$CDM cosmology with $H_{0} = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{M} = 0.3$ and $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "We have presented the first detailed study of the structural properties of cluster gas in a large, representative sample of nearby galaxy clusters. As the sample was selected by X-ray luminosity, it includes clusters of all dynamical states, allowing us to investigate the effect that the inclusion of less regular systems has on results obtained with previous studies of regular clusters. We found the following results: \\begin{itemize} \\item The 1-D gas density profiles scale self-similarly, with a scatter ranging from $\\sim 100\\%$ in the inner regions to $\\sim 20 \\%$ at a radius of $0.3R_{500}$ when expected evolution with redshift is taken into account. \\item Gas density normalisation at $0.3 R_{500}$ is correlated with global cluster temperature, with a scaling of $n_{e} (0.3R_{500}) \\propto T_{X}^{0.5}$, consistent with the expectation of modified entropy-temperature scaling models. Using this scaling reduced the scatter in the gas density profiles at $0.3 R_{500}$ to $\\sim 15 \\%$; however, the scatter at larger radii is slightly increased, which indicates that the entropy excess is much less significant beyond $\\sim 0.5 R_{500}$. \\item The gas density slope continues to increase with radius in the region $0.5R_{500} - R_{500}$, as found by others, which is of importance for cluster mass estimates. \\item The outer gas density slope is correlated with X-ray temperature, primarily due to a lack of hot clusters with flat gas distributions. The flatter slope in lower temperature systems, combined with the entropy excess at intermediate radii and the steep slope of the $M_{gas}$--$T$ relation suggests that gas has been displaced from the centres of the lower temperature systems to larger radii. \\item Based on a characterisation of the cluster dynamical state using power ratios and centre shifts, there is evidence of a correlation between cluster dynamical state and outer gas density slope, and no correlation with inner gas density slope. There is also evidence for a correlation between dynamical state and the central gas properties (gas density normalisation at $0.007 R_{500}$ and cooling time at $0.03 R_{500}$). \\item There is no evidence for bimodality in the central gas density, gas density slope or cooling times for this sample, suggesting that X-ray clusters form a single population with a continuous distribution of central gas properties. \\item The gas-mass temperature relation for the REXCESS sample is in good agreement with predictions of self-similar scaling modified by the presence of an entropy excess, and with previous work on samples of regular clusters; however, the intrinsic scatter is a factor of $\\sim 2.5$ times higher than for the relaxed cluster population. \\item The scaling properties of the gas density profiles appear to be in broad agreement with those of simulated cluster samples (Borgani et al. 2004) at intermediate radius, with a slightly flatter slope in the outer regions. However, in contrast to the observational data there is no correlation between gas density normalisation and temperature in the simulated sample, or between the outer slope and temperature. These discrepancies suggest that the non-gravitational heating of the intra-cluster medium may be too centrally concentrated in these models.  \\end{itemize}"
    },
    "0801/0801.2570_arXiv.txt": {
        "abstract": "We use Chandra observations of 13 nearby groups of galaxies to investigate the hot gas content of their member galaxies.  We find that a large fraction of near-IR bright, early-type galaxies in groups have extended X-ray emission, indicating that they retain significant hot gas halos even in these dense environments.  In particular, we detect hot gas halos in $\\sim$80\\% of $L_K > L_{\\ast}$ galaxies.  We do not find a significant difference in the $L_K-L_X$ relation for detected group and cluster early-type galaxies.  However, we detect X-ray emission from a significantly higher fraction of galaxies brighter than $L_{\\ast}$ in groups compared to clusters, indicating that a larger fraction of galaxies in clusters experience significant stripping of their hot gas.  In addition, group and cluster galaxies appear to be X-ray faint compared to field galaxies, though a Chandra based field sample is needed to confirm this result.  The near-IR bright late-types galaxies in clusters and groups appear to follow the $L_K-L_X$ relation for early-type galaxies, while near-IR fainter late-type galaxies are significantly more X-ray luminous than this relation likely due to star formation.  Finally, we find individual examples of ongoing gas stripping of group galaxies.  One galaxy shows a 40-50 kpc X-ray tail, and two merging galaxy systems show tidal bridges/tails of X-ray emission.  Therefore, stripping of hot galactic gas through both ram pressure and tidal forces does occur in groups and clusters, but the frequency or efficiency of such events must be moderate enough to allow hot gas halos in a large fraction of bright galaxies to survive even in group and cluster cores. ",
        "introduction": "X-ray observations of the hot gas halos ($\\sim 10^7$ K) surrounding galaxies offer a new window on galaxy evolution and the effects of environment.  In dense environments, like clusters of galaxies, the galaxy populations are known to be quite different from galaxies in the field: they typically have early-type morphologies and little active star formation (e.g. Dressler 1980; Balogh et al. 1997).  Several mechanisms have been proposed to transform galaxies in dense environments which can also act to remove their hot gas halos.  Broadly, these include mergers/tidal interactions with other galaxies (e.g. Spitzer \\& Baade 1951; Richstone 1976; White 1978) or the cluster potential (e.g. Merritt 1984; Miller 1986; Byrd \\& Valtonen 1990), ram-pressure or viscous stripping through interaction with the intracluster medium (ICM; Gunn \\& Gott 1972; Nulsen 1982), or evaporation by the hot ICM (Cowie \\& Songaila 1977).  Typically, ram-pressure stripping is thought to be more effective in high density cluster cores (e.g. Gunn \\& Gott 1972), while galaxy-galaxy mergers are thought to be more common in less massive galaxy groups where the relative galaxy velocities are lower (e.g. Barnes 1985; Aarseth \\& Fall 1980; Merritt 1984).  However, recent simulations and observations show that ram pressure can be effective at stripping the hot gas, though not the cold gas, from some galaxies even in low-mass groups (Kawata \\& Mulchaey 2007; Rasmussen, Ponman, \\& Mulchaey 2006).  In spiral galaxies, the hot gas halo provides a reservoir of gas which can cool and replenish the cold gas supply to fuel star formation.  If this gas is stripped from galaxies in clusters and groups, star formation will be quenched after a few gigayears when the galaxy exhausts its gas supply.  Referred to as ``strangulation'' or ``starvation'', this mechanism has been proposed to explain the observed lack of star-forming galaxies and the transformation of spirals to S0 galaxies (as the disk fades) in clusters (e.g. Larson et al. 1980; Bekki et al. 2002).  In elliptical galaxies, hot gas halos are formed and enriched through stellar mass lost from evolved stars and possibly through gas accreted from the local environment (e.g. Mathews 1990; Mathews \\& Brighenti 2003).  Observations with \\textit{Einstein} showed that hot gas halos surrounding ellipticals are common (e.g. Forman, Jones, \\& Tucker 1985; Kim, Fabbiano, \\& Trinchieri 1992).  A correlation was found between the X-ray luminosity of early-type galaxies and optical or near-IR luminosity, but with surprisingly large scatter (e.g. Forman et al. 1985; Trinchieri \\& Fabbiano 1985; Canizares, Fabbiano, \\& Trinchieri 1987; O'Sullivan, Forbes, \\& Ponman 2001; Ellis \\& O'Sullivan 2006).  In addition to several intrinsic galaxy properties (dark matter profile, rotation, supernovae rate), environment and ram-pressure stripping in dense environments are among the proposed factors to explain this scatter (see review by Mathews \\& Brighenti 2003).  The stripping of gas from cluster and group galaxies may also be an important contribution to the enrichment of the ICM (Byrd \\& Valtonen 1990; White 1991; Schindler et al. 2005; Domainko et al.2006; Kapferer et al.2007).  How effective are stripping mechanisms at removing the hot gas from typical galaxies in clusters?  In lower density groups?  % Based on the Millenium Simulations, Br\\\"{u}ggen \\& De Lucia (2007) predict that nearly all galaxies in the cores of clusters experienced strong ram pressure in their histories.  In addition, observational evidence for hot gas stripping has been found for $\\sim 10$ galaxies in clusters or groups in the form of long X-ray tails trailing the galaxy (Forman et al. 1979; Rangarajan et al. 1995; Irwin \\& Sarazin 1996; Biller et al. 2004; Trinchieri et al. 1997; Kim et al. 2007; Sakelliou et al. 2005; Wang et al. 2004; Scharf et al. 2005; Machacek et al. 2005a, 2005a, 2006, 2007; Rasmussen et al. 2006; Sun \\& Vikhlinin 2005; Sun et al. 2006).  However, a recent \\textit{Chandra} archival survey of galaxies in 25 clusters by Sun et al. (2007; S07) found that most bright cluster galaxies ($>$ 60\\% of galaxies above 2$L_{\\ast}$) host hot halos.  Supporting these observations, recent hydrodynamic simulations by McCarthy et al. 2007 show that for typical galaxy structures and orbits, galaxies falling in to a massive group/poor cluster maintain $\\sim$30\\% of their hot halo gas after 10 Gyr, and the amount of stripping depends signficantly on the orbital history.  The superb resolution of \\textit{Chandra} is invaluable in the study of galaxy halos, allowing a separation between extended X-ray halos and point source AGN.  Previous studies which did not separate these two components have given mixed results regarding the relationship of the X-ray emission from galaxies to environment.  For example, some \\textit{ROSAT} studies did not find a significant trend in the $L_K-L_X$ or $L_B-L_X$ relations with environment (O'Sullivan et al. 2001; Helsdon et al. 2001; Ellis \\& O'Sullivan 2006), while other authors found that galaxies in regions of high local galaxy density have lower $L_X/L_B$ (White \\& Sarazin 1991; Henriksen \\& Cousineau 1999), and Brown \\& Bregman (2000), whose sample included several brightest group galaxies (BGG), found the opposite trend for galaxies in higher density environments to have higher $L_X/L_B$.  Recent \\textit{XMM-Newton} and \\textit{Chandra} observations of the X-ray luminosity function of Coma cluster galaxies indicate that these are X-ray underluminous compared to field galaxies (Finoguenov et al. 2004; Hornschemeier et al. 2006).  In this paper, we extend the \\textit{Chandra} study of hot galaxy halos to lower density environments through a systematic study of archival observations of 13 galaxy groups in order to re-examine the effects of environment on hot gas halos.  % Our sample is presented in \\S2, the data reduction in \\S3, and our results in \\S4. These include the $L_K-L_X$ relation and halo temperatures of early-type galaxies (\\S4.1), the detection rate (\\S4.2), detections of late-type galaxies (\\S4.3), and examples of X-ray tails and mergers (\\S4.4).  Throughout the paper, we assume $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "Using archival \\textit{Chandra} observations of 13 groups, we search for hot gas halos surrounding group member galaxies.  Here we only consider satellite galaxies or BGGs which do not lie at the center of a luminous group-size X-ray halo.  We find that a large fraction of bright, early-type galaxies in groups have extended X-ray emission, indicating that they retain significant hot gas halos even in these dense environments.  In particular, we detect hot gas halos in $\\sim$80\\% of $L_K > L_{\\ast}$ early-type galaxies.  We compare the X-ray halos of group galaxies to cluster galaxies, a number of which are also found to host significant hot gas halos (S07).  Within the errors, we do not find a significant difference in the $L_K-L_X$ relation for detected group and cluster early-type galaxies, indicating that there are not large differences in the size of the hot halos for those galaxies which host significant halos in these environments.  However, we detect a significantly higher fraction of $L_K > L_{\\ast}$ early-type galaxies in groups compared to clusters, indicating that a larger fraction of galaxies in clusters experience significant stripping of their hot gas.  Ram-pressure stripping is, in fact, expected to be more effective in clusters versus groups where the ICM density is higher and the galaxy velocities are larger, while the detected galaxies are likely those who, due to their orbits and accretion times, have experienced only mild ram pressure.  We also find that group galaxies tend to have lower halo temperatures than cluster galaxies at the same near-IR luminosity or stellar velocity dispersion indicating typically less heating of the gas relative to the stars, particularly for the higher velocity dispersion galaxies in groups.  Our sample includes significantly fewer late-type galaxies, but we do detect 4 out of 11.  Combined with the late-type galaxies in the cluster sample of S07, the X-ray emission of late-type galaxies does not correlate as strongly with $K_s$-band luminosity (S07) as it does for early-type galaxies.  Brighter late-types ($L_K > L_{\\ast}$) appear to follow the $L_K-L_X$ relation for early-type galaxies, while the fainter late-type galaxies are significantly more X-ray luminous than this relation.  In particular, clusters, at least, contain a population of $K_s$-band faint, X-ray luminous, star-forming galaxies.  The X-ray luminosity of group and cluster late-type galaxies correlates much more strongly with B-band luminosity, supporting the theroy that the X-ray emission is related to star formation.  Intriguingly, the detected group and cluster late-type galaxies do not appear to have singificantly lower $L_X/L_B$ than the overall local spiral population as observed by \\textit{Einstein}.  The upper limits for the undetected late-types do not appear to be significantly underluminous compared to the detected galaxies, and deeper observations are needed to establish if these galaxies are X-ray deficient.  The fact that many galaxies in groups and clusters retain significant hot gas halos is in qualitative agreement with recent hydrodynamic simulations which show that for typical galaxy structures and orbits, galaxies falling in to groups and clusters maintain $\\sim$30\\% of their hot halo gas after 10 Gyr (McCarthy et al. 2007), with the amount of stripping depending signficantly on the orbital history.  Our result that a significantly higher fraction of galaxies are detected in groups versus clusters is also in qualitative agreement with the predictions of recent semi-analytical simulations based on the Millenium Simulations, which find that a higher fraction of galaxies in massive clusters versus poor clusters experience strong ram pressure at some point in their history (Br\\\"{u}ggen \\& De Lucia 2007).  However, these simulations also predict that nearly all galaxies in cluster and group cores have experienced strong ram pressure, seemingly at odds with the large number of detected X-ray halos.  These models also predict that there should be a strong radial trend in the stripping experienced by group and cluster galaxies, but we find no significant trend of the $L_K-L_X$ relation or detections versus non-detections with radius.  Here we face limitations of both the sample size and the radial coverage of the \\textit{Chandra} data, which may limit our ability to detect such a trend.  Our detected group galaxies are limited in radius to $r \\leq 0.3 r_{500}$.  Comparing to field galaxies in the \\textit{ROSAT} observed sample of Ellis \\& O'Sullivan (2006), we do find that field early-type galaxies have a significantly higher X-ray luminosity at fixed $L_K$ than group and cluster galaxies.  We find a significant offset in the $L_K-L_X$ relation between group and field galaxies even when we assume that all of the X-ray emission from detected sources originates from a thermal component, as was done by Ellis \\& O'Sullivan (2006).  This result may confirm theoretical expectations that the hot halos of galaxies in dense environments are significantly stripped compared to field galaxies, but analysis of a field sample with \\textit{Chandra} is needed to confirm that the \\textit{ROSAT} luminosities are not significantly enhanced by point source or background contamination.  It is, however, clear that ram pressure and tidal stripping of hot galactic gas does occur in at least some galaxies in dense environments.  We find that $\\sim$ 5\\% of galaxies detected with extended halos and $\\sim$ 2\\% of all $\\geq L_{\\ast}$ galaxies in groups and clusters show long X-ray tails indicative of ram-pressure/viscous stripping.  We additionally find two systems of merging galaxies in our groups that show indications of tidal interaction in their X-ray gas.  These cases add to the $\\sim$ 10 other cluster/group galaxies in the literature found to show X-ray signatures of ram pressure and/or tidal stripping (Forman et al. 1979; Rangarajan et al. 1995; Irwin \\& Sarazin 1996; Biller et al. 2004; Trinchieri et al. 1997; Kim et al. 2007; Sakelliou et al. 2005; Wang et al. 2004; Scharf et al. 2005; Machacek et al. 2005a, 2005a, 2006, 2007; Rasmussen et al. 2006).  These systems include both early and late-type galaxies.  We also find galaxies in both groups and clusters which are significantly underluminous (both detections and non-detections) compared to the $L_K-L_X$ relation, with a higher fraction of these galaxies in clusters.  A picture is therefore emerging that stripping of hot galactic gas through both ram pressure and tidal forces does occur in groups and clusters, but the frequency or efficiency of such events must be moderate enough to allow the hot gas halos in a large fraction of bright galaxies to survive even in group and cluster cores.  Our results also indicate that a larger fraction of cluster galaxies have experienced significant gas stripping compared to galaxies in groups.  The question remains whether these processes contribute significantly to galaxy evolution or ICM enrichment.  This question can be addressed with future simulations using these observations as a baseline for the hot gas content of galaxies in groups and clusters."
    },
    "0801/0801.2746_arXiv.txt": {
        "abstract": "{} {Coronal mass ejections or CMEs are large dynamical solar-corona events. The mass balance and kinematics of a fast limb CME, including its prominence progenitor and the associated flare, will be compared with computed magnetic structures to look for their origin and effect.} {Multi-wavelength ground-based and spaceborne observations are used to study a fast W-limb CME event of December 2, 2003, taking into account both on and off disk observations. Its erupting prominence is measured at high cadence with the Pic du Midi full H$\\alpha$ line-flux imaging coronagraph. EUV images from SOHO/EIT and CORONAS-F/SPIRIT space instruments are processed including difference imaging. SOHO/LASCO images are used to study the mass excess and motions. Computed coronal structures from extrapolated surface magnetic fields are compared to observations.} {A fast bright expanding coronal loop is identified in the region recorded slightly later by GOES as a C7.2 flare, followed by a brightening and an acceleration phase of the erupting material with both cool and hot components. The total coronal radiative flux dropped by $\\sim$ 7\\% in the 19.5 nm channel and by 4\\% in the 17.5 nm channel, revealing a large dimming effect at and above the limb over a 2 hour interval. The typical 3-part structure observed 1 hour later by the Lasco C2 and C3 coronagraphs shows a core shaped similarly to the eruptive filament/prominence. The total measured mass of the escaping CME ($\\sim 1.5\\cdot10^{16}$~g from C2 LASCO observations) definitely exceeds the estimated mass of the escaping cool prominence material although assumptions made to analyze the H$\\alpha$ erupting prominence, as well as  the corresponding EUV darkening of the filament observed several days before, made this evaluation uncertain by a factor of 2. This mass budget suggests that the event is not confined to the eruption region alone. From the current free extrapolation we discuss the shape of the magnetic neutral surface and a possible scenario leading to an instability, including the small scale dynamics inside and around the filament.    } {} ",
        "introduction": "Coronal Mass Ejections (CMEs) were originally defined as large coronal dynamical phenomena (coronal transients) propagating outward in the field of view of white-light (W-L) externally occulted space-borne coronagraphs and often showing three parts (see Wagner \\cite{Wagner}). Their primary property is the mass transport of a presumably coronal mass originally situated behind the occulting device, toward the external part, without specifying the origin of this mass revealed by the W-L event due to the scattering on free electrons. Modelers had to address the fundamental question: from where does this large amount of coronal plasma originate? Although many statistical studies of CMEs were published in recent decades (among the latest, see Vourlidas et al. \\cite{Vourlidas}; Cremades \\& Bothmer \\cite{ Cremades}; Gopalswamy \\cite{Gopalswamy06}) this question has not been systematically considered and no clear answer exists. One possibility is that the mass of the CME comes from a related erupting prominence (EP) and its immediate surroundings. This assumption was seriously considered recently and prompted another question: what is the mass of the EP (see Gilbert et al. \\cite{Gilbert05, Gilbert06})? However, the general consensus is that only the core of the 3-part CME corresponds to the prominence.  \\begin{figure} \\centering \\includegraphics[width=9.9cm]{Fig1.eps} \\caption { The BBSO H$\\alpha$ image (a) showing the zoom box  corresponding to  the MDI magnetic map (b) of the AR 10508 region with superimposed contours of filaments (red) from the BBSO H$\\alpha$ image; (c) EIT image of the box region at 195 \\AA. All images are taken on November 26, 2003.  } \\label{Fig1}% \\end{figure}  A distinction has been made in considering the presumably more important fast CMEs, usually related to a flare and more likely to be geoeffective and with space weather implications (see e.g. Bao et al. \\cite{Bao}). The EP connection is then sometimes omitted in favor of a discussion of the origin of solar energetic particles that contribute only a little to the mass budget.  CMEs are often associated with erupting filaments or prominences even when considered far from the Sun (see Bothmer \\& Schwenn \\cite{Bothmer}). In the last few years many investigations were undertaken to study the relationship between prominence activity and CMEs in different wavelengths. Based on H$\\alpha$ observations at the Mauna Loa Solar Observatory, Gilbert et al. (\\cite{Gilbert00}) proposed a classification of active and eruptive prominences according to their kinematical properties (radial height, velocity and acceleration). A separation between escaping prominence material lifting away from the Sun and material returning toward the solar surface would occur in the height range from 1.20~$R_\\odot$ to 1.35~$R_\\odot$ which could be related to the formation of an X-type neutral line in this region.  Gopalswamy et al. (\\cite{Gopalswamy03}) studied the temporal and spatial relations between the signatures of prominence eruptions detected in microwaves with the Nobeyama Radioheliograph and white-light CMEs. It was shown that in most of the studied events the prominence moved predominantly in the radial direction, with CMEs and  EPs starting at nearly the same time.  Jing et al. (\\cite{Jing}) presented a statistical study of filament eruptions observed in H$\\alpha$ at the Big Bear Solar Observatory with other phenomena of solar activity. They found that most eruptions ($>$ 50\\%) were associated with white-light CMEs, active region filament eruptions being more associated with flares than quiescent filament eruptions. Because of their close relationship to CMEs, EPs were regarded as one of the main sources of material in the mass balance of CMEs. Indeed, from the theoretical modeling point of view, the mass of the prominence could be an issue when considering the launch of the CME (Low et al. \\cite{Low}) or, more generally, its associated magnetic structure (see Chen et al. \\cite{Chen}). Further out, there are strong indications that the core in the three-part CME structure is associated with a prominence but this relation is not direct; multi-wavelength observations are needed.  Evaluations of the EP mass are traditionally based on observations of the H$\\alpha$ absorption of the corresponding filament on the disk prior to eruption, where they are seen as dark objects. The destabilization of filaments appears as a sudden disappearance in H$\\alpha$ or {\\it a disparition brusque} (DB) (Tandberg-Hanssen \\cite{Tandberg-Hanssen}; Soru-Escaut \\& Mouradian \\cite{Soru-Escaut}) which can be caused by an increase of the plasma ionization degree (thermal effect) or which can also be a result of Doppler - Fizeau effects (dynamical DB) including a darkening produced by rapid mass motion (Heinzel \\& Rompolt \\cite{Heinzel87}). With TRACE usually a large, high but faint $\\sim$1 MK loop moves outward before anything else happens.   \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{Fig2.eps}} \\caption{Partial frame from an H$\\alpha$ filtergram taken before the eruption from the International H$\\alpha$ Network and, at left, the deduced set of calibrated isophotes used to evaluate the mass of the off-disk prominence. Note the levels of scattered parasitic light outside the solar disc (some typical intensities are shown near the bottom part).  } \\label{Fig2}% \\end{figure}   The approach to estimating the hydrogen density in filaments from H$\\alpha$  absorption was developed by Mein et al. (\\cite{Mein96}) and Heinzel et al. (\\cite{Heinzel99}) who used the non-LTE cloud model with optical thickness, Doppler width and velocity as input parameters. Because the hydrogen ionization degree in the non-LTE model is a free parameter, the total mass of the filament cannot be determined unambiguously. When the prominence is seen in emission outside the disk, a more direct evaluation can be attempted using both narrow passband W-L coronagraphic measurements and H$\\alpha$  line-flux emission imaging (see Koutchmy \\& Nikolsky \\cite{Koutchmy81}). To make filtergrams, a narrow passband Lyot  filter is usually used and the data need additional calibrations to take into account the broader line profile.  This problem was partly overcome by combining H$\\alpha$ observations with observations of filaments in coronal EUV lines, where filaments are seen as dark objects on the background of the underlying corona, on the disk or even outside the disk. EUV observations of filaments were started with the Skylab mission and then continued by SOHO (EIT and CDS) and TRACE. Batchelor \\& Schmahl (\\cite{Batchelor}) interpreted the EUV absorption of cool objects as a continuum absorption of the coronal radiation by neutral H, and by neutral and ionized He depending on the wavelength. Kucera et al. (\\cite{Kucera}) and Penn (\\cite{Penn}) used CDS data for prominence diagnostics and obtained a column density of neutral hydrogen \\ion{H}{i} of $\\sim 4\\cdot10^{17} - 10^{18}$~cm$^{-2}$. Mein et al. (\\cite{Mein01}) have shown that a combination of H$\\alpha$, \\ion{Ca}{ii} (8542 \\AA) and TRACE 171 \\AA\\ data may yield a more self-consistent determination of the neutral hydrogen column density, of the electron density and of the electron temperature.   \\begin{figure*} \\centering \\includegraphics[width=\\textwidth]{Fig3.eps} \\caption{Active phase of the prominence eruption seen with HACO in  H$\\alpha$ (a)-(c) and in EIT 195 \\AA\\ (d)-(f). The arrow marks the possible place of emergence of the magnetic loop. In the bottom panels, the time variations of the total flux in H$\\alpha$ measured with HACO (g) and, over the full day time interval, in X-rays 1-8 \\AA\\ measured with GOES-12 (h). } \\label{Fig3}% \\end{figure*}  Heinzel et al. (\\cite{Heinzel01}) and Schmieder et al. (\\cite{Schmieder03, Schmieder04}) noticed that filaments are more extended in EUV spectral lines with wavelengths below the hydrogen Lyman-continuum edge (912 \\AA) than in H$\\alpha$ and they introduced the term ``EUV filament channel\" (EFC) to describe their observations. A spectroscopic model of the EUV absorption of filaments and prominences was suggested by Heinzel et al. (\\cite{Heinzel03a, Heinzel03b}),  Anzer \\& Heinzel (\\cite{Anzer}), and Schwartz et al. (\\cite{Schwartz}). The method takes the absorption of EUV-line radiation by the Lyman continuum self-consistently into account, as well as the volume-blocking effect that is potentially important for coronal lines. A similar technique  for deriving prominence mass from EUV absorption features was developed by Gilbert et al. (\\cite{Gilbert05, Gilbert06}). However, when the EUV filament is seen as a prominence on the off-disk coronal background, often an edge-brightening around the prominence is observed, strongly suggesting that heated coronal material is concentrated at the periphery of the prominence. De Boer et al. (\\cite{Boer}) analyzed the periphery of a prominence using transition region (TR) line emission and describes this enhanced emission with non-thermal broadenings and velocities reaching values up to 45 km/s, which is considerably higher than what is observed using the profiles of cool lines inside the prominence. During the onset of our EP we observed a similar coronal edge brightening.  Here we report on a large W-limb event, which occurred on December 2, 2003, 9-14$^{\\mbox{\\small h}}$ UT, in the vicinity of the old multi-polar region AR 10508 which was well observed earlier but was still on the disk at the time of the CME. We analyse the observations of the prominence eruption taken in H$\\alpha$ full-line images with the HACO coronagraph of the Pic du Midi Observatory (Romeuf  et al. \\cite{Romeuf}), images of the surrounding corona in EUV with SOHO/EIT (Delaboudiniere et al. \\cite{Delaboudiniere}),  and CORONAS-F/SPIRIT telescopes (Zhitnik et al. \\cite{Zhitnik}), and W-L images of the associated CME taken with the LASCO (SOHO) C2 and C3 coronagraphs. We discuss in detail the spatial and the temporal development of the eruption, evaluate the kinematics and especially the mass budget of the CME-related phenomenon and we propose a possible scenario of the event. This approach is quite similar to the early work of Plunkett et al. (\\cite{Plunkett}) where a fast limb CME event with an EP was studied, although no flare was reported (possibly because the region was already beyond the limb). Here we pay special attention to the mass budget, including the measurement of the EUV dimming effect, and we propose a model applicable to this particular event based on a computed magnetic field context to describe the whole event, instead of proposing a universal 2.5 D model. ",
        "conclusions": "We documented a large and fast coronal limb event, which occurred on December 2, 2003, in the interval 9-14$^{\\mbox{\\small h}}$ UT, in the vicinity of an old multi-polar region. Using observations obtained on the ground and in space with two EUV telescopes, an H$\\alpha$ coronagraph and the LASCO white-light coronagraphs, we analysed a spectacular dynamical phenomenon passing through different stages: a quiescent filament initially, then a prominence with EUV brightening at its periphery, its dynamical and thermal activation, a coronal eruption with a small flare and the launch of the eruptive prominence (EP) and finally the formation of a classical three-part fast CME. This CME fits rather well  within the framework of the so-called ``standard model\", (see e.g. Lin \\cite{Lin}). Here we have added several features that we believe could shed some light on the whole process, emphasizing the mass transport processes related to the EP and its coronal surroundings.  From what we saw, the eruption process apparently started with the emergence of a new magnetic loop striking the prominence and destabilizing it. This led to a rapid heating of the gas around the prominence body, which spatially corresponds to a part of the EUV filament channel. The brightness of the body of the prominence in H$\\alpha$ at first increased to a maximum value (at 09:35 UT) possibly due to the growth of turbulence inside which produces a large Doppler brightening effect and, accordingly, magnifies the H$\\alpha$ flux without the need to increase the densities. This brightening even precedes the lifting of the cool prominence, which seems to be an important new feature to be taken into account in future work on CMEs. We note that this observation was made possible by the use of a coronagraph having a broad band filter and not a narrow band H$\\alpha$ filter as is usually done with a solar telescope to reject the parasitic light. After 09:41~UT the brightness decreased in H$\\alpha$ and increased in EUV, due to the enhanced ionization process and possibly a dissipation of the mechanical energy from waves or counter-flow motions. The prominence began to move outward at 09:48 UT, which coincides with the X-ray flare peak and would precede the CME extrapolated onset time.  A temporal analysis of the eruptive prominence motion and of the dimming intensity light curves has shown that both the frontal structure and the core of a CME were initiated simultaneously within 5-10 min around the peak of an X-ray flare.  Evaluations based on photometric data coming from the W-L coronagraphic coverage of the CME and from the photometric data provided by the H$\\alpha$  line-flux images of the prominence before its launch show that the total mass of the CME almost certainly surpasses the mass of the ejected prominence material seen in H$\\alpha$ before its eruption: \\begin{itemize} \\item   the mass of the CME core is 2-2.5 times greater than the mass of the quiescent prominence, including the mass of an EUV channel (although taking into account possible turbulence this difference may vanish). This clearly shows that the mass of a fast CME is mainly due to the surrounding low corona swept out during the eruption, as confirmed by the large dimming effect we measured. This part of the CME is not treated in this paper but new ways already exist to consider it, e.g.  Delann\\'{e}e et al. (\\cite{Del}). \\item   the dimming effect is observed over a large part of the corona, on the disk and also outside the disk when the line of sight integration reaches a 2-fold increased value right above the limb with a corresponding increase of the dimming effect. The dimming is in phase with the onset of the CME. Approximately 5\\% or more of the corona is depleted, neglecting the depletion of the mass of the corona at very low (TR) and/or very high coronal temperatures. The corona seen in the EIT 195 \\AA\\ channel should indeed correspond to the bulk of the mass because it corresponds to the most probable temperature of the corona, so that our estimate should not be far from the truth. \\end{itemize}  We also found that the eruption led to a spatial deflection and a brightening of the neighboring South-West streamer, which may give rise to an overestimate of the CME mass by  $\\sim~20$\\%. Further out toward both polar regions the deflections of structures such as  extended plumes are clearly visible as a back and forth quasi-coherent transverse motion. This phenomenon would need special study which is beyond the scope of this paper.  Finally our results suggest that greater attention should be paid to regions around the EP, bringing added justification to the theoretical studies trying to describe the behavior of the coronal magnetic field leading to CMEs. However only the field inside prominences can be easily measured with cool lines and the Hanle effect; an improved precision H$\\alpha$ coronagraph could routinely be used to perform such measurements at high speed."
    },
    "0801/0801.0928_arXiv.txt": {
        "abstract": "We compare stellar models produced by different stellar evolution codes for the {\\corot}/{\\esta} project, comparing their global quantities, their physical structure, and their oscillation properties. We discuss the differences between models and identify the underlying reasons for these differences. The stellar models are representative of  potential {\\corot} targets. Overall we find very good agreement between the five different codes, but with some significant deviations. We find noticeable discrepancies (though still at the per cent level) that result from the handling of the equation of state, of the opacities and of the convective boundaries. The results of our work will be helpful in interpreting future asteroseismology results from {\\corot}. ",
        "introduction": "\\label{sec:intro}  The goals of \\esta-{\\small TASK}s~1 and 3 are to test the numerical tools used in stellar modelling, with the objective to be ready to interpret safely the asteroseismic data that will come from the {\\corot} mission. This consists in quantifying the effects of different numerical implementations of the stellar evolution equations and related input physics on the internal structure, evolution and seismic properties of stellar models. As a result, we aim at improving the stellar evolution codes to get a good agreement between models built with different codes and same input physics. For that purpose, several study cases have been defined that cover a large range of stellar masses and evolutionary stages and stellar models have been calculated without (in {\\task}~1) or with (in {\\task}~3) microscopic diffusion \\citep[see][]{2006corm.conf..363M,yl1-apss}.  In this paper, we present the results of the detailed comparisons of the internal structures and seismic properties of {\\small TASK}s~1 and 3 target models. The comparisons of the global parameters and evolutionary tracks are discussed in \\cite{mm2-apss}. In order to ensure that the differences found are mainly determined by the way each code calculates the evolution and the structure of the model and not by significant differences in the input physics, we decided to use and compare models whose global parameters (age, luminosity, and radius) are very similar. Therefore, we selected models computed by five codes among the ten  participating in \\esta: {\\astec} \\citep{jcd1-apss}, {\\cesam} \\citep{pm-apss}, {\\cles} \\citep{rs1-apss}, {\\garstec} \\citep{aw-apss}, and {\\starox} \\citep{ir1-apss}.  In Section \\ref{sec:tasks} we recall the specifications of  {\\small TASK}s~1 and 3 and present the five codes used in the present paper. We then present the comparisons for {\\task}~1 in Sect.~\\ref{sec:task1} and for {\\task}~3 in Sect.~\\ref{sec:task3}. For each case in each task, we have computed the relative differences of the physical quantities between pairs of models. We display the variation of the differences between the more relevant quantities inside the star and we provide the average and extreme values of the variations. We then compare the location of the boundaries of the convective regions as well as their evolution with time. For models including microscopic diffusion we examine how helium is depleted at the surface as a function of time. Finally, we analyse the effect of internal structure differences on seismic properties of the model. ",
        "conclusions": "We have presented detailed comparisons of the internal structures and seismic properties of stellar models in a range of stellar parameters -- mass, chemical composition and evolutionary stage -- covering those of the {\\corot} targets.  The models were calculated by 5 codes ({\\astec}, {\\cesam}, {\\cles}, {\\starox}, {\\garstec}) which have followed rather closely the specifications for the stellar models (input physics, physical and astronomical constants) that were defined by the {\\esta} group, although some differences remain, sometimes not fully identified. The oscillation frequencies were calculated by the {\\small LOSC} code (see Sect.~\\ref{sec:tools}).  \\begin{figure*}[htbp!] \\centering \\resizebox*{\\hsize}{!}{\\hspace*{-0.2cm}\\includegraphics*{Ys1.0.eps}\\hspace*{1.5cm}\\includegraphics*{Ys1.2.eps}\\hspace*{1.8cm}\\includegraphics*{Ys1.3.eps}} \\caption{{\\task}~3: Helium content $Y_{\\rm s}$ in the convection envelope for the different codes and cases considered (A and B are for middle and end of the {\\MS} respectively while C is for {\\SGB}, see Table~\\ref{tab:task3}) for $1.0${\\msol} (left),  $1.2${\\msol} (centre) and  $1.3${\\msol} (right).} \\label{fig:task3-ys}       % \\end{figure*}  \\begin{figure*}[htbp!] \\centering \\resizebox*{\\hsize}{!}{\\hspace*{-0.2cm}\\includegraphics*{dnu_3.1.A.eps}\\hspace*{1.5cm}\\includegraphics*{dnu_3.1.B.eps}\\hspace*{1.8cm}\\includegraphics*{dnu_3.1.C.eps}} \\resizebox*{\\hsize}{!}{\\hspace*{-0.2cm}\\includegraphics*{dnu_3.2.A.eps}\\hspace*{1.5cm}\\includegraphics*{dnu_3.2.B.eps}\\hspace*{1.8cm}\\includegraphics*{dnu_3.2.C.eps}} \\resizebox*{\\hsize}{!}{\\hspace*{-0.2cm}\\includegraphics*{dnu_3.3.A.eps}\\hspace*{1.5cm}\\includegraphics*{dnu_3.3.B.eps}\\hspace*{1.8cm}\\includegraphics*{dnu_3.3.C.eps}}  \\caption{{\\task}~3: p-mode frequency differences between models produced by different codes, for Case~3.1 (top row), 3.2 (centre) and 3.3 (bottom). {\\cesam} model is taken as reference, and the frequencies have been scaled to remove the effect of different stellar radii. For each code, we plot two curves corresponding to modes with degrees $\\ell=0$ and $\\ell=1$.} \\label{fig:task3-freq}       % \\end{figure*}   In a first step, we have examined {\\esta}-{\\task}~1 models, calculated for masses in the range $0.9-5$\\msol, with different chemical compositions and evolutionary stages from {\\PMS} to {\\SGB}. In all these models microscopic diffusion of chemical elements has not been included while one model accounts for overshooting of convective cores. In a second step, we have considered the {\\esta}-{\\task}~3 models, in the mass range $1.0-1.3$\\msol, solar composition, and evolutionary stages from the middle of the {\\MS} to the {\\SGB}. In all these models, microscopic diffusion has been taken into account.  For both tasks we have discussed the maximum and average differences in the physical quantities from centre to surface (hydrogen abundance $X$, pressure $P$, density $\\rho$, luminosity $L_r$, opacity $\\kappa$, adiabatic exponent $\\Gamma_1$ and gradient $\\nabla_{\\rm ad}$, specific heat at constant pressure $C_p $ and sound speed $c$). % We have found that the average differences are in general small. Differences in $P$, $T$, $L_r$, $r$, $\\rho$ and $\\kappa$ are on the percent level while differences in the thermodynamical quantities are often well below $1$\\%. Concerning the maximal differences, we have found that they are mostly located in the outer layers and in the zones close to the frontiers of the convective zones. As expected, differences generally increase as the evolution proceeds. They are larger in models with convective cores, in particular in models where the convective core increases during a large part of the {\\MS} before receding. They are also higher in models including microscopic diffusion or overshooting of the convective core.  \\begin{table*}[htbp!] \\caption{{\\task}~3: values of the fractional radius and mass at the base of the convective envelope ($r_{\\rm cz}/R$, $m_{\\rm cz}/M$ ) and border of the convective core ($r_{\\rm cc}/R$, $m_{\\rm cc}/M$) and hydrogen abundance in the convection zones} \\centering \\label{tab:task3-zc} \\begin{tabular}[h]{llcccccc} \\hline\\noalign{\\smallskip} {\\bf Code} & {\\bf Case} & \\boldmath$r_{\\rm cz}/R$ & \\boldmath$m_{\\rm cz}/M$ & \\boldmath$X_{\\rm cz}$& \\boldmath$r_{\\rm cc}/R$ & \\boldmath$m_{\\rm cc}/M$ & \\boldmath$X_{\\rm cc}$\\\\[3pt] \\tableheadseprule\\noalign{\\smallskip} \\astec   & 3.1A  &   0.7348 &   0.9835 &   0.7436 &--&--&--\\\\ \\cesamB  & 3.1A  &   0.7326 &   0.9831 &   0.7446&--&--&-- \\\\ \\cesamMP & 3.1A  &   0.7321 &   0.9830 &   0.7449 &--&--&--\\\\ \\cles    & 3.1A  &   0.7315 &   0.9829 &   0.7470&--&--&-- \\\\ \\garstec & 3.1A  &   0.7357 &   0.9835 &   0.7475&--&--&-- \\\\ \\hline\\noalign{\\smallskip} \\astec   & 3.1B  &   0.7202 &   0.9837 &   0.7630&--&--&-- \\\\ \\cesamB  & 3.1B  &   0.7153 &   0.9832 &   0.7657 &--&--&--\\\\ \\cesamMP & 3.1B  &   0.7143 &   0.9830 &   0.7656&--&--&-- \\\\ \\cles    & 3.1B  &   0.7151 &   0.9831 &   0.7695&--&--&-- \\\\ \\garstec & 3.1B  &   0.7159 &   0.9832 &   0.7705 &--&--&--\\\\ \\hline\\noalign{\\smallskip} \\astec   & 3.1C  &   0.6676 &   0.9720 &   0.7627&--&--&-- \\\\ \\cesamB  & 3.1C  &   0.6644 &   0.9716 &   0.7657 &--&--&--\\\\ \\cesamMP & 3.1C  &   0.6631 &   0.9713 &   0.7661 &--&--&--\\\\ \\cles    & 3.1C  &   0.6643 &   0.9715 &   0.7684 &--&--&--\\\\ \\garstec & 3.1C  &   0.6643 &   0.9712 &   0.7691&--&--&-- \\\\ \\hline\\noalign{\\smallskip} \\astec   & 3.2A  &   0.8489 &   0.9990 &   0.7754 &  0.0456 &   0.0170 &   0.3500 \\\\ \\cesamB  & 3.2A  &   0.8451 &   0.9999 &   0.7835  &   0.0451 &   0.0169 &   0.3499 \\\\ \\cesamMP & 3.2A  &   0.8443 &   0.9990 &   0.7817& 0.0451 &   0.0170 &   0.3500 \\\\ \\cles    & 3.2A  &   0.8444 &   0.9990 &   0.7888     &   0.0447 &   0.0165 &   0.3504 \\\\ \\garstec & 3.2A  &   0.8450 &   0.9990 &   0.7883   &  0.0433 &   0.0149 &   0.3499 \\\\ \\hline\\noalign{\\smallskip} \\cesamB  & 3.2B  &   0.7957 &   0.9969 &   0.7723  &   0.0376 &   0.0325 &   0.0099 \\\\ \\cesamMP & 3.2B  &   0.7927 &   0.9968 &   0.7723 &  0.0375 &   0.0328 &   0.0100 \\\\ \\cles    & 3.2B  &   0.7956 &   0.9969 &   0.7796 &0.0373 &   0.0319 &   0.0101 \\\\ \\garstec & 3.2B  &   0.7953 &   0.9969 &   0.7767 &   0.0372 &   0.0318 &   0.0100 \\\\ \\hline\\noalign{\\smallskip} \\cesamB  & 3.2C  &   0.7921 &   0.9972 &   0.7768 &--&--&-- \\\\ \\cesamMP & 3.2C  &   0.7913 &   0.9972 &   0.7765 &--&--&-- \\\\ \\cles    & 3.2C  &   0.7903 &   0.9971 &   0.7843 &--&--&-- \\\\ \\garstec & 3.2C  &   0.7931 &   0.9972 &   0.7814 &--&--&-- \\\\ \\hline\\noalign{\\smallskip} \\cesamB  & 3.3A  &   0.8916 &   0.9999 &   0.8840 &   0.0641 &   0.0575 &   0.3500 \\\\ \\cesamMP & 3.3A  &   0.8893 &   0.9998 &   0.8631 &   0.0641 &   0.0572 &   0.3490 \\\\ \\cles    & 3.3A  &   0.8893 &   0.9998 &   0.8947 &   0.0633 &   0.0557 &   0.3503 \\\\ \\garstec & 3.3A  &   0.8920 &   0.9999 &   0.8955 &   0.0631 &   0.0554 &   0.3500 \\\\ \\hline\\noalign{\\smallskip} \\cesamB  & 3.3B  &   0.8336 &   0.9990 &   0.7866 &0.0389 &   0.0404 &   0.0099 \\\\ \\cesamMP & 3.3B  &   0.8342 &   0.9990 &   0.7886&   0.0385 &   0.0396 &   0.0100 \\\\ \\cles    & 3.3B  &   0.8335 &   0.9990 &   0.7948&   0.0387 &   0.0399 &   0.0101 \\\\ \\garstec & 3.3B  &   0.8334 &   0.9989 &   0.7923 &   0.0384 &   0.0391 &   0.0100 \\\\ \\hline\\noalign{\\smallskip} \\cesamB  & 3.3C  &   0.8534 &   0.9996 &   0.7902 &--&--&-- \\\\ \\cesamMP & 3.3C  &   0.8517 &   0.9995 &   0.7926 &--&--&-- \\\\ \\cles    & 3.3C  &   0.8526 &   0.9996 &   0.7988 &--&--&-- \\\\ \\garstec & 3.3C  &   0.8514 &   0.9995 &   0.7963 &--&--&-- \\\\ \\hline\\noalign{\\smallskip} \\end{tabular} \\label{tab:task3-conv} \\end{table*}   We have then discussed each case individually and tried to identify the origin of the differences.  The way the codes handle the {\\small OPAL-EOS} tables has an impact on the output thermodynamical properties of the models. In particular, the choice of the thermodynamical quantities to be taken from the tables and of those to be recalculated from others by means of thermodynamic relations is critical because it is known that some of the thermodynamical quantities tabulated in the {\\small OPAL} tables are inconsistent \\citep{2003ApJ...583.1004B}. In particular, it has been shown by one of us (IW) during one of the {\\esta} workshops that it is better not to use the tabulated $C_V$-value. Some further detailed comparisons of {\\cles} and {\\cesam} models by \\citet{jm-apss}, have demonstrated that these inconsistencies lead to differences in the stellar models and their oscillation frequencies substantially dominating the uncertainties resulting from the use of different interpolation tools. Similarly, the differences in the opacity derived by the codes do not come from the different interpolation schemes but mainly from the differences in the opacity tables themselves. Discrepancies depend on stellar mass and for the cases considered in this study the maximum differences in opacities are of the order of 2\\% for 2{\\msol} models except in a very narrow zone close to $\\log T\\simeq 4.0$ where there are at a few percents level for all models due to the way the {\\small OPAL95} and {\\small AF94} tables are combined.  In unevolved models, differences have been found that pertain either to the lack of a detailed calculation of the {\\PMS} phase or to the simplifications in the nuclear reactions in the CN cycle. In evolved models we have  identified differences which are due to the method used to solve the set of equations governing the temporal evolution of the chemical composition. The shape and position of the $\\mu$ gradient and the numerical handling of the temporal evolution of the border of the convection zones are critical as well. The star keeps the memory of the displacements of the convective core (either growing or receding) through the $\\mu$ gradient. Models with microscopic diffusion show differences in the He and metals distributions which result from differences in the diffusion velocities and that affect in turn thermodynamic quantities and therefore the oscillation frequencies. The situation is particularly thorny for models that undergo semiconvection, either below the convective zone or at the border of the convective core (for models with diffusion), because none of the codes treats this phenomenon.  We found that differences in the radius at the bottom of the convective envelope are small, lower than $0.7$\\%. The differences in the mass of the convective core are sometimes large as in models that undergo semiconvection or in {\\PMS} and low-mass {\\ZAMS} models (up to $17-30$\\%). In other models, the mass of the convective core differs by $0.5$ to $5$\\%. Differences in the surface helium abundance in models including microscopic diffusion are of a few per cent except for the $1.3${\\msol} model on the {\\MS} where they are in the range $15-30$\\% due to the different formalisms used to treat diffusion (see Sect.~\\ref{sec:task3}).  We have examined the differences in the oscillation frequencies of p and g modes of degrees $\\ell= 0, 1, 2, 3$. For solar-type stars we also calculated and examined the large frequency separation for $\\ell=0, 1$ and the frequency separation ratios defined by \\cite{1999ASPC..173..257R}.  For solar-type stars on the {\\MS}, the differences in the frequencies calculated by the different codes are in the range $0.05-0.1\\mu$Hz (for $\\ell=0$, no microscopic diffusion) and $0.1-0.2\\mu$Hz (models with diffusion). For advanced models ({\\TAMS} or {\\SGB}) differences are larger (up to $1\\mu$Hz for $\\ell=0$ modes). We find that the frequency separation ratios are in excellent agreement which confirms that the differences found in the frequencies and in the large frequency separation have their origin in near-surface effects. Differences are larger in models including microscopic diffusion where the sound speed differences are larger (in the centre and at the borders of convection zones). In addition, in evolved models, at low frequency for $\\ell=1$ modes, we found differences of up to $4\\mu$Hz that result from differences in the Brunt-V\\\"ais\\\"al\\\"a frequency and from the mixed character of the modes. For stars of $2.0$\\msol, frequency differences ($\\ell=0, 1$ modes) are lower than $0.5\\mu$Hz ({\\PMS}) and may reach $1\\mu$Hz for the evolved model with overshooting. They are due to structure differences in regions close to the surface, in the region of second He-ionisation and close to the centre.  In the evolved model some modes are mixed modes sensitive to the location of and features in the $\\mu$ gradient. Finally, we found that the frequency differences of the massive models ($3$ and $5$\\msol) are generally smaller than $0.2\\mu$Hz.   This thorough comparison work has proven to be very useful in understanding in detail the methods used to handle the calculation of stellar models in different stellar evolution codes. Several bugs and inconsistencies in the codes have been found and corrected. The comparisons have shown that some numerical methods had to be improved and that several simplifications made in the input physics are no longer satisfactory if models of high precision are needed, in particular for asteroseismic applications. We are aware of the weaknesses of the models and therefore of the need for further developments and improvements to bring to them and we are able to give an estimate of the precision they can reach. This gives us confidence on their ability to interpret the asteroseismic observations which are beginning to be delivered by the {\\corot} mission where an accuracy of a few $10^{-7}\\ {\\mathrm{Hz}}$ is expected on the oscillation frequencies \\citep{michel06} as well as those that will come from future missions as {\\small NASA}'s Kepler mission to be launched in 2009 \\citep{2007CoAst.150..350C}."
    },
    "0801/0801.3663_arXiv.txt": {
        "abstract": "We present a systematic X-ray study of eight active galactic nuclei (AGNs) with intermediate mass black holes ($M_{BH} \\sim 8-95\\times 10^4{\\rm~M_{\\odot}}$) based on 12 \\xmm{} observations. The sample includes the two prototype AGNs in this class -- NGC~4395 and POX~52 and six other AGNs discovered with the Sloan Digitized Sky Survey. These AGNs show some of the strongest X-ray variability with the normalized excess variances being the largest and the power density break time scales being the shortest observed among radio-quiet AGNs. The excess variance -- luminosity correlation appears to depend on both the BH mass and the Eddington luminosity ratio. The break time scale -- black hole mass relations for AGN with IMBHs are consistent with that observed for massive AGNs. We find that the FWHM of the H$\\beta$/H$\\alpha$ line is uncorrelated with the BH mass, but shows strong anticorrelation with the Eddington luminosity ratio.  Four AGNs show clear evidence for soft X-ray excess emission ($kT_{in} \\sim 150-200\\ev$). X-ray spectra of three other AGNs are consistent with the presence of the soft excess emission. NGC~4395 with lowest $L/L_{Edd}$ lacks the soft excess emission.  Evidently small black mass is not the primary driver of strong soft X-ray excess emission from AGNs.  The X-ray spectral properties and optical-to-X-ray spectral energy distributions of these AGNs are similar to those of Seyfert 1 galaxies. The observed X-ray/UV properties of AGNs with IMBHs are consistent with these AGNs being low mass extension of more massive AGNs; those with high Eddington luminosity ratio looking more like narrow-line Seyfert 1s while those with low $L/L_{Edd}$ looking more like broad-line Seyfert 1s. ",
        "introduction": "The existence of astrophysical black holes in two mass ranges, stellar mass black holes with $\\sim 10M_{\\odot}$ in X-ray binaries and supermassive black holes (SMBHs) with masses in the range of $\\sim10^6 - 10^{9}M_{\\odot}$ at the centers of active galaxies, is well supported by observations. Intermediate mass black holes (IMBHs), bridging the gap between the stellar mass and supermassive black holes, remain relatively unexplored.  There was no observational evidence for such black holes for many years.  Recent optical and X-ray observations have revived the possibility of existence of IMBHs. There are indirect evidence for IMBHs with masses of $\\sim 10^4-10^6M_{\\odot}$ at the center of some galaxies based on radiative signatures or dynamical measurements \\citep{2003ApJ...588L..13F,2004ApJ...607...90B, 2004ApJ...610..722G,2007arXiv0707.2617G}. The AGNs with IMBHs form an important class for a number of reasons.  First, the growth of SMBHs in Seyfert galaxies and quasars is still not understood.  The nature of seed black hole (BH) is the major challenge to any cosmological BH growth model.  Observations of IMBHs at the dynamical centers of nearby galaxies would provide significant constraints on the nature of seed black holes for cosmological growth models.  The $M_{BH}-\\sigma_{*}$ relation strongly suggests co-evolution of galaxies and BHs, but it is not known if this relation is established early in the evolution process. If this relation is established late in the galaxy evolution, as suggested by \\cite{2003ApJ...593...56D}, then we should see deviations in the $M_{BH}-\\sigma_{*}$ relation in a population of IMBHs that is not yet fully grown. In this growing phase, the accretion process and the disk-corona geometry around the central IMBH could be markedly different than that around SMBHs in luminous AGNs. Therefore, it is crucial to study accretion onto the least massive central BHs in galaxies. AGNs with IMBHs also provide the missing link between the stellar mass and SMBHs and enable us to test if the accretion physics is the same at all scales. They will also help to understand the accretion physics as it would be easier to isolate certain characteristics arising solely from an extreme physical parameter, in this case, low mass BH. AGNs with IMBHs may reveal new accretion phenomena depending on the mode of accretion at low BH masses.  It is also possible that they pass through an active phase of accretion and contribute to the observed background radiation at some level, particularly to the X-ray background. The mergers of intermediate mass black holes with masses $10^5M_{\\odot}$ are expected to provide strong signals of gravitational waves for the {\\it Laser Interferometric Space Antenna} \\citep[e.g.,][]{2002MNRAS.331..805H}.  The two prototypes examples of IMBH AGNs are the late-type spiral galaxy NGC~4395 \\citep{2003ApJ...588L..13F} and the dwarf elliptical galaxy POX~52 \\citep{2004ApJ...607...90B}. \\cite{2004ApJ...610..722G} discovered 19 such AGNs using the Sloan Digital Sky Survey (SDSS). They have increased their sample by an order of magnitude to 174 using the fourth data release of the SDSS \\citep{2007arXiv0707.2617G}. NGC~4395 has the emission properties of a type 1 AGN, with broad optical and UV emission lines and a point-like hard X-ray source. Its X-ray emission is extremely rapidly variable \\citep[see e.g.,][]{2005MNRAS.356..524V}. \\cite{2003ApJ...588L..13F} argued that the true mass of the black hole is likely to be $\\sim 10^4-10^5M_{\\odot}$, while reverberation mapping of NGC~4395 has provided a BH mass of $(3.6\\pm1.1)\\times10^{5}M_{\\odot}$ \\citep{2005ApJ...632..799P}.  Another recently discovered AGN with an IMBH is POX~52, a dwarf elliptical galaxy at $z=0.021$. It was discovered in an objective-prism search for emission line objects by \\cite{1981A&AS...44..229K} who also noted its star-like appearance in the Palomar Sky Survey images.  Follow-up imaging and spectroscopic observations by \\cite{1987AJ.....93...29K} revealed a dwarf galaxy with an AGN spectrum. The object was classified as a Seyfert 2 galaxy by \\citet{1981A&AS...44..229K} based on the flux-ratios of narrow emission lines although a weak broad component of the H$\\beta$ line was seen with a full width at half maximum, FWHM $\\sim 840\\kms$.  The galaxy is indeed quite small with a diameter $\\sim 6.4{\\rm~kpc}$ and its absolute magnitude is only $M_{V} = -16.9$. The luminosity of the nucleus is comparable to that of a mild Seyfert 2 nucleus but the ratio of nuclear to host-galaxy luminosity is 1.1, a typical value for quasars \\citep{1987AJ.....93...29K}.  \\cite{2004ApJ...607...90B} obtained a high quality optical spectra of POX~52 at the Keck that revealed an emission line spectrum very similar to that of the dwarf Seyfert 1 galaxy NGC~4395, with broad components to the permitted lines and reclassified POX~52 as a Seyfert 1 galaxy.  This galaxy has a central velocity dispersion of $36\\kms$, which yields a mass of $1.4\\times 10^{5}{\\rm~M_{\\odot}}$, again consistent with that derived from the H$\\beta$ line width \\citep{2004ApJ...607...90B}.  \\begin{table*} \\centering \\caption{General properties of AGN with intermediate mass black holes observed with \\xmm{} \\label{t1}} \\begin{tabular}{lcccccccc} \\tableline\\tableline Object  &  \\multicolumn{2}{c}{Position} & Redshift & Galactic $N_H$ & FWHM(${\\rm km~s^{-1}}$) & $M_{BH}$(${\\rm M_{\\odot}}$) & $\\frac{L_{bol}}{L_{Edd}}$ & Reference \\\\ & $\\alpha$(J2000.0) & $\\delta$(J2000.0) & ($z$) & ($10^{20}{\\rm~cm^{-2}}$) & ($H\\alpha$ or $H\\beta$) &    &  & \\\\ \\tableline NGC~4395 &  12 25 48.93 & +33 32 47.8 & $0.001$ & $1.85$ & $1500$ & $3.6\\times10^{5}$ & $1.2\\times10^{-3}$ & (1) \\\\ POX~52   &  12 02 56.90 & --20 56 03.0  & $0.021$  & $3.85$ & $760$ & $1.6\\times10^{5}$ & $0.5-1$ & (2)  \\\\ SDSS~J010712.03+140844.9 & 01 07 12.03 & +14 08 44.9 & $0.0768$ & $3.44$ & $830$ & $7.2\\times10^{5}$ & $1.24$  & (3)  \\\\ SDSS J024912.86-081525.6 & 02 49 12.86 & --08 15 25.6 & $0.0295$ & $3.67$ & $732$ &  $8.3\\times10^{4}$ & $0.71$ & (3) \\\\ SDSS J082912.67+500652.3 & 08 29 12.67 & +50 06 52.3 & $0.0434$ & $4.08$ & $870$ & $5.2\\times10^{5}$ & $0.95$ & (3) \\\\ SDSS J114008.71+030711.4 & 11 40 08.71 & +03 07 11.4 & $0.0811$ & $1.91$ & $591$ & $5.9\\times10^{5}$ & $2.98$ & (3) \\\\ SDSS J135724.52+652505.8 & 13 57 24.52 &  +65 25 05.8 & $0.106$ & $1.36$ & $872$ &  $9.5\\times10^{5}$ & $1.21$ & (3)  \\\\ SDSS J143450.62+033842.5 & 14 34 50.71 & +03 38 40.4 & $0.0284$ & $2.43$ & $1089$ & $2.0\\times10^{5}$ & $0.33$ & (3) \\\\ \\tableline \\end{tabular} References: (1) \\citet{2005ApJ...632..799P}, (2) \\cite{2004ApJ...607...90B} (3) \\cite{2004ApJ...610..722G}. \\end{table*}  AGNs with IMBHs are rare compared to luminous AGNs.  One explanation is that they are short-lived. Presumably, all AGNs grow fast in their early phase until their BH mass increases to $10^6{\\rm~M_{\\odot}}$.  During this fast-growth phase, the accretion rate may be close to Eddington or super-Eddington. The relative accretion rates of $\\sim 1$ for the SDSS AGNs with IMBH, derived by \\cite{2004ApJ...610..722G} from the optical data, favors the above scenario. However, all AGNs with IMBH are not accreting at high rates e.g., NGC~4395 is underlumnious for its BH mass with $L/L_{Edd} \\sim 1.2\\times 10^{-3}$ \\citep{2005ApJ...632..799P} or at most $0.2$ \\citep{2005MNRAS.356..524V} In fact, very little is known about the accretion process onto IMBHs in AGNs. The only AGN that has been well studied in X-rays is NGC~4395. \\cite{2000MNRAS.318..879I} presented the ASCA spectrum that showed a power-law continuum of photon index $\\Gamma=1.7\\pm0.3$ with a Fe K line marginally detected at $\\sim 6.4\\kev$. The soft-X-ray emission below $3\\kev$ is strongly attenuated by absorption. The X-ray spectrum in this absorption band showed a dramatic change in response to the variation in continuum luminosity. \\cite{2000MNRAS.318..879I} concluded that the nuclear source of NGC~4395 is consistent with a scaled-down version of higher-luminosity Seyfert nuclei.  \\cite{2005MNRAS.356..524V} have studied the exceptional X-ray variability of NGC~4395. The variations observed are among the most violent seen in an AGN to date, with the fractional rms amplitude exceeding $100\\%$ in the softest band.  \\cite{2007ApJ...656...84G} have performed snapshot ($5\\ks$ exposure) \\chandra{} ACIS X-ray observations of IMBH AGNs discovered with SDSS. They detected 8 of the 10 AGNs with IMBH with a significance $\\ge 3\\sigma$. They measured the $0.5-2\\kev$ photon indices in the range of $1-2.7$, consistent with that for luminous AGNs, implying that the BH mass is not the fundamental driver for the soft X-ray spectral shape. However, this result is not robust as some of the measurements are based on hardness ratios. Moreover, with $5\\ks$ \\chandra{} exposures, it was not possible to make detailed spectral and temporal analysis of any of the AGNs.  \\begin{table*} \\centering \\caption{Log of \\xmm{} observations \\label{t2}} \\begin{tabular}{lcccc} \\tableline\\tableline Object &  ObsID & Start date & pn rate & pn exposure  \\\\ &        &            & (${\\rm~counts~s^{-1}}$) & (${\\rm~ks}$) \\\\ \\tableline NGC~4395 & 0112522001(1) & 2002-06-12 & \t \\\\ & 0112521901(2) & 2002-05-31 &  $0.56\\pm0.007$ & $12.5$ \\\\ & 0142830101(3) & 2003-11-30 & $1.13\\pm0.004$ & $93.0$  \\\\ & 0112522701(4) & 2003-01-03 & $1.14\\pm0.014$ & $5.9$\\\\ POX~52   & 0302420101(1) & 2005-07-08 & $0.05\\pm0.001$ & $81.1$ \\\\ J0107+1408 & 0305920101(1) & 2005-07-22  & $0.27\\pm0.003$ & $24.9$ \\\\ J0249--0815 & 0303550101(1) & 2006-02-16 & $0.23\\pm0.008$ & $4.0$ \\\\ J0829+5006 & 0303550301(1) & 2006-03-27 &   \\\\ & 0303550801(2) & 2006-03-28 &  \\\\ & 0303550901(3) & 2006-04-26 & $0.95\\pm0.02$ & $2.3$ \\\\ & 0303551001(4) & 2006-04-26 &  \\\\ J1140+0307 & 0305920201(1) & 2005-12-03 & $0.66\\pm0.004$ & $35.1$ \\\\ J1357+6525 & 0305920301(1) & 2005-04-04 & $0.45\\pm0.005$ & $19.2$  \\\\ & 0305920501(2) & 2005-05-16 & $0.55\\pm0.01$ & $1.9$ \\\\ & 0305920601(3) & 2005-06-23 & $0.42\\pm0.006$ & $12.0$  \\\\ J1434+0338 & 0305920401(1) & 2005-08-18 & $0.10\\pm0.002$ & $22.0$ \\\\ \\tableline \\end{tabular} \\end{table*}  In this paper, we present X-ray temporal and spectral study of a sample of 8 AGNs with IMBHs.  These AGNs with $M_{BH} < 10^6M_{\\odot}$ were selected based on the availability of \\xmm{} data. In Table~\\ref{t1}, we list these AGNs along with their general properties.  The least massive AGN in our sample has only $M_{BH} \\sim 83000M\\odot$. The accretion rates of these AGNs, estimated from optical observations, has a large range ($L_{bol}/L_{Edd} \\sim 10^{-3} - 3$), suggesting a large diversity in the X-ray properties of these AGNs. We also present optical/UV emission observed with \\xmm{} optical monitor (OM).  We describe \\xmm{} observations and data reduction in Section 2. In Section 3, we present temporal analysis of the EPIC-pn data, followed by spectral analysis in Section 4. We present the analysis of OM data and derive optical-to-X-ray spectral indices ($\\alpha_{ox}$) in Section 5.  We discuss the results in Section 6, followed by a summary of our study in Section 7.   We assume the following cosmological parameters to calculate distances; $H_0 = 71{\\rm~km~s^{-1}~Mpc^{-1}}$, $\\Omega_m = 0.27$, and $\\Omega_{\\Lambda} = 0.73$. In the following, we abbreviate the names of the AGNs with IMBHs discovered with the SDSS e.g., we mention SDSS~J082912.67+500652.3 as J0829+5006. ",
        "conclusions": "We have examined X-ray/UV properties of eight AGNs with IMBHs using 12 \\xmm{} observations. These constitute the first X-ray study of seven of these AGNs excluding NGC~4395. The primary purpose of this study is to investigate dependance of X-ray spectral and tempral properties on BH mass as these AGNs have the lowest BH masses among all AGNs studied in X-rays to date (see Table~\\ref{t1}). All AGNs are strong X-ray/UV sources with $0.3-10\\kev$ X-ray luminosities in the range of $4.5\\times10^{40} - 2.3\\times10^{43} {\\rm~ergs~s^{-1}}$ and UV luminosity at $2500{\\rm \\AA}$ in the range of $1.1\\times10^{40} - 3.4\\times10^{43}  {\\rm~ergs~s^{-1}}$.   \\begin{figure} \\centering \\includegraphics[width=10cm]{fig8.ps} \\caption{Strength of the soft X-ray excess emission in the $0.6-10\\kev$ band  relative to the power-law in the same band plotted as a function of $\\Gamma$ for AGNs with IMBH (filled circles) and NLS1s (open circles) from \\cite{1999MNRAS.309..113V}.  The photon indices were derived from the best-fit models with an MCD component for all IMBH AGNs except NGC~4395. } \\label{soft_excess_gamma} \\end{figure}  \\subsection{X-ray continuum and soft X-ray excess emission} The $0.3-10\\kev$ primary X-ray continua of all AGNs with IMBHs consists of either of a single power-law or a soft X-ray excess component and a power-law. Only four AGNs show clear evidence for soft excess emission that is well described by multicolor disk blackbody with $kT_{in} \\sim 150-200\\ev$. The soft excess emission from the four contrinute only $ltsim50\\%$ in the $0.3-2\\kev$ band. X-ray spectra of three other AGNs are also consistent with the presence of soft excess emission with similar temperatures but the component is not required statistically. NGC~4395 is the only AGN that clearly lacks strong soft X-ray emission.  Figure~\\ref{soft_excess_gamma} compares the strength of soft excess and the photon indices of AGNs with IMBHs and NLS1 galaxies studied by \\cite{1999MNRAS.309..113V}.  The strength of the soft excess emission is shown as the ratio of the soft excess luminosity ($L_{excess}$) to that of the power-law component ($L_{PL}$) in the $0.6-10\\kev$ band. $L_{excess}$ is the luminosity of the MCD component in case of AGNs with IMBH, while it is the luminosity of the blackbody component for NLS1 galaxies as \\cite{1999MNRAS.309..113V} parameterized the soft excess emission with a blackbody component. We have plotted only the upper limits for the three SDSS AGNs whose X-ray spectra are equally well described by a simple PL or PL+MCD.  NLS1s show a large range in the strength of their soft excess emission, contributing $0-68\\%$ to the $0.6-10\\kev$ band, while AGNs with IMBHs appear to show only weak soft excess emission, contributing only $<35\\%$ in the same energy band. NGC~4395 with the lowest $L/L_{Edd}$ has the weakest or no soft excess emission in our sample.  Evidently, black hole mass is not the primary driver of the strength of soft X-ray excess emission from AGNs. High $L/L_{Edd}$ is likely the necessary condition for the strong soft X-ray excess emission.  The  power-law photon indices of eight AGNs range from $\\sim 1.6-2.4$ which is similar to that observed from more massive AGNs. In particular, AGNs with IMBHs show similar range in thier photon indices as that of NLS1 galaxies (see Fig.~\\ref{soft_excess_gamma}). The temperature of the soft excess emission detected from the AGNs with IMBHs is also similar to that observed from NLS1 galaxies. In Figure~\\ref{gamma_kt}, we have plotted $kT_{in}$ and $\\Gamma$ for seven AGNs with IMBHs for which the data are consistent with the presence of soft excess emission. We also show a similar plot in Figure~\\ref{gamma_kt} for 19 NLS1 galaxies studied by \\cite{1999MNRAS.309..113V} based on $0.6-10\\kev$ \\asca{} data. For NLS1 galaxies and AGNs with IMBHs, the blackbody temperature appears to be uniform in the $100-300\\ev$ range irrespective of the photon index.  \\begin{figure} \\centering \\includegraphics[width=8.5cm]{fig9.ps} \\caption{$\\Gamma$ Vs MCD $kT_{in}$ for AGNs with IMBHs (filled circles) or blackbody $kT_{BB}$ for NLS1 galaxies (open circles) from Vaughan et al. (1999).} \\label{gamma_kt} \\end{figure}  \\subsection{X-ray Variability} Strong X-ray variability appears to be a general property of AGNs with IMBH, albeit at different amplitudes.  IMBH AGNs are among the most extreme variable radio-quiet AGNs as suggested by the comparison of the normalized excess variances in Fig.~\\ref{exc_var_lum}. The excess variances and X-ray luminosities of AGNs with the smallest black holes are also consistent with the anticorrelation found for BLS1 and NLS1s \\citep{1997ApJ...476...70N,1999ApJS..125..297L}. In Fig.~\\ref{exc_var_lum}, IMBH AGNs and BLS1s seem to form a strong antocorrelation and most of the dispersion in the $\\sigma^2_{NXS} - L_{2-10\\kev}$ relation is caused by the NLS1 galaxies. It is likely that the dispersion in the $\\sigma^2_{NXS} - L_{2-10\\kev}$ relation is due to the variation in the accretion rates ($\\dot{m_E} = L_{bol}/L_{Edd}$) as NLS1s have high accretion rates. This trend is indeed expected from the power density spectra. \\cite{2006Natur.444..730M} have shown a strong correlation between the break frequency, black hole mass and the accretion rate relative to the Eddington rate such that break time scale $T_{B} \\approx M_{BH}^{1.12}/\\dot{m}_E^{0.98}$. This implies that at constant black hole mass, the break frequency increases as $\\dot{m}_E^{0.98}$, thus increasing the integrated power over certain frequency range covering the high frequency part of the PDS even if the PDS normalization is the same. An increased integrated power also means more excess variance. Thus both the lower black hole mass and higher accretion lead to higher break frequencies and hence more excess variance, explaining Fig.~\\ref{exc_var_lum}. From optical studies, AGNs with IMBHs show a wide range of $L/L_{Edd}$ (Table\\ref{t1}). Accordingly, they also show a range in excess variance, with POX 52, having the highest $L/L_{Edd}$ also showing the highest excess variance.  \\begin{figure} \\centering \\includegraphics[width=9cm]{fig10.ps} \\caption{The relation between black hole mass and PDS break time scale for AGNs and BHBs. The AGNs with IMBHs, NGC~4395 and J1140+0307 are marked. The dotted lines represent linear scaling, $M_{BH} \\propto T_{br}$, based on the typical break times observed from the high/soft (shorter time scales) and low/hard states (longer time scales) of Cyg~X-1. See \\cite{2005MNRAS.359.1469M} for details.  } \\label{mbh_tb} \\end{figure}   The power density spectra of NGC~4395, derived by \\cite{2005MNRAS.356..524V}, and J1170+0307 are consistent with a flat power-law below a break and steep power-law above the break. The derived break frequencies for IMBH AGNs are the highest observed among all AGNs. The form of the PDS are broadly similar to that of more massive AGNs. In Figure~\\ref{mbh_tb}, we have plotted the break time scales and black holes masses of the two AGNs with IMBHs along with their massive counterparts. The break frequency and black hole mass for NGC~4395 were taken from \\cite{2005MNRAS.356..524V} and \\cite{2005ApJ...632..799P}, respectively. For massive AGNs, the data were taken from \\cite{2005MNRAS.363..586U}. The two dotted lines represent assumed linear scaling of black hole mass with break time scale based on the high/soft (shorter time scale at the same black hole mass) and low/hard states of Cyg X-1 \\citep[see][]{2005MNRAS.363..586U}. The data points located left to the dotted lines represent highly accreting AGNs. The break time scales and black hole mass of J1140+0307  is consistent with the $M_{BH} - T_{br}$ relation for AGNs and BHBs. However, NGC~4395 with low accretion rate ($\\dot{m}_E \\simeq 1.2\\times10^{-3}$; Table~\\ref{t1} and Table~\\ref{t7}) is inappropriately placed in Fig.~\\ref{mbh_tb}. This is due to an order of magnitude uncertainty in the black hole mass of NGC~4395.  \\cite{2003ApJ...588L..13F} argued that the true black hole mass of NGC~4395 is in the range of $10^4 - 10^5M_{\\odot}$ which is consistent with several other measurements \\citep[see][]{2005MNRAS.356..524V}. However, the black hole mass derived from the reverberation mapping by \\cite{2005ApJ...632..799P} is an order of magnitude larger than the best guess value of \\cite{2003ApJ...588L..13F}. When this lower mass of $10^4 - 10^5M_{\\odot}$ is used NGC~4395 data are consistent with other AGNs as was mentioned by \\cite{2005MNRAS.356..524V}.  With the $\\sigma^2_{NXS} - L_{2-10keV}$ and $M_{BH} - T_{br}$ relations for IMBH AGNs both being consistent with that for massive AGNs, the variability process of these low mass AGNs are intrinsically the same to that operating in massive AGNs. The extreme variability of AGNs with IMBHs is mainly due to their lower black hole masses.      \\subsection{$FWHM_{H\\beta} - \\Gamma_{X}$ relation} Seyfert 1 galaxies show a large range in the width of their optical permitted emission lines e.g., full width at half maximum (FWHM) of the ${\\rm H}\\beta$ line is found to be in the range $\\sim$ $1000 - 10000~{\\rm ~km~s^{-1}}$. NLS1 galaxies lie at the lower end of the line width distribution, they were originally identified with FWHM $({\\rm H}\\beta) \\ltsim 2000{\\rm ~km~s^{-1}}$ \\citep{1985ApJ...297..166O}.  By this definition, all the AGNs with IMBH qualify as NLS1 galaxies.  NLS1 galaxies also show distinct X-ray properties e.g., rapid variability and generally steep X-ray spectrum \\citep{1996A&A...305...53B,1999ApJS..125..297L, 1999ApJS..125..317L,2004AJ....127..156G}. NLS1 galaxies show a large dispersion in their soft ($0.1-2.4\\kev$) and hard X-ray ($2-10\\kev$) photon indices than the BLS1 galaxies with larger H$\\beta$ FWHM \\citep{1996A&A...305...53B,1997MNRAS.285L..25B}. The dispersion in the photon indices of IMBH AGNs is similar to that of NLS1 galaxies. The difference is that IMBH AGNs have lower H$\\beta$ FWHM on an average than that of NLS1 galaxies. The current best explanation of distinct properties of NLS1 galaxies is that they host low mass black holes that are accreting at high fraction of their Eddington rate. In this picture, the Balmer lines are relatively narrow due to low black hole masses. AGNs with IMBHs play here an important role to investigate if the narrower permitted lines arise solely due to the low black hole mass. In Figure~\\ref{hbeta}, we have plotted the black hole masses as a function of the FWHM of the Balmer lines H$\\beta$ or H$\\alpha$ line for the AGN with IMBHs.  If the black hole mass were the single physical parameter responsible for the width of Balmers H$\\beta$ or H$\\alpha$ line, we expect a strong correlation between black hole mass and FWHM of the H$\\beta$ or H$\\alpha$ line. However, there is no correlation between the two quantities.  We have used the reverberation mass for NGC~4395 in Fig~\\ref{hbeta}.  As discussed earlier, the measured PDS break frequency for NGC~4395 is consistent with a black hole mass in the range $10^4 - 10^5M\\odot$. However, changing the black hole mass of NGC~4395 in Fig.~\\ref{hbeta} does not alter our conclusion. In fact, NGC~4395 should have the narrowest H$\\beta$ among the AGNs with IMBH if the black hole mass were the only driving parameter. However, H$\\beta$ is broadest for NGC~4395 among the AGNs with IMBHs (see Table~\\ref{t1}). Evidently, the black hole mass alone is not responsible for the width of permitted lines.    \\begin{figure*} \\centering \\includegraphics[width=8.5cm]{fig11a.ps} \\includegraphics[width=8.5cm]{fig11b.ps} \\caption{Black hole mass ({\\it left panel}) and X-ray luminosity relative to the Eddington luminosity ({\\it right panel}) as a function of FWHM of the Balmer H$\\alpha$ or H$\\beta$ line. } \\label{hbeta} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[width=8.5cm]{fig12a.ps} \\includegraphics[width=8.5cm]{fig12b.ps} \\caption{X-ray photon index $\\Gamma$ as a function of X-ray luminosity relative to the Eddington luminosity ({\\it left panel}) and Eddington ratio ({\\it right panel}). } \\label{gamma} \\end{figure*}   It is well known from the reverberation mapping of AGNs that the size of the broad line region (BLR) scales with the optical luminosity \\citep[$R_{BLR} \\propto L^{0.70\\pm0.03}$; see e.g.,][]{2000ApJ...533..631K}. Thus, luminosity is another parameter in addition to the black hole mass that determines the width of the permitted lines. This is confirmed in Figure~\\ref{hbeta}, showing that X-ray luminosity relative to the Eddington luminosity is well correlated with the FWHM of the H$\\beta$ or H$\\alpha$ line. Similar correlation is also observed between the Balmer line widths and Eddington ratio ($L_{bol}/L_{Edd}$). The anti-correlation between the FWHM of the H$\\beta$ line and the X-ray photon index in NLS1s and BLS1s then implies that steep X-ray spectrum is also caused by $L/L_{Edd}$ which is a function of both black hole mass and luminosity. However, we know that the photon index of the X-ray spectrum does not depend on black hole mass as BHBs, massive AGNs and AGNs with IMBHs, all show similar range of photon indices.  The photon index, however, depends on the X-ray or bolometric luminosities relative to the Eddington luminosity or the accretion relative to the Eddington rate (see Figure~\\ref{gamma}). Thus low black hole masses and high absolute accretion rates resulting in high $\\dot{m}_E = L_{bol}/L_{Edd}$ are necessary criteria for AGNs to show strong anticorrelation between the $H\\beta$ line width and X-ray photon index. Properties of AGNs with low black hole mass and low accretion rates resulting in low $\\dot{m}_E$ e.g., NGC~4395, cannot be similar to those of NLS1 galaxies. The accretion processes in the low mass AGNs with low $\\dot{m}_E$ are likely similar to the massive AGNs with low $\\dot{m}_E$. This explains the departure of NGC~4395 from the anticorrelation between the FWHM$_{H\\beta}$ line and X-ray $\\Gamma$ and similarity with BLS1s in the antocorrelation between normalized excess variance and $2-10\\kev$ luminosity. This is indeed consistent with the fact that all NLS1s do not show the steep X-ray spectrum. \\cite{2002ApJ...581L..71D} showed NLS1s with low and high accretion rates relative to the Eddington rate differ in their X-ray characteristics and NLS1s with low $\\dot{m}_E$ are similar to the BLS1 galaxies in terms of soft X-ray excess and shape of their X-ray spectra. Based on \\chandra{} observations of optically selected, X-ray weak NLS1 galaxies, \\cite{2004ApJ...610..737W} showed that strong soft X-ray excess emission and steep X-ray spectrum is not a universal characteristic of NLS1s and many NLS1s have low $\\dot{m}_E$.  \\subsection{Intrinsic absorption} Among the eight AGNs with IMBHs, NGC~4395 showed the most complex X-ray spectrum. The primary continuum of NGC~4395 observed on 2003 November is modified by multiple absorbers -- (i) fully covering neutral absorber ($N_H \\sim 5\\times10^{20}{\\rm~cm^{-2}}$), (ii) partially covering neutral absorber ($N_H \\sim 3.6\\times 10^{23}{\\rm~cm^{-2}}$ with a covering fraction of $\\sim 50\\%$), and ($iii$) three warm absorbers with columns in the range $N_W \\sim 0.3-5.8)\\times10^{21}{\\rm~cm^{-2}}$ and ionization parameter $log\\xi \\sim 0-3.5$).  The X-ray spectrum of POX~52 also shows evidence for a partially covering neutral absorber ($N_H \\sim 10^{23}{\\rm~cm^{-2}}$, covering fraction $\\sim 90\\%$) and warm absorber ($N_W \\sim 7\\times 10^{21}{\\rm~cm^{-2}}$, $log \\xi \\sim 2.1$). The columns and ionization parameters of the warm absorbers in NGC~4395 and POX~52 are well within the range observed for Seyfert 1 galaxies studied by \\cite{2005A&A...431..111B}. Hence the presence of warm absorbers and physical conditions do not depend on black hole mass or accretion rate.  The dramatic spectral variability of NGC~4395 is well accounted by variation in the absorbing components. The covering fraction of the absorber changed from $\\sim 80\\%$ in 2002 May to $\\sim 50\\%$ in 2003 November while the absorption column increased in the same period from $N_H = 8_{-1}^{+2}\\times10^{22}{\\rm~cm^{-2}}$ in 2002 May to $3.6_{-0.3}^{+0.2}\\times10^{23}{\\rm~cm^{-2}}$ in 2003 November. In addition to the higher covering fraction of the neutral absorber, the warm absorber columns were also higher in 2002 May than 2003 November, though the highly ionized warm absorber was not detected in 2002 May. The decreased covering fraction and warm absorber columns in 2003 November resulted almost an order of magnitude increase in the observed $0.3-2\\kev$ flux.  Partial covering of primary X-ray source by cold matter is not uncommon in massive AGNs. The model was first invoked in the early days \\citep{1980ApJ...241L..13H}. Recently, the partial covering or patchy absorber has been identified in several NLS1 galaxies \\citep{2002MNRAS.329L...1B, 2004MNRAS.353.1064G, 2006AN....327.1071B, 2004AJ....127.3161G, 2007ApJ...668L.111G}. However, the origin, location, geometry and physical conditions of the partial covering absorbers are not well known. Variations in partial covering absorption has been observed on timescales of weeks to years \\citep{2004MNRAS.353.1064G, 2007ApJ...668L.111G}. In this respect, NGC~4395 is not unusual except that it is an AGN with low accretion rates. Hence it is unlikely that accretion disk winds act as patchy absorber in NGC~4395 as some models invoke that a partial covering obscuration may arise if our line of sight passes through winds that are launched at intermediate radii of the disk \\citep[e.g.,][]{2000ApJ...545...63E,2007ApJ...661..693P}. In case of NGC~4395, the time scale of variations is long that is apparently not consistent with the toy model of \\cite{2000A&A...356..475A} suggesting thick clouds at $10-100$ Schwarzschild radii partially obscure the central source. However, we note that the variability of the partial covering absorption on short time scales, days to weeks, is unknown. Monitoring observations of NGC~4395 with \\xmm{} will be very useful to test models for partial covering absorption.  \\subsection{Broadband continuum properties}  One of the great advantages of \\xmm{} is the availability of the optical monitor (OM), together with X-ray detectors. As a result, it is possible to get optical-to-X-ray spectral energy distribution of sources with simultaneous observations, thus bypassing the problems arising due to variability. The OM fluxes of our targets are listed in Table~\\ref{tab8} while the X-ray fluxes are given in Tables~\\ref{t5} and \\ref{t6}. The optical to X-ray spectral energy distribution is often parametrized by $\\alpha_{ox}$ as defined above (\\S 5); the $\\alpha_{ox}$ values of our targets are listed in Table~\\ref{tab9}.  It is useful to note here that the OM observations are taken with one or more of optical/UV filters with central wavelengths at 2070\\AA, 2298\\AA, 2905\\AA. Thus the actual observations are taken at wavelengths close to 2500\\AA, at which the monochromatic optical flux is calculated in the standard $\\alpha_{ox}$ definition. For this reason our measurements of $f_{\\nu}$(2500\\AA) are far more accurate than those calculated from extrapolations of longer wavelengths \\citep[e.g.,][]{2005AJ....130..387S}. It is well known that AGNs are increasingly X-ray faint relative to UV, for higher luminosities. In Figure~\\ref{alphaox} we have plotted the $\\alpha_{ox}$ vs. $L_{\\nu}$(2500\\AA) (erg s$^{-1}$ Hz$^{-1}$). The \\citet{2005AJ....130..387S} relation for the optically selected AGNs is shown by the solid line and our targets with IMBHs are shown as solid circles. Ours are all low luminosity sources; all AGNs in the Strateva et al. sample lie to the right of the vertical dotted line. Our sample of IMBH AGNs extends the relation to lower luminosities.  The horizontal line marks the lowest $\\alpha_{ox}$ for $\\log L_{\\nu} <28$ in the Strateva et al. data.  While six of our sources appear to be consistent with the Strateva et al. relation, SDSS J1434+0338 and POX 52 are clearly X-ray week. While broad absorption line quasars are often undetected in X-rays, and thus observed as X-ray faint, their intrinsic X-ray fluxes, corrected for absorption are usually similar to non-BALQSOs; \\citealt{1996ApJ...462..637G}), X-ray weakness of the two AGNs with IMBHs is not due to intrinsic absorption as the $\\alpha_{ox}$ values were derived after correctiong for the Galactic as well any intrinsic neutral/warm absorption. It is worth noting that POX~52 and J1434+0338 are only mildly X-ray weak. NLS1 galaxies studied by Gallo et al. (2006) show a large range $-1.76 \\le \\alpha_{ox} \\le -0.94$. Hence the optical-to-X-ray spectral energy distribution of all eight AGNs with IMBHs are entirely consistent with that of NLS1 galaxies.  NGC~4395 appears to be weak in the UV and bright in X-rays as it has largest $\\alpha_{ox}$ among the IMBH AGNs. This is also evident from the $L_{X}/L_{Edd}$ and $L_{bol}/L_{Edd}$ values. While the contribution of $0.3-10\\kev$ X-ray luminosity to the bolomtertic luminosity is only $6-21\\%$ for the six SDSS AGNs and POX~52, the contribution is $\\sim 80\\%$ in the case of NGC~4395. This may suggest different disk-corona geometries for NGC~4395 and other AGNs with IMBHs. If the accretion disk in NGC~4395 is truncated at large inner radii, then the disk is expected to be cooler and hence weaker optical/UV emission. Such truncated disks are thought to exist in the hard state of BHBs when the accretion rate is low \\citep[see e.g.,][]{2006ARA&A..44...49R}. The lack of soft X-ray excess emission, low $L/L_{Edd}$, extremely rapid variability and the lack of UV emission are all consistent with a truncated disk and NGC~4395 is likely in a low state similar to the low/hard state of BHBs.    \\begin{figure} \\centering \\includegraphics[width=10cm]{fig13.ps} \\caption{$\\alpha_{ox}$ as a function of monochromatic luminosity at $2500{\\rm \\AA}$ for AGNs with IMBHs. The solid line represents the primary dependance of $\\alpha_{ox}$ on $log(L_{2500{\\rm \\AA}})$ ($\\alpha_{ox} = -0.136 logL_{UV} + 2.616$) for a large sample of 195 AGNs \\citep{2005AJ....130..387S}. The horizontal dash-dotted line represents the lower bounds to scatter in the $\\alpha_{ox}-L_{\\nu}$ relation for the 195 AGNs. The vertocal dotted line marks the lower end of the $L_{\\nu}(2500{\\rm \\AA})$ range in the large sample of \\citet{2005AJ....130..387S} } \\label{alphaox} \\end{figure}  \\subsection{Nature of AGNs with IMBHs} The main objective of this study is to investigate the dependance of X-ray properties on BH mass and thereby to identify any departure in the accretion process that may be related to the growing phase in the early evolution of BHs and galaxies.  As shown above, though AGNs with IMBHs show strong X-ray variability, the excess variance -- luminosity  and the PDS break -- BH mass relations are consistent with the variability properties of AGNs with SMBHs. The short break time scales are consistent with the low BH mass of AGNs with IMBHs. The primary X-ray continuum and the optical-to-X-ray spectral energy distributions are also similar to that observed from massive AGNs. The propereties of neutral and warm absorbers in NGC~4395 and POX~52 are also not distinct when compared to that observed from Seyfert galaxies. Thus we conclude that the observed X-ray and UV properties of AGNs with IMBHs are consistent with the low BH mass extension of SMBHs in Seyfert galaxies and quasars. The accretion process in AGNs with IMBHs is similar to that in more massive AGNs. There is no clear indication of any departure in the observational characteristics that may be related to the early evolution of black holes and galaxies if the co-evolution of BHs and galaxies is not established in the early phases. This conclusion is consistent with the fact that two of AGNs NGC~4395 and POX~52 follow the same $M_{BH}- \\sigma^{*}$ relation known for the massive AGNs and normal galaxies."
    },
    "0801/0801.1450_arXiv.txt": {
        "abstract": "This paper provides an analytical description of the transport of ultrahigh energy cosmic rays in an inhomogeneously magnetized intergalactic medium. This latter is modeled as a collection of magnetized scattering centers such as radio cocoons, magnetized galactic winds, clusters or magnetized filaments of large scale structure, with negligible magnetic fields in between. Magnetic deflection is no longer a continuous process, it is rather dominated by scattering events. We study the interaction between high energy cosmic rays and the scattering agents. We then compute the optical depth of the Universe to cosmic ray scattering and discuss the phenomological consequences for various source scenarios. For typical parameters of the scattering centers, the optical depth is greater than unity at $5\\times10^{19}\\,$eV, but the total angular deflection is smaller than unity.  One important consequence of this scenario is the possibility that the last scattering center encountered by a cosmic ray be mistaken with the source of this cosmic ray. In particular, we suggest that part of the correlation recently reported by the Pierre Auger Observatory may be affected by such delusion: this experiment may be observing in part the last scattering surface of ultrahigh energy cosmic rays rather than their source population.  Since the optical depth falls rapidly with increasing energy, one should probe the arrival directions of the highest energy events beyond $10^{20}\\,$eV on an event by event basis to circumvent this effect. ",
        "introduction": "The problem of the origin of ultrahigh energy cosmic rays has generally been expressed as the conjunction of two questions: {\\it (i)} how can particles be accelerated to energies in excess of $10^{20}\\,$eV?  {\\it (ii)} why is the source not seen in the arrival directions of the highest energy events? Progress on the former question has certainly been hindered by our relative lack of knowledge on acceleration mechanisms and high energy processes in the most powerful astrophysical objects. Regarding the latter question, progress has been mostly limited by the scarcity of experimental data at the highest energies, at least until very recently.  Indeed the first results of the Pierre Auger Observatory, which have just been published, report a significant correlation of the arrival directions of the highest energy events with a catalog of active galactic nuclei (AGN) closer than 75~Mpc~\\cite{Auger1,Auger2}. This observation certainly marks an important step in the search for the source of ultrahigh energy cosmic rays. However one should not overinterpret the significance of these results. In particular, the likelihood of the reported coincidence rests on the comparison with isotropic arrival directions, yet the large scale structure is known to be highly inhomogeneous at least up to 75~Mpc. Since AGN are known to cluster with the large scale structure, one cannot exclude at present that the observed correlation remains a coincidence if the source itself clusters with the large scale structure~\\cite{Auger2}. More will be said on these data in Section~\\ref{sec:disc} of the present paper.  Furthermore, there exist other (and sometimes contradictory) claims in the literature on the existence of correlations of ultrahigh energy cosmic ray arrival directions with various source catalogs~\\cite{Tea95}, the strongest being the association with BL Lacertae objects reported in Refs.~\\cite{TT01,GTTT04,TT04} (see also Refs.~\\cite{EFS04,HiRes06}). Since the existing data is so scarce at the highest energies, the assessment of the statistical significance remains a difficult task. Finally, the reported evidence for multiplets of events tends to suggest that the source lies in the arrival direction of the events clusters. However some of these clusters show interacting galaxies as the sole peculiar objects on the line of sight~\\cite{Uea00}, while a more recent multiplet appears correlated with interacting clusters of galaxies~\\cite{Fea05,HiRes05}. Taken at face value, all these claims do not allow to draw a clear and consistent picture of the source of ultrahigh energy cosmic rays.  It is admitted that cosmic magnetic fields must play a key role in this puzzle, although which role exactly is also a question that is still seeking for an answer. And this source of uncertainty is in turn related to our poor knowledge of the strength and the distribution of extragalactic magnetic fields (see Ref.~\\cite{K94,V97} for detailed reviews of existing data).  There exists a rather large body of literature on the relation between cosmic magnetic fields and ultrahigh energy cosmic rays. Most studies have constructed models of extragalactic magnetic fields and then resorted to Monte Carlo simulations in order to quantify the influence of these fields on the time, energy and angular images expected in large scale detectors. One must however underline the analytical works of Refs.~\\cite{WM96,1999astro.ph..7060A} on cosmic ray transport in tangled extragalactic magnetic fields of homogeneous power, those of Refs.~\\cite{2002JHEP...03..045H,2002JHEP...07..006H} which discuss the particular effect of magnetic lensing and finally Refs.~\\cite{1979Natur.281..356W,1990A&A...232..582B,1999PhRvD..59b3001B} which discuss diffusive transport in a magnetized supercluster.  Earlier numerical studies have addressed the phenomenology of ultrahigh energy proton propagation in tangled magnetic fields of homogeneous power ~\\cite{1997APh.....7..213L,1997APh.....6..337M,LSOS97,1997PhRvD..56.4470S, 1998APh.....9..221C,2000PhRvD..62i3005S,2003ApJ...586.1211Y,2004APh....21..609D}. There has since been a trend toward more realistic magnetic field configurations. For instance, Refs.~\\cite{1998ApJ...505L..79M,1999APh....10..141S,1999astro.ph..3124L,2001PASJ...53.1153I,2002PhRvD..65b3004I,2002PhRvD..66h3002I} have studied the diffusive or non-diffusive propagation in a magnetized local supercluster and Ref.~\\cite{2003PhRvD..68j3004S} has brought to light the spectral distortions induced by the interaction of ultrahigh energy cosmic rays with a supercluster harboring large scale regular magnetic fields. More recently, several studies have attempted to model a realistic configuration in which the magnetic field follows the matter density and then studied the transport of ultrahigh energy cosmic rays in the resulting structure. In order to construct the magnetic field, Refs.~\\cite{2003PhRvD..68d3002S,2004PhRvD..70d3007S, 2005PhRvD..72d3009A,DGST04,DGST05,KRC07,2007PhRvD..75j3001S} have used numerical simulations of large scale structure formation involving a passive magnetic field whose strength was normalized to the value measured in clusters of galaxies. Refs.~\\cite{1997astro.ph..7054M,2006ApJ...639..803T,TS07,KL07} have rather reconstructed the extragalactic magnetic field by scaling the field strength to the underlying density field.   In general, these studies have assumed the magnetic field to be all pervading (albeit, with a more or less pronounced degree of inhomogeneity) so that magnetic deflection has been modeled as a continuous process. This assumption has been relaxed in Ref.~\\cite{MTE01} which provides numerical simulations of cosmic ray arrival directions after scattering with fossils of radio-galaxy lobes. Similarly, Ref.~\\cite{2002astro.ph.10095B} has mentioned the possibility of discrete cosmic ray interactions with localized regions of enhanced magnetic fields, their discussion pointing toward clusters of galaxies as the main scattering agents.  This picture in which ultrahigh energy cosmic ray transport occurs through random discrete events is indeed more likely to be valid on distance scales up to a few hundreds of Mpc as a consequence of the high degree of clustering of matter in the Universe. For instance, even if the magnetic field were produced in a uniform manner at high redshift (see Ref.~\\cite{W02} for a review of models of the origin of large scale magnetic fields), then the present-day magnetic field should be highly inhomogeneous, as a result of the amplification of the magnetic field in the shear and compressive flows associated with the formation of non-linear structures~\\cite{1998AA...335...19R,SME04,DGST05,TS07,KRC07} (see also~\\cite{D06} for a general discussion). In these simulations, voids in the large scale structure are essentially deprived of magnetic field.  Furthermore, if one attributes the origin of the extragalactic magnetic field to pollution by a sub-class of galaxies, for instance starburst galaxies \\cite{KLH99,Bea00,BVE06} or radio-galaxies ~\\cite{RS68,FL01,GKW01}, the magnetic field configuration should resemble that of a percolating process (see Ref.~\\cite{BSW05} for a clear illustration). As explained further below, if the filling factor of the polluted regions becomes comparable to that of the filaments of large scale structure, the filaments themselves become the scattering agents as in Refs.~\\cite{SME04,DGST05,TS07,KRC07}.   The goal of the present paper is to provide an analytical description of ultrahigh energy cosmic ray transport in such an inhomogeneous medium, which is modeled by scattering centers embedded in an unmagnetized intergalactic medium.  These scattering centers comprise the filaments just as the clusters of galaxies but also all possible regions of locally enhanced magnetic fields, such as galactic winds, groups of galaxies, large scale structure shocks and fossil radio-galaxy cocoons. One motivation of the present work is thus to make progress toward a more realistic magnetic field configuration which takes into account those localized regions of intense magnetic activity. In order to do so, we first sketch a census of relevant scattering centers~(Section~\\ref{sec:od}) then analyse their respective influence.  The present work is further motivated by the fact that Refs.~\\cite{SME04,DGST05,TS07,KRC07} diverge as to the conclusions they draw on the influence of the extragalactic magnetic fields on ultrahigh energy cosmic rays, even though they try to construct ab initio predictions for the distribution of these large scale magnetic fields. This difference stems from the uncertainty on the origin of these magnetic fields, not withstanding the complexity of modeling accurately the evolution of magnetic fields in the formation of large scale structure. Analytical tools become useful in this context as they allow to parametrize the influence of such magnetic fields on the images and spectra of ultrahigh energy cosmic rays. This in turn will help to deconvolve this effect from existing and upcoming data, and therefore to infer useful constraints on these magnetic fields.    In the present description, magnetic deflection is no longer a continuous process, but is instead dominated by scattering events. We thus use the notion of the optical depth of the Universe to ultrahigh energy cosmic ray scattering and discuss the phenomenological consequences. In particular, we show that the optical depth decreases very abruptly as the energy increases, because the source distance scale decreases due to increasing energy losses, and because the influence of cosmic magnetic fields diminishes with increasing energy.  We argue that the energy beyond which the Universe becomes translucent or transparent to cosmic ray scattering may be tantalizingly close to the threshold beyond which experiments search for counterparts, $E\\simeq 4-6\\times10^{19}\\,$eV. This could have profound consequences for our interpretation of existing data. For instance, if most sources lie beyond the last scattering surface, one could mistake the scattering centers on the last scattering surface (such as starbursts, old radio-galaxies or giant shock waves) with the source of ultrahigh energy cosmic rays.  The phenomenological consequences thus differ widely from the case of continuous deflection in an all-pervading medium. We thus discuss in some detail the expected effects and their relation to current and future observations of cosmic ray arrival directions.  This paper is laid out as follows. In Section~\\ref{sec:od}, we sketch a census of possible scattering centers and their influence on the optical depth of the Univers to cosmic ray scattering. We also calculate the distance to the last scattering surface and compare it to the expected source distance scale. In Section~\\ref{sec:crt}, we discuss the transport of cosmic rays in this strongly inhomogeneous medium and the expected observational consequences. We notably provide sky maps of the expected optical depth up to different distances for our local Universe. Finally, in Section~\\ref{sec:disc}, we summarize our findings and comment on the existing data in the framework of the present model. The physics of the interaction of cosmic rays with scattering centers is discussed in Appendix~\\ref{sec:appA}. ",
        "conclusions": "\\label{sec:disc}  \\subsection{Summary of present results} The present work has provided an analytical description of ultrahigh cosmic ray transport in highly structured extragalactic magnetic fields. The corresponding configuration of the extragalactic magnetic field is that of a collection of scattering centers, such as halos of radio-galaxies or starburst galaxies, or magnetized filaments, with a negligible magnetic field in between. Such a configuration is generally expected in scenarios in which the magnetic field is produced and ejected by a sub-class of galaxies, or generated at the accretion shock waves of large scale structure. Even if the magnetic field is rather generated at high redshift, subsequent amplification in the shear and compressive flows of large scale structure formation tends to produce a highly structured configuration, with strong fields in the filaments of galaxies and weak fields in the voids~\\cite{SME04,DGST05,TS07,KRC07}.  In our description, transport of cosmic rays is modeled as a sequence of interactions with the scattering centers, during which the particle acquires a non-zero deflection angle and time delay (with respect to straight line crossing of the magnetized region), see Appendix~\\ref{sec:appA}. In Section~\\ref{sec:od}, we have sketched a list of possible scattering centers and their characteristics (mean free path to scattering, magnetic field, coherence length and extent). We have then computed the optical depth $\\tau$ of the Universe to cosmic ray scattering as a function of energy and distance to the source, as well as the effective optical depth $\\tau_{\\rm eff}$ (which is defined in such a way as to become unity when the total angular deflection becomes unity). As discussed in Section~\\ref{sec:od}, the Universe can be translucent to cosmic ray scattering if $\\tau>1>\\tau_{\\rm eff}$, meaning that the total deflection is smaller than unity but non-zero, opaque if $\\tau>\\tau_{\\rm eff}>1$, or even transparent if $1>\\tau>\\tau_{\\rm eff}$. For typical values of the scattering centers parameters, it is expected that the Universe be translucent or opaque on the source distance scale and at energies close to the pion production threshold. Since this energy is that generally used by experiments as a threshold for the search for counterparts, the above may have important phenomenological consequences.  In particular, in the translucent or opaque regime, the closest object lying in the cosmic ray arrival direction should be a scattering center. Since these scattering centers are sites of intense magnetic activity (radio-galaxies, starburst galaxies, shock waves, ...), they might be mistaken with the source. This peculiar feature does not arise in models in which magnetic deflection is a continuous process in an all-pervading magnetic field.  One could thus conceive an ``ironic'' scenario, in which cosmic rays are accelerated in gamma-ray bursts, but scatter against radio-galaxies magnetized lobes, so that one interpret these latter as the source of cosmic rays because they are the only active objects seen on the line of sight. If such counterfeiting is taking place, one should observe that the apparent distance scale to the source (actually the distance to the last scattering surface) is smaller than the expected distance scale to the source (as determined by the energy losses). This offers a simple way to test for the above effect.  In the translucent regime, the source image is broadened by an angle $\\delta\\alpha$ which takes values of order of a degree at energy $10^{20}\\,$eV for the fiducial values of the scattering structures that we considered: interaction length $d_i\\,\\simeq\\, 30\\,$Mpc, extent $r_i\\,\\simeq\\,1\\,$Mpc, magnetic field $B\\,\\simeq\\, 10^{-8}\\,$G, and coherence length $\\lambda_i\\,\\simeq\\,0.1\\,$Mpc. The average optical depth at distance $100\\,$Mpc is thus $\\tau\\,\\simeq\\,3$ for these values. Due to the uncertainties surrounding these parameters, the deflection could however be larger or smaller by about an order of magnitude. In Section~\\ref{sec:angim}, we have discussed effects related to the shape of angular images when the discreteness of the scattering centers is taken into account.  The inhomogeneous distribution of matter in the local Universe implies that this optical depth to cosmic ray scattering should vary with the direction of observation. In Section~\\ref{sec:direcdep}, we have provided sky maps of the integrated baryonic matter density up to different distances, using the PSCz catalog of galaxies. These maps allow to estimate the fluctuation of the optical depth in different directions, hence that of the deflection angle, since $\\delta\\alpha\\,\\propto\\,\\tau^{1/2}$.  In our discussion, we have taken into account the inhomogeneous distributions of the scattering centers, see Sections~\\ref{sec:inhom}, \\ref{sec:inhom-num}. We have shown numerically that on path lengths longer than $\\,\\sim200\\,$Mpc, the effect of inhomogeneity is negligible, as expected for a Universe that is homogeneous and isotropic on these scales. The path length to the first interaction is generally higher by about 40\\% than in the homogeneous case if the scattering centers distribute according to the dark matter density. Since scattering centers tend to concentrate in filaments of large scale structure, a particle may also experience multiple interactions upon crossing a filament, as discussed and quantified in Section~\\ref{sec:inhom}. This explains why the number of interactions in the inhomogeneous case converges toward that of the homogeneous case on long path lengths.   \\subsection{Recent data from the Pierre Auger Observatory} In its first years of operation, the Pierre Auger Observatory has already achieved the largest aperture (in ${\\rm km}^2\\cdot{\\rm str}\\cdot{\\rm yr}$)~\\cite{Auger1}, and it has recently released the largest catalog of events above $5.7\\cdot10^{19}\\,$eV~\\cite{Auger2}. In this catalog, 20 out of 27 events originate from within 3 degrees of an active galactic nucleus located within 75~Mpc.  The most straightforward interpretation is to infer that active galactic nuclei are the sources of ultrahigh energy cosmic rays. However, only one of the observed counterparts is of the FR-I type (Centaurus~A), all others are more common Seyfert galaxies. From a theoretical point of view, this is unexpected, since these common active galactic nuclei do not seem to offer the required characteristics for the acceleration to ultrahigh energies~\\cite{NMA95}. Even Centaurus~A, as far as its jets are concerned, does not appear to be a likely source of ultrahigh energy cosmic rays~\\cite{CLP02}.  Furthermore, on a purely experimental level, Gorbunov and co-authors~\\cite{GTTT07} have recently pointed out an anomaly in this observed correlation. Assuming that the AGN seen in the arrival directions of these high energy events are the source of ultrahigh energy cosmic rays, these authors have computed the expected flux using the known distances to these AGN. They have observed that the Pierre Auger Observatory has collected zero event in the direction to the Virgo cluster, whereas at least six should be expected on the basis of the large concentration of AGN in this direction and the small distance scale (assuming that the cosmic rays coming from Centaurus~A indeed originate from this object).  Ref.~\\cite{GTTT07} thus argues that this observation rules out the possibility that AGN are the sources of ultrahigh energy cosmic rays, unless the cosmic rays seen in the direction to Centaurus~A come from further away. However, as pointed out to us during the refereeing process, it could also be that the absence of AGN-like source in Virgo is a statistical fluctuation due to the small number of sources in the local Universe, or that all AGN-like sources do not have the same cosmic ray luminosity. One may also ponder on the possibility that the Galactic magnetic field would exhibit a particular configuration in the direction to Virgo (which lies toward the Galactic North Pole), which would prevent cosmic rays from penetrating from this direction. Hence at present, one cannot exclude formally that AGN are the source of ultrahigh energy cosmic rays, but the data of the Pierre Auger Observatory cannot be argued to sustend this hypothesis strongly either.  Another interpretation suggests that sources of ultrahigh energy cosmic rays cluster with the large scale structure, as AGN do, hence the observed correlation with AGN is a coincidence. This hypothesis deserves to be more carefully studied, for instance by performing cross-correlations of the observed arrival directions with galaxy catalogs, or by following the method introduced in Ref.~\\cite{WFP96}. However, assuming that the sources are located close to the AGN which have been seen in the arrival directions should not resolve the flux anomaly noted in Ref.~\\cite{GTTT07}, also it might mitigate it somewhat.  A third interpretation is to assume that at least part of the observed correlation is accidental because the scattering centers on the last scattering surface cluster with the large scale structure, hence with AGN. This would alleviate this flux anomaly, since the sources would no longer have to be associated with the AGN distribution. In particular, the events seen to arise from the Centaurus complex might have been deflected in its vicinity.  As mentioned previously, this scenario can be tested by comparing the expected source distance scale with the counterpart distance scale. Interestingly, both do not match, as the source distance scale for particles with observed energy $6\\times 10^{19}\\,$eV is of the order of 200~Mpc, significantly larger than the maximum distance of 75~Mpc for the observed counterparts.  This fact has been noted in Ref.~\\cite{Auger2}; it remained mostly unexplained, although it was suggested in this work that both distance scales would agree if the energy scale were raised by 30\\%.  More quantitatively, one can calculate the probability that a given event with a given observed energy originates from a certain distance, using the fraction of the flux contributed by sources within a certain distance at a certain energy. This probability law can be calculated using the techniques developed in Ref.~\\cite{BGG02}, then tabulated. It is then possible to calculate the probability of seeing 20 out 27 events from a source located within 75~Mpc using the events energies reported in Ref.~\\cite{Auger2}. This probability is small, about 3\\%; the mean lies at 15 events out of 27 coming from within 75~Mpc. If one restricts the set of events to those that lie outside the Galactic plane ($\\vert b\\vert > 12^\\circ$), with 19 out of 21 seen to correlate, the probability becomes marginal, of order 0.1\\% (the mean lies at 12 out 21 within 75~Mpc). Finally, if one restricts oneself to the second set of events collected after May 27 2006, and on those which lie outside of the Galactic plane, with 9 out of 11 seen to correlate, the probability becomes of order 10\\%, with a mean at 7 out 11 within 75~Mpc. In this latter case, the signal is less significant, but the statistics is also smaller.  Since the above estimates do not take into account the uncertainty on the energy, and since they assume continuous instead of stochastic energy losses, these numbers should be taken with caution. Nonetheless, the above estimates agree with those of Ref.~\\cite{2006JCAP...11..012H}, which indicate that 50\\% of protons with energy $E>6\\times 10^{19}\\,$eV should come from distances less than $100\\,$Mpc and 90\\% from distances less than $200\\,$Mpc.  The above discussion suggests that, unless the energy scale is too low or an experimental artefact is present, the inferred distance scale to the source appears smaller than the expected source distance scale. In light of the analysis developed in the present paper, this suggests that part of the correlation may actually pinpoint scattering centers correlating with AGN rather than the source of ultrahigh energy cosmic rays. Said otherwise, the Pierre Auger Observatory may be seeing, at least partly, the last scattering surface of ultrahigh energy cosmic rays, rather than the source population.   In order to estimate the fraction of events that are likely to be contaminated by such pollution, one may proceed as follows. Assume first that the total deflection imparted to the particles with energy $>6\\times 10^{19}\\,$eV is less than the $3^\\circ$ radius used by the Pierre Auger Observatory for their search. One may then calculate the the fraction of galaxies in the PSCz catalog up to a distance $l=200\\,$Mpc, weighted appropriately, which lie within $3^\\circ$ of an AGN which is itself located closer than 75~Mpc. The distance $l=200\\,$Mpc is motivated by the fact that 90\\% of events with energy $>6\\times10^{19}\\,$eV originate from a distance smaller than 200~Mpc~\\cite{2006JCAP...11..012H}. One should weigh each galaxy with the selection function of the PSCz catalog at the distance $l$ of this galaxy in order to correct for the incompleteness of the catalog; one should also weigh each galaxy with a factor $1/l^2$ to account for flux dilution during propagation.  In the above estimate, the PSCz catalog is used as a tracer of the cosmic ray source population, and one simply calculates the probability of angular coincidence with the AGN sample. The number obtained is $0.31$, which suggests that 31\\% of events above $6\\times10^{19}\\,$eV could correlate with the AGN, assuming that the PSCz galaxies provide an unbiased tracer of the cosmic ray source population and that the magnetic deflection is much smaller than the search radius of $3^\\circ$. Note that this estimate does not take into account the effect of the magnetic field; if one were to restrict the angular radius to $2^\\circ$ in order to account for further possible Galactic deflection, the above fraction would become 25\\%. For reference, the probability that a random direction on the sky falls within $3^\\circ$ of an AGN (located closer than 75~Mpc) is $0.22$ (becoming $0.11$ for a radius of $2^\\circ$), which therefore gives the covering factor on the sky of these AGN.  In order to account for magnetic deflection, one may repeat the above procedure and calculate the probability of coincidence to within $3^\\circ$ of an AGN assuming that the event is displaced randomly by an angle $\\delta\\alpha$ from the location of the galaxy drawn from the PSCz catalog. Of course, one recovers the above result $0.31$ for $\\delta\\alpha\\rightarrow 0$, and the probability $0.22$ corresponding to isotropic source distribution for $\\delta\\alpha \\,\\sim\\, 1$ (in practice, $\\delta\\alpha\\,\\gtrsim\\,45^\\circ$ gives a probability $0.22$). Interestingly, the fraction of contaminated events increases as $\\delta\\alpha$ becomes of order of a few degrees: it equals 39\\% for $\\delta\\alpha=1^\\circ$, 48\\% for $\\delta\\alpha=3^\\circ$, then decreases, being 45\\% for $\\delta\\alpha=5^\\circ$ and 43\\% for $\\delta\\alpha=7^\\circ$, etc. If the radius of the correlation with the AGN is restricted to $2^\\circ$ to allow for further deflection in the Galactic magnetic field, these numbers become 21\\% for $\\delta\\alpha=1^\\circ$, 29\\% for $\\delta\\alpha=3^\\circ$ and 25\\% for $\\delta\\alpha=5^\\circ$.  The above estimates indicate that, within the assumptions of the above discussion, the delusion should not affect all events of the Pierre Auger Observatory, but a significant fraction nonetheless, possibly as high as $\\simeq 50\\,$\\%. Moreover, it also indicates that intergalactic magnetic deflection could be larger than $3^\\circ$ and yet produce a relatively significant false correlation with AGN. If further data from cosmic ray experiments strengthen the observed correlation, then the present interpretation would fail, unless some other effects artifically enhance this false correlation.  For instance, one should point out that the above fraction of contaminated events is likely to be enhanced if ultrahigh energy cosmic rays originate from gamma-ray bursts. Indeed, as discussed in Section~\\ref{sec:direcdep}, one expects in this case the number of events in regions of low foreground density to be smaller by a factor of order $\\tau$ ($\\tau$ being the optical depth measured in such directions) when compared to that coming from regions of optical depth greater than unity. The main reason is that a given source has a probability $\\sim\\,\\tau$ of being located behind a scattering center which would provide sufficient time delay for the source to become observable wih reasonable probability. On the contrary, a nearby gamma-ray burst with no scattering center on the line of sight has a negligible probability of being seen during a time span of a few years as a result of the small ocurrence rate. Although it is difficult to give a simple estimate of the magnitude of this effect on the amount of false correlations, one can easily see that it would tend to increase this fraction by providing more weight to regions of high foreground density (in which AGN are more numerous).    In Ref.~\\cite{Auger2}, the Pierre Auger Observatory has discussed the evolution of the probability of null hypothesis for an isotropic distribution of sources with a varying search radius, maximum AGN redshift and minimum energy (see Fig.~3 of Ref.~\\cite{Auger2}). The minimum probability (which indicates a maximal correlation with the AGN) corresponds to a search radius $3.2^\\circ$. This minimum can be interpreted as an estimate of the amount of Galactic and intergalactic magnetic deflection if one assumes that the source exactly correlates with the AGN. Interestingly, our above discussion suggests that this number may be a biased estimate and that the intergalactic deflection could be slightly larger. The increase of the probability of null hypothesis at larger search radii in the Pierre Auger data corresponds to the fact that the covering factor of the search area increase rapidly with search radius, being already 0.50 at $6^\\circ$. Concerning the redshift evolution, one would expect in the present model that the correlation would persist to distances as large as $200\\,$Mpc if the search radius is larger than the typical intergalactic deflection. Unfortunately, Ref.~\\cite{Auger2} does not plot this correlation beyond $100\\,$Mpc. It would be interesting to also carry out this test for different search radii.  One should emphasize that in the present interpretation, the correlation with AGN should not persist as the threshold energy is decreased. Indeed, the maximum propagation distance of particles of observed energy $4\\times10^{19}\\,$eV is of order $500\\,$Mpc, on which scale the Universe appears isotropic. Therefore, at these energies the incoming flux is increasingly isotropic, and the presence of scattering centers on the line of sight cannot induce anisotropies on an isotropic sky distribution (see discussion in Section~\\ref{sec:angim} as well as the discussion on the application of the Liouville theorem in Ref.~\\cite{1999JHEP...08..022H}).  The fraction of flux contributed by the isotropic background has been estimated in Ref.~\\cite{2006JCAP...01..009C} in the absence of extragalactic magnetic field; it reaches 83\\% for $E>3\\times 10^{19}\\,$eV, and 3.6\\% for $E>5\\times 10^{19}\\,$eV. The strong rise toward isotropy as the threshold energy decreases is thus clear. This effect is present, at least qualitatively, in the data of the Pierre Auger Observatory (see Figure 3 of Ref.~\\cite{Auger2}).  Finally, it appears that comparing the apparent source distance scale with the expected one, as we have done above, remains the most direct and simple test of the present interpretation.  Since there is a non-negligible degeneracy between the expected distance scale and the energy calibration, it seems mandatory to obtain a calibration through other methods that is as accurate as possible.  It also appears imperative to probe the arrival directions on an event by event basis, focussing on the most energetic events. In the catalog reported in Ref.~\\cite{Auger2}, there is only one event above $10^{20}\\,$eV, whose arrival direction has a relatively small super-Galactic latitude, $b_{\\rm SG}\\,\\simeq\\,-6.5^\\circ$. In the above scenario, one should expect to find a scattering center or the source on the line of sight, hence it should prove useful to perform a deep search in this direction in the radio domain, looking for traces of synchrotron emission that would attest of the presence of a locally enhanced intergalactic magnetic field. Many more events at higher energies, as expected from future detectors such as Auger North~\\footnote{{\\tt http://www.augernorth.org}}, would certainly help in this regard.  A last word should be added concerning the amount of magnetic deflection and the source models. In particular, it would be interesting to examine whether (and to what cost) the current data could be reconciled with ultrahigh energy cosmic rays being accelerated in the most powerful AGN, which offer stronger ground than Seyfert galaxies for acceleration. Such a study can only be conducted through detailed Monte-Carlo simulations which allow for substantial scattering angles in inhomogeneous magnetic fields.  As explained in Ref.~\\cite{W01}, gamma-ray bursts are probably the most elusive of possible ultrahigh energy cosmic ray sources, as the strongest predictions are that no counterpart should be detected, that the flux should show significant variations around the mean at sufficiently high energies (a few $10^{20}\\,$eV), and that multiplets of events should be clustered in energy. Current data do not violate any of these predictions, but it is clear that experiments with much larger aperture at the highest energies will be needed to test such effects.  It is certain that much physics and astrophysics of cosmic ray sources and large scale magnetic fields remain to be unveiled by ongoing and future detectors."
    },
    "0801/0801.0510_arXiv.txt": {
        "abstract": "{ The chemistry of molecular clouds has been studied for decades, with an increasingly general and sophisticated treatment of the reactions involved. Yet  the treatment of turbulent diffusion has remained extremely sketchy, assuming simple Fickian diffusion with a scalar diffusivity $D$. However, turbulent flows similar to those in the interstellar medium are known to give rise to anomalous diffusion phenomena, more specifically superdiffusion (increase of the diffusivity with the spatial scales involved). This paper considers to what extent and in what sense superdiffusion modifies molecular abundances in interstellar clouds. For this first exploration of the subject we employ a very rough treatment of the chemistry and the effect of non-unifom cloud density on the diffusion equation is also treated in a simplified way. The results nevertheless clearly demonstrate that the effect of superdiffusion is quite significant, abundance values at a given radius being modified by order of unity factors. The sense and character of this influence is highly nontrivial.} ",
        "introduction": "The chemistry of molecular clouds is a complex system affected by different factors.  The first models laid down the foundations of the chemical reaction network: following the initial steady state gas-phase chemistry model of Herbst \\& Klemperer (1973) several (pseudo-)time-dependent models, with fixed profiles of physical parameters were developed, considering more and more chemical reactions, while neglecting photodissociation (Leung et al.~1984; Herbst \\& Leung 1989; Millar \\& Herbst 1990). The agreement of the mo\\-dels with observed fractional abundances improved when the effect of turbulent diffusion was taken into account, resulting in smoother density distribution profiles for the more important species (Xie et al.~1995, hereafter XAL95). Subsequent models further refined the chemistry, considering details of the ion-neutral reaction scheme, and taking into account the gas-grains interaction, $\\mathrm{H_2}$ and CO self-shield\\-ing, grain accretion effects and adsorption onto grains (Will\\-acy et al.~2002; Yate \\& Millar 2003).  In contrast to the enormous efforts made in order to improve the representation of chemical processes, the treatment of turbulent diffusion has remained extremely sketchy, assuming simple Fickian diffusion with a scalar diffusivity $D$. Under the conditions prevailing in interstellar space this is surely wrong. Interstellar turbulence is mainly driven by supernova shocks on spatial scales far exceeding that of the molecular clouds. From those scales, kinetic energy cascades down to the scale of clouds, cloud clumps and even cloud cores (Boldyrev 2002; Scalo \\& Elmegreen 2004; Ryan Joung et al.~2006). Indeed, in current thinking clouds are hardly more than positive density fluctuations in this large-scale compressible turbulent flow. In such a turbulent flow, the increase of the r.m.s. separation $\\Delta$ of two tracer particles will only be due to eddies smaller than $\\Delta$, leading to a {\\it scale-dependent eddy diffusivity.} As a result, the increase of $\\Delta$ with time will deviate from the Fickian square-root law: such {\\it non-Fickian} or {\\it anomalous} diffusive pro\\-cess\\-es commonly occur in various fields of physics (Avellaneda \\& Majda 1992), though astrophysical applications have so far been limited to solar physics (Petrovay 1999). As in the turbulent cascade the diffusivity increases for larger spatial scales, the transport of passive scalars by inertial-range turbulence belongs to the subclass of {\\it superdiffusive} processes.  The work presented here is a first exploration of the effect of superdiffusion on the chemistry of molecular clouds. Our approach is described in Section~2. In order to get a first impression of the importance of superdiffusive effects, we neglect the intricacies of the chemical reaction network and treat their net effect as a simple parametric source term in the diffusion equation. As anomalous diffusivity cannot be described by any partial differential equation in physical space, we are forced to work in Fourier space, with $D=D(k)$ a function of the wavenumber $k$. To alleviate the ensuing difficulties, in Section~3.1 we first consider a case where the density of the main component is uniform throughout the cloud. A semianalytical solution is possible in this case; the results indicate that there are significant and nontrivial differences between the diffusive and superdiffusive cases.  Section~3.2 then considers what changes are to be expected if the assumption of a uniform cloud is dropped. Section~4 concludes the paper. ",
        "conclusions": "In this paper we considered to what extent and in what sense superdiffusion modifies molecular abundances in interstellar clouds. The results clearly demonstrate that the effect is quite significant, abundance values at a given radius being modified by order of unity factors. The sense and character of this influence is highly nontrivial.  In order to make a meaningful comparison with observational data possible, this model needs to be developed further to include a realistic treatment of the chemical reactions and to include a more rigorous calculation of the effect of non-uniformity of the cloud in the diffusion equation. Work in this direction is in progress."
    },
    "0801/0801.3275_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\subsection{About \\lzebra}  \\zebra{} estimates redshifts and template types of galaxies using a photometric catalog. It is based on: \\\\ $(1.)$ An automatic iterative technique to correct the original set of galaxy templates to best represent the SEDs of real galaxies at different redshifts;  \\\\ $(2.)$ A training set of spectroscopic redshifts for a small fraction of the photometric sample to improve the robustness of the photometric redshift estimates; and \\\\ $(3.)$ An iterative technique for  Bayesian redshift estimates, which extracts  the full two-dimensional  redshift and template likelihood function for each galaxy.  \\paragraph{Input}  The user can choose from, or has otherwise to provide, an initial set of templates and filter transmission curves, see Table \\ref{tab:providedFilters} and Table \\ref{tab:providedTemplates}. In addition, \\zebra{} needs a catalog containing the magnitudes of the individual galaxies in the various filter bands (at least 3 filter bands are needed to run \\zebra{} in {\\it Maximum-Likelihood mode} or {\\it photometry-check mode}). The different input files are described in detail in section \\ref{sect:BasicFiles}.  \\paragraph{Output}  \\zebra{} offers a variety of output information depending on its run mode.  \\begin{description} \\item \\zebra{} prints on-screen information during its execution. The level of verbosity of this output can be adjusted with the \\texttt{-v}/\\texttt{--verbosity} option.  \\item When run in {\\it photometry-check mode} \\zebra{} returns a catalog with calibrated photometry together with detailed information about the applied changes.  \\item In {\\it template-optimization mode} \\zebra{} returns the corrected templates as wavelength -- flux density (per unit wavelength) tables.  \\item If \\zebra{} is run in {\\it Maximum-Likelihood mode}, it will return the best fit redshift and template type together with their confidence limits estimated from constant $\\chi^2$ boundaries. Additionally, \\zebra{} returns: (i) the minimum $\\chi^2$, (ii) the normalization factor of the best fit template, (iii) the rest-frame B-band magnitude and (iv) the luminosity distance for the given cosmological Friedmann-Robertson-Walker model. If specified by the user, \\zebra{} provides further information, e.g. the likelihood functions for all galaxies in several output formats or the residuals between best fit template magnitude and measured magnitude for each galaxy in each filter band.  \\item In the {\\it Bayesian mode} \\zebra{} calculates the 2D-prior in redshift and template space in an iterative fashion. This prior (and, if specified, the interim prior of each iteration step) is returned. The final prior is used by \\zebra{} to compute a posterior for each galaxy. The posterior can be saved as full 2D-table or in marginalized form. \\zebra{}'s output includes the most probable redshift and template type for each galaxy as defined by (i) the maximum of the posterior or (ii) after marginalizing over templates types and redshifts, respectively. The errors are calculated directly from the posterior.  \\item \\zebra{} can also be employed to derive template-based k-corrections using the specified templates and filters.  \\end{description} All input and output files of \\zebra{} are ASCII-files.  \\subsection{The software license} This program is free software; you can redistribute it and/or modify it under the terms of the GNU General Public License (GPL) as published by the Free Software Foundation; either version 3 of the License, or (at your option) any later version.  This program is distributed in the hope that it will be useful, but WITHOUT ANY WARRANTY; without even the implied warranty of MERCHANTABILITY or FITNESS FOR A PARTICULAR PURPOSE.  See the GNU General Public License 3 for more details.  A copy of the GPL 3 is provided at the end of this manuscript.  \\subsection{Usage and questions} It is kindly requested that the use of \\zebra{} should be acknowledged with an explicit reference to \\cite{feldmann06} in the bibliographic list of any resulting publication.  Any problems, comments and suggestions should be sent via e-mail to zebra@phys.ethz.ch.  \\clearpage ",
        "conclusions": ""
    },
    "0801/0801.4455_arXiv.txt": {
        "abstract": "A variety of magnetohydrodynamic mechanisms that may play a role in magnetic, chemically peculiar (mCP) stars is reviewed. These involve dynamo mechanisms in laminar flows as well as turbulent environments, and magnetic instabilities of poloidal and toroidal fields as well as combinations of the two. While the proto-stellar phase makes the survival of primordial fields difficult, the variety of magnetic field configurations on mCP stars may be an indication for that they are instability remnants, but there is no process which is clearly superior in explaining the strong fields. ",
        "introduction": "A considerable fraction of chemically peculiar A and B stars (CP stars) show strong surface magnetic fields of about 500~G up to over 10~kG. The evolution of these fields from star formation to their current presence is unknown. This paper compiles a number of physical processes which may or may not play a role in the whole scenario of CP star magnetism.  The fraction of main-sequence stars with radiative envelopes which show magnetic fields was found to increase from F stars to late B stars (Wolff 1968). The result was recently confirmed by Power et al. (2007) who found an increase of mCP star fractions with stellar mass up to $3.6M_\\odot$ for nearby stars. A further increase towards higher masses was hidden in the small number of higher-mass stars. While the average fraction of mCP stars among normal stars of the same mass range is generally quoted as being roughly 10\\%, the fraction is less than 2\\% in the solar neighbourhood of 100~pc radius (Power et al. 2007). Why only a fraction of intermediate-mass stars shows strong magnetic fields is one of the key issues in the search for their origin. The other is that mCP stars form a distinct sub-group with slow rotation among A and B stars.  There are indications of a decrease of magnetic field strength with stellar age for $M > 3M_\\odot$ (Kochukhov \\& Bagnulo 2006; Landstreet et al. 2007). Although the increase of stellar radius may cause the reduction at constant magnetic flux, it appears that even the flux is decreasing during the main-sequence life of mCP stars. At first glance this decay may be attributed to ohmic decay in the radiative envelope of the star. The radial profile of the magnetic diffusivity $\\eta(r)$ throughout the star according to Spitzer (1956) suggests that the decay time is of the order of $10^{11}$~yr everywhere in the radiation zone. The decay time inferred from Landstreet et al. (2007) is about $10^7$ to $10^8$~yr (as is $\\eta$ in cm$^2$/s accidentally). Figure~\\ref{decay_times} shows $\\eta$ as a function of radius and the corresponding ohmic diffusion time $\\tau_{\\rm diff}=r^2/\\eta$. The above decay time deduced from observations is also indicated -- without radial dependence for obvious reasons. It appears as if the magnetic diffusivity is higher than the purely microscopic value hinting on very mild turbulence even without convection.  \\begin{figure}[t] \\centerline{\\includegraphics[width=9cm,clip=]{decay_time.eps}} \\caption{Magnetic diffusivities (lower part of diagram) and the corresponding decay times for magnetic structures (upper part of diagram). The diffusivity depends on the location in the star at radius $r$ and so does the decay time of a structure as large as $r$. The ``small-scale'' curve refers to structures with a scale of $0.1r$ which decay 100~times faster. A rough estimate of the magnetic diffusivity according to the age dependence of magnetic field strengths of CP stars is given in middle.} \\label{decay_times} \\end{figure}   Another puzzling fact is the decrease of obliquity of the magnetic field  axis against the rotation axis with rotation period (Landstreet \\& Mathys 2000; Bagnulo et al. 2002). MHD processes tend to prefer axisymmetric solutions for fast rotators, if a small degree of differential rotation is present in the star. This is the very opposite of what is observed. When comparing the rotational time-scale with the diffusive time-scale, even the slowest CP stars have rotation periods many orders of magnitude shorter than $\\tau_{\\rm diff}$ and are fast rotators in that sense. Since in view of the time-scales normal A stars and Ap stars are not vastly different, we tend to place the magnetic fields first -- the reduced rotation is a consequence of their presence. In other words, the reason for mCP stars to be magnetic is not their initially slow rotation.  There is not only the gradual decrease of magnetic field strengths with the time which the $M>3M_\\odot$ stars spent on the main sequence, but also a possible late emergence of surface fields for stars with $M<2M_\\odot$. Hubrig et al. (2000) found that these stars show magnetic fields only after they have spent about 30\\% of their life on the main sequence. Meanwhile, counter-examples were found, and the appearance of fields is certainly not as sharp as it seemed at an earlier research stage.  It should be noted that in general all these characteristics are based on quite scattered individual stars. The scatter is very likely not a result of observing uncertainties. Any physical process leading to the observed features must also allow for a certain bandwidth of results or sufficient input of stochastic behaviour. ",
        "conclusions": "A compilation of a few possible mechanism in CP star magnetism with their pros and contras are listed in Table~\\ref{summary}. When searching for mechanisms explaining considerable magnetic fields in radiative envelopes, one always returns to the question why only a small fraction of A stars shows such fields. The effect of different star formation environments on the existence of strong fields on main-sequence stars now needs to be investigated. Bridging collapse and pre-main-sequence evolution is not a trivial task though.   \\begin{table} \\caption{Four scenarios of CP star magnetism.\\label{summary}} \\begin{tabular}{|l|l|} \\hline {\\bf Background field amplified at}         & {\\bf Dynamo in radiative zone}\\\\ {\\bf star formation}                        & $+$ straight-forward mechanism \\\\ $+$ random orientation                      & $-$ $m=1$ modes cause regular obliquity \\\\ $+$ field strength decreases with age       & $-$ field strength does not decreases\\\\ $-$ convective phase or MRI produce         & \\hspace{4mm}with age\\\\ \\hspace{4mm}small-scale field               & $-$ rather independent of star formation\\\\ $-$ contradicts possible late emer-         & \\hspace{4mm}$\\rightarrow$ no distinction between A and Ap\\\\ \\hspace{4mm}gence of field                  & $-$ dynamo difficult to excite\\\\ \\hline {\\bf Early-life dynamo remnants}            & {\\bf Early-life dynamo + field instab.}\\\\ $+$ field strength decreases with age       & $+$ random orientations at emergence\\\\ $-$ fully convective phase questionable     & $+$ field strength decreases slowly \\\\ $-$ non-axisymmetry only very shortly       & \\hspace{4mm}with age\\\\ \\hline \\end{tabular} \\end{table}"
    },
    "0801/0801.3569_arXiv.txt": {
        "abstract": "{Young stars are born with envelopes, which in the early stages obscure the central (proto)star and circumstellar disk. In the Class I stage, the disks are still young, but the envelopes are largely dispersed. This makes the Class I sources ideal targets for studies of the early stages of disks.} {We aim to determine the masses of the envelopes, disks, and central stars of young stellar objects (YSOs) in the Class I stage.} {We observed the embedded Class I objects IRS~63 and Elias~29 in the $\\rho$~Ophiuchi star-forming region with the Submillimeter Array (SMA) at 1.1~mm.} {IRS~63 and Elias~29 are both clearly detected in the continuum, with peak fluxes of 459 and 47~mJy/beam, respectively. The continuum emission toward Elias~29 is clearly resolved, whereas IRS~63 is consistent with a point source down to a scale of 3$\\arcsec$ (400~AU). The SMA data are combined with single-dish data, and both disk masses of 0.055 and $\\le 0.007$~M$_\\odot$ and envelope masses of 0.058 and $\\le 0.058$~M$_\\odot$ are empirically determined for IRS~63 and Elias~29, respectively. The disk+envelope systems are modelled with the axisymmetric radiative-transfer code RADMC, yielding disk and envelope masses that differ from the empirical results by factors of a few. HCO$^+$ $J$ = 3--2 is detected toward both sources, HCN $J$ = 3--2 is not. The HCO$^+$ position-velocity diagrams are indicative of Keplerian rotation and allow an estimate of the mass of the central stars. For a fiducial inclination of 30$^\\circ$, we find stellar masses of $0.37 \\pm 0.13$ for IRS~63 and $2.5 \\pm 0.6$~M$_\\odot$ for Elias~29.} {The sensitivity and spatial resolution of the SMA at 1.1~mm allow a good separation of the disks around Class~I YSOs from their circumstellar envelopes and environments, and the spectral resolution makes it possible to resolve their dynamical structure and estimate the masses of the central stars. The ratios of the envelope and disk masses $M_{\\rm env}/M_{\\rm disk}$ are found to be 0.2 for IRS~63 and 6 for Elias~29. This is lower than the values for Class~0 sources, which have $M_{\\rm env}/M_{\\rm disk} \\ge 10$, suggesting that this ratio is a tracer of the evolutionary stage of a YSO.} ",
        "introduction": "\\label{sect: introduction}  Low- and intermediate-mass stars are formed from the graviational collapse of molecular cloud cores. In the earliest stages, the newly-formed protostar remains embedded in the remnants of this core, a cold envelope of dust and gas, which is gradually accreted by the young star (e.g., Shu 1977). Due to the angular momentum initially present in the core, most of the envelope material does not fall directly onto the central protostar but is piled-up in a circumstellar disk (e.g., Terebey et al. 1984). Understanding this interplay between star, disk, and envelope is crucial in order to be able to relate the initial conditions of star formation such as the mass of the protostellar core to the end-product -- namely the properties of the young star, and the mass and thus potential of the disk for forming planets. Some of the key questions include: Where does most of the mass reside at a given time? Will all the mass that is seen in prestellar cores or in the envelopes around deeply embedded sources end up in the star, or will a large fraction be dispersed from the system? How do the masses of the circumstellar envelope, the disk, and that of the central star evolve over time, and how long does it take for the circumstellar matter to be accreted onto the star?  Young stellar objects (YSOs) are usually classified according to their slopes in the infrared (IR) wavelength regime. Originally, the LW classification \\citep{lada:1984, lada:1987} ran from the embedded Class I, via the optically visible Class II or classical T Tauri stars, to the Class III spectral energy distributions of post-T Tauri stars. Later, the Class 0 stage was added to this classification \\citep[see, e.g.,][]{andre:1993}, where the Class 0 sources are distinguished from the Class I sources through their high relative luminosity at submillimetre (submm) wavelengths. This classification roughly reflects the evolutionary stage of the YSOs under consideration. The most deeply embedded Class 0 sources are thought to evolve through the Class I stage while dissipating their circumstellar envelopes. Eventually they become optically visible as pre-main sequence T Tauri stars with circumstellar disks.  The Class 0 and Class II YSOs have been studied quite extensively with high-resolution (sub)millimetre interferometers \\citep[e.g.,][and references therein]{jorgensen:2007, andrews:2007}. Studies of deeply embedded Class~0 YSOs have shown that circumstellar disks are\tformed early \\citep{harvey:2003, jorgensen:2004}, but in these systems, it is difficult to separate the emission from the disk from that of the envelope. The Class~I sources, in which a large part of the original circumstellar envelopes has been dissipated, are ideal to study young disks in YSOs. Interferometric studies of Class I objects have so far been largely limited to studies at around 3~mm \\citep[e.g.,][]{ohashi:1997, hogerheijde:1997a, hogerheijde:1998, looney:2000}. The Submillimeter Array (SMA) allows observations around 1~mm, where the thermal dust continuum emission is an order of magnitude stronger than at 3~mm. Also at these shorter wavelengths it is possible to detect the higher rotational transitions of the molecules that trace the dense gas in the disks and inner envelopes of the young systems, rather than lower-density extended envelope emission.  We here present SMA observations of two Class I objects that appear to be in an evolutionary stage where the kinematics of the circumstellar material are no longer dominated by infall. \\object{IRS 63} (\\object{WLY 2-63}, \\object{GWAYL 4}) and \\object{Elias 29} (\\object{Elia 2-29}, \\object{WLY 1-7}) are located in the $\\rho$~Ophiuchi cloud, taken to be at a distance D~$= 125 \\pm 25$~pc \\citep{degeus:1989}. Spitzer photometry from the ``Cores to Disks'' Legacy program \\citep{evans:2003} was used to determine the source's infrared colours and their bolometric luminosities and temperatures. IRS~63 and Elias~29 have similar bolometric temperatures, $T_{\\rm bol} = 351$~K and 391~K, respectively, which places them in the LW Class~I regime. The bolometric luminosities were calculated to be $L_{\\rm bol} = 0.79$~L$_\\odot$ for IRS~63 and 13.6~L$_\\odot$ for Elias~29. According to their infrared slopes of $\\alpha_{2-24 \\mu m} = 0.15$ and 0.42, IRS~63 is the slightly more evolved, falling into the ``flat-spectrum'' class that separates the Class~I from the Class~II sources \\citep{greene:1994}. The values for $\\alpha_{\\rm IR}$ quoted here are from un-dereddened observations. Correction for the extinction toward the $\\rho$~Ophiuchi star-forming region \\citep[e.g.,][]{flaherty:2007} would result in lower values for $\\alpha_{\\rm IR}$; in other words, the sources may be slightly more evolved than the raw values of $\\alpha_{\\rm IR}$ suggest. In terms of environment, the two sources are quite each other's opposites. The SCUBA 850~$\\mu$m map \\citep{johnstone:2004} from the COMPLETE survey \\citep{ridge:2006} shows that IRS~63 is an isolated compact source. Elias~29, on the other hand, is located in a dense ridge of molecular material, which contains several more YSOs. It is likely that these YSOs were formed from condensations in this molecular ridge.  We present the masses of all main components of Class~I YSOs, i.e., the central star, the disk, and the envelope, for the first time. IRS~63 and Elias~29 are the first two sources in a larger survey of Class~I objects studied with the SMA at 1.1~mm. These observations complement the survey of Class~0 sources in the PROSAC programme \\citep{jorgensen:2007} and will allow us to trace the similarities and differences of these evolving protostars. The results of the complete campaign will be presented in a future paper. In \\S\\ref{sect: observations}, the observations are presented, and the results and implications are discussed in \\S\\S\\ref{sect: results} and \\ref{sect: discussion}. We summarise the main conclusions in \\S\\ref{sect: conclusions}. ",
        "conclusions": "\\label{sect: conclusions} We used the SMA to observe the Class I YSOs Elias~29 and IRS~63 at 1.1~mm. The main results are as follows. \\begin{itemize} \\item Both sources are detected in the continuum and in the HCO$^+$ $J$ = 3--2 line, but not in the HCN $J$ = 3--2 line. The HCO$^+$ emission toward IRS~63 is compact and associated with the YSO, whereas a significant part of the molecular emission in the direction of Elias~29 is due to an enhancement in the ridge from which Elias~29 formed. \\item Assuming that the continuum emission toward the sources is optically thin, and that any circumstellar envelope is resolved out by the interferometer on the longest baselines, we find empirical disk masses of 0.055 for IRS~63 and $\\le 0.007$~M$_\\odot$ for Elias~29. By comparing our interferometer data with single-dish observations, envelope masses of 0.058~M$_\\odot$ are derived for both sources, within 30$\\arcsec$ radius for IRS~63 and 20$\\arcsec$ radius for Elias~29, where the value for Elias~29 should be treated as an upper limit because of the presence of the ridge enhancement. \\item Using a radiative-transfer code to create SEDs and amplitude vs. ($u, v$)-distance plots for different combinations of disk and envelope masses, we find disk masses of 0.13~M$_\\odot$ for IRS~63 and 0.004~M$_\\odot$ for Elias~29, and envelope masses of 0.022~M$_\\odot$ and 0.025~M$_\\odot$. \\item The modelling results yield rather different ratios of M$_{\\rm env}$/M$_{\\rm disk}$: 0.2 for IRS~63 and 6 for Elias~29, with uncertainties of a factor of a few. From this it is concluded that IRS~63 is well on its way to become an optically visible Class~II source, rather than an embedded Class~I source. Elias~29, on the other hand, is more like the deeply embedded Class~0 sources. \\item Velocity gradients in the HCO$^+$ $J$ = 3--2 emission are interpreted as signs of Keplerian rotation, and indicate central masses of $2.5 \\pm 0.6$~M$_\\odot$ for Elias~29 and $0.37 \\pm 0.13$~M$_\\odot$ for IRS~63, for inclinations of 30$^\\circ$. These masses correspond to ages of a few $10^5$~yr. \\end{itemize}  With a larger set of sources treated in a similar manner as is presented for IRS~63 and Elias~29 here, it will be possible to constrain models for the evolution of the envelope, disk, and stellar masses through this important stage of YSOs."
    },
    "0801/0801.0871_arXiv.txt": {
        "abstract": "Fitting whole spectra at intermediate spectral resolution (R = 1000 -- 3000), to derive physical properties of stellar populations, appears as an optimized alternative to methods based on spectrophotometric indices: it uses all the redundant information contained in the signal. This paper addresses the validation of the method and it investigates the quality of the population models together with the reliability of the fitting procedures. Our method compares observed optical spectra with models to derive the age, metallicity, and line broadening due to the internal kinematics. It is insensitive to the shape of the continuum and the results are consistent with Lick indices but three times more precise. We are using two algorithms: STECKMAP, a non-parametric regularized program and NBURSTS a parametric non-linear minimization.  We compare three spectral synthesis models for single stellar populations: \\phr, Galaxev (BC03) and \\vaz, and we analyse spectra of Galactic clusters whose populations are known from studies of color-magnitude diagrams (CMD) and spectroscopy of individual stars.  We find that: (1) The quality of the models critically depends on the stellar library they use, and in particular on its coverage in age, metallicity and surface gravity of the stars. \\phr\\ and \\vaz\\ are consistent, while the comparison between \\phr\\ and BC03 shows some systematics reflecting the limitations of the stellar library (STELIB) used to generate the latter models; (2) The two fitting programs are consistent; (3) For globular clusters and M67 spectra, the method restitutes metallicities in agreement with spectroscopy of stars with a precision of 0.14~dex; (4) The spectroscopic ages are very sensitive to the presence of a blue horizontal branch (BHB) or of blue stragglers. A BHB morphology results in a young SSP-equivalent age. Fitting a free amount of blue stars in addition to the SSP model to mimic the BHB improves and stabilizes the fit and restores ages in agreement with CMDs studies. This method is potentially able to disentangle age or BHB effects in extragalactic clusters. Full spectrum fitting is a reliable method to derive the parameters of a stellar population. ",
        "introduction": "The goal of this paper is to validate a method for determining the characteristics of a stellar population using an entire medium resolution spectrum in the rest-frame optical range (rather than concentrating on specific features). Special attention is given to make an optimal use of the information, to look for systematics and to compare the different models of stellar populations.  Stellar populations keep a record of most of the evolution of the stellar systems, and by studying Hertzprung-Russell (HR) diagrams (or colour-magnitude diagrams, CMDs) of  galaxies, it is possible to constrain the history of the star formation and enrichment since their formation \\citep[see eg.][]{grebel05}. Unfortunately, these star-resolved observations are restricted to nearby galaxies (Local Group and a little further) because of the shallow magnitude limits and crowding effect. At large distances, only the spectral energy distributions (SED) and the line-of-sight (LOS) integrated spectra are available.  Many methods were devised to determine the age and the metallicity of stellar systems from their integrated light. In this paper we distinguish three approaches: \\begin{enumerate} \\item SED fitting  - sensitive to the general shape of the continuum, (usually) made in broad wavelength ranges; \\item Spectrophotometric indices - concentrates on the strength or equivalent width of specific features, not sensitive to the flux calibration; \\item Full spectrum fitting - use all the information by fitting all the pixels, but independently from the shape of the continuum; \\end{enumerate}  SED fitting  \\citep{cid05,pan,sol05}  can be made using colours, low or medium resolution spectra. It is performed over large wavelength ranges and compares the observations to models like {\\sc P\\'egase}.2 \\citep{pegase,pegase2},~{\\sc Galaxev} \\citep{BC03} or others. These models should take into account the various sources of extinction (internal and Galactic). SED fittig constrains not only the stellar population but it also informs about the dust content.  Using low- and medium resolution spectra, the classical approach is to measure spectrophotometric indices, like Lick \\citep{faber73,wor94} or \\citet{rose84} indices and compare them to model predictions. Spectrophotometric indices measure the strength of some spectral features expected to trace the abundance of particular elements or to correlate with the mean age of stellar populations. However, primarily because a galaxy spectrum is naturally broadened by the internal kinematics, the spectrophotometric indices are defined at low resolution (a few {\\AA}) and are therefore blends of many spectral lines. That is why their responses to age and metallicity are quite complex and require careful calibrations \\citep{K05}. Attempts to define indices at a higher resolution \\citep[see][]{PHR} require delicate velocity-dispersion corrections or must be designed carefully to reduce their sensitivity to the velocity dispersion. A good example of high-resolution index is the H$\\gamma$ index defined in \\citet{va99}. Another possibility to improve the sensitivity of an index to either age or metallicity is to use a composite definition, like for the CaT$*$ (calcium triplet corrected for Paschen lines, \\citealt{cenarro01}) or MgFe (insensitive to [$\\alpha$/Fe], \\citealt{G93,tmb03}) indices.  Lick indices are widely used for more than a decade and the methods continue to improve with introducing the effects of non-solar abundances. However, they use only a few small passbands out of the full observed spectral range, and although the information in the spectrum is probably redundant, this does not optimize the usage of the collected light. Another limitation of the indices is that they are sensitive to missing data or bad pixels \\citep{Chil07}: The interpolation of spectral bins affected by detector defects or cosmic hits may bias the measurement of the index if the observations are close to the sampling limit. The indices are also affected by weak emission lines \\citep{GE96}.  The full spectrum fitting that we propose combines the advantages of spectrophotometric indices and SED fitting: It uses all the information, it is insensitive to extinction or flux calibrating errors and it is not limited by the physical broadening, since the internal kinematics is determined simultaneously with the population parameters. >From the point of view of measuring the internal kinematics, it is an extension of the classical methods (eg. optimal template fitting, \\citealt{vdmarel94}).  The method has been presented in several places \\citep{pru2003,ocv06a,ocv06b} and its reliability assessed on the basis of various simulations. It is three times more sensitive, but still consistent with Lick indices \\citep{Chil07,mich07}.  Full spectrum fitting is well suited when the resolution is comparable with the physical broadening (10~km\\,s$^{-1} <  \\sigma_{ins} < $ 200~km\\,s$^{-1}$ ). At lower resolution (R = 500 -- 1000) the spectral features are diluted and the sensitivity of the method would rely on the strong spectral features and therefore would become equivalent to Lick indices. At even lower resolution, only strong features, like the Balmer break provides information on the stellar population. In this case, SED fitting using also the shape of the continuum is the good solution.  In its intention to be insensitive to the flux calibration and extinction our full spectrum fitting is close to spectrophotometric indices, but it uses all the spectrum. A possible disavantage compared to Lick indices is that the method may, in principle, be sensitive to the wavelength range and to spectral resolution. However, \\citet{kol06} have shown that this sensitivity is not critical: the precision of the results mostly depends on the total S/N (i.e. integrated over all the spectral range), and even excluding the Balmer lines does not bias the age determination  (only precision is reduced).  To check the reliability of a spectroscopically determined SSP equivalent age/metallicity, it is wise to test it on globular clusters (GC) and to compare the results with the results from CMDs determinations. It has been shown \\citep{wolf07} that the full spectrum fitting reproduces better the results from the CMDs than the other methods.  This paper is organised as follows: In Sect. 2 we present the population models and the two algorithms implementing the full spectrum fitting. In Sect. 3, we use these algorithms to compare different synthetic population models and in Sect. 4 we invert spectra of Galactic clusters and compare the results with age determinations from HR diagrams. Section 5 discusses the prospects for analyzing stellar populations and gives conclusions. ",
        "conclusions": ""
    },
    "0801/0801.0040_arXiv.txt": {
        "abstract": " ",
        "introduction": "As modern cosmological observations become more and more precise, study of the non-Gaussianity of CMB temperature fluctuations has become a more and more pressing issue in recent years. In the WMAP convention \\cite{Komatsu:2001rj,Spergel:2006hy}, the primordial non-Gaussianity is parameterized by $f_{NL}$ assuming the ansatz \\begin{equation}\\label{wmap} \\zeta=\\zeta_L+\\frac{3}{5}f_{NL}\\zeta_L^2, \\end{equation} where $\\zeta$ is the scalar perturbation and $\\zeta_L$ is its linear Gaussian part\\footnote{Note that Maldacena's convention is also popularly used in the literature \\cite{Maldacena:2002vr,Chen:2006nt}, which assumes $\\zeta=\\zeta_L-\\frac{3}{5}f_{NL}\\zeta_L^2$ instead of (\\ref{wmap}). That convention is different in sign of $f_{NL}$ from the WMAP convention. Please refer to \\cite{Komatsu:2002db,Komatsu:kitpc} for clarification. In this paper, we will use the WMAP convention.}. However, this ansatz only corresponds to a restricted shape of non-Gaussianity. Theoretically, during inflation, the non-Gaussianity is usually produced in different shapes, and the estimator $f_{NL}$ can be defined  in terms of the 3-point function $\\langle\\zeta_{k_1}\\zeta_{k_2}\\zeta_{k_3}\\rangle$. Two limits of $f_{NL}$ are of most interest. One is the local, squeezed limit $k_1\\ll k_2\\simeq k_3$, for which we will use the notation $f_{NL}^{local}$. The other is the non-local, equilateral limit $k_1\\simeq k_2\\simeq k_3$, for which the notation will be $f_{NL}^{equil}$.  In a recent work \\cite{Yadav:2007yy}, utilizing the fast estimator of primordial non-Gaussianity \\cite{Yadav:2007ny}, Yadav and Wandelt claimed that data from two channels of WMAP3 reject $f_{NL}^{local}=0$ at the $2.89\\sigma$ level, or $99.6\\%$ significance. They also showed that $26.91<f_{NL}^{local}<146.71$ at $95\\%$ C.L., with a central value of $f_{NL}^{local}=86.8$. If this result is confirmed by future observations, it will have a great impact on our study of the early universe, because a large class of inflation models will be ruled out. For example, the simplest model of inflation is slow rolling, driven by a single scalar field. Ignoring the non-Gaussianity, previous observational data fit well with single-field slow-roll inflation \\cite{Spergel:2006hy,Alabidi:2006qa}. On the other hand, in the conventional single-field slow-roll inflation, it has been found for both local and equilateral forms that $|f_{NL}|<1$, which is too small to detect in the near future \\cite{Maldacena:2002vr,Acquaviva:2002ud}. Therefore, the confirmation of a large $f_{NL}^{local}$ observationally will exclude almost all models of single-field slow-roll inflation.  Confronted with the evidence for $f_{NL}^{local}\\gg1$, is single-field slow-roll inflation dying? Not necessarily. There are at least two ways to save it. First, a large non-Gaussianity may arise in stochastic inflation due to non-linear effects between inflaton and metric perturbations \\cite{Gangui:1993tt}. Second, the curvaton mechanism provides an elegant way to produce a large positive $f_{NL}^{local}$ \\cite{Lyth:2002my}.  Although there is no evidence for $|f_{NL}^{equil}|$ to be large at present, several inflation models with $|f_{NL}^{equil}|\\gg1$ have already appeared in the past few years. In single-field slow-roll inflation, this is realized by the introduction of non-canonical kinetic terms, such as $k$-inflation \\cite{ArmendarizPicon:1999rj,Garriga:1999vw}, ghost inflation \\cite{ArkaniHamed:2003uz} DBI inflation \\cite{Silverstein:2003hf,Alishahiha:2004eh,Chen:2004gc,Chen:2005ad}, and some other mechanisms \\cite{berera}. Translated into the WMAP convention, most of the models predict $f_{NL}^{equil}\\ll-1$ and a small $f_{NL}^{local}$. It is remarkable that the value of $f_{NL}^{local}$ favored by \\cite{Yadav:2007yy} is of positive sign in the WMAP convention. For $f_{NL}^{equil}$, the present constraint is not stringent enough to make a conclusion. So one expects that there are three possibilities in the future: \\begin{enumerate} \\item $f_{NL}^{local}\\gg1$, $-1<f_{NL}^{equil}<1$. This can be explained by the conventional single-field slow-roll inflation + curvaton mechanism. \\item $f_{NL}^{local}\\gg1$, $f_{NL}^{equil}<-1$. This corresponds to DBI inflation (or most known $k$-inflation/ghost inflation) + curvaton mechanism. \\item $f_{NL}^{local}\\gg1$, $f_{NL}^{equil}>1$. This is more challenging to explain. \\end{enumerate}  For both the second and the third possibility, one should keep in mind that additional fine-tuning is needed in order that both the curvaton and the inflaton produce perturbations of comparable magnitude.  The purpose of this paper is to search for $k$-inflation models with $f_{NL}^{equil}\\gg1$. If this type of models are constructed, one can combine them with the curvaton mechanism to meet the challenge posed by the third possibility.  Indeed, starting from the action (\\ref{action}), we have found several examples of general single field inflation in which $f_{NL}^{equil}\\gg1$. As will be shown in section \\ref{run}, even if the inflation is driven by the potential of inflaton, a large positive $f_{NL}^{equil}$ can be generated by non-conventional kinetic terms. In some other models, the non-conventional kinetic terms not only give rise to the desired $f_{NL}^{equil}$, but also drive the inflation. Examples for this type of models are constructed in power-law $k$-inflation. In all our models, typically the desired non-Gaussianity stems from high order terms in $X$, with $X$ defined in (\\ref{X}).  The paper is organized as follows. In section \\ref{review}, we review briefly the general single field inflation. In section \\ref{nogo}, we prove a no-go theorem for the $p(X)$ models. We show that if the matter Lagrangian depends only on $X$, but not on $\\phi$ directly, then one can not obtain a large and positive $f^{equil}_{NL}$. In section \\ref{run}, we construct generalized slow roll inflation models with $f^{equil}_{NL}\\gg 1$. In section \\ref{pl}, we construct power-law inflation models with $f^{equil}_{NL}\\gg 1$. We conclude in section \\ref{conclusion}. ",
        "conclusions": "In this paper, we have constructed several models with a large positive $f_{NL}$ in the WMAP convention \\cite{Komatsu:2001rj,Spergel:2006hy}. These models are general single field inflation with higher order kinetic terms. In one class of models, the inflation is driven by the potential term of the inflaton, as shown in section \\ref{run}. In section \\ref{pl}, we have given another class of models, in which inflation is driven by non-conventional kinetic terms \\cite{ArmendarizPicon:1999rj,Garriga:1999vw}. In both classes of models, due to the appearance of non-conventional kinetic terms, we can arrange the parameters to produce $f_{NL}^{equil}\\gg1$. These are first examples to generate $f_{NL}^{equil}\\gg1$ in general single field $k$-inflation. They are different from ghost inflation \\cite{ArkaniHamed:2003uz} since we restricted our discussion in $k$-inflation and did not introduce terms like $\\nabla^2\\phi$ in our Lagrangian.  The common features of these models are as following: \\begin{enumerate} \\item Typically we need to introduce four parameters in a model, although there are only three data to fit, the COBE normalization of the two point function, the spectral index $n_s$ and the non-Gaussianity parameter $f_{NL}$. Introduction of four parameters is not strictly necessary. For example, for models studied in section 4, we introduced two parameters $\\alpha$ and $\\beta$ in order to have a model easy to solve. In section 5, we introduced four parameters $c_i$ since we try to also fit a parameter $r$ whose value is not fixed by experiments yet. However, a upper bound on $r$ is set, so it is necessary to make sure that our model does not violate this bound. \\item High order terms in $X$ are absolutely necessary, and luckily these terms can be viewed as operators of high dimensions in an effective field theory with a mass cut-off lower than the Planck scale. The bad news is that, the highest order term in general triggers an instability if we do not introduce more high order terms. \\end{enumerate}"
    },
    "0801/0801.2273_arXiv.txt": {
        "abstract": "{Classical T Tauri stars (hereafter CTTS) show a plethora of in- and outflow signatures in a variety of wavelength bands.}{In order to constrain gas velocities and temperatures, we analyse the emission in the hot ion lines.}{We use all available archival \\emph{FUSE} spectra of CTTS to measure the widths, fluxes and shifts of the detected hot ion lines and complement these data with \\emph{HST/GHRS} and \\emph{HST/STIS} data. We present theoretical estimates of the temperatures reached in possible emission models such as jets, winds, disks and accretion funnels and look for correlations with X-ray lines and absorption properties.}{We find line shifts in the range from -170~km~s$^{-1}$ to +100~km~s$^{-1}$. Most linewidths exceed the stellar rotational broadening. Those CTTS with blue-shifted lines also show excess absorption in X-rays. CTTS can be distinguished from main sequence (hereafter MS) stars by their large ratio of the \\ion{O}{vii} to \\ion{O}{vi} luminosities.}{No single emission mechanism can be found for all objects. The properties of those stars with blue-shifted lines are compatible with an origin in a shock-heated dust-depleted outflow.} ",
        "introduction": "\\label{introduction}  T Tauri stars (TTS) are young ($<10\\;\\mathrm{Myr}$), low mass ($M_*<3M_{\\sun}$), pre-main sequence stars exhibiting strong H$\\alpha$ emission.  The subclass of Classical TTS (CTTS) are accreting material from a surrounding disk.  This disk is truncated close to the corotation radius by interaction with the magnetic field. At the inner rim, the material is ionised by the stellar radiation and loaded onto the field lines \\citep{1984PASJ...36..105U,1991ApJ...370L..39K} and accelerated to nearly free-fall velocity along an accretion funnel before hitting the stellar surface where an accretion shock forms. Angular momentum may be transferred from the star to the disk by magnetic torque \\citep{1994ApJ...429..781S}, at the same time this process is expected to drive outflows possibly in the form of a disk wind or bipolar jets. \\citet{2004ApJ...610..920R} performed full magneto-hydrodynamic (MHD) calculations of the accretion geometry, while the accretion shock and its post-shock cooling zone were simulated by different authors \\citep{calvetgullbring,lamzin, acc_model}.  Observational support for the scenario described above is found in data from different wavelength bands. The shocked material is heated up to temperatures in the MK range and should produce copious amounts of X-ray emission.  Density-sensitive X-ray line ratios strongly suggest a post-shock origin \\citep{acc_model} of the cool parts of the plasma observed at X-ray wavelengths, and the cooling matter veils absorption lines in the optical and infrared \\citep{calvetgullbring}. The geometry of the pre-shock accretion funnel can be traced by H$\\alpha$ line profile analysis \\citep{2000ApJ...535L..47M}. Undoubtedly CTTS also show outflows \\citep[a review is given by][]{2007prpl.conf..215B}, however, the precise origin and the physical driving mechanism(s) are uncertain.  Theoretical models propose a variety of collimated stellar outflows and disk winds to remove angular momentum from the system \\citep{1994ApJ...429..781S,2002ApJ...574..232M,2005MNRAS.362..569G,2007prpl.conf..277P}. The outflows can take the form of winds observed from radio to the UV \\citep[e.g.][]{2000AJ....119.1881A,2001ApJ...551.1037B,2004AstL...30..413L,2005ApJ...625L.131D,2006ApJ...646..319E}, or collimated, often bipolar jets, again observed at different wavelength bands \\citep{1995RMxAC...1....1R,2004ApJ...604..758C,2005ApJ...626L..53G}. The dynamics of the surrounding gas have been successfully probed by UV spectroscopy with \\emph{HST/GHRS}, \\emph{HST/STIS} and \\emph{FUSE}, especially in the H$_2$ lines, which show outflows and fluorescence on the hot disk surface \\citep{2002ApJ...566.1100A,2002ApJ...567.1013A, 2002ApJ...572..310H, 2005AJ....129.2777H, 2006ApJS..165..256H}. In the specific case of TW Hya, \\citet{2005ApJ...625L.131D} claim the existence of a warm/hot wind. Their conclusion assumes that the CIII and OVI lines have an intrinsic Gaussian shape, with centroid matching the stellar rest-frame. The observed asymmetry of the CIII and OVI line-profiles can then be explained by the presence of spherically-symmetric and smoothly-accelerated hot wind absorbing the radiation in the blue wing of these particular lines. However, \\citet{2007ApJ...655..345J} argue that this model is incompatible with \\emph{HST/STIS} observations, especially for the \\ion{C}{iv} 1550~\\AA{} doublet. The existence of a hot wind in TW Hya, therefore, remains an open issue.  In this paper we analyse the shifts and widths of hot-ion spectral lines, particularly \\ion{C}{iii} 977~\\AA{} and the \\ion{O}{vi} doublet at 1032~\\AA{} and 1038~\\AA{}, observed with \\emph{FUSE}. This complements earlier studies of all pre-main sequence stars observed with \\emph{IUE} \\citep{2000ApJS..129..399V}, \\emph{GHRS} \\citep{2002ApJ...566.1100A,2002ApJ...567.1013A} and of the $H_2$ emission in CTTS \\citep{2006ApJS..165..256H}. In Sect.~\\ref{stellarproperties} we briefly summarise the properties of the sample stars, and in Sect.~\\ref{observations} we describe the observations and data that we use in our present analysis.  In Sect.~\\ref{results} the line profile analysis is presented, several possible formation regions for the observed properties are discussed in Sect.~\\ref{discussion}, and in Sect.~\\ref{summary} we summarise our results. ",
        "conclusions": "\\label{summary} We have presented the hot ion lines originating from \\ion{C}{iii} and \\ion{O}{vi} in all available CTTS spectra observed with \\emph{FUSE}. The fitted centroids range from about -170~km~s$^{-1}$ to +100~km~s$^{-1}$ and the shifts of the \\ion{C}{iii} 977~\\AA{} line and the \\ion{O}{vi} 1032~\\AA{} and 1038~\\AA{} doublet are consistent. Most, if not all, lines are superrotationally broadened. The blue-shifted lines could originate in a stellar outflow, maybe a jet; the red-shifted lines are incompatible with current models of magnetospheric accretion. The presence of boundary layers, winds and extended chromospheres is inconsistent with the observed line profiles.  Furthermore, the absorbing column densities observed in X-rays are incompatible with the optical reddening measured assuming a standard gas-to-dust ratios and  interstellar extinction law. This situation holds for the most heavily-veiled objects AA~Tau, DF~Tau, T~Tau, SU~Aur and, although for lower absolute values of reddening, RU~Lup. In cases where data are available from both UV and X-ray, the corresponding objects demonstrate blue-shifted emission in the \\emph{FUSE} observations. We interpret this as a sign of a dust-free absorber, which may consist of shock-heated jet material. In the ratio between the \\ion{O}{vii} luminosity and the \\ion{O}{vi} luminosity, all CTTS display a clear excess of soft X-ray emission in the \\ion{O}{vii} lines. Together with X-ray results, this illustrates that the signatures of a hot accretion spot are most evident in the temperature range 1-2~MK."
    },
    "0801/0801.0923_arXiv.txt": {
        "abstract": "In this paper we review the possible mechanisms for production of non-thermal electrons which are responsible for the observed non-thermal radiation in clusters of galaxies. Our primary focus is on non-thermal Bremsstrahlung and inverse Compton scattering, that produce hard X-ray emission. We first give a brief review of acceleration mechanisms and point out that in most astrophysical situations, and in particular for the intracluster medium, shocks, turbulence and plasma waves play a crucial role. We also outline how the effects of the turbulence can be accounted for. Using a generic model for turbulence and acceleration, we then consider two scenarios for production of non-thermal radiation. The first is motivated by the possibility that hard X-ray emission is due to non-thermal Bremsstrahlung by nonrelativistic particles and  attempts to produce non-thermal tails by accelerating the electrons from the background plasma with an initial Maxwellian distribution.  For acceleration rates smaller than the Coulomb energy loss rate, the effect of energising the plasma  is to primarily heat the plasma  with little sign of a distinct non-thermal tail. Such tails are discernible only for acceleration rates comparable or larger than the Coulomb loss rate.  However, these tails are accompanied by significant heating and they are present for a short  time of $<10^6$ yr, which is also the time that the tail will be thermalised. A longer period of acceleration at such rates will result in a runaway situation with most particles being accelerated to very high energies. These more exact treatments confirm the difficulty with this model, first pointed out by Petrosian (2001). Such non-thermal tails, even if possible, can only explain the hard X-ray but not the radio emission which needs GeV or higher energy electrons. For these and for production of hard X-rays by the inverse Compton model, we need the second scenario where there is injection and subsequent acceleration of relativistic electrons. It is shown that a steady state situation, for example arising from secondary electrons produced from cosmic ray proton scattering by background protons, will most likely lead to flatter than required electron spectra or it requires a short escape time of the electrons from the cluster. An episodic injection of relativistic electrons, presumably from galaxies or AGN, and/or episodic generation of turbulence and shocks by mergers can result in an electron spectrum consistent with observations but for only a short period of less than one billion years. ",
        "introduction": "\\label{intro2}  In this paper we attempt to constrain the acceleration models based on the observations described in the  papers by \\citet{Durret2008}, \\citet{Rephaeli2008} and \\citet{Ferrari2008} - Chapters $4-6$, this volume, and the required spectrum of the accelerated electrons shown in Fig.~6 of \\citealt{Petrosian2008} - Chapter 10, this volume. However, before addressing these details  we first compare various acceleration processes and stress the importance of plasma waves or turbulence (PWT) as an agent of scattering and acceleration, and then describe the basic scenario and equations for treatment of these processes. As pointed out  below there is growing evidence that PWT plays an important role in acceleration of particles in general, and in clusters of galaxies in particular. The two most commonly used acceleration mechanisms are the following.  \\subsection{Electric field acceleration}  Electric fields parallel to magnetic fields can accelerate charged particles and can arise as a result of magnetic field reconnection in a current sheet or other situations. For fields less than  the so-called Dreicer field, defined as $E_{\\rm D} = {\\rm k}T/({\\rm e}\\lambda_{\\rm Coul})$, where \\begin{equation} \\lambda_{\\rm Coul}\\sim 15\\ {\\rm kpc}\\left(\\frac{T}{10^8\\,{\\rm K}}\\right)^{2}\\left(\\frac{10^{-3}\\,{\\rm cm}^{-3}}{n}\\right) \\label{meanfp} \\end{equation} is the collision (electron-electron or proton-proton) mean free path\\footnote{The proton-proton or ion-ion mean free path will be slightly smaller because of the larger value of the Coulomb logarithm $\\ln \\Lambda\\sim 40$ in the ICM.}, the rate of acceleration is less than the rate of collision losses and only a small fraction of the particles can be accelerated into a non-thermal tail of energy $E<L{\\rm e}E_{\\rm D}$. For the ICM $E_{\\rm D}\\sim 10^{-14}$~V\\,cm$^{-1}$ and $L\\sim 10^{24}$~cm so that sub-Dreicer fields  can only accelerate particles up to 100's of keV, which is far below the  10's of GeV electrons required by observations.  Super-Dreicer fields, which seem to be present in many simulations of reconnection \\citep{Drake2006,Cassak2006,Zenitani2005},  accelerate particles at a rate that is faster than the collision or thermalisation time $\\tau_{\\rm therm}$. This can lead to a runaway and an unstable electron distribution which, as shown theoretically, by laboratory experiments and by the above mentioned simulations, most probably will give rise to PWT  \\citep{Boris1970,Holman1985}.  {\\sl In summary the electric fields arising as a result of reconnection cannot be the sole agent of acceleration in the ICM, because there are no large scale magnetically dominated cosmological flows, but it may locally produce an unstable particle momentum distribution which will produce PWT that can then accelerate particles.}  \\subsection{Fermi acceleration}  Nowadays this process has been divided into two kinds. In the original Fermi process particles of velocity $v$ moving along magnetic field lines (strength $B$) with a pitch angle $\\cos\\mu$ undergo random scattering by moving agents with a velocity $u$. Because the head (energy gaining) collisions are more probable than trailing (energy losing) collisions, on average, the particles gain energy at a rate proportional to $(u/v)^2D_{\\mu\\mu}$, where $D_{\\mu\\mu}$ is the pitch angle diffusion rate. This, known as a {\\sl second order Fermi process} is what we shall call stochastic acceleration. In general, the most likely agent for scattering is PWT. An alternative process is what is commonly referred to as a {\\sl first order Fermi process}, where the actual acceleration occurs when particles cross a  shock or any region of converging flow. Upon crossing the shock  the fractional gain of momentum $\\delta p/p\\propto u_{\\rm sh}/v$. Ever since the 1970's, when several authors demonstrated that a very simple version of this process leads to a power law spectrum that agrees approximately with observations of the cosmic rays, shock acceleration is commonly invoked in space and astrophysical plasmas. However, this simple model, though very elegant, has some shortcomings specially when applied to electron acceleration in non-thermal radiating sources. Moreover,  some of the features that make this scenario for acceleration of cosmic rays attractive are not present in most radiating sources where one needs efficient acceleration of electrons to relativistic energies from a low energy reservoir.  The original, test particle theory of diffusive shock acceleration (DSA), although very elegant and independent of geometry and other details (e.g. \\citealt{Blandford1978}) required several conditions such as injection of seed particles and of course turbulence. A great deal of work has gone into addressing these aspects of the problem and there has been a great deal of progress. It is clear that nonlinear effects  (see e.g. \\citealt{Drury1983,Blandford1987,Jones1991,Malkov2001}) and losses (specially for electrons)  play an important role and modify the resultant spectra and efficiency of acceleration. Another important point is the source of the turbulence or the scattering agents. A common practice is to assume Bohm diffusion (see e.g. \\citealt{Ellison2005}). Second order acceleration effects could modify the particle spectra accelerated by shocks (see e.g. \\citealt{Schlickeiser1993,Bykov2000}). Although there are indications that turbulence may be generated by the shocks and the accelerated particle upstream, many details (e.g. the nature and spectrum of the turbulence) need to be addressed more quantitatively. There has been progress on the understanding of generation of the magnetic field and turbulence on strong shocks \\citep{Bell2001,Amato2006,Vladimirov2006} as required in recent observations of supernova remnants (see e.g. \\citealt{Volk2005}). There is also some evidence for these processes from observations of heliospheric shocks (see e.g. \\citealt{Kennel1986,Ellison1990}). Basic features of particle acceleration by cosmological shocks were discussed by \\citealt{Bykov2008a} - Chapter 7, this volume,  so we will concentrate here on the stochastic acceleration perspective.  \\subsection{Stochastic acceleration}  The PWT needed for scattering can also accelerate particles stochastically with a rate $D_{EE}/E^2$, where $D_{EE}$ is the energy diffusion coefficient, so that shocks may not be always necessary.  In low beta plasmas, $\\beta_{\\rm p}=2(v_{\\rm s}/v_{\\rm A})^2<1$, where the  Alf\\'ven velocity $v_{\\rm A}=\\sqrt{B^2/4\\pi \\rho}$, the sound velocity $v_{\\rm s}=\\sqrt{{\\rm k}T/m}$, $\\rho=nm$ is the mass density and $n$ is the number density of the gas, and for relativistic particles the PWT-particle interactions are dominated by Alf\\'venic turbulence, in which case the rate of energy gain $D_{EE}/E^2=(v_{\\rm A}/v)^2D_{\\mu\\mu}\\ll D_{\\mu\\mu}$, so that the first order Fermi process is more efficient. However, at low energies and/or in very strongly magnetised plasmas, where $v_{\\rm A}$  can exceed ${\\rm c}$, the speed of light\\footnote{Note that the Alf\\'ven group velocity $v_{\\rm g}= {\\rm c}\\sqrt{v_{\\rm A}^2/(v_{\\rm A}^2+{\\rm c}^2)}$ is always less than ${\\rm c}$.}, the acceleration rate may exceed the scattering rate (see \\citealt{Pryadko1997}), in which case low energy electrons are accelerated more efficiently by PWT than by shocks.\\footnote{In practice, i.e. mathematically, there is little difference between the two mechanisms \\protect\\citep{Jones1994}, and the acceleration by turbulence and shocks can be combined (see below).}  Irrespective of which process dominates the particle acceleration, it is clear that PWT has a role in all of them. Thus, understanding of the production of PWT and its interaction with particles  is extremely important. Moreover, turbulence is expected to be present in most astrophysical plasmas including the ICM and in and around merger or accretion shocks, because the ordinary and magnetic Reynolds numbers  are  large. Indeed turbulence may be the most efficient channel of energy dissipation. In recent years there has been a substantial progress in the understanding of MHD turbulence \\citep{Goldreich1995,Goldreich1997,Lithwick2003,Cho2002,Cho2006}. These provide new tools for a more quantitative investigation of turbulence and the role it plays in many astrophysical sources. ",
        "conclusions": "We have given a brief overview of particle acceleration in astrophysical plasmas in general, and acceleration of electrons in the ICM in particular. We have pointed out the crucial role plasma waves and turbulence play in all acceleration mechanisms and outlined the equations that describe the generation, cascade and damping of these waves and the coupling of these processes to the particle kinetics and energising of the plasma and acceleration in both relativistic and nonrelativistic regimes.  We have applied these ideas to the ICM of clusters of galaxies with the aim of production of electron spectra which can explain the claimed hard X-ray emission either as  non-thermal bremsstrahlung emission by nonrelativistic electrons or as inverse Compton emission via scattering of CMB photons by a population of relativistic electrons. It is shown that the first possibility which can come about by accelerating background  electrons into a non-thermal tail is not a viable mechanism, as was pointed earlier in P01. The primary reason for this difficulty is due to the short Coulomb collision and loss timescales. Quite generally, it can be stated that at low rates of acceleration one obtains a hotter plasma and an insignificant non-thermal tail. Discernible tails can be obtained at higher rates of acceleration but only for short periods of time. For periods on the order of a billion year such rates will also cause excessive heating and will lead to runaway conditions where most of the electrons are accelerated to relativistic energies, at which they are no longer bound to the cluster, unless there exists a strong scattering agent.  This leads us to the model where hard X-rays are produced by the inverse Compton process and relativistic electrons. Moreover, even if the hard X-ray radiation turns out to be not present, or one finds a way to circumvent the above difficulties, we still require the presence of relativistic electrons to explain the radio emission. These electrons must be injected into the ICM by some other means. They can come from galaxies, specially when they are undergoing an active nuclear (or AGN) phase. Or they may be due to interactions of cosmic ray protons with thermal protons and the resultant pion decays. We have shown that just injection may not be sufficient, because for reasonable injected spectra the transport effects in the ICM modify the spectrum such that the effective radiating spectrum is inconsistent with what is required. Thus, a re-acceleration in the ICM is necessary and turbulence and merger shocks may be the agent of this acceleration. In this case, it also appears that a steady state scenario, like the hadronic mechanism described above, will in general give a flatter than required spectrum unless the electrons escape the ICM unhindered. This requirement is not reasonable because the expected tangled magnetic field will increase this time. But, more importantly, the presence of turbulence necessary for re-acceleration will result in a short mean free path and a much longer escape time.  A more attractive scenario is if the injection of electrons and/or the production of turbulence is episodic. For example for a short lived electron injected phase (from say an AGN) but a longer period of presence of turbulence one can determine the spectral evolution of the electrons subject to acceleration and losses. We have shown that for some periods of time lasting several times the acceleration timescales one can obtain electron spectra consistent with what is required by observations. The same will be true for a hadronic source if there is a short period of production of turbulence. In either case we are dealing with periods on the order of several  hundreds of million years to a billion years, which is comparable with timescales expected from merging of Mpc size clusters with velocities of several thousands of km\\,s$^{-1}$ which are theoretically reasonable and agree with observations (see e.g. \\citealt{Bradac2006})."
    },
    "0801/0801.0112_arXiv.txt": {
        "abstract": "The  magnetic Reynolds number $R_M$, is defined as the product of a characteristic  scale and associated flow  speed  divided by the microphysical magnetic diffusivity. For laminar flows, $R_M$ also  approximates the ratio of advective to dissipative terms in the total  magnetic energy equation, but  for  turbulent flows this latter ratio  depends on the energy spectra and approaches unity in a steady state. {To generalize for flows of arbitrary spectra we define an effective magnetic dissipation number, $R_{M,e}$, as the ratio of the advection to microphysical dissipation terms in  the total magnetic energy equation, incorporating the full spectrum of scales, arbitrary magnetic Prandtl numbers,  and  distinct pairs of inner and outer scales for   magnetic and kinetic spectra.  As expected,  for a substantial  parameter range $R_{M,e}\\sim {O}(1) << R_M$. We also distinguish  $R_{M,e}$ from ${\\tilde R}_{M,e}$ where the latter is an effective magnetic Reynolds number  for the mean magnetic field equation when a turbulent diffusivity is explicitly imposed as a closure. That  $R_{M,e}$ and ${\\tilde R}_{M,e} $  approach unity even if $R_M>>1$ highlights that, just as in hydrodynamic turbulence, energy dissipation  of large scale structures in turbulent flows via a cascade can be  much faster than the dissipation of large scale structures in laminar flows. This illustrates that the rate of energy dissipation by magnetic reconnection is much faster in turbulent flows, and much less sensitive to microphysical reconnection rates compared to laminar flows.} ",
        "introduction": "Magnetic fields play an observable and dynamical role in a range of astrophysical and laboratory plasmas.  When the scales of interest are larger than the particle mean free path, magnetohydrodynamics provides a suitable approximation (\\cite{moffatt}, \\cite{p79}, \\cite{biskamp97}). Derived from Maxwell's equations and Ohm's law, the induction equation describing the time evolution of the magnetic field $\\bfB$ in non-relativistic magnetohydrodynamics (MHD) is given by \\beq \\partial_t\\bfB= \\curl(\\bfv\\ts\\bfB) +\\lambda\\nabla^2\\bfB, \\label{0} \\ee where $\\lambda$ is the magnetic diffusivity, and $\\bfv$ is the velocity, determined from the fluid momentum equation. For incompressible flows, the momentum equation incorporating the Lorentz force is \\beq \\partial_t\\bfv=-\\bfv\\cdot\\nabla \\bfv-\\nabla P/\\rho+ (\\curl\\bfB)\\ts\\bfB/4\\pi\\rho +\\nu \\nabla^2\\bfv, \\label{01} \\ee where $P$ is the pressure and $\\nu$ is the kinematic viscosity, and $\\rho$ is the density.  Taking the approximate  ratio of magnitudes of the second to third terms in (\\ref{0})  introduces the magnetic Reynolds number \\beq R_{M,l}\\equiv l v/\\lambda, \\label{oldrm} \\ee where  $l$ is a macroscopic scale of field variation for the system being studied. When $l$ is the largest macroscopic scale of the system,  we write $R_M$, eliminating the subscript $l$.  For laminar flows in which $R_M>>1$,  the last term of (\\ref{0}) is often ignored.  An important consequence of ignoring this term is magnetic flux freezing, $d_t\\int\\bfB\\cdot d{\\bf S}=0$ (\\cite{p79}) which implies that the integrated product of the magnetic field and the surface area it penetrates, or the magnetic flux moving with the plasma, is  conserved. Dotting (\\ref{0}) with $\\bfB$ gives the same $R_M$ for the corresponding terms in the magnetic energy equation so hat  $R_M>>1$ would also imply that the resistive dissipation term is small compared to the other dynamical terms. When most of the magnetic energy is on the largest scale, the implication is that the dissipation of the total magnetic energy is negligible.  The above physical implications of  $R_M>>1$  do not  apply if the flow is turbulent. When an initially laminar flow becomes turbulent, the field and flow evolve to develop non-trivial turbulent spectra (e.g. \\cite{biskamp97,biskamp03}). Eqn. (\\ref{oldrm}), with $l$ taken as the largest macroscopic scale of the system, is no longer a good approximation to the ratio of the advection to dissipation term in the total magnetic energy equation even if the energy is dominated by structures on that scale. The dissipation term is  not necessarily  negligible because it depends more strongly on high wavenumber components of the turbulent spectrum than do the advection terms.  The importance of  dissipation in MHD turbulence just described, is directly analogous to the case of  hydrodynamic turbulence. For the latter, the Kolmogorov steady-state assumption (Kolmogorov 1941; see also Davidson 2004) that the energy dissipation rate per mass is constant on all scales through the cascade implies that the energy dissipation rate at the dissipation scales exactly equals that at the input scale. Therefore, the dissipation term in the total energy evolution equation exactly balances the nonlinear energy transfer terms when all scales are included and a steady-state is assumed. Extracting dimensionless measures of hydrodynamic turbulence dissipation rates have confirmed this basic conceptual picture (e.g. Pearson et al. 2004).  In making contact with these concepts from hydrodynamic turbulence, we find it particularly important to emphasize that dissipation is non-negligible in MHD turbulence even though $R_M>>1$. This very much contrasts the physical implications commonly attributed to laminar $R_M>>1$  MHD flows. Characterizing the MHD energy depletion time scales for large $R_M$ turbulent MHD flows such as those of stellar convection zones, accretion disks, and galactic interstellar media, is important for understanding the  magnetic field evolution, conversion of magnetic energy into particles, and requirements for magnetic field sustenance. Since these flows are characterized by a spectrum of scales (\\cite{biskamp97}, \\cite{biskamp03}), the ratio of advection to dissipation terms in the total magnetic energy equation  is not  in general $R_M$ but  a generalized magnetic dissipation  number $R_{M,e}$ which depends on the magnetic and kinetic energy spectra.  Quantifying $R_{M,e}$, requires incorporating the spectrum of magnetic and velocity fields along with  the magnetic Prandtl number $Pr_M\\equiv\\nu/\\lambda$. {Here we calculate $R_{M,e}$ as a function of the magnetic spectral index and $Pr_M$, the latter of which depends on the microphysical diffusivities and thus on $R_M$. Although $R_M$ (as defined with the largest scale) does help determine  whether the flow would have become turbulent in the first place, it is not a good indicator of the time scale  for overall  energy dissipation when the flow is turbulent. For regimes in which  $R_{M,e}$ is largely independent of $R_M$, the total energy depletion rate  becomes largely independent of the specific small scale mechanism of dissipation; turbulence cascades the energy to scales where the dissipation time is short. }     In section 2, we derive the approximations for the corresponding second and third terms of (\\ref{0}) in the magnetic energy equation, incorporating the turbulent spectra and triple correlations for the case without a mean field. We then plot $R_{M,e}$ as a function of the magnetic spectral index and $Pr_M$. In section 3, we  distinguish $R_{M,e}$  from a different  quantity ${\\tilde R}_{M,e}$ for the mean field equation. In section 4, we conclude with  a discussion of the implications of ${\\tilde R}_{M,e}<<1$ and $R_{M,e} <<R_M$. ",
        "conclusions": "{ We have  calculated a dimensionless measure of the relative importance of advective vs. dissipative terms for the total magnetic energy equation in MHD that applies for arbitrary magnetic and kinetic spectra, and magnetic Prandtl number. We have discussed how the standard magnetic Reynolds number $R_M$ defined using the largest gradient scale of a system does not  accurately  approximate this ratio for turbulent systems, even if the magnetic energy is dominated by large scale structures. Instead, this ratio is best approximated by  $R_{M,e}$, a quantity largely independent of $R_M$ for a range of turbulent spectra (though reverting to $R_M$ for a laminar flow) and typically satisfies $R_{M,e}\\sim O(1) <<R_M$. The latter results because small scale structures augment the relative importance of the microphysical dissipation terms for a turbulent system compared to a laminar system with the same microphysical diffusion coefficient. In a steady state,  such as those  reached when MHD turbulence is steadily forced, $R_{M,e}\\sim 1$.   We have also discussed how the effective dissipation number ${\\tilde R}_{M,e}$ for the mean magnetic field equation is defined differently from $R_{M,e}$ because the mean field equation picks out the gradient scale of the mean field only.  A low value of ${\\tilde R}_{M,e}<<R_M$ results not from small scale gradients, but because macroscopic turbulent diffusion coefficients exceed  microscopic diffusivities in closures that approximate  nonlinear cascade terms by diffusion.   While the results that ${\\tilde R}_{M,e}<<R_M$, $ R_{M,e}<<R_M$, are  not at all surprising when interpreting them analogously to steady-state Kolmogorov hydrodynamic turbulence, they do highlight the much greater role of dissipation in large $R_M$ turbulent  MHD flows compared to large $R_M$ laminar  flows. In particular, an important implication  of  ${\\tilde R}_{M,e}<<R_M$, $ R_{M,e}<<R_M$, and the very weak dependence of  $R_{M,e}$ and ${\\tilde R}_{M,e}$ on $R_M$  is that the rate at which large scale MHD energy is dissipated can be much less sensitive to the actual microphysical dissipation mechanism in turbulent flows when compared to laminar flows. On a macroscopic dynamical time, the MHD energy cascades to scales small enough such that  it can be drained at multiple sites. If the turbulence is steadily forced, then the system adjusts to produce enough small scale structures to accommodate the driving energy input rate.  The overall energy dissipation rate can  be insensitive to the microphysical reconnection  (or other microphysical energy conversion rate) at any one site. Instead, the aggregate of dissipation sites emerges to drain the total magnetic energy  on macroscopic dynamical time scales. Since many astrophysical settings are MHD turbulent, this highlights the importance of being precise when asking and studying the question ``Is magnetic reconnection fast in astrophysics?''   The principle that small scale structures increase the effective reconnection rate is consistent with  rapid reconnection mechanisms in the MHD  models of Hendrix et al. (1996), Galasgaard \\& Nordlund (1996), Lazarian \\& Vishniac (1999), and in the rapid reconnection phase of  Fan et al. (2004). Reconnection studies that start with a large scale laminar field and reconnect on collisionless scales (e.g. \\cite{birn,shay}) address a different regime.  The latter can however apply to each of the small scale multiple dissipation sites within a globally turbulent MHD system. In some systems, collisionless individual structures are observationally resolved (e.g. the Earth's magnetosphere) and studying the plasma instabilities and plasma turbulence for different conditions on such scales is needed to understand when reconnecting structures  incur a reconnection  rate that depends only on its macroscopic properties (\\cite{kulsrud}). That being said, if a stellar or accretion disk corona  is considered as a global entity, its overall energy dissipation rate in a steady state is determined by the dynamical rate of magnetohydrodynamic energy injection from below; the small scale structures develop to accommodate this rate.  The overall dissipation rate is best characterized by ${\\tilde R}_M<<R_M$ and $R_{M,e}<<R_M$ even if the rate of each individual dissipation site were to depend on  $R_M$.                    \\ni {\\bf Acknowledgments:} EGB acknowledges support from NSF grant AST-0406799 and NASA grant ATP04-0000-0016."
    },
    "0801/0801.3327_arXiv.txt": {
        "abstract": "We present the results of a 25-year program to monitor the radio flux evolution of the planetary nebula NGC\\,7027. We find significant evolution of the spectral flux densities.  The flux density at 1465~MHz, where the nebula is optically thick, is increasing at a rate of $0.251\\pm0.015\\%\\,\\rm yr^{-1}$, caused by the expansion of the ionized nebula. At frequencies where the emission is optically thin, the spectral flux density is changing at a rate of $-0.145\\pm0.005\\%$ per year, caused by a decrease in the number of ionizing photons coming from the central star. A distance of $980\\pm100$\\,pc is derived. By fitting interpolated models of post-AGB evolution to the observed changes, we find that over the 25-yr monitoring period, the stellar temperature has increased by $3900\\pm900$\\,K and the stellar bolometric luminosity has decreased by $1.75\\pm0.38$\\%.  We derive a distance-independent stellar mass of $0.655\\pm0.01\\,\\rm M_\\odot$ adopting the Bl\\\"ocker stellar evolution models, or about 0.04\\,M$_\\odot$ higher when using models of Vassiliadis \\&\\ Wood which may provide a better fit. A Cloudy photoionization model is used to fit all epochs at all frequencies simultaneously. The differences between the radio flux density predictions and the observed values show some time-independent residuals of typically 1\\%. A possible explanation is inaccuracies in the radio flux scale of Baars {\\sl et al.}  We propose an adjustment to the flux density scale of the primary radio flux calibrator 3C\\,286, based on the Cloudy model of NGC\\,7027. We also calculate precise flux densities for NGC\\,7027 for all standard continuum bands used at the VLA, as well as for some new 30\\,GHz experiments. ",
        "introduction": "Planetary nebulae trace one of the fastest phases of stellar evolution. They form during the transformation from a low to intermediate mass star (1--8\\,M$_\\odot$) to its final white dwarf stage. The star first ejects its hydrogen envelope during a phase of catastrophic mass loss.  This mass loss ceases once the hydrogen envelope is reduced to $\\sim 0.02\\,\\rm M_\\odot$. The star subsequently heats up, briefly reaching surface temperatures in excess of $10^5$\\,K, before the cessation of hydrogen burning when the remnant C/O core enters the white dwarf cooling track. The evolution from the end of the AGB mass loss to the onset of cooling lasts between $10^3$ and $10^5$\\,yr: the duration is a steep function of white dwarf mass \\citep{Bloecker1995}. The time scales are short enough that one may expect, in some cases, to observe significant changes in stellar temperature and luminosity within decades.  Once the star has reached temperatures above $2 \\times 10^4\\,$K, it ionizes the ejecta forming a bright planetary nebula (PN). The strong emission lines and radio continuum from the PN provides a tracer of the evolution of the parent star.  \\citet{Masson1986, Masson1989} detected evidence for expansion of the brightest planetary nebula, NGC\\,7027, using radio images over a four-year baseline: this provided proof of the fast evolution. As the expansion velocity of the nebular gas is known, the result provides both a measure of its distance and of its age since ejection.  Masson found a distance of 880\\,pc. The same method has been applied by \\citet{Hajian1993} who found a somewhat faster angular expansion and therefore a lower distance, 770\\,pc.  The radio method has been applied to three other PNe \\citep{Seaquist1991, Christianto1998, Hajian1993, Kawamura1996}. Expansion-based distances have also been measured using optical HST images \\citep{Palen2002}.  The evolution of the temperature and luminosity of the star can be derived through observations at optical or radio wavelengths. In the optical, the absolute photometric accuracy is set by the Earth's atmosphere as well as instrumentational limitations. At radio frequencies, the required accuracies are currently achievable in the centimeter wavelength range, due to the limited to negligible effect of the atmosphere, the stability of radio telescopes and electronics, and the availability of constant flux calibrators.  In this paper, we utilize these to investigate variations in the the radio flux of NGC\\,7027 over a 25-year period.  NGC\\,7027 is one of the secondary flux calibrators used in the \\citet{Baars1977} flux scale. Its rapid evolution may raise doubts on its suitability as a calibrator.  However, the radio spectrum, due to free-free emission, can be well understood and modeled, and hence predicted.  We use this to investigate whether NGC\\,7027 can be used to improve the internal consistency of the radio flux calibration scale. ",
        "conclusions": "We have presented the results of a 25-year monitoring program of NGC\\,7027 at radio wavelengths. These results provide conclusive evidence for on-going evolution of this object. At low frequencies, where the nebula is optically thick, the radio flux is increasing approximately linearly with time, at a rate of $0.251\\pm0.015\\%\\,\\rm yr^{-1}$. This is shown to be directly caused by expansion of the nebula. From the observed increase and the known velocity field, we derive an expansion distance of $d=980\\pm100\\,$pc. The distance includes several correction factors: the most important one of these is the difference between the velocity of the gas and the velocity of the ionization front. The uncertainties in these correction factors dominate the final result: improved measurements of the rate of expansion will therefore not a-priori yield better distances.  At higher frequencies, where the nebula is optically thin, the data shows a steady decline of the radio flux, at a rate of $-0.145\\pm0.005\\%\\,\\rm yr^{-1}$.  We show that this can only be caused by evolution of the exciting star.  The number of ionizing photons is dropping, due to an {\\it increase} in stellar temperature and a decrease in stellar luminosity.  Evolutionary models of post-AGB stars are used to fit the observed decrease in ionizing photons. Using interpolated evolutionary tracks of \\citet{Bloecker1995}, we find a well-constrained mass for the central star of NGC\\,7027, of $0.655\\pm0.01\\,\\rm M_\\odot$. However, the alternative models of \\citet{Vassiliadis1994} evolve slower, and would indicate masses higher by about 0.04\\,M$_\\odot$. The discrepancy is related to a difference in the relation between envelope mass and effective temperature. For the Bl\\\"ocker model, we find a temperature increase of $135\\pm35 \\,\\rm K\\, yr^{-1}$, and a luminosity decrease of $- 0.075\\pm0.025 \\,\\rm \\% \\, yr^{-1}$.  The mass is consistent with those of white dwarfs in open clusters with initial masses around 3\\,M$_\\odot$, or slightly below.  The photoionization model of NGC\\,7027 reveals  small residuals which are frequency-dependent but constant with time. Correcting for these gives a large improvement on the fit. We propose that they are caused by inaccuracies in the relative fluxes of 3C286, based on the Baars flux scale. If this is correct, it is possible to improve the internal consistency of the Baars scale. A set of revised flux densities is proposed for 3C\\,286 which takes this correction into account. Using the models for NGC\\,7027 developed in this paper, we  calculate its time-dependent radio flux densities for the standard continuum bands used at the VLA, as well as some high frequency bands used by a number of new experiments."
    },
    "0801/0801.1052_arXiv.txt": {
        "abstract": "Because of their deep gravitational potential wells, clusters of galaxies retain all the metals produced by the stellar populations of the member galaxies. Most of these metals reside in the hot plasma which dominates the baryon content of clusters. This makes them excellent laboratories for the study of the nucleosynthesis and chemical enrichment history of the Universe. Here we review the history, current possibilities and limitations of the abundance studies, and the present observational status of X-ray measurements of the chemical composition of the intra-cluster medium. We summarise the latest progress in using the abundance patterns in clusters to put constraints on theoretical models of supernovae and we show how cluster abundances provide new insights into the star-formation history of the Universe. ",
        "introduction": "\\label{Introduction}  Clusters of galaxies are excellent astrophysical laboratories, which allow us to study the chemical enrichment history of the Universe. They have the deepest known gravitational potential wells which keep the metals produced in the stellar populations of the member galaxies within the clusters. About 70--90~\\% of the baryonic mass content of clusters of galaxies is in the form of hot ($10^7-10^8$~K) X-ray emitting gas \\citep{ettori1999}. In this hot intra-cluster medium (ICM) the dominant fraction of cluster metals resides. To the extent that the stellar populations where the cluster metals were synthesised can be considered representative, the metal abundances in the ICM provide constraints on nucleosynthesis and on the star formation history of the Universe.  It is remarkable that all the abundant elements, that were synthesised in  stars after the primordial nucleosynthesis, have the energies of their  K- and L-shell transitions in the spectral band accessible to modern X-ray telescopes.  Most of the observed emission lines in the ICM arise from the well understood  hydrogen and helium like ions and their equivalent widths can be, under the reasonable assumption of collisional equilibrium, directly converted into the elemental abundance of the corresponding element. Complications arising from optical depth effects, depletion into dust grains, extinction, non-equilibrium ionisation are minimal or absent. Therefore, abundance determinations of the hot ICM are more robust than those of stellar systems, \\ion{H}{ii} regions or planetary nebulae. This relatively uncomplicated physical environment makes the ICM an attractive tool for studies of the chemical enrichment (for the theoretical progress in metal enrichment processes see \\citealt{schindler2008} - Chapter~17, this volume). ",
        "conclusions": ""
    },
    "0801/0801.1578_arXiv.txt": {
        "abstract": "We investigated the nature of the hitherto unresolved elliptical infrared emission in the centre of the $\\sim$20\\,000\\,AU disc silhouette in M\\,17. We combined high-resolution {\\jhklm}~band imaging carried out with NAOS/CONICA at the VLT with [Fe\\,{\\sc ii}] narrow~band imaging using SOFI at the NTT. The analysis is supported by Spitzer/GLIMPSE archival data and by already published SINFONI/VLT Integral Field Spectroscopy data. For the first time, we resolve the elongated central infrared emission into a point-source and a jet-like feature that extends to the northeast in the opposite direction of the recently discovered collimated H$_2$ jet. They are both orientated almost perpendicular to the disc plane. In addition, our images reveal a curved southwestern emission nebula whose morphology resembles that of the previously detected northeastern one. Both nebulae are located at a distance of 1500\\,AU from the disc centre. We describe the infrared point-source in terms of a protostar that is embedded in circumstellar material producing a visual extinction of $60 \\leq A_V \\leq 82$. The observed $K\\!s$~band magnitude is equivalent to a stellar mass range of $2.8\\,M_{\\sun} \\leq M_\\star \\leq 8\\,M_{\\sun}$ adopting conversions for a main-sequence star. Altogether, we suggest that the large M\\,17 accretion disc is forming an intermediate to high-mass protostar. Part of the accreted material is expelled through a symmetric bipolar jet/outflow. ",
        "introduction": "The processes forming high-mass stars are still not well understood. Although recent theoretical models require the presence of and the accretion from a circumstellar disc (\\citealt{bonnell01}; \\citealt*{yorke02}; \\citealt*{krumholz07disks}), the direct observation of such a configuration is still missing.  During a systematic infrared study of the young star cluster NGC~6618 inside the H\\,{\\sc ii} region M\\,17 (Omega Nebula), \\citet{chini04nat} discovered an opaque silhouette, shaped like a flared disc at $JHK$ with a diameter of $\\sim$20\\,000\\,AU against the bright background of the H\\,{\\sc ii} region. It is associated with an optically visible hourglass-shaped nebula perpendicular to the silhouette plane. The optical spectrum of the nebula exhibits emission lines with blue-shifted absorption features indicating disc accretion.  \\citet{steinacker06} successfully modelled the optical depth and the photon scattering of the silhouette at $2.2\\,\\umu$m and obtained a large circumstellar disc seen at an inclination angle of $12^\\circ$ (almost edge-on). They found disc masses between 0.02 and $5\\,M_{\\sun}$ depending on the assumed dust model and the distance. The mass range is in good agreement with the distance-scaled values by \\citet{sako05} derived from their NIR (near infrared) extinction map. The distance has recently been determined to $2.1 \\pm 0.2$\\,kpc on the basis of about 50 spectroscopically classified early-type stars \\citep{hoffmeister08} ruling out earlier estimates below 2\\,kpc that treated several unresolved high-mass binaries as single stars. The new spectro-photometric distance agrees well with previous studies based on multi-colour photometry \\citep*[$2.2\\pm0.2$\\,kpc;][]{chini80} and radio data \\citep[$2.4\\pm^{0.4}_{0.5}$\\,kpc;][]{russeil03}.  Independent of its mass, the M\\,17 silhouette is the largest disc found to date that is exhibiting accretion activity. In fact, \\citet{nuernberger07} discovered a H$_2$ jet emerging from the disc centre, also suggesting ongoing accretion.  At the disc centre, \\citet{chini04nat} found an elliptical $K$~band emission of $240 \\times 450$\\,AU$^2$. Likewise, \\citet{sako05} describe the central object as an elongated compact infrared source. A faint emission of unknown nature is also visible in their [Ne\\,{\\sc ii}] narrow~band MIR (mid infrared) image at $12.8\\,\\umu$m.  Assuming that the central $K$~band flux found by \\citet{chini04nat} originates from a single star at a distance of 2.1\\,kpc, the stellar mass may range between a few and $45\\,M_{\\sun}$, depending on the extinction model \\citep{steinacker06}. \\citet{sako05} explained their results in terms of an intermediate-mass object with a stellar mass of $2.5$ to $8\\,M_{\\sun}$. Their analysis was part of a larger survey of candidate silhouette objects in M\\,17 \\citep{ito08}.  \\begin{figure*} \\centering \\resizebox{\\hsize}{!}{\\includegraphics[angle=270]{disk2.eps}} \\caption{\\label{f:hk_image} Infrared images of the inner region ($5\\arcsec \\times 5\\arcsec$) of the M\\,17 disc silhouette centred on the newly detected point-source at R.A. = $18^{\\rm h}20^{\\rm m}26\\fs191$ and Dec. = $-16^\\circ12'10\\farcs46$ (J2000). The inserts enlarge the central $0\\farcs6\\times0\\farcs6$. Left: $H$ image showing the symmetric bipolar nebula. The central point-source is not detected. To visualise the symmetry of the two nebulae, we have superimposed their contours after rotating them by $180^\\circ$ around the central point-source which we define as the symmetry centre; solid contours refer to the southwestern lobe, dashed contours represent the northeastern one. Right: $K\\!s$ image resolving the central NIR emission into a point-source and a fainter tail extending to the northeast. The hourglass structure above and below the disc hosts the bipolar nebula, of which the southwestern lobe is not detected. The contours denote the continuum-free H$_2$ $1-0$\\,S(1) line emission at $\\lambda = 2.12\\,\\umu$m as published by \\citet{nuernberger07}; the dotted box delineates the extent of the H$_2$ image. The dashed blue lines mark the suggested jet axis as determined by the orientation of the detected tail with respect to the point-source.} \\end{figure*}  So far, both the mass and the temperature of the forming star within this spectacular disc remain under discussion. In this work, we present new data that provide further insight into the nature of the central protostellar object and constrain the possible mass range. We also discuss the detection of a possible counter-jet headed in the opposite direction of the H$_2$ jet and a curved southwestern emission nebula that is geometrically symmetric to the northeastern one. ",
        "conclusions": "With the new observations, we have further clarified the nature of the emission in the centre of the large M\\,17 disc and its associated nebula in the following aspects:  \\begin{enumerate} \\item The faint central elliptical emission is resolved for the first time into a point-source and a tail that extends to the northeast. This morphology suggests a single protostar accompanied by a jet.  \\item The jet-like feature is headed in the opposite direction of a collimated H$_2$ jet.  \\item A southwestern emission nebula is turning the outflow pattern from the disc centre into a fairly symmetric morphology. The intensity ratio of the bipolar emission, however, is highly asymmetric, probably due to the inclination of the dusty disc.  \\item The bipolar nebula might be scattered light which is reddened by extinction caused by the disc material.  \\item Various estimates of the extinction towards the embedded protostar yields a range of $60\\leq A_V \\leq 82$.  \\item The observed 2.2\\,$\\umu$m flux is equivalent to a protostar with a mass range of $2.8\\,M_{\\sun} \\leq M_\\star \\leq 8\\,M_{\\sun}$. \\end{enumerate}  The new data strengthen the evidence for an intermediate to high-mass protostellar object being formed by accretion."
    },
    "0801/0801.4441_arXiv.txt": {
        "abstract": "The millimeter-wave rotational emission lines ($4_{04}-3_{03}$ and $5_{05}-4_{04}$) of protonated carbon dioxide, HCO$_2^+$(HOCO$^+$), has been detected toward the low-mass class 0 protostar IRAS 04368$+$2557 in L1527 with the IRAM 30 m telescope.  This is the first detection of HCO$_2^+$ except for the Galactic Center clouds.  The column density of HCO$_2^+$ averaged over the beam size (29$^{\\prime\\prime}$) is determined to be $7.6 \\times 10^{10}$ cm$^{-2}$, assuming the rotational temperature of 12.3 K.  The fractional abundance of gaseous CO$_2$ relative to H$_2$ is estimated from the column density of HCO$_2^+$ with an aid of a simplified chemical model.  If the HCO$_2^+$ emission only comes from the evaporation region of CO$_2$ near the protostar ($T$\\raisebox{0.4ex}{$>$}\\hspace{-0.75em}\\raisebox{-.7ex}{$\\sim$}50 K), the fractional abundance of CO$_2$ is estimated to be higher than $6.6 \\times 10^{-4}$.  This is comparable to the elemental abundance of carbon in interstellar clouds, and hence, the direct evaporation of CO$_2$ from dust grain is unrealistic as a source of gaseous CO$_2$ in L1527.  A narrow line width of HCO$_2^+$ also supports this.  On the other hand, the fractional abundance of CO$_2$ is estimated to be $2.9 \\times 10^{-7}$, if the source size is comparable to the beam size.  These results indicate that gaseous CO$_2$ is abundant even in the low-mass star-forming region.  Possible production mechanisms of gaseous CO$_2$ are discussed.\\\\ ",
        "introduction": "Carbon dioxide (CO$_2$) is an abundant and important constituent in planetary atmospheres as well as cometary comae, and is considered to be a key species which links between interstellar chemistry and planetary chemistry.  So far solid CO$_2$ was found ubiquitously toward intermediate- to high mass star-forming regions through observations of its vibrational bands with Infrared Space Observatory (ISO) and Spitzer Space Telescope (e.g. de Graauw et al. 1996; Grakines et al. 1999; Nummelin et al. 2001; Sonnentrucker et al. 2006).  A typical abundance of solid CO$_2$ relative to H$_2$ is as high as 10$^{-5}$ to 10$^{-6}$ (Gerakines et al. 1999).  Furthermore, Dartoris et al. (2005) discovered huge amount of solid CO$_2$ even toward the low-mass class 0 protostar, L723, by observations with Spitzer Space Telescope.  Gaseous CO$_2$ was also detected in hot regions around several massive young stellar objects, where a typical abundance relative to H$_2$ was reported to be an order of 10$^{-7}$ (e.g. van Dishoeck et al. 1996; Boonman et al. 2003b; Sonnentrucker et al. 2006).  These observations clearly indicate that CO$_2$ is a major chemical component in molecular clouds.  However, the formation processes of CO$_2$ are not well understood.  In particular, a contribution of the gas phase production of CO$_2$ is quite uncertain.  Although the gas-phase chemical models predict the CO$_2$ abundance relative to H$_2$ of about 10$^{-7}$ (e.g. Lee et al. 1997), this has not been confirmed observationally in cold molecular clouds and low-mass star-forming regions.  A difficulty in measuring the abundances of gaseous CO$_2$ in these regions comes from a lack of a permanent dipole moment of CO$_2$.  Because of this, radio observations of the rotational emission lines are impossible.  Infrared observations of its vibration-rotation transitions would also be difficult because of insufficient brightness of background infrared sources or even of a lack of the sources.  An alternative way is to observe the rotational emission lines of protonated carbon dioxide, HCO$_2^+$, in the millimeter-wave region, because its abundance is directly related to the abundance of CO$_2$.  However, HCO$_2^+$ has never been detected in any cold molecular clouds and any star-forming regions except for the Galactic Center clouds like Sgr B2 and Sgr A (Thaddeus et al. 1981; DeFrees et al. 1982; Minh et al. 1991).  It has been thought that HCO$_2^+$ is produced only under a specific condition in the Galactic Center region, where shock chemistry would play an important role (e.g. Minh et al. 1988; Charnley et al. 2000).  An important breakthrough to this situation came from our preliminary line survey toward a low-mass protostar IRAS 04368+2557 in L1527 with the Nobeyama 45 m radio telescope\\footnote{Nobeyama Radio Observatory is a branch of the National Astronomical Observatory of Japan, National Institutes of Natural Sciences, Japan} (NRO 45 m), by which peculiar carbon-chain chemistry in this source (Sakai et al. 2007, 2008a, 2008b) was being explored.  With a very sensitive observation, we recently found a weak emission line at the frequency of the $4_{04}-3_{03}$ line of HCO$_2^+$.  IRAS 04368+2557 is a low-mass protostar in a transient phase from class 0 to class I, which has extensively been studied by many researchers (e.g. Ohashi et al. 1997; Hogerheijde et al. 1997, 1998).  Hence, this discovery, if confirmed, is very important for understanding of the behavior of CO$_2$ during chemical evolution from molecular clouds to protoplanetary disks.  Then we carried out more sensitive observations with the IRAM 30 m telescope (IRAM 30 m) and firmly confirmed our detection of HCO$_2^+$.  In this Letter, we report the first detection of HCO$_2^+$ in the star-forming region, L1527, and discuss its astrochemical implication.\\\\ ",
        "conclusions": "\\subsection{\\it Fractional Abundance of CO$_2$} In this section, we derive the fractional abundance of the gaseous CO$_2$ with an aid of a simplified chemical model.  The abundance of HCO$_2^+$ is related to that of CO$_2$ through the following reactions; $$ {\\rm H}_3^+ + {\\rm CO}_2  \\rightarrow  {\\rm HCO}_2^+ + {\\rm H}_2:  \\eqno(1) $$ $$ {\\rm HCO}_2^+ + {\\rm CO}  \\rightarrow  {\\rm HCO}^+ + {\\rm CO}_2,  \\eqno(2) $$ where H$_3^+$ is formed by the cosmic-ray ionization of H$_2$ and subsequent reaction of H$_2^+$ + H$_2$, and is destructed by $$ {\\rm H}_3^+ + {\\rm CO}  \\rightarrow  {\\rm HCO}^+ + {\\rm H}_2.  \\eqno(3) $$ The electron recombination is less important as destruction processes of HCO$_2^+$ and H$_3^+$, and is ignored.  When we assume the steady-state condition, ratio of the number density, $n$(HCO$_2^+$)/$n$(CO$_2$), can be written as $$ \\frac{n({\\rm HCO}_2^+)}{n({\\rm CO}_2)} \\simeq \\frac{{\\it k}_1 \\zeta n({\\rm H}_2)}{{\\it k}_2{\\it k}_3n({\\rm CO})^2} = 1.9 \\times 10^{-8} \\times \\frac{1}{f_{\\rm CO}^{ 2} n({\\rm H}_2)},  \\eqno(4) $$ where $k_1$, $k_2$, and $k_3$ represent the rate coefficients for reactions (1), (2), and (3), respectively, and $\\zeta$ represents the cosmic ray ionization rate, $1.3 \\times 10^{-17}$ s$^{-1}$ (e.g. Aikawa et al. 2008).  Here the unit of $n$(X) is cm$^{-3}$, and the fractional abundance of CO relative to H$_2$ is denoted as $f_{\\rm CO}$.  We employ the $k_1$, $k_2$, and $k_3$ values of $1.9 \\times 10^{-9}$ cm$^3$ s$^{-1}$ (Burt et al. 1970), $7.8 \\times 10^{-10}$ cm$^3$ s$^{-1}$ (UMIST data base), and $1.7 \\times 10^{-9}$ cm$^3$ s$^{-1}$ (Kim et al. 1975), respectively.  The numerical factor in equation (4) could be larger or smaller by a factor of 3.  It should be noted that the $n$(HCO$_2^+$)/$n$(CO$_2$) ratio is inversely proportional to the H$_2$ density, as far as $f_{\\rm CO}$ is constant.  In other words, the fractional abundance of CO$_2$ relative to H$_2$ ($f_{{\\rm CO}_2}$) is simply proportional to $n$(HCO$_2^+$); $$ f_{{\\rm CO}_2} \\simeq 5.4 \\times 10^7 \\times f_{\\rm CO}^{ 2} \\times n({\\rm HCO}_2^+).  \\eqno(5) $$ If CO$_2$ exists only within the sphere with a radius of $R$ around the protostar, with the constant $f_{{\\rm CO}_2}$, $n$(HCO$_2^+$) is also constant and is related to the beam averaged column density $N_{ave}$(HCO$_2^+$) as $$ N_{ave}({\\rm HCO}_2^+) = n({\\rm HCO}_2^+) \\times 2R \\int _{0}^{(\\frac{R}{R_{b}})^2}  e^{-t}\\sqrt{1-(\\frac{R_{b}}{R})^2t} \\ dt,  \\eqno(6) $$ assuming the Gaussian-shaped telescope beam with an e-fold beam radius of $R_b$.  Note that the right-hand side of equation (6) approaches to $2Rn$(HCO$_2^+$) for $R \\gg R_b$ and to $n$(HCO$_2^+$)$(4 \\pi R^3/3)/(\\pi R_b^2)$ for $R \\ll R_b$.  With equation (6), we can derive $f_{{\\rm CO}_2}$ directly from the observed $N_{ave}$(HCO$_2^+$) without any assumptions on the H$_2$ density and the H$_2$ column density.  We calculate $f_{{\\rm CO}_2}$ as a function of $R$, as shown in Figure 2, where we employ the $f_{\\rm CO}$ value of $3.9 \\times 10^{-5}$ reported by J$\\o$rgensen et al. (2002).  The half-power beam radius ($R_b\\sqrt{ln2}$) for IRAM 30 m is 2000 AU, where the distance to L1527 is assumed to be 140 pc.  According to Maret et al. (2004), the size of the hot region where CO$_2$ can be evaporated (\\raisebox{0.4ex}{$>$}\\hspace{-0.75em}\\raisebox{-.7ex}{$\\sim$}50 K) is estimated to be smaller than 140 AU for L1527.  If gaseous CO$_2$ traced by our HCO$_2^+$ observation are confined to this small radius, $f_{{\\rm CO}_2}$ should be evaluated to be higher than $6.6 \\times 10^{-4}$ by using Figure 2.  In the ice evaporation region(\\raisebox{0.4ex}{$>$}\\hspace{-0.75em}\\raisebox{-.7ex}{$\\sim$}100 K), H$_2$O and NH$_3$ would become very abundant and contribute to further destruction of HCO$_2^+$ in the innermost region.  This effect can be involved as an effectively larger $f_{\\rm CO}$ value in equations (4) and (5).  Then $f_{{\\rm CO}_2}$ is comparable to the elemental abundance of carbon in interstellar clouds ($7 \\times 10^{-4}$).  On the other hand, Furlan et al. (2007) recently reported the mass fractional abundance of solid CO$_2$ relative to the gas in L1527 to be $1 \\times 10^{-4}$ with the Spitzer observation, which corresponds to the fractional abundance of solid CO$_2$ relative to H$_2$ of $5 \\times 10^{-6}$.  This is much lower than the $f_{{\\rm CO}_2}$ value derived from our observation.  Hence, the direct evaporation of CO$_2$ is very unrealistic as a major source of gaseous CO$_2$ traced by HCO$_2^+$ in L1527.  This would further be supported by the narrow line width (0.57 km s$^{-1}$) of HCO$_2^+$ in comparison with the line width of the very high excitation line of HC$_5$N ($J=32-31$, 85.2 GHz, $E_u$=67 K) observed with IRAM 30 m in L1527 (Sakai et al. in preparation).  If we assume $f_{{\\rm CO}_2}$ of $\\sim 5 \\times 10^{-7}$, which is close to the value reported in Orion-KL IRc2 by the infrared observation (Boonman et al. 2003a), $R$ is derived to be 1600 AU.  On the other hand, $f_{{\\rm CO}_2}$ is still as high as $2.9 \\times 10^{-7}$, even if HCO$_2^+$ is extended over the beam size of IRAM 30 m.  Therefore, substantial amount of CO$_2$ exists even in the low temperature region where thermal evaporation of CO$_2$ from grain mantles is very inefficient.  It should be noted that the interaction between outflows and ambient gas would not be responsible for the CO$_2$ production in L1527, because the line width of HCO$_2^+$ is narrow and no velocity shift is observed.  \\subsection{\\it Origin of Gaseous CO$_2$} As described above, most of gaseous CO$_2$ in L1527 does not originate from direct evaporation of solid CO$_2$, but from gas-phase reactions.  Since the gas phase production of CO$_2$ is not well understood observationally (e.g. van Dishoeck et al. 2004), we discuss a few possibilities.  First we compare our result with the chemical model calculation by Lee et al. (1997).  According to their 'new standard' model, the gas phase abundance of CO$_2$ is expected to be $2.8 \\times 10^{-7}$ and $1.4 \\times 10^{-7}$ at $10^{5.5}$ yr and the steady state, respectively, for the H$_2$ density of 10$^5$ cm$^{-3}$.  This is almost consistent with our result in L1527, if the CO$_2$ distribution is extended over the beam size of IRAM 30 m.  Secondly, CO$_2$ might be produced in a hot region through a reaction of OH and CO.  Since this reaction has activation barrier of 176 K (UMIST data base; See also Talbi and Herbst 2002 for the potential energy surface.), it is efficient in a small region near the protostar, like a hot core or a hot corino.  In such a hot region, the abundance of H$_2$CO is enhanced due to mantle evaporation, which would supply HCO through the electron recombination reaction of H$_2$COH$^+$ formed from H$_2$CO and H$_3^+$.  Then, CO$_2$ will be produced through the HCO $+$ O reaction (e.g. Snyder et al. 1985).  However, these high temperature mechanisms occur in a much smaller region than the CO$_2$ evaporation region, and hence, they can not be major contributors.  Since L1527 shows warm carbon-chain chemistry (WCCC) (Sakai et al. 2008a), we finally consider a possible CO$_2$ formation related to the WCCC.  In the WCCC, CH$_4$ evaporated from grain mantles drives the carbon-chain chemistry.  At the same time, HCO would be produced by various reactions, which successively forms CO$_2$ by the reaction with the atomic oxygen.  The warm region, where CH$_4$ can be evaporated, is substantially larger than the region where CO$_2$ can be evaporated, and hence, CO$_2$ would be distributed over the size of WCCC (Sakai et al. 2008a).  Rather compact distribution of HCO$_2^+$ inferred from the beam dilution effect, if correct, favors this scenario rather than the production of CO$_2$ in a cold cloud.  In this relation, the HCO ($1_{01}-0_{00}$, 86.7 GHz) line is detected toward L1527 in our preliminary line survey, which may also support the above idea.  Furthermore, an enhancement of gaseous CO$_2$ in the warm region, where CH$_4$ is evaporated, is predicted in a chemical model of a dynamically evolving cloud by Aikawa et al. (2008, Figs. 6 and 10).  In order to establish the formation pathway of CO$_2$, distribution of HCO$_2^+$ in L1527 is a key.  A high spatial resolution observation with an interferometer would be necessary.  By comparing the result with the detailed chemical model calculations, we will be able to investigate the contributions of cold gas phase chemistry and WCCC to the formation of CO$_2$.  Sensitive observations of HCO$_2^+$ toward various clouds would also be needed to understand the origin of gaseous CO$_2$."
    },
    "0801/0801.4488_arXiv.txt": {
        "abstract": "We investigate the diffusive shock acceleration in the presence of the non-resonant streaming instability introduced by Bell \\cite{bell04}. The numerical MHD simulations of the magnetic field amplification combined with the analytical treatment of cosmic ray acceleration permit us to calculate the maximum energy of particles accelerated by high-velocity supernova shocks. The estimates for Cas A, Kepler, SN1006, and Tycho historical supernova remnants are given. We also found that the amplified magnetic field is preferentially oriented perpendicular to the shock front downstream of the fast shock. This explains the origin of the radial magnetic fields observed in young supernova remnants. ",
        "introduction": "The instabilities produced by energetic particles are important phenomena accompanying the diffusive shock acceleration (Krymsky \\cite{krymsky77}; Axford et al. \\cite{axford77}; Bell \\cite{bell78}; Blandford and Ostriker \\cite{blandford78}) of cosmic rays in supernova remnants (SNRs). The scattering of energetic particles both upstream and downstream of a supernova shock is supplied by magnetic inhomogeneities existing in the shock vicinity. It was suggested that this may be the result of a resonant streaming instability that develops due to the presence of diffusive streaming of accelerated particles (Bell \\cite{bell78}). The total magnetic field may be amplified if the energy of unstable magnetohydrodynamic (MHD) waves becomes comparable with the energy of the background magnetic field.  Such an amplification seems quite possible because quasilinear theory of the resonant streaming instability allows it (see e.g. McKenzie \\& V\\\"olk \\cite{mckenzie82}).  The presence of amplified magnetic fields in young SNRs is well established now. It was indicated by the radio observations of SNRs and was attributed to the Rayleigh-Taylor instability of the contact discontinuity between the gas of supernova ejecta compressed at the reverse shock  and the circumstellar gas compressed at the forward shock  (Gull \\cite{gull75}). The radial magnetic fields inferred from measurements of the radio-polarization in young shell type SNRs (see e.g. Milne \\cite{milne87}) may appear as this instability develops.  However, the discovery of thin X-ray filaments coinciding with the position of the forward shock in young galactic SNRs (Gotthelf et al. \\cite{gotthelf01}, Hwang et al. \\cite{hwang02}, Vink \\& Laming \\cite{vink03}, Long et al. \\cite{long03}, Bamba et al. \\cite{bamba03}, Bamba et al. \\cite{bamba05}) has led to the conclusion that the magnetic field is amplified just at the forward shock. This conclusion does not depend on the nature of the mechanism which produces such filaments: the fast synchrotron cooling of electrons accelerated at the forward shock (see e.g. the analysis by Berezhko et al. \\cite{berezhko02}), or the dissipation of the MHD turbulence and the corresponding decrease of the magnetic field, as was suggested by Pohl et al. \\cite{pohl05}.  Recently Bell \\cite{bell04}, using the dispersion relation for collisionless MHD waves derived by Achterberg \\cite{achterberg83}, found a new regime of a non-resonant streaming instability. In the presence of the strong electric current of accelerated particles a non-oscillatory purely growing MHD mode appears at spatial scales smaller than the gyroradius of the particles. Bell \\cite {bell04} performed MHD simulations and showed that the magnetic field may be strongly amplified.  This instability is  investigated in more detail in our companion paper (Zirakashvili et al. \\cite{zirakashvili08}, Paper I) via numerical MHD simulations. Here we  combine these simulations with the analytical treatment of diffusive acceleration at a plane steady-state shock. It allows us to estimate the maximum energy of the accelerated particles and to obtain the value of the amplified magnetic field.  The paper is organized as follows.  The analytical model of  diffusive acceleration at the plane shock is considered in Sect.2. The MHD simulations are described in Sect.3 and 4. The maximum energies of accelerated particles are estimated in Sect.5.  Sect. 6 contains the discussion of obtained results. The summary is given in the last Sect.7. ",
        "conclusions": "We have investigated the acceleration of particles at the fast plane parallel shock. The generation of MHD turbulence by the non-resonant streaming instability (Bell \\cite{bell04}) was taken into account. We solved this problem using only first principles and with a minimum of simplifying assumptions. We combined the analytical solution for the particle acceleration and the numerical MHD calculations for the evolution of the MHD turbulence. The following results were obtained:  1) For the relatively fast shocks when the  condition (22) is satisfied, the particles at the high-energy end of the spectrum are scattered by small-scale random magnetic fields generated by the non-resonant streaming instability. Their spectrum has the universal shape given by Eq. (8).  2) The MHD turbulence is mainly generated by the streaming of run-away particles at large distance from the shock. The level of the MHD turbulence is the  highest in the shock vicinity. The accelerated particles are concentrated in the same region. This means that the acceleration of particles can be considered in the one-dimensional approximation even for a three-dimensional system. The characteristic width of the particle distribution in our simulations is not larger than $0.1\\ L $ (see Fig.3 and 10th line of Table 1).  3) It is important, that the non-resonant instability produces strong density fluctuations upstream of the shock (see Sect.4 and Paper I). These fluctuations produce a strong deformation of the shock front and fast vortex motions downstream of the shock. That is why the magnetic amplification in the shock transition region is not reduced to a simple compression of the magnetic field in the direction perpendicular to the shock front. The magnetic field is also stretched by the flow motions in the direction perpendicular to the shock front. As a result, the magnetic field component which is perpendicular to the shock front is a factor of 1.4 larger than the parallel components downstream of the shock. This naturally explains the preferable radial orientation of magnetic fields in young SNRs.  The characteristic time 1 year of the shock deformation and of the corresponding MHD fluctuations found here (see Sect.6) is of particular interest in the light of the last results on variability of X-ray emission observed in RXJ1713 SNR (Uchiyama et al. \\cite{uchiyama07}).  4) The dissipation of the magnetic field downstream of the shock is relatively slow in our simulation. If so, the origin of X-ray filaments, observed in young SNRs, is related to the fast synchrotron cooling of accelerated electrons but not to the decay of the MHD turbulence.  5) The values of the calculated amplified magnetic field are similar to those, observed in historical SNRs, if the energy flux of the run-away particles is not low: $\\eta _{esc}>0.05$ (see Table 2).  6) The magnetic field growth is only linear in time for the fast shocks with the normalized velocities higher than about ten thousand km s$^{-1}$. This reduces the maximum energy of accelerated particles compared to the case of the exponential growth. For these shocks the energy of the magnetic field amplified upstream is a small fraction $\\sqrt{V_a/c}$ of the energy density of the highest energy particles (see Eq.(20)). The magnetic field amplification is relatively weak for slow shocks with the normalized velocities 1230 km s$^{-1}$ and smaller (see Eq.(18)).  7) We calculated numerically the maximum energy of accelerated particles (see Fig.2). This energy may be described by the analytical formula (19). The maximum energies of particles are higher than the energies obtained in the Bohm limit in the background magnetic field (see the 7th line of the Table 1) but lower than the energies obtained using the Bohm limit in the amplified field. The significantly long time of the random magnetic field growth is the main  factor that limits the maximum energy of accelerated particles. We calculated the maximum energies for four historical SNRs (Table 2).  8) Using the last result we found that the maximum energy of cosmic ray protons accelerated by Ia and IIP supernovae is about $100\\div 300$ TeV. Only Ib/c and presumably IIb supernovae may accelerate protons up to PeV energies. Since it is expected that the explosion rate is the highest for IIP supernovae, we should observe a change of the slope of the galactic cosmic ray spectrum at the energies of the order of 100 TeV.  9) The MHD turbulence generated by the non-resonant streaming instability has a non-zero magnetic helicity. The helical Lorentz force produces corresponding plasma motions and the mean electric field ${\\bf E}_0$ that is in the opposite direction to the electric current of energetic particles. This electric field modifies the cosmic ray transport equation (see Paper I). This effect is significant for very fast shocks with velocities larger than 30 thousands km s$^{-1}$. The presence of this field should result in the steepening of the spectrum of particles which produce the  non-resonant instability (presumably nucleons) and the flattening of the spectrum of oppositely charged particles (presumably electrons)."
    },
    "0801/0801.1264_arXiv.txt": {
        "abstract": "It is been known for more than a decade that BALQSOs (broad absorption line quasars) are highly attenuated in the X-ray regime compared to other quasars, especially in the soft band ( $< $ 1 keV). Using X-ray selection techniques we have found ``soft X-ray loud\" BALQSOs that, by definition, have soft X-ray (0.3 keV) to UV ($3000 \\AA$) flux density ratios that are higher than typical nonBAL radio quiet quasars. Our sample of 3 sources includes one LoBALQSO (low ionization BALQSO) which are generally considered to be the most highly attenuated in the X-rays. The three QSOs are the only known BALQSOs that have X-ray observations that are consistent with no intrinsic soft X-ray absorption. The existence of a large X-ray luminosity and the hard ionizing continuum that it presents to potential UV absorption gas is in conflict with the ionization states that are conducive to line driving forces within BAL winds (especially for the LoBALs). ",
        "introduction": "One of the biggest challenges to our understanding of broad absorption line quasars (BALQSOs, hereafter) is how lithium like species (those species producing the resonant UV absorption lines) can form within an out-flowing wind in the presence of the hard ionizing continuum of a quasar (EUV and X-ray). Thus, it was no surprise that BALQSOs turned out to be X-ray weak due to thick absorption columns \\citep{gre96}. The necessity of this X-ray absorption screen between the central quasar X-ray source and the BAL wind was anticipated by \\citet{mur95}. Every deep X-ray observation of a BALQSO has shown significant absorption with the least absorbed sources having neutral hydrogen absorption columns $N_{H} \\gtrsim 10^{22} \\mathrm{cm}^{-2}$ and most BALQSOs have absorption columns of $ 10^{23} \\mathrm{cm}^{-2}\\leq N_{H} \\leq 10^{25} \\mathrm{cm}^{-2}$ \\citep{pun06}. As a qualifier, it is possible that the modest radio jet seen in some BALQSOs can also contribute to the X-ray flux and some of these sources might not appear X-ray suppressed due to a secondary source of X-rays from the relativistic jet $\\sim$ 10 - 100 pc from the central quasar \\citep{bro05}. Thus, we make a distinction between radio quiet (RQBALQSOs), $\\log R<1$, and radio loud BALQSOs (RLBALQSOs), $\\log R>1$, where R is the k-corrected 5 GHz to $2500\\AA$ flux density ratio (the RQBALQSOs considered here have $\\log R<0$). The few RQBALQSOs with modest absorption, UM425 and CSO755, $N_{H} \\gtrsim 10^{22} \\mathrm{cm}^{-2}$ have fairly typical hard X-ray to UV flux density ratios for radio quiet quasars, but are highly attenuated in the soft X-rays \\citep{ald03,she05,mur95}. To this date, all known RQBALQSOs are highly absorbed in the soft X-rays. The lone exception is the X-ray spectrum of IRAS 07598+059 that is consistent with no absorption. However, the X-ray flux relative to the UV is suppressed by a factor of more than 20 compared to a typical radio quiet quasar, thus it is believed that the central quasar X-ray source is completely obscured and we are seeing a secondary weak source, likely the central source seen in reflection off an electron scattering mirror \\citep{ima04}. \\par We have searched the ROSAT database to look for evidence of soft X-ray loud BALQSOs. Previously, the only BALQSO ROSAT detection in the literature is 1245-067 and it is highly absorbed, $N_{H} \\approx 10^{23} \\mathrm{cm}^{-2}$ \\citep{gre96}. We have found approximately 40 likely detections of BALQSOs by ROSAT and are preparing a catalog of these. However, most of these either are not necessarily strong compared to the UV flux or have poor photon statistics, so it is difficult to determine the X-ray luminosity with much certainty. In sorting through the catalog, we have segregated three AGN that are the most convincing examples of RQBALQSOs that are soft X-ray loud (defined by soft X-ray (0.3 keV) to UV ($3000 \\AA$) flux density ratios that are higher than typical nonBAL radio quiet quasars), one of which has low ionization UV absorption lines (LoBALQSO). This is even more surprising since the the LoBALQSOs require the maximal amount of X-ray screening, otherwise the gas will be over-ionized for the formation of Li-like low ionization species \\citep{mur95}. These BALQSOs could either be intrinsically exceptional objects or objects that are fortuitously configured to reveal some unforseen physical features of the QSO central engine. ",
        "conclusions": "In this Letter, we describe 3 BALQSOs that have soft X-ray to UV ratios that are larger than typical quasars (see the $\\alpha_{os}$ values in Table 2). Hence, the designation as soft X-ray loud. However, the $\\alpha_{ox}$ values are not exceptional for quasars due to the steep X-ray spectral indices \\citep{str05}. These are also the first known BALQSOs with X-ray absorption consistent with pure Galactic absorption. The tentative discovery of soft X-ray loud BALQSOs is completely unexpected based on theoretical treatments of BALQSOs. The conundrum posed by this new class of AGN is how can a BAL region coexist with a powerful source of X-rays, since even a modest flux of X-rays will over-ionize the gas making it impossible to form Li-like atoms \\cite{mur95}. On a speculative note, a relativistic X-ray jet beamed towards earth and away from the BAL gas is a possibility \\citep{pun99,pun00}. It is now known that there are BALQSOs with relativistic jets beamed toward earth and a jet similar to the one in Figure 1 of \\cite{gho07} that is radio weak, but X-ray strong would conform to the properties of the BALQSOs in Tables 1 and 2. There are extragalactic jets that have X-ray and radio properties in-line with theses sources. Extreme high frequency peaked BL Lac objects such as PG1553+113, at $z \\approx 2$, would have an apparent X-ray luminosity of $\\gtrsim 10^{45}$ ergs/s, a steep X-ray spectrum and would have a radio flux at 1.4 GHz of about 0.1 - 0.5 mJy in agreement with the properties of the BALQSOs presented here \\citep{ost06}."
    },
    "0801/0801.3782_arXiv.txt": {
        "abstract": "We present photometric analysis of deep mid-infrared observations obtained by Spitzer/IRAC covering the fields Q1422+2309, Q2233+1341, DSF2237a,b, HDFN, SSA22a,b and B20902+34, giving the number counts and the depths for each field. In a sample of 751 LBGs lying in those fields, 443, 448, 137 and 152  are identified at 3.6$\\mu$m, 4.5$\\mu$m, 5.8$\\mu$m, 8.0$\\mu$m IRAC bands respectively, expanding their spectral energy distribution to rest-near-infrared and revealing that LBGs display a variety of colours. Their rest-near-infrared properties are rather inhomogeneous, ranging from those that are bright in IRAC bands and exhibit $[R]-[3.6]>1.5$ colours to those that are faint or not detected at all in IRAC bands with $[R]-[3.6]<1.5$ colours and these two groups of LBGs are investigated. We compare the mid-IR colours of the LBGs with the colours of star-forming galaxies and we find that LBGs have colours consistent with star--foming galaxies at z$\\sim$3. The properties of the LBGs detected in the 8$\\mu$m IRAC band (rest frame K-band) are examined separately, showing that they exhibit redder$[R]-[3.6]$ colours than the rest of the population and that although in general, a multi--wavelength study is needed to reach more secure results, IRAC 8$\\mu$m band can be used as a diagnostic tool, to separate high z, luminous AGN dominated objects from normal star-forming galaxies at z$\\sim$3. ",
        "introduction": "Observation and study of high-redshift galaxies is essential to constrain the history of galaxy evolution and give us a systematic and quantitative picture of galaxies in the early universe, an epoch of rigorous star and galaxy formation. Large samples of high-z galaxies that have recently become available, play a key role to that direction and have revealed a zoo of different galaxy populations at z$\\sim$3. There are various techniques for detecting high-z galaxies involving observations in wavelengths that span from optical to far-IR. Among these, the three most efficient are 1)sub-mm blank field observations, using the sub-mm Common User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope (JCMT) (e.g., Hughes et al. 1998) or the Max Plank Millimeter Bolometer array (MAMBO, e.g., Bertoldi et al. 2000), taking advantage of the strong negative K-correction effect and revealing the population of the sub-mm galaxies at z$>$2 (Chapman et al. 2000, Ivison et al. 2002, Smail et al. 2002) and 2)the U-band-dropout technique (Steidel $\\&$ Hamilton 1993, Steidel et al. 1999, 2003 Franx et al. 2003, Daddi et al. 2004), sensitive to the presence of the 921{\\AA} break, designed to select z${\\approx}$3 galaxies and revealing the population of the Lyman Break Galaxies (LBGs), 3) the Near-IR colour  (J$-$K $>$ 2.3 (Franx et al. 2003) and BzK ((z$-$K)$_{\\rm AB}$ $-$ (B$-$z)$_{\\rm AB}$$ >-$0.2) (Daddi et al. 2004) selection for old and dusty galaxies at 2$<$z$<$3 (Distant Red Galaxies (DRGs) and BzK galaxies).  With recent deep optical and IR observations of these 3 types of galaxies, there has been considerable progress in understanding the relation between LBGs, SMGs, DRGs and BzK galaxies. Chapman et al. 2005 confirmed that some of the SCUBA galaxies have rest-frame-UV colours typical of the LBGs. Using IRAC and MIPS observations, Huang et al. 2006 has shown that LBGs detected in MIPS 24$\\mu$m band, (Infrared Luminous LBGs) and cold SCUBA sources share similar [8.0]$-$[24] colours, while Rigopoulou et al. 2006 suggest that ILLBGs and SCUBA galaxies tend to have similar stellar masses and dust amount. A possible scenario is one in which sub-mm galaxies and LBGs form a continuum of objects with SMGs representing the reddest dustier and more intensively star-forming LBGs, but further investigation is required to establish a more secure link between these two populations.  LBGs constitute at the moment the largest and most well studied galaxy population at  z$\\sim$3 (Steidel et al. 2003). Based on observations of the UV continuum emission, the predicted mean star formation rate of the LBGs is 20--50$M_{\\odot}$/yr (assuming $H_{o}$$=70 Kms^{-1}Mpc^{-1}$ and $q_{o}$=0.5, Pettini et al. 2001). This star formation rate increases significantly if corrected for dust attenuation, to a mean value of $\\sim$100$M_{\\odot}$/yr. Correction for dust attenuation must be taken into account as there is clear evidence for the presence of significant amounts of dust in the galactic medium of LBGs (e.g., Sawicki $\\&$ Yee 1998, Vijh 2003).  To investigate the amounts of dust in the LBGs, various techniques and observations spanning from optical to X-rays have been used. The techniques range from studies of optical line ratios (Pettini et al. 2001) to formal fits of the overall SED of LBGs based on various star formation history scenarios (e.g., Shapley et al. 2001, 2003, Papovich et al. 2001) and X-ray stacking studies (e.g., Nandra et al. 2002, Reddy$\\&$Steidel 2004). All approaches agree that LBGs with higher star formation rates contain more dust (e.g., Adelberger $\\&$ Steidel et al. 2000, Reddy et al. 2006).  One of the most important properties of the population is the stellar mass of the galaxies. Stellar masses play a central role to the understanding of the evolution of the LBGs, as the co-moving stellar mass density at any redshift is the integral of the past star-forming activity and suffers fewer uncertainties than the star formation rate. Most of the attempts to estimate the stellar content of the LBGs were based on ground based observations. As the observed optical band corresponds to rest-UV for galaxies at z$\\sim$3, most of the light emitted from the galaxy at these wavelengths is radiated from young and massive OB stars. Therefore, optical observations cannot be regarded as a robust tool to probe the stellar mass of the galaxy as they are sensitive only to recent star formation episodes rather than the stellar population that has accumulated over the galaxy's life-time. Observations at longer wavelengths, i.e., rest-near-infrared (e.g. Bell \\& deJong 2001, where the bulk of the stellar population radiates are essential to constrain the stellar content of the z$\\sim$3 galaxy population.  With the advent of the Spitzer Space Telescope (Werner et al. 2004) we now have access to longer wavelengths. The four IRAC bands on board \\textsl{Spitzer} (Fazio et. al 2004) and more particularly the fourth band at 8.0$\\mu$m allows us to observe the rest frame K-band where the bulk of the the stellar population of a galaxy radiates and expand the SED of the galaxies from rest-UV to NIR. Thus, with IRAC data in hand we have a powerful tool to probe in a more secure way the stellar content of the LBGs. Recent results of adding IRAC photometry to stellar mass estimates have been presented in e.g., Barmby et al. 2004  for z$\\sim$3, Shapley et al. 2005 for BX/BM objects and Reddy et al. 2006  for z$\\sim$2--3. Also, IRAC colours combined with ground based data, can enlighten the diversity of the population revealing the existence of different sub-classes of LBGs, and provide an insight to their physical properties such as the energy source that powers the LBGs.  This paper is organised as follows: Section 2 reviews the data of this study and presents the source extraction and the photometric analysis, as well as, the differential number counts for each field. Section 3 focus on the mid-infrared identifications of LBGs, while in Section 4 the rest frame near-infrared photometric properties of the LBGs are discussed. Finally, Section 5 summarises the main results of this study. All magnitudes appearing in this study are in AB magnitude system. ",
        "conclusions": "Through the photometric analysis of 6 fields  covered by all four IRAC bands on board Spitzer our conclusions are as follows:  \\begin{enumerate} \\item The excellent agreement of the number counts between all the fields in the bright end and between observations of equal exposure time in the faint end, shows that our photometric technique and source extraction has treated all fields in the same manner.  \\item Out of $\\sim700$ LBGs that were covered by our data, 443, 448, 137 and 152 LBGs were identified at 3.6$\\mu$m, 4.5$\\mu$m, 5.8$\\mu$m, 8.0$\\mu$m IRAC bands respectively, creating the largest existing rest-near-infrared sample of high-redshift galaxies.  \\item The SED of the LBGs were expanded to NIR and show that the near-infrared colours of the population spans over 6 magnitudes. The addition of IRAC bands reveals for the first time that LBGs display a variety of colours and their rest-near-infrared properties are rather inhomogeneous, ranging from : \\begin{itemize} \\item Those that are bright in IRAC bands and exhibit $R-[3.6]>1.5$ colours with steeply rising SEDs towards longer wavelengths  to \\item Those that are faint or not detected at all in IRAC bands with $R-[3.6]<1.5$ colour whose SEDs are rather flat from the far-UV to the NIR with marginal IRAC detections. \\end{itemize}  \\item Out of the whole sample, $\\sim20$$\\%$ of the LBGs are detected at 8.0$\\mu$m. We refer to them as the 8.0$\\mu$m sample of LBGs (It is equivelent to a rest frame K--selected sample). Those LBGs tend to have redder R$-$[3.6] colours when compared to the rest population with median values of 1.81 ($\\pm0.16$) and 1.09 ($\\pm0.26$) respectively.  \\item The infrared colours of LBGs are consistent with those of z$\\sim3$ galaxies and indicating that their SEDs are can be fitted with various stellar synthesis population models. The mid-infrared properties of the LBG (i.e., masses, dust, age, link to other z$\\sim3$ galaxy populations) will be presented in detail in a forthcoming paper as well as the full photometric catalogues.  \\item Based on on results for a few z$\\sim$3 optically identified AGN, IRAC 8$\\mu$m band can be used as a diagnostic tool to separate luminous, high z, AGN dominated objects from star--forming galaxies with AGNs being brighter in [8.0] band when compared to the LBG population.  \\end{enumerate} This work is based on observations made with the Spitzer Space Telescope, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with NASA. Support for this work was provided by NASA through an award issued by JPL/Caltech."
    },
    "0801/0801.0674_arXiv.txt": {
        "abstract": "{ The spectrum of IRAS 16594$-$4656 shows shock-excited H$_2$ emission and collisionally excited emission lines such as [O\\,{\\sc i}], [C\\,{\\sc i}], and [Fe\\,{\\sc ii}]. } {The goal is to determine the location of the H$_2$ and [Fe\\,{\\sc ii}] shock emission, to determine the shock velocities, and to constrain the physical properties in the shock.  } {High resolution spectra of the H$_2$~1-0~S(1), H$_2$~2-1~S(1), [Fe\\,{\\sc ii}], and Pa$\\beta$ emission lines were obtained with the near infrared spectrograph Phoenix on Gemini South. } {The position-velocity diagrams of H$_2$ 1-0 S(1), H$_2$ 2-1 S(1), and [Fe\\,{\\sc ii}] are presented.  The H$_2$ and [Fe\\,{\\sc ii}] emission is spatially extended. The collisionally excited [O\\,{\\sc i}] and [C\\,{\\sc i}] optical emission lines have a similar double-peaked profile compared to the extracted H$_2$ profile and appear to be produced in the same shock.  They all indicate an expansion velocity of $\\sim$8~km~s$^{-1}$ and the presence of a neutral, very high-density region with $n_{\\rm e}$ about $3\\,\\times\\,10^6$ to $5\\,\\times\\,10^7$\\,cm$^{-3}$. However, the [Fe\\,{\\sc ii}] emission is single-peaked. It has a Gaussian FWHM of 30\\,km\\,s$^{-1}$ and a total width of 62\\,km\\,s$^{-1}$ at 1\\% of the peak. The Pa$\\beta$ profile is even wider with a Gaussian FWHM of 48 km\\,s$^{-1}$ and a total width of 75\\,km\\,s$^{-1}$ at 1\\% of the peak.  } {The H$_2$ emission is excited in a slow 5 to 20 km\\,s$^{-1}$ shock into dense material at the edge of the lobes, caused by the interaction of the AGB ejecta and the post-AGB wind. The 3D representation of the H$_2$ data shows a hollow structure with less H$_2$ emission in the equatorial region. The [Fe\\,{\\sc ii}] emission is not present in the lobes, but originates close to the central star in fast shocks in the post-AGB wind or in a disk.  The Pa$\\beta$ emission also appears to originate close to the star.  } ",
        "introduction": "\\label{introduction}  Post-Asymptotic Giant Branch (AGB) stars represent an important transition phase in the evolution of low and intermediate-mass stars, between the AGB and the planetary nebula (PN) phases. During this period, the detached circumstellar envelope of gas and dust is expanding away from the star.  Meanwhile the star itself is increasing in temperature at nearly constant luminosity.  This phase lasts a few hundred to a few thousand years depending upon the star's core mass.  When the temperature is high enough and the star photoionizes the nebula, it has entered the PN phase (e.g. Kwok \\cite{Kwok93}). In spite of extensive study, the evolution from the AGB toward the PN stage is still poorly understood.  The drastic changes observed in circumstellar structure and kinematics are particularly puzzling. During the late AGB or early post-AGB evolutionary stages, the geometry of the circumstellar material changes from more or less spherically symmetric to axially symmetric, with the result that most PNe exhibit axisymmetric structures, ranging from elliptical to bipolar (e.g., Balick \\& Frank \\cite{Balick02}).  Bipolar PNe frequently possess molecular envelopes that are readily detectable in the near-infrared ro-vibrational lines of H$_2$ (Kastner et al.\\ \\cite{Kastner96}; Kelly \\& Hrivnak \\cite{Kelly05}). The available data suggest that the onset of near-infrared H$_2$ emission in PNe can be traced back to the proto-planetary nebula (PPN) phase but not back to the AGB phase of evolution (Weintraub et al.\\ \\cite{Weintraub98}; Davis et al.\\ \\cite{Davis05}). These observations suggest that further studies of H$_2$ emission from PPNe may offer insight into the transition from AGB star to PN and from spherical to axisymmetric mass loss.  The study of transition objects showing H$_2$ emission at an early stage is crucial for understanding the hydrodynamic processes shaping the nebulae. The H$_2$ lines can reveal details about the physical conditions in the shocks associated with these hydrodynamic processes and thereby help constrain models of the interaction of the central star with the AGB remnant.  IRAS 16594$-$4656 is classified as a post-AGB star for several reasons.  In the IRAS color-color diagram it has the colors of a PN (Van de Steene \\& Pottasch \\cite{VdSteene93}), but it has not been detected in the radio continuum at 3 or 6 cm. It has a large infrared excess due to dust with a color temperature of 173~K. It displays a double-peaked spectral energy distribution, with the peak in the mid-infrared much brighter than the peak in the near-infrared (Van de Steene et al.\\ \\cite{VdSteene00a}).  It possesses a CO envelope with an expansion velocity of 14 km\\,s$^{-1}$ (Woods et al.\\ \\cite{Woods05}).  The chemistry of IRAS 16594$-$4656 appears to be carbon-rich. This is based on the detection of unidentified IR emission features at 3.3, 6.2, 7.7, 8.6, 11.3, 12.6, and 13.4~$\\mu$m, as well as the 21~$\\mu$m feature (Garc\\'{\\i}a-Lario et al. \\cite{Garcia-Lario99}), all commonly associated with a carbon-rich chemistry. The optical spectrum of IRAS 16594$-$4656 shows a spectral type B7 with significant reddening. The H$\\alpha$ emission has a P-Cygni type profile indicative of a stellar wind (Van de Steene et al.\\ \\cite{VdSteene00b}). In Van de Steene \\& van Hoof (\\cite{VdSteene03}, hereafter Paper\\,I) we derived the total extinction value (i.e., interstellar and circumstellar extinction combined) using the extinction law of Cardelli et al. (\\cite{Cardelli89}). We found A$_{V}=7.5\\,\\pm\\,0.4$\\,mag with R$_V\\,=\\,4$. With this extinction value and the flux calibration from the Kurucz model, we determined a distance of ($2.2\\,\\pm\\,0.4$)\\,L$_4$$^{1/2}$\\, kpc, which was in good agreement with the distance of 2.5\\,L$_4$$^{1/2}$\\,kpc derived by Su et al. (\\cite{Su01}) (with L$_4$ the luminosity in units of $10^4$\\,L$_{\\sun}$).  In the optical the nebulosity is dominated by scattered light.  HST WFPC2 images showed that the object is a multipolar reflection nebula with 3 extensions on each side (e.g., Hrivnak et al.\\ \\cite{Hrivnak99}). It has been interpreted as a bipolar outflow episodically channeled from a rotating/precessing torus. The outer nebula also shows concentric arcs (Hrivnak et al.\\ \\cite{Hrivnak01}).  High resolution mid-infrared images in the N and Q-band show a bright equatorial torus viewed almost edge-on and a pair of bipolar lobes, which show a close correspondence with the H$_2$ HST map. The shape of the bipolar lobes shows that the fast outflow is still confined by the remnant AGB shell (Volk et al.\\ \\cite{Volk06}). The elongation of the inner nebula is along the symmetry axis of this torus at $\\sim 80^{\\circ}$~PA, and it seems likely that the material ejection is channeled into the bipolar lobes by the torus and/or some other collimation mechanism (Garc\\'{\\i}a-Hern\\'andez et al.\\ \\cite{Hernandez04}).  Ueta et al. (\\cite{Ueta05}, \\cite{Ueta07}) found, based on polarized-flux images, that the inner structure of the nebula is clearly elongated in the east- west direction (5\\farcs0 $\\times$ 2\\farcs0 at PA~81\\degr). Their polarized flux maps uncover the bipolar cusp structure of the shell, corresponding to the main elongation in intensity, and a hollow shell structure which delineates the wall of the elongated bipolar cavities.  In Paper\\,I we examined the near-infrared spectrum of IRAS 16594$-$4656.  It shows strong H$_2$ emission lines and some typical metastable shock-excited lines such as [Fe\\,{\\sc ii}] 1.257~\\& 1.644~$\\mu$m. We argued that the molecular hydrogen emission is mainly collisionally excited in a C-type shock. However, H$_2$ and [Fe\\,{\\sc ii}] emission don't usually coexist.  The goal of this paper is to clarify the location of the H$_2$ and [Fe\\,{\\sc ii}] shock emission in this object and to determine the shock velocities.  High resolution near-infrared spectra of IRAS 16594$-$4656 have been obtained with the near infrared high resolution spectrograph Phoenix on Gemini South.  The observations and data reduction are discussed in Section \\ref{observations}. The results of the analysis are presented in Section \\ref{results}, and the morphology, kinematics, and shock properties are discussed in Section \\ref{discussion}, and finally we present some conclusions in Section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  The H$_2$ emission is excited in slow 5 to 20 km\\,s$^{-1}$ shocks into dense material at the edge of the lobes, caused by the interaction of the AGB ejecta and the post-AGB wind.  The 3D representation of the H$_2$ emission shows a hollow structure.  There is less H$_2$ emission in the equatorial region.  The collisionally excited [O\\,{\\sc i}] and [C\\,{\\sc i}] optical emission lines have a similar profile compared to the extracted H$_2$ profile and appear to be produced in the same shock.  They all indicate an expansion velocity of $\\sim$~8~km~s$^{-1}$ and the presence of a neutral, very high density region of about $3\\,\\times\\,10^6$ to $5\\,\\times\\,10^7$~ cm$^{-3}$.  The [Fe\\,{\\sc ii}] emission is not present in the lobes, but originates close to the central star. It originates in fast shocks in the post-AGB wind, or in a circumstellar or circumbinary disk. The Pa$\\beta$ emission also appears to originate close to the star. The total width of the [Fe\\,{\\sc ii}] and Pa$\\beta$ emission line profiles is 62 and 75\\,km~s$^{-1}$ respectively, at 1\\% of the peak."
    },
    "0801/0801.0990_arXiv.txt": {
        "abstract": "The hierarchical growth of mass in the Universe is a pillar of all cold dark matter (CDM) models.   In this paper we demonstrate that this principle leads to a robust, falsifiable prediction of the stellar content of groups and clusters, that is testable with current observations and is relatively insensitive to the details of baryonic physics or cosmological parameters.  Since it is difficult to preferentially remove stars from dark-matter dominated systems, when these systems merge the fraction of total mass in stars can only increase (via star formation) or remain constant, relative to the fraction in the combined systems prior to the merger.  Therefore, hierarchical models can put strong constraints on the observed correlation between stellar fraction, \\fstar, and total system mass, $M_{500}$.  In particular, if this relation is fixed and does not evolve with redshift, CDM models predict $b=\\b\\gtrsim-0.3$.  This constraint can be weakened if the \\fstar--\\m500\\ relation evolves strongly, but this implies more stars must be formed in situ in groups at low redshift. Conservatively requiring that at least half the stars in groups were formed by $z=1$, the constraint from evolution models is $b\\gtrsim-0.35$.  Since the most massive clusters ($M\\sim10^{15}M_\\odot$) are observed to have \\fstar$\\sim 0.01$, this means that groups with $M=5\\times10^{13}M_\\odot$ must have \\fstar$\\leq0.03$.  Recent observations by \\citet{GZZ} indicate a much steeper relation, with \\fstar$>0.04$ in groups leading to $b\\approx -0.64$.  If confirmed, this would rule out hierarchical structure formation models: today's clusters could not have been built from today's groups, or even from the higher-redshift progenitors of those groups.  We perform a careful analysis of these and other data to identify the most important systematic uncertainties in their measurements. Although correlated uncertainties on stellar and total masses might explain the steep observed relation, the data are only consistent with theory if the observed group masses are systematically underestimated. ",
        "introduction": "An inescapable prediction of all cold dark matter (CDM) models is that mass in the Universe, in the form of CDM dominated ``haloes'', builds up {\\it hierarchically}, with low-mass systems merging to form progressively more massive galaxies and clusters of galaxies \\citep[e.g.][]{BBKS,DEFW}.  Moreover, the rate of this mass growth is precisely determined for a given set of cosmological parameters \\citep[e.g.][]{GW07}.  In practice, the complex and non-linear nature of baryonic physics \\citep[particularly cooling and heating processes, ][and many others]{WR78,WF91,Cole2000,bower06,Croton05} means that this is a difficult prediction to test through observations of galaxies alone. In fact, observations show that galaxies form in the opposite way, with the most massive systems having their stars in place first \\citep[e.g.][]{Cowie+96,Juneau+04,Pozzetti+07}.  This is not considered a falsification of the cold dark matter model, because the effect can be qualitatively explained by improving the physical description of baryonic processes in the models \\citep[e.g.][]{Croton05,bower06}, in a way that leaves untouched the prediction of hierarchical growth in the dark matter component.  However, an interesting and robust test of the theory can be obtained from observations of the stellar fraction of dark-matter dominated structures.  Unlike gas, it is very difficult to separate stars from dark matter, since they are both collisionless forms of matter that interact only via gravity.  And while new stars can be formed from gas (a process which is very poorly understood) they can only be destroyed through normal stellar evolution processes (which are quite well understood).  The latter effect only removes 10-30 per cent of the total stellar mass, with this range reflecting a weak dependence on star formation history and initial mass function \\citep[e.g.][]{JCP,BC03}.  Therefore,  when two similar systems merge, the mass fraction in visible stars must be at least as large as the fraction in the combined system prior to the merger, and the simplest expectation is that \\fstar\\ will either be constant or increase with total system mass.  Stellar fractions are most reliably measured for galaxy clusters, where the total mass in dark matter can be determined in various, independent ways (e.g. gravitational lensing, X-ray gas, or galaxy dynamics).  Interestingly, numerous studies have consistently shown that \\fstar\\ of clusters and groups {\\it decreases} with increasing mass \\citep[e.g.][]{H+00,MH02,Eke-groups3,G+02,RBGMR,LMS,cnoc2_ir}.  The usual explanation is that clusters are built not only from groups, but also from the accretion of low mass galaxies, where \\fstar\\ must be very low to explain the shallow faint-end slope of the luminosity function \\citep[e.g.][]{WF91,MH02}.  Another possibility is that low-mass groups form a significant number of stars, but only {\\it after} most clusters have been assembled. Either possibility allows theory to accommodate a mildly decreasing \\fstar\\ on cluster scales; it is our goal in this paper to use conservative constraints on these effects to put a robust limit on just how steep this decrease can be.  An important omission in many of the observational studies above, however,  has been the contribution from intracluster light (ICL), a low-surface brightness distribution of stars in groups and clusters that is very difficult to measure. Most studies of rich clusters find that the ICL contribution is relatively small, contributing less than 30 per cent to the total stellar light \\citep[e.g.][]{Durrell+02,Feld+04,Covone+06,KB07}, with at most a weak dependence on system mass \\citep{Zibetti}. Recently, using deep $I-$band observations of 23 nearby systems, \\citet{GZZ1} have made careful measurements of the ICL component, and come to the surprising conclusion that both \\fstar\\ and the relative ICL contribution depends much more strongly on mass than has been found previously, with the ICL actually dominating the total stellar mass in groups \\citep[][hereafter GZZ]{GZZ}. This has motivated us to consider whether or not these observations are able to falsify the hierarchical structure growth model.  We will begin by reanalyzing the observational data of GZZ, and complementary data from \\citet{LM04} in \\S~\\ref{sec-data}. In \\S~\\ref{sec-models} we use theoretical predictions for the growth of dark matter structure to put robust, falsifiable limits on the mass-dependence of \\fstar.   This prediction is then directly confronted with the observational data in \\S~\\ref{sec-test}, where we also discuss the implications of our findings, and the effect of possible biases and uncertainties in the measurements. Throughout this paper we generally assume a cosmology with $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, and $H_\\circ=70$\\kms.  However we explicitly consider how our results depend on cosmological parameters, in \\S~\\ref{sec-cosmo}. ",
        "conclusions": "The stellar fraction \\fstar\\ as a function of cluster mass \\m500\\ is an important test of hierarchical structure formation models.  We find that such models can make a robust, falsifiable prediction that the power-law slope relating these two quantities is $\\b>-0.35$, for systems with $M_{500}>5\\times10^{13}M_\\odot$. A steeper slope can only be accommodated if the \\fstar--$M_{500}$ relation evolves strongly, such that galaxy groups formed most of their stars in situ since $z=1.0$, which is not supported by observations. Since the most massive clusters today (with $M_{500}\\sim10^{15}M_\\odot$) are robustly measured to have \\fstar$\\sim0.01$, hierarchical models therefore require that galaxy groups (with $M_{500}\\sim5\\times 10^{13}M_\\odot$) have \\fstar$<0.03$.  This constraint is a conservative limit; ab-initio models predict a much flatter relationship with $b>-0.1$ \\citep{bower06}.  Recent observations by \\citet{GZZ} appear to conflict with this prediction.  In particular, their data follow a power-law relation over almost two orders of magnitude in mass, with $b=-0.64$, and  the lowest-mass systems in their sample have stellar fractions of 30 per cent, exceeding even the limits on the baryon fraction from WMAP.  If confirmed, these observations definitively rule out hierarchical structure formation.  $K-$band observations from \\citet{LM04}, on the other hand, are just consistent with the model constraints, but may still be inconsistent with the ab-initio model of \\citet{bower06} if the ICL contribution is as dominant in groups as claimed by \\citet{GZZ}.  These data do not extend to such low mass systems, nor are their data deep enough to directly measure the important ICL contribution.  More observations are needed to resolve the discrepancy between these two studies, and to thereby test the viability of cold dark matter models.  Particularly valuable would be X-ray images (and temperature maps) of the \\citet{GZZ} sample, to measure the gas content and total mass, and deep $I$ or near-infrared images of the \\citet{LM04} systems to directly measure the ICL component.  We end by noting however that the four lowest-mass systems in the Gonzalez et al. sample have poorly calibrated masses.  If we accept that their masses are underestimated by a factor of ten, then the remainder of their data are consistent with a universal stellar fraction of a few percent, and the apparent correlation with system mass is due to the correlated error bars on dynamical and stellar masses. Furthermore, removing these four systems greatly reduces the significance of their claim that the intracluster light fraction is a strong function of mass and, in this case, the \\citet{LM04} data (which have much smaller error bars) are also consistent with a constant stellar fraction \\fstar$\\approx 0.01$. This conclusion challenges the claim by GZZ that the stellar component of galaxy groups is dominated by the BCG and ICL, and that the low gas fractions in these groups is attributable to an increased stellar fraction."
    },
    "0801/0801.0442_arXiv.txt": {
        "abstract": "Massive stars are very rare, but their extreme luminosities make them both the only type of young star we can observe in distant galaxies and the dominant energy sources in the universe today. They form rarely because efficient radiative cooling keeps most star-forming gas clouds close to isothermal as they collapse, and this favors fragmentation into stars $\\ltsim 1$ $\\msun$ (solar mass; ref. \\cite{larson05,jappsen05,bonnell06d}). Heating of a cloud by accreting low-mass stars within it can prevent fragmentation and allow formation of massive stars\\cite{krumholz06b, krumholz07a}, but what properties a cloud must have to form massive stars, and thus where massive stars form in a galaxy, has not yet been determined. Here we show that only clouds with column densities $\\gtsim 1$ g cm$^{-2}$ can avoid fragmentation and form massive stars. This threshold, and the environmental variation of the stellar initial mass function (IMF) that it implies, naturally explain the characteristic column densities of massive star clusters\\cite{plume97,mueller02,shirley03,mckee03} and the difference between the radial profiles of H$\\alpha$ and UV emission in galactic disks\\cite{martin01, boissier06}. The existence of a threshold also implies that there should be detectable variations in the IMF with environment within the Galaxy and in the characteristic column densities of massive star clusters between galaxies, and that star formation rates in some galactic environments may have been systematically underestimated. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.0433.txt": {
        "abstract": "High resolution spectra of 110 selected red giant stars in the globular cluster M15 (NGC~7078) were obtained with Hectochelle at the MMT telescope in 2005 May, 2006 May, and 2006 October. Echelle orders containing H$\\alpha$ and \\ion{Ca}{2}~H~$\\&$~K are used to identify emission and line asymmetries characterizing motions in the extended atmospheres. Emission in H$\\alpha$ is detected to a luminosity of $log (L/L_{\\odot})=2.36$, in this very metal deficient cluster, comparable to other studies, suggesting that appearance of emission wings is independent of stellar metallicity. The faintest stars showing H$\\alpha$ emission appear to lie on the asymptotic giant branch (AGB) in M15. A line-bisector technique for H$\\alpha$ reveals outflowing velocities in all stars brighter than $log (L/L_{\\odot})=2.5$, and this outflow velocity increases with stellar luminosity, indicating the mass outflow increases smoothly with luminosity. Many stars lying low on the AGB show exceptionally high outflow velocities (up to 10$-$15 km s$^{-1}$) and more velocity variability (up to 6$-$8~km~s$^{-1}$), than red giant branch (RGB) stars of similar apparent magnitude. High velocities in M15 may be related to the low cluster metallicity. Dusty stars identified from \\textit{Spitzer} Space Telescope infrared photometry as AGB stars are confirmed as cluster members by radial velocity measurements, yet their H$\\alpha$ profiles are similar to those of RGB stars without dust. If substantial mass loss creates the circumstellar shell responsible for infrared emission, such mass loss must be episodic. ",
        "introduction": "The well-known second parameter problem in globular clusters \\citep{sandage02}, in which a parameter other than metallicity, affects the morphology of the horizontal branch, remains unresolved. Metallicity, as first noted by \\citet{sandage01}, remains the principal parameter, but pairs of clusters, with the same metallicity, display quite different horizontal branch morphologies thus challenging the canonical models of stellar evolution and leading to the need for a `Second Parameter'. Cluster ages have been  examined in many studies \\citep{searle01, lee02, stetson01, lee01, sarajedini01, sarajedini02} and in addition, many other suggestions for the `second parameter(s)' have been proposed, including: total cluster mass; stellar environment (and possibly free-floating planets); primordial He abundance; post-mixing surface helium abundance; CNO abundance; stellar rotation; and mass loss \\citep{catelan01, catelan02, sills01, soker01, sweigart01, buonanno01, peterson03, buonanno02, recio01}. Many authors \\citep{vandenberg01, lee02, catelan01} have proposed that more than one second parameter may exist in addition to age. One parameter may be mass loss which, as \\citet{catelan01} notes, remains an `untested second-parameter candidate'.  This spectroscopic study addresses the presence of mass loss from luminous stars in the globular cluster M15. Subsequent papers will include other clusters. Although stellar evolution theory predicts that low-mass Population II stars ascending the red giant branch (RGB) for the first time must lose mass \\citep{renzini01, sweigart02}, few observations have identified the ongoing mass loss process. Evidence from the period-luminosity relation for RR Lyrae stars suggests that the luminosity variations can be accommodated theoretically if mass loss $\\sim 0.2-0.4 \\ M_{\\odot}$ has occurred \\citep{fusi01, christy01}. Circumstellar CO emission in M-type irregular and semiregular asymptotic giant branch (AGB)-variables implies mass loss rates on the AGB $\\sim 10^{-7} - 10^{-8} \\ M_{\\odot} \\ yr^{-1}$ \\citep{olofsson01}. Indirect evidence of mass loss processes would be detection of an intracluster medium. These efforts have been marginally successful. Diffuse gas ($< \\ 1 \\ M_{\\odot}$) was suggested in NGC~2808 through the detection of 21$-$cm H line emission \\citep{faulkner01}, but has remained unconfirmed. Ionized intracluster gas was found in the globular cluster 47 Tucanae by measuring the radio dispersion of millisecond pulsars in the cluster \\citep{freire01}. The central electron density was derived ($n_e=0.067 \\ \\pm 0.0015$~cm$^{-3}$) and found to be two orders of magnitude higher than the ISM in the vicinity of 47 Tuc \\citep{taylor01}. \\citet{freire01} determined the electron density in M15 using four millisecond pulsars to be higher ($n_e\\sim0.2$~cm$^{-3}$) than in 47 Tuc.  Indirect evidence of mass loss processes comes also from infrared observations. \\citet{origlia01} using ISOCAM images found a mid-IR excess associated with giants in several globular clusters and attributed to dusty circumstellar envelopes. The first detection of intracluster dust in M15 was made by \\citet{evans01} from the analysis of far infrared imaging data obtained with the ISO instrument ISOPHOT. \\citet{loon01} also presented a tentative detection of 0.3 $M_{\\odot}$ of neutral hydrogen in M15. \\citet{smith02} placed an upper limit of 0.4 $M_{\\odot}$ for the molecular gas in M15 from CO observations with the 15-m James Clerk Maxwell Telescope on Mauna Kea. Using the \\textit{Spitzer} Space Telescope, \\citet{boyer01} detected a population of dusty red giants near the center of M15. Observations from \\textit{Spitzer} with the Multiband Imaging Photometer for \\textit{Spitzer} (MIPS) also revealed the intracluster medium (ICM) discovered by \\citet{evans01} near the core of the globular cluster. As \\citet{origlia01} noted, the infrared detections may only be tracing the outflowing gas and may not be related to the driving mechanisms for the wind. More recently \\citet{origlia02} identified dusty RGB stars in 47 Tuc and derived an empirical mass loss law for Population II stars. Mass loss rates derived from these observations showed that the mass loss increases with luminosity and possibly episodic.  High resolution stellar spectroscopy allows the direct detection of mass outflow from the red giants themselves. Emission in the wings of H$\\alpha$ lines in the spectra of globular cluster red giants was first described in detail by \\citet{cohen01}. Later observations revealed that emission in H$\\alpha$ is common in globular clusters and night-to-night variations can occur \\citep{mallia01, peterson01, peterson02, cacciari02, gratton01}. These studies have shown that most of the stars brighter than $log (L/L_{\\odot}) \\approx 2.7$ exhibit H$\\alpha$ emission wings. The emission itself is likely not a direct indicator of mass loss, because emission can arise from an optically thick stellar chromosphere surrounding the star \\citep{dupree01}. Variation of the strength of emission can also be affected by stellar pulsation \\citep{smith01}. Better mass flow indicators in the optical are the line coreshifts or asymmetries of the H$\\alpha$ or \\ion{Ca}{2}~H$\\&$K profiles and emission features. Red giants in globular clusters (M22 and Omega Centauri) were found to have velocity shifts less than 14~km~s$^{-1}$ in the cores of H$\\alpha$ relative to the photospheric lines \\citep{bates01, bates02}. These results were similar to metal-poor field giants, where only giants brighter than $M_V=-1.7$ have emission wings and the line shifts were $< \\ 9$~km~s$^{-1}$ \\citep{smith01} indicating very slow outflows and inflows in the chromosphere. For globular clusters, \\citet{lyons01} discussed the H$\\alpha$ and Na~I~D line profiles for a sample of 63 RGB stars in M4, M13, M22, M55, and $\\omega$~Cen. The coreshifts were less than 10~km~s$^{-1}$, much smaller than the escape velocity from the stellar atmosphere at $2 \\ R_{*}$ ($\\approx 50-70$~km~s$^{-1}$). \\citet{dupree02} studied 2 RGB stars in NGC~6752 and found that the \\ion{Ca}{2}~K and H$\\alpha$ coreshifts were also low (less than 10~km~s$^{-1}$). However, asymmetries in the Mg~II lines showed a stellar wind with a velocity of $\\approx 150$~km~s$^{-1}$, indicative of a strong outflow. Mg~II lines are formed higher in the atmosphere than H$\\alpha$ and \\ion{Ca}{2}~K, which suggests that the stellar wind becomes detectable near the top of the chromosphere. A detailed study was carried out by \\citet{cacciari01}, who observed 137 red giant stars in NGC~2808. Most of the stars brighter than $log (L/L_{\\odot})=2.5$ clearly showed emission wings in H$\\alpha$. The velocity shift of the H$\\alpha$ line core compared to the photosphere is less than $\\approx 9$~km~s$^{-1}$. Outward motions were also found in both Na~I~D and \\ion{Ca}{2}~K profiles.  This paper discusses high-resolution spectroscopy of the H$\\alpha$ and \\ion{Ca}{2}~H$\\&$K lines of red giant stars in M15. Our deep sample of M15 giants in this metal poor cluster ([Fe/H]=$-$2.26) offers a good comparison to other studies of more metal rich clusters. Also, the high radial velocity of M15 ($-$105~km~s$^{-1}$) minimizes contamination by interstellar absorption in the profiles of resonance lines. ",
        "conclusions": "Differences found in the profiles of the H$\\alpha$ line give insight into the atmospheric structure and dynamics. Stars which are physically larger show more emission. At a given luminosity, the M15 stars exhibit a consistently lower {\\it fraction} of stars with emission, than found in NGC~2808, a cluster which is more metal rich. This suggests the M15 chromospheres may have less material than stars at the same luminosity in NGC~2808 as demonstrated by detailed modeling of emission wings and consistent with the decrease in radius with metallicity.  Stars in M15 have H$\\alpha$ emission wings that vary in time so that the magnitude of the faintest giant showing emission changes among the different dates of observation. The lower limit to the presence of H$\\alpha$ emission in M15 [$log (L/L_{\\odot})=2.36$] is comparable to that found in NGC~2808 \\citep{cacciari01}.  M15 exhibits continuous outflows at lower luminosities and with higher velocities than the more metal rich clusters, NGC~2808 and M4, hinting at a metallicity dependence. These outflows may decrease the chromospheric material and account for the lower fraction of stars with emission wings in M15. Detailed modeling is necessary to evaluate the mass loss rate from the line profiles. To identify mass loss as the solution to the second parameter problem will also require more complete measurements of clusters with varying metallicities.  The bisector velocity of the H$\\alpha$ core along the RGB indicates outflow (negative velocities) or no motion for stars brighter than $log (L/L_{\\odot})=2.5$; the outflow velocity increases with increasing stellar luminosity. However, AGB stars near $log (L/L_{\\odot})\\sim2.4$ have bisector velocities comparable in value to those at the tip of the RGB and also exhibit larger changes in velocity between observation than the RGB stars. We take this as evidence of more substantial and episodic mass outflow on the AGB.  \\ion{Ca}{2}~K emission is detected at least to the limits of H$\\alpha$ emission; however, the asymmetry in the \\ion{Ca}{2}~K core, where measurable, may differ from the asymmetry measured in the H-alpha wings perhaps due to time variability or different line-forming regions.  Twelve stars identified in Spitzer observations as dusty IR sources and AGB stars \\citep{boyer01} have radial velocities consistent with cluster membership. The similarities in H$\\alpha$ line profile characteristics between the \\textit{Spitzer} sources and other red giants in M15 suggests the IR emission attributed to circumstellar dust must be produced by an episodic process."
    },
    "0801/0801.1996_arXiv.txt": {
        "abstract": "We investigate the scattering and absorption of light by random ballistic aggregates of spherical monomers. We present general measures for the size, shape, and porosity of an irregular particle. Three different classes of ballistic aggregates are considered, with different degrees of porosity. Scattering and absorption cross sections are calculated, using the discrete dipole approximation (DDA), for grains of three compositions ($50\\%$ silicate and $50\\%$ graphite; $50\\%$ silicate and $50\\%$ amorphous carbon; and $100\\%$ silicate, where percentages are by volume), for wavelengths from $0.1\\ \\mu$m to $4\\,\\mu$m. For fixed amount of solid material, increased porosity increases the extinction at short wavelengths, but decreases the extinction at wavelengths long compared to the overall aggregate size. Scattering and absorption cross sections are insensitive to monomer size as long as the constituent monomers are small compared with the incident wavelength. We compare our accurate DDA results with two other approximations: the analytical multi-layer sphere (MLS) model and effective medium theory (EMT). For high porosity and/or absorptive materials, the MLS model does not provide a good approximation for scattering and absorption by ballistic aggregates. The EMT method provides a much better approximation than the MLS model for these aggregates, with a typical difference $\\lesssim 20\\%$ in extinction and scattering cross sections compared with DDA results, for all types, compositions and wavelengths probed in this study. ",
        "introduction": "The appearance of star-forming galaxies is strongly affected by interstellar dust.  Starlight is absorbed and scattered by dust particles, and the absorbed energy is predominantly reradiated at infrared wavelengths. The optical properties of the dust must be characterized in order to determine the spectra and spatial distribution of the stars (i.e., to ``correct'' for reddening and extinction), and to interpret the observed infrared emission. Characterizing the dust is also important for more fundamental reasons: dust grains play an important role in the thermodynamics and chemistry of the interstellar medium (e.g., heating from photoelectric emission, and grain catalysis of H$_2$), in interstellar gas dynamics (e.g., radiation pressure on dust grains), and in the formation of stars, planets, and planetesimals.  Many interplanetary dust particles (IDPs) are irregular, fluffy aggregates, sometimes described as having a ``fractal'' appearance. Direct evidence comes from such IDPs collected by high-flying aircraft \\citep{Brownlee_1985,Warren+Barrett+Dodson+etal_1994}. Laboratory and microgravity experiments of dust particle interactions, which may mimic the conditions in the early solar system, also suggest that particles form fractal assemblies through ballistic aggregation \\citep{Wurm+Blum_1998, Blum+Wurm_2000, Krause+Blum_2004}. Similar aggregation processes may well occur in interstellar environments, where small dust grains coagulate in dense molecular clouds \\citep{Dorschner+Henning_1995}. A number of authors have proposed that interstellar grains consist primarily of such aggregate structures, with a random, porous geometry \\citep[e.g.,][]{Mathis+Whiffen_1989}.  Aggregates are irregular, porous, particle assemblies, possibly incorporating multiple materials. Early attempts to estimate the optical properties of such structures used a combination of Mie theory and effective medium theory (EMT), where an aggregate is approximated by a homogeneous sphere with an ``effective'' refractive index intended to allow for the effects of porosity \\citep[e.g.,][]{Mathis_1996, Li+Greenberg_1998}. Alternatively, \\citet{Voshchinnikov+Mathis_1999} proposed a multi-layer sphere (MLS) model to account for porous, composite dust grains. The advantage of this MLS model is that, just as for homogeneous spheres, the solution is exact and can be obtained via a fast algorithm \\citep{Wu+Guo+Ren+etal_1997}. The limitations are: the grains are assumed to be spherical, and all compositions (including vacuum) are assumed to be well mixed inside the sphere. Comparisons of the MLS with EMT-Mie theory are done by \\citet{Voshchinnikov+Ilin+Henning_2005, Voshchinnikov+Ilin+Henning+Dubkova_2006}, suggesting that the MLS model is accurate for spherical, well-mixed dust grains, for certain compositions and porosities. So far there has been no study of whether or not the MLS prescription is also suitable for irregular aggregates.  Solving Maxwell's equations becomes increasingly challenging when the geometry departs from spheres. Several methods have been developed to compute light scattering by non-spherical particles. The ``extended boundary condition method'' (EBCM) introduced by \\citet{Waterman_1971}, and developed by \\citet{Mishchenko+Travis_1994} and \\citet{Wielaard+Mishchenko+Macke+Carlson_1997} can be applied to solid targets with relatively smooth surfaces, and exact series expansions have been developed for spheroids \\citep{Asano+Yamamoto_1975,Asano+Sato_1980,Voshchinnikov+Farafonov_1993, Farafonov+Voshchinnikov+Somsikov_1996}, but these techniques cannot be applied to the complex geometries of interest here. The special case of clusters of spheres can also be treated by superposition of vector spherical harmonics \\citep{Mackowski_1991,Xu_1997}, sometimes referred to as the generalized multisphere Mie (GMM) solution. If the complete T-matrix can be found for the individual particles, the T-matrix for a cluster can, in principle, be found using the superposition T-matrix method \\citep[TMM, e.g.,][]{Mackowski+Mishchenko_1996}. However, these techniques, although formally exact, can prove very computationally-demanding when applied to spheres that are numerous and in close proximity (e.g., in contact), particularly when the refractive index differs appreciably from vacuum.  The discrete dipole approximation (DDA) method \\citep[e.g.,][]{Purcell+Pennypacker_1973, Draine+Flatau_1994} is the most flexible among all the methods, although substantial computational resources may be needed in order to achieve the desired accuracy. Using either GMM, TMM or DDA, the optical properties of different kinds of aggregates have been investigated extensively during the past decade \\citep[e.g.,][]{West_1991, Lumme+Rahola_1994, Petrova+Jockers+Kiselev_2000, Kimura+Kolokolova+Mann_2006, Bertini+Thomas+Barbieri_2007}; most of these investigations are applied to cometary dust.   There is increasing interest in the possibility that interstellar and circumstellar grains may be porous, random aggregates. To confront this hypothesis with observational data, it is necessary to calculate the optical properties of such aggregates. The calculations must be carried out for a wide range of wavelengths - from infrared to vacuum ultraviolet - and for aggregates with potentially realistic geometry and composition. Massive computations are required to attain these goals.   In this paper, we use the DDA to calculate scattering and absorption for a large sample of aggregates with different compositions, sizes, and porosities, over a wide range of wavelength. We explore the dependence of various optical properties on wavelength, grain size, aggregate geometry, as well as compositions, and we investigate the applicability of multi-layer sphere calculations and the EMT-Mie model to these random aggregates.  Characterization of the size, gross shape, and porosity of an irregular particle is discussed in \\S \\ref{sec:geom}. \\S \\ref{sec:model} describes procedures for generating random aggregates (``BAM1'' and ``BAM2'') that are less porous than the standard ballistic aggregates, and \\S \\ref{sec:DDA} discusses the applicability of the DDA method to these aggregates.  Extinction and absorption behaviors are presented in \\S \\ref{sec:extinction}, with comparison to the MLS model in \\S \\ref{subsec:mls}, and comparison to the EMT-Mie model in \\S \\ref{subsec:emt}. Our results are summarized and discussed in \\S\\ref{sec:discussion}. Applications to circumstellar debris disks and cometary dust will be presented in a companion paper \\citep[][hereafter Paper II]{Shen+Draine+Johnson_2007b}. ",
        "conclusions": "\\label{sec:discussion} The principal results of this study are the following: \\begin{enumerate}  \\item Two new algorithms for generating random aggregates are introduced: ballistic aggregation with one migration (BAM1), and ballistic aggregation with two migrations (BAM2).  BAM1 and BAM2 aggregates are less porous, and more mechanically robust, than conventional BA aggregates.  \\item A measure $\\poro$ of the porosity of a structure is proposed (eq. \\ref{eq:fill factor}), as well as a measure $\\Rabc$ (eq. \\ref{eq:Rabc}) for the ``characteristic size'' of the structure.  \\item Monte-Carlo simulations are used to determine the statistical properties of $\\poro$ and $\\Rabc$ for ballistic aggregates (see Tables 1-2 and Figs.\\ \\ref{fig:frchar}-\\ref{fig:ffill}).  \\item We confirm (see Figure \\ref{fig:converg}) that the accuracy of the DDA scales as the interdipole spacing $d$ as $d\\rightarrow0$, or, equivalently as $N_{\\rm dip}^{-1/3}$ as $N_{\\rm dip}\\rightarrow\\infty$, where $N_{\\rm dip}$ is the total number of dipoles.  \\item Scattering, absorption, and extinction cross sections are calculated for the $N=256$, $a_0=0.02\\ \\micron$ ($\\aeff=0.127\\ \\micron$) BA, BAM1, and BAM2 clusters, for three different compositions: 100\\% silicate, 50\\% silicate/50\\% amorphous carbon, and 50\\% silicate/50\\% graphite, for wavelengths $0.1\\le\\lambda\\le3.981\\,\\micron$. The BA clusters (with the highest porosity $\\poro$) have the largest extinction cross sections at short wavelengths, but at optical and near-IR wavelengths, the BAM2 clusters (with the lowest $\\poro$) provide more extinction per unit solid material. At constant porosity and same amount of solid material, the monomer size has no significant effect provided the monomers are small compared to the incident wavelength.  \\item We compared the DDA results with the analytical MLS model and EMT-Mie theory. We found the MLS model does not provide a good approximation for absorptive and/or very porous grains; the EMT-Mie model provides much better agreement with the DDA results. For computing total cross sections ($Q_{\\rm ext}$, $Q_{\\rm abs}$, $Q_{\\rm sca}$), the EMT-Mie method provides results accurate to $\\sim10\\%$ if the vacuum fraction $f_{\\rm vac}$ is taken to be $0.94\\poro$.  \\end{enumerate}  The effects of porosity on extinction cross sections have important implications for the abundance budget problem in interstellar dust models. The recent observed decrease of solar abundances \\citep[e.g.,][]{Asplund+Grevesse+Sauval_2005b} and the claim that interstellar abundances might be better represented by those of B stars \\citep[e.g.,][]{Snow+Witt_1996} have imposed a challenge to dust extinction models. Porous dust grains have been thought to be a solution to this abundance budget problem \\citep[e.g.,][]{Mathis_1996}, because they were expected to result in greater extinction per unit solid material than compact grains. However, our results show that porosity actually {\\it decreases} the opacity at wavelengths long compared to the overall grain size. Hence caution must be paid when dealing with the abundance budget problem. Until detailed models have been constructed using random aggregates to reproduce the observed interstellar extinction (and polarization), we will not know if such grain models will alleviate the interstellar abundance problem.  Work on this problem is underway \\citep{Johnson+Draine_2008}.  Ballistic aggregates are promising candidates for interstellar and circumstellar dust grains. In the companion Paper II, we will discuss the scattering properties of ballistic aggregates and present examples that can reproduce the observations of light scattered by dust in debris disks and comets."
    },
    "0801/0801.1358_arXiv.txt": {
        "abstract": "{It has been suggested by Roukema and coworkers (hereafter R04) that the topology of the Universe as probed by the ``matched circles'' method using the first year release of the WMAP CMB data, might be that of the Poincar\\'e dodecahedral space (PDS) model. An excess in the correlation of the ``identified circles'' was reported by R04, for circles of angular radius of $\\sim 11^\\circ$ for a relative phase twist $-36^\\circ$, hinting that this \\langed{could be} due to a Clifford translation, if the hypothesized model \\langed{were} true. R04 did not however specify the statistical significance of the correlation signal.  We investigate the statistical significance of the signal using Monte Carlo CMB simulations in a simply connected Universe, and present an updated analysis using the three-year WMAP data. We find that our analyses of the first and three year WMAP data provide results that are consistent with the simply connected space at a confidence level as low as 68\\%. ",
        "introduction": "\\langed{If the topology of the Universe were multiply connected, as opposed to simply-connected, and if the comoving size of the fundamental domain (FD) were smaller than the comoving distance to the surface-of-last-scattering (SLS), then it should be possible to detect repeating patterns in the CMB fluctuations using full-sky data of sufficient signal-to-noise ratio.} These fluctuations would be those lying along pairs of circles defined by points of intersection between different copies of the SLS in the covering space \\citep{Corn98b}. These patterns although found in different directions of the sky, would constitute so-called ``matched circles'', as they would represent the same physical points, but observed from different directions due to topological lensing.\\\\ While this principle is true for any 3-manifold model of space, the number of pairs of ``matched circles'' or their sizes and relative spatial orientations, as well as their handedness, or phase shift, depend significantly on the assumed 3-manifold and its topological properties, thus providing a way to observationally distinguish between models.  \\par While the positive correlation signal from matched pairs is expected directly from the metric perturbations, via the Sachs-Wolfe Effect \\citep{SW67}, there are many other cosmological effects (e.g. the Doppler effect, the Integrated Sachs-Wolfe Effect (ISW)) \\citep{2008PhRvD..77b3525K}, astrophysical foregrounds \\citep{WMAPforegrounds} and instrumental effects that constitute noise, from the point-of-view of a matched circles search, and the magnitude of the effects depends on the angular scale. \\par Although the CMB data have been analyzed to detect topological lensing signals since the availability of the COBE data \\citep{Roukema00-3}, the release of the WMAP observations has provided full-sky data of unprecedented accuracy and resolution, \\langed{opening up more promise for direct tests of the topology of the Universe.  } Although the ``matched circles'' test is straightforward, it is limited due to noise and FD size constraints. Additional theoretical predictions can be used as independent tests that involve predictions of CMB temperature and polarization fluctuations \\langed{for the case that the Universe is multiply connected}, both in real and spherical harmonic spaces, or topological effects on the CMB power spectrum \\citep{2004CQGra..21.4901A,2003MPLA...18.2099W,2003PhLA..311..319G,2005MNRAS.358.1285D,2004PhRvD..69j3514R, LLU98,Inoue99, 2007PhRvL..99h1302N,Corn98a,deOliv95,2006PhRvD..73b3511K,2003Natur.425..593L, 2006AIPC..848..774N,2006ApJ...645..820P, 2007AA...476..691C}. Although a successful ``matched circles'' test would provide strong support for \\langed{the Universe being multiply connected}, no statistically-significant evidence has been found \\citep{2004PhRvL..92t1302C,2006astro.ph..4616S}.  \\par In \\cite{RLCMB04} we performed a ``matched circles'' search using the first year WMAP ILC map \\citep{WMAPforegrounds} and found an excess correlation, which one would expect under the PDS hypothesis for circles of angular radii $\\alpha\\sim 11^\\circ$ with centers towards \\thedodec and their opposites (Fig.~\\ref{fig:thedodec}).  \\par In this present work we have two key objectives. We revisit those results, verify the existence of the excess correlations and quantify their statistical significance. Secondly, we update the search with the WMAP three year data release, extend it to probe three different resolutions (smoothing lengths) and define the detection confidence thresholds. We also discuss the effects of underlying 2-point correlations, smoothing length and incomplete sky coverage on the value of correlation coefficient.   \\begin{figure}[!hbt] \\renewcommand{\\figurename}{Fig} \\includegraphics[width=0.5\\textwidth]{DODfig1.eps} \\caption{Visualization of the matched circles solution reported in \\protect\\cite{RLCMB04} and reproduced in this paper to constrain its statistical significance, over-plotted on the first year ILC map masked with Kp2 sky mask.} \\label{fig:thedodec} \\end{figure}    In section~\\ref{sec:data_and_sims} we introduce the datasets used in the analysis, provide details of their preprocessing, and describe simulations that we use to complete a statistical significance analysis. In section~\\ref{sec:statistics} we introduce the details of the statistics being performed and our confidence-level analysis. Results are presented in section~\\ref{sec:results}. We conclude in section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} \\par In \\cite{RLCMB04} it has been suggested that the shape of the space might be consistent with \\langed{the} Poincar\\'e dodecahedral space (PDS) model (Fig.~\\ref{fig:thedodec}). This suggestion was due to an \\langed{excess positive correlation} in the matched circles test \\citep{Corn98b} of \\langed{the} first year WMAP ILC map, however the statistical significance of this excess was not specified.  We have revisited those results and found consistent correlation excess corresponding to the same orientation of the fundamental dodecahedron using independent software.  We extended and updated the matched circles search with the WMAP three year ILC data and the three year foreground reduced, inverse noise co-added map, and tested these at three different smoothing scales $(1^\\circ,2^\\circ,4^\\circ)$.  We performed \\langed{an} analysis of the statistical significance of the reported excess, based on realistic and very conservative MC GRF CMB simulations of the datasets.  We find that \\langed{under} ``matched circles'' tests, both the first and three year WMAP data are consistent with the simply-connected topology hypothesis, for all smoothing scales, at a confidence level as low as 68\\%, apart from the first year ILC data smoothed to $4^\\circ$, which \\langed{are} consistent at 95\\% CL."
    },
    "0801/0801.1672_arXiv.txt": {
        "abstract": "{% In the solar convection zone, rotation couples with intensely turbulent convection to build global-scale flows of differential rotation and meridional circulation.  Our sun must have rotated more rapidly in its past, as is suggested by observations of many rapidly rotating young solar-type stars. Here we explore the effects of more rapid rotation on the patterns of convection in such stars and the global-scale flows which are self-consistently established.  The convection in these systems is richly time dependent and in our most rapidly rotating suns a striking pattern of spatially localized convection emerges.  Convection near the equator in these systems is dominated by one or two patches of locally enhanced convection, with nearly quiescent streaming flow in between at the highest rotation rates.  These active nests of convection maintain a strong differential rotation despite their small size.  The structure of differential rotation is similar in all of our more rapidly rotating suns, with fast equators and slower poles.  We find that the total shear in differential rotation, as measured by latitudinal angular velocity contrast, $\\Delta \\Omega$, increases with more rapid rotation while the relative shear, $\\Delta \\Omega/ \\Omega$, decreases.  In contrast, at more rapid rotation the meridional circulations decrease in both energy and peak velocities and break into multiple cells of circulation in both radius and latitude.} ",
        "introduction": " ",
        "conclusions": "The complex coupling between rotation and convection drives a strong differential rotation in these more rapidly rotating stars.  The meridional circulations in contrast are much weaker, decreasing in both magnitude and relative contribution with more rapid rotation.  As such, more rapidly rotating stars may have dynamos which differ substantially from the solar dynamo, where meridional circulations are thought to play an important role.  These mean circulations are achieved despite the presence of strong spatial modulation in the convection of the equatorial regions.  If these active nests of convection persist in stellar convection zones in the presence of magnetism, they may lead to features at the surface which propagate at rates distinct from that of the surface convection or the zonal flows of differential rotation.  Such strong and persistent modulation in convective structures may have important observational consequences."
    },
    "0801/0801.0596_arXiv.txt": {
        "abstract": "We consider the effect of mass segregation on the observable integrated properties of star clusters. The measurable properties depend on a combination of the dynamical age of the cluster and the physical age of the stars in the cluster. To investigate all possible combinations of these two quantities, we propose an analytical model for the mass function of segregated star clusters that agrees with the results of N-body simulations, in which any combination can be specified.  For a realistic degree of mass segregation and a fixed density profile, we find with increasing age an increase in the measured core radii and a central surface brightness which decreases in all filters more rapidly than what is expected from stellar evolution alone. Within a Gyr the measured core radius increases by a factor of two and the central surface density in all filters of a segregated cluster will be overestimated by a similar factor if mass segregation is not taken into account during the conversion from light to mass. We find that the $V-I$ colour of mass segregated clusters decreases with radius by about 0.1-0.2 mag, which could be observable. From recent observations of partially resolved extra-galactic clusters, a decreasing half-light radius with increasing wavelength was observed, which was attributed to mass segregation. These observations can not be reproduced by our models. In addition, we provide physical arguments based on the evolution of individual stars that one should not expect strong dependence of core radius as a function of the wavelength.  We find that the differences between measured radii in different filters are always smaller than 5\\%. ",
        "introduction": "\\label{sect:introduction} Observations of star clusters are often used as test-beds for theories of star formation, the stellar initial mass function (IMF) and dynamical evolution. An important ingredient in most of these theories is the differential distribution of stellar masses within a cluster, or {\\it mass segregation}.  In several resolved clusters, evidence of mass segregation has been claimed on the basis of observed variations of the stellar mass function (MF) with distance to the cluster centre (see for example \\citealt{1996ApJ...466..254B, 1998ApJ...492..540H,2002MNRAS.331..245D,2002A&A...394..459S} for R136 in 30 Doradus, the Orion Nebula Cluster, clusters in the LMC and the Arches cluster, respectively). These observations claim an overabundance of massive stars in the cluster centre. This stratification of stellar masses is expected from dynamical evolution, since dynamical friction slows down the most massive stars. As a result, these stars sink to the cluster centre on a time-scale that is inversely proportional to their mass. Because the aforementioned clusters are dynamically young, the observations advocate a primordial origin for this segregation of stellar masses.  From the ``competitive accretion'' star formation model (see for example \\citealt{1997MNRAS.285..201B}), it is expected that the most massive stars form in the highest density environments, corresponding to the inner parts of the clusters. The preferential formation of massive stars in the centre of the cluster is often used to explain the observations of dynamically young, but mass segregated clusters.  The determination of the stellar MF in different annuli around the cluster centre, which is the most common technique used to ``detect\" mass segregation in resolved clusters, is hampered by several observational difficulties. First, crowding and blending of stars in the core can mimic a shallower MF at that location \\citep{2008ApJ...677.1278M, ascenso}. Second, the determination of stellar masses from the observed luminosities depends on the adopted age, which is usually taken constant for all stars in the cluster. However, only a small spread in age is enough to cause misinterpretations (see \\citealt{1995ApJ...448..179H}, who show that the MF of the R136 cluster is consistent with Salpeter when this effect is taken into account).  Alternatively, integrated properties, such as the surface brightness profile in different filters can be employed to study radial variations of the stellar MF. In this way one does not have to rely on individual star counts, thus avoiding possible biases; moreover, this method can be used for clusters at larger distances. It was shown that metal-rich (red) star clusters appear to have smaller half-light radii than their metal-poor (blue) counterparts \\citep{2004ApJ...613L.117J}. In addition, one could expect that segregated star clusters appear larger in the ultra-violet (UV) than in the near infra-red (NIR) as most of the light at these red wavelengths comes from the massive stars, whereas the bluer wavelengths are dominated by intermediate mass stars. Tentative evidence for this is given by \\citet{2005ApJ...621..278M} who find a decreasing half-light radius, or effective radius (\\reff), with increasing wavelength for the young massive cluster M82-F. \\citet{2007arXiv0710.0547L} find for NGC~1569-B that the \\reff\\ measured in the $U$-band is around 30-50\\% larger than the \\reff\\ measured in the $I$-band. A smaller radius in the redder filters is qualitatively what one would expect when the massive (red) stars are more centrally concentrated, but it has thus-far not been quantified how large the expected difference is.  Although the integrated properties are free of the biases that are encountered in methods that rely on individual star counts, there are other problems such as variations of the PSF between the different filters \\citep{2007arXiv0710.0547L}, differential foreground extinction \\citep{2007MNRAS.379.1333B} and intra-cluster extinction \\citep{2002A&A...394..459S}, that make it a challenging task to accurately determine intrinsic differences between the cluster properties in different filters.  To quantify the expected variations in different filters, we present a method to rapidly simulate integrated luminosity profiles of mass segregated star cluster in different filters. We choose to do this analytically to avoid statistical fluctuations which one has to deal with when considering (realistic) $N$-body systems, and therefore we do not include the dynamical evolution. Instead, we apply our method to clusters with different concentrations, which may result from dynamical evolution.  In \\sect{\\ref{sect:model}} we introduce the model of the mass function of a segregated star cluster and in \\sect{\\ref{sect:observations}} we present simulated observational properties of such clusters. A discussion and our conclusions are presented in \\sect{\\ref{sect:conclusions}}.    \\begin{figure*} \\begin{center} \\includegraphics[scale=0.8]{plots/profile_w5.ps}\\\\ \\includegraphics[scale=0.8]{plots/profile_w7.ps}\\\\ \\includegraphics[scale=0.8]{plots/profile_w9.ps}\\\\ \\end{center} \\caption[]{% Surface brightness as a function of a distance to the cluster centre, which is given in units of core radius of non-segregated star cluster. The dotted blue line shows a surface brightness profile in $U$-band, dashed green line is for $V$-band and dash-dot red line is for $I$-band. The black line represents surface brightness profile of a non-segregated cluster in $V$-band. The mass density profile for all the clusters of different ages is kept the same. The density profiles are King models with $W_0=5$, $7$ and $9$ for top, middle and bottom panels, respectively.} \\label{fig:profiles} \\end{figure*} ",
        "conclusions": "\\label{sect:conclusions} We have simulated observational properties of mass segregated star clusters with the aim to quantify the imprint of mass segregation on integrated cluster properties. We choose to model the segregated mass functions analytically based on results from observations and $N$-body simulations.  In order to calculate upper limits of the imprint of mass segregation, we only considered clusters with the maximum possible degree of segregation which can be achieved in our model. While this may not necessary be reached through dynamical evolution, one may think of this setup as the result of a combination of primordial segregation combined with dynamical evolution.  In young ($\\sim10\\,$Myr) star clusters, we find only small differences ($\\lesssim5\\%$) between the core radius (\\rc) found in different filters, and the differences becomes smaller for older clusters. The explanation for these small differences is that the most massive stars dominate the light in all filters at all ages. It, therefore, appears that the comparison of the measure \\rc\\ in different filters is not a particularly suitable tool to look for mass segregation in star clusters and, therefore, we do not expect that one could study whether the young star cluster is in core-collapse phase by using the light alone \\citep{2006ApJ...653L.113K,2007MNRAS.tmpL..45P}.  The observed \\rc\\ is underestimated for young segregated star clusters. The difference with the real \\rc\\ decreases with age and is only 10\\% at 1 Gyr compared to nearly a factor of two at 10 Myr and this does not depend on the cluster concentration.  The same factor was found for the half-light radius in gaseous models of mass segregated clusters by \\cite{2005ApJ...620L..27B}.  The underestimation of \\rc\\ results in an overestimation of the observed central density. This is because the light-to-mass ratio in the centre of a segregated cluster can be an order of magnitude higher than that of a simple stellar population without segregation and, therefore, central 3D mass density might be overestimated by nearly an order of magnitude. However, due to the projection the measured central surface density is overestimated by a factor of three, and this result slightly depends on the choice of $W_0$. In particular, for clusters with high concentration the central density might be overestimated by nearly an order of magnitude. These results are consistent with the finding of \\cite{2005ApJ...620L..27B}.  The observed central surface and core radius approached its true value around 1 Gyr. We therefore conclude that light is not a good tracer of mass in young ($<100\\,$Myr) star clusters that are segregated and this effect should be taken into account when trends of \\rc\\ or central surface density with age are discussed \\citep{2007A&A...469..925S, 2007AJ....134..912N, bastian08}. For clusters with ages above 100\\,Myr these differences are smaller than 20\\%.  It is possible to observe mass segregation by looking at colour differences between the inner and outer parts of star clusters. We find that our simulated clusters have $V-I$ differences of roughly 0.1-0.2 magnitude between the centre and the outer part of the cluster. This effect has been observed by \\cite{2007arXiv0710.0547L} in the $\\sim20\\,$Myr old massive cluster NGC~1569-B. The colour difference decreases at older ages.    We would like to stress that our analytic results for the surface brightness profiles represent an observational best case scenario. From the observational data, the accuracy of the \\rc\\ determination will be limited by photon noise, and perhaps more importantly, by stochastic effects in the stellar IMF. Several studies have quantified the effect of stochastic fluctuations (see for example \\citealt{2000ASPC..211...34L, 2002IAUS..207..616B, 2004A&A...413..145C}). \\citet{2000ASPC..211...34L} determine the age dependent minimum mass a cluster should have, such that the relative fluctuations around the mean flux, $\\sigma_L/L$, is less than 10\\%, corresponding to roughly 0.1 mag photometric uncertainty\\footnote{Photometric uncertainty is calculated form the flux uncertainty in the following way, $\\sigma_V = -2.5\\log(1 + \\sigma_L/L)$ mag.}. When flux is determined in the $V$-band, this minimum mass at ages of [10, 50, 200, 1000]\\,Myr is [$10^5, 2\\,10^4, 10^4, 6\\,10^3$]\\,$\\msun$.  We predict that the most prominent feature of mass segregation in the integrated properties is a 0.1 mag difference in $V-I$ colour (Fig.~\\ref{fig:colours}) between the inner part and the outer part. In order to be able to report a detection of this difference, one needs an uncertainty of $\\sigma_{V-I} <<0.1$ mag and we adopt  $\\sigma_L/L \\simeq 0.01$.   \\citet{2000ASPC..211...34L} show that $\\sigma_L/L$ scales with $(\\sqrt{N}L)^{-1} \\propto N^{-3/2}$, since $L \\propto N$. This implies that the minimum masses quoted above have to be a factor $(0.1/0.01)^{2/3}\\simeq4.6$ higher to be able to distinguish a radial colour variations due to mass segregation from stochastic fluctuations due to IMF sampling. For the 10\\,Myr case this implies a minimum mass of $\\sim 5 \\times10^5\\,\\msun$. This means it will be difficult to detect mass segregation in young massive Galactic clusters, such as the Arches clusters or Westerlund 1. Even the cluster R136 in the 30 Doradus region in the Large Magellanic Cloud is probably of too low mass to detect mass segregation. R136 is probably a special case anyway, since it does not have red evolved stars, which all of our models do have. Young clusters more massive than $\\sim 5 \\times10^5\\,\\msun$ are known, for example in the Antennae galaxies \\citep{1995AJ....109..960W} and M51 \\citep{2000MNRAS.319..893L, 2005A&A...431..905B}, but these are too distant ($\\gtrsim10\\,$Mpc) to be able to determine a colour gradient in the light profile.  There are probably only a handful of candidate clusters that are resolved enough such that a colour gradient can be observed (for example some of the young massive cluster in M82 \\citep{1995ApJ...446L...1O}, the massive ``young globular cluster\" in NGC~6946 \\citep{1967PASP...79...29H, 2002ApJ...567..896L} and a few clusters in M83 \\citep{2004A&A...427..495L}.  An alternative method to detect differences in the IMF between the inner and outer parts would be to determine the spectral properties. If this can be done without changing instrument set-up, then no problems with changing PSFs or weather conditions should affect the observations. It would, therefore, be interesting to investigate which spectral range would be most sensitive to changes of the slope in the IMF."
    },
    "0801/0801.0243_arXiv.txt": {
        "abstract": "We present the mass functions of actively accreting supermassive black holes over the redshift range 0.3 $\\leq z \\leq$ 5 for a well-defined, homogeneous sample of 15,180 quasars from the Sloan Digital Sky Survey Data Release 3 (SDSS DR3) within an effective area of 1644 deg$^2$. This sample is the most uniform statistically significant subset available for the DR3 quasar sample. It was used for the DR3 quasar luminosity function, presented by Richards et al., and is the only sample suitable for the determination of the SDSS quasar black hole mass function. The sample extends from $i$ = 15 to $i$ = 19.1 at $z \\lsim 3$ and to $i$ = 20.2 for $z \\gsim 3$. The mass functions display a rise and fall in the space density distribution of active black holes at all epochs. Within the uncertainties the high-mass decline is consistent with a constant slope of $\\beta \\approx -3.3$ at all epochs. This slope is similar to the bright end slope of the luminosity function for epochs below $z = 4$. Our tests suggest that the down-turn toward lower mass values is due to incompleteness of the quasar sample with respect to black hole mass. Further details and analysis of these mass functions will be presented in forthcoming papers. ",
        "introduction": "We live in exciting times. Since their discovery, the power-house of quasars was believed to be actively accreting supermassive black holes (\\eg Salpeter 1964; Zeldovich \\& Novikov 1964). Owing to recent advances, not only have we been able to confirm the existence of such dense, powerful objects (e.g., Ghez \\et 2005; Genzel \\et 2003; Harms \\et 1994) the determination of their mass has become a common focus of many studies, in spite of the typically non-trivial nature of this process. The reasoning is quite simple: (a) for local galaxies, the early studies revealed a close connection between the black hole mass and the mass, luminosity, and velocity dispersion of the spheroidal component of the host galaxy (\\eg Magorrian \\et 1998; Ferrarese \\& Merritt 2000; Gebhardt \\et 2000a, 2000b) suggesting an intimate connection between their evolutions and (b) recent semi-analytical and hydrodynamical simulations (\\eg Granato \\et 2004; Springel \\et 2005; Cattaneo \\et 2006; Kang \\et 2006) show that black hole evolution and activity may have an important influence on the properties of massive elliptical galaxies. Mapping the demography and mass of the black holes in the centers of galaxies across the history of the Universe is a first step to understanding this connection. Since a fraction of the matter accreted onto black holes as they grow is converted into radiation, the cosmic evolution of the quasar luminosity function helps constrain the accretion and growth history of black holes and can therefore also help us understand the connection with galaxy evolution. Yet certain degeneracies limit what we can learn from the luminosity functions (\\eg Wyithe \\& Padmanabhan 2006) owing to the unknown degree of the radiative efficiency and the mass accretion rate and how these parameters evolve and depend on black hole properties.  In principle, these degeneracies can be broken by combining the luminosity and mass functions.  Recently, Richards \\et (2006, hereafter R06) presented the luminosity function and its cosmic evolution from redshift 5 to redshift 0.3 for a well-defined, homogeneous subset of 15,180 quasars from the Sloan Digital Sky Survey (SDSS) Data Release 3 (DR3) quasar catalog of over 46,400 quasars for which the quasar selection function is well-determined. Therefore, the determination of the mass function for this specific quasar sample is of high interest. In this Letter we present this mass function of actively accreting black holes and its cosmic evolution from redshift 5 to 0.3.  To determine the mass of the black holes we adopt the scaling relations based on the broad emission line widths and nuclear continuum luminosities (\\eg Vestergaard 2002; Warner \\et 2003; Vestergaard \\& Peterson 2006) because (a) they are anchored in robust black hole mass determinations of active black holes in the nearby Universe (Peterson \\et 2004; Onken \\et 2004), (b) several lines of evidence exist in favor of our application of this method to the most distant active black holes (see Vestergaard 2007 for details), and (c) the associated uncertainties on the mass estimates are among the lowest for the available mass estimation methods for distant sources. A cosmology of H$_0$ = 70 ${\\rm km~ s^{-1} Mpc^{-1}}$, $\\Omega_{\\Lambda}$ = 0.7, and $\\Omega_m$ = 0.3 is used throughout. ",
        "conclusions": "For each redshift bin in Figure~\\ref{mfm.fig} the mass function displays a rise and fall in the mass distribution.  The decline toward higher masses is expected to be intrinsic to the quasar population (\\S~\\ref{mbh}); while the uncertainty in the mass estimates will affect the accuracy of the slope value, it is too small to be responsible for the decline. We determine\\footnote{The slope of the mass functions at $z < 2$ is determined for masses above 9.2\\,dex and at $z > 2$ for \\mbh{} $\\geq$ 9.6\\,dex, since these mass bins are deemed (essentially) completely populated. To account somewhat for the uncertainties in the \\mbh{} values and that the mass bin population varies, the uncertainty in a given mass bin was estimated as the propagated uncertainty in the absolute mass estimates (\\eg Table~5 of Vestergaard \\& Peterson 2006) given the Poisson population statistics. } a typical high-end slope $\\beta$ of about $-$3.3 ($\\Psi \\propto M^{\\beta}$) with uncertainties between 0.4 to 1.3 for the mass functions below a redshift of 4. At the mean redshifts of 4.25 and 4.75 the high-end slope `flattens' to $\\beta \\approx -2.9 \\pm 1.1$ and $\\beta \\approx -1.9 \\pm 1.7$, respectively.  The mass function at redshift 4.75 is based on very few black holes at each mass bin, especially for \\mbh{} $\\geq 10^{10}$\\Msol, so the shape is somewhat uncertain. The mass functions are clearly consistent with a constant slope across all epochs.  This is also easily verified in Figure~\\ref{mfm.fig} by the reference curves for the mass functions at $z$=0.49 (black dotted) and at $z$= 2.01 (red dashed)  which also show the space density of high mass black holes slightly increases at $z \\gsim$ 1.  The slope of the mass functions is consistent with the bright end slope of the quasar luminosity function ($\\beta \\approx -3.3$; Croom \\et 2004; R06) to within the uncertainties with the exception that the latter flattens above a redshift of 4 (to $\\beta = - 2.5$; Schmidt \\et 1995; Fan \\et 2001). Another difference is the relative normalizations. The mass functions exhibit a change in space density for a given mass of only a factor of a few between $z$ = 0.5 and $z$ = 2 (cf.\\ the dotted and dashed curves in Fig~\\ref{mfm.fig}), while the space density of quasars brighter than $\\sim$\\,$-$27 mag drops 2 orders of magnitude or more between redshifts 2 and 0.5 (Fig.~18, R06).  While there is no one-to-one correspondence between a given mass value and a given luminosity value since black holes radiate at a range of (Eddington) rates, the different space density changes show that the cumulative volume density of $L/M$ ($ \\propto L/L_{\\rm Edd}$) decreases toward low redshift. This confirms what we already know based on the declining space density of active nuclei at later epochs (\\eg R06): low redshift black holes are typically less active.   In Figure~\\ref{mfm.fig}, a characteristic peak shift to higher mass values with redshift is seen. It is important to establish the reality of these peaks and their shifts, since this would have intriguing cosmological implications. For one, it would be a clear manifestation of cosmic down-sizing: the most massive black holes were more actively accreting early in the universe and with time progressively less massive black holes dominate the population of actively accreting black holes. Alternatively, the down-turn toward lower mass values may simply be due to incompleteness of the sample, as the source selection is based on the source luminosity and broad-band colors, not on black hole mass. As a result, we will not obtain a sharp lower boundary in mass (as we do in the luminosity distribution) but a decrease in space density below the survey flux limits, owing to the distribution of black hole masses at a given luminosity. Unfortunately, the SDSS does not go deep enough for us to make meaningful cuts in black hole mass for the current sample, as will be clear later.   The best way to assess which part of the down-turn is intrinsic requires detailed simulations that include the uncertainties in the mass estimates, the selection function, and the flux limits. Since this is non-trivial we will address this in a forthcoming paper (B.\\ Kelly et al., in preparation). Here, we instead perform crude checks as follows. Firstly, in Figure~\\ref{mfm.fig} we mark the masses of a fiducial source at the mean bin redshift with a continuum luminosity equivalent to the flux limit at that redshift and a broad line width of 2000 \\kms{} (green dashed vertical line) and 3500 \\kms{} (blue dotted vertical line), respectively; 2000 \\kms{} is the canonical lower line width for which a source is a bona-fide quasar and 3500 \\kms{} is the median linewidth for nearby quasars. There are objects below the (green) dashed line since the observed $i$-band magnitude contains emission contribution in addition to the power-law continuum luminosity used for the mass estimates (see \\S~\\ref{data}). Nonetheless, these fiducial masses do yield a guideline location of the SDSS flux limits in the mass parameter space, and their location close to the peak of the mass distribution suggests that the peak shift is simply due to the survey flux limits.  Given the non-trivial relationship between the $i$-band magnitude and the monochromatic luminosity, we also performed a simple simulation with the aim of minimizing selection effects. We generated a mock catalog of quasars covering redshifts between 0.2 to 5.0 and bolometric luminosities between $10^{42}$ \\ergs{} and $10^{48}$ \\ergs{}. For each ($L,z$) grid point we computed the distribution of black hole masses by applying the typical {\\em observed} line width distribution for these $L$ and $z$ values.  The mass functions based on this mock catalog (not shown) do not display the narrowly peaked distributions seen in Figure~\\ref{mfm.fig}, but display a slower continued rise in the distribution below the high mass fast decline and a sharp downturn at the lowest masses which coincides with the lower luminosity bound.  No downturn is seen at the low mass end if a sharp cut in black hole mass is imposed. Hence, this mock catalog confirms the earlier indication that the low-mass downturn in the observed mass functions is a consequence of the relatively narrow luminosity range probed by SDSS broadened by the line width distribution, especially at higher redshift. We thus conclude that the peaks and their shifts with redshift are mainly due to incompleteness of active black holes near the survey flux limits. These crude checks are, however, unable to verify which part of the turnover is trustworthy. The mass functions in the lower redshift bins display a slower turnover, part of which may well be real. B.\\ Kelly et al., (in preparation) will present more advanced simulations that address these and related issues.  The space density decline at $z$\\,$\\approx$\\,2.8 relative to $z$\\,=\\,2.4 and $z$\\,=\\,3.25 in Figure~\\ref{mfm.fig} is caused by a much lower selection efficiency (by a factor of 8; Fig.~6 of R06) as quasars at these epochs coincide with the stellar loci in the selected color-spaces. This suggests that the completeness level is slightly overestimated at these epochs. This situation is unfortunate since the space density of quasars is known to peak between redshifts of 2 and 3 (\\eg Osmer 1982; Warren \\et 1994; Schmidt, Schneider, \\& Gunn 1995; Fan \\et 2001).  It is thus not possible with the current sample to get a complete picture of how the space density distribution of black holes behave during this important epoch. The Large Bright Quasar Survey (hereafter LBQS; Hewett \\et 1995) of about 1000 quasars, extending to a redshift of 3, offers an opportunity to address these issues further, since this sample is not selected based on broad-band color. The LBQS quasar mass function will be presented by Vestergaard \\& Osmer (in preparation) and discussed in concert with the DR3 mass function by M.\\ Vestergaard et al.\\ (in preparation), where the mass functions will also be compared to previous work."
    },
    "0801/0801.2300_arXiv.txt": {
        "abstract": "{We present the characteristics of the X-ray variability of stars in the  cluster \\N as derived from  XMM-Newton/EPIC/pn data. The  X-ray variations on short (hours), medium (months), and long (years) time scales have been explored. We detected  303 distinct X-ray sources by analysing six EPIC/pn observations; 194 of them are members of the cluster. Stars of all spectral types, from the early-types to the late-M dwarfs, were detected.  The Kolmogorov-Smirnov test applied to the X-ray photon time series shows that, on short time scales, only a relatively small fraction (ranging from 6\\%  to 31\\% for dG and dF, respectively) of the members of \\N are variable with a confidence level $\\geq$99\\%; however, it is possible  that the fraction is small only because of the poor statistics. The time X-ray amplitude distribution functions (XAD) of a set of dF7-dK2 stars, derived on short (hours) and medium (months) time scales, seem to suggest that medium-term variations, if present, have a much smaller amplitude than those on short time scales; a similar result is also obtained  for dK3-dM stars. The amplitude variations of late-type stars in \\N are consistent with those of the coeval Pleiades stars. Comparing these data with those of ROSAT/PSPC, collected 7-8 years earlier, and of ROSAT/HRI, just 4-5 years earlier, we find no evidence of significant \\v on the related time scales, suggesting that long-term variations due to activity cycles similar to the solar cycle are not common among young stars. Indications of spectral \\v was found in one star whose spectra at three epochs were available. ",
        "introduction": "Observations using the {\\em Einstein} observatory, nearly  30 years ago, showed the ubiquity of stellar X-ray emission throughout the HR diagram \\citep{Va81}. X-ray emission from stars has been attributed to several physical mechanisms depending on the mass of the star. Going from high to low mass, they include winds, % dynamo, and turbulent dynamo. Variability studies are one of the best ways to explore these mechanisms and to infer physical conditions of the regions where X-ray emission originates and their evolution along stellar life  \\citep[e.g.][]{Mont83, Amb87, Ster95, Mar03, Mar03a, Pi05}. In this context, open clusters, providing large, chemically homogeneous, and precisely dated samples of stars, are ideal laboratories for studying coronal emitters and for constraining the X-ray-generating mechanisms. Comparative studies of the \\v properties of homogeneous classes are excellent for diagnosing and  determining the origin of the X-ray emission.  In this perspective we have analysed the X-ray variability properties of  homogeneous samples (with respect to mass and age): dF7-dK2 and  dK3-dM stars in the solar  neighbourhood, dF7-dK2, and dK3-dM of Pleiades observed with ROSAT-PSPC \\citep{Mar00, Mar02, Mar03a}. These studies  show that the short-term \\v of solar-type stars decreases with age (together with luminosity), while long-term \\v increases. In  contrast, the short-term \\v of low mass stars is present in all stellar life phases while long-term variations have never been observed.  Also called the \"Southern Pleiades\", NGC 2516 is a rich and young open cluster in the constellation of Carina, that has been observed several times at the XMM-Newton observatory, thereby allowing us the study of X-ray variability properties on multiple time scales. At a distance  of less than 400 pc (387 pc, \\citealt{Jef97}; 346 pc, \\citealt{Robi99}), the  cluster contains about 1300 known members spanning all spectral types and enabling  simultaneous study of the different processes driving X-ray emission from stars with different internal structures. The metallicity of the cluster is controversial; several photometric studies have suggested a metal underabundance by a factor of 2 with respect to the Sun \\citep[e.g.][]{Jef97, Jef98, Jeff01, Pin98}, while \\citet{Ten02}  have found  solar-like metal abundance on the basis of high-resolution spectra. Since metallicity affects the depth of the convection zone, which in turn should influence the dynamo's efficiency, observable differences in the X-ray emission levels may be expected between solar and sub-solar abundances \\citep{Mice00}. The ROSAT  observations have shown that the G and K members in \\N  were underactive in X-rays with respect to the Pleiades G and K stars, while no difference was found for M stars.  As discussed by \\citet{Jef97} and \\citet{Mice00}, these results could be attributed to a low-metallicity effect, although when taking the recent optical data into account \\citep{Ten02} these explanations are questioned. Also a recent  analysis of summed XMM/EPIC data \\citep{Pil05} shows that late-types stars in \\N are significantly less luminous than those of the Pleiades.  \\N has been observed several times with Chandra and the analysis of the observations was made by \\citet{Har01} and \\citet{Dam03}. \\citet{Wo04} used Chandra data for a timing analysis of \\N stars, finding that the stochastic \\v  rate is similar for all sources in their sample, while the time scale of \\v is shorter for later-type stars.   In this paper we present a study of the X-ray variability  of the NGC 2516 stars analysing six XMM-Newton/EPIC observations spanning  $\\sim$ 19 months. We explore X-ray variability properties on short (hours), medium (months), and long (years) time scales, we compare our data with those obtained for the coeval Pleiades, and we search for spectral variability.   The structure of the paper is the following:  Sect. 2 describes the XMM-Newton/EPIC observation set  and the data  analysis; Sect. 3 and 4 present the time and the spectral  analysis, respectively; Sect. 5 summarises the main results. ",
        "conclusions": ""
    },
    "0801/0801.1629_arXiv.txt": {
        "abstract": "{An analytical solution of the generalized diffusive and convective transport equation is derived to explain the transport of cosmic ray protons within elliptical galaxies.} {Cosmic ray transport within elliptical galaxies is an interesting element in understanding the origin of high energetic particles measured on Earth. As probable sources of those high energetic particles, elliptical galaxies show a dense interstellar medium as a consequence of activity in the galactic nucleus or merging events between galaxies. Thus it is necessary for an appropriate description of cosmic ray transport to take the diffusive and convective processes in a dense interstellar environment into account. Here we show that the transport equations can be solved analytically with respect to the given geometry  and boundary conditions in position space, as well as in momentum space.} {From the relativistic Vlasov equation, which is the most fundamental equation for a kinetic description of charged particles within the interstellar medium in galaxies, one finds a generalized diffusion-convection equation in quasilinear theory. This has the form of a `leaky box' equation, meaning particles are able to escape the confinement region by diffusing out of the galaxy. We apply here the `diffusion approximation', meaning that diffusion in gyrophase and pitch angle are the fastest particle-wave interaction processes. An analytical solution can be obtained using the `scattering time method', i.e. separation of the spatial and momentum problems. } {The spatial solution is shown using a generalized source of cosmic rays. Additionally, the special case of a jet-like source is illustrated. We present the solution in momentum space with respect to an escape term for cosmic ray protons depending on the spatial shape of the galaxy. For a delta-shape injection function, the momentum solution is obtained analytically. We find that the spectral index measured on Earth can be obtained by appropriately choosing of the strength of Fermi I and Fermi II processes. From these results we calculate the gamma-ray flux from pion decay due to proton-proton interaction to give connection to observations. Additionally we determine the escape-spectrum of cosmic rays. The results show that both spectra are harder than the intrinsic power-law spectrum for cosmic rays in elliptical galaxies. } {} ",
        "introduction": "Since their discovery by Viktor Hess in 1912, cosmic rays have been one of the biggest fields of interest in astrophysics, and yet the origin of these particles is still an open question. Fully ionized atomic nuclei reach the Earth coming from outside the solar system with very high energies up to $10^{20}$ eV. Most of them with energies $< 10^{17}$ eV seem to originate in the Milky Way, while the highest energetic ones are considered to have an extragalactic origin \\citep[for a review see][]{Hoer}.  The most accepted model for the origin of ultrahigh energy cosmic rays (UHECRs) is acceleration in shock fronts due to Fermi processes \\citep[]{Fermi49}; hence, the main sources are gamma ray bursts (GRBs), active galaxies like active galactic nuclei (AGN) \\citep[]{Tav}, or colliding galaxies. The last two are specially interesting for two reasons. First UHECRs can be generated in elliptical galaxies. Second the increase in the interstellar medium density in objects of these types influences cosmic ray transport \\citep[see][]{BekkiShioya}. But since all elliptical galaxies have an interstellar medium due to star winds, we conclude that the study of transport processes is interesting in general \\citep[see][]{Knapp}.  In particular, as a result of the GZK-effect, very close AGN are the most probable candidates for UHECR sources \\citep[see][]{Bier}. One of these nearby AGN is the giant elliptical galaxy M87. It is proposed that this galaxy is responsible for acceleration of cosmic ray protons due to Fermi I processes \\citep{Bland,Rieger} in shockfronts within the jet \\citep{Reimer}. For an overview of proton acceleration in jets \\citep[see][]{Mannheim}. Since particle acceleration in a jet is located within active galaxies surrounded by an interstellar medium, the high energetic protons undergo physical transport processes before they escape out of the galaxy and reach the detectors on Earth, making a model for cosmic ray transport in elliptical galaxies inevitable. This helps give an answer to the question about the origin of cosmic rays.  Progress has been made in the field of modeling cosmic ray transport with the numerical description of transport processes by \\citet{OwJok} and \\citet{StrMos}. Nevertheless we follow the basic ideas presented in the underlying papers of \\citet{LS85}, \\citet{WangSchlick}, and \\citet{LS88}, who used an analytical description of cosmic ray transport. In relation to our work such, a treatment has the following advantages: the model is adequate for cosmic ray transport within any kind of elliptical galaxy including arbitrary cosmic ray sources and the physical parameters involved in our model can be easily fitted to measurements. After all, our analytical model can serve as a test case for more profound numerical models.  In this paper we solve the cosmic ray transport equation analytically with respect to a kinetic description of the interstellar plasma in elliptical galaxies. Special attention is paid to the spatial transport of charged nuclei. In addition, the solution of the momentum equation is derived to explain general properties of this model. As a result, we present illustrative examples of spatial, as well as momentum, cosmic ray transport for given sources of charged nuclei. To show the connection to observations, we calculate the gamma-ray flux from neutral pion decay. These mesons are produced by inelastic scattering processes between cosmic ray protons. The resulting power-law spectrum is slightly harder than the intrinsic one for cosmic rays. Finally, we present the escape spectrum of charged particles leaving elliptical galaxies. Similar to the gamma-ray flux, this spectrum is flatter than the intrinsic one. ",
        "conclusions": "We showed that an analytical treatment of cosmic ray transport in elliptical galaxies based on the diffusion approximation is possible in general. The formal solution we found, combined with appropriate boundary conditions and source functions, can be used to study transport processes in elliptical galaxies. This model is valid for the complete physical parameter space as long as the diffusion approximation holds.  The first test case where we applied our model to is a jet-like injection shape like in M87. As we can separate our problem into spatial and momentum problems, this test case probes the spatial problem. For this model we found that for very long timescales cosmic rays are distributed throughout the galaxy almost isotropically.  We also provided a test case for the momentum problem. Under the assumption of a resulting power-law spectrum matching the observed power law-spectrum on Earth, we could identify the governing eigenfunctions and therefore the governing processes for cosmic ray acceleration and momentum diffusion in  elliptical galaxies. It turned out that in elliptical galaxies adiabatic losses are responsible for the high energy cutoff, whereas the slope of the spectrum is given by the ratio of the strength of Fermi I and Fermi II processes. As a result of the basic possibility of explaining this power law-with physical parameters that might be found in M87, our calculation gives rise to the theory of M87 as a source of ultra-high energy cosmic rays. In this context the gamma-ray flux from pion decay and the escaping cosmic ray spectrum are essential for a more adequate view of cosmic ray astrophysics. We find that the gamma-ray spectrum with a power law index of $-2.69$ is a bit harder than the intrinsic cosmic ray spectrum (power law index $-2.8$). In the context of understanding the sources of high-energy cosmic rays, the flux of charged particles escaping from the confinement region of elliptical galaxies is also very interesting. For the same intrinsic spectrum as above, we find $-2.13$ as the spectral index for escaping particles. Again the spectrum is harder than the value within elliptical galaxies. \\\\ Our model may also serve as a testbed for some more advanced questions: \\begin{itemize} \\item Are elliptical galaxies, and especially AGN within these objects, reasonable sources of the ultrahigh energy cosmic rays? \\item What is the total flux leaving a single elliptical galaxy with respect to a measured source of cosmic rays in elliptical galaxies? \\item What is the total cosmic ray flux we can expect to measure on Earth in the special case of the nearby active elliptical galaxy M87 or Centaurus A? \\end{itemize} Clearly, to answer these questions, a detailed quantitative analysis of the properties of cosmic ray transport within elliptical galaxies is required. Therefore we need to determine from observations the exact physical parameters of the interstellar medium and of the emission processes within the source. With our analytical model at hand, the task of constraining the possible parameter space gets easier. Especially by assuming that UHECR are coming mostly from sources like M87, we may be able to derive the total cosmic ray content of this type of elliptical galaxies.  Nevertheless this work is considered as a starting point for more sophisticated (numerical) models. Another interesting point is the expansion of this model to cosmic ray electrons. While the spatial description may be derived analogously, we need to add the physical processes involved in the transport of electrons in momentum space. Synchrotron losses and the inverse Compton effect will then play a major role \\citep[see, e.g.,][]{Casadei}. The output of this model can be easily tested with the radio data of elliptical galaxies since the electrons provide the main contribution to this radiation."
    },
    "0801/0801.3649.txt": {
        "abstract": "{} % aims heading (mandatory) {We present the first X-ray observations with the XMM-Newton and INTEGRAL satellites of the recently discovered cataclysmic variable 1RXS\\,J173021.5-055933, together with simultaneous UV and coordinated optical photometry aiming at characterising its broad-band  temporal and spectral properties and classifying this system as a magnetic one.} % methods heading (mandatory) {We performed a timing analysis of the X-ray, UV, and optical light curves to identify and to study the energy dependence of the fast 128\\,s  pulsation over a wide energy range. X-ray spectral analysis in the  broad 0.2-100\\,keV X-ray range was performed to characterise the peculiar emission properties of this source. } % results heading (mandatory) {We find that the X-ray light curve is dominated by the spin period of the accreting white dwarf in contrast to the far-UV range, which turns out to be  unmodulated at a 3$\\sigma$ level. Near-UV and optical pulses are instead detected at twice the spin frequency. We identify the contributions from  two accreting poles that imply a moderately inclined dipole field allowing, one pole to dominate at energies at least up to 10\\,keV, and  a secondary that instead is negligible above 5\\,keV. X-ray spectral analysis reveals the presence of multiple emission components consisting of optically thin plasma  with temperatures ranging from 0.17\\,keV to 60\\,keV and a hot blackbody  at $\\sim$90\\,eV. The spectrum is also strongly affected by peculiar absorption components consisting of two high-density ($\\rm \\sim 3\\times 10^{21}\\,cm^{-2}$ and $\\rm 2\\times 10^{23}\\,cm^{-2}$)  intervening columns, plus a warm absorber. The last is detected from an  OVII absorption  edge at 0.74\\,keV, which suggests that photoionization of pre-shock material is also occurring in this system. } % conclusions heading (optional), leave it empty if necessary {The observed properties indicate that the accretor in 1RXS\\,J173021.5-055933 is a white dwarf with a likely weak magnetic field, thus confirming this cataclysmic variable as an intermediate polar (IP) with one of the most extreme spin-to-orbit period ratios. This system also joins the small group of IPs showing a soft X-ray reprocessed component, suggesting that this characteristics is not uncommon in these systems. } ",
        "introduction": "Cataclysmic variables (CVs) are close binaries containing an accreting white dwarf (WD) from a late type Roche-lobe filling (main sequence or subgiant) secondary star (for a comprehensive review see Warner \\cite{Warner95}). Among them, the magnetic systems (mCVs) nowdays constitute a rather conspicuous group ($\\sim$25$\\%$ of all CVs).  The current census of mCVs, including confirmed ones and candidates, amounts to about 150 systems (Ritter \\& Kolb (\\cite{RitterKolb06}). They are further divided in two classes:  the polars, containing a strongly magnetised ($>$10\\,MG) WD, whose rotation is synchronised with the orbital period and the intermediate polars (IPs),  which contain accreting WDs that are rotating much faster than the orbital period. The presence of X-ray pulses at the WD rotation but the lack of detectable polarized circular radiation in the optical/near-infrared in most systems, led to believe that IPs possess weakly magnetised WDs (B$\\lesssim $10\\,MG) (see review by Patterson \\cite{patterson}). However this is still the subject of great debate (see below). As of today, the polars consist of $\\sim$58$\\%$ of the known mCVs. Because of their strong soft X-ray ($\\sim$40\\,eV) emission component, the polars increased in number after the ROSAT era (Beuermann \\cite{Beuermann99}).  This was not the case for the IPs, which are  dominated  by a hard 10-20\\,keV optically thin emission. However in the last few years a large number ($\\sim$34) of IP candidates were  discovered from ground-based observations (Warner \\& Woudt \\cite{WarnerWoudt04}, Rodriguez-Gil et al. \\cite{Rodriguez-Gil05}, G\\\"ansicke et al. \\cite{Gaensicke05}, Bonnet-Bidaud et al. \\cite{BonnetBidaud06}, Southworth et al. \\cite{Southworth06}, Southworth et al. \\cite{Southworth07}).   Also, hard X-ray surveys such as those carried out by the {\\em INTEGRAL} and {\\em SWIFT} satellites are providing new and interesting results on CVs (Wheatley et al. \\cite{Wheatley2006}, Masetti et al. \\cite{masetti06a}, Masetti et al. \\cite{masetti06b}). In particular, from the current {\\em INTEGRAL} catalogue (Bird et al. \\cite{Bird07}) about 5$\\%$ of the detected sources are CVs, most of them are IPs (Barlow et al. \\cite {Barlow06},  Falanga et al. \\cite{falanga}). This has  further stimulated  optical and X-ray follow-ups to search for new IPs (e.g. Bonnet-Bidaud et al. \\cite{BonnetBidaud07}). The detection of high energy emission of IPs extending up to  $\\sim$90\\,keV was already identified from {\\em BeppoSAX} observations of a few bright systems (de Martino et al. \\cite{demartino01}, de Martino et al. \\cite{demartino04}), which makes the new discoveries in the hard X-rays very promising to understand the role of this, formerly, small group of mCVs.  The link between polars  and IPs, and hence the evolution of mCVs, is still uncertain. Polars are typically found at short orbital periods, most of them below the so-called 2-3\\,hr orbital period gap. IPs instead populate the orbital period distribution at longer periods with a handful of systems below the gap. The majority of IPs have rotational periods ($\\rm P_{spin}$) about 7-10$\\%$ the binary orbital periods ($\\rm P_{orb}$). However, including the new optical candidates  a  wide range of degree of asynchronism,  $\\rm P_{spin}/P_{orb}$ from $\\sim$0.002 to $\\sim 0.7$, is encountered (Woudt \\& Warner \\cite{WarnerWoudt04}, de Martino et al. \\cite{demartino06a}).  This might suggest that the former  separation between IPs and polars could be only apparent and reinforces the debate on the hypothesis that IPs possess similar magnetic fields to those of the low-field polars, and will synchronise while evolving towards short orbital periods (Cumming \\cite{Cumming}, Norton et al. \\cite{norton04}, Southworth et al. \\cite{Southworth07}). The confirmation of new optically discovered candidates is therefore essential in this respect. Differently from the optical range, heavily affected by X-ray reprocessing, a secure identification comes from the detection of X-ray pulses at the WD spin period, which arise from the post-shock region of the magnetically confined accretion flow onto the WD.  In the framework of a programme with {\\em XMM-Newton} aiming at unambiguously confirming the IP nature of new optically identified systems, we here present the first X-ray study of 1RXS\\,J173021.5-055933 (hereafter RXJ\\,1730). Identified as a magnetic CV by  G\\\"ansicke et al. (\\cite{Gaensicke05}), its very fast (128\\,s)  optical pulsations and very long (15.4\\,hr)  orbital period make it an extreme object in terms of degree of asynchronism.  RX\\,J1730 has the second fastest spinning WD preceeded only by AE\\,Aqr with a spin period of 33\\,s. Its orbital period is the second longest ever observed in mCVs, being headed by GK\\,Per  with a period of 47.9\\,hr.  While GK\\,Per and AE\\,Aqr likely harbour  K subgiant secondary stars (Morales-Rueda et al. \\cite{morales-rueda}, Casares et al. \\cite{casares}) the  donor star in RX\\,J1730 was suggested  to be a G-type  main sequence star, and differently from these two, the optical spectrum is dominated by the accretion flow rather than by  the donor star (G\\\"ansicke et al. \\cite{Gaensicke05}). The  nature and evolutionary path of very long orbital period systems is still unknown though of great importance to understand the evolution of mCVs as a whole.  With the purpose of identifying unambiguously the rotational period of the WD primary   and  to study the emission and accretion properties of this system, the analysis presented here include a temporal and spectral analysis of the {\\em XMM-Newton} observation. It was also detected in the hard X-rays during the IBIS/ISGRI galactic plane survey as IGR\\,J17303-0601=INTEGRAL1 56 (Bird et al. \\cite{Bird04}) and formerly proposed to  be a low mass X-ray binary (LMXB) by Masetti et al. (\\cite{masetti04}). Hence, we included in the  spectral  analysis also the  archival data from the {\\em INTEGRAL} satellite. We further   complement the X-ray study with simultaneous far-UV photometry  from the  {\\em XMM-Newton} Optical Monitor as well as coordinated quasi-simultaneous near-UV/optical photometry acquired at the {\\em William Herschel Telescope} in La Palma.     %__________________________________________________ one column table \\begin{table*}[t!] \\caption{Summary of the observations of RX\\,J1730.} \\label{obslog} \\centering \\begin{tabular}{l  c c r c c} \\hline \\hline \\noalign{\\smallskip} Instrument     &  Date & UT(start) & Duration (s) & Net Count Rate (cts\\,s$^{-1}$)\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} {\\em XMM-Newton} & & & & \\\\ EPIC PN &  2005 Aug 29 & 07:05 & 11271 & 3.73\\\\ EPIC MOS & & 06:37 & 13275 & 1.02 \\\\ RGS    & & 06:36 & 13507 & 0.12  \\\\ OM UVW1 & & 06:45 & 2319 & 2.57\\\\ & & 07:29 & 2320 & \\\\ & & 08:13 & 2320 & \\\\ & & 08:58 & 2320 & \\\\ & & 09:42 & 2319 & \\\\ \\hline \\noalign{\\smallskip} Instrument     &  Date & N. Science Window & Duration$^a$ (s) & Net Count Rate (cts\\,s$^{-1}$)\\\\ \\hline \\noalign{\\smallskip} {\\em INTEGRAL}/IBIS & & & & 0.57\\\\ & 2003 Apr 23-30 & 75 & 135000 & \\\\ & 2003 May 1 & 4 & 9182 & \\\\ & 2003 Aug 30 & 5 & 17757 & \\\\ & 2003 Sept 3 & 3 & 8637 &\\\\ & 2003 Oct 10-17 & 95  & 172000 & \\\\ & 2004 Feb 15-25 & 28 & 61800 & \\\\ & 2004 Mar 11-20 & 35 & 66700 & \\\\ & 2004 Apr 2     &  6 & 19875 & \\\\ \\hline \\noalign{\\smallskip} Instrument     &  Date & UT(start) & Duration (s) & r'\\,(mag)\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} {\\em William Herschel Telescope} & & & & \\\\ ULTRACAM/u',g',r' & 2005 Aug 27 & 20:13 & 2567 & 15.6 \\\\ & 2005 Aug 28 & 20:30 & 1925 & 15.6 \\\\ & 2005 Aug 29 & 20:15 & 3476 & 15.4 \\\\ & 2005 Aug 30 & 20:19 & 1255 & 15.3 \\\\ & 2005 Aug 31 & 20:26 & 992  & 15.3 \\\\ & 2005 Sep 1 &  20:11 & 1640 & 15.3 \\\\ \\noalign{\\smallskip} \\hline \\hline \\end{tabular} \\begin{flushleft} $^a$: Total exposure time \\end{flushleft} \\end{table*}   %__________________________________________________________________ ",
        "conclusions": "We  presented the first  X-ray study of RX\\,J1730 based on an {\\em XMM-Newton} observation, complemented with {\\em INTEGRAL} and optical fast photometric data that confirm this CV as a magnetic  system of the IP type.  The rapid 128\\,s pulsation is energy dependent with remarkable differences from the X-rays, the UV to the optical ranges. We identified two emitting poles with different contributions. A  main emitting pole contributes from soft to the hard X-rays, while the secondary pole contributes most in the soft bands. The WD magnetic dipole field should be moderately inclined, thus allowing the visibility of both poles, with the lower pole better represented. The optical spin pulse is mainly due to  reprocessing in the accretion flows onto the WD and dominated by the outer regions. The lack of detectable UV pulsation is an unexpected property, which might indicate that reprocessing at both the WD surface and in the inner regions of the accretion flow is masked by the contribution of the outer flow regions.  X-ray spectral analysis reveals a hard X-ray spectrum as well as remarkable soft X-ray complexities. These can be identified in a warm absorber, which makes this system the second IP where an absorption edge from ionised oxygen has been revealed, together with multiple optically thin emissions and a blackbody  component. The last adds RX\\,J1730 to the small group of soft X-ray IPs. We however find a large luminosity for this component that is difficult to explain. The observed X-ray properties are consistent with a system hosting a weakly magnetised WD rather than a NS. The long 15.4\\,hr orbital period and rapid 128\\,s WD rotation could suggest that RX\\,J1730 is a young IP born with a rapidly rotating WD. To understand its evolutionary  status, this system deserves further observations especially  addressing the donor star parameters."
    },
    "0801/0801.1135_arXiv.txt": {
        "abstract": "As part of the Center for Adaptive Optics  (AO) Treasury Survey (CATS) we imaged a set of 15 intermediate redshift ($z\\sim0.8$) luminous infrared galaxies (LIRGs) with the Keck Laser Guide Star (LGS) AO facility.  These galaxies were selected from the Great Observatories Origins Deep Survey (GOODS) southern field, allowing us to combine the high spatial resolution \\HST\\ optical ($B$, $V$, $i$, and $z$-bands) images with our near-infrared ($K'$-band) images to study the LIRG morphologies and spatially resolved spectral energy distributions (SEDs).  Two thirds of the LIRGs are disk galaxies, with only one third showing some evidence for interactions, minor, or major mergers. In contrast with local LIRG disks (which are primarily barred systems), only 10\\% of the LIRG disks in our sample contain a prominent bar. While the optical bands tend to show significant point-like substructure, indicating distributed star formation, the AO $K$-band images tend to be smooth.  They lack point-like structures to a $K \\sim 23.5$ limit  This places an upper bound on the number of red super giants per blue-knot at less than 4000.  The SEDs of the LIRGs are consistent with distributed dusty star formation, as exhibited by optical to IR colors redder than allowed by old stellar populations alone.   This effect is most pronounced in the galaxy cores, possibly indicating central star formation.  We also observed a set of 11 intermediate redshift comparison galaxies, selected to be non-ellipticals with apparent $K$-band magnitudes comparable to the LIRGs. The ``normal'' (non-LIRG) systems tended to have lower optical luminosity, lower stellar mass, and more irregular morphology than the LIRGs.   Half of the ``normal'' galaxies have SEDs consistent with intermediate aged stellar populations and minimal dust.  The other half show evidence for some dusty star formation, usually concentrated in their cores. Our work suggests that the LIRG disk galaxies are similar to large disk systems today, undergoing self regulated star formation, only at 10 - 20 times higher rates. ",
        "introduction": "\\begin{deluxetable*}{ccccccc} \\tabletypesize{\\small} \\tablecaption{Keck Observation Summary \\label{table:ao_obs}} \\tablehead{\\colhead{object} & \\colhead{R.A. } & \\colhead{Dec.} & \\colhead{obs. date} & \\colhead{exptime } & \\colhead{FWHM}& \\colhead{Strehl} \\\\ & \\colhead{(J2000)}& & & \\colhead{(m)} & \\colhead{of PSF ($\\arcsec$)} &\\colhead {estimate}} \\startdata LIRG 1&              03:32:13.86&              -27:42:48.9&         Sept. 29,  2005 &      44&    0.10&    0.30\\\\ LIRG 2&              03:32:13.24&              -27:42:40.9&         Sept. 29,  2005 &      44&    0.10&    0.30\\\\ LIRG 3&              03:32:12.81&              -27:42:01.5&          Nov. 25,  2005 &      55&    0.08&    0.35\\\\ LIRG 4&              03:32:32.99&              -27:41:17.0&          Dec. 11,  2006 &      50&    0.11&    0.25\\\\ LIRG 5&              03:32:32.90&              -27:41:23.9&          Dec. 11,  2006 &      50&    0.14&    0.18\\\\ LIRG 6&              03:32:31.48&              -27:41:18.1&          Dec. 11,  2006 &      50&    0.11&    0.25\\\\ LIRG 7&              03:32:31.05&              -27:40:50.1&          Dec. 11,  2006 &      50&    0.11&    0.20\\\\ LIRG 8&              03:32:32.46&              -27:40:56.6&          Dec. 11,  2006 &      50&    0.11&    0.25\\\\ LIRG 9&              03:32:52.88&              -27:51:19.8&          Dec. 11,  2006 &      53&    0.12&    0.18\\\\ LIRG 10&              03:32:52.87&              -27:51:14.7&          Dec. 11,  2006 &      53&    0.12&    0.18\\\\ LIRG 11&              03:32:53.02&              -27:51:04.7&          Dec. 11,  2006 &      53&    0.11&    0.20\\\\ LIRG 12&              03:32:07.98&              -27:42:39.4&         Sept. 28,  2005 &      60&    0.14&    0.14\\\\ LIRG 13&              03:32:55.88&              -27:53:27.9&          Nov. 12,  2005 &      30&    0.12&    0.15\\\\ LIRG 14&              03:32:38.12&              -27:39:44.8&          Dec. 11,  2006 &      42&    0.10&    0.24\\\\ LIRG 15&              03:32:55.69&              -27:53:45.9&          Nov. 12,  2005 &      30&    0.09&    0.30\\\\ COMPGAL 1&              03:32:13.06&              -27:42:04.9&          Nov. 25,  2005 &      55&    0.08&    0.35\\\\ COMPGAL 2&              03:32:30.93&             -27:40:58.43&          Dec. 11,  2006 &      50&    0.11&    0.20\\\\ COMPGAL 3&             03:32:53.428&             -27:50:51.06&          Dec. 11,  2006 &      53&    0.10&    0.24\\\\ COMPGAL 4&              03:32:52.33&              -27:50:51.8&          Dec. 11,  2006 &      53&    0.11&    0.20\\\\ COMPGAL 5&              03:32:09.40&              -27:42:36.1&         Sept. 28,  2005 &      60&    0.12&    0.15\\\\ COMPGAL 6&              03:32:08.43&              -27:42:43.5&         Sept. 28,  2005 &      60&    0.12&    0.15\\\\ COMPGAL 7&              03:32:55.95&              -27:53:57.1&          Nov. 12,  2005 &      30&    0.10&    0.25\\\\ COMPGAL 8&             03:32:13.458&             -27:41:49.89&          Nov. 25,  2005 &      55&    0.08&    0.35\\\\ COMPGAL 9&              03:32:36.17&              -27:39:55.7&          Dec. 11,  2006 &      42&    0.12&    0.20\\\\ COMPGAL 10&              03:32:36.68&              -27:39:54.6&          Dec. 11,  2006 &      42&    0.10&    0.30\\\\ COMPGAL 11&             03:32:55.003&             -27:50:51.63&          Dec. 11,  2006 &      53&    0.11&    0.20\\\\ \\enddata \\end{deluxetable*}  Luminous infrared galaxies \\citep[LIRGs; for review see ][]{SandersMirabel96} are defined to have total IR luminosities ($8 - 1000 \\mu m$) in excess of $10^{11}L_{\\odot}$.  At low redshift, LIRGs are rare, comprising less than 5\\% of the total IR energy density of local galaxies \\citep{Soifer91}.  However, by $z=1$, LIRGs account for $\\sim70\\%$ of the IR background \\citep{LeFloch05}.  The high IR luminosities of LIRGs result from thermal dust heating in starbursts, or active galactic nuclei (AGN).  Dust absorbs optical and UV radiation, and reemits that energy in a broad thermal peak in the mid to far IR.  In order for a star forming galaxy to emit at a LIRG level it must have a  very high star formation rate, in excess of $\\sim17 M_{\\odot} \\; yr^{-1}$.  In the local universe, these high star formation rates are primarily triggered by galaxy-galaxy interactions or mergers \\citep{Ishida04}.  However, the rise in the number density of LIRGs with redshift appears to be unrelated to any change in the merger rate.  For instance \\citet{Bell05,MKL05, Lotz07} all show that roughly 50\\% of intermediate redshift LIRGs are disk galaxies with little sign of recent merger activity.  While morphology rules out major mergers as triggers for the majority of intermediate redshift LIRGs, interactions with neighbors, minor mergers, or bar instabilities may contribute.  In addition, intermediate redshift disks may contain higher gas fractions than their local counterparts, allowing otherwise normal galaxies to reach higher levels of star formation.  Some star formation triggers have observable signatures.  For instance interactions, minor mergers and bar instabilities will tend to drive central star formation.  Centrally concentrated star formation is a feature of local LIRGs \\citep{SandersMirabel96}.  Minor mergers may also show elevated star star formation in a dwarf companion which may appear as a single bright blue knot at these redshifts \\citep{Ferreiro2004}.  Minor mergers and interactions may also produce tidal features \\citep{Knierman05} or drive bar instabilities.  Low redshift LIRG disks are found to uniformly contain bars \\citep{Wang06}.  Interestingly, \\citet{Zheng05} found that intermediate redshift LIRGs are relatively free of bars.  Their study, however, was done with optical (rest-frame UV) imaging and therefore may not be sensitive to bars, which in the presence of dust will be more easily observed at rest frame red to near-IR wavelengths.   While centrally concentrated star formation is a feature of low redshift LIRGs \\citep{SandersMirabel96}, it may be less important for intermediate redshift LIRGs.  If the intermediate redshift LIRG disks are experiencing smoothly distributed star formation, rather than centrally concentrated star formation, then there may be no need to invoke a ``trigger'' for the event.  These galaxies may tend to high star formation rates simply because of higher gas fractions.  If that is the case, then higher star formation rates may actually be normal for the large disks of that epoch.  It is currently an open question, whether the intermediate redshift LIRGs are ``triggered'' or experiencing ``normal'', if elevated, star formation.  In a series of 3 papers we discuss the modes of star formation in a set of 15 intermediate redshift LIRGs and 11 comparison galaxies.  These galaxies are located in the Great Observatories Origins Deep Survey  \\citep[GOODS]{giavalisco04}, in fields for which we have obtained high spatial resolution IR imaging from the Keck laser guide star (LGS) adaptive optics (AO) facility .    This paper will present the \\HST\\ (from GOODS) and AO images, basic photometric measurements, and will discuss the morphologies of the LIRG and comparison galaxy samples.  It will seek to identify ``triggers'' (e.g. mergers, interactions, bars, or AGN) for the LIRG events, and  will present the spectral energy distributions of subcomponents of galaxies in an attempt to identify sites of dusty star formation.     We compare the results from the LIRG sample with the set of ``normal'' (i.e. non-LIRG) galaxies drawn from the same fields.  Paper 2 will attempt to ``fit'' the SEDs of the galaxy sub-components with specific stellar population synthesis models.  Paper 2 will compare the stellar populations of the central and outer regions of each galaxy in an effort to identify where the bulk of the star formation is occurring.  It will also attempt to locate the bulk of the dust and compare to the global dust measurements from \\emph{Spitzer}\\ \\tf\\ images.  Paper 3 will expand our work on the morphologies of the LIRGs.  It will present analysis of the Sersic profiles of each galaxy.  It will then look for substructures hidden beneath the Sersic profiles, specifically looking for small central bars, or tidal features not immediately obvious in the images.  While the sample size is small, 15 LIRGs and 11 ``normal'' galaxies, it is a unique data set.  It is the only such set of intermediate redshift galaxies that has been observed in 4 optical bands with \\HST\\ and with equally high spatial resolution in the near-IR.  The point-spread-function of our $K$-band AO images has a  full-width half-max of $0.1\\arcsec$ comparable to \\HST\\ in the optical, and a factor of 3 improvement over \\HST\\ NICMOS.  The high spatial resolution is necessary to study galaxy subcomponents individually.  The Keck IR band is necessary to distinguish between old and young-dusty stellar populations, which are degenerate cases given the optical \\HST\\ data alone \\citep{Melbourneetal05}.  \\begin{figure*} \\includegraphics{f1.ps} \\caption{\\label{fig:ao_field} A typical CATS pointing in GOODS-S.  The left panel shows the Keck AO pointing.  The tip-tilt star is in the bottom, left-hand corner of the image.  North is up, east is to the left.  During each observation the laser was positioned at the center of the frame.  As the telescope dithered between frames the laser moved with the telescope.  The right hand panel is the same as the left only it is a three color image. combining the Keck AO $K$-band(red) and HST $i$ and $B$-band (green and blue) images.  The blow up shows LIRG 3 which contains a prominent red bulge and blue star forming regions.  It also shows COMPGAL 1, an  inclined spiral. } \\end{figure*}  While the AO images are vital for this project, they are difficult to obtain.  Adaptive optics systems require guide stars in order to track and correct the turbulent layers in the atmosphere.  The first generation adaptive optics systems required very bright ($R<13$) natural guide stars (NGS) nearby the science target ($d<40\\arcsec$).  NGS stars are rare, especially at the high galactic latitudes of the extragalactic treasury fields.  GOODS-S contains no suitable NGS guide stars.  Laser guide stars improve the situation.  The Keck sodium LGS AO system produces a bright ($R\\sim10$) artificial star in the sodium layer of the Earth's atmosphere.  The LGS AO system uses the laser spot to correct the high-order aberrations of the wavefront.  LGS AO systems still require a natural guide star to correct for image motion (``tip-tilt'' correction).  These tip-tilt stars, however, can be as faint as $R\\sim18$ making $\\sim 30\\%$ of the GOODS field accessible to the Keck AO facility.  The images used in our study were obtained as part of the Center for Adaptive Optics Treasury Survey \\citep[CATS;][; Larkin et al. in preparation]{Melbourneetal05}.   During Spring 2008, the images will be made available for download to the public through the CATS web interface \\footnote{http://www.ucolick.org/$\\sim$jmel/cats\\_database/cats\\_search.php}. CATS NGS images in the COSMOS \\citep[Cosmological Evolution Survey]{Scoville06} and GEMS \\citep[Galaxy Evolution From Morphology And SEDs, ]{Rix04}  regions are already available for download as a service to the community.  This paper is divided into the following sections.  Section two describes the basic data used in this study, including a discussion of the AO observing strategy, data reduction, and PSF estimation. Section 3 presents the 5-band high resolution images, and a discussion of galaxy morphologies.  Section 4 gives global photometric measurements and derived galaxy properties, such as absolute magnitude, rest-frame color, stellar mass, and IR luminosity.  This section compares the broad properties of the LIRG and ``normal'' galaxy samples.  Section 5 provides photometry and SED's of galaxy sub-components and compares these with single burst stellar population synthesis models.  Section 6 discusses the findings of the paper and puts them into the context of previous results from other groups.  Section 7 summarizes.  We assume a flat cosmology with $H_0= 70$ km/s/Mpc, $\\Omega_m=0.3$, and $\\Omega_{\\lambda}=0.7$. ",
        "conclusions": ""
    },
    "0801/0801.4526_arXiv.txt": {
        "abstract": " ",
        "introduction": "}  The present-day universe is undergoing exponential expansion on a timescale set by today's Hubble parameter $H_0 \\sim 10^{-33} \\, {\\rm eV}$.  One of the biggest challenges in physics is understanding why $H_0$ is so small relative to the energy scales of particle physics. This is the puzzle of the cosmological constant, reviewed in \\cite{Weinberg:1988cp,Weinberg:1996xe,Carroll:2000fy,Weinberg:2000yb, Polchinski:2006gy}.  Many solutions have been proposed \\cite{Nobbenhuis:2004wn}.  An approach which is currently popular is to invoke landscape / anthropic ideas: in a potential with enough minima, at least one should have the observed vacuum energy \\cite{Bousso:2000xa}.  Although this may ultimately prove to be the right explanation, it's still important to explore alternatives.  In fact, the existence of a landscape may require us to explore alternatives.  To explain this point, let us note that in a simple model of flux compactification one has the potential \\cite{Bousso:2000xa,Douglas:2006es} \\[ V \\sim -V_0 + \\sum_{i = 1}^N c_i n_i^2\\,. \\] Here $-V_0$ is a large fixed negative energy density, $N$ is the number of non-trivial cycles on some compactification manifold, the $c_i$ are order one constants, and the integers $n_i$ measure the fluxes through the various cycles.  The number of flux vacua then grows with the vacuum energy roughly as \\be\\label{minima} \\hbox{\\# vacua} \\sim (V + V_0)^{N/2}\\,. \\ee For large $N$ the number of vacua grows extremely rapidly with energy. Now suppose there is a mechanism for (nearly) cancelling a large positive vacuum energy, leaving a small effective cosmological constant.  We do not expect to find such a mechanism without introducing small parameters \\cite{Weinberg:1988cp}.  But in view of the growth (\\ref{minima}), such a mechanism could be statistically favored in the landscape, even though it requires carefully adjusted couplings.  This leads us to consider theories which nearly cancel a large underlying vacuum energy.  The basic mechanism we shall explore is the following.  Suppose there is a potential $V$ for a collection of fields $\\psi$.  We assume that $V(\\psi)$ vanishes somewhere in field space, but (as should be generic) we assume the derivatives of $V$ are non-zero at that point: \\[ V \\big\\vert_{\\psi_0} = 0 \\quad {\\rm but} \\quad \\left.{\\partial V \\over \\partial \\psi_i}\\right\\vert_{\\psi_0} \\not= 0 \\quad \\hbox{\\rm for some $i$.} \\] Then in a neighborhood of $\\psi_0$ we can redefine fields, introducing a new field $\\phi$ such that \\be \\label{IntroPotential} V = \\rv - \\phi\\,. \\ee Here $\\rv$ is a constant representing a large underlying vacuum energy.  Clearly at the level of the potential (\\ref{IntroPotential}), $\\rv$ has no meaning: it is a redundant coupling which can be eliminated by a field redefinition \\be \\label{IntroPseudo} \\phi \\rightarrow \\phi + a \\qquad\\quad \\rv \\rightarrow \\rv + a\\,. \\ee However generically this pseudo-symmetry will be explicitly broken by other terms in the Lagrangian.  By choosing appropriate breaking terms, we will be able to write down models that have de Sitter solutions of arbitrarily small curvature, even in the presence of an underlying Planck-scale vacuum energy.  The necessary breaking terms will involve extremely small couplings.  In practice this allows us to shuffle fine tunings away from the cosmological constant and into other sectors of the theory.  Such tunings may be expected to arise in the landscape, as we emphasized above.  Also by shifting small parameters around, we will be able to construct models in which a small cosmological constant is technically natural (stable under radiative corrections).  Interestingly, these models also exhibit a kind of see-saw mechanism, in that a large underlying cosmological constant gives rise to multiple solutions, some with high curvature, some with low curvature.  An outline of this paper is as follows.  In section \\ref{sect:redundant} we develop the idea of making vacuum energy nearly redundant.  In section \\ref{sect:F(R)} we present a model-building recipe based on fields with non-minimal couplings to the scalar curvature.  In section \\ref{sect:Rphi2} we discuss a particular example in more detail: a model with an $R\\phi^2$ coupling that can be thought of as a generalization of $R + 1/R$ gravity \\cite{Carroll:2003wy}.  In section \\ref{sect:GB} we discuss another example, based on Gauss-Bonnet gravity, which provides a technically natural explanation for dark energy.  In the appendices we show that all models we consider can be made compatible with cosmological and solar system tests of gravity, and give further details on issues of classical and radiative stability.  Readers who are interested in the concrete examples could skip directly to sections \\ref{sect:Rphi2} or \\ref{sect:GB}, which can be read independently of the rest of the paper.  Conventions: our metric signature is $(-+++)$.  Quantities which appear in the underlying Lagrangian include the bare gravitational coupling $\\kappa = 1/\\mp$ and the bare vacuum energy $\\rv$.  Quantities measured at low energies include the effective gravitational coupling $\\ke = 1/\\mpe = (8 \\pi G_N)^{1/2} \\sim 10^{18} \\, {\\rm GeV}$ and the effective vacuum energy $\\rve \\sim (10^{-3} \\, {\\rm eV})^4$. ",
        "conclusions": "}  In this paper we explored the idea that the underlying vacuum energy could be large, of order the (effective) Planck scale, with the observed slow Hubble expansion due to the dynamics of a scalar field with non-minimal couplings.  We showed that the idea could be realized in the $R\\phi^2$ model and its $F(R)$ generalizations.  However all such theories, involving only the scalar curvature, require severe fine-tuning and are unstable with respect to radiative corrections.  We went on to consider the Gauss-Bonnet model in which the scalar field couples to the Euler density.  The Gauss-Bonnet model requires that one combination of parameters be small, namely $J^2 \\ll \\xi\\kappa^6\\rv^3$.  However we argued that for small $\\xi$ this parameter is stable under radiative corrections, and in this sense the Gauss-Bonnet model provides a technically natural explanation for dark energy.  The crucial feature that makes all this possible is the fact that the bare vacuum energy is a redundant coupling, and therefore unobservable, when $\\xi = 0$.  This allows us to shuffle the necessary fine-tunings among various couplings in the Lagrangian.  The way in which tunings are shuffled will play an important role in determining how likely these models are in the landscape.  Also, by shifting some of the tuning to the Gauss-Bonnet term, pseudo-redundancy made it possible to construct radiatively stable models for dark energy.  This leaves many open questions: \\begin{itemize} \\item Can one envisage a cosmological history in which the scalar field naturally evolves to take on its required expectation value? Or does a realistic cosmology require modifying the model in some way?  If one does keep $\\xi$ small throughout cosmological evolution then the quantum diffusion studied in \\cite{Linde,Garriga:2003hj} will play an important role. \\item Some models, in particular Gauss-Bonnet, realize a see-saw mechanism in which a large underlying vacuum energy can drive either a fast or a slow de Sitter expansion.  Can one find interesting cosmological solutions in which the universe spends some time inflating near the high-curvature solution before evolving to low curvature? \\item The models we studied require super-Planckian vevs which are often viewed as problematic \\cite{McAllister:2007bg}.  But are these large vevs really necessary?  Or could the models be modified in some way to eliminate them? \\item Can the necessary non-minimal couplings be realized in a UV-complete theory such as string theory?  Or is there some fundamental obstacle to achieving this? \\item One intriguing feature of the $R\\phi^2$ model is that cosmological stability and solar system constraints put roughly similar requirements on the coupling $\\xi$. Therefore, if $\\xi$ takes a value such that the dark energy equation of state deviates from $-1$ by an observable amount, it is likely that one would also see violations of general relativity in precision solar system measurements. This offers an interesting way to test the model. It would be useful to check if the same property holds in the Gauss-Bonnet model. \\end{itemize}  \\bigskip \\goodbreak \\centerline{\\bf Acknowledgements} \\noindent We are grateful to Matt Kleban, Ali Masoumi, Alberto Nicolis, Massimo Porrati and Iggy Sawicki for valuable discussions. PB, KH, LH and DK are supported by DOE grant DE-FG02-92ER40699 and by a Columbia University Initiatives in Science and Engineering grant.  \\appendix"
    },
    "0801/0801.4656_arXiv.txt": {
        "abstract": "s{ The Alpha Magnetic Spectrometer (AMS) is a particle physics detector designed to measure charged cosmic ray spectra with energies up to the TeV region and with high energy photon detection capability up to few hundred GeV. It will be installed on the International Space Station (ISS) in 2008 and will operate for more than three years. Due to its large acceptance, the flight duration and the state-of-art of particle identification techniques, AMS will have a remarkable sensitivity on antimatter and dark matter searches.\\\\ The addition of different detector systems provide AMS with complementary and redundant electric charge and velocity measurements. The velocity of singly charged particles is expected to be measured with a precision of 0.1$\\%$ and charge separation up to iron is attainable. The AMS capability of measuring a large range of electric charges and accurate velocities, will largely contribute to a better understanding of cosmic ray production, acceleration and propagation mechanisms in the galaxy.} \\vspace{-1.0cm} ",
        "introduction": " ",
        "conclusions": "AMS is a spectrometer designed for antimatter and dark matter searches and for measuring relative abundances of nuclei and isotopes.The velocity of singly charged particles is expected to be measured with a precision of 0.1$\\%$ and charge separation up to iron is attainable. Real data analysis was done with data collected with prototypes of TOF, tracker and RICH in test beams at CERN, in October 2002 and 2003\\\\ \\vspace{-0.5cm}"
    },
    "0801/0801.3595_arXiv.txt": {
        "abstract": "We have obtained multi-epoch, high-resolution spectroscopy of \\Ntarget\\ candidate low-mass stars and brown dwarfs in the young clusters around \\sori\\ and \\lori. We find that \\Nmember\\ targets are cluster members based on their radial velocity, the equivalent width of their Na\\,I\\,8200 lines and the spectral type from their TiO band strength. We have identified \\Nrvvmem\\ new binary stars among the cluster members based on their variable radial velocity and an additional binary from the variation in its line width and shape. Of these, 6 are double-lined spectroscopic binaries (SB2) where the components of the binary are of comparable brightness. The others are single-lined binaries (SB1) in which the companion is faint or the spectra of the stars are blended. There are  3 narrow-lined SB1 binaries in our sample for which the companion is more than 2.5 magnitudes fainter than the primary. This suggests that the mass ratio distribution for the spectroscopic binaries in our sample is broad but that there may be a peak in the distribution near $q=1$. The sample covers the magnitude range $\\Ic = 14$\\,--\\,18.9 (mass $\\approx 0.55 \\-- 0.03\\Msolar$), but all of the binary stars are brighter than $\\Ic=16.6 $ (mass $\\approx 0.12\\Msolar$) and 10 are brighter than $\\Ic = 15.5$ (mass $\\approx 0.23\\Msolar$). There is a significant lack of spectroscopic binaries in our sample at faint magnitudes even when we account for the decrease in sensitivity with increasing magnitude. We can reject the hypothesis that the fraction of spectroscopic binaries is a uniform function of \\Ic\\ magnitude with more than 99 percent confidence. The spectroscopic binary fraction for stars more massive than about $0.1\\Msolar$ ($\\Ic < 16.9$) is $f_{\\rm bright} = 0.095^{+0.012}_{-0.028}$. The 90 percent confidence upper limit to the spectroscopic binary fraction for very low mass (VLM) stars (mass $< 0.1\\Msolar$) and brown dwarfs (BDs) is $f_{\\rm faint} < 7.5$\\,percent. The hypothesis that $f_{\\rm bright}$ and $f_{\\rm faint}$  are equal can be rejected with 90\\,percent confidence.  The average detection probability for our survey is 50\\,percent or more for binaries with separations  up to 0.28\\,au for stars with $\\Ic < 16.9$ and 0.033\\,au for the fainter stars in our sample. We conclude that we have found strong evidence for a change in the fraction of spectroscopic binaries  among young VLM stars and brown dwarfs when compared to more massive stars in the same star-forming region. This implies a difference in the total binary fraction between VLM stars and BDs compared to more massive stars or a difference in the distribution of semi-major axes, or both. ",
        "introduction": "The origins of very low-mass stars (VLMS) and brown dwarfs (BD) are proving difficult to understand, despite being more common than stars of higher mass. Ideas include ejection from protostellar aggregates \\citep{2001AJ....122..432R}, formation within convergent flows generated by turbulence \\citep{2004ApJ...617..559P}, the photo-erosion of pre-stellar cores \\citep{2004A&A...427..299W} or fragmentation within the outer parts of circumstellar discs \\citep{2006A&A...458..817W}.  The frequency and separation distribution of binary systems is an important constraint on the likely formation process. There is strong evidence that the binary properties of the lowest mass stars and brown dwarfs are quite different to those of higher mass objects (see the review of \\citealt{2007prpl.conf..427B}). Resolved imaging of nearby VLMS and BDs in the field show that about 15--20 percent of systems are binaries with separations greater than 1--2\\,au, but that very few have separations greater than 20\\,au (e.g. \\citealt{2003ApJ...587..407C}; \\citealt{2003AJ....126.1526B}). This contrasts with higher mass stars where overall binary frequencies are 30--60 percent, with a much broader spread of possible separations (\\citealt{1992ApJ...396..178F}, \\citealt{1991A&A...248..485D}).  A missing part of the picture is how many VLMS and BDs are binary systems with separations less than about 1\\,au, where the imaging observations cannot reach. \\citet{2005MNRAS.362L..45M} used previously published radial velocity (RV) results to show that an overall binary frequency (at all separations) of 32--45 percent was needed to explain the presence of several binaries detected by RV variations, with most of them at small separations. This high frequency found support from a small RV survey by \\citet{2006MNRAS.372.1879K} who found a frequency of 11\\,--\\,40 percent for separations less than 0.1\\,au in a young association. On the other hand \\citet{2006AJ....132..663B} found few RV variables in their survey of field VLMS/BDs and concluded that the overall binary frequency was 16--36 percent, with very few binaries at separations below 1\\,au  The situation is unresolved and it is quite likely that the apparent discrepancies between these various authors arise from biases within the samples considered; differences in analysis technique and that in terms of absolute numbers, very few short-period VLMS/BD binary systems have yet  been found, thus limiting the statistical precision possible.  In this paper we present the results of an RV survey of a large number of low-mass stars and brown dwarfs in the \\sori\\ and \\lori\\ clusters. These clusters are young and nearby (3--5 Myr, 330--450\\,pc) and contain large populations of VLMS and BDs (\\citealt{2004AN....325..705B}; \\citealt{2004ApJ...610.1064B}; \\citealt{2005MNRAS.356...89K}).  We have observed more than 200 objects using the {\\sc flames} multi-fibre spectrograph on the VLT-Kueyen telescope to measure radial velocities at several epochs and searched for close binary systems at a range of masses. ",
        "conclusions": "We have conducted a large radial velocity survey for spectroscopic binary stars among low mass stars and brown dwarfs in the young clusters around \\lori\\ and  \\sori. We have identified \\Nmember\\ members of these clusters based on their radial velocity and spectral type. Of these, 6 are SB2 binaries and 6 are SB1 binaries. All the spectroscopic binaries we have detected are brighter than \\Ic=16.6  (mass $\\approx$ 0.12\\Msolar). We conclude that the frequency of spectroscopic binaries in these clusters among very low mass stars (mass $<0.1\\Msolar$) and brown dwarfs is significantly lower ($<7.5$ percent) than that for more massive stars ($9\\pm 2$ percent). The change in binary properties with mass that we have discovered may be due to a change in the total binary frequency with mass  or a change in the period distribution of the binaries with mass or both.  The number of SB2 binaries in this sample suggests there may be a peak in the mass ratio distribution for spectroscopic binaries in these clusters near $q\\approx 1$. There is also clear evidence from the properties of the SB1 binaries that the mass ratio distribution for spectroscopic binaries in these clusters is broad, extending down to $q\\approx 0.25$."
    },
    "0801/0801.1842_arXiv.txt": {
        "abstract": "We present mid-infrared spectra of thirty two high redshift ultraluminous infrared galaxies, selected via the stellar photospheric feature at rest-frame 1.6$\\mu$m, and an observed-frame 24$\\mu$m flux of $>$500$\\mu$Jy. Nearly all the sample reside in a redshift range of $\\langle z \\rangle=1.71\\pm0.15$, and have rest-frame 1-1000$\\mu$m luminosities of 10$^{12.9}$ - 10$^{13.8}$L$_{\\odot}$. Most of the spectra exhibit prominent polycyclic aromatic hydrocarbon emission features, and weak silicate absorption, consistent with a starburst origin for the IR emission. Our selection method appears to be a straightforward and efficient way of finding distant, IR-luminous, star-forming galaxies in narrow redshift ranges. There is however evidence that the mid-IR spectra of our sample differ systematically from those of local ULIRGs; our sample have comparable PAH equivalent widths but weaker apparent silicate absorption, and (possibly) enhanced PAH 6.2$\\mu$m/7.7$\\mu$m and 6.2$\\mu$m/11.2$\\mu$m flux ratios. Furthermore, the composite mid-IR spectrum of our sample is almost identical to that of local starbursts with IR luminosities of $10^{10}-10^{11}$L$_{\\odot}$ rather than that of local ULIRGs. These differences are consistent with a reduced dust column, which can plausibly be obtained via some combination of (1) star formation that is extended over spatial scales of 1-4Kpc, and (2) star formation in unusually gas-rich regions. ",
        "introduction": "First discovered in the 1970s \\citep{rie72}, Ultraluminous Infrared Galaxies (ULIRGs, those objects with 1-1000$\\mu$m luminosities in excess of $10^{12}$L$_{\\odot}$) are among the brightest objects in the low redshift Universe, with luminosities rivalling the bolometric output of quasars. Their immense luminosities, implying astrophysical processes operating on extreme scales, have made ULIRGs the object of intense study over the last thirty years, and in the last decade or so a nascent consensus has developed on their nature. Local ULIRGs are probably `composite' objects; most are powered mainly by a starburst, but a significant fraction ($\\sim45\\%$) also contain an IR-luminous AGN. This is suggested from several lines of evidence, including optical/UV spectroscopy \\citep{vei99,far05}, mid-infrared spectroscopy \\citep{gen98,lut98,rig99,arm07,far07b,ima07}, and X-ray observations \\citep{fra03,pta03}. Local ULIRGs are also associated almost exclusively with galaxy mergers \\citep{far01,bus02,vei02}, and may be involved to some degree in triggering QSOs \\citep{san88,tac02,kaw06,zau07}. Reviews of the properties of ULIRGs can be found in \\citet{san96}, and more recently in \\citet{lfs06}.  Initially, studies of ULIRGs were motivated mostly by their luminosities, as they are rare at low redshifts (with less than 50 examples known over the whole sky at $z<0.2$), and so were not thought to play a fundamental role in the broader picture of galaxy evolution. This perception changed abruptly however, when it was realized that ULIRGs are vastly more numerous at high redshift. First hinted at by spectroscopic followup of {\\it Infrared Astronomical Satellite} (IRAS) surveys \\citep{hac87,lon90,sau90}, which showed strong evolution in the ULIRG luminosity function with redshift, this was confirmed by surveys with the {\\it Infrared Space Observatory} (ISO \\citealt{rr97,dol,ver05}) and in the sub-mm \\citep{hug,eal,scot,bor,mor}, which showed that there were several hundred ULIRGs {\\it per square degree} at $z\\gtrsim1$. Later surveys with the {\\it Spitzer} space telescope \\citep{wer04} showed that ULIRGs contributed an increasing fraction of the comoving bolometric energy density with increasing redshift, reaching a level comparable to the contribution from lower luminosity systems at $z>1$ (\\citealt{lef05}, their figure 14). Distant ULIRGs seem to be similar to their low redshift counterparts, in that they appear to be powered by both starburst and AGN activity \\citep{far02b,sma03,sma04,ale05,tak06,val07}, and are probably mostly mergers \\citep{far02a,cha03,bri07}.  The much higher space density of ULIRGs in the high redshift Universe compared to locally implies that they play an increasingly important role in galaxy evolution with increasing redshift; for example, their high star formation rates make distant ULIRGs excellent candidates for being the rapid growth phases of massive ($\\gtrsim2L^{*}$) elliptical galaxies \\citep{scot,roc,swi06}. The sheer number of ULIRGs at high redshift however also raises some serious problems, notably for theories for the formation of large-scale structures. These theories usually invoke `biased' hierarchical buildup, in which overdense halos in the dark matter distribution undergo successive mergers to build haloes of increasing mass, with galaxies forming from the baryonic matter in these haloes \\citep{col,gran,hatt}. In this scenario, ULIRGs, being the likely antecedents of massive systems, should be found in overdense peaks in the underlying DM distribution, and observations seem to bear this out \\citep{bla04,far06a,far06b}. Despite this, the prevalence of ULIRGs at high redshift, and their very high inferred star formation rates, cause major problems for models as they find it difficult to concentrate enough baryons in sufficiently small volumes, and a diverse variety of solutions have been proposed to overcome this \\citep{vank,gran2,bau,bow}.  There is therefore a pressing need for both an efficient census of ULIRGs at $z>1$, and a thorough understanding of how these distant ULIRGs form stars. {\\it Spitzer} is ideal for undertaking such studies; the Infrared Array Camera (IRAC, \\citealt{faz04}) and Multiband Imaging Photometer for Spitzer (MIPS, \\citealt{rie04}) imaging instruments, and the Infrared Spectrograph (IRS, \\citealt{hou04}) all offer dramatic improvements in sensitivity and resolution compared to previous generation facilities. In this paper, we use the IRS to study 32 ULIRGs at $z\\gtrsim1$, selected using IRAC and MIPS photometry. We quantify the efficiency of the selection method and use mid-IR spectral features to set constraints on the nature of their star formation. An overview of the sample selection and a description of the observations is given in \\S\\ref{sectobs}. Results are summarized in \\S\\ref{takeaguess} with discussion and conclusions given in \\S\\ref{discuss} and \\S\\ref{conc} respectively. We assume a spatially flat cosmology with $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega=1$, and $\\Omega_{m}=0.3$. Unless otherwise stated, the term `Infrared (IR) luminosity' refers to the total luminosity integrated over 1-1000$\\mu$m in the rest-frame. ",
        "conclusions": "\\label{conc} We have presented mid-infrared spectra, taken using the IRS onboard Spitzer, of 32 objects selected on (1) an optical band magnitude fainter than $m_{r}=23$, (2) excess flux in the IRAC 4.5$\\mu$m passband compared to the other three IRAC channels, and (3) a 24$\\mu$m flux in excess of 0.5mJy. Our conclusions are:  \\noindent 1 - All of the sample lie within $0.8<z<2.3$, with more than 90\\% lying in a much tighter redshift range, characterized by $\\langle z \\rangle=1.71\\pm0.15$. The rest-frame IR luminosities are consistent with the range $10^{12.9}-10^{13.8}$L$_{\\odot}$. Our selection method is therefore a straightforward and efficient way of finding large samples of distant ULIRGs confined to narrow redshift ranges. It avoids all of the observational challenges inherent in finding distant ULIRGs via their far-IR emission, and produces samples with redshift distributions akin to those of optically selected samples such as Lyman-Break galaxies.  \\noindent 2 - Nearly all of the sample show two or more strong PAH emission features in their mid-IR spectra, and systematically lack strong silicate absorption. The individual spectra closely resemble the IRS spectra of high-redshift sub-mm selected ULIRGs. The star formation rates computed from the PAH luminosities are all extremely high, of order $\\gtrsim1000$M$_{\\odot}$yr$^{-1}$, sufficient to power all of the IR emission. Coupled with the lack of a power law continuum in the IRAC channels, we infer that star formation is likely to be the dominant power source in most of the sample, though it is possible that some also contain a luminous AGN.  \\noindent 3 - Our sample occupy a different region on the PAH 6.2$\\mu$m EW vs silicate strength plane to most local ULIRGs, and there is marginal evidence for enhanced PAH 6.2$\\mu$m/7.7$\\mu$m and 6.2$\\mu$m/11.2$\\mu$m ratios. Furthermore, the composite mid-IR spectrum of our sample resembles that of local starburst galaxies with IR luminosities of $10^{10}-10^{11.5}$L$_{\\odot}$, rather than that of local ULIRGs. Though selection bias likely plays a role, we propose that the most likely reasons for this are that the star formation in distant ULIRGs is extended over scales of a few Kpc rather than the sub-Kpc starbursts seen in local ULIRGs, and/or occurs in unusually gas-rich environments."
    },
    "0801/0801.4183_arXiv.txt": {
        "abstract": "The dynamic spectrum of a radio pulsar is an in-line digital hologram of the ionised interstellar medium. It has previously been demonstrated that such holograms permit image reconstruction, in the sense that one can determine an approximation to the complex electric field values as a function of Doppler-shift and delay, but to date the quality of the reconstructions has been poor. Here we report a substantial improvement in the method which we have achieved by simultaneous optimisation of the thousands of coefficients that describe the electric field.  For our test spectrum of PSR B0834+06 we find that the model provides an accurate representation of the data over the full 63~dB dynamic range of the observations: residual differences between model and data are noise-like. The advent of interstellar holography enables detailed quantitative investigation of the interstellar radio-wave propagation paths for a given pulsar at each epoch of observation; we illustrate this using our test data which show the scattering material to be structured and highly anisotropic. The temporal response of the medium exhibits a scattering tail out to beyond $100\\,{\\rm\\mu s}$ and a pulse arrival time measurement at this frequency and this epoch of observation would be affected by a mean delay of $15\\,{\\rm\\mu s}$ due to multipath propagation in the interstellar medium. ",
        "introduction": "\\label{sec:intro}  With modern instrumentation pulsar dynamic spectra can be recorded at high spectral and temporal resolution yielding a dataset with a large information content from just one observation. For example: if we observe a pulsar for one hour, sampling the spectrum with 1~kHz channels every 10 seconds, over a total bandwidth of 100~MHz, then we have $\\sim4\\times10^7$ independent flux measurements. And if the pulsar is bright and the telescope is large then each of these measurements can have signal-to-noise ratio of order unity implying a total information content of potentially $\\sim$40~Mbit. This information relates primarily to multipath scattering of the radio-waves by the ionised Interstellar Medium (ISM), as it is this scattering which gives rise to the observed interference fringes.  It may be that these paths are determined by random irregularities -- e.g.  caused by turbulence -- in the ISM. In that case any given dynamic spectrum contains information about those random irregularities, and there is no strong motivation to study the individual spectra in detail as they reflect particular realisations of a stochastic process.   But some pulsar dynamic spectra exhibit a high level of order in their fringe patterns --- see Rickett (1991) for an overview. This fact has been emphasised by consideration of the two-dimensional power spectra of the dynamic spectra, wherein power is often seen to be concentrated in parabolic arcs and inverted arclets (Stinebring et al 2001). There is a consensus that this phenomenon arises directly from the geometry of the scattering process (Stinebring et al 2001; Cordes et al 2006;  Walker et al 2004), with waves scattered through an angle $\\vec{\\theta}$ experiencing a Doppler-shift proportional to one component of $\\vec{\\theta}$ and a delay proportional to $\\vec{\\theta}\\cdot\\vec{\\theta}$. In cases where the parabolae are very sharp it has been argued that the scattering is highly anisotropic (Walker et al 2004; Cordes et al 2006). Sharply defined arcs/arclets also require that the scattering arises in a region of small extent along the line-of-sight (Stinebring et al 2001), so it is not distributed turbulence but discrete, localised  structures which are responsible for these features. However, the physical nature of the scattering media remains obscure. Consequently there is now some incentive to explore the information content in individual dynamic spectra, in order to build a detailed picture of the scattering structures. Further motivation for investigating  individual dynamic spectra is provided by studies of the pulsars themselves: if the properties of the scattering screen are accurately known it is possible to use the screen for very high resolution imaging of the pulsar magnetosphere (Wolszczan and Cordes 1987; Gwinn et al 1997; Walker and Stinebring 2005, WS05 henceforth). Precision pulsar timing programs (e.g. Manchester 2007) also provide an incentive for understanding the particular interstellar propagation paths which contribute to individual observations --- if the propagation delays remain uncorrected in the data they constitute a potentially large systematic error in pulse arrival time measurements.  It has previously been demonstrated that  one can iteratively construct a model of the electric field as a function of radio-frequency and time, starting from the observed dynamic spectrum (WS05). The procedure for achieving this is largely equivalent to determining a phase for the electric field in each pixel of the dynamic spectrum, because a noisy estimate of the field amplitude is given directly by forming the square root of the observed intensity. This situation is common to the broad class of problems known as ``phase-retrieval problems'', which are well known in the optics literature (e.g. Fienup 1982). However, the method of solution demonstrated for pulsar dynamic spectra appears to be new; it exploits sparseness of the power distribution in the Fourier domain, and a solution is obtained iteratively by adding discrete new field components in such a way as to minimise the differences between the model dynamic spectrum and the data. Conceptually the process has a strong similarity to the CLEAN algorithm (H\\\"ogbom 1974) which is commonly used in  radio astronomical imaging for deconvolving the synthesised telescope beam from a ``dirty'' image of the sky; CLEAN usually works well if the image is only sparsely populated with emission. The connection between the two algorithms is reinforced when we recognise that determining the electric field from the dynamic spectrum is also equivalent to a deconvolution. The dynamic spectrum as a function of radio-frequency and time is simply $I(\\nu, t) = U^*(\\nu, t)\\,U(\\nu, t)$, where $U$ is the electric field; in the Fourier domain this relationship becomes a convolution $\\widetilde{I} = \\widetilde{U}^*\\otimes\\widetilde{U}$ and so we are deconvolving the Fourier transform of the electric field from its complex conjugate.  Some comments about nomenclature are appropriate at this point. Any process which allows us to reconstruct the electric field which gave rise to a recorded fringe pattern is sensibly termed ``holography'', and the recorded fringe pattern (i.e. the dynamic spectrum in our case) is the hologram. In the case considered here it is ``digital holography'' because the reconstruction is done in software, and ``in-line'' because the object (i.e. the scattering screen in our case) is transparent and sits in the beam which forms the reference wave. In-line holography is sometimes called ``Gabor holography'' because it is the arrangement which was originally conceived by the inventor of holography, Dennis Gabor.  To date the fidelity of electric field reconstructions from pulsar dynamic spectra has been poor. Although the model of WS05 successfully reproduces the general appearance of their test dynamic spectrum there are substantial quantitative differences between the model and the data. These errors are seen when the model ``secondary spectrum'' (i.e the power spectrum of the dynamic spectrum) is compared to the data: the data exhibit a dynamic range of 63~dB, but the {\\it difference\\/} between the WS05 model secondary spectrum and the secondary spectrum of the data has a dynamic range of 47~dB, implying that the reconstruction has correctly captured only the top 16~dB of the data. Here we report modifications to the reconstruction process which have yielded dramatic improvements to the accuracy of the electric field model; on the same test dataset as used by WS05 we find that our improved model captures the full dynamic range of the data. These gains  were achieved by simultaneous optimisation of the thousands of parameters describing the wave interference phenomenon, and by simultaneous optimisation of the hundreds of parameters which describe the intrinsic temporal flux variations of the pulsar.  This paper is organised as follows: in the next section we detail the improvements we have made to the reconstruction algorithm described by WS05; in \\S3 we present the results we have obtained, using the same test data employed by WS05; and in \\S4 we consider what our test data tell us about the interstellar medium, with emphasis on the temporal response of the scattering medium. ",
        "conclusions": "Interstellar holography is a precise new technique which affords detailed insights into the interstellar propagation of radio pulsar signals. The holographic image reconstructed from our test data reveals a complex scattering structure whose physical nature is unknown at present. Holographic imaging can be used to determine the influence of multi-path propagation on pulse arrival time measurements and thus to correct for these propagation delays. Interstellar propagation delays are unpredictable and can potentially make large contributions  to the systematic errors in pulse arrival time measurements. It is therefore prudent to make provision for holography in every instance where an accurate measurement of the unperturbed arrival time is desired."
    },
    "0801/0801.3847_arXiv.txt": {
        "abstract": "We study the degeneracies between dark energy dynamics, dark matter and curvature using a non-parametric and non-perturbative approach. This allows us to examine the knock-on bias induced in the reconstructed dark energy equation of state, $w(z)$, when there is a bias in the cosmic curvature or dark matter content, without relying on any specific parameterisation of $w$. Even assuming perfect Hubble, distance and volume measurements, we show that for $z > 1$, the bias in $w(z)$ is up to two orders of magnitude larger than the corresponding errors in $\\Omega_k$ or $\\Omega_m$. This highlights the importance of obtaining unbiased estimators of all cosmic parameters in the hunt for dark energy dynamics. ",
        "introduction": "}  Since 1998 \\cite{perlm, riess} evidence has been mounting in support of an accelerated expansion of the Universe. Nearly 10 years on, the puzzle of the origin of this acceleration -- dubbed dark energy -- remains one of the most intriguing enigmas in modern day science. Much activity has come from both the theoretical and observational sectors of the physics community in an attempt to pin down its origin. The current drive in dark energy studies is focused on trying to establish its dynamical behaviour as a function of redshift, $w(z)$. While the simplest explanation remains a $\\Lambda$CDM universe with $w=-1$ for all redshift, dynamics in $w(z)$ would provide a window into new physics. Therefore, uncovering the dynamics of dark energy as described by the ratio of its pressure to density, $w(z) = p_{DE}/\\rho_{DE}$, has become the focus of multi-billion dollar proposed experiments using a wide variety of methods, with several planned surveys at redshifts above unity, as high-redshift measurements are useful to constrain dark energy parameters and test for deviation from the concordance $\\Lambda$CDM model (see e.g. \\cite{DETF}). Unfortunately the search for dynamical behaviour in $w$ is a mani-fold problem. The nature of dark energy is elusive: cosmic observations depend not only on dark energy but also on other cosmic parameters such as the cosmic curvature, $\\Omega_k$, and the total matter content, $\\Omega_m$, leading to degeneracies between these and $w(z)$ parameters, an issue which has recently been under intense scrutiny by the community \\cite{m1, CCB,linder, huang}. Kunz \\cite{m1} argues that observations are only sensitive to the full energy-momentum tensor and thus cannot see beyond a combination of the ``dark component\" -- dark matter \\textit{plus} dark energy. The degeneracy between the geometry of the universe and the equation of state of dark energy has also been discussed in light of the well-known result that a cosmological constant in the presence of spatial curvature can mimic a dynamical dark energy \\cite{CCB}.  In this work we review current constraints on cosmic curvature and extend the approach in \\cite{CCB} in two ways. First we include the reconstruction of $w(z)$ that would follow from measurements of the rate of change of cosmic volume with redshift, $dV/dz$. Secondly we discuss the $w(z)$ that would be incorrectly reconstructed from perfect Hubble, distance and volume data if the wrong value of $\\Omega_m$ were used. We assume perfect data for all observations, which allows us to probe fundamental, ``in-principle\" degeneracies that are not due to finite errors and incomplete redshift-coverage. This implies that given a specific bias in a cosmological parameter, the degeneracies will be true no matter what progress is made in improving future cosmic surveys. Furthermore, the key point in our reconstruction of $w(z)$ is that it is performed in a fully non-parametric manner, and so does not rely on the validity of any particular parameterisation of $w(z)$. To illustrate the power of this  non-parametric approach, we compare our method with a standard equation of state parameterisation \\cite{cp,lin}, which cannot fully resolve the above degeneracies.   \\subsection{Degeneracies in Dark Energy Studies}  The success of the inflationary scenario for the early Universe and its standard prediction of flatness to high precision ($\\Omega_k < 10^{-10}$) is perhaps the main reason why curvature has traditionally been left out in analyses of dark energy. However, possible scenarios in which inflation is consistent with non-zero spatial curvature have recently been investigated \\cite{freivogel}. It is also interesting to note that the backreaction of cosmological fluctuations may cause effective non-zero curvature that may yield practical limits on our ability to measure $w(z)$ accurately at $z > 1$ (see e.g. \\cite{backreact}). Since measurements of the Cosmic Microwave Background (CMB) have so far proved consistent with flatness (e.g. \\cite{wmap3}) statistical quantities that measure the necessity of introducing extra parameters (such as Bayesian Evidence or information criteria \\cite{BCK, m2, trotta}) do not favour the inclusion of curvature as a parameter in current analyses \\cite{liddle2}. However this is more a symptom of the inability of current data to constrain an extra parameter than a conclusive case for a flat universe. Further, Bayesian evidence or information criteria do not take into account the power of the biases that may be introduced by falsely neglecting a parameter. We will show below that the biases introduced in neglecting curvature are very significant at $z > 1$.  In general constraints on curvature are very fragile to assumptions about the dark energy since they are primarily derived from distance measurements ($d_L$ or $d_A$) which are completely degenerate with curvature \\cite{weinberg70}. One way to illustrate the degeneracy between curvature and dynamics is as follows. Let us assume that we know all cosmic parameters perfectly other than the curvature $\\Omega_k$ and the dark energy equation of state, $w(z)$. At any redshift, $z_*$, a perfect measurement of $d_L(z_*)$ (or $d_A(z_*)$) allows us to measure a single quantity. If we know $w(z_*)$ then that quantity can be $\\Omega_k$. However, if $w(z)$ is truly a free function, then its value at $z_*$ is completely free and we are left trying to find two numbers from a single observation, which is impossible.  Only when we start to correlate the values of $w(z)$ at different redshifts can we begin to use distance measurements alone to constrain the curvature. The standard way to do this is to assume that $w(z)$ can be compressed onto a finite-dimensional subspace described by $n$ parameters, e.g. \\begin{equation} w(z) =  \\Sigma_j^n w_j z^j \\end{equation} In this case perfect distance measurements at $n+1$ different redshifts will allow a complete solution of the problem and will yield the $w_j$ and $\\Omega_k$. The most extreme version of this is to assume $\\Lambda$CDM, $w(z)=-1$. Within this context it is of course possible to derive very stringent constraints on the curvature. For example, combining the WMAP 3 year data and the SDSS DR5 Luminous Red Galaxy (LRG)  sample leads to $\\Omega_k = -0.003\\pm 0.010$ assuming $w = -1$ \\cite{sdsslrg}. The addition of extra data is crucial since the WMAP data alone provides only the constraint $\\Omega_k = -0.3040 + 0.4067 \\Omega_{\\Lambda}$ \\cite{wmap3}.  It is a highly non-trivial statement that flat $\\Lambda$CDM models provide such a good fit to all the data, but we must be aware that such constraints on the curvature are artificially strong in the sense that adding more dark energy parameters will lead to an almost perfect degeneracy with the curvature. This is visible in Fig. 17 of the WMAP3 \\cite{wmap3} (here Fig.~\\ref{wmap17}) which shows the correlation between a constant $w$ and $\\Omega_k$.   \\begin{figure} \\begin{center} \\includegraphics{f17.eps} \\caption{{\\bf The curvature-dark energy degeneracy} Contours showing the 2D marginalized contours for $w$ and $\\Omega_k$ based on combined data from WMAP3, 2dFGRS, SDSS and supernova surveys. While the slope of the degeneracy differs for this combination of data, the sign of the degeneracy is consistent with the $w_0$ term in Eqs. (\\ref{degenhub}),(\\ref{degend}). Taken from \\cite{wmap3}.\\label{wmap17}} \\end{center} \\end{figure}   Hence we can currently say very little about the true value of the $\\Omega_k$ and the belief that the spatial curvature is small is essentially based on Occam's Razor. Although one could fit distance measurements with any value of $\\Omega_k$, the required $w(z)$ functions would be disfavoured by Bayesian model selection which penalize models with extra parameters that do not significantly improve the fit. We show in detail later the required $w(z)$ functions to do precisely this.  At present a well-defined program for measuring the spatial curvature of the cosmos does not exist. To illustrate this, consider fixing the dark energy to be described by only $n$ parameters. One would hope that given this restriction the resulting constraints on $\\Omega_k$ would be independent of the precise choice of the $n$ parameters, i.e. independent of the parameterisation. Unfortunately a little thought makes it clear that this cannot be true. A parameterisation of $w(z)$ which does not allow mimicry of curvature will provide good, decorrelated constraints on the curvature (which does not mean the corresponding best-fit will be a good fit to the data) while a model which allows perfect mimicry of the dynamics of the curvature (i.e. $(1+z)^2$) will show highly correlated constraints (although for negative $\\Omega_k$ mimicry of distance data is only possible up to a critical redshift as we show later).  \\begin{figure}[htbp!] \\begin{center} \\includegraphics[width=0.5\\textwidth]{ok.eps} \\caption{{\\bf The curvature-dark energy degeneracy} likelihood for $\\Omega_k$ for different parameterisations of dark energy. Assuming a constant model for $w$, allows $\\Omega_k$ to be tightly constrained at $2\\sigma$ to be near 0. However introducing dynamics reduces these constraints significantly. Here $X(z) = \\rho_X(z)/\\rho_X(0)$ is the dark energy density, which \\cite{wang} assume is a free function below some cut-off redshift $z_{cut}$. The value of $X$ at redshifts $z_i = z_{cut}(i/n), i = 1,2..,n$ are treated as $n$ independent model parameters that are estimated from the data. A specific functional form for $X$ is assumed above the cut-off redshift. The likelihoods are given for two such forms of $X(z)$;  namely a power law, $X \\propto (1+z)^{\\alpha}$ for $ z > z_{cut} $, and an exponential function $X \\propto e^{\\alpha z}$. In these figures there are $n = 3$ independent redshifts below a cut-off redshift of $z_{cut} = 1.4$. Taken from \\cite{wang}\\label{wangfig}.} \\end{center} \\end{figure}  \\begin{figure*}[htbp!] \\centering \\mbox{ \\subfigure{\\scalebox{0.8}{\\includegraphics{w0wa1.eps}}} \\quad \\subfigure{\\scalebox{0.8}{\\includegraphics{w0wa2.eps}}}} \\caption{{\\bf Left} - $1\\sigma$ error contours assuming flatness for the dark energy parameters $w_0$ and $w_a$ for the CPL parameterisations. {\\bf Right} - as on the left but with curvature left free and marginalised over. Note how pure distance measurements suffer strongly even with the very limited $w(z)$ parameterisation but that when all the surveys are combined the final error ellipse is essentially unaffected. This is to be expected from Equation (\\ref{curv}) which shows how $\\Omega_k$ can be determined from simultaneous Hubble rate and distance measurements. Figure from Knox {\\em et al.} \\cite{knox}.\\label{knoxfig}} \\end{figure*}  This dilemma is visible in various recent studies attempting to constrain cosmic curvature in the presence of multiple dark energy $w(z)$ parameters \\cite{groupbb, wang, huang}. For some popular parameterisation constraints on $\\Omega_k$ are of order $|\\Omega_k| < 0.05$ at $2\\sigma$. For other parameterisation the constraints evaporate and even $\\Omega_k \\sim 0.2$ cannot be ruled out (see Fig.~\\ref{wangfig}, which is taken from \\cite{wang}).  Alternative cosmic measurements sensitive to curvature include the Integrated Sachs Wolfe (ISW) effect, which is sensitive to the growth of the metric fluctuations $\\Phi$, which is in turn sensitive to both dark energy and curvature. Recent work to investigate the ISW effect as a function of redshift uses the combination of CMB data with information on large scale structure \\cite{isw}. Combining WMAP with such suitable tracers of large scale structure shows that $\\Phi$ has been decreasing with cosmic time \\cite{nolte}, which rules out a large positive curvature which would have predicted the opposite trend.  Another measurement sensitive to the growth function is differential number counts $dN/dz$, e.g. of clusters. This is a potentially sensitive test which, given a constant comoving number of objects, reduces to a test of the rate of change of cosmic volume with redshift, $dV/dz$. We discuss in detail below how perfect measurements of $dV/dz$ allow reconstruction of $w(z)$, and we discuss the resulting errors on dark energy when systematic biases in cosmic parameters are present.  Measurements of the power spectrum from CMB data and from measurement of Baryon Acoustic Oscillations (BAO) provide estimates of the matter content of the universe. While constraints on $\\Omega_m$ are sharpened by combining data from many observations, the best-fit value is often derived on the assumption of flatness \\cite{sdsslrg, perciv}. Unlike the case for cosmic curvature the degeneracy between observables and the matter content is perfect and we show that incorrectly assuming a particular value for $\\Omega_m$ can also mimic deviations from $\\Lambda$CDM.  \\subsection{Future surveys}  We will show in equation (\\ref{curv}) that simultaneous measurements of the Hubble rate $H(z)$, distance $D \\propto d_A,d_L$ and $D'(z)$ allow for a perfect measurement of $\\Omega_k$. BAO allow for the simultaneous measurement of both distance and Hubble rate at the central redshift \\cite{bao}. For a flat universe $D'(z) \\propto 1/H(z)$, but in a curved universe this is not true: the curved geodesics mean that $D'(z)$ contains extra information encoded in $\\Omega_k$.  Measuring $D'(z)$ is in principle possible with future BAO, weak lensing and supernova surveys. In particular, cross-correlation tomography of deep lensing surveys appears to be a very powerful probe of curvature when combined with BAO surveys \\cite{bernstein}, assuming that self-calibration is possible. In principle it should be possible to measure the cosmic curvature to an accuracy of about $\\sigma(\\Omega_k) \\simeq 0.01$ for an all-sky weak lensing and BAO survey out to $z=10$. In principle such a survey would be able to measure distances to about $10^{-4} f_{sky}^{-1/2}$ in redshift bins of width $\\Delta z = 0.1$ out to $z=2.5$ \\cite{bernstein}. This relies critically on the combination of weak lensing and BAO data since constraints from either observations alone are significantly degraded. This can also be seen in Fig. (\\ref{knoxfig}) which shows the error ellipses for the parameters in the CPL parameterisation $w(z) = w_0 + w_a \\left(\\frac{z}{1+z}\\right)$ assuming flatness (left) and leaving $\\Omega_k$ free (right). Note that although individual error ellipses are significantly degraded, the combined data sets have an almost unchanged error ellipse. \\newline  Our work is organized as follows: we illustrate the dependence of the background observables on the cosmological parameters $\\Omega_m, \\Omega_k$ in section \\ref{bg_expand} and discuss obtaining the dark energy equation of state $w$ from observables in section \\ref{obs_eos}. The process of reconstructing $w$ via a non-parametric approach is described in section \\ref{recon_w}. Finally we link this non-parametric approach to other standard approaches to dark energy degeneracies in section \\ref{param} and conclude in section \\ref{conc}. ",
        "conclusions": "}  We have explored the degeneracies between the dark energy equation of state $w(z)$ and cosmic parameters using a non-parametric approach. This means we are able to write down the precise $w(z)$ that will be reconstructed from perfect data if slightly wrong or biased values for the cosmic parameters $\\Omega_k, \\Omega_m$ are assumed. This is complementary to traditional methods which typically use an aggressive compression of the $w(z)$ function onto a couple of parameters (usually $w_0,w_a$) and then study the degeneracy between these and other cosmic parameters. Our approach is superior in one way however: degeneracies between $w(z)$ and some cosmic parameters such as $\\Omega_k$ can appear to be quite weak in the parameterised approach. However, in the case of distance measurements this is completely artificial and due to strong assumptions about the allowed form of $w(z)$ since the degeneracy is perfect if $w(z)$ is allowed to be totally free.  We extend the work of \\cite{CCB} to show the reconstructed $w(z)$ from measurements of volume for both wrongly assumed $\\Omega_k$ and $\\Omega_m$. As with Hubble and distance measurements we show that the errors in $w(z)$ that result from uncertainty in the cosmic parameters are {\\em much} larger than the uncertainty in $\\Omega_k$ or $\\Omega_m$,  especially at large redshifts. We have shown that curvature affects measurements of $H(z)$ and $D(z)$ in complementary ways, with the error at high redshift having opposite signs for an error in $\\Omega_k$. In the case of an $\\Omega_m$ error, Hubble, distance and volume measurements all lead to the same erroneously reconstructed $w(z)$, a manifestation of the dark matter-dark energy degeneracy highlighted in \\cite{m1}.  In this review we have assumed perfect data for Hubble rate, distance and volume at all redshifts. It would be interesting to extend our non-parametric approach to the case of imperfect data which has incomplete redshift coverage and errors on the observables. This is left to future work but will allow contact with the approaches in \\cite{arman,pilar}.  {\\em Acknowledgments -- }we thank Luca Amendola, Chris Blake, Thomas Buchert, Daniel Eisenstein, George Ellis, Martin Kunz,  Roy Maartens, Bob Nichol and David Parkinson for useful comments and insights. MC thanks Andrew Liddle and FTC for support. BB and CC acknowledge support from the NRF and RH acknowledges funding from KAT."
    },
    "0801/0801.3488.txt": {
        "abstract": "%Long baseline interferometric observations of the symbiotic long-period variable %R Aquarii (R Aqr) in the near-infrared atmospheric window at 1.65 $\\mu$m are reported. %Observations were taken with % We report imaging observations of the symbotic long--period variable R % Aquarii (R Aqr)  at near--infrared and radio wavelengths.  The % near--infrared observations were made with the IOTA interferometer in % three narrow band filters centered at 1.51 , % 1.64  and 1.78 $\\mu$m. The 1.64 $\\mu$m filter probes % the pseudo-continuum region, the other two filters probe the water molecule forming % regions in the upper atmosphere of this Mira variable. % The measured fringe visibilities and closure phases are analyzed using % models, varying from simple to more complex, as follows. %%The observations were carried out at least at two photometric phases of these %%Mira variables to study the phase dependency of the complex atmospheres of the %%Mira variables. %(a) A uniform disk model with wavelength dependent sizes fail to fit %the squared visibility  data satisfactorily, %% and we obtain larger sizes at 1.51 $\\mu$m and 1.78 $\\mu$m than at 1.64 $\\mu$m. However, %and in addition, the derived sizes are inconsistent with %our closure phase measurements. %(b) A three-component model that consists %of a central Mira star  surrounded by a water shell and an off--axis %point source provide a good fit to the visibility and %closure phase data, pointing to the need for both water shell and an asymmetric component. %(c) A constrained image reconstruction analysis provides more insight into the nature of the asymmetry, %suggesting that our R Aqr data can be completely understood in terms of a highly non-uniform (clumpy) %water shell around the star.  In addition, we made high resolution VLBA observations of the SiO masers %in the outer parts of the molecular envelope of R Aqr. These images show evidence of turbulence in the %molecular envelope with jet-like features containing velocity gradients. The relation of the systematic rotation-like %velocity seen in SiO masers only at certain epochs and any companions to the AGB star is unclear.  However it is clear %from our near-infrared and radio data that the water shell and SiO maser \"shell\" are both extremely clumpy in appearance, %although the physical connection between these two concentric regions is still elusive.  We report imaging observations of the symbotic long-period Mira variable R Aquarii (R Aqr) at near-infrared and radio wavelengths. The near-infrared observations were made with the IOTA imaging interferometer in three narrow-band filters centered at 1.51, 1.64, and 1.78 $\\mu$m, which sample mainly water, continuum, and water features, respectively.  Our near-infrared fringe visibility and closure phase data are analyzed using three models. (a) A uniform disk model with wavelength-dependent sizes fails to fit the visibility data, and is inconsistent with the closure phase data. (b) A three-component model, comprising a Mira star, water shell, and an off-axis point source, provide a good fit to all data.  (c)  A model generated by a constrained image reconstruction analysis provides more insight, suggesting that the water shell is highly non-uniform, i.e., clumpy.  The VLBA observations of SiO masers in the outer molecular envelope show evidence of turbulence, with jet-like features containing velocity gradients. %Taken together, our near-infrared and radio data suggest that the water shell and SiO maser features are both extremely % clumpy. %, but the physical connection between these two concentric regions remains elusive. ",
        "introduction": "R Aqr is the brightest known symbiotic system \\citep{Whitelock83} consisting of a Mira primary and a white dwarf companion. The companion has not been directly detected yet because of the presence of nebular emissions and circumstellar material around R Aqr itself. The pulsation period of R Aqr is 387 days. The light curves show reduced amplitude lasting for several cycles \\citep{Mattei79}.  Miras with close companion white dwarfs usually are classified as symbiotic systems \\citep{Allen84, Whitelock87, Luthardt92, Bel00}. A few Miras are known to have companions but are not (or are only very mildly) symbiotic systems; this includes {\\it o} Cet = Mira, with a probable WD companion in a multi-century orbit \\citep{Reimers85, Wood04}. %Statistics for the binarity of Miras are otherwise quite uncertain, in part because the expected orbital velocity amplitudes for a close companion, 30 km s$^{-1}$ at 1 AU and 10 km s$^{-1}$ around 5 AU, are very similar to the shock amplitudes of 20-30 km s$^{-1}$ produced by the Mira pulsation itself \\citep{Hinkle84}.  The occasional appearance of bright emission lines in the visible spectrum signifies the %possible presence of a hot companion around R Aqr \\citep{Merrill21,Merrill50}. The observed wide and flat minima of the light curve of R Aqr was interpreted as the signature of %possible a companion around this target \\citep{Merrill56}. Ultraviolet observations from IUE have been interpreted as evidence of a %possible white dwarf companion \\citep{Michal80}. The observed broad infrared excess at 12$\\mu$m \\citep{Kenyon88} could also be interpreted as evidence for the presence of dust heated by a hot companion; alternatively, dust shells with different particle sizes or composition could also explain the observed broad infrared excess at 12$\\mu$m.  An early suggestion for an orbital period of 27 yrs from radial velocity measurements \\citep{Merrill50} was not supported by later observations. The observed decrease in the brightness of the visible and near-infrared light curves of R Aqr during 1975-78 and on two previous occasions was interpreted in terms of an obscuring circumstellar dust clump eclipsing the target \\citep{Willson81}. These authors suggested that R Aqr is an eclipsing binary system with an orbital period of 44 years. Subsequently, \\citet{Whitelock83} found a similar eclipsing signature on 6 symbiotic stars including R Aqr, however, the frequency of occurrence of these decreases in symbiotic stars (3 or 4 out of 6 well studied cases) with such long periods suggests that eclipsing by a secondary body probably is not the likely mechanism causing the observed long term amplitude changes of a few tens of years \\citep{Feast83}. However, if the Mira wind geometry is altered by the presence of the white dwarf, or by two winds colliding, then the \u0093eclipse\u0094 could correspond to observing through the unmodified Mira wind on the side away from the white dwarf. Alternatively, episodic dust emission observed around Mira and AGB stars \\citep{Danchi95, Bester96} could explain these observed brightness modulations. A nova explosion on the surface of the white dwarf companion followed by dust condensation could also explain the observed long term amplitude changes. R Aqr could very well be a potential candidate of becoming a slow nova such as RR Tel which is a symbiotic system with a Mira primary and infrared excess.  In this paper we use the word ``asymmetry'' to mean that part of the 2-dimensional brightness distribution which cannot be made symmetric with respect to a reflection through a point. Thus, for example, an elliptical uniform disc or an equal-brightness binary system are both symmetric, but a binary with unequal brightness or a star with an off-centered bright/dark spot is asymmetric \\citep{Ragland06}.  Departures from circular symmetry are known in AGB stars from various high angular resolution observations \\citep{Karovska91, Wilson92, Haniff92, Richichi95, Ragland96, Weigelt96, Karovska97, Tuthill97, Lattanzi97, Wittkowski98, Tuthill99, Tuthill00, Hofmann00, Thompson02, Monnier04, Weiner06, Ragland06}. However, no consensus exists as to the mechanism that would cause such departures from apparent circular symmetry \\citep{Ragland06}.  Observations of dust shells around AGB stars in the infrared \\citep{Danchi94, Lopez97, Tuthill00, Monnier00, Weiner06} and SiO maser shells at radio/millimeter wavelengths \\citep{Diamond94, Green95, Diamond03, Cotton04, Soria04, Cotton06} also show asymmetries. Again, the connection between apparent surface features (or companions) and the morphology of the dust or SiO shells has not been established.  Recent studies of astrophysical jets around a few AGB stars \\citep{Kellogg01, Imai02, Sahai03, Sokoloski03, Brocksopp04} %and  pre-planetary nebula \\citep{Bujarrabal02, Sahai03b} from radio, X-ray or Hubble Space Telescope (HST) observations suggest that those stars showing substantial asymmetry may all have a low mass stellar companion accreting mass from the AGB primary.  The R Aqr system has complex nebular structures at arcsecond and arcminute scales. In addition, X-ray, UV, visible and radio observations show astrophysical jets \\citep{Kellogg01, Hollis91, Solf85, Hollis85}. These large scale structures have no influence on our high angular resolution multi-wavelength observations of the central region (size $\\sim 0.4$ arcsec) of this complex system reported in this paper.  Recently, we reported the results from the first phase of our program in which we surveyed 56 evolved giants including 35 Mira stars, 18 SR variables and 3 Irr variables looking for asymmetry in their brightness profiles (\\citep{Ragland06}, hereafter Paper I). We found that about 29\\% of the AGB stars show non-zero (or non-$\\pi$) closure phase signifying the asymmetric nature of the brightness profiles. R Aqr is the lone symbiotic star with positive asymmetry detected by our AGB survey. In this paper we report followup observations on R Aqr to characterize the observed asymmetry in this symbiotic system. A second objective is to detect a possible water shell around R Aqr and its possible role in the asymmetry detection. The reasons for combining our near-infrared observations with SiO maser observations are the following: (1) probably the water shell emissions detectable in the near-infrared and the SiO emissions detectable in the radio arise from the same circumstellar region \\citep{Cotton04, Cotton06}, (2) the SiO maser observations give us some ideas about the distortions from spherical symmetry we might expect in the near-infrared, and also some velocity indications. %Our multi-wavelength closure-phase measurements would support or rule-out clumpy molecular %clouds as the cause for observed asymmetry. ",
        "conclusions": "In order to characterize the observed asymmetry reported in our earlier work \\citep{Ragland06}, we carried out follow-up observations of selected Mira stars at near-infrared wavelengths. In this article, we presented observations of the nearest known symbiotic Mira, R Aqr, from the IOTA imaging interferometer and from the VLBA.  The reconstructed near-infrared images of R Aqr are reported for the first time at three wavelengths namely, 1.51, 1.64 and 1.78 $\\mu$m. The near-infrared images suggest a clumpy molecular shell around the central Mira star. The observed data could also be modeled by a three component model consisting of a symmetric central star surrounded by a water shell with a radius of about 2.25R$_*$, and an off-axis compact feature at about 2R$_*$ at a position angle of 131$^{\\rm o}$ (an alternate model exists wherein the off-axis feature is at about R$_*$ at a position angle of -94$^{\\rm o}$). %The compact model feature has a flux ratio (feature/total) of $\\sim$ 2\\% in the 1.51-1.78$\\mu$m wavelength range. %The observed relatively large flux ratio constrains the model feature to be an accretion disk around a white dwarf %companion (the white-dwarf itself is likely too faint to be detectable in our observations) or a bright ``hot %spot'' on the surface of AGB star. However, the companion feature solution is inconsistent with the orbital parameters reported in the literature \\citep{Hollis97}. We conclude that our observations are best explained with a clumpy, extended $H_2O$ circumstellar envelope. %We have also constrained the mass of the water shell and the accretion disk from our measurements.  The ratio of the size of water shell to the size of the central star is 2.25. The ratio of the size of SiO shell to that of water shell is 1.25. Thus, the masers appear at the outer edge of the molecular envelope as reported for other Mira stars in the literature.  The SiO maser emission around R Aqr shows large scale turbulence in the molecular envelope with prominent jet--like features towards the north--east. These features are roughly parallel to the large scale jet from the more distant companion star discussed by \\citet{Hollis97}. However, similar features seen in 2001 \\citep{Cotton04,Hollis01} have a very different orientation from either the jet or the inferred orbit of the companion so any connection between the SiO features and the companion or its jet is doubtful. The magnetic fields appear to be oriented parallel to the features but due to the range of orientations seen in these features, their magnetic fields are not part of a persistent, large scale dipole field. %The multi-epoch SiO observations (\\citet{Hollis01}, \\citet{Cotton04}, \\citet{Cotton06}, Figure \\ref{geometry}) % show systematic rotation-like velocity at certain epochs, but not always, thus supporting the hypothesis of a % close white dwarf companion orbiting around the Mira primary.  Further investigation and observations will be required to understand the details of this complex and interesting system. %Indeed, we are reaching a stage where interferometry can help to elucidate the details of Roche lobe overflow and %the complexities of phenomena associated with the accretion disks around compact objects such as white dwarf stars."
    },
    "0801/0801.4901_arXiv.txt": {
        "abstract": "The recent interferometric study of Achernar, leading to the conclusion that its geometrical oblateness cannot be explained in the Roche approximation, has stirred substantial interest in the community, in view of its potential impact in many fields of stellar astrophysics. It is the purpose of this paper to reinterpret the interferometric observations with a fast rotating, gravity darkened central star surrounded by a small equatorial disk, whose presence is consistent with contemporaneous spectroscopic data. We find that we can only fit the available data assuming a critically rotating central star. We identified two different disk models that simultaneously fit the spectroscopic, polarimetric, and interferometric observational constraints: a tenuous disk in hydrostatic equilibrium (i.e., with small scaleheight) and a smaller, scaleheight enhanced disk. We believe that these relatively small disks correspond to the transition region between the photosphere and the circumstellar environment, and that they are probably perturbed by some photospheric mechanism. The study of this interface between photosphere and circumstellar disk for near-critical rotators is crucial to our understanding of the Be phenomenon, and the mass and angular momentum loss of stars in general. This work shows that it is nowadays possible to directly study this transition region from simultaneous multi-technique observations. ",
        "introduction": "}  Interferometry greatly increased our knowledge of the Be Star Achernar ($\\alpha$ Eri, HD10144). \\citet[][hereafter D03]{dom03} observed Achernar during fall 2002 with the ESO-VLTI/VINCI instrument \\citep{ker03} in the $K$ band and found that the star is highly oblate. By converting the individual visibilities into equivalent uniform disk angular diameters they derived a ratio between the maximum and minor elongation axis $a/b=1.56\\pm0.05$. Later, \\citet{ker06} detected a tenuous polar wind and very recently \\citet{ker07} found that Achernar has a main-sequence, lower mass companion.  The available interferometric data, together with the abundance of data from the literature, forms a body of information that is probably unrivaled by that of any other Be star. There is, however, much that is still unknown about the star and its circumstellar environment. As pointed out by D03, the determination of the actual stellar oblateness from the interferometric data is not a simple task, since there are so many unknowns involved. In their original work D03, based on simultaneous spectroscopic data, assumed that there was no circumstellar material at the time of the observations, and that the observed shape was purely photospheric in origin. Their modeling of the observations with a rotationally deformed, gravity darkened star, assuming the Roche approximation of uniform rotation and centrally condensed mass, led to the conclusion that the stellar flattening required to explain the observations exceeds that of a critical rigid rotator .  Recent theoretical developments in the theory of differentially rotating stars (e.g., \\citealt{jac04}) show that they can have equatorial radii larger than the Roche limit, which is physically compatible with the interferometric results. However, \\citet{vin06} demonstrated the existence of a tiny, yet non negligible emission in H$\\alpha$ at the time of the VLTI observations, an indication that there was some circumstellar material around Achernar, probably associated with a small rotating disk. Furthermore, \\citet{vin06} showed, from temporal analysis of profile variations of the line He I $\\lambda6678$ \\AA, that orbiting gas clouds are a very frequent feature of Achernar, even during its quiescent phase.  The presence of this circumstellar material raises the question of how much it can contribute to the observed visibilities. In this paper we investigate how the presence of a small disk around Achernar may alter to the interferometric signal, and the effects this may have on determination of fundamental parameters such as stellar rotation rate and stellar flattening. ",
        "conclusions": "}  The presence of a small disk around Achernar is a realistic possibility to explain the interferometric observations. As shown by D03 (see also Figure~\\ref{fig:sed}), Roche models without a disk cannot account for the observed aspect ratio.  The combined constraints imposed by the interferometry, spectroscopy and the adopted upper limit for the polarization have allowed us to impose very narrow limits on the model parameters, within our assumption of a rigidly rotating star in the Roche approximation surrounded by a small disk. The main result is that the star must be rotating very close to critical, since all of our models with $\\Omega/\\Omega_\\mathrm{crit} < 0.992$ that successfully reproduced both the visibilities and the $H\\alpha$ EW (e.g. model 5)  had too large polarization levels, for the reasons discussed above.  We have found two critical models that reproduce equally well the observations: a large and dense disk in hydrostatic equilibrium (i.e., geometrically thin; model 2) and a small and more tenuous disk with enhanced scaleheight (i.e., geometrically thick; model 3). Because the lack of contemporaneous polarization measurements, the model parameters shown in Table~\\ref{tab:models} have a degree of uncertainty that is difficult to estimate. This stresses the importance of having simultaneous multi-technique data.  The current paradigm for Be star disks is that of a geometrically thin viscous Keplerian disk in vertical hydrostatic equilibrium. This paradigm has been corroborated by several recent studies \\citep[e.g.][]{car06b,mei07}.  When a star is close to critical rotation, the stellar photospheric material near the equator, which is weakly bound to the star because of the strong centrifugal force, can eventually escape the star provided it is given an extra energy by some other mechanism (e.g. stellar pulsation, interaction in a binary system, photospheric activity). This material is likely to have, initially, a very complicated density and velocity distribution. As the gas diffuses outwards as a result of viscosity, the density and velocity distributions will tend to relax and become in hydrostatic equilibrium.  Thus, it is reasonable to assume that there is a \\emph{transition region} between the photosphere (which is characterized by large densities of the order of $10^{-10}\\;\\rm g\\; cm^3$) and the disk itself, which has densities at least 10 times lower. Many unknowns exist concerning the size and physical properties (e.g. density, temperature, velocity field) of this smooth transition region.  With the available data we were able to narrow down substantially the range of possible values for Achernar's model parameters, but we cannot distinguish yet between model 2 and 3, i.e., between a model in hydrostatic equilibrium and a model with enhanced scaleheights that might result from the perturbation of the gas by some photospheric mechanism.  As shown in Figure~\\ref{fig:sed}, high-precision spectropolarimetry might add invaluable information, since models 2 and 3 have different values for the size of Balmer and Pashen jumps in polarization. Also, with simultaneous spectro-interferometry and spectropolarimetry one can spatially resolve the velocity field, thus establishing whether the material is in equilibrium or not. The observed residual emission profile (Figure~\\ref{fig:sed}) is indicative of velocities fields more complex than the simple Keplerian rotation that was assumed in this work.  In any case, one important consequence of this work is that we have shown that with current observing techniques and state-of-the-art modeling it is already possible to study the properties of the very inner layers of the disk in Achernar and possibly other nearby Be stars. Future simultaneous spectropolarimetry, spectro-interferometry and photometry will allow us to study the properties of this region and determine, for instance, the size of the transition region and the point at which the disk becomes in vertical hydrostatic equilibrium."
    },
    "0801/0801.1653_arXiv.txt": {
        "abstract": "Properties and structure of neutron stars are determined by the equation of state (EOS) of neutron-rich stellar matter. While the collective flow and particle production in relativistic heavy-ion collisions have constrained tightly the EOS of symmetric nuclear matter up to about five times the normal nuclear matter density, the more recent experimental data on isospin-diffusion and isoscaling in heavy-ion collisions at intermediate energies have constrained considerably the density dependence of the nuclear symmetry energy at subsaturation densities. Although there are still many uncertainties and challenges to pin down completely the EOS of neutron-rich nuclear matter, the heavy-ion reaction experiments in terrestrial laboratories have limited the EOS of neutron-rich nuclear matter in a range much narrower than that spanned by various EOSs currently used in astrophysical studies in the literature. These nuclear physics constraints could thus provide more reliable information about properties of neutron stars. Within well established formalisms using the nuclear constrained EOSs we study the momenta of inertia of neutron stars. We put the special emphasis on the component A of the extremely relativistic double neutron star system PSR J0737-3039. Its moment of inertia is found to be between $1.30$ and $1.63$ $(\\times10^{45}g$ $cm^2)$. Moreover, the transition density at the crust-core boundary is shown to be in the narrow range of $\\rho_t=[0.091-0.093](fm^{-3})$. ",
        "introduction": "Neutron stars exhibit a large array of extreme characteristics. Their properties and structure are determined by the equation of state (EOS) of neutron-rich stellar matter at densities up to an order of magnitude higher than those found in ordinary nuclei, see, e.g.~\\citet{Weber:1999a} and \\citet{Lattimer:2004pg}. Therefore, the detailed knowledge about the EOS of neutron-rich nuclear matter over a wide range of densities is necessary for the study of neutron stars. For isospin asymmetric nuclear matter, various theoretical studies have shown that the energy per nucleon can be well approximated by \\begin{eqnarray} E(\\rho ,\\delta )=E(\\rho ,\\delta =0)+E_{\\mathrm{sym}}(\\rho )\\delta ^{2}+O(\\delta ^{4}), \\end{eqnarray} in terms of the baryon density $\\rho =\\rho _{n}+\\rho _{p}$, the isospin asymmetry $\\delta=(\\rho _{n}-\\rho _{p})/(\\rho _{n}+\\rho _{p})$, the energy per nucleon in symmetric nuclear matter $ E(\\rho ,\\delta =0)$, and the bulk nuclear symmetry energy $E_{\\mathrm{sym}}(\\rho )$. Here we report the results of a study on the moment of inertia and the core-crust transition density of neutron stars within well established formalisms in the literature using several EOSs constrained by the latest terrestrial heavy-ion reaction experiments.  Presently, besides the possibilities of phase transitions into various non-nucleonic states the behavior of nuclear matter under extreme densities, pressures and/or isospin-asymmetry is still highly uncertain and relies upon, often, rather controversial theoretical predictions. This circumstance introduces corresponding uncertainties in the EOS of neutron-rich nuclear matter and thus limits our ability to understand many key issues in astrophysics. While astrophysical observations can also limit the EOS of neutron-rich nuclear matter, terrestrial laboratories experiments provide complementary information and have their unique advantages. In this regard, it is especially interesting to mention that the collective flow and particle production in relativistic heavy-ion collisions have constrained the EOS of symmetric nuclear matter $E(\\rho ,\\delta =0)$ up to about five times the normal nuclear matter density to a narrow range~\\citep{Danielewicz:2002pu}. However, there are still many challenges and uncertainties in pinning down precisely the EOS of neutron-rich nuclear matter. One of the major remaining uncertainties is the density dependence of the nuclear symmetry energy $E_{sym}(\\rho)$, see e.g. Refs. ~\\citep{Lattimer:2004pg,Steiner:2004fi,Chen:2007} for recent reviews. To constrain the density dependence of the symmetry energy, many terrestrial nuclear experiments have been carried out recently or planned. Depending on the techniques used, some experiments are more useful for exploring the symmetry energy at low densities while others are more effective at high densities. For instance, heavy-ion reactions, especially those involving radioactive beams, provide a unique means to probe the $E_{sym}(\\rho)$ over a broad density range ~\\citep{Li:1997px,Li:2001a,Baran:2004ih,Chen:2007}. In fact, some significant progress has been made very recently by studying the isospin diffusion~\\citep{Shi:2003np,Tsang:2004,Chen:2005a,Li:2005jy} and isoscaling~\\citep{Tsang:2001,Shetty:2007} in heavy-ion reactions at intermediate energies. The analysis of these phenomena based on transport theories of heavy-ion reactions and thermodynamical models of nuclear multifragmentation has limited the $E_{sym}(\\rho)$ in a range much narrower than that spanned by various forms of the $E_{sym}(\\rho)$ currently used in astrophysical studies in the literature. Moreover, the lower bound of the $E_{sym}(\\rho)$ extracted from the heavy-ion reactions is consistent with the RMF prediction using the FSUGold interaction that can reproduce not only saturation properties of nuclear matter but also structure properties and giant resonances of many finite nuclei~\\citep{Piek07}.  It is also well known that the sizes of neutron skins in heavy nuclei are sensitive to the symmetry energy at subsaturation densities, see, e.g., Refs.~\\citep{Brown00,Horowitz:2000xj,Horowitz:2002mb,Die03,Fur02,Steiner:2004fi,ToddRutel:2005fa,Steiner:2005rd,Chen05b}. However, available data of neutron-skin thickness obtained using hadronic probes are not accurate enough to constrain significantly the symmetry energy. Interestingly, the parity radius experiment (PREX) at the Jefferson Laboratory aiming to measure the neutron radius in $^{208}Pb$ via parity violating electron scattering (Jefferson Laboratory Experiment E-00-003) \\citep{Horowitz:2001} hopefully will provide much more precise data and thus constrain the symmetry energy at low densities more tightly in the near future. On the other hand, at supranormal densities, a number of potential probes of the symmetry energy have been proposed~\\citep{Li:1997rc,Li:2000bj,Li:2002qx,Li:1997px,Chen:2007}. Moreover, several experiments to probe the high density behavior of the symmetry energy with high energy radioactive beams have been planned at the CSR/Lanzhou, FAIR/GSI, RIKEN and the NSCL/MSU.  While the EOS of neutron-rich nuclear matter has not been completely determined yet, it is still very interesting to examine astrophysical implications of the EOS constrained by the latest terrestrial laboratory experiments mentioned above. Global properties of spherically symmetric static (non-rotating) neutron stars have been studied extensively over many years, for recent reviews, see, e.g., Refs.~\\citep{Lattimer:2000kb,Lattimer:2004pg,Prakash:2001rx,Yakovlev:2004iq,Heiselberg:2000dn,Heiselberg:1999mq,Steiner:2004fi}. However, properties of (rapidly) rotating neutron stars have been investigated to lesser extent. Models of (rapidly) rotating neutron stars have been constructed only by several research groups with various degree of approximation \\citep{Hartle:1967he,Hartle:1968si,Friedman:1986tx,Bombaci:2000rc,1990ApJ...355..241L,1989MNRAS.237..355K,1994ApJ...424..823C,Stergioulas:1994ea,Stergioulas:1997ja,1993A&A...278..421B,1998PhRvD..58j4020B,Weber:1999a,2002A&A...381L..49A} (see \\citet{Stergioulas:2003yp} for a review). In a recent work~\\citep{KLW2} we have reported predictions on the gravitational masses, radii, maximal rotational (Kepler) frequencies, and thermal properties of (rapidly) rotating neutron stars. In this work, using the nuclear constrained EOSs we calculate the momenta of inertia for both spherically-symmetric (static) and (rapidly) rotating neutron stars using well established formalisms in the literature. Such studies are important and timely as they are related to the astrophysical observations in the near future. In particular, the moment of inertia of pulsar $A$ in the extremely relativistic neutron star binary PSR J0737-3039 \\citep{Burgay2003} may be determined in a few years through detailed measurements of the periastron advance \\citep{BBH2005}. ",
        "conclusions": "\\begin{figure}[h!] \\centering \\includegraphics[totalheight=2.8in]{f3.eps} \\vspace{5mm} \\caption{(Color online) Mass-radius relation for static and maximally rotating neutron stars. Solid lines correspond to static and broken lines to maximally rotating stellar models. Taken form \\citep{KLW2}.} \\label{fig.3} \\end{figure}  We calculate the neutron star moment of inertia applying several nucleonic EOSs (see Section 2) considering both slow- and rapid-rotation regimes. In Fig.~\\ref{fig.3} we show stellar sequences computed with the $RNS$ code for spherical and maximally rotating models. As seen in the figure, rapid rotation alters significantly the mass-radius relation of rapidly rotating stars (with respect to static configurations). Generally, for a given EOS it increases the maximum possible mass by $\\sim 15\\%$, while reducing/increasing the polar/circumferential radius by several kilometers, leading to an overall oblate shape of the rotating star. The degree to which the neutron star properties are impacted by rapid rotation depends on the details of the EOS: it is greater for models from stiffer EOS which produce less centrally condensed and gravitational bound neutron stars \\citep{1984Natur.312..255F}. In view of these considerations, one should expect similar changes in the moment of inertia of rapidly rotating neutron stars (with respect to static models). We address these and other implications next.  \\subsection{Slow rotation} \\begin{figure}[b!] \\centering \\includegraphics[totalheight=3.1in]{f4.eps} \\vspace{5mm} \\caption{(Color online) Total moment of inertia of neutron stars estimated with Eq.~(\\ref{eq.14}).} \\label{fig.4} \\end{figure} If the rotational frequency is much smaller than the Kepler frequency, the deviations from spherical symmetry are negligible and the moment of inertia can be calculated applying the slow-rotation approximation discussed briefly in Section 3. For this case \\citet{Lattimer:2005} showed that the moment of inertia can be very well approximated by Eq.~(\\ref{eq.14}). In Fig.~\\ref{fig.4} we display the moment of inertia as a function of stellar mass for slowly rotating neutron stars as computed with the empirical relation (\\ref{eq.14}). As shown in Fig.~\\ref{fig.3}, above $\\sim 1.0M_{\\sun}$ the neutron star radius remains approximately constant before reaching the maximum mass supported by a given EOS. The moment of inertia ($I\\sim MR^2$) thus increases almost linearly with stellar mass for all models. Right before the maximum mass is achieved, the neutron star radius starts to decrease (Fig.~\\ref{fig.3}), which causes the sharp drop in the moment of inertia observed in Fig.~\\ref{fig.4}. Since $I$ is proportional to the mass and the square of the radius, it is more sensitive to the density dependence of the nuclear symmetry energy, which determines the neutron star radius. Here we recall that the $x=-1$ EOS has much stiffer symmetry energy (with respect to the one of the $x=0$ EOS), which results in neutron star models with larger radii and, in turn, momenta of inertia. For instance, for a ``canonical'' neutron star ($M=1.4M_{\\sun}$), the difference in the moment of inertia is more than $30\\%$ with the $x=0$ and the $x=-1$ EOSs. In Fig.~\\ref{fig.5} we take another view of the moment of inertia where $I$ is scaled by $M^{3/2}$ as a function of the stellar mass (after \\citep{Lattimer:2005}).  \\begin{figure}[!t] \\centering \\includegraphics[totalheight=3.2in]{f5.eps} \\vspace{5mm} \\caption{(Color online) The moment of inertia scaled by $M^{3/2}$ as a function of the stellar mass $M$. The shaded band illustrates a 10\\% error of hypothetical $I/M^{3/2}$ measurement of 50 $km^2$ $M_{\\sun}^{-1/2}$. The error bar shows the specific case in which the mass is $1.34M_{\\sun}$ (after \\citep{Lattimer:2005}).} \\label{fig.5} \\end{figure}  The discovery of the extremely relativistic binary pulsar PSR J0737-3039A,B provides an unprecedented opportunity to test General Relativity and physics of pulsars \\citep{Burgay2003}. \\citet{Lattimer:2005} estimated that the moment of inertia of the A component of the system should be measurable with an accuracy of about 10\\%. Given that the masses of both stars are already accurately determined by observations, a measurement of the moment of inertia of even one neutron star could have enormous importance for the neutron star physics \\citep{Lattimer:2005}. (The significance of such a measurement is illustrated in Fig.~\\ref{fig.5}. As pointed by \\citet{Lattimer:2005}, it is clear that very few EOSs would survive these constraints.) Thus, theoretical predictions of the moment of inertia are very timely. Calculations of the moment of inertia of pulsar A ($M_A=1.338M_{\\sun}$, $\\nu_A=44.05Hz$) have been reported by \\citet{Morrison:2004} and \\citet{BBH2005}. \\begin{table}[t!] \\caption{Numerical results for PSR J0737-3039A ($M_A=1.338M_{\\sun}$, $\\nu_A=44.05Hz$).} \\begin{center} \\begin{tabular}{lcccc}\\label{tab.3} EOS &  $\\epsilon_c(\\times 10^{14}g$ $cm^{-3})$ & $R_{eq}(km)$ & $I(\\times 10^{45}g$ $cm^2)$ & $I^{LS}(\\times 10^{45}g$ $cm^2)$\\\\ \\hline\\hline MDI(x=-1)    & 7.04 & 13.75 (13.64) & 1.63 & 1.67\\\\ DBHF+Bonn B  & 7.34 & 12.56 (12.47) & 1.57 & 1.43\\\\ MDI(x=0)     & 9.85 & 12.00 (11.90) & 1.30 & 1.34\\\\ APR          & 9.58 & 11.60 (11.52) & 1.25 & 1.26\\\\ \\hline \\end{tabular} \\end{center} {\\small The first column identifies the equation of state. The remaining columns exhibit the following quantities: central mass-energy density, equatorial radius (the numbers in the parenthesis are the radii of the spherical models; the deviations from sphericity due to rotation are $\\sim 1\\%$), total moment of inertia, total moment of inertia $I^{LS}$ as computed with Eq.~(\\ref{eq.14}).} \\end{table} In Table~\\ref{tab.3} we show the moment of inertia and (other selected quantities) of PSR J0737-3039A computed with the $RNS$ code using the EOSs employed in this study. Our results with the APR EOS are in very good agreement with those by \\citet{Morrison:2004} ($I^{APR}=1.24\\times 10^{45}g$ $cm^2$) and \\citet{BBH2005} ($I^{APR}=1.23\\times 10^{45}g$ $cm^2$). In the last column of Table~\\ref{tab.3} we also include results computed with the empirical relation (Eq.~(\\ref{eq.14})). From a comparison with the results from the exact numerical calculation we conclude that Eq.~(\\ref{eq.14}) is an excellent approximation for the moment of inertia of slowly-rotating neutron stars. (The average uncertainty of Eq.~(\\ref{eq.14}) is $\\sim 2\\%$, except for the DBHF+BonnB EOS for which it is $\\sim 8\\%$.) Our results (with the MDI EOS) allowed us to constrain the moment of inertia of pulsar A to be in the range $I=(1.30-1.63)\\times 10^{45}(g$ $cm^2)$.  \\subsection{Rapid rotation}  \\begin{figure}[!t] \\centering \\includegraphics[totalheight=2.8in]{f6.eps} \\vspace{5mm} \\caption{(Color online) Total moment of inertia for Keplerian models. The neutron star sequences are computed with the $RNS$ code.} \\label{fig.6} \\end{figure}  \\begin{figure}[!t] \\centering \\includegraphics[totalheight=3.4in]{f7.eps} \\vspace{5mm} \\caption{(Color online) Total moment of inertia as a function of stellar mass for models rotating at 716Hz (upper frame) and 1122 Hz (lower frame).} \\label{fig.7} \\end{figure}  In this subsection we turn our attention to the moment of inertia of rapidly rotating neutron stars. In Fig.~\\ref{fig.6} we show the moment of inertia as a function of stellar mass for neutron star models spinning at the mass-shedding (Kepler) frequency. The numerical calculation is performed with the $RNS$ code. We observe that the momenta of inertia of rapidly rotating neutron stars are significantly larger than those of slowly rotating models (for a fixed mass). This is easily understood in terms of the increased (equatorial) radius (Fig.~\\ref{fig.3}).  We also compute the momenta of inertia of neutron stars rotating at 716 \\citep{Hessels:2006ze} and 1122Hz \\citep{Kaaret:2006gr} which are the rotational frequencies of the fastest pulsars of today. The numerical results are presented in Fig.~\\ref{fig.7}. As demonstrated by \\citet{Bejger:2006hn} and most recently by \\citet{KLW2}, the range of the allowed masses supported by a given EOS for rapidly rotating neutron stars becomes narrower than the one for static configurations. The effect becomes stronger with increasing frequency and depends upon the EOS. This is also illustrated in Fig.~\\ref{fig.7}, particularly in the lower panel. Additionally, the moment of inertia shows increase with rotational frequency at a rate dependent upon the details of the EOS. This is best seen in Fig.~\\ref{fig.8} where we display the moment of inertia as a function of the rotational frequency for stellar models with a fixed mass ($M=1.4M_{\\sun}$). \\begin{figure}[!t] \\centering \\includegraphics[totalheight=2.8in]{f8.eps} \\vspace{5mm} \\caption{(Color online) Total moment of inertia as a function of rotational frequency for stellar models with mass $M=1.4M_{\\sun}$.}\\label{fig.8} \\end{figure} The neutron star sequences shown in Fig.~\\ref{fig.8} are terminated at the mass-shedding frequency. At the lowest frequencies the moment of inertia remains roughly constant for all EOSs (which justifies the application of the slow-rotation approximation and Eq.~(\\ref{eq.14})). As the stellar models approach the Kepler frequency, the moment of inertia exhibits a sharp rise . This is attributed to the large increase of the circumferential radius as the star approaches the ``mass-shedding point''. As pointed by \\citet{1984Natur.312..255F}, properties of rapidly rotating neutron stars display greater deviations from those of spherically symmetric (static) stars for models computed with stiffer EOSs. This is because such models are less centrally condensed and gravitationally bound. This also explains why the momenta of inertia of rapidly rotating neuron star configurations from the $x=-1$ EOS show the greatest deviation from those of static models.  \\subsection{Fractional moment of inertia of the neutron star crust}  As it was discussed extensively by \\citet{Lattimer:2000kb} (and others), the neutron star crust thickness might be measurable from observations of pulsar glitches, the occasional disrupts of the otherwise extremely regular pulsation from magnetized, rotating neutron stars. The canonical model of \\citet{Link:1999} suggests that glitches are due to the angular momentum transfer from superfluid neutrons to normal matter in the neutron star crust, the region of the star containing nuclei and nucleons that have dripped out of nuclei. This region is bound by the neutron drip density at which nuclei merge into uniform nucleonic matter. \\citet{Link:1999} concluded from the observations of the Vela pulsar that at least $1.4\\%$ of the total moment of inertia resides in the crust of the Vela pulsar. For slowly rotating neutron stars, applying several realistic hadronic EOSs that permit maximum masses of at least $\\sim 1.6M_{\\sun}$  \\citet{Lattimer:2000kb} found that the fractional moment of inertia, $\\Delta I/I$, can be expressed approximately as \\begin{equation}\\label{eq.15} {\\Delta I\\over I}\\simeq{28\\pi P_t R^3\\over 3Mc^2}{(1-1.67\\beta-0.6\\beta^2)\\over\\beta}\\left[1+{2P_t(1+5\\beta-14\\beta^2) \\over \\rho_tm_bc^2\\beta^2}\\right]^{-1}\\ \\end{equation} In the above equation $\\Delta I$ is the moment of inertia of the neutron star crust, $I$ is the total moment of inertia, $\\beta = GM/Rc^2$ is the compactness parameter, $m_b$ is the average nucleon mass, $\\rho_t$ is the transition density at the crust-core boundary, and $P_t$ is the transition pressure. The determination of the transition density itself is a very complicated problem. Different approaches often give quite different results. Similar to determining the critical density for the spinodal decomposition for the liquid-gas phase transition in nuclear matter, for uniform $npe$-matter, \\citet{Lattimer:2000kb} and more recently \\citet{Kubis:2007} have evaluated the crust transition density by investigating when the incompressibility of $npe$-matter becomes negative, i.e \\begin{equation}\\label{eq.16} K_{\\mu}=\\rho^2{d^2E_0\\over d\\rho^2}+2\\rho{dE_0\\over d\\rho}+\\delta^2\\left[\\rho^2{d^2E_{sym}\\over d\\rho^2}+2\\rho{dE_{sym}\\over d\\rho}-2E_{sym}^{-1}\\left(\\rho{dE_{sym}\\over d\\rho}\\right)^2\\right]<0 \\end{equation} (see Fig.~\\ref{fig.9}) where $E_0(\\rho)$ is the EOS of symmetric nuclear matter, $E_{sym}$ is the nuclear symmetry energy, and $\\delta=(\\rho_n-\\rho_p)/(\\rho_n+\\rho_p)$ is the asymmetry parameter. Using this approach and the MDI interaction, \\citet{Kubis:2007} found the transition density of $0.119, 0.092, 0.095$ and $0.160 fm^{-3}$ for the $x$ parameter of $1, 0, -1 $ and $-2$, respectively. Similarly, we have calculated the transition densities and pressures for the EOSs employed in this work. Our results are summarized in Table~\\ref{tab.4}. We find good agreement between our results and those by \\citet{Kubis:2007} with the MDI interaction. It is interesting to notice that the transition densities predicted by all EOSs are in the same density range explored by heavy-ion reactions at intermediate energies. The MDI interaction with $x=0$ and $x=-1$ constrained by the available data on isospin diffusion in heavy-ion reaction at intermediate energies thus limits the transition density rather tightly in the range of $\\rho_t=[0.091-0.093](fm^{-3})$.  \\begin{figure}[t!] \\centering \\includegraphics[totalheight=2.4in]{f9.eps} \\vspace{5mm} \\caption{(Color online) The incompressibility, $K_{\\mu}$, as a function of baryon density $\\rho$.}\\label{fig.9} \\end{figure}  \\begin{table}[b!] \\caption{Transition densities and pressures for the EOSs used in this paper.} \\begin{center} \\begin{tabular}{lcccc}\\label{tab.4} EOS &  MDI(x=0) & MDI(x=-1) & APR & DBHF+Bonn B\\\\ \\hline\\hline $\\rho_t(fm^{-3})$     & 0.091 (0.095) & 0.093 (0.092) & 0.087 & 0.100\\\\ $P_t(MeV$ $fm^{-3})$  & 0.645 & 0.982 & 0.513 & 0.393\\\\ \\hline \\end{tabular} \\end{center} {\\small The first row identifies the equation of state. The remaining rows exhibit the following quantities: transition density, transition pressure. The numbers in the parenthesis are the transition densities calculated by \\citet{Kubis:2007}.} \\end{table}  \\begin{figure}[t!] \\centering \\includegraphics[totalheight=2.5in]{f10.eps} \\vspace{5mm} \\caption{(Color online) The fractional moment of inertia of the neutron star crust as a function of the neutron star mass (left panel) and radius (right panel) estimated with Eq.~(\\ref{eq.15}). The constraint from the glitches of the Vela pulsar is also shown.}\\label{fig.10} \\end{figure}  The fractional momenta of inertia $\\Delta I/I$ of the neutron star crusts are shown in Fig.~\\ref{fig.10} as computed through Eq.~(\\ref{eq.15}) with the parameters listed in Table~\\ref{tab.4}. It is seen that the condition $\\Delta I/I>0.014$ extracted from studying the glitches of the Vela pulsar does put a strict lower limit on the radius for a given EOS. It also limits the maximum mass to be less than about $2M_{\\sun}$ for all of the EOSs considered. Similar to the total momenta of inertia the ratio $\\Delta I/I$ changes more sensitively with the radius as the EOS is varied."
    },
    "0801/0801.3656_arXiv.txt": {
        "abstract": "\\vskip .5in  \\centerline{\\bf Abstract}  \\vskip .2in  We review the uncertainties in the spin-independent and -dependent elastic scattering cross sections of supersymmetric dark matter particles on protons and neutrons. We propagate the uncertainties in quark masses and hadronic matrix elements that are related to the $\\pi$-nucleon $\\sigma$ term and the spin content of the nucleon. By far the largest single uncertainty is that in spin-independent scattering induced by our ignorance of the $\\langle N | {\\bar q} q | N \\rangle$ matrix elements linked to the $\\pi$-nucleon $\\sigma$ term, which affects the ratio of cross sections on proton and neutron targets as well as their absolute values. This uncertainty is already impacting the interpretations of experimental searches for cold dark matter. {\\it We plead for an experimental campaign to determine better the $\\pi$-nucleon $\\sigma$ term.}  Uncertainties in the spin content of the proton affect significantly, but less strongly, the calculation of rates used in indirect searches. ",
        "introduction": "Introduction}  The most convincing way to confirm the existence and nature of dark matter particles would be to observe directly their scattering on nuclei in low-background underground experiments. The sensitivities of these experiments are currently improving rapidly, and beginning to cut into the parameter space of plausible supersymmetric scenarios~\\cite{osusy07}.  In order to evaluate accurately the impacts of these experiments, it is important to understand and minimize the hadronic uncertainties in the elastic scattering matrix elements for any given supersymmetric model. The rates for elastic scattering also control the rates for the capture of dark matter particles by celestial bodies such the Sun or Earth. There are good prospects for increasing significantly the sensitivities of experiments looking indirectly for astrophysical dark matter via the products of their annihilations in such bodies, adding to the motivations for understanding and reducing their uncertainties.  Beyond the interpretation of upper limits on dark matter scattering may lie the interpretation of any eventual detection of a signal and the task of identifying the nature of the dark matter particle. In principle, there are four observables that could contribute to such an analysis, namely the spin-independent and -dependent cross sections on protons and neutrons, respectively. Part of the strategy for identifying the nature of any detected dark matter signal would be the comparison of the measured rates for scattering on different targets, with the comparison of spin-independent and -dependent scattering rates playing a particularly important role as a diagnostic tool~\\cite{Savage:2004fn,Bertone:2007xj}. As we see later, currently there are considerable uncertainties also in such comparisons, related principally to uncertainties in the hadronic matrix elements of higher-dimensional effective interactions.  In this paper, we consider only hadronic uncertainties in the elastic scattering rates.  There are also potentially important uncertainties related to the supersymmetric model itself, namely how accurately one can estimate the coefficient of a given higher-dimensional effective interaction in a given model, and also in the astrophysical density of dark matter particles. Reducing the model uncertainty would require, \\eg, a complete calculation of radiative corrections to the effective scattering operator, which lies beyond the scope of this work. As for the local density of cold dark matter, it is usually taken to be 0.3~GeV/cm$^{3}$, but lower values have occasionally been advocated.  The hadronic uncertainties we consider are listed in \\reftab{params}. They include those in the quark masses, expressed as $m_{\\rm{d,c,b,t}}$ and the ratios $\\mumd$ and $\\msmd$, those in the matrix elements $\\langle N | {\\bar q} q | N \\rangle$, which are related to the change in the nucleon mass due to non-zero quark masses, denoted by $\\sigma_0$, and therefore to the $\\pi$-nucleon $\\sigma$ term, $\\SigmapiN$ as discussed later, and the axial-current matrix elements $\\langle N | {\\bar q} \\gamma_\\mu \\gamma_5 q | N \\rangle$, which are related to the quantities $\\Deltaps, \\athree$ and $\\aeight$, as also discussed later. We find that the uncertainties in the elastic scattering cross section induced by the uncertainties in the quark masses, apart from the top quark, are negligible. However, cross section uncertainties induced by the uncertainties in the matrix elements $\\langle N | {\\bar q} q | N \\rangle$ and $\\langle N | {\\bar q} \\gamma_\\mu \\gamma_5 q | N \\rangle$ are very important, as we discuss below. In particular, the uncertainties induced by our ignorance of $\\SigmapiN$ are particularly important.  \\begin{table} \\addtolength{\\tabcolsep}{1em} \\begin{tabular}{lc@{$\\,\\pm\\,$}cc} \\hline \\hline \\mumd & 0.553 & 0.043    &~\\cite{Leutwyler:1996qg} \\\\ \\md   & 5     & 2 MeV    &~\\cite{Yao:2006px} \\\\ \\msmd & 18.9  & 0.8      &~\\cite{Leutwyler:1996qg} \\\\ \\mc   & 1.25  & 0.09 GeV &~\\cite{Yao:2006px} \\\\ \\mb   & 4.20  & 0.07 GeV &~\\cite{Yao:2006px} \\\\ \\mt   & 171.4 & 2.1 GeV  &~\\cite{Heinson:2006yq} \\\\ \\hline $\\sigma_0$ & 36 & 7 MeV &~\\cite{Borasoy:1996bx} \\\\ \\SigmapiN  & 64 & 8 MeV &~\\cite{Pavan:2001wz,Ellis:2005mb} \\\\ \\hline \\athree  & 1.2695 & 0.0029 &~\\cite{Yao:2006px} \\\\ \\aeight  & 0.585  & 0.025  &~\\cite{Goto:1999by,Leader:2002ni} \\\\ \\Deltaps & -0.09  & 0.03   &~\\cite{Alekseev:2007vi} \\\\ \\hline \\hline \\end{tabular} \\caption[Parameters]{ Hadronic parameters used to determine neutralino-nucleon scattering cross-sections, with estimates of their experimental uncertainties. } \\label{tab:params} \\end{table}    For illustration, we work within the minimal supersymmetric extension of the Standard Model (MSSM) with conserved $R$ parity, and assume that the astrophysical cold dark matter is provided by the lightest neutralino $\\chi$~\\cite{EHNOS}. We further assume a constrained MSSM (CMSSM) framework, in which the soft supersymmetry-breaking parameters $m_{1/2}, m_0$ and $A_0$ are assumed to be universal at the GUT input scale, and restrict our attention to scenarios with the Higgs mixing parameter $\\mu > 0$ and specific values of the ratio of supersymmetric Higgs v.e.v.s $\\tanb$~\\cite{cmssm}.  We illustrate our observations by studies of some specific CMSSM benchmark scenarios~\\cite{bench}, and also by surveys along strips in the $(m_{1/2}, m_0)$ plane for $\\tan \\beta = 10, 50$ along which ${\\tilde \\tau} - \\chi$ coannihilation maintains the relic neutralino density within the range favoured by WMAP and other experiments~\\cite{cmssmwmap}.  We find that the spin-independent cross section may vary by almost an order of magnitude for 48~MeV $< \\SigmapiN < 80$~MeV, the $\\pm 2$-$\\sigma$ range according to the uncertainties in \\reftab{params}. This uncertainty is already impacting the interpretations of experimental searches for cold dark matter. Propagating the $\\pm 2$-$\\sigma$ uncertainties in $\\Deltaps$, the next most important parameter, we find a variation by a factor $\\sim 2$ in the spin-dependent cross section. Since the spin-independent cross section may now be on the verge of detectability in certain models, and the uncertainty in the cross section is far greater, {\\it we appeal for a greater, dedicated effort to reduce the experimental uncertainty in the $\\pi$-nucleon $\\sigma$ term $\\SigmapiN$}. This quantity is not just an object of curiosity for those interested in the structure of the nucleon and non-perturbative strong-interaction effects: it may also be key to understanding new physics beyond the Standard Model. ",
        "conclusions": "Summary and Discussion}  We have analyzed in this paper the principal hadronic uncertainties in the spin-independent (SI) and spin-dependent (SD) cross sections for supersymmetric relic scattering on protons and neutrons, using three benchmark points and two coannihilation strips as illustrations. We have found that the principal hadronic uncertainty in the SD cross sections is due to our lack of knowledge of the $\\pi$-nucleon $\\sigma$ term. In comparison, uncertainties in quark masses and their ratios are much less important. In the case of the SD cross sections, the dominant uncertainty is due to our ignorance of the strange-quark contribution to the nucleon spin, though this uncertainty is relatively less important than that induced in the SI cross sections by the $\\pi$-nucleon $\\sigma$ term.  This uncertainty in the $\\pi$-nucleon $\\sigma$ term clouds very significantly the interpretation of searches for (and eventually measurements of) dark matter scattering on nuclei, preventing precise answers to the key questions: How do present unsuccessful searches constrain the supersymmetric model parameter space? How accurately could a possible future measurement be used to refine the model parameters? This hadronic uncertainty is much larger than that generates by uncertainties in supersymmetric model calculations of the effective LSP-quark interactions, and also much larger than the astrophysical uncertainty in the local cold dark matter density.  One of the great hopes in supersymmetric phenomenology is that one will eventually be able to use measurements at accelerators such as the LHC and/or a linear $e^+ e^-$ collider to calculate the relic LSP density in the Universe, and the rates for dark matter scattering. The hadronic uncertainty in the latter that is induced by our ignorance of the $\\pi$-nucleon $\\sigma$ term limits severely the prospects for completing the second part of this programme. Specifically, this uncertainty is much larger than the uncertainty in calculating the relic LSP density that could be expected from LHC measurements in at least one benchmark model.  {\\it We therefore plead for an experimental campaign to determine better the $\\pi$-nucleon $\\sigma$ term.} This quantity is certainly interesting and important in its own right and as a measure of the importance of strange quarks in the nucleon. However, as argued in this paper, it is potentially also a key ingredient in the effort to understand one of the most important aspects of possible new physics beyond the Standard Model."
    },
    "0801/0801.4289_arXiv.txt": {
        "abstract": "{} {We undertake an optical and ultraviolet spectroscopic analysis of a sample of 20 Galactic B0 -- B5 supergiants of luminosity classes Ia, Ib, Iab, and II. Fundamental stellar parameters are obtained from optical diagnostics and a critical comparison of the model predictions to observed UV spectral features is made.} {Fundamental parameters (e.g., \\teff, \\loglstar, mass-loss rates and CNO abundances) are derived for individual stars using CMFGEN, a nLTE, line-blanketed model atmosphere code. The impact of these newly derived parameters on the Galactic B supergiant \\teff\\ scale, mass discrepancy, and wind-momentum luminosity relation is examined.} {The B supergiant temperature scale derived here shows a reduction of about 1\\,000 -- 3\\,000 K compared to previous results using unblanketed codes. Mass-loss rate estimates are in good agreement with predicted theoretical values, and all of the 20 B0~ -- B5 supergiants analysed show evidence of CNO processing. A mass discrepancy still exists between spectroscopic and evolutionary masses, with the largest discrepancy occurring at \\logllsolar\\ $\\sim$ 5.4. The observed WLR values calculated for B0 -- B0.7 supergiants are higher than predicted values, whereas the reverse is true for B1 -- B5 supergiants. This means that the discrepancy between observed and theoretical values cannot be resolved by adopting clumped (i.e., lower) mass-loss rates as for O stars. The most surprising result is that, although CMFGEN succeeds in reproducing the optical stellar spectrum accurately, it fails to precisely reproduce key UV diagnostics, such as the \\ion{N}{v} and \\ion{C}{iv} P Cygni profiles. This problem arises because the models are not ionised enough and fail to reproduce the full extent of the observed absorption trough of the P Cygni profiles. } {Newly-derived fundamental parameters for early B supergiants are in good agreement with similar work in the field. The most significant discovery, however, is the failure of CMFGEN to predict the correct ionisation fraction for some ions. Such findings add further support to revising the current standard model of massive star winds, as our understanding of these winds is incomplete without a precise knowledge of the ionisation structure and distribution of clumping in the wind. } ",
        "introduction": "The study of luminous, massive stars is fundamental to improving our understanding of galactic evolution, since the radiatively driven winds of these stars have a tremendous impact on their host galaxies. This huge input of mechanical energy is responsible for creating H II regions, making a significant contribution to the integrated light of starburst galaxies and providing star formation diagnostics at both low and high redshifts. They substantially enrich the local ISM with the products of nucleosynthesis via their stellar winds and supernovae explosions and are a possible source of gamma ray bursts. It is therefore imperative to obtain accurate fundamental parameters for luminous massive stars since they contribute to many currently active areas of astrophysical research.\\\\  \\noindent However there are still some uncertainties regarding the post-main sequence evolution of massive stars since their evolution is controlled by variable mass loss from the star as well as rotation, binarity and convective processes, the latter leading to surface enrichment as the products of nuclear burning are brought to the surface. Until recently, stellar evolution models failed to predict the correct amount of CNO processing in massive stars. However, new evolutionary tracks that account for the effects of rotation \\citep{meynet2001} show better agreement between predicted and observed amounts of CNO enrichment in massive stars. A far greater problem in stellar astrophysics is the determination of {\\em accurate} observed mass loss rates. Recent research (e.g., \\citealt{fullerton2006,prinja2005,massa2003,puls2006,repolust2004}), has shown that current OB star mass loss rates might be over-estimated by at least a factor of 10. Such large uncertainties in the mass loss rates of massive stars suggest that our understanding of their winds is incomplete. It is now widely accepted that the winds of both O and B stars are highly structured and clumped and therefore our assumptions that they are smooth and homogeneous are invalid. Evidence to support this claim has come from hydrodynamical, time-dependent simulations of stellar winds (e.g. \\citealt{owocki1988, owocki2002}), the latter of which proposed the idea that instabilites in the line-driving of the wind can produce small-scale, stochastic structure in the wind. Further evidence for the inhomogeneity of stellar winds comes in the form of various observational studies \\citep[e.g.,][]{puls2006,bouret2005,massa2003, prinja2002,bianchi2002}. Time-series analyses of both Balmer and metal spectral lines \\citep[e.g.,][]{prinja2004} in OB stars highlight clear, periodic patterns of variability that correspond to the evolution of structure in the wind. \\cite{ puls2006} demonstrated that the discrepancy between values of \\mdothalpha\\ and \\mdotradio\\ implies the presence of different amounts of clumping at the base of and further out in the wind.  \\cite{massa2003} showed that for a sample of O stars in the LMC, the empirical ionisation fractions derived were several orders of magnitude lower than expected, indicating a lack of dominant ions in the wind. Similar results were found by Prinja et al. (2005) for early B supergiants. More recently, \\cite{fullerton2006} demonstrated that the ionisation fraction of \\ion{P}{v}, which is dominant over a given range in \\teff\\ in the O star spectral range, never approaches a value of unity. They subsequently showed that a reduction in mass loss rate of at least a factor of 100 is required to resolve the situation. We intend to re-address the issue of the ionisation structure of early B supergiants, following on from \\cite{prinja2005}, in a forthcoming paper (Searle et al., 2007b, in preparation; hereafter Paper II). Such drastic reductions in OB star mass loss rates would have severe consequences for the post-main-sequence evolution of these stars; in particular it would affect the numbers of Wolf-Rayet stars produced and the ratio of neutron stars to black holes produced in the  final stages of massive star evolution. \\\\   \\noindent Early type B supergiants are particularly important since they are the most numerous massive luminous stars and are ideal candidates for extra-galactic distance indicators, essential for calibrating the Wind-Momentum-Luminosity Relation (WLR) \\citep[e.g.,][]{kudritzki1999}.  Research into this WLR calibration has highlighted a spectral type dependence for Galactic OBA type stars \\cite[e.g.,][]{kudritzki1999,repolust2004,markova2004}, whilst others have explored the effect of metallicity on the WLR by studying OB stars in the metal-poor environment of the Magellanic Clouds \\cite[e.g.,][]{kudritzki2000,trundle2004,evans2004a}. Accurately derived mass loss rates are essential in calibrating the WLR, yet discrepancies still exist between observed mass loss rates obtained from different wavelength regions (i.e.  optical, UV or IR). Furthermore acknowledged discrepancies of up to 30~\\% have been found between observational and theoretically predicted mass loss rates \\citep{vink2000}. Vink has remarked that these discrepancies for the mass loss rates of B stars can be attributed to systematic errors in the methods employed to derive the observed values. Good agreement was found between observed and predicted mass loss rates for O stars in \\cite{vink2000}. Additionally, \\cite{puls2006} recently derived values of both \\mdothalpha\\ and \\mdotradio, highlighting a discrepancy of roughly a factor of two between both values. \\\\  \\noindent The layout of this paper is as follows. \\S \\ref{obs} introduces the sample of 20 Galactic B0 -- B5 supergiants upon which this analysis is based. \\S \\ref{param} describes the methodology used in deriving fundamental parameters for this sample. The results are presented in \\S \\ref{results} and a critical examination of the {\\sc CMFGEN} model fit to the UV spectra of the 20 B supergiants is made in \\S \\ref{uv}. Finally the conclusions are given in \\S \\ref{conc}. ",
        "conclusions": "\\label{conc}  A quantitative study of the optical and UV properties of B0~--~B5 Ia, Iab, Ib/II supergiants has been carried out, using the nLTE, line-blanketed stellar atmosphere code of \\cite{hillier1998}. A revised B supergiant \\teff\\ scale (derived using a stellar atmosphere code that includes the effects of line blanketing) has been presented, giving a range of 14.5\\,000~K $\\le \\teff \\le$ 30\\,000~K for these stars. This scale shows a drop of up to 10\\,000K from B0 Ia/b to B1 Iab and a difference of up to 2\\,500 K between Ia and Ib stars. It also shows that on average the effect of including line blanketing in the model produces a modest reduction of up to 1\\,000 K for B0~--~B0.7 and B3~--~B5 supergiants, whereas a larger reduction of up to 3\\,000 K is seen for B1 -- B2 supergiants (see Table \\ref{teffsp}). The 20 Galactic B supergiants also displayed a range of 2.1 $\\le$ log $g \\le$ 3.4 in surface gravity. These results, together with those of \\cite{crowther2006}, have been used to construct a new set of averaged fundamental parameters for B0~--~B5 supergiants, according to spectral type. Mass loss rates derived from \\ha\\ proved B supergiant winds to be generally weaker than those of O supergiants (as expected since they are lower-luminosity objects) with \\mdot\\ ranging from $-7.22 \\le \\logmdot \\le -5.30$. All 20 B supergiants also shown signs of CNO processing, with the largest nitrogen enrichments being seen for B1-B2 supergiants. Evidence for a mass discrepancy is found between estimates of \\mspec\\ and \\mevol, with the largest differences peaking at a value of \\logllsolar\\ $\\sim$ 5.4. \\\\  \\noindent A Wind-Momentum--Luminosity relation has also been derived for our sample, which is lower in value for B1 -- B5 supergiants than that predicted by \\cite{vink2000}, but greater than predicted values for B0 -- B0.7 supergiants. For this reason it is not possible to reconcile this difference in observed and theoretical WLRs over the whole B supergiant spectral range by adopting clumped \\mdot\\ as is the case for O stars. A severe problem exists in the form of the optical-UV discrepancy, where the model fails to reproduce some of the P Cygni profiles accurately. This highlights a failure in the model to generate enough high-velocity absorption to succeed in reproducing the observed P Cygni profile and more crucially highlights that the models are not predicting the correct ionisation structure. Given that B supergiants, along with other massive stars, have their peak flux in the UV, it is imperative that this discrepancy is resolved if we are to have confidence that fundamental parameters derived by this method are a true representation of the star's properties. Furthermore it underlines the incompleteness of our current understanding of the physics of massive star winds and the necessity to review the standard model. A more thorough analysis of the ionisation structure of early B supergiant winds will be presented in Paper II. \\\\"
    },
    "0801/0801.3460_arXiv.txt": {
        "abstract": "We present radial velocities for 2045 stars in the Small Magellanic Cloud (SMC), obtained from the 2dF survey by \\citet{eh04}.  The great majority of these stars are of OBA type, tracing the dynamics of the young stellar population.  Dividing the sample into {\\it ad hoc} `bar' and `wing' samples (north and south, respectively, of the line: $\\delta = -77^{\\circ}50^{\\prime} + [4\\alpha]^{\\prime}$, where $\\alpha$ is in minutes of time) we find that the velocities in the SMC bar show a gradient of $26.3 \\pm 1.6$~\\kms~deg$^{-1}$ at a position angle of $126 \\pm 4^\\circ$.  The derived gradient in the bar is robust to the adopted line of demarcation between the two samples.  The largest redshifts are found in the SMC wing, in which the velocity distribution appears distinct from that in the bar, most likely a consequence of the interaction between the Magellanic Clouds that is predicted to have occurred 0.2~Gyr ago.  The mean velocity for all stars in the sample is $+172.0 \\pm 0.2$~\\kms (redshifted by $\\sim$20~\\kms\\/ when compared to published results for older populations), with a velocity dispersion of 30~\\kms. ",
        "introduction": "\\label{intro}  By studying the kinematic and chemical characteristics of the Milky Way and nearby galaxies, past interactions and the history of mass assembly can be explored.  The Magellanic Clouds are of particular interest in this context as their dynamical evolution and star-formation histories appear to be closely entwined \\citep[e.g.,][]{yn03}.  For instance, the Magellanic Stream has generally been thought to arise from tidal disruption by the Milky Way \\citep{gn96,put98,ckg06}, although new proper motion results \\citep{kma06} suggest that the Clouds may be on their first passage about the Milky Way, calling into question the origins of the Stream \\citep{b07}.  From observations of the neutral-hydrogen gas, \\citet{mf84} advanced the concept that the SMC is comprised of two components, with a `Mini-Magellanic Cloud' being effectively torn off by a past interaction with the LMC, leaving the remnant of the SMC.  These components had been identified previously by \\citet{mn81}, who also noted additional features at larger and smaller velocities. Observations of the H$\\;$\\1 gas at better spatial resolution \\citep{ss97} have revealed a much more complex structure, comprised of hundreds of expanding shells, presumably driven by stellar-wind outflows and supernovae, which led the authors to describe the SMC as `distinctly frothy and filamentary'.  Subsequent observations of three large supershells were reported by \\citet{stan99}, with the most dominant of these, 304A, potentially accounting for the twin components seen previously in the main body of the SMC.  To date, the most comprehensive survey of stellar kinematics in the SMC is the study of 2046 red giants by \\citet{hz06}, who concluded that this old population is supported by its velocity dispersion (rather than rotation).  Previous, more limited, studies have focussed on luminous supergiants \\citep{am77,am79,maa87}, Cepheids \\citep{dsm88}, and on the older stellar populations, e.g., carbon stars \\citep{dh97,kdi00} and red-clump stars \\citep{hch93}.  The proximity and significantly sub-solar metallicity of the Clouds have also made them popular targets for spectroscopic studies of stellar evolution.  This was the primary motivation for our own spectroscopic survey of the young, blue population of the Small Magellanic Cloud \\citep[][hereinafter Paper~I]{eh04}, which resulted in spectral classifications for 4161 stars, dominated by early (OBA) types.  Here we present radial velocities for 2045 objects from that sample, covering a relatively wide spatial area (including the `wing' region) and providing kinematic information for a component of the SMC population which has previously gone largely unexplored.   \\begin{figure*} \\begin{center} \\includegraphics[scale=0.65]{SMCfig.pdf} \\caption[]{The approximate boundaries of the 2dF SMC survey, superimposed on a $V$-band image by M.~Bessel (left; contours show H$\\;$\\1\\ column densities), and on an H$\\;$\\1 velocity map (right; small circles show fields observed by \\citealt{hz06}).  The dense optical region constitutes the SMC bar, with the Wing being the lower-density eastern region.  Adapted, with permission, from \\citet{stan04}.} \\label{fig:stan} \\end{center} \\end{figure*} ",
        "conclusions": "We have presented stellar radial velocities measured for 2045 stars from our 2dF survey of the SMC.  We find a local velocity dispersion that is approximately gaussian, with $\\sigma \\simeq 30$, similar to recent results for red giants ($\\sigma \\simeq 27.5$~\\kms; \\citealt{hz06}), but the global velocity range is significantly larger, with the SMC wing receding some $\\sim$20~\\kms\\ faster than the bar.  We find a continuous distribution of velocities across the SMC, but those in the wing appear to have a distinct distribution compared to those in the bar (left-hand panel of Fig.~\\ref{fig:rotcurv}), presumably a consequence of the last close-interaction between the SMC and LMC $\\sim$0.2 Gyr ago \\citep{mf80,yn03,ckg06}.  Within the bar we find a mean gradient of $26.30 \\pm 1.6$~\\kms~deg$^{-1}$ at PA~=~126$^\\circ$ (cp. 8.3 ~\\kms~deg$^{-1}$ at PA~$=$~23$^\\circ$ for the red giants), although, as the H$\\;$\\1 maps show, this simple parameterization of the data almost certainly disguises a more complex dynamical behaviour."
    },
    "0801/0801.1471_arXiv.txt": {
        "abstract": "% The backbone of double bars is made out of double-frequency orbits, and loops, their maps, indicate the bars' extent, morphology and dynamics. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.4915_arXiv.txt": {
        "abstract": "Observations with the Hinode space observatory led to the discovery of predominantly horizontal magnetic fields in the photosphere of the quiet internetwork region. Here we investigate realistic numerical simulations of the surface layers of the Sun with respect to horizontal magnetic fields and compute the corresponding polarimetric response in the \\ion{Fe}{1}~630~nm line pair. We find a local maximum in the mean strength of the horizontal field component at a height of around 500~km in the photosphere, where it surpasses the vertical component by a factor of 2.0 or 5.6, depending on the initial and boundary conditions. From the synthesized Stokes profiles we derive a mean horizontal field component that is, respectively, 1.6 and 4.3 times stronger than the vertical component. This is a consequence of both the intrinsically stronger flux density of, and the larger area occupied by the horizontal fields. We find that convective overshooting expels horizontal fields to the upper photosphere, making the Poynting flux positive in the photosphere, while this quantity is negative in the convectively unstable layer below it. ",
        "introduction": "\\label{sect1}  Recent observations with the spectropolarimeter of the Solar Optical Telescope (SOT) onboard the Hinode space observatory \\citep{kosugi+al07} indicate that the quiet internetwork region (the inner regions of supergranular cells of the quiet Sun) harbors a photospheric magnetic field whose  mean flux density of the horizontal component considerably surpasses that  of the vertical component \\citep{lites+al07,orozco+al07,lites+al08}. According to these papers, the vertical fields are concentrated in the intergranular lanes, whereas the stronger, horizontal fields  occur most commonly at the edges of the bright granules, aside from the vertical fields. In a gravitationally stratified atmosphere,  vertical magnetic flux concentrations naturally develop a horizontal component as they expand with height in a funnel-like manner. Indeed, \\cite{rezaei+al07} found funnel shaped magnetic elements in the internetwork from the same Hinode data. However, the newly discovered horizontal fields also occur apart from vertical flux concentrations and seem to cover a larger surface fraction than the vertical fields.  Regarding numerical simulations, \\citet{ugd+al98}, note: ``we find in all simulations also strong horizontal fields above convective upflows'', and \\citet{schaffenberger+al05,schaffenberger+al06} find frequent horizontal fields in their three-dimensional simulations, which they describe as ``small-scale canopies''. Also the 3-D simulations of \\citet{abbett07} display ``horizontally directed ribbons of magnetic flux that permeate the model chromosphere'', not unlike the figures shown by \\citet{schaffenberger+al06}. More recently, \\citet{schuessler+voegler08} find in a three-dimensional surface-dynamo simulation ``a clear dominance of the horizontal field in the height range where the spectral lines used for the Hinode observations are formed''.  Here we report on the analysis of existing and new three-dimensional magnetohydrodynamic computer simulations of the internetwork magnetic field aiming at following questions: Does a realistic simulation of the surface layers of the Sun intrinsically produce horizontal magnetic fields and can their mean flux density indeed surpass the mean flux density of the vertical field component? What is the polarimetric signal of this field and how does it compare to measurements with Hinode?  In the following we explain in Sect.~2 the details of two simulations and present results that answer the first two of the above questions in Sect.~3. In Sect.~4, we synthesize Stokes profiles and compare them to measurements from Hinode. Conclusions follow in Sect.~5. ",
        "conclusions": "\\label{sect5}  We have carried out  two simulations of magnetoconvection in the surface layers of the quiet internetwork region of the solar atmosphere. The simulations greatly differ in their initial state and boundary conditions for the magnetic field, but otherwise they both equal faithfully reproduce properties of normal granulation (as the magnetic field is weak). The top boundary is placed in the middle chromosphere (at a height of 1400~km) far away from the photospheric layers.  Both simulations intrinsically produce a horizontal magnetic field throughout the photosphere and lower chromosphere with a mean field strength that exceeds the mean strength of the vertical field component at the same height by up to a factor of 5.6. The strength of the horizontal field component shows a local maximum close to the classical temperature minimum near 500~km height (which largely escapes measurements with the \\ion{Fe}{1} 630~nm line pair). Fields with a horizontal component exceeding a certain limit in strength occupy a significantly larger surface area than fields with a vertical component exceeding this limit.  This horizontal field can be considered a consequence of the flux expulsion process \\citep{galloway+weiss81}: in the same way as magnetic flux is expelled from the granular interior to the intergranular lanes, it also gets pushed to the middle and upper photosphere by overshooting convection, where it tends to form a layer of horizontal field of enhanced flux density, reaching up into the lower chromosphere. Below the surface of $\\tau_{\\rm c} \\approx 0.1$, convective plumes pump horizontal magnetic field in the downward direction. Hence, this surface acts as a separatrix for the Poynting flux, which is mainly directed upwards above it and in the downward direction below it.  The response of this field in linear and circular polarization of the two neutral iron lines at 630~nm yields a ratio $\\langle  B^{\\mathrm{T}}_{\\mathrm{app}} \\rangle/ \\langle |B^{\\mathrm{L}}_{\\mathrm{app}}|\\rangle = 4.3 $ in case of run h20 (which, according to Sect.~2, we deem better to represent the conditions of internetwork regions). This is close to the  measurements of \\cite{lites+al08}, which indicate a factor of 5. Errors may come from the straylight produced by the spectrograph and polarization optics that was not taken into account with our PSF, the difference in mean absolute flux density (run h20, has only about half the measured value [see Sect.~\\ref{sect4}]), the frequency-independent treatment of radiative transfer, lacking spatial resolution, but also natural fluctuations. The predominance of the horizontal component, may possibly only exist on a scale comparable to or less than the spatial resolution of SOT. At full spatial resolution of the simulation we obtain a ratio of 2.8 instead of 4.3."
    },
    "0801/0801.1537_arXiv.txt": {
        "abstract": "We have analyzed various characteristic temperatures and energies of hole-doped high-$T_c$ cuprates as a function of a dimensionless hole-doping concentration ($p_u$). Entirely based on the experimental grounds we construct a unified electronic phase diagram (UEPD), where three characteristic temperatures ($T^*$'s) and their corresponding energies ($E^*$'s) converge as $p_u$ increases in the underdoped regime. $T^*$'s and $E^*$'s merge together with the $T_c$ curve and 3.5$k_BT_c$ curve at $p_u$ $\\sim$ 1.1 in the overdoped regime, respectively. They finally go to zero at $p_u$ $\\sim$ 1.3. The UEPD follows an asymmetric half-dome-shaped $T_c$ curve in which $T_c$ appears at $p_u$ $\\sim$ 0.4, reaches a maximum at $p_u$ $\\sim$ 1, and rapidly goes to zero at $p_u$ $\\sim$ 1.3. The asymmetric  half-dome-shaped $T_c$ curve is at odds with the well-known symmetric superconducting dome for La$_{2-x}$Sr$_x$CuO$_4$ (SrD-La214), in which two characteristic temperatures and energies converge as $p_u$ increases and merge together at $p_u$ $\\sim$ 1.6, where $T_c$ goes to zero. The UEPD clearly shows that pseudogap phase precedes and coexists with high temperature superconductivity in the underdoped and overdoped regimes, respectively. It is also clearly seen that the upper limit of high-$T_c$ cuprate physics ends at a hole concentration that equals to 1.3 times the optimal doping concentration for almost all high-$T_c$ cuprate materials, and 1.6 times the optimal doping concentration for the SrD-La214. Our analysis strongly suggests that pseudogap is a precursor of high-$T_c$ superconductivity, the observed quantum critical point inside the superconducting dome may be related to the end point of UEPD, and the normal state of the underdoped and overdoped high temperature superconductors cannot be regarded as a conventional Fermi liquid phase. ",
        "introduction": "The unique hallmark of high temperature superconductors (HTSs) is a pseudogap phase characterized by the observation of a multiple pseudogap temperatures ($T^*$'s) and pseudogap energies ($E^*$'s) by a large number of different experimental probes. While the pseudogap phase precedes the high temperature superconducting phase characterized by the superconducting transition temperature ($T_c$) and superconducting gap energy ($\\Delta_c$), it is not clear how $T^*$, $T_c$, $E^*$, and $\\Delta_c$ are related to each other. Specifically, how are $T^*$ and $E^*$ related to the occurrence of the high-$T_c$ superconductivity is still unclear. Is pseudogap a sufficient and/or necessary condition for high $T_c$ or is it just a complication of specific material systems? Is it collaborating or competing with superconductivity? For instance, it is argued that the pseudogap is a competing order that may have nothing to do with high $T_c$.\\cite{tal01} On the other hand it is also suggested that the pseudogap is intimately related to high $T_c$.\\cite{bat96,eme97} To distinguish these two contradictory pictures that are critical to the mechanism of high-$T_c$ superconductivity requires a comparison of various characteristic temperatures and energies in a universal phase diagram for all HTSs. Any systematic behavior derived from this kind of phase diagram will provide true intrinsic properties of HTS that are free from material-specific complications. However, up until now there is no such a comparison made and no such phase diagram available. We have analyzed numerous published data in the literature. We carefully select 27 HTSs: 11 single-layer, 11 double-layer and five triple-layer HTSs as summarized in the Table~\\ref{tab:table1}. The selection criteria will follow when we discuss the construction of the figures. There are 15 different experimental probes used for these 27 HTSs which are summarized in Table~\\ref{tab:table2}. In this paper we unify the characteristic temperatures of all these data of 27 HTSs on one single phase diagram entirely based on our proposed universal hole concentration scale that itself is also based on experimental results.   In the single-layer SrD-La214, where the hole-doping concentration can be unambiguously determined from the Sr-content ($x$),\\cite{tor88} $T_c$($x$) exhibits a well-known symmetric bell-shaped curve, i.e., the so-called superconducting dome, with a maximum $T_c$ ($T_c^{max}$) located at $x$ $\\sim$ 0.16.\\cite{pre91} The symmetrical dome-shaped $T_c$ curve or the superconducting dome is approximately represented by the following parabola, \\begin{eqnarray} 1 - \\frac{T_c}{T_c^{max}} = \\left. \\begin{array}{ll} 82.6  \\times (x-0.16)^2.  & \\label{form1} \\end{array} \\right. \\end{eqnarray} Assuming that all HTSs have the identical symmetric superconducting dome, $x$ can be replaced with the hole-doping concentration ($P_{T_c}$). Then, this relation could be used to determine the hole-doping concentration for many other HTSs.\\cite{pre91,naq05,obe92,ber96,wat00,sut03,haw07,tak01,cam99,din01,tan06,miy98,ozy00,moo93,blu97,cap07,gas97,ken95,sug03} Using this hole-scale based on the superconducting dome, the $P_{T_c}$-scale, various phase diagrams have been constructed.\\cite{tal01} A distinct feature in one of such phase diagrams is that $T^*$ crosses the superconducting dome and reaches zero at a quantum critical point (QCP) inside the dome.\\cite{tal01,naq05} On the other hand, without using the $P_{T_c}$ scale, some qualitative experimental observations seem to support another picture, where $T^*$ touches the superconducting dome at around $T_c^{max}$ and merges into the superconducting dome with no QCP inside the dome.\\cite{bat96} To distinguish these two fundamentally different pictures, we need a hole scale that can reveal the true intrinsic doping dependences of $T^*$, $T_c$, $E^*$, and $\\Delta_c$, which have been already observed.   \\begin{table}[b] \\caption{\\label{tab:table1}The chemical formula and the notation for the HTSs used in the present work.} \\begin{ruledtabular} \\begin{tabular}{ll} Chemical formula & Notation \\\\ \\hline (Single-layer HTS) \\\\ La$_{2-x}$Sr$_x$CuO$_4$ & SrD-La214 \\\\ La$_{2-x}$Ba$_x$CuO$_4$ & BaD-La214 \\\\ La$_2$CuO$_4$ & OD-La214 \\\\ (Nd$_{1.6-x}$Ce$_x$Sr$_{0.4}$)CuO$_4$ & CeD-NdSr214 \\\\ (La$_{1.6-x}$Nd$_{0.4}$Sr$_x$)CuO$_4$ & SrD-LaNd214 \\\\ Tl$_2$Ba$_2$CuO$_{6+\\delta}$ & OD-Tl2201 \\\\ Bi$_2$Sr$_{2-x}$La$_x$CuO$_{6+\\delta}$ & CD-Bi2201 \\\\ (Bi$_{1.74}$Pb$_{0.38}$)Sr$_{1.88}$CuO$_{6+\\delta}$ & OD-BiPb2201 \\\\ (Bi$_{1.35}$Pb$_{0.85}$)(Sr$_{1.47-x}$La$_{0.38+x}$)CuO$_{6+\\delta}$ & CD-BiPb2201 \\\\ HgBa$_2$CuO$_{4+\\delta}$ & OD-Hg1201 \\\\ Tl$_{1-x}$Pb$_x$Sr$_2$CuO$_{5-\\delta}$ & CD-TlPb1201 \\\\ \\hline (Double-layer HTS) \\\\ Y$_{1-x}$Ca$_x$Ba$_2$Cu$_3$O$_6$ & CaD-Y1236 \\\\ YBa$_2$Cu$_3$O$_{6+\\delta}$ & OD-Y123 \\\\ Y$_{1-x}$Ca$_x$Ba$_2$Cu$_3$O$_{6+\\delta}$ & CD-YCa123 \\\\ (Ca$_{1-x}$La$_{x}$)(Ba$_{1.75-x}$La$_{0.25+x}$)Cu$_3$O$_{6+\\delta}$ & CLBLCO \\\\ CaLaBaCu$_3$O$_{6+\\delta}$ & CLBCO \\\\ Bi$_2$Sr$_2$CaCu$_2$O$_{8+\\delta}$ & OD-Bi2212 \\\\ Bi$_2$Sr$_2$(Ca$_{1-x}$Y$_x$)Cu$_2$O$_{8+\\delta}$ & CD-Bi2212 \\\\ HgBa$_2$CaCu$_2$O$_{6+\\delta}$ & OD-Hg1212 \\\\ (Hg$_{0.5}$Fe$_{0.5}$)Ba$_2$(Ca$_{1-x}$Y$_x$)Cu$_2$O$_{6+\\delta}$ & CD-HgFe1212\\\\ Tl(BaSr)CaCu$_2$O$_{6+\\delta}$ & CD-Tl1212 \\\\ (Tl$_{0.5+x}$Pb$_{0.5-x}$)Sr$_2$(Ca$_{1-y}$Y$_y$)Cu$_2$O$_{6+\\delta}$ & CD-TlPb1212 \\\\ \\hline (Triple-layer HTS) \\\\ Bi$_2$Sr$_2$CaCu$_3$O$_{10+\\delta}$ & OD-Bi2223 \\\\ HgBa$_2$Ca$_2$Cu$_3$O$_{8+\\delta}$ & OD-Hg1223 \\\\ TlBa$_2$Ca$_2$Cu$_3$O$_{8+\\delta}$ & OD-Tl1223 \\\\ (Cu$_{1-x}$Ca$_x$)Ba$_2$Ca$_2$Cu$_3$O$_{8+\\delta}$ & CD-CuCa1223 \\\\ (Cu$_{1-x}$C$_x$)Ba$_2$Ca$_2$Cu$_3$O$_{8+\\delta}$ & CD-CuC1223 \\\\ \\end{tabular} \\end{ruledtabular} \\end{table}   \\begin{table}[b] \\caption{\\label{tab:table2}The experimental probes and their notations for the present work.} \\begin{ruledtabular} \\begin{tabular}{ll} Experimental probe & Notation \\\\ \\hline resistivity & $\\rho$ \\\\ $a$-axis resistivity & $\\rho_a$ \\\\ $c$-axis resistivity & $\\rho_c$ \\\\ in-plane resistivity & $\\rho_{ab}$ \\\\ inflection point of $\\rho$($T$) & $d^2$$\\rho$/$dT^2$ \\\\ thermoelectric power & TEP or $S$ \\\\ TEP at 290 K & $S^{290}$ \\\\ $a$-axis TEP & $S_a$ \\\\ in-plane TEP & $S_{ab}$ \\\\ susceptivility & $\\chi$ \\\\ susceptivility (H//$c$) & $\\chi_c$ \\\\ susceptivility (H//$ab$) & $\\chi_{ab}$ \\\\ nuclear magnetic resonance  & NMR \\\\ nuclear quardruple resonance & NQR \\\\ spin-lattice relaxation rate (NQR) & $(T_1T)^{-1}$ \\\\ NMR knight shift (H//$c$) & $K_c$ \\\\ angle-resolved photoemission spectroscopy & ARPES \\\\ angle-integrated photoemission spectroscopy & AIPES \\\\ superconductor-insulator-superconductor \\\\ tunneling & SIS \\\\ superconductor-insulator-normal metal \\\\ tunneling & SIN \\\\ near edge x-ray absorption fine structure & NEXAFS \\\\ electronic specific heat coefficient & $\\gamma$ \\\\ thermal conductivity & $\\kappa$ \\\\ neutron scattering & $neutron$ \\\\ electronic Raman scattering & $ERS$ \\\\ quasiparticle relaxation rate & $QPR$ \\\\ polar angular magnetoresistance oscillations & $AMRO$ \\\\ \\end{tabular} \\end{ruledtabular} \\end{table}   The common structural features of HTS are CuO$_2$-planes that host the doped holes and the block layers that supply the holes into the planes through oxygen-doping and/or cation-doping. While the doped hole-carriers are initially confined in the CuO$_2$-planes sandwiched between the insulator-like block layers, the holes are partially deconfined from the planes with doping. Therefore, the lightly doped HTS generally shows strongly two-dimensional (2D) properties. However, as the hole-doping increases, some physical properties are 2D and some, although built on the 2D carriers, will nominally be three-dimensional (3D) in nature. Therefore, it is necessary to use 2D and 3D carrier-doping concentrations to address 2D and 3D physical properties, respectively. To quantitatively study such dimensionality-dependent physical properties we have proposed a universal \\textit{\\textbf{planar}} hole scale ($P_{pl}$-scale) for determining the hole-doping content per CuO$_2$ plane ($P_{pl}$).\\cite{hon04} In this scale, the $P_{pl}$ is uniquely determined from  $S^{290}$.\\cite{hon04} We showed that, in Ref.\\ \\onlinecite{hon04}, the $P_{pl}$-scale is independent of the nature of the dopant, the number of CuO$_2$-plane layers per formula unit cell ($n_{layer}$), the structure and the sample quality, namely, single crystal or not. This universal $S^{290}$($P_{pl}$)-relationship is built on the sound experimental observations, which is similar to the situation of the most popular $P_{T_c}$-scale, although it is still mainly empirical and waited to be theoretically justified. Since the average area per copper in the CuO$_2$-plane is almost independent of the HTS materials, therefore, $P_{pl}$ is essentially equal to 2D hole-doping concentration defined as the hole-doping content per unit area. Using the 2D $P_{pl}$-scale, it was found in the phase diagram for all major HTSs plotted as a function of $P_{pl}$ that the $T^*$-curves are independent of the $n_{layer}$ while the $T_c$-curve strongly depends on it.\\cite{hon04} Therefore the $P_{pl}$-scale is intrinsically consistent with the pseudogap energy scale.\\cite{hon04} We can also extend the hole-doping content per CuO$_2$ plane to an effective 3D hole-doping content per CuO$_2$ block, which includes the oxygen coordination around the plane, ($P_{3D}$) by a simple conversion formula $P_{3D}$ $\\equiv$ $P_{pl}$ $\\times$ ($n_{layer}$/$V_{u.c.}$), where $V_{u.c.}$ is the unit cell volume.\\cite{hon06} Since $P_{3D}$ is essentially the hole-doping content per unit volume, therefore, this natural extension of $P_{pl}$-scale to $P_{3D}$ ($P_{3D}$-scale) has allowed us to address the corresponding 3D properties. \\cite{hon06} For instance, in the case of the single-layer HTS, the Hall number per ``$cm^3$'', calculated from the in-plane Hall coefficient, is not scaled with $P_{pl}$ but $P_{3D}$.\\cite{hon06} The $\\tau_c$($T_c$), a reduced temperature-scale defined as $\\tau_c$($T$) $\\equiv$ $T$/$T_c^{max.}$, of the single-layer HTS universally appears at 6 $\\times$ 10$^{20}$ cm$^{-3}$ and reaches the $T_c^{max.}$ at 1.6 $\\times$ 10$^{21}$ cm$^{-3}$ as shown in Fig.\\ \\ref{fig1}(b),\\cite{hon06} although their critical hole-doping concentrations on the $P_{pl}$-scale depend on the materials as shown in Fig.\\ \\ref{fig1}(a).\\cite{hon04} Thus, it was shown that various normal and superconducting properties for many different material systems can be consistently compared by using either $P_{pl}$ or $P_{3D}$.\\cite{hon04,hon06,hon07}   \\begin{figure} \\includegraphics[scale=0.5]{Figure1.pdf} \\caption{\\label{fig1} (Color online) For the single-layer HTSs, (a) the superconducting transition temperature ($T_c$) as a function of $P_{pl}$, (b) the reduced superconducting transition temperature $\\tau_c$($T_c$) ($\\equiv$ $T_c$/$T_c^{max}$) as a function of the effective 3D hole-doping concentration $P_{3D}$, and (c) the $\\tau_c$($T_c$) as a function of $P_{3D}$. The plotted data are summarized in the Table~\\ref{tab:table3}. The broken line comes from the equation (\\ref{form1}). The solid line is our superconducting dome.} \\end{figure}  \\begin{table} \\caption{\\label{tab:table3}The $T_c^{max.}$ and $P_{pl}^{opt.}$ for single-layer HTSs plotted in Figs.\\ \\ref{fig1}(a) -\\ \\ref{fig1}(c).} \\begin{ruledtabular} \\begin{tabular}{lcll} \\qquad HTS & $T_c^{max.}$ (K) & $P_{pl}^{opt.}$ & Ref(s). \\quad \\\\ \\hline \\quad  SrD-La214 & 39.4 & 0.16 & \\onlinecite{nag93} \\\\ \\quad  SrD-La214 & 37 & 0.16 & \\onlinecite{rad94} \\\\ \\quad  SrD-La214 & 36 & 0.16 & \\onlinecite{kak99} \\\\ \\quad  SrD-La214 & 38 & 0.16 & \\onlinecite{kom05} \\\\ \\quad  OD-Hg1201 & 97 & 0.235 & \\onlinecite{yam00,wil00} \\\\ \\quad  CD-Bi2201 & 35.5  & 0.28 & \\onlinecite{and00} \\\\ \\quad  CD-Bi2201 & 33  & 0.28 & \\onlinecite{oka06} \\\\ \\quad  OD-Tl2201 & 93$^{\\text{a}}$  & 0.25$^{\\text{a}}$ & \\onlinecite{obe92,shu93,mar95} \\\\ \\end{tabular} \\end{ruledtabular} \\footnotetext[1]{We use the reported highest $T_c$ = 93 K as $T_c^{max.}$ (Ref.\\ \\onlinecite{wag97}). From the plot of $T_c$ vs $P_{pl}$ in Fig.\\ \\ref{fig1}(a), the optimal $P_{pl}$ is estimated to be $\\sim$0.25. The detail is in the text.} \\end{table}   In order to reveal the intrinsic generic electronic properties of all HTSs, it is necessary to be able to put both 2D and 3D physical properties on a single phase diagram. To achieve this goal, we need a carrier-scale that is not only independent of the material system but also independent of the dimensionality of the physical properties. This can be achieved if, for each material system, we scale $P_{pl}$ and $P_{3D}$ with their corresponding optimal doping concentrations, $P_{pl}^{opt.}$ and $P_{3D}^{opt.}$, respectively. Here, we introduce a dimensionless unified hole-doping concentration, $p_u$ ($p_u$ $\\equiv$ $P_{pl}$/$P_{pl}^{opt.}$ = $P_{3D}$/$P_{3D}^{opt.}$). This unified hole scale ($p_u$ scale) can be used for all physical properties, which is independent of their dimensionality, in all HTSs. Indeed, the identical doping dependent behaviors are preserved even though $\\tau_c$($T_c$) of the single-layer HTS plotted as a function of $P_{3D}$ in Fig.\\ \\ref{fig1}(b) was replotted as a function of $p_u$ in Fig.\\ \\ref{fig1}(c).\\cite{rad94,nag93,kom05,yam00,wil00,and00,oka06,obe92,shu93,mar95,kak99} Here, each $P_{pl}^{opt.}$ was determined from the plot of $T_c$ vs $P_{pl}$ for the each compound in the present work or Refs.\\ \\onlinecite{hon04,hon06,hon07}. For the OD-Tl2201 there was few reports on the optimally doped samples because the optimally doped OD-Tl2201 is hard to prepare. In this case we use the highest $T_c$ = 93 K among the published data as $T_c^{max.}$.\\cite{wag97} From the plot of $T_c$ vs $P_{pl}$ in Fig.\\ \\ref{fig1}(a), the optimal $P_{pl}$ is estimated to be $\\sim$0.25. They are summarized in Table~\\ref{tab:table3}. Essentially we can view $p_u$ as a scaled dimensionality- and material- independent universal carrier-doping concentration that preserves the intrinsic doping dependency for \\textit{\\textbf{any}} physical property for \\textit{\\textbf{all}} HTSs. In this paper, we have analyzed the characteristic temperatures and energies observed in the 27 HTSs by 15 different experimental probe as a function of $p_u$. We find a dopant-specific unified electronic phase diagram for HTS. The dominate phase diagram is an asymmetric half-dome-shaped $T_c$-curve for the cation and anion (oxygen) co-doped (CD) HTS. $T_c$ for the purely oxygen-doped (OD) HTS also follows the half-dome-shaped $T_c$ curve with some indication of the influence of the thermally induced oxygen redistribution. ",
        "conclusions": "\\subsection{Universal hole-doping scale}   \\begin{figure}[t] \\includegraphics[scale=0.5]{Figure3.pdf} \\caption{\\label{fig3} (Color online) Hole-doping concentration determined by various techniques as a function of $P_{pl}$. The plotted data are are summarized in Table~\\ref{tab:table5}.} \\end{figure}   First of all, we demonstrate how the hole-doping scale based on the $S^{290}$ is effective and universal. In Figs.\\ \\ref{fig2}(a) and (b), we plot $S^{290}$ of sintered sample and $S_{ab}^{290}$ of the single crystal as a function of $P_{pl}$, together with previously reported data.\\cite{hon04} $S^{290}$ ($\\geq$ 7 $\\mu$V/K) on the upper panel is plotted on a logarithmic scale, while $S^{290}$ ($<$ 7 $\\mu$V/K) on the lower panel is plotted on a linear scale. In Figs.\\ \\ref{fig2}(a) and (b), the five single-layer, one double-layer, and one triple-layer HTSs are the newly added data points. They have been plotted with the previous reported SrD-La214 and CaD-Y1236. The plotted data are listed in Table~\\ref{tab:table4}. $P_{pl}$ of SrD-La214 without excess oxygen is equal to Sr-content.\\cite{tor88} The $P_{pl}$ of CaD-Y1236, which the oxygen-content was determined to be 6 by the iodometric titration in Ar gas,\\cite{hon04} can be ambiguously and directly determined as a half of Ca-content, since the CaD-Y1236 has the isolated Cu layer in stead of CuO chain. In fact, it is shown by the O $1s$ and Cu $2p$ NEXAFS experiment that the holes introduced by replacing Y$^{3+}$ with Ca$^{2+}$ appear solely in the CuO$_2$ planes without affecting the isolated Cu layers in the CaD-Y1236.\\cite{mer98} $P_{pl}$ for the other materials were determined from the copper valency measured by the iodometric titration for the OD-La214 \\cite{yu96} and CD-Bi2201,\\cite{dev90,smi92} and the double iodometric titration for the OD-Hg1212,\\cite{fuk97} CD-HgFe1212, \\cite{kan05} and OD-Hg1223.\\cite{fuk97} The error of $P_{pl}$ is mainly coming from the oxygen-deficient ($\\delta$). For the double- and triple-layer HTSs, the error of $\\delta$ was below 0.01.\\cite{hon04,fuk97,kan05} For the SrD-La214, the oxygen deficient is estimated to be $\\sim$0.005, according to the result of Radaelli $et$ $al$.\\cite{rad94} These error of $P_{pl}$ can be estimated to be below 0.01. For the CD-Bi2201, the error of $\\delta$ is $\\sim$0.02,\\cite{dev90,smi92} and therefore the error of $P_{pl}$ is $\\sim$0.04. Noticed that the plotted data follow the universal $S^{290}$($P_{pl}$)-curve proposed in Ref.\\ \\onlinecite{hon04}, which is irrespective of the nature of dopant, $n_{layer}$, the structure and the sample quality, namely, single crystal or not. It is also independent from whether the CuO$_2$ plane is surrounded by the octahedral or pyramidal oxygen coordination. For the SrD-La214, there is the upward deviation from the universal line at $P_{pl}$ $>$ 0.25. This deviation is considered to be due to the oxygen-deficient that was reported to be significant over $x$ = $P_{pl}$ $\\sim$ 0.25.\\cite{rad94} In the CeD-NdSr214, the upward deviation over $P_{pl}$ $\\sim$ 0.15 from the universal line can be explained by the oxygen deficiency generating the hole deficient of $\\sim$0.05 as pointed out in Ref.\\ \\onlinecite{amb02}. Accordingly, all plotted data lie in a shaded area around our universal $S^{290}$($P_{pl}$)-curve with the $P_{pl}$-accuracy of $\\pm$0.01 within the reported error. Therefore, the proposed universal $S^{290}$($P_{pl}$)-curve that purely based on the experimental grounds works well as the empirical intrinsic hole-scale for the HTS in the range of 0.01 $<$ $P_{pl}$ $<$ 0.34. In Fig.\\ \\ref{fig2}(c), we compare $P_{pl}$ with $P_{T_c}$. The solid line shows $P_{pl}$ as a function of $P_{T_c}$. The broken line shows $P_{pl}$ = $P_{T_c}$. The quantitative difference between the $P_{pl}$ scale and $P_{T_c}$ scale becomes clear in Fig.\\ \\ref{fig2}(c). In addition we used this relation for the conversion from $P_{T_c}$ into $P_{pl}$ when the data plotted here have the $P_{T_c}$ without TEP.   Next, we compare our universal scale based on the $S^{290}$ to that determined by other techniques. The hole-doping concentration by ARPES ($P_{ARPES}$)  is deduced from the area of the experimental Fermi surface (FS). The planar hole-doping concentration is also determined by NEXAFS ($P_{NEXAFS}$), by NQR ($P_{NQR}$) and by AMRO ($P_{AMRO}$). In Fig.\\ \\ref{fig3}, we plot the $P_{ARPES}$, $P_{NEXAFS}$, $P_{NQR}$ and $P_{AMRO}$ as a function of $P_{pl}$. The plotted data are summarized in Table~\\ref{tab:table5}. It can be clearly seen that the $P_{pl}$ determined by TEP is quite consistent with $P_{NEXAFS}$ and $P_{NQR}$. Although there is a slight scattering, $P_{pl}$ is also consistent with $P_{ARPES}$ and $P_{AMRO}$. Thus, our $P_{pl}$-scale is consistent with above other scales. Accordingly, the present $p_u$ scale is also intrinsically consistent with the hole-doping concentrations determined by the above techniques.   \\begin{table}[t] \\caption{\\label{tab:table5}The data plotted in Fig.\\ \\ref{fig3}.} \\begin{ruledtabular} \\begin{tabular}{llcl} \\quad Probe & \\quad HTS & $P_{pl}$ & Ref. \\qquad \\\\ \\hline \\quad ARPES & SrD-La214 (a) &  & \\onlinecite{ino02} \\\\ & SrD-La214 (b) &  & \\onlinecite{yos06} \\\\ & OD-BiPb2201 &  & \\onlinecite{kon05} \\\\ & CD-BiPb2201 &  & \\onlinecite{kon05} \\\\ \\hline \\quad NEXAFS & CaD-Y1236 &  & \\onlinecite{mer98} \\\\ & OD-Y123 & I & \\onlinecite{mer98} \\\\ & CD-YCa123 ($x$=0.1) & I & \\onlinecite{mer98} \\\\ \\hline \\quad NQR    & OD-Y123 & I & \\onlinecite{kot01} \\\\ & OD-Tl2201 & I & \\onlinecite{kot01} \\\\ \\hline \\quad AMRO & OD-Tl2201 & I &\\onlinecite{hus03} \\\\ \\end{tabular} \\end{ruledtabular} \\end{table}   \\begin{figure}[b] \\includegraphics[scale=0.5]{Figure4.pdf} \\caption{\\label{fig4} (Color online) $T_c$ and $T_{\\rho}^*$ as a function of $p_u$ for (a) the OD-Y123-related materials and (b) Y$_{0.8}$Ca$_{0.2}$Ba$_2$(Cu$_{1-y}$Zn$_y$)$_3$O$_{6+\\delta}$. The plotted data are summarized in Table~\\ref{tab:table6}. The solid line is a half-dome-shaped $T_c$-curve with $T_c^{max}$ = 86 K. The dotted line comes from the equation (\\ref{form1}) with $T_c^{max}$ = 86 K.} \\end{figure}   \\begin{table}[b] \\caption{\\label{tab:table6}The data plotted in Figs.\\ \\ref{fig4}(a) and (b).} \\begin{ruledtabular} \\begin{tabular}{llcl} \\quad Fig. & \\quad HTS & $P_{pl}$ & Ref(s). \\quad \\\\ \\hline \\quad 4(a)  & CD-YCa123 ($x$=0.2) &  & \\onlinecite{coo00,ber96} \\\\ & CD-YCa123 ($x$=0.05) & II & \\onlinecite{naq05} \\\\ & CLBLCO ($x$=0.4) & & \\onlinecite{kni99} \\\\ & CLBCO & & \\onlinecite{hay96} \\\\ \\hline \\quad 4(b)  & Y$_{0.8}$Ca$_{0.2}$Ba$_2$(Cu$_{1-y}$Zn$_y$)$_3$O$_{6+\\delta}$ & II & \\onlinecite{coo00} \\\\ & (0 $\\leq$ $y$ $\\leq$ 0.04) & & \\\\ \\end{tabular} \\end{ruledtabular} \\end{table}   \\subsection{Asymmetric half-dome-shaped $T_c$ curve}   In Fig.\\ \\ref{fig4}(a), we plot the $T_c$ and $T_{\\rho}^*$ of all the OD-Y123 related materials which do not have significant contribution of CuO chain as a function of $p_u$. It includes CD-YCa123,\\cite{naq05,coo00,ber96} CLBLCO,\\cite{kni99} and CLBCO.\\cite{hay96} First, we note that the $T_c$($p_u$) does not follow the well-known superconducting dome, as shown as a dotted line in Fig.\\ \\ref{fig4}(a), instead, it follows an asymmetric half-dome-shaped curve shown as a solid line. Although $T_c$ in the underdoped regime basically follows the superconducting dome, $T_c$ in the overdoped regime decreases much more rapidly. In Fig.\\ \\ref{fig4}(a), $T_{\\rho}^*$ decreases with doping, smoothly merges into the half-dome-shaped $T_c$ curve and finally tends to reach an end point located at ($p_u$,$T$) = (1.3,0). Therefore, in contrast to the proposal that the $T_{\\rho}^*$ curve crosses the $T_c$ curve,\\cite{tal01} the $T_{\\rho}^*$ curve smoothly merges into the $T_c$ curve in the overdoped regime. In Fig.\\ \\ref{fig4}(b), we plot $T_c$ and $T_{\\rho}^*$ as a function of $p_u$ for Y$_{0.8}$Ca$_{0.2}$Ba$_2$(Cu$_{1-y}$Zn$_y$)$_3$O$_{6+\\delta}$ for 0 $\\leq$ $y$ $\\leq$ 0.04.\\cite{naq05}, Although  $T_{\\rho}^*$($p_u$) slightly depends on Zn-content in the overdoped regime, $T_{\\rho}^*$($p_u$) again tends to merge into $T_c$($p_u$) at the overdoped regime. This should be compared to the original plot, based on the $P_{T_c}$-scale, in which $T_{\\rho}^*$ crossed the superconducting dome and reached zero at a proposed QCP ($P_{Tc}$ = 0.19).\\cite{naq05} Accordingly, the crossing was a artifact that came from two sources; one is the use of a hole-scale that failed to taking into account the differences in dimensionality of different physical properties, namely, the two-dimensional $T_{\\rho}^*$ vs\\ three-dimensional \\textit{\\textbf{$T_c$}}, {\\bf and} the other is that the $T_c$ curve for the majority of HTS follows the asymmetric dome-shaped curve, \\textit{\\textbf{only}} SrD-La214 follows the symmetric dome-shaped $T_c$curve or superconducting dome.   $\\tau_c$($T_c$) vs $p_u$ plot for the cation and oxygen co-doped HTS and the purely oxygen-doped HTS are shown in Figs.\\ \\ref{fig5}(a) and (b), respectively. For comparison, $\\tau_c$($T_c$) vs $P_{pl}$ curve of OD-Y123 reported in Ref.\\ \\onlinecite{hon07} is also plotted in Fig.\\ \\ref{fig5}(b). $T_c^{max.}$ and $P_{pl}^{opt.}$ are summarized in Table~\\ref{tab:table7}. The CD-HTSs follow the present asymmetric half-dome-shaped $T_c$-curve. $\\tau_c$($T_c$) vs $p_u$ curve of the single-layer OD-Tl2201, which behaves differently from that of the other in the plot of $T_c$ vs $P_{3D}$ as shown in Fig.\\ \\ref{fig1}(b),\\cite{hon06} actually follows the asymmetric half-dome shaped $T_c$-curve. The other overdoped OD-HTSs also follow the half-dome-shaped $T_c$ curve. Note that $T_c$ of the underdoped OD-HTS is slightly enhanced from the half-dome shaped $T_c$-curve and $T_c$ appears at a lower $p_u$. The OD-Y123 also shows the similar trend, although it is influenced by the CuO chain ordering.\\cite{hon07} We attribute this to the influence of the soft oxygen dopants.\\cite{lor02} Thus, opposite to the common belief, the $\\tau_c$($T_c$) vs $p_u$ phase diagram of the majority of HTSs follow the asymmetric half-dome shaped $T_c$ curve. Noticed that the asymmetric half-dome-shaped $T_c$ curve goes to zero at $p_u$ $\\sim$ 1.3. It is interesting to point out that if we take $P_{pl}^{opt.}$ to be universally equal to 0.16, as assumed in the $P_{T_c}$ scale, then the critical $p_u$ $\\sim$ 1.3 corresponds to $P_{T_c}$ $\\sim$ 0.2 in the $P_{T_c}$ scale. This value is very close to the proposed QCP ($P_{T_c}$ = 0.19) identified by various experiments on the $P_{T_c}$-scale.\\cite{tal01,naq05} Therefore this critical doping concentration is not located inside the superconducting phase and, physically, it is the doping concentration where all the phenomenology of high $T_c$ ceases to exist and the ground state becomes a conventional Fermi liquid (FL) for $p_u$ $>$ 1.3.   \\begin{figure}[b] \\includegraphics[scale=0.5]{Figure5.pdf} \\caption{\\label{fig5} (Color online) Extended unified electronic phase diagram plotted as $\\tau_c$($T_c$) vs $p_u$ for (a) the cation and oxygen co-doped HTS's and (b) the purely oxygen doped HTS. The plotted data are summarized in Table~\\ref{tab:table7}. The solid and broken lines are an asymmetric half-dome-shaped $T_c$ curve and our superconducting dome, respectively. The dotted line is the $T_c$ curve for OD-Y123.\\cite{hon07}} \\end{figure}   \\begin{table} \\caption{\\label{tab:table7}The $T_c^{max.}$ and $P_{pl}^{opt.}$ for the HTSs plotted in Figs.\\ \\ref{fig5}(a) and\\ \\ref{fig5}(b).} \\begin{ruledtabular} \\begin{tabular}{llclcl} Fig. & \\qquad HTS & $T_c^{max.}$(K) & $P_{pl}^{opt.}$ & $P_{pl}$ & Ref(s). \\\\ \\hline 5(a) & (Single-layer HTS) \\\\ & CD-TlPb1201 & 50 & 0.25 & & \\onlinecite{sub94} \\\\ & CD-Bi2201 & 35.5 & 0.28 & &  \\onlinecite{and00} \\\\ & CD-Bi2201 & 33 & 0.28 & &  \\onlinecite{oka06} \\\\ \\cline{2-6} & (Double-layer HTS) \\\\ & CD-YCa123 ($x$=0.2) & 85 & 0.237 & &  \\onlinecite{coo00,lor98} \\\\ & CD-YCa123 ($x$=0.2) & 81 & 0.238 & II &  \\onlinecite{naq05} \\\\ & CD-YCa123 ($x$=0.2) & 85.5 & 0.238 & &  \\onlinecite{ber96} \\\\ & CLBLCO ($x$=0.4) & 81 & 0.235 & &  \\onlinecite{kni99} \\\\ & CLBLCO ($x$=0.1) & 57.7 & 0.205 & &  \\onlinecite{kni99} \\\\ & CD-TlPb1212 ($x$=0) & - $^{\\text{a}}$ & 0.235 & &  \\onlinecite{ber96} \\\\ & CD-TlPb1212 & 94 & 0.235 & &  \\onlinecite{sub92} \\\\ & CD-Tl1212 & 90 & 0.235 & &  \\onlinecite{mar95} \\\\ & CD-Bi2212 & 92 & 0.236 & &  \\onlinecite{ako98} \\\\ & CD-Bi2212 & 81 & 0.238 & &  \\onlinecite{man96} \\\\ & CD-HgFe1212 & 73 & 0.227 & &  \\onlinecite{kan05} \\\\ \\cline{2-6} & (Triple-layer HTS) \\\\ & CD-CuCa1223 & 122 & 0.248 & &  \\onlinecite{cao97} \\\\ & CD-CuC1223 & 110 & 0.215 & &  \\onlinecite{cao97} \\\\ \\hline 5(b) & (Single-layer HTS) \\\\ & OD-Tl2201 & 93$^{\\text{b}}$ & 0.25$^{\\text{b}}$ & &  \\onlinecite{obe92,shu93,mar95} \\\\ & OD-Hg1201 & 97 & 0.235 & &  \\onlinecite{yam00,wil00} \\\\ \\cline{2-6} & (Double-layer HTS) \\\\ & OD-Bi2212 & 92 & 0.238 & &  \\onlinecite{obe92} \\\\ & OD-Hg1212 & 127 & 0.227 & &  \\onlinecite{fuk97} \\\\ & OD-Hg1212 & 125 & 0.227 & &  \\onlinecite{coh99} \\\\ \\cline{2-6} & (Triple-layer HTS) \\\\ & OD-Bi2223 & 108 & 0.215 & &  \\onlinecite{fuj02} \\\\ & OD-Hg1223 & 135 & 0.215 & &  \\onlinecite{fuk97} \\\\ & OD-Hg1223 & 138 & 0.215 & &  \\onlinecite{sub95} \\\\ & OD-Tl1223 & 128 & 0.23 & &  \\onlinecite{mik06} \\\\  \\end{tabular} \\end{ruledtabular} \\footnotetext[1]{The $T_c$/$T_c^{max.}$ was reported.} \\footnotetext[2]{We use the reported highest $T_c$ = 93 K as $T_c^{max.}$.\\cite{wag97} From the plot of $T_c$ vs $P_{pl}$ in Fig.\\ \\ref{fig1}(a), the optimal $P_{pl}$ is estimated to be $\\sim$0.25. The detail is in the text.} \\end{table}   In the overdoped triple-layer HTS, the charge density of the inner and outer planes were reported to be inhomogeneous.\\cite{tok99} This is consistent with the $\\tau_c$($T_c$) vs $p_u$ behavior of OD-Bi2223, black triangles in Fig.\\ \\ref{fig5}(b), that $T_c$ shows a flat region in the overdoped regime.\\cite{fuj02} However, the $\\tau_c$($T_c$) vs $p_u$ behaviors of CD-CuCa1223,\\cite{cao97} and CD-CuC1223 \\cite{cao97} plotted into Fig.\\ \\ref{fig5}(a) show the same trend as that of the single- and double-layer CD-HTSs. $\\tau_c$($T_c$) vs $p_u$ of OD-Tl1223,\\cite{mik06} and OD-Hg1223 \\cite{fuk97,sub95} plotted in Fig.\\ \\ref{fig5}(b) also show the identical trend as that of the single- and double-layer OD-HTSs. Accordingly, although counter-intuitive, the charge density of the inner and outer planes of these materials is expected to be the same.   \\subsection{Unified electroni phase diagram}   \\begin{figure*} \\includegraphics[scale=0.6]{Figure6ab.pdf}  \\caption{\\label{fig6} (Color online) Unified electronic phase diagram for the single- and double-layer HTS's with $T_c^{max}$ $\\sim$ 90 K. The temperature- and energy-scale for the pseudogap and superconducting gap obtained from transport properties are summarized in Figs.\\ \\ref{fig6}(a) and\\ \\ref{fig6}(b), from the spectroscopy properties in Figs.\\ \\ref{fig6}(c) and (d), and from NMR, QPR and scattering properties in Figs.\\ \\ref{fig6}(e) and\\ \\ref{fig6}(f), respectively. For Fig.\\ \\ref{fig6}(a) and\\ \\ref{fig6} (b), the plotted data are summarized in the Table~\\ref{tab:table8}. For Figs.\\ \\ref{fig6}(c) -\\ \\ref{fig6}(f), the plotted data are summarized in the Table~\\ref{tab:table9}. The $E_{hump}^*$, $T_{up}^*$, $T_{lp}^*$, and $T_c$ curves are directly determined from the plotted data. The $T_{hump}^*$, $E_{up}^*$, $E_{lp}^*$, and $\\Delta_c$ curves are calculated from the $E_{hump}^*$, $T_{up}^*$, $T_{lp}^*$, and $T_c$-curves using a relation of $T$ = $E$/$zk_B$ or $E$ = $zk_BT$, respectively. The broken lines show the $T_c$ curve or $\\Delta_c$ curve the for OD-Y123. In the energy-scale, the solid lines corresponds to $z$ = 3.5 and gray zone shows the energy range from 3$k_BT$ to 4$k_BT$. } \\end{figure*}   \\begin{figure*} \\includegraphics[scale=0.6]{Figure6cd.pdf} \\includegraphics[scale=0.6]{Figure6ef.pdf} \\end{figure*}   \\begin{table} \\caption{\\label{tab:table8}The HTSs plotted in Figs.\\ \\ref{fig6}(a) and\\ \\ref{fig6}(b).} \\begin{ruledtabular} \\begin{tabular}{lllcl} Fig. & Probe & HTS & $P_{pl}$ & Ref(s). \\\\ \\hline 6(a) & $\\rho$      & OD-Y123           & I & \\onlinecite{wuy96} \\\\ &             & CD-YCa123 ($x$=0.2) & II & \\onlinecite{naq05} \\\\ &             & CLBCO            &  & \\onlinecite{hay96} \\\\ &             & OD-Hg1201         &     & \\onlinecite{yam00} \\\\ \\cline{2-5} & $\\rho_a$    & CD-Bi2212         &     & \\onlinecite{wat00} \\\\ \\cline{2-5} & $\\rho_{ab}$ & OD-Y123           & I & \\onlinecite{ito93} \\\\ &             & OD-Bi2212         & I & \\onlinecite{oda97} \\\\ \\cline{2-5} & $\\rho_c$    & OD-Y123           & I & \\onlinecite{tak94,bav99} \\\\ &             & OD-Bi2212         & I & \\onlinecite{wat00} \\\\ \\cline{2-5} & $d^2\\rho_a/dT^2$ & OD-Y123         & I & \\onlinecite{and04} \\\\ \\cline{2-5} & TEP         & OD-Y123           &     & \\onlinecite{obe92} \\\\ &             & CD-YCa123 ($x$=0.2) &     & \\onlinecite{ber96} \\\\ &             & CLBCO             &     & \\onlinecite{hay96} \\\\ &             & CLBLCO ($x$=0.4)  &     & \\onlinecite{kni99} \\\\ &             & OD-Bi2212         &     & \\onlinecite{obe92} \\\\ &             & CD-Bi2212         &     & \\onlinecite{man96,ako98} \\\\ &             & OD-Hg1201         &     & \\onlinecite{yam00}\\\\ \\cline{2-5} & $\\gamma$    & OD-Y123         & I & \\onlinecite{lor93} \\\\ \\cline{2-5} & QPR         & OD-Y123           & I & \\onlinecite{kab99} \\\\ \\hline 6(b)  & $\\gamma$    & CD-YCa123 ($x$=0.2) & I & \\onlinecite{low97} \\\\ &             & OD-Y123           &     & \\onlinecite{coo96} \\\\ &             & OD-Bi2212         & I & \\onlinecite{low00} \\\\ \\cline{2-5} & $\\kappa$     & OD-Y123           & I & \\onlinecite{sut03} \\\\ &              & OD-Bi2212         & I & \\onlinecite{cha00} \\\\ &             & OD-Tl2201         & I & \\onlinecite{haw07} \\\\ \\cline{2-5} & QPR         & OD-Y123           & I & \\onlinecite{kab99} \\\\ &             & CD-YCa123 ($x$=0.2) & I & \\onlinecite{dem99} \\\\ \\end{tabular} \\end{ruledtabular} \\end{table}   \\begin{table} \\caption{\\label{tab:table9}The HTSs plotted in Figs.\\ \\ref{fig6}(c) -\\ \\ref{fig6}(f).} \\begin{ruledtabular} \\begin{tabular}{lllcl} Fig. & Probe & HTS & $P_{pl}$ & Ref(s). \\\\ \\hline 6(c)  & ARPES       & OD-Bi2212         & I & \\onlinecite{cam99,din96,sat02} \\\\ \\cline{2-5} & AIPES       & OD-Bi2212         & I & \\onlinecite{tak01} \\\\ \\cline{2-5} & SIS         & OD-Bi2212         & I & \\onlinecite{dip02} \\\\ \\hline 6(d)  & ARPES       & OD-Y123           & I & \\onlinecite{lu01} \\\\ &             & OD-Bi2212         & I & \\onlinecite{cam99,sat02,tan06,din01,dam03} \\\\ &             & OD-Tl2201         & I & \\onlinecite{pee07} \\\\ \\cline{2-5} & SIS         & OD-Y123           & I & \\onlinecite{wei98} \\\\ &             & OD-Bi2212         & I & \\onlinecite{ren98,miy98,zas01,dip02,ozy00} \\\\ &             & OD-Tl2201         & I & \\onlinecite{ozy99} \\\\ \\cline{2-5} & SIN         & OD-Bi2212         & I & \\onlinecite{ozy00} \\\\ \\hline 6(e)  & ($T_1T$)$^{-1}$  & OD-Y123      & I & \\onlinecite{tak91,got96} \\\\ &             & OD-Bi2212         & I & \\onlinecite{ish98} \\\\ &             & OD-Hg1201         & I & \\onlinecite{ito98} \\\\ \\cline{2-5} & $K_c$       & OD-Bi2212 & I & \\onlinecite{ish98} \\\\ &             & OD-Hg1201         & I & \\onlinecite{ito98} \\\\ \\cline{2-5} & neutron     & OD-Y123           & I & \\onlinecite{faq06,moo93} \\\\ \\cline{2-5} & $\\chi$      & OD-Bi2212         & I & \\onlinecite{oda97} \\\\ &             & OD-Tl2201         & I & \\onlinecite{kub91} \\\\ \\cline{2-5} & $\\chi_{ab}$ and $\\chi_c$ & OD-Bi2212  & I & \\onlinecite{wat00} \\\\ \\hline 6(f)  & $B_{1g}$ ERS & OD-Y123          & I & \\onlinecite{pai06,sug03,mak94} \\\\ &             & CD-YCa123 ($x$=0.2) & I & \\onlinecite{pai06} \\\\ &             & OD-Bi2212         & I & \\onlinecite{cap07,blu97,hew02,ken95,ven02} \\\\ &             & CD-Bi2212         & I & \\onlinecite{sug00} \\\\ &             & OD-Tl2201         & I & \\onlinecite{kan96} \\\\ &             & OD-Hg1201         & I & \\onlinecite{tac06} \\\\ \\cline{2-5} & $B_{2g}$ ERS & OD-Y123          & I & \\onlinecite{sug03} \\\\ &             & OD-Bi2212         & I & \\onlinecite{hew02,ken95,ven02} \\\\ &             & CD-Bi2212         & I & \\onlinecite{sug03} \\\\ &             & OD-Tl2201         & I & \\onlinecite{gas98} \\\\ &             & OD-Hg1201         & I & \\onlinecite{tac06} \\\\ \\cline{2-5} & neutron     & OD-Y123           & I & \\onlinecite{moo93} \\\\ \\end{tabular} \\end{ruledtabular} \\end{table}   We now examine various characteristic temperatures and energies for HTSs that fall into the asymmetric half-dome-shaped $T_c$ curve. $T_c$($P_{pl}$) depends on the $n_{layer}$, while the $T^*$($P_{pl}$) is independent of it.\\cite{hon04} Therefore, we group the single- and double-layer HTSs with same $T_c^{max}$ $\\sim$ 90 K together. We found in Ref.\\ \\onlinecite{hon04} that the various characteristic temperatures or pseudogap temperatures can be separated into two groups of the lower pseudogap temperature ($T_{lp}^*$) and upper pseudogap temperature ($T_{up}^*$). $T_c$ and major characteristic temperatures, which includes $T_{lp}^*$ and $T_{up}^*$, are plotted on the reduced temperature-scale as a function of $p_u$ in Figs.\\ \\ref{fig6}(a), (c) and (e). The characteristic energies are plotted on a reduced energy-scale $\\mu_c$($E$) $\\equiv$ $E$/$k_BT_c^{max}$, where $k_B$ is Boltzmann's constant, as a function of $p_u$ in Figs.\\ \\ref{fig6}(b), (d) and (f). We will call the four solid curves from the top to bottom the hump temperature ($T_{hump}^*$), $T_{up}^*$, $T_{lp}^*$ and $T_c$ curves in the temperature-scale, and the hump energy ($E_{hump}^*$), the upper pseudogap energy ($E_{up}^*$), the lower pseudogap energy ($E_{lp}^*$) and $\\Delta_c$-curves in the energy-scale, respectively. The $E_{hump}^*$, $T_{up}^*$, $T_{lp}^*$ and $T_c$-curves are directly determined from the plotted data. The $T_{hump}^*$, $E_{up}^*$, $E_{lp}^*$ and $\\Delta_c$ curves are converted from the $E_{hump}^*$, $T_{up}^*$, $T_{lp}^*$ and $T_c$ curves using a relation of $T$ = $E$/$zk_B$ or $E$ = $zk_BT$ for each characteristic energy or temperature. In the energy-scale, the solid curves correspond to $z$ = 3.5 and gray zone shows the energy range from 3$k_BT$ to 4$k_BT$.   First, we summarize the characteristic temperatures and energies derived from the transport and thermodynamic properties in Figs.\\ \\ref{fig6}(a) and (b). Here, the plotted data are summarized in Table~\\ref{tab:table8}. $T^*$'s determined from $TEP$ and $\\gamma$ lie on the $T_{lp}^*$-curves, while $T^*$'s determined from $\\rho$ and QPR lie on the $T_{up}^*$ curves, as reported in Ref.\\ \\onlinecite{hon04}. Accordingly, the upper pseudogap is identified by the $\\rho$ and QPR, and the lower pseudogap is identified by $TEP$ and $\\gamma$ experiments. However, $T^*$ determined from the $\\rho_c$ tends to be higher than $T_{up}^*$, although the doping range is restricted. This may suggest a third pseudogap as already pointed out in Ref.\\ \\onlinecite{hon04}. This suggestion is further supported by the similar behavior derived by other probes in the temperature- and energy-scales. We plot the upper and lower inflection points of $\\rho_a$ of OD-Y123 into the Fig.\\ \\ref{fig6}(a).\\cite{and04} The lower and upper inflection points seem to lie on the $T_{lp}^*$ and $T_{up}^*$ curves, respectively.   For the characteristic energies, we use the data reported in the $\\gamma$,\\cite{low97,coo96,low00} $\\kappa$,\\cite{sut03,haw07,cha00} and QPR experiment.\\cite{kab99,dem99} $E^*$'s determined from the QPR lie on the $E_{up}^*$ urve. $E^*$'s determined from $\\kappa$ show up on either $E_{lp}^*$ or $E_{up}^*$ curve. In the overdoped side, these $E^*$'s clearly merge into the $\\Delta_c$-curve. This indicates that there is no QCP inside the superconducting phase. In Fig.\\ \\ref{fig6}(b), the normal state gap, $E_{sh}^*$(110 K) and zero temperature superconducting gap, $E_{sh}^*$(0 K) determined by the specific heat measurement are plotted as the open symbols and solid circles, respectively.\\cite{low97,coo96,low00} The $E_{sh}^*$(110 K) at $p_u$ $<$ 0.85 and $E_{sh}^*$(0 K) follows the $E_{lp}^*$- or $zk_BT_{lp}^*$-curve. However, $E_{sh}^*$(110 K) at $p_u$ $>$ 0.85 deviates downward from the $E_{lp}^*$-curve, crosses the $\\Delta_c$-curve and finally goes to zero inside the $\\Delta_c$-curve. The temperature of $\\sim$110 K ($\\tau_c$(110 K) = 110/90 $\\sim$ 1.2) corresponds to the $T_{lp}^*$ at $p_u$ $\\sim$ 0.85. The influence of the lower pseudogap on the extraction of $E^*$'s is clearly seen. The plotted $E_{sh}^*$(110 K) is the same data set used to support the existence of the QCP \\textit{\\textbf{inside}} the superconducting phase on the $P_{T_c}$ scale.\\cite{tal01} Accordingly, the existence of the QCP inside the superconducting phase is extrinsic to high $T_c$.   We summarize the characteristic temperatures and energies derived from the spectroscopic measurements in Figs.\\ \\ref{fig6}(c) and (d). The plotted data are summarized in Table~\\ref{tab:table9}. $T^*$ determined in the AIPES lies on $T_{hump}^*$.\\cite{tak01} $T^*$ determined from the ARPES,\\cite{din96,cam99,sat02} and SIS \\cite{dip02} cannot be grouped into either $T_{up}^*$ or $T_{lp}^*$ curve, since they lie between $T_{up}^*$ and $T_{lp}^*$-curves. The $E^*$ determined in the ARPES,\\cite{cam99,sat02,din01,lu01,tan06,pee07,dam03} and tunneling,\\cite{dip02,ren98,miy98,wei98,zas01,ozy99,ozy00} are plotted in Fig.\\ \\ref{fig6}(d): the peak and hump energies observed in ARPES and tunneling lie on the $E_{lp}^*$ and $E_{hump}^*$ curves, respectively. It is clearly seen that there is a third energy scale corresponding to the hump structure observed in the ARPES and tunneling spectroscopy.   We summarize the characteristic temperatures and energies derived from the spin  and charge probes in Figs.\\ \\ref{fig6}(e) and (f), respectively. The plotted data are summarized in Table~\\ref{tab:table9}. For the characteristic temperatures, $T^*$ determined from the ($T_1T$)$^{-1}$ lies on the $T_{lp}^*$. $T^*$ determined from the neutron lies between $T_{up}^*$- and $T_{lp}^*$ curves. $T_{mK}^*$ and $T_K^*$ observed in $K_c$ lie on the $T_{hump}^*$ and $T_{up}^*$, respectively.\\cite{ish98,ito98} For the characteristic energies, the half of the coherent peak energy of $B_{2g}$ ERS,\\cite{ken95,sug03,hew02,gas98,tac06,ven02} half of the coherent peak energy of $B_{1g}$ ERS,\\cite{pai06,blu97,cap07,gas97,mak94,sug00,kan96,ken95,sug03,hew02,gas98,tac06,ven02} and half of the two-magnon peak energy of $B_{1g}$ ERS \\cite{blu97,sug03,mak94} lie on the $\\Delta_c$, $E_{lp}^*$ and $E_{hump}^*$ curves, respectively.   \\begin{figure*}[t] \\includegraphics[scale=0.6]{Figure7.pdf} \\caption{\\label{fig7} (Color online) Electronic phase diagram for the single-layer SrD-La214. The temperature- and energy-scale for the characteristic temperatures and energies are summarized in Figs.\\ \\ref{fig7}(a) and (b), respectively. The plotted data are summarized in Table~\\ref{tab:table10}. } \\end{figure*}    From Figs.\\ \\ref{fig6}(a) $-$\\ \\ref{fig6}(f), we can conclude that the phase diagram fundamentally reproduces the $T$ vs $P_{pl}$ plot in Ref.\\ \\onlinecite{hon04}. Their characteristic temperatures $T^*$ lie on either the $T_{up}^*$ or the $T_{lp}^*$ curve in Ref.\\ \\onlinecite{hon04}. Furthermore, the third characteristic energy, the ``hump'' energy, does exist, although it is hard to detect as the corresponding characteristic temperature or $T_{hump}^*$. All four characteristic temperatures ($T_c$, $T_{lp}^*$, $T_{up}^*$, and $T_{hump}^*$) and the corresponding energies ($\\Delta_c$, $E_{lp}^*$, $E_{up}^*$, and $E_{hump}^*$) do not cross each other. The four temperatures and energies tend to converge with increasing $p_u$, merge at $p_u$ $\\sim$ 1.1, and finally vanish at $p_u$ $\\sim$ 1.3.   Some $T^*$'s ($E^*$'s) have relatively large scattering, and are hard to group into either $T_{up}^*$ ($E_{up}^*$) -curve or $T_{lp}^*$- ($E_{lp}^*$) -curve. For example, the $E^*$ from $\\kappa$ and $T^*$ from ARPES, tunneling spectroscopy and neutron scattering are scattered. These scattering may come from the differences in the characteristic time-scale and length-scale specific to different experimental probes for observing the intrinsically inhomogeneous electronic states, as discussed by Mihailovic and Kabanov.\\cite{mih04} Indeed, similar to $T_{up}^*$- and $T_{lp}^*$-curves or $E_{up}^*$- and $E_{lp}^*$-curves, they all become smaller and closer in the magnitude with increasing doping in the underdoped regime. They merge into $T_c$ or $\\Delta_c$-curve in the overdoped regime and universally vanish at $p_u$ = 1.3. The pseudogap, manifested either as the characteristic energy or characteristic temperature and independent of its origin, universally disappears at $p_u$ $\\sim$ 1.3 together with the superconductivity. This strongly suggests that the pseudogap phase is the precursor of the superconducting phase.   \\begin{table}[b] \\caption{\\label{tab:table10}The data plotted in Fig.\\ \\ref{fig7} for SrD-La214.} \\begin{ruledtabular} \\begin{tabular}{lll} \\quad  Fig. & Probe & Ref(s). \\\\ \\hline \\quad  7(a) & TEP      & \\onlinecite{kak99,nis94,zho95,coo96}  \\qquad \\\\ \\cline{2-3} & $\\chi$   & \\onlinecite{nak94,yos90,joh89,nak98} \\\\ \\cline{2-3} & $\\rho$   & \\onlinecite{nak94,nak98} \\\\ \\cline{2-3} & $d^2\\rho_{ab}/dT^2$ & \\onlinecite{and04} \\\\ \\cline{2-3} & ARPES    & \\onlinecite{sat99,ino98,ino00} \\\\ \\cline{2-3} & AIPES    & \\onlinecite{has07} \\\\ \\cline{2-3} & specific heat & \\onlinecite{lor98} \\\\ \\cline{2-3} & ($T_1T$)$^{-1}$ & \\onlinecite{ito04} \\\\ \\hline \\quad 7(b)  & ARPES    & \\onlinecite{sat99,ino98,ino00} \\\\ \\cline{2-3} & QPR      & \\onlinecite{kus05} \\\\ \\cline{2-3} & $\\gamma$ & \\onlinecite{lor98} \\\\ \\cline{2-3} & $B_{1g}$ ERS & \\onlinecite{sug03} \\\\ \\cline{2-3} & $B_{2g}$ ERS & \\onlinecite{sug03} \\\\ \\end{tabular} \\end{ruledtabular} \\end{table}   \\begin{figure*}[t] \\includegraphics[scale=0.6]{Figure8.pdf} \\caption{\\label{fig8} (Color online) Sketches of (a) the UEPD for HTS with $T_c^{max}$ $\\sim$ 90 K and (b) the phase diagram for the SrD-La214. In Fig.\\ \\ref{fig8}(a), the superconducting (SC) and antiferromagnetic (AF) phases represented by the dotted lines are coming from the OD-Y123.\\cite{hon07,san04} In Fig.\\ \\ref{fig8}(b), the AF phase for the SrD-La214 is cited from Refs.\\ \\onlinecite{nie98} and \\onlinecite{mat02}.} \\end{figure*}    In Fig.\\ \\ref{fig8}(a), we present a sketch of the unified electronic phase diagram for HTSs purely based on experimental grounds. The characteristic features of the unified electronic phase diagram for single- and double-layer HTSs with $T_c^{max}$ $\\sim$ 90 K are: (i) the asymmetric half-dome-shaped $T_c$-curve (SC phase), (ii) there are three characteristic temperatures, $T_{hump}^*$, $T_{up}^*$, and $T_{lp}^*$ in the underdoped region ($p_u$ $<$ 1), (iii) all three characteristic temperatures and $T_c$ come together at $p_u$ = 1.1 in the overdoped region and vanish at $p_u$ = 1.3, (iv) $T_{hump}^*$ changes into the rapid decrease at $p_u$ $\\sim$ 1, (v) $T_{hump}^*$ and $T_{lp}^*$ are concave upward, while the $T_{up}^*$ is concave downward, and (vi) the electronic phase diagram on the temperature-scale can be translated into that on the energy-scale through $E$ = $zk_BT$ with $z$ = 3.5 $\\pm$ 0.5. Although we use HTSs with $T_c^{max}$ $\\sim$ 90 K as our model system, we should emphasize that (i) $-$ (vi) are salient features for all, except of SrD-La214 discussed in Sec. III E, HTSs. We will call this phase diagram the ``unified electronic phase diagram (UEPD)'' of HTS. Finally the $T_{lp}^*$- and $T_{up}^*$-curves tend to merge into the N\\'{e}el temperature ($T_N$) curve with undoping.\\cite{nie98,mat02}   \\subsection{Phase diagram for the SrD-La214}   Now we discuss the HTS with symmetric $T_c$ curve, i.e., the phase diagram of SrD-La214. For the SrD-La214, the characteristic temperatures on the reduced temperature-scale and  characteristic energies on the reduced energy-scale are plotted as a function of $p_u$ in Figs.\\ \\ref{fig7}(a) and (b), respectively. The plotted data are summarized in Table~\\ref{tab:table10}. First, noticed that the characteristic temperatures are separated into not three curves of $T_{lp}^*$, $T_{up}^*$ and $T_{hump}^*$, but two curves. Also in the energy-scale, the characteristic energies are separated into two curves. In the UEPD, the $T_{lp}^*$- or $E_{lp}^*$- curve was defined by $TEP$, the peak structure of ARPES and tunneling, $\\gamma$, and $T_{hump}^*$ or $E_{hump}^*$ curve was defined by the $\\rho_c$, susceptibility, the hump structure of ARPES and tunneling, and two-magnon peak of $B_{1g}$ ERS. In the SrD-La214, the lower curve of the temperature- or energy-scale is identified by $TEP$, the peak structure of ARPES and tunneling, $\\gamma$, and the upper curve is identified by the $\\chi$, the hump structure of ARPES and tunneling, and two-magnon peak of $B_{1g}$ ERS. Accordingly, the lower and upper curves of SrD-La214 are designated to be $T_{lp}^*$ and $T_{hump}^*$ ($E_{lp}^*$ and $E_{hump}^*$), respectively. In the SrD-La214, the $T^*$ or $E^*$ defined by the $\\rho$ lie on either of the two curves, suggesting no $T_{up}^*$- or $E_{up}^*$ curve. Because, in the UEPD, the $T^*$ or $E^*$ defined by the $\\rho$ and $QPR$ lied on $T_{up}^*$ or $E_{up}^*$ curve. In fact, $T_{\\rho}^*$'s for the SrD-La214 with $x$ $<$ 0.16 ($p_u$ $<$ 1) lie on the $T_{hump}^*$-curve,\\cite{nak94} while $T_{\\rho}^*$'s for $x$ $>$ 0.14 ($p_u$ $>$ 0.875) lie on the $T_{lp}^*$ curve.\\cite{nak98b} We also plot the upper and lower inflection points of $\\rho_{ab}$ into the Fig.\\ \\ref{fig7}(a).\\cite{and04} The upper and lower inflection points seem to correspond to the $T_{hump}^*$ and $T_{lp}^*$ curves, respectively. In the SrD-La214, the usual upper pseudogap temperatures identified in UEPD tend to lie on either lower pseudogap temperature or the hump temperature.   In Fig.\\ \\ref{fig8}(b), we present a sketch of the phase diagram for SrD-La214. The $T_c$($p_u$) follows a symmetric dome-shaped $T_c$-curve (SC phase). This is quite different from feature (i) of the UEPD. There are two characteristic temperatures, i.e., $T_{hump}^*$ and $T_{lp}^*$, in the range from the underdoped regime. This is also different from feature (ii) above. $T_{lp}^*$ of SrD-La214 seems to be a combination of $T_{up}^*$ and $T_{lp}^*$ of the UEPD. Although the $T_c$, $T_{hump}^*$ and $T_{lp}^*$ decreases with doping, there is no merging until the end point. This is also different from the feature (iii) above. $T_c$, $T_{hump}^*$, and $T_{lp}^*$ fall down to ($p_u$,$\\tau_c$) = ( 1.6, 0 ) in contrast to ( 1.3, 0 ) in the UEPD. There is a slight change in curvature in the $T_{hump}^*$($p_u$) at $p_u$ $\\sim$ 1. This may share the same origin as the feature (iv) above, although it is much weaker in the SrD-La214 system. Both features (v) and (vi) are similar to that of UEPD.   \\subsection{Comparison between the unified electronic phase diagram and the other phase diagram}   The present UEPD is different from the phase diagrams that were proposed and discussed in Refs.\\ \\onlinecite{tal01,bat96,eme97} and \\onlinecite{ric99}. The phase diagram in Ref.\\ \\onlinecite{tal01} suggests that the single $T^*$ curve crosses the dome-shaped $T_c$ curve or superconducting phase at around the optimal doping level, and fall down to $T$ = 0 at the QCP inside the superconducting phase. This phase diagram implies that there is no correlation between the pseudogap phase and high-$T_c$ phase, and therefore the pseudogap is a pure competing order. The phase diagram of Ref.\\ \\onlinecite{bat96} suggests that the dome-shaped $T_c$ curve intercepts the double $T^*$ curves at around the optimal doping level. The upper and lower $T^*$ curves are concave downward and upward, respectively. These phase diagram is based on the $P_{T_c}$-scale, since $T_c$ follows the superconducting dome. The phase diagram on Ref.\\ \\onlinecite{eme97} shows a tendency that the double $T^*$ curves merge into the asymmetric $T_c$ curve at around the slightly overdoped level and go to zero with $T_c$ at the end point. However, both $T^*$ curves are concave downward. In Refs.\\ \\onlinecite{tal01,bat96,eme97}, the pseudogap does not merge into the $T_N$ curve with undoping. The phase diagram discussed in Ref.\\ \\onlinecite{ric99} shows that the single $T^*$-curve smoothly merges into the $T_N$-curve with undoping, and smoothly merges into the asymmetric $T_c$ curve at the end point with doping. However, the single $T^*$ curve is concave downward. Thus, without alluding to the microscopic picture for the high-$T_c$ mechanism, all the previously proposed phase diagrams are different from our UEPD, except that the asymmetric $T_c$ curve in Refs.\\ \\onlinecite{eme97} and \\onlinecite{ric99} is similar to the present half-dome shaped $T_c$ curve.  The present UEPD clearly appears that pseudogap exists above $T_c$ for $p_u$ $<$ 1.1, while for $p_u$ $>$ 1.1 pseudogap appears at $T_c$. Even experimental data that supported a QCP inside the superconducting dome on the $P_{T_c}$-scale followed the UEPD. The phase diagram for the SrD-La214 shows that pseudogap temperatures and corresponding characteristic energies always exist above the superconducting phase, until the pseudogap disappears together with the superconducting phase at $p_u$ = 1.6. These results indicate that for all the HTSs the pseudogap phase always co-exists with the superconducting phase up to the end point and does not intersect the superconducting phase. Furthermore, the overdoped HTS with superconductivity cannot be regarded as a conventional FL phase, since there is always the pseudogap phase with superconducting phase. Both phase diagrams suggest no QCP inside the superconducting phase. Actually, it has reported that QCP may exist at around the end point of the superconducting phase when superconductivity has completely disappeared in CLBLCO,\\cite{wat06} SrD-La214,\\cite{ris07} OD-Tl2201,\\cite{kru07} and OD-Bi2212.\\cite{kru07}   The UEPD is consistent with the idea that the pseudogap phase is, if not sufficient, necessary for the high $T_c$. It also implies that at least two distinct energy scales, i.e., pseudogap and superconductivity, are required to realize high $T_c$. If we adopt a scenario that superconducting pairing is realized in the pseudogap phase and the global phase coherence occurs at $T_c$, then the smooth merging of $T^*$'s and $T_c$ in the overdoped regime suggests that cuprates become a ``more conventional'' superconductor. Because paring and phase coherence occurs at the same temperature ($T_c$). However since pseudogap still exists, it simply merges with $T_c$ and changes with $T_c$, therefore the superconducting state as well as the normal state are still ``unconventional'' up to the end point as reported in some studies.\\cite{haw07,pee07} This also explains why the pseudogap phase was never observed in the overdoped regime except SrD-La214. Even in the SrD-La214, the observation of the pseudogap in the overdoped regime strongly depends on the experimental probe. For example, it is not observed in the resistivity measurements but can be clearly seen by magnetic susceptibility and TEP measurements, as shown in Fig.\\ \\ref{fig7}.   In the previous paper, we pointed out that the observed $T_{lp}^*$ and $T_{up}^*$ are coming from not one pseudogap, but two pseudogaps,\\cite{hon04} because the temperature where the TEP has the broad peak, corresponding to $T_{lp}^*$, was different from the temperature where the TEP starts to depend on the Zn-doping, corresponding to $T_{up}^*$.\\cite{hon04} However, according to the idea by Mihailovic and Kabanov,\\cite{mih04} we can not completely rule out a possibility that three characteristic temperatures, including $T_{hump}^*$, are of the same physical origin. The different characteristic temperatures may come from the differences in the characteristic time-scale and length-scale specific to different experimental probes for observing the intrinsically inhomogeneous electronic states or pseudogap phase."
    },
    "0801/0801.4124_arXiv.txt": {
        "abstract": "We present high precision, time-resolved visible and near infrared photometry of the large (diameter $\\sim$ 2500 km) Kuiper belt object (136108) 2003$\\,$EL$_{61}$.  The new data confirm rapid rotation at period $P$ = 3.9155$\\pm$0.0001$\\,$hr with a peak-to-peak photometric range $\\Delta m_R$ = 0.29$\\pm$0.02$\\,$mag.\\ and further show subtle but reproducible color variations with rotation.  Rotational deformation of 2003$\\,$EL$_{61}$ alone would give rise to a symmetric lightcurve free of color variations. The observed photometric deviations from the best-fit equilibrium model show the existence of a large surface region with an albedo and color different from the mean surface of 2003$\\,$EL$_{61}$.  We explore constraints on the nature of this anomalous region set by the existing data. ",
        "introduction": "The known Kuiper belt objects (KBOs) extend in size from bodies so small as to be at the limits of sensitivity of the largest telescopes up to Pluto-sized bodies large enough to have body shapes controlled by self-gravity.  The large objects are particularly amenable to physical study and a number of intriguing results have already been secured.  The case of Pluto is well known: the main object rotates slowly (period $\\sim$6 days) but the massive satellite Charon carries enough angular momentum that the system as a whole is near the critical threshold for breakup \\citep{1989ApJ...344L..41M}. Surface maps derived from mutual occultation events show a spatially variegated surface, with a wide range of local albedos \\citep[from 0.1 to 0.9:][]{1992Icar...97..211B,1999AJ....117.1063Y} that may be related to surface deposition of frosts from Pluto's thin atmosphere \\citep{1989GeoRL..16.1213T}. Other KBOs have less well known physical properties but new data are beginning to reveal a startling range of surface types \\citep{2004Natur.432..731J,2007AJ....133..526T,2007ApJ...655.1172T} and rotational \\citep{2006AJ....131.2314L,2007AJ....134..787S} properties. Notable examples of the latter include $\\sim$900 km diameter (20000) Varuna, whose 6$\\,$hr period and 0.4$\\,$mag.\\ photometric range are best explained as products of a rotationally deformed body shape and a bulk density of 1000\\kgm\\ \\citep{2002AJ....123.2110J,2004PASJ...56.1099T,2007AJ....133.1393L}. The large amplitude (1.14$\\pm$0.04 mag.), long rotation period (13.7744$\\pm$0.0004$\\,$hr) and eclipsing binary-like lightcurve of $\\sim$240$\\,$km diameter 2001$\\,$QG$_{298}$ suggest an even more extreme interpretation as a contact or near-contact binary \\citep{2004AJ....127.3023S,2007AJ....133.1393L}.  One of the most remarkable of the large KBOs yet to be identified is (136108) 2003$\\,$EL$_{61}$ (hereafter ``EL61''), whose rapid rotation (period$\\sim$3.9154$\\pm$0.0002$\\,$hr), and lightcurve range ($\\Delta m_R \\sim 0.4\\,$mag) and near-symmetric morphology together suggest a rotationally deformed body of density $\\sim$2500\\kgm\\ \\citep{2006ApJ...639.1238R,2007AJ....133.1393L}. EL61 is further interesting in its own right, as an extreme example of a large KBO with a rapid spin and also as the possible parent of a reported dynamical cluster of KBOs, perhaps produced by an ancient, shattering collision (Brown et al.  2007, Ragozzine and Brown 2007). Some members of this dynamical cluster share spectral features with EL61. Nearly all Pluto-sized KBOs have methane rich surfaces. EL61 is unusual in that it is covered in almost pure crystalline H$_2$O ice (see Fig.~\\ref{Fig.Spectrum}; Trujillo et al. 2007).  In this paper, we present new high-precision, time-resolved photometry taken to further explore the nature of EL61.  \\begin{deluxetable*}{cccccccc} \\tablecaption{Journal of Observations. \\label{Table.Journal}} \\tablewidth{0pt} \\tablehead{ \\colhead{UT Date} & \\colhead{$R$\\tna} & \\colhead{$\\Delta$\\tnb} & \\colhead{$\\alpha$\\tnc} & \\colhead{Tel.\\tnd} & \\colhead{Filt.\\tne} & \\colhead{Seeing\\tnf} & \\colhead{Exp. Time\\tng} \\\\ \\colhead{} & \\colhead{$[AU]$} & \\colhead{$[AU]$} & \\colhead{[\\degr]} & \\colhead{} & \\colhead{} & \\colhead{[\\arcsec]} & \\colhead{[s]} } \\startdata 2007 Jun 11 & 51.1570 & 50.8037 & 1.07 &UH2.2m & $R$      & 0.9 & 80 \\\\ 2007 Jun 13 & 51.1567 & 50.8296 & 1.08 &UH2.2m & $R$      & 1.0 & 80 \\\\ 2007 Jun 15 & 51.1565 & 50.8576 & 1.09 &UH2.2m & $B$      & 0.9 & 260 \\\\ 2007 Jul 07 & 51.1536 & 51.1800 & 1.14 &UKIRT  & $J$      & 1.0 & 60 \\\\ 2007 Jul 08 & 51.1535 & 51.1949 & 1.14 &UKIRT  & $J$      & 1.2 & 60 \\\\ 2007 Jul 22 & 51.1517 & 51.4001 & 1.10 &UH2.2m & $B$      & 1.5 & 300 \\\\ 2007 Jul 24 & 51.1514 & 51.4286 & 1.09 &UH2.2m & $B$, $R$ & 1.5 & 260, 80 \\\\ \\enddata \\tablenotetext{a}{Heliocentric distance in AU;} \\tablenotetext{b}{Geocentric distance in AU;} \\tablenotetext{c}{Phase angle in degrees;} \\tablenotetext{d}{Telescope used;} \\tablenotetext{e}{Filters used;} \\tablenotetext{f}{Typical seeing in arcseconds;} \\tablenotetext{g}{Typical integration time per frame in seconds.} \\end{deluxetable*} ",
        "conclusions": "Photometric measurements were obtained first relative to field stars, to provide protection from transient changes in the transparency of the Earth's atmosphere and variations in the seeing (which can impact the accuracy of photometry obtained through discrete apertures).  The resulting measurements were calibrated against standard stars using large aperture photometry. Internal accuracy of the lightcurve data in the $B$ and $R$ filters is good to about $\\pm$0.01$\\,$mag.\\ while, in the infrared, scatter in the photometry shows that the accuracy is at the $\\pm$0.03$\\,$mag.\\ level.  We did not apply a phase angle correction to the data.  The phase angle changed from 1.07 to 1.10 degrees in our $B$ and $R$ observations.  In this 0.03 degree phase angle range, with a phase coefficient of $\\sim$0.1 mag/deg \\citep{2007AJ....133...26R}, the effect of phase is only 0.003$\\,$mag.\\ and therefore unimportant compared to the 0.01$\\,$mag.\\ photometric accuracy.  Furthermore, the color dependence of the phase coefficient is small for EL61 according to these authors, and the expected change in the color resulting from phase angle is only about 0.001 mag, which is again negligible.  The best-fit lightcurve period was determined from the $R$-band data using phase-dispersion minimization (PDM; Stellingwerf 1978) as $P$ = 3.9155$\\pm$0.0001$\\,$hr (two-peaked lightcurve).  This period is in close agreement with $P$ = 3.9154$\\pm$0.0002$\\,$hr determined independently \\citep{2006ApJ...639.1238R}. Photometry in the other filters was scaled to the $R$-band lightcurve by subtracting the median colors $B-R=0.972$ and $R-J=0.88$, as determined from our data.  The $B-R$ color is again in good agreement with $B-R$ = 0.969$\\pm$0.030 reported by Rabinowitz et al; these authors did not measure $R-J$.  \\FigBRJLightcurve  \\FigBRColorCurve  The resulting phased $B$-, $R$- and $J$-band lightcurves of EL61 are shown in Fig.~\\ref{Fig.BRJLC}. Two main features are immediately apparent from the lightcurves. Firstly, the two peaks of the lightcurve are unequal.  The total range (peak-to-peak) is 0.29$\\pm$0.02$\\,$mag.\\ but the second peak is smaller by roughly 0.08$\\,$mag.  This asymmetry in the lightcurve peaks cannot be matched by simple equilibrium shape models of the type proposed by \\citet{2006ApJ...639.1238R}, since the latter are symmetric \\citep{1969efe..book.....C}. Secondly, we note that, in the interval roughly from 0.7 to 1.0 in rotational phase (Fig.~\\ref{Fig.BRJLC}) the $B$ data fall systematically below the $R$ data. Although small, this effect appears in measurements from three different nights and hence we regard it as observationally secure.  The $B-R$ color curve, computed from the data in Fig.~\\ref{Fig.BRJLC}, is shown separately in Fig.~\\ref{Fig.BRColorCurve}.  There, the $R$ magnitudes were interpolated to the rotational phases of $B$ photometry and were subtracted from the $B$ data points. The resulting color curve was smoothed using a running median filter to show a reddening of up to 0.035 magnitudes. The $J$ data are of lower signal-to-noise but, in the region near 0.75 rotational phase, $R-J$ is also redder than near the 0.25 rotational phase peak (Fig.~\\ref{Fig.BRJLC}).  We quantitatively assess the significance of the red feature in the $B-R$ color curve by noting that, in the interval from 0.7 to 1.0 in rotational phase, 21 of the 23 consecutive phased measurements fall above the median $B-R$ for EL61. The probability of this result is the same as the probability of obtaining at least 21 ``tails'' in 23 tosses of an unbiased coin. Assuming a binomial distribution, this probability is roughly $p=3.3\\times10^{-5}$, corresponding to $\\sim$4$\\sigma$.  Furthermore, at least 9 measurements in that same interval lie $>3\\sigma$ above the median, corresponding to $\\sqrt{9}\\times 3\\sigma=9\\sigma$.  In this sense, the redder region in Fig.~\\ref{Fig.BRColorCurve} is unlikely to be due to chance.  This, plus the fact that the red spot is confirmed by observations on different nights together strongly suggest that the feature is real.  Lightcurve asymmetry of the type observed in Fig.~\\ref{Fig.BRJLC} could be caused by strength-supported topography. However, the associated color variations (Fig.~\\ref{Fig.BRColorCurve}) cannot be so explained.  Instead, the data are best explained by the presence of wavelength-dependent albedo markings on the surface of EL61, perhaps analogous to the ones already mapped on Pluto.  Specifically,  given that the body shape of EL61 is close to a figure of equilibrium, the multi-wavelength lightcurve data show the existence of a region, near the second peak in Fig.~\\ref{Fig.JacSpotCurves}, that is darker and redder than elsewhere.  For want of a better label, we refer to this as the ``dark red spot'' (DRS).  \\FigSpotCurves  The time-resolved measurements of the 1.5$\\,\\mu$m water-ice band from UT 2007 July 07 are plotted in Fig.~\\ref{Fig.WaterBandDepth}.  The data provide no compelling indication of variability, except that the ratio of the 1.6$\\,\\mu$m flux density to the continuum flux density at 1.25$\\,\\mu$m appears lower (the water-ice band deepens) near phase $\\sim$0 than at other phases.  An attempt to repeat the 1.6$\\,\\mu$m photometry on UT 2007 July 08 was thwarted by unstable atmospheric opacity.  In the absence of confirming data from a second night, we regard the variation seen in Fig.~\\ref{Fig.WaterBandDepth} as interesting but inconclusive.  We cannot determine whether the water-band depth varies with the rotation of EL61.  What constraints can be set on the albedo markings present on the surface of EL61?  We first address the spatial extent of the DRS.  In principle, the DRS could be very small compared to the instantaneous projected cross-section of EL61 but would then need to be very red and very dark relative to the surroundings in order to give rise to the observed lightcurve differences.  At the other extreme, the DRS could be large, possibly even hemispheric in extent, in which case its albedo and color contrast relative to the surroundings would be minimal (see Fig.~\\ref{Fig.SpotModels}).  To explore the range of possibilities we computed models in which the area of the surface of EL61 occupied by the DRS was taken as a free parameter.  The models we used are described in detail in \\citet{2007AJ....133.1393L}. In short, we render 3-dimensional models of EL61 at different rotational phases which are used to generate the synthetic lightcurve. In this paper we adopt a Lambert scattering law, appropriate for high-albedo icy surfaces. The spot was simulated as a region of different reflectivity and color curves were generated by subtracting the lightcurves of two spots of equal sizes but different reflectivities. The shape of EL61 was modeled by a Jacobi ellipsoid with axis ratios $b/a=0.87$ and $c/a=0.54$, which provides the best match to the lightcurve data if no albedo variegation is present (see Fig.~\\ref{Fig.JacSpotCurves}).  The size of the spot is parametrized by its sky-plane cross-section area relative to the maximum cross-section of the Jacobi ellipsoid, $\\pi ac$. We assumed that the rotation axis of EL61 was inclined relative to the line-of-sight by 90$^{\\circ}$, consistent with the large measured rotational lightcurve range,  and that the DRS is located on the equator of EL61. The observed sequence of brighter and fainter extrema indicates that longitude of the DRS must lie in a leading quadrant with respect to one of the semi-major axes. This prediction is corroborated by our models, which further show that a longitudinal separation of 45\\degr between the DRS and the long axis of EL61 produces a better match to the data than 30\\degr or 60\\degr. The ability to fit the shape of the lightcurves in different filters was used as a metric for the models.  Three of the best-fit examples are shown in Figs.~\\ref{Fig.JacSpotCurves} and \\ref{Fig.SpotModels}.  \\FigWaterBandDepth  \\FigSpotModels  The results, which confirm the qualitative expectations outlined above, are shown in Fig.~\\ref{Fig.AlbedoAndColorVsArea}. Figure~\\ref{Fig.ColorVsAlbedo} combines the color and albedo constraints and allows comparison with real surfaces. Also marked in Fig.~\\ref{Fig.ColorVsAlbedo} are the ranges of color and albedo for established outer Solar system materials, including the dark regions on Pluto and Saturn's satellite Iapetus.  The EL61 data are incompatible with very small patches of dark, red matter like that found on the low-albedo side of Iapetus, or even with the darker material on the surface of Pluto.  Indeed, if the spot is to have a $B-R$ color within the range observed for Solar system objects ($B-R\\lesssim2$) then it must be larger than $\\sim$20\\% of the maximum cross-section of EL61 (see Fig.~\\ref{Fig.AlbedoAndColorVsArea}).  Instead, Iapetus' and Pluto's bright areas match the DRS in term of albedo and $B-R$ color. The surfaces of Eris (a large KBO) and 2005 FY9 are also consistent with the DRS, even if these objects have highly uncertain albedos \\citep{2007astro.ph..2538S}. All matching surface types would imply a DRS size of 35\\% to 50\\%.  Another possibility is that the DRS is simply terrain contaminated by dirt. This would account for both the darkening and the reddening, but the suspected deepening of the 1.5$\\,\\mu$m water band close to the DRS in rotational phase (see Fig.~\\ref{Fig.WaterBandDepth}) would be harder to explain; a weaker, i.e. less deep, water feature would be expected if that were the case. Alternatively, the DRS could be a region depleted in a spectrally neutral substance, both brighter and bluer than water ice.  A more contrived explanation involves the presence of larger water-ice grains on the DRS which would lower the albedo and redden the surface \\citep{1982Icar...49..244C}, and produce deeper water-ice absorption bands \\citep{2005LPI....36.2061S}. On Enceladus, larger grains are found on the region often referred to as the tiger stripes, where cryovolcanism is thought to be happening \\citep{2007LPI....38.1747S}.  \\FigAlbedoAndColorVsArea  \\FigColorVsAlbedo  What might be the origin of the DRS?  On Pluto, the light and dark albedo markings may be self-sustaining and caused by the mobility of surface volatiles, partly driven by the seasons (Hansen and Paige 1996).  There, dark regions are heated by the Sun leading to higher sublimation rates and the migration of volatiles towards brighter, cooler regions.  In this way the volatile ices may naturally migrate to restricted regions of the surface.  The dominant volatile species on Pluto is the highly volatile solid nitrogen, N$_2$, with methane (CH$_4$) mixed-in as an optically active tracer.  In contrast, the surface of EL61 appears to be water-ice dominated, with no evidence for the diagnostic $N_2$ band at 2.15 $\\mu$m (Fig. \\ref{Fig.Spectrum}).  Water ice is utterly refractory at the $\\sim$30 K temperatures on EL61, and this albedo instability mechanism seems unlikely to apply.   It has been suggested that EL61 is the source of an impact-produced dynamical family of water-rich KBOs.  It is tempting to speculate that the DRS could mark the scar of the impact from which the family members were purportedly excavated, although such an explanation could hardly be unique."
    },
    "0801/0801.4638_arXiv.txt": {
        "abstract": "We discuss short wavelength (inertial wave) instabilities present in the standard two-fluid neutron star model when there is sufficient relative flow along the superfluid neutron vortex array. We demonstrate that these instabilities may be triggered in precessing neutron stars, since the angular velocity vectors of the neutron and proton fluids are misaligned during precession. The presence of such an instability would render the standard, solid body rotation, model for free precession inconsistent. Our results suggest that the standard (Eulerian) slow precession that results for weak drag between the vortices and the charged fluid (protons and electrons) is not seriously constrained by the existence of the instability. In contrast, the fast  precession, which results when vortices are strongly coupled to the charged component, is generally unstable. This implies that fast precession may not be realised in astrophysical systems. ",
        "introduction": "\\label{sec:intro}  Over the last few years, the modelling of neutron star free precession has attracted considerable attention from theorists.  This has been caused in large part by the reported observations of free precession in PSR B1828-11 by \\citet{1828}. There exist several other precession candidates, most notably PSR B1642-03 \\citep{shabanov}, but none are as convincing as PSR B1828-11, whose pulse variations extend over multiple precession periods and are visible in multiple channels (pulse time-of-arrival and beam shape).  This precessional motion is of great interest in several regards. Firstly, PSR B1828-11 is the only clear precession candidate out of the several thousand known pulsars, which begs the questions why is precession so rare, and what has set PSR B1828-11 into precession in the first place?  Secondly, the free precession timescale as extracted from the observations is approximately 500 days (modulo a possible factor of two up or down, depending on which of the observed harmonics in the frequency spectrum is the fundamental one).  This is extremely uncomfortable from the theoretical point of view, as the canonical model of a neutron star containing one or more superfluid components would suggest a much shorter precession timescale. Instead, this long period free precession period seems to be consistent with the simple wobble of a star deformed away from sphericity by strains in its crust or by a magnetic field, with superfluidity playing no role whatsoever \\citep{akgun}.  It is this second question regarding the apparent lack of the effects of superfluidity with which we will mainly be concerned in this paper.  The basic picture is as follows.  Both on theoretical grounds, and in order to explain the glitch phenomena, neutron stars are believed to contain superfluid components whose rotation is achieved by their forming arrays of vortices pointing along the rotation axis.  In a pioneering paper \\citet{shaham77} showed that in the limit where these vortices attach or `pin' themselves tightly to the radio-emitting crustal phase (as is required to explain the glitches), the classical free precession of the star is greatly modified by the gyroscopic effect of the superfluid, whose angular momentum vector would be fixed in the crust's frame.  The resulting free precession period is then approximately $P/ (I_s/I_1 + \\epsilon)$, where $P$ is the spin period, $I_s$ and $I_1$ the moments of inertia of the superfluid and crustal phases, and $\\epsilon$ the dimensionless non-sphericity in the crust's moment of inertia tensor.  Note that by `crust' we really mean the crystalline crust itself and any other part of the star coupled to it on sufficiently short timescale to participate in its precession, which probably includes the core proton fluid.  One region where neutron vortices are expected to interact strongly with the crust is the \\emph{inner crust}, where neutron and crustal matter coexist.  The interaction here is expected to be sufficiently strong that the vortices pin tightly to the crustal lattice, resulting in a very short free precession period, in conflict with the PSR 1828-11 observations.  A possible resolution to this problem was suggested by \\citet{cutler}, who argued that forces created by the precessional motion itself would un-pin the vortices, restoring long period free precession, consistent with the observations.  More problematically, neutron vortices in the \\emph{core} of the neutron star are also expected to interact with the charged component. Specifically, the interior magnetic field is composed of a large number of fluxtubes (a consequence of the protons being condensed in a type-II superconductor), which can interact with the vortices via a process known as \\emph{mutual friction}.  As was shown in detail by \\citet{swc}, even if this coupling force isn't strong enough to completely pin the vortices, the free precession timescale in the coupled system is likely to be much shorter than the observed precession timescale.  The significance of this result was emphasised recently by \\citet{link03}, who argued that the observed slow precession necessitated a profound change in our view of neutron star interiors, requiring either that one or other (or both) of the posited superfluid/superconducting phases do not exist, or else that there are no regions in the interior where they exist simultaneously, perhaps because the protons form a type-I superconductor.  It is this line of argument we address in this paper.  The results outlined above, leading to the prediction of fast precession, have all relied implicitly on treating the superfluid and non-superfluid phases as rigid bodies, each with a well defined angular velocity vector. However, in reality the phases are better described by the fluid Euler equations, the ansatz of rigidity having been made mainly for reasons of tractability.  In this paper, which expands on initial results discussed by \\cite{preclett}, we demonstrate that the relative motion between the two phases that rigid body free precession generates can give rise to a superfluid two-stream instability.  The instability itself occurs in the inertial modes of the rotating fluids, and is related to the experimentally observed instability of Helium~II explained long ago by Glaberson and collaborators \\citep{glab}. It is also intimately linked to the neutron star instability described very recently by \\citet{sidery07}.  As we will argue in detail, this instability renders the free precession rigid rotation ansatz inconsistent and makes all subsequent conclusions concerning the interior composition of neutron stars unsafe.  The plan of this paper is as follows.  In Section 2 we describe the basic formulation of the two-fluid coupled system. Section 3 provides  the plane-wave analysis and the instability criteria for inertial modes in various regimes.  In Section 4 we discuss the relevance of these results for freely precessing neutron stars. The detailed precession modes are analysed in Appendix~A drawing heavily on the previous work of \\citet{swc}.  Our conclusions can be found in Section 5. ",
        "conclusions": "Within the standard two-fluid model for superfluid neutron stars we have discussed a new class of instabilities that set in through short wavelength inertial waves, provided there is sufficient relative flow along the neutron vortices. As a demonstration of the astrophysical relevance of these instabilities we have considered freely precessing neutron stars, in which the two angular velocity vectors are naturally misaligned. Our analysis is based on the usual form for the superfluid mutual friction, which depends on the strength of the drag that determines the friction between neutron vortices and the proton/electron fluid. Our results show that precession may trigger instabilities both in the weak- and strong drag regimes. Based on the analogy with the so-called Donnelly-Glaberson instability in superfluid Helium \\citep{glab,turbook}, we argue that this instability is likely to lead to the formation of vortex tangles and generate turbulence in superfluid neutron star interiors.   The standard model \\citep{swc} predicts two possible precession regimes. Weak drag leads to slowly damped long period precession, while strong drag implies that precession should be fast. Assessing the stability of each precession mode we conclude that slow precession is generally stable. In particular, this should be true for the few known candidate precessors, like PSR B1828-11. An interesting exception could be slow precession of millisecond-period neutron stars. These may become unstable above a wobble angle $\\theta_w \\sim 1.5^\\circ $ assuming a crust deformation $\\epsilon \\sim 5 \\times 10^{-7} $, well within the theoretically allowed range \\citep{cutler,haskell}.  In contrast, we find the fast precession solution to be unstable under  generic conditions. This suggests that the solid-body rotation assumption that leads to the precession mode \\citep{swc} is inconsistent. The upshot of this could be that fast precession is not realised in an astrophysical system. It should certainly be the case that one cannot ignore the small scale hydrodynamical degrees of freedom in the neutron star core in an analysis of the precession problem. This is an interesting conclusion since it adds significant complexity to the modelling. It also calls into question some well established ideas.  Since the seminal discussion of \\citet{shaham77}, it has been known that precession is sensitive to the neutron vortex dynamics. This is natural since the vortices provide the bulk of the star's rotation. If the vortices interact strongly with the charged component (e.g. the crust) then the gyroscopic nature of the vortex array will inevitably lead to fast precession. This corresponding motion can either be described in terms of a ``pinned'' component affecting the usual Eulerian precession \\citep{dij} or as a strong drag fast precession solution in the framework of \\citet{swc}. The observed long period precession of PSR B1828-11 is thus consistent only with weak mutual friction.  However, one can argue that the interaction between neutron vortices and the (much more numerous) magnetic fluxtubes in a type~II proton superconductor ought to lead to strong drag \\citep{sauls,ruderman}. If this leads to the core vortices being effectively pinned to the charged component, then fast precession should result (as in the Shaham model). If on the other hand, the precession motion is able to force the vortices through the fluxtubes then the motion should be highly dissipative and there would be no lasting precession. This led \\citet{link03} to the suggestion that the apparent conflict with the observations can be avoided if the protons instead form a type~I superconductor. There would then be no fluxtubes and slow precession may be possible. However, this would not accord with the generally held view that the protons in the outer neutron star core will condense into a type~II superconductor.  Our results add yet another twist to this story. We have demonstrated that the fast precession solution, that would result if the vortices and the type~II fluxtubes interact strongly, is generically unstable. Although we cannot claim to understand the dynamical repercussions of this, it seems reasonable to assume that the fast solid-body precession motion of the \\citet{swc} model will be significantly affected. If the instability leads to a state of superfluid turbulence, as the analogy with the Helium problem suggests, then it could well be that the precession does not survive. To explore this, a meaningful analysis must venture beyond the rigid-body precession model and account for detailed superfluid hydrodynamics. This may be prohibitively difficult. Anyway, it is clear that we cannot yet use the precession observations to draw any reliable conclusions concerning the state of the proton superconductor.  The present work can (and needs to) be extended in a number of ways. Our uniform density neutron star model is too simplistic and one would certainly want to account for interior stratification etcetera. If one wants to consider superfluid-superconducting mixtures seriously then the two-fluid model probably needs to be upgraded into a more realistic setup with three fluids (after relaxing the assumption of comoving charged particles), plus the magnetic field. The dynamics of such a model for neutron star cores is largely unexplored \\citep{mendell98}, yet the additional degrees of freedom could lead to new features. Our calculation was perturbative. Hence, little can be said about the non-linear behaviour of the unstable waves. Most likely this will be complicated, due to the small scales involved. Numerical simulations of this problem may be challenging, but should be encouraged. In fact they may be required if we are to understand the fate of a precessing neutron star in the strong-drag regime. Some of the necessary tools have already been developed by \\citet{peralta05,peralta06}, and there is analogous work for superfluid Helium \\citep{turbook}. Finally, at the level of linearised theory the present two-fluid analysis could be modified to consider the inner layers of the crust, where superfluid neutrons coexist with an nuclear lattice. One may expect the nature of any superfluid instabilities in that region to be different.  These are all interesting problems that we expect to return to in the future."
    },
    "0801/0801.2534_arXiv.txt": {
        "abstract": "We present a new analysis of the diffuse gas in the Magellanic Bridge (RA\\,$\\ga 3^{\\rm h}$) based on {\\em HST}/STIS E140M and {\\fuse}\\ spectra of 2 early-type stars lying within the Bridge and a QSO behind it. We derive the column densities of \\hi\\ (from \\lya), \\nni, \\oi, \\ari, \\siii, \\sii, and \\feii\\ of the gas in the Bridge. Using the atomic species, we determine the first gas-phase metallicity of the Magellanic Bridge, $[{\\rm Z/H}] = -1.02 \\pm 0.07$ toward one sightline, and $-1.7 < [{\\rm Z/H}] < -0.9$ toward the other one, a factor 2 or more smaller than the present-day SMC metallicity. Using the metallicity and $N($\\hi$)$, we show that the Bridge gas along our three lines of sight is $\\sim$70--90\\% ionized, despite high \\hi\\ columns, $\\log N($\\hi$) \\simeq 19.6$--20.1. Possible sources for the ongoing ionization are certainly the hot stars within the Bridge, hot gas (revealed by \\ovi\\ absorption), and leaking photons from the SMC and LMC. From the analysis of \\cii*, we deduce that the overall density of the Bridge must be low ($<0.03$--0.1 cm$^{-3}$). We argue that our findings combined with other recent observational results should motivate new models of the evolution of the SMC-LMC-Galaxy system. ",
        "introduction": "The evolution of galaxies is closely coupled to the interactions between them. Encounters between galaxies are frequent and were certainly habitual in the early epoch of galaxy formation \\citep[e.g.,][]{barnes92,larson78,maller06}. These encounters reshape the galaxies and transfer mass, energy, metals between the galaxies, and may even create new sites of star formation within the newly produced gaseous features between the galaxies or new burst of star-formations in the colliding galaxies \\citep{larson78,barton07,demello07}. These galactic interactions are not believed to be a major contributor to the pollution of the intergalactic medium \\citep[e.g.,][]{aguirre01}, but they nonetheless affect the metal content and physics of the galaxies themselves and their halos, and therefore the evolution of the galaxies and their surroundings.  In our galactic neighborhood, interactions between the Magellanic clouds and the Galaxy are believed to have produced several large gaseous features: The Magellanic Bridge linking the Small Magellanic Cloud (SMC) and the Large Magellanic Cloud (LMC), the Magellanic Stream and the Leading Arm apparenly linking the Clouds and our Galaxy \\citep[e.g.,][]{mathewson74,putman00,bruns05}. The Magellanic Bridge is believed to have been produced through an interaction between the SMC and LMC, while the Stream and Leading arm may have arisen through an interaction between the Galaxy and the Clouds \\citep[e.g.][]{gardiner96}, although the origin of these latter gaseous features is far from settled \\citep{nidever07,besla07}. The proximity of these gaseous features and the global view of the Magellanic System with little line-of-sight confusion allow us to study the result of interacting galaxies in a way not otherwise possible. The Magellanic Bridge (hereafter Bridge) is the focus of the present work.  In the past few years, multi-wavelength observations of the Bridge have revealed some key characteristics. Radio \\hi\\ observations have shown that approximately 2/3 of the \\ion{H}{1} gas surrounding the Clouds is found in the Magellanic Bridge, while only 25\\% and 6\\% are found in the Stream and the Leading Arm, respectively \\citep{bruns05}.  Some of the Bridge gas appears even to build-up and feed the Stream via an Interface Region \\citep{bruns05}. Far-ultraviolet  observations have shown multiple gas phases in the Bridge including not only the well known neutral gas, but also a significant amount of ionized gas and a small H$_2$ content \\citep[][this paper]{lehner01,lehner02}. Dense molecular clouds were also subsequently identified in form of CO \\citep{muller03,mizuno06} (although these CO surveys were realized in the SMC Wing, a much denser region in stars and gas content of the Bridge). Optical studies have revealed  massive hot stars throughout the Bridge \\citep{demers98}. The ages of these stars are $\\sim$10 to 40 Myr, implying that star formation is ongoing within the Bridge gas since these stars could not migrate from the SMC during their lifetimes, although \\citet{harris07} suggests that star formation in the Bridge may  have been more important 200--300 Myr ago.  The N-body numerical  simulation of the Galaxy-SMC-LMC system by \\citet{gardiner96} have reproduced some of the observed characterics of the Bridge (principally the \\hi\\ column density and velocity distributions), along with those of the Stream. In this paper, we still use their numerical simulation as a basis, although we note that the hypotheses and validity of this model and other recent numerical simulations \\citep[e.g.,][]{yoshizawa03} has been put recently in question in view of the new {\\em Hubble Space Telescope (HST)}\\ measurements of the proper motions of the LMC and SMC \\citep[][and see \\S\\ref{sec-per}]{kallivayalil06,kallivayalil06b,besla07}. Gardiner \\& Noguchi's model  predicts that the Bridge was formed from tidally-stripped gas pulled from the SMC during a close encounter between the SMC and the LMC some 200 Myr ago. The new calculations of the LMC and SMC orbits still suggest that the closest approach between these two galaxies occured some 200 Myr ago \\citep{kallivayalil06}, and therefore the new proper motions may only affect the interpretation of the origins of the Stream and Leading Arm. Nevertheless, there appear to be some challenges to these models.  Tidal models predict that both gas and stars are pulled from the SMC, and yet search for an old stellar population in the Bridge has so far failed \\citep{harris07}. Furthermore, material pulled from the SMC should have a somewhat similar metallicity that the SMC some 200 Myr ago. But this appears in conflict with the B-type star abundances in the Bridge that imply an extremely low Bridge metallicity, $-1.1$ dex from solar \\citep{rolleston99,lee05}, a factor of 3 and 5 metal deficiency with respect to the SMC and LMC, respectively. The metallicity in the sparse regions of the Bridge is also at odds with those measured in the SMC wing, also known as the Western end of the Bridge, where the metallicity of B-type stars is found to be SMC-like \\citep{lee05}. These metallicity estimates complicate the origin of the Bridge and suggest different evolutions or origins for the sparse and dense stellar regions of the Bridge.  An estimate of the present-day metallicy of the Bridge that is independent of stellar measurements is therefore critical. In this work, we present an estimate of the absolute gas-phase abundances of the Bridge toward two stars situated in relatively low \\hi\\ column density regions. We compare the column densities of the neutral species (\\nni, \\oi, \\ari) observed in absorption in the Space Telescope Imaging Spectrograph (STIS) and {\\em Far Ultraviolet Spectroscopic Explorer} ({\\fuse}) spectra to those for \\hi\\ Ly$\\alpha$ lines toward two early-type stars, i.e. our metallicity estimate is mostly independent of ionization correction.  The Bridge Ly$\\alpha$ absorption is heavily blended with the Galactic component, but we show that it is feasible to retrieve the Bridge \\hi\\ column density. Using the metallicity and the amount of \\hi\\ along these sightlines, we can for the first time quantify the amount of ionized gas in the Bridge; ionized gas is revealed to be a major component of the Bridge, yet this gas-phase is mostly ignored in numerical simulations.   Our paper is organized as follows. In \\S\\ref{sec-obs} we briefly discuss the previous STIS and {\\fuse}\\ observations, and the new {\\fuse}\\ observations of DGIK\\,975 and DI\\,1388. In \\S\\ref{sec-analysis} we discuss our measurements of the metals and the \\hi\\ colum densities. We present in \\S\\ref{sec-result} the first metallicity estimate of the Bridge gas and the fraction of neutral, ionized, and molecular gas. From the \\cii*\\ diagnostic, we estimate the electron density and cooling rate in the Bridge. In \\S\\ref{sec-discuss} we discuss our results (in particular address the possible origins for the low metallicity and ionization sources of Bridge) and the current observational challenges to numerical simulations. Finally, in \\S\\ref{sec-sum} we summarize our key findings. ",
        "conclusions": "\\label{sec-discuss} In Table~\\ref{t-prop}, we summarize the properties of the Bridge derived from this work. The most recent estimate of the \\hi\\ mass of the Bridge was derived by \\citet{bruns05}, and to estimate the mass of H of the Bridge we assume that the large ionization fraction in the Bridge observed in our 3 lines of sight is characteristic of the Bridge as a whole.  Since our sightlines probe very different regions of the Bridge, the hypothesis that the ionization fraction is large in most regions of the Bridge appears reasonable. This mass is  $\\sim$5 times larger than usually estimated in current models \\citep[e.g.][]{muller07a}. However, since the gas mass fraction of the SMC and LMC and the ratio of the gas disk to the stellar disk of SMC can be adjusted in models, and since the SMC mass is not known to better than a factor two, it is possible that simulations may be able to accommodate a larger Bridge mass.    \\begin{deluxetable}{lc} \\tabcolsep=6pt \\tablecolumns{2} \\tablewidth{0pt} \\tablecaption{Summary  of the Properties in the Diffuse Gas of the Magellanic Bridge \\label{t-prop}} \\tablehead{\\colhead{Parameter}    &  \\colhead{Value} } \\startdata $[{\\rm Z/H}]$$^a$\t\t    & $-1.05 \\pm 0.06 $  \\\\ $[$\\ovi$/$H$]$\t\t\t    & $-3.8$ to $-3.3$  \\\\ $[$\\ovi$/$\\hi$]$\t\t    & $-3.1$ to $-2.3$  \\\\ Depletion of Si\t\t\t    & $-0.45 \\pm 0.06$  \\\\ Depletion of Fe\t\t\t    & $-0.61 \\pm 0.06$  \\\\ Fraction of total \\hi\\\t\t    & $\\sim$20\\%   \\\\ Fraction of total \\hii\\\t\t    & $\\sim$80\\%   \\\\ Fraction of H$_2$\t\t    & $\\la 0.004$\\% \\\\ Density $n_e \\simeq n_p$$^b$\t    & $<0.03$--$0.1\\sqrt{T_4}$  cm$^{-3}$ \\\\ Total H mass$^c$\t\t    & $\\sim 9\\times 10^8$ M$_\\odot$ \\enddata \\tablecomments{$a$: Average metallicity from the interstellar and B-type stellar estimates. $b$: Density of ionized gas. $T_4$ is the temperature in units of $10^4$ K. $c$: Adopting a total \\hi\\ mass of $1.8 \\times 10^8$ M$_\\odot$ \\citep{bruns05} and assuming that the fraction of ionized is 80\\% throughout the Bridge. Note that the {\\em total} mass of the Bridge should include the H$+$He mass. } \\end{deluxetable}    \\subsection{Sources of Ionization of the Magellanic Bridge}\\label{sec-discuss-ion} One likely source of ionization of the Bridge is the stars within the Bridge. A young population of stars was discovered throughout the Bridge \\citep{irwin90,bica95,demers98}. The density of OB associations is, however, not uniform: the OB-type stars are highly concentrated in the SMC Wing where the \\hi\\ column densities are the largest (roughly delimited by the box in Fig.~\\ref{fig-map}) and are far more sparse farther away from the SMC \\citep{irwin90,battinelli92}. The rectangle shown in Fig.~\\ref{fig-map} also corresponds to the CO survey undertaken by \\citet{mizuno06}, where CO emission was discovered, suggesting that conditions might be conducive to star formation.  We note that the star-formation history must be quite different from the less dense region of Bridge (see \\S\\ref{sec-per}), since in the SMC Wing, an SMC-like metallicity is derived from the analysis of B-type stars \\citep{lee05}.\\footnote{There might be some confusion in the literature with the nomenclature ``SMC Wing'' and ``Bridge'' (also called InterCloud Region -- ICR). In some cases, the Wing is an environment by itself (usually in the studies of stars), in other cases (usually in  the studies of CO and \\hi\\ emission) it is a direct component of the Magellanic Bridge. The latter is possible and seems to be justified in view of the \\hi\\ column density and velocity maps. But  the SMC Wing has a much higher concentration of OB associations \\citep{irwin90} and an SMC-like metallicity \\citep{lee05}, which seem to imply a different evolution from the Bridge at RA(J2000)\\,$\\ga 2^{\\rm h}\\,30^{\\rm m}$ (see \\S\\ref{sec-per}). In particular, in view of these differences, it does not seem to be justified to extrapolate the properties derived in the SMC Wing to the whole Bridge and vice versa, at least not before a better understanding of the connection between these two regions.} Yet, the presence of early-type stars of age 20 Myr old or younger well inside and throughout the Bridge \\citep{demers98,rolleston99} implies ongoing star-formation at some level.  The H$\\alpha$ intensity of a cloud can be used to estimate the incident Lyman continuum flux. In our Galaxy, \\citet{reynolds84,reynolds93} has used this approach to calculate the power required to ionize the Galactic warm ionized medium up to 2--3 kpc above the plane, comparing it with the power available from OB-type stars. Following \\citet{tufte98}, the incident Lyman continuum flux is $F_{\\rm LC} = 2.1\\times 10^5 I[{\\rm H}\\alpha]$ photons\\,s$^{-1}$\\,cm$^{-2}$, where $I[{\\rm H}\\alpha]$  in Rayleigh is assumed to arise solely from photoionization. Using the 3$\\sigma$ upper limit derived from \\cii*, we find for the diffuse ionized gas of the Bridge (outside denser \\hii\\ regions around OB-type stars) $F_{\\rm LC}(\\rm MB) < 0.4 \\times 10^5$ photons\\,s$^{-1}$\\,cm$^{-2}$. This rate is at least 100 times smaller than the one found in the Galaxy at high galactic latitudes \\citep{reynolds84}. This rate would increase by a factor four if we used the \\cii*\\ detection but as discussed above the origin of the excitation is uncertain as both cold neutral and warm ionized could contribute to the observed absorption. We therefore concentrate in this section on the origin of the nearly fully ionized gas.  The flux of Lyman continuum photons from stars is highly dependent on their types \\citep{panagia73}. For example, an O6-type star will produce a factor 10 more ionizing photons than an O9. But more importantly the rate of photons drops dramatically for B1 and later types: a B1 compared to an O9 produces 600 times fewer ionizing photons per second. In our Galaxy, surveys of OB-type stars have shown that O-type stars produce 93\\% of ionizing photons, even though they are much less numerous than B-type stars \\citep{terzian74,abbott82}.  Unfortunately, little is known about the OB-type distribution in the Bridge, especially outside the SMC Wing. \\citet{irwin90} reported an O8 star outside the SMC Wing, and therefore there is likely to be O-type stars in the Bridge. In the SMC Wing, \\citet{pierre86} estimated that the number of ionizing photons over a very small area (0.1 square degree) is $ 70 \\times 10^6$ photons\\,s$^{-1}$\\,cm$^{-2}$. Such a value would only require 0.1\\% escaping Lyman continuum photons from the SMC Wing to ionize the Bridge. However, the sample of \\citet{pierre86} is small and may be not be representative of the OB distribution in the Wing. In particular, 20\\% of the stars in their sample with types earlier than B0 are O7, which seems quite a larger number (in our Galaxy this number is about 4\\%, e.g., Terzian 1974); the three O7 stars in their sample contribute to the majority  of ionizing photons. A better characterization of the spectral types of the stars in the SMC Wing and Bridge would clearly help to better understand this important source of ionizing photons.   OB stars in the Bridge and Wing will also create an  environment that is favorable for the escape of Lyman continuum photons and for sustaining a high level of ionization throughout the Bridge. Indeed, OB associations through the winds and death of massive stars can produce large bubbles and chimneys that can help ionizing photons to travel large distances \\citep[e.g.,][]{dove94,dove00,norman89}. Such \\hi\\ shells and even possibly blow-outs and chimneys were surveyed by \\citet{muller03}, but only in the Wing at RA\\,$< 3^{\\rm h}$ and some may be only due to projection effect \\citet{muller07a}. Furthermore, in an overall low density medium, it is likely that signatures of supernova remnants and stellar winds will be difficult to decipher since rather than forming shells as observed in dense \\hi\\ regions, they might just dilute themselves in their surroundings.  Supernovae could also play a role in the ionization of the Bridge. Since stellar formation has occured for  200--300 Myr \\citep{harris07}, it is likely that there have been several supernovae in and near the Bridge. Assuming steady ionization of H (with $\\sim$13.6 eV per ionization), our limits on H$\\alpha$ require $< 9 \\times 10^{-7}$ erg\\,s$^{-1}$\\,cm$^{-2}$. Over the entire Bridge (assuming a thickness of 5 kpc), this corresponds to a constant power input to the gas of $<7\\times 10^{38}$ erg\\,s$^{-1}$. This power requires a supernova rate of about one every 50,000 yr  \\citep[see][]{abbott82}. The supernova rate is unknown and, as discussed by \\citet{reynolds84}, only a small fraction ($\\sim 17$\\%) of the kinetic energy injected in the interstellar medium may be converted into ionizing hydrogen. Yet the presence of young, massive stars, supernovae have likely contributed to the ionization and heating of the Bridge.  Highly ionized gas found in the Bridge \\citep{lehner01,lehner02} is  an indirect signature of hot gas that is still present in significant amounts (e.g. where the observed high ions are in an interface between the hot and cooler gas) or that was present and is in the process of cooling. Both interfaces and cooling of hot gas can produce photons that may be important sources of photoionization of the environment \\citep{slavin00,borkowski00,knauth03}. For example, \\citet{slavin00} modeled the contribution to the photoionization from cooling of hot supernova remnants in the disk of our Galaxy and found that this source was adequate to account for the observed ionization of the Galactic halo. Therefore supernovae in the Bridge and in the SMC and LMC near the Bridge could not only inject energy and heat the Bridge but also their cooling remnants may provide a source of ionization.  While the escaping of ionizing photons from the SMC Wing could be an important source of ionization, the escape of ionizing photons from the Galactic, LMC, and SMC disks \\citep[see][]{bm99,dove00} is also likely to contribute to the ionization of the Bridge. Figure~8 in \\citet{putman03} shows predicted $I[{\\rm H}\\alpha]$ near the LMC caused by the escape of ionizing photons from the LMC and the Galactic stellar bulge. The H$\\alpha$ emission estimate produced from the leaking photons of the LMC (and Galaxy, but the latter is rather negligible) could reach 0.1 to  0.25 R near the LMC. Thus, such leakage alone is likely an  important source of ionizing photons for the Bridge.  While it is not clear which ionizing source is dominant in the Bridge, there appear to be enough photons to maintain the high level of ionization seen in the Bridge. The presence of massive stars in the rarefied environment of the Bridge is likely sufficient in itself for this environment to be largely ionized. Since cooling and recombination times are expected to be longer than in the Galactic environment because the Bridge gas has an extremely low metallicity and density, the ionization is likely to be sustainable for a long time. For example, for a $10^6$ K gas with a density $10^{-3}$ cm$^{-3}$, the radiative cooling is $t_{\\rm cool} \\sim 2$--$3\\times 10^9$ yr for a 0.1 solar metallicity \\citep{gnat07}, about 10 times longer than for a solar metallicity environment. Furthermore, according to \\citet{harris07}, star formation has occured in the Bridge for the last 200--300 Myr, with a lower rate over the last $\\sim$40 Myr. Therefore, several generations of OB stars could have ionized the Bridge over a long time period.   We finally note that the ionization fraction toward DGIK\\,975 is not only larger than toward DI\\,1388, but the gas is also more highly-ionized. \\citet{lehner02} reported the measurements of \\ovi, \\siiv, and \\civ\\ toward these stars. In particular, $N_{\\rm DGIK\\,975}($\\ovi$) \\simeq 13 N_{\\rm DI\\,1388}($\\ovi$)$, while $N_{\\rm DGIK\\,975}($\\siii$) \\simeq 2 N_{\\rm DI\\,1388}($\\siii$)$. One can refer to Fig.~11 in \\citet{lehner02} to see the dramatic difference in the \\ovi\\ absorption profiles toward DI\\,1388 and DGIK\\,975. Recently, \\citet{lehner07} argued for the presence of outflows and even possibly a hot galactic wind from the LMC. Since DGIK\\,975 is in the outskirts of the LMC and DI\\,1388 is farther away from the LMC (see Fig.~\\ref{fig-map}), the larger amount of ambient hot gas toward DGIK\\,975 might be related to the ouflows from the LMC. Perhaps the feedback processes occurring in the LMC inject energy into the Bridge. We note that the amount of metals incorporated into the Bridge gas from the LMC at large distance must be extremely small since neither the interstellar nor stellar abundances \\citep[$-1.2 \\pm 0.2$ dex solar for DGIK\\,975, see][]{rolleston99,lee05} suggest an increase of the metallicity in this region.    \\subsection{Metallicity of the Magellanic Bridge}  The metallicity of the Bridge is difficult to accommodate with most models of the SMC-LMC-Galaxy interactions (but see below). The present-day low metallicity of the Bridge has, however, a simple explanation in the context of the chemical evolution of the SMC if the Bridge is much older than 200 Myr. Indeed the bursting model for the star formation rate in the SMC described by \\citet{pagel98} suggests that its metallicity only slowly increased from $-1.4$ to $-1.1$ dex between about 12 and 2.5 Gyr ago. In this model \\citep[see also][]{harris04,idiart07}, it is only in the last $\\sim$2.5 Gyr that the SMC metallicity has exceeded the modern Bridge value. This is a large window for forming the Bridge from SMC gas, and suggests that the Magellanic Bridge could be 10 times older than current  N-body simulations suggest. \\citet{harris07} showed that there were several bursts (2.5, 0.4, 0.06 Gyr ago) of star formation in the SMC in the last  $\\sim$3 Gyr, but the most important one being 2.5 Gyr ago. New bursts of star-formation are believed to be often caused by the effects of close tidal interactions between galaxies \\citep[e.g.,][]{larson78,barton07}. While \\citet{harris07} noted that the bursts at 2.5 Gyr and 0.4 Gyr are temporally coincident with past perigalactic passages of the SMC and Galaxy, these epochs also correspond to close encounters between the SMC and LMC \\citep{kallivayalil06}. Due to its proximity, the LMC tidal perturbations on the SMC are larger than those of the Galaxy \\citep{lin95}. However, calculations of the orbits of the SMC and LMC with the updated proper motions of the SMC and LMC still predict two more perigalactic approaches between the SMC and LMC since the 2.2 Gyr close encounter at $\\sim$1.4, 0.2 Gyr, the latter being the closest approach \\citep{kallivayalil06}. It is therefore unclear if the Bridge could be stable over 2--3 Gyr (Gardiner \\& Noguchi 1994 argued it could not be stable over a very long time but principally because of the Galaxy's gravitational field; yet they did not quantify it further).  On the other hand, \\citet{gardiner96} predicted from their model that the Bridge gas originated primarily from the halo of the SMC, not its disk. Under the assumption that the tidal models are correct, this would mean that despite the steep increase of star formation rate in the SMC disk, the metallicity of the SMC halo would not have been enhanced in the last 2.5 Gyr by any disk material via outflows or else was diluted by an external low-metallicity source. Assuming that the initial mass of the Bridge is $M_{\\rm MB}^i$ and metallicity is $Z_{\\rm MB}^i$, and a ``diluting'' gas with  $M_{\\rm dil}$ and $Z_{\\rm dil}$, the  metallicity of the mixed-up gas in the Bridge, $Z_{\\rm MB}^f$, can be expressed as: \\begin{equation} Z_{\\rm MB}^f = \\frac{Z_{\\rm MB}^i M_{\\rm MB}^i + Z_{\\rm dil} M_{\\rm dil}}{M_{\\rm MB}^i + M_{\\rm dil}}\\,; \\end{equation} and assuming that $Z_{\\rm MB}^i M_{\\rm MB}^i \\gg Z_{\\rm dil} M_{\\rm dil}$, we have \\begin{equation} Z_{\\rm MB}^f = Z_{\\rm MB}^i \\left(1+ \\frac{M_{\\rm dil}}{M_{\\rm MB}^i}\\right)^{-1}\\,. \\end{equation} Therefore, if the metallicity of the tidally disrupted material was present-day SMC like, $Z_{\\rm MB}^i = 0.25 Z_\\odot$, one would need  $M_{\\rm dil} \\approx 2 M_{\\rm MB}^i$ (with $Z_{\\rm dil} \\la 0.01 Z_\\odot$) to achieve $Z_{\\rm MB}^f \\approx 0.09 Z_\\odot$. In contrast, $Z_{\\rm MB}^i M_{\\rm MB}^i \\ll Z_{\\rm dil} M_{\\rm dil}$, $Z_{\\rm dil}^i = 0.09 Z_\\odot$ and $Z_{\\rm MB}^i = 0.25 Z_\\odot$ would require $M_{\\rm dil} \\ga 200 M_{\\rm MB}^i$ to have $Z_{\\rm MB}^f = 0.09 Z_\\odot$. A determination of the ratio the SMC halo to SMC disk particles that were pulled into the Bridge in simulations would be valuable. This scenario might be plausible since dwarf galaxies can have the \\hi\\ gas occupying radii quite larger than the stellar component \\citep{salpeter96}. This scenario was also proposed by \\citet{harris07} to explain the absence of evidence of tidally stripped stars in the Bridge. If this envelope was extended enough, it could have reduced the impact of feedbacks from massive stars in the SMC at large distances from the SMC for $\\sim$2--3 Gyr.  If the very low metallicity material did not come from an extended halo of gas around the SMC stellar disk (and assuming the Bridge is young),  the Bridge could have swept up metal-poor ambient matter over time. For example, measurements of the proper motion of the Fornax dwarf spheroidal galaxy suggest that the orbit of the Magellanic Clouds was crossed by Fornax some 200 Myr ago \\citep{dinescu04}, a time coincidentally similar to the believed formation epoch of the Bridge. The Fornax has a wide range of metallicity from $-2.0$ to $-0.4$ dex solar, and peak near $-0.9$ dex \\citep{pont04}. While this is speculative, low-metallicity material from the Fornax or other sources may have mixed-up with Bridge gas.  Bekki \\& Chiba (2007a,2007b) have recently produced self-consistent chemodynamical models of the LMC-SMC-Galaxy interactions and found a metallicity distribution that peaks at $-0.8$ dex but could extend down to $-0.9$ dex. To the best of our knowledge, this is the only model that attempts to explain the lower metallicity of Bridge. In this model, a low metallicity is reached because the Bridge is mostly formed from the SMC halo gas, where the metallicity is significantly smaller because of an assumed metallicity gradient in the SMC, although it is not clear that there is or has been such a gradient since the present-day metallicity in the SMC Wing \\citep{lee05} is similar to the SMC.  \\subsection{Concluding Remarks}\\label{sec-per} Our analysis provides the first metallicity of the Magellanic Bridge gas and shows for the first time that the Bridge is substantially ionized. As our sample of sightlines is small, it is, however, not certain if large chemical and/or ionization inhomogeneties exist. As we argued above, our three sightlines probe the low \\hi\\ column density regions of the Bridge (RA\\,$\\ga 3^{\\rm h}$) at very different locations and depths. The present-day metallicity of the Bridge derived from 3 B-type stars (including DGIK\\,975) also suggests some chemical homogeneity \\citep{rolleston99,lee05}. We note that two of the Rolleston et al. stars are close to the SMC Wing, but outside the high \\hi\\ column density regions. In contrast, \\citet{lee05} found an SMC-like metallicity for 4 B-type stars well inside the SMC Wing (and in the high \\hi\\ column density region of the Wing). While the samples are small in each region, it seems very unlikely that current samples probe each an unrepresentative population.  The difference of 0.4--0.5 dex in the metallicity between these two adjacent regions is another puzzle, although we note that \\citet{lee05} quoted errors of $\\pm0.3$ dex in their abundances and the difference may actually be smaller. These authors also used the same method for the SMC Wing and Bridge stars and found a systematic difference of about 0.5 dex in the abundance of these stars. This difference may be explained (assuming chemical inhomogeneities in each region are small) if the Wing and the low \\hi\\ column density regions of the Bridge have a different origin or  the star-formation rate and initial-mass function were quite different in each region, and/or the Wing was not diluted with extremely-low metallicity gas. In the former scenario, the Wing must be younger than the rest of the Bridge, so that it was formed recently from gas and stars stripped from the SMC that have a present-day SMC metallicity. This would imply that the Bridge itself must be much older and that the star-formation was  inefficient over a long period of time or only started in the last 200--300 Myr as the results of \\citet{harris07} suggest. In the second scenario, at the epoch the Bridge and the Wing were formed, the metallicity should have been around  $\\la -1.1$ dex in both regions. However, it would seem very unlikely that the metallicity of the Wing has increased by 0.5 dex in the last 200 Myr assuming this is the age of these structures. In the SMC where star formation is more important than in the Wing, about 2--3 Gyr were needed to increase the metallicity from $-1.1$ dex to the present-day metallicity \\citep{pagel98}. Therefore the second scenario is not likely if these structures are only 200 Myr old. In the third scenario, the Bridge outside the Wing regions could have been diluted with low metallicity gas, while the Wing was not or at a neglible level, possibly because the star-formation was more efficient or the Wing material comes from deeper regions of  the SMC (where the gas has present-day SMC metallicity) than the rest of the Bridge gas. The first and third scenarios are consistent with the findings of evolved stars in the Wing and not anywhere else in the Bridge \\citep{harris07}.  Hence despite the fact that current tidal simulations reproduce well the observed \\hi\\ structures (i.e. the Bridge, Stream, and Leading Arm), they seem to be in conflict with the above conclusions regarding the metallicities. As already mentioned, the velocities of the LMC and SMC adopted in the models of \\citet{gardiner96} and subsequent models differ from those derived from the recent proper motion estimates of the SMC and LMC by  over 100 \\km\\ \\citep{kallivayalil06,kallivayalil06b}. It is also unclear how bound the SMC and LMC are and were \\citep{kallivayalil06}. With those updated proper motions, \\citet{besla07} strongly suggest that existing numerical models of the Clouds may no longer be appropriate, and in particular cannot explain anymore the Stream and Leading Arm.  As discussed by \\citet{nidever07}, current simulations may also miss important physics, including feedback processes such as outflows from the LMC that may interact with the Bridge gas. In particular, it is quite important to understand feedback and mixing processes in the SMC disk and halo for the last 3 Gyr if the Bridge is only 200 Myr old and its main origin is the halo of the SMC.  Future theoretical investigations using the updated proper motions of the LMC and SMC should not only try to reproduce the \\hi\\ properties but also accommodate the low metallicity of the Bridge, the apparent discrepancy of the metallicity between the Bridge and the SMC Wing, its ionization structure, and the apparent lack of old stars in the Bridge \\citep{harris07}. Future models should in particular investigage if the Bridge can be stable over a long time (2--3 Gyr) or despite an important burst of star formation within the SMC 2.5 Gyr ago, the SMC halo gas could maintain a very low metallicity. If none of these conditions are satisfied, an important (depending on the chemical homogeneity in the Bridge) amount of matter in the Bridge must come from somewhere else. If the Bridge could survive several perigalactic approaches between the SMC and LMC, it would be interesting to know if the low density and low metallicity regions of the Bridge could have been formed some 2--3 Gyr ago, while the SMC Wing would have been produced at a latter time, some 200 Myr ago during the last and closest approach between the SMC and LMC.  From an observational point of view, a larger stellar and interstellar sample where the abundances can be estimated in the Bridge and the SMC Wing would be extremely valuable. The soon to be installed Cosmic Origins Spectrograph (COS)  will be particularly useful for the study of the gas-phases in the Bridge since fainter targets can be observed with it. In particular, to better characterize the ionization structure, future COS and deep H$\\alpha$ observations will be invaluable. A better characterization of the stellar population in the Bridge will also help to constrain the origin of the ionization, the star-formation history, and the inital-mass function in the Bridge. Since the Bridge is the closest gaseous and stellar tidal remnant, and since interactions between galaxies are rather common, future observational and theoretical efforts offer the unique opportunity to comprehend the physical processes and origin of such structures."
    },
    "0801/0801.2158_arXiv.txt": {
        "abstract": "We show that a large-scale, weak magnetic field threading a turbulent accretion disk tends to be advected inward, contrary to previous suggestions that it will be stopped by outward diffusion.  The efficient inward transport is a consequence of the diffuse, magnetically-dominated surface layers of the disk, where the turbulence is suppressed and the conductivity is very high.  This structure arises naturally in three-dimensional simulations of magnetorotationally unstable disks, and we demonstrate here that it can easily support inward advection and compression of a weak field.  The advected field is anchored in the surface layer but penetrates the main body of the disk, where it can generate strong turbulence and produce values of $\\alpha$ (i.e., the turbulent stress) large enough to match observational constraints; typical values of the vertical magnetic field merely need to reach a few percent of equipartition for this to occur.  Overall, these results have important implications for models of jet formation which require strong, large-scale magnetic fields to exist over a region of the inner accretion disk. ",
        "introduction": "\\label{sec:intro}  Early theoretical work on accretion disks argued that a large-scale magnetic field (of, for example, the interstellar medium) would be dragged inward and greatly compressed by the accreting plasma \\citep{Bisnovatyi1974,Bisnovatyi1976,Lovelace1976}.  Figure \\ref{fig:hourglass} illustrates this concept by showing a sketch of an ordered magnetic field threading an accretion disk, in which inward advection has caused the magnetic field lines to bunch together into an ``hourglass'' shape.  This was thought to be a simple mechanism for generating dynamically significant fields in the inner disk.  In the present paper, we revisit this issue, building off the recent work of \\citet{BKL2007}.  Our motivation for doing so is that in the intervening years, the early theoretical arguments have been challenged.  More detailed models of turbulent disks suggested that a large-scale, weak magnetic field such as that shown in Figure \\ref{fig:hourglass} in fact will diffuse outward rapidly \\citep*{vanball1989,Lubow1994} if the turbulent magnetic diffusivity and turbulent viscosity are of similar order of magnitude, as they are expected to be \\citep{Parker1971,Bisnovatyi1976,Canuto1988}---the turbulence responsible for driving the accretion also leads to enhanced reconnection of the large-scale radial field across the thickness of the disk, thereby causing the vertical field to diffuse away.  This cast doubt on the idea that weak fields could be dragged inward and compressed by advection.  At the same time, it was known that the angular momentum loss to magnetohydrodynamic (MHD) outflows from a disk threaded by a sufficiently {\\it strong} large-scale field could more than offset the outward diffusion and lead to a rapid, implosive increase of the field in the central region of the disk \\citep*{Lovelace1994}.  However, it seemed to be the case that growth of a strong magnetic field ``from scratch,'' due to continual advection of a weak field, was impossible in a thin disk.  Although this conclusion has been occasionally challenged \\citep[e.g.,][]{OgilvieLivio2001}, it is still generally accepted, which has led to the recent suggestion that special conditions (extremely nonaxisymmetric regions of strong field in an otherwise weakly-magnetized disk) are required for the field to be advected inward \\citep{Spruit2005}.  At the same time, recent three-dimensional MHD simulations have been performed that allow this issue to be addressed computationally. These simulations resolve the largest scales of magnetorotational turbulence and therefore self-consistently include the turbulent viscosity and diffusivity (without having to prescribe their values {\\it a priori}). Most simulations performed to date have investigated conditions in which the accreting matter does not contain any net magnetic flux and where no magnetic field is supplied at the boundary of the computational domain. However, in  one simulation, weak poloidal flux injected at the outer boundary was clearly observed to be dragged into the central region of the disk, leading to the buildup of a strong central magnetic field \\citep*{Igumenshchev2003}. A similar process, albeit transient, may occur in simulations without a net magnetic flux; there, radial stretching of locally poloidal field lines in the initial configuration often leads to large-scale poloidal fields and jet structures in the inner disk \\citetext{e.g., \\citealt{Hirose2004,DeVilliers2005,HawleyKrolik2006}; see also the discussion in \\citealt{Igumenshchev2003} and \\citealt{McKinney2007} and especially the simulations of \\citealt{McKinney2004} and \\citealt*{Beckwith2007}, which explore the effect of different initial field geometries on the formation of jets}. The extent to which any of the advection of magnetic field lines seen in numerical simulations requires the presence of a thick disk or nonaxisymmetric conditions is unclear.  In light of these numerical results, we return to the question of inward advection of magnetic fields in this paper, allowing for the possibility that the disk is thin and axisymmetric and asking once again whether advection of a weak field is possible under these conditions.  The mechanisms we discuss here can occur in sufficiently ionized regions of any accretion disk; although they are perhaps most widely applicable to disks around black holes (where the large-scale magnetic field arises entirely within the accreting plasma), they are relevant for disks around many other types of accreting objects as well.  \\begin{figure} \\epsscale{0.95} \\plotone{f1.eps} \\epsscale{1.0} \\caption{Sketch of the magnetic field threading an accretion disk, showing the increase of the field due to its assumed inward advection with the gas (as proposed in early theoretical models).} \\label{fig:hourglass} \\end{figure}  The organization of this paper is as follows.  In \\S \\ref{sec:advection}, we analyze the advection of a large-scale field in an accretion disk and point out the importance of the vertical structure of the disk, which was not taken into account in most previous studies.  Based on an earlier suggestion \\citep{BKL2007}, we show that the thin, highly conducting surface layer of the disk, where turbulence is suppressed, allows a large-scale magnetic field to be advected inward and compressed.  In \\S \\ref{sec:alpha}, we argue that the resulting magnetic flux through the main body of the disk (due to the large-scale field being advected inward) can produce values of the turbulent $\\alpha$ parameter that are in accord with observational data.  This is in contrast with numerical simulations of turbulent disks without a net imposed magnetic flux, which are unable to generate large enough turbulent stress.  Finally, in \\S \\ref{sec:conditions}, we derive detailed conditions on the field strength, geometry and ionization fraction that are required for the field to be advected inward and show that these are typically weak constraints.  Conclusions of this work are summarized in \\S \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  This paper reanalyzes the advection of a large-scale, weak magnetic field in an accretion disk.  We consider the vertical structure of the disk, which strongly influences the vertical profile of the conductivity, as pointed out by \\citet{BKL2007}.  In the thin, diffuse surface layers of the disk, the magnetic energy density is large enough compared to the thermal energy density that magnetorotational turbulence is suppressed.  As a consequence, magnetic field lines threading the surface layer can be advected inward with the main body of the disk, without being opposed by turbulent diffusion.  No special conditions are required for the field to be advected inward except that it meet the rather modest constraints in \\S \\ref{sec:conditions}.  The required field strengths are relatively weak, and the primary constraint on the field geometry is simply that it must help the nonturbulent surface layer accrete inward (i.e., the vertical magnetic stress must extract angular momentum from this layer, in virtually any amount).  This can be accomplished either via coupling between the main, turbulent body of the disk and the surface, or via a wind or jet.  The simplest way for this condition to be met is if the accretion flow is provided with a weak, large-scale vertical seed field threading the outer region of the disk (which could come from the interstellar medium or a companion star), although in some cases, the proper geometry may be attained entirely as a result of fields produced via the local magnetorotational dynamo.  Once a weak, large-scale field with the proper geometry is in place, the field will be advected inward along the disk's surface layer and strengthened as it is compressed along with the accretion flow.  The presence of a large-scale magnetic field anchored in the surface layer will drive strong magnetorotational turbulence in the main body of the disk, which can consequently produce values of the $\\alpha$ parameter (i.e., the turbulent stress) large enough to match observational constraints.  We find that typical vertical fields on the order of a few percent of equipartition are required for this to occur.  Because the field is advected inward and anchored in the surface layer, there is no need to worry about maintaining the required vertical fields via internal magnetorotational fluctuations \\citep[as suggested by][]{Pessah2007}.  We propose that long-term changes in $\\alpha$ (in response to the history of magnetic field advection) should be explored as a possible source of the long-term evolution in the variability patterns seen in the light curves of X-ray binaries such as GRS~1915+105 \\citep[see, e.g.,][]{Tagger2004}.  The mechanisms we discuss in this paper are relevant to many different kinds of accreting objects.  In the outer part of the disk (far away from the central star or black hole), our work should be applicable provided only that the disk is sufficiently ionized (see \\S \\ref{sec:strength}).  Closer to the inner part of the disk, our work is most obviously applicable to black holes, where the large-scale magnetic field arises entirely within the accreting plasma.  However, it may also be relevant to the case where a large-scale field from a magnetized central star penetrates the disk; previous suggestions that outward radial diffusion of the field may be an important process in such systems \\citep*[e.g.,][]{Lovelace1995,BardouHeyvaerts1996,Agapitou2000,Uzdensky2002} should be examined in light of our work.  Finally, the applicability of our work to radiation-dominated regions of the disk may require further analysis, because it is unclear whether the vertical structure of these regions can support the inward advection of magnetic fields (although we argue here that it can).  Future numerical simulations may eventually be able to address this point.  Also, we have assumed axisymmetric disks in this paper, but the equations we have derived apply more generally if they are averaged over azimuth (with the introduction of appropriate correction factors).  Thus, a magnetic field with an {\\it azimuthally-averaged} strength and {\\it azimuthally-averaged} geometry that roughly meets the criteria in \\S \\ref{sec:conditions} is likely to be able to advect inward in an accretion disk.  Nonaxisymmetric advection of a large-scale magnetic field \\citep[e.g., as envisioned by][]{Spruit2005} is clearly consistent with the mechanism we have proposed in this paper, but we have also shown that advection is equally plausible in an axisymmetric disk, without having to assume any special conditions.  More numerical simulations that investigate disks in which a large-scale magnetic field is supplied at the outer boundary are clearly needed.  Finally, our work has important implications for models of jet formation which require strong, large-scale magnetic fields to exist over a region of the inner accretion disk.  As suggested more than 30 years ago, these magnetic fields can arise in the inner disk via a very simple process: advection of a weak field from outside.  This opens up the possibility for magnetically-dominated outflows (i.e., Poynting jets) to exist in the inner regions of disks around a wide variety of accreting objects.  In addition, the radial stretching of field lines produced by advection may allow winds accelerated by the magnetocentrifugal effect \\citep{BlandfordPayne1982} to exist over a wide range of radii.  This is in contrast to the results of numerical simulations in which no magnetic fields are supplied at the outer boundary, and jets only form in an extremely narrow inner region where the energetics may be dominated by relativistic effects that require the presence of a rotating black hole."
    },
    "0801/0801.0477_arXiv.txt": {
        "abstract": "The present scenario for high-luminosity long $\\gamma$-ray bursts is strongly influenced by the paper of Fruchter et al. (2006). Whereas the main contention of this paper that these GRBs occur in low-metallicity irregular galaxies is based on a considerable collection of observational results and although the main thesis is doubtless correct, the paper does not explain the dynamics that produces such GRBs and much of the discussion not directly concerning the main thesis is wrong.  We propose a dynamics and elucidate how the \\citet{Fruchter06} results may be tested, in our neighborhood in the LMC, suggesting that LMC X-3 is a relic of a high luminosity explosion, probably accompanied by a GRB and hypernova explosion.  The way to test our suggestion is to measure the system velocity of the present black hole.  We correct errors of the Fruchter et al. paper in stellar evolution, so that the study of GRBs is consistent with it.  We show that the subluminous GRB 060218 had a low-mass black hole as central engine. ",
        "introduction": "\\citet{Fruchter06} collect an important amount of data on long $\\gamma$-ray bursts.  They propose that the long-duration $\\gamma$-ray bursts are associated with the most extremely massive stars and that they may be restricted to galaxies of limited chemical evolution.  Also that long $\\gamma$-ray bursts are relatively rare in galaxies such as our own Milky Way.  In recent papers \\citep{Mor07, BLMM07} we have developed the Blandford-Znajek model of GRBs into quantitative calculations of the angular momentum energy that can be delivered for GRBs and hypernova explosions.  This was possible because \\citet{Lee02} showed how to work out the Kerr parameters of the rotating black holes.  We showed that the Galactic soft X-ray transient sources were relics of Galactic explosions and constructed the energy of all 15 of the known Galactic sources.  It turns out that most of the Galactic sources underwent subluminous  GRBs, not because they did not possess sufficient rotational energy, but mostly because the rotational energy was so large that it destroyed the accretion disk so quickly that the central engine was dismantled before the GRB could properly develop.  We give as examples the transient sources Nova Sco (GRO J1655$-$40) and Il Lupi (4U 1543$-$47) for which \\citet{Lee02} had predicted Kerr parameters of $a_\\star=0.8$ and which were measured, by \\citet{Sha06}, at the present time, to be $a_\\star=0.65-0.75$ and $a_\\star=0.75-0.85$, respectively.  The natal rotational energy was $430$ bethes (one bethe$=10^{51}$ ergs) and the final (measured) energy is indistinguishable, within errors, from the natal energy.  \\citet{Bro00} had reconstructed the GRB and hypernova explosion for Nova Sco.  The nature of the explosion could be reconstructed from the donor, which accepted a number of $\\alpha$-particle nuclei, especially $^{32}$S which is special for hypernova explosions, but rare in supernova explosions.  Thus, from the near equality of the natal and present rotational energies, only a few percent of the available rotational energy could have been accepted in the GRB and hypernova explosion.  \\citet{Mor07} and \\citet{BLMM07} showed from population syntheses that the soft X-ray transient sources were sufficient in number to account for all of the subluminous bursts in our neighborhood.  What about the cosmological bursts which are the long $\\gamma$-ray bursts considered by \\citet{Fruchter06}?  The Woosley Collapsar model was invented to describe these.  In terms of numbers these are only a few percent of the subluminous bursts \\citep{Lia07}.  We will not take issue with the Woosley Collapsar model describing these, because our binary evolution begins by the donor spinning the He star, progenitor of the black hole, to whatever angular momentum is needed. Then the donor decouples, acting only as a witness to the explosion, with the He star collapsing into a black hole in the same way as in the Woosley Collapsar model, the tidal locking between the donor and the He star being transferred to one between the donor and the black hole.  From the latter, \\citet{Lee02} predicted the Kerr parameter of the black hole.  The main point developed by \\citet{BLMM07} was that the rotational energy of the binary is roughly inversely proportional to the mass of the donor. This follows from Kepler's law and from the fact that in Case C mass transfer (following the He burning), the initial $a_i$s (distance between the giant and the companion) of the binaries are roughly equal.  The fact that, according to \\citet{Fruchter06}, the long $\\gamma$-ray bursts are in low-metallicity galaxies, does not elucidate the dynamics which produce the long bursts.  The dynamics result from the fact that low-metallicity galaxies tend to have stars of higher mass than Galactic.  The higher mass of the donors slows the binary down sufficiently that the rotational energy can be accepted by the central engine.  In other words, the question of energy is a ``Goldilocks\" one.  It must be not too much, because in that case the central engine will be dismantled, and not too little, because that would only be sufficient for a subluminous burst, but for a long high luminosity $\\gamma$-ray burst it must be just right.  We have learned enough about Kerr parameters from our calculations and from the Smithsonian-Harvard measurements to construct a ``guesstimate\".  Namely, we believe that LMC X$-$3 underwent the closest (in energy) explosion in our Galaxy, to a cosmological GRB.  (Of course, one can say that LMC X$-$3 is not in our Galaxy, but in the LMC.)  We believe that in its $\\sim1/3$ solar metallicity \\citep{Russ90}, it tends towards the low-metallicity stars considered by \\citet{Fruchter06}, so that at the least it has somewhere between Galactic and the low metallicity favored there.  There is uncertainty in the masses, but \\citet{Davis06} have a value of $a_\\star\\simeq0.26$ for the present Kerr parameter.  They took $7\\msun$ as the mass of the black hole, which would imply a donor of $\\sim4\\msun$ from the lower end of the measurements of \\citet{Cow83}.  We take these to be representative; other investigators have found other masses, so we suggest our evolution as only a possible one.  All of the binaries in \\citet{BLMM07} and \\citet{Mor07} which had so much energy that they dismantled the black hole accretion disk had donor masses of $(1-2)\\msun$, so the donor in LMC X$-$3 is at least double those masses.  The donor mass is close to the $\\sim5\\msun$ that \\citet{Mor07} estimated would give the energy of a cosmological GRB.  In any case, LMC X$-$3 is the closest ``nearby\" relic of binaries similar to the progenitors of cosmological GRBs.  We can, therefore, use it as an example to try to reconstruct the explosion and model the energy of the explosion.  We estimate the mass loss in the explosion, finding it to be substantial.  It is likely that the estimated system velocity can be measured, at least the radial component of it, which should test our prediction.  In this paper we wish to also summarize results of earlier calculations which have a direct bearing on the 5 papers in the 31 August 2006 Nature \\citep{N0,N1,N2,N3,N4}, in order to show that useful evolutions of black holes have been carried out in the past, of which the 119 astronomers who signed these articles were unaware.  We show that these previous calculations have a direct bearing on the measurements of GRB 060218/SN 2006aj; namely that the central engine was a black hole, not the magnetar conjectured by most of the authors, and that the black hole was one of minimum black hole mass with an $\\sim(18-20)\\msun$ ZAMS (Zero Age Main Sequence) progenitor.  We also correct a number of errors in Galactic black hole evolution in the \\citet{Fruchter06} article. ",
        "conclusions": "We suggest that LMC X-3 may be similar to relics of cosmological GRBs, to the extent that some or most of these latter result from binaries.  The number of the latter is certainly sufficient to produce the GRBs and the binary nature takes care of the necessary angular momentum, which can be achieved by choosing the donor mass.  \\citet{Fruchter06} have made a case that high-luminosity, long GRBs came from irregular low-metal galaxies.  We suggest LMC X-3 as the closest nearby binary from a region of $1/3$ solar metallicity in an irregular galaxy.  The lower metallicity environment has more massive stars, and the donor in LMC X-3 is probably at least twice as massive as the donors of the relics of the Galactic GRBs, which should slow the binary down to a rotational energy that can be accepted.  We believe that LMC X-3 brings us in the metallicity towards the irregular low-metallicity binaries considered by \\citet{Fruchter06}.  Working in this region of energies we feel that we can make predictions, because our calculation of Kerr parameters makes it possible to make quantitative calculations.  Our predictions are, however, at best, order of magnitude, because the necessary properties of the binary LMC X-3 have not been accurately measured.  As they are measured we will probably have to readjust our numbers.  In particular, it would be most valuable to obtain a reliable value for the explosion energy.  Whereas the hypernova energy is probably measured with reasonable accuracy, the kinetic energy required by the work performed by the ram pressure to clear the way for the GRB is mostly hidden, and the net GRB energy is a small difference between large kinetic energy and a large amount of work done by the ram pressure.  We believe that there will be great variations in GRBs depending upon these small differences which must be sensitive to stellar properties.  Our present estimate of $\\sim 60$ bethes for the total cosological GRB energy is the highest recent estimate that we have seen.  For this estimate we invoke equipartition of kinetic and thermal energies.  Finally, we wish to point out that our ability to calculate Kerr parameters and the confirmation of these by the Smithsonian-Harvard measurements makes it possible to calculate in a reliable way the amount of angular momentum carried by the black-hole binary.  This angular momentum is conserved, but the questions to be answered concern how it is to be distributed:  how much remains in the binary, how much energy goes into the explosion and what is the fraction of the latter which goes into kinetic and into heat energy?  Progress in answering these questions is necessary to make a quantitative study out of the GRBs, and to put order into their classification."
    },
    "0801/0801.2969_arXiv.txt": {
        "abstract": "The results of a search for sub-km Kuiper Belt Objects (KBOs) with the method of serendipitous stellar occultations are reported. Photometric time series were obtained on the 1.8m telescope at the Dominion Astrophysical Observatory (DAO) in Victoria, British Columbia, and were analyzed for the presence of occultation events. Observations were performed at 40 Hz and included a total of 5.0 star-hours for target stars in the ecliptic open cluster M35 ($\\beta=0.9\\arcdeg$), and 2.1 star-hours for control stars in the off-ecliptic open cluster M34 ($\\beta=25.7\\arcdeg$).  To evaluate the recovery fraction of the analysis method, and thereby determine the limiting detectable size, artificial occultation events were added to simulated time series (1/f scintillation-like power-spectra), and to the real data.  No viable candidate occultation events were detected.  This limits the cumulative surface density of KBOs to $3.5\\times10^{10}$ deg$^{-2}$ (95\\% confidence) for KBOs brighter than m$_R$=35.3 (larger than $\\sim$860m in diameter, assuming a geometric albedo of 0.04 and a distance of 40 AU).  An evaluation of TNO occultations reported in the literature suggests that they are unlikely to be genuine, and an overall 95\\%-confidence upper limit on the surface density of $2.8\\times10^{9}$ deg$^{-2}$ is obtained for KBOs brighter than m$_R$=35 (larger than $\\sim$1 km in diameter, assuming a geometric albedo of 0.04 and a distance of 40 AU) when all existing surveys are combined. ",
        "introduction": "\\label{sec:intro}  The physical and orbital properties of the Kuiper Belt Objects (KBOs) are believed to provide valuable information about the formation of the solar system.  To date, directly observed KBOs include objects with diameters ranging from $\\sim$25km (eg. 1999 DA$_8$, \\citealt{gladman01}; and 2003 BH$_{91}$, \\citealt{bernstein04}) to 2400km \\citep[eg. Eris,][]{brown06}.  KBOs with diameters below a few tens of kilometers are beyond the detection of current ground and space-based telescopes.  The KBO size distribution is the outcome of accretion and collisional evolution in the outer solar system \\citep{stern96b,davis97,stern97,kenyon99b,kenyon04,pan05}.  The cumulative luminosity function (CLF) of the KBOs provides an indirect means of determining their size distribution, and is observed to have the form: \\mbox{$\\sum N(<m_R; m_R<25) = 10^{\\alpha(m_R-R_0)}$}.  The KBO size distribution is then a power-law of the form: \\mbox{$\\dif N \\propto r^{-q} \\dif r$}, with slope $-q$, where $q$ is related to $\\alpha$ via: $q=5\\alpha+1$ \\citep[][G01]{gladman01}. Current observations indicate $\\alpha\\approx0.68$ \\citep{fraser07}, or equivalently, $q=4.6$.  HST observations by \\citet{bernstein04} suggest a lower CLF slope for $m_R>25$.  Current theoretical models of the accretion and fragmentation processes in the early solar system suggest that $q$ remains constant only for larger objects, and that the slope of the size distribution changes at the `break radius', $r_b<$100 km \\citep{kenyon99a, kenyon99b, kenyon04, pan05}.  Objects smaller than the break radius would be in collisional equilibrium and their size distribution would have the so-called Dohnanyi equilibrium slope of $q\\approx3.5$ \\citep{dohnanyi69}.  Object densities and bulk-moduli can influence the slope for the small objects, and the position of the break-radius \\citep{kenyon04}.  The deepest direct observations to date \\citep{bernstein04} were performed with HST and reached a limiting magnitude of R=28.5 (corresponding to an object diameter of $\\sim$20km for an albedo of 0.04 at 40AU).  The number of observed objects was $\\sim$25$\\times$ lower than the extrapolation of the current CLF slope.  This absence of faint objects indicates a decrease in the density of the Kuiper Belt beyond 50 AU; and a dearth of sub-100km KBOs, in general.  Direct observation is not a feasible option for pursuing the much fainter km-sized KBOs.  Some indirect methods of inferring the slope in the range of the km-sized objects have been explored.  \\citet{stern00} used Voyager 2 images of Neptune's moon Triton to analyze cratering and infer a slope of $q\\approx3$ for the size distribution of the km-sized impactors. \\citet{kenyon01} have placed upper limits on small object densities by examining the optical and FIR sky surface brightnesses of the ecliptic. They found that a straight extrapolation of the CLF to R$\\approx$40-50 mag would yield an ecliptic surface brightness higher than that observed.  One promising method of indirect detection for the smaller KBOs is that of serendipitous stellar occultations \\citep{bailey76,dyson92,brown97,roques00,cooray03a,cooray03b,gaudi04,nihei07}. A km-sized KBO passing through the line of sight to a star will partially obscure the light from that star to reveal the presence of the otherwise invisible object.  Such an occultation cannot be predicted and the practice relies upon serendipity.  A 1 km diameter KBO at 40AU subtends and angle of only 34 $\\mu$as, and the lightcurve resulting from an occultation by such an object is expected to exhibit a series of increases and decreases due to diffraction \\citep{roques87} (henceforth R87).  At opposition, the orbital velocities of the Earth (30 \\kmps) and a typical KBO at 40 AU ($\\sim$4 \\kmps) produce a retrograde relative velocity of $v\\approx$26 \\kmps for the occulter, and the characteristic diffraction shadow cast by a 1 km KBO will pass over a terrestrial observer in a fraction of a second.  Observations must be performed at $\\geq$40 Hz to sample the flux oscillation in an observed lightcurve.  We have performed a KBO occultation search with the 1.8m Plaskett Telescope at the Dominion Astrophysical Observatory (DAO) using a 40Hz camera.  Models of occultation lightcurves are presented in Section~\\ref{sec:model}.  Observation details are given in Section~\\ref{sec:observations}, with data reduction and analysis methods described in Sections~\\ref{sec:reduction}, and~\\ref{sec:detection}, respectively.  Our results are presented in Section~\\ref{sec:results}, and reported KBO occultation detections are evaluated in Section~\\ref{sec:critique}.  A surface density upper limit is presented in Section~\\ref{sec:upperlimit}, with a discussion and summary provided in Sections~\\ref{sec:discussion} and~\\ref{sec:summary}, respectively. ",
        "conclusions": "Non-spherical objects originating in the outer solar system (eg. peanut-shaped comets) suggest that the spherical symmetry assumed in our lightcurve models is not to be expected for TNOs.  However, diffraction shadows cast by objects with R$_{KBO}\\lesssim$0.5 Fsu are circularly symmetric, regardless of the shape of the occulter. Occultation shadows for different-sized elliptical TNOs ($a/b = 2.5$, where $a/b$ is the ratio of the semi-major and semi-minor axes) were generated to illustrate this (see Figure~\\ref{fig:symmetry}).  The circular symmetry of the shadows is preserved for $a_{KBO}\\lesssim$0.5 Fsu ($\\sim$650m for visible light at 40 AU), despite the asymmetry of the occulting masks.  The shadow profiles for objects larger than this do not require cross-correlation to detect; they would have been identifiable by visual inspection of the time series, had any been present.  \\begin{figure}[htbp] \\plotone{f11.eps} \\caption[Occulting masks and diffraction shadows for non-spherical KBOs.] {Occulting masks (upper panels) and diffraction shadows (lower panels) for different-sized non-spherical KBOs (a/b=2.5 in each case, where a and b are the KBO's semi-major and semi-minor axes.) The diffraction shadows are circularly symmetric for a$\\lesssim$0.5 Fsu ($\\sim$650m for visible light at 40 AU), despite the asymmetry of the occulting masks.  The grayscales were chosen individually to cover the range of each shadow's intensity.} \\label{fig:symmetry} \\end{figure}   Plant and recover tests were performed with profiles for 500 elliptical objects having uniformly distributed random parameters of: 200m$<a<$1000m, 0m$<$impact-parameter$<$4000m, 1.0$<(a/b)<$3.0, and $0<\\theta<\\pi$; where $a$, $b$, and $\\theta$ are the semi-major and semi-minor axes of the object, and $\\theta$ is the position angle of the rotated object. R$_{KBO}$ was taken to be $0.5(a+b)$ in order to compare the results to Figure~\\ref{fig:recovered_kbos}. The observed detection limits did not change and are shown in Figure~\\ref{fig:ellip_recovery}. Thus, cross-correlation detection kernels based on circularly-symmetric occulting masks successfully detect occultations by sub-Fresnel-scale irregularly-shaped KBOs.  \\begin{figure}[htbp] \\plotone{f12.eps} \\caption[The recovered sizes and impact parameters of planted occultation events for elliptical objects] {The recovered sizes and impact parameters of planted occultation events for elliptical objects.  Squares represent planted occultation events; filled if recovered, and open if not recovered.  The 80\\% contour from the theoretical recovery tests for {\\itshape circular} 40AU objects (Section~\\ref{ssec:efficiency}) is shown with a solid line.  The planted events had parameters of: 200m$<a<$1000m, 0m$<$impact-parameter$<$4000m, $1.0<(a/b)<3.0$, and $0<\\theta<\\pi$; where $a$, $b$, and $\\theta$ are the semi-major and semi-minor axes of the object, and $\\theta$ is the position angle of the rotated object. R$_{KBO}$ was taken to be $0.5(a+b)$.} \\label{fig:ellip_recovery} \\end{figure}  Our method of generating artificial time series did not incorporate the effects of variable noise levels in a time series.  Although many of our recovery tests were performed with constant-noise artificial time series, they were verified in our final detection run by planting artificial occultation events in the real data.  The results were found to be consistent with simulations we performed with the artificial data, and we believe that the presented estimates for the limiting detectable KBO size, $r_{min}$, and the maximum impact parameter, $b_{max}$, are accurate.  The presence of a period of increased noise within a time series alters the statistical properties of the variates.  The overall {\\itshape measured} standard deviation of the variates will lie between those measured for the low-noise and high-noise segments, which will lead to increases in the rates of false negatives and false positives in the respective segments.  Our time series were broken into segments having approximately constant noise levels to mitigate this problem.      Analysis of our 5.0 star-hours of photometric time series on two B9V stars in the ecliptic open cluster, M35, revealed no viable candidate occultations.  For the first time, the recovery of artificial occultation lightcurves was used to determine the limiting detectable object size and impact parameter for a KBO occultation.  We place an upper limit on the surface density of KBOs of \\mbox{$\\sum N (m_R<35.3)$} = \\mbox{$3.5\\times 10^{10}$ deg$^{-2}$} (m$_R\\simeq$35.3 corresponds to a $\\sim$860m (diameter) object with a geometric albedo of 0.04).  Having evaluated positive detections of TNO occultation events \\citep{chang06,chang07,roques06}, we believe that the published events are likely to be spurious and that no serendipitous stellar occultation by a TNO has yet been reported.  Comparison of the expected lightcurve structure for the C07 events with the observed structures show that the decreases in flux are too brief to be consistent with occultations by small bodies at 40 AU; and evaluation of the variability index, VI, used by R06 revealed that the {\\itshape VI variates} are not normally distributed, but are drawn from a positively skewed distribution, prone to yielding high values.  Also, the two $\\sim$600m (diameter) TNOs discovered in 10 star-hours of R06 data suggest $\\sim$4 such events should be present in the 89 star-hour x-ray time series observed at a similar ecliptic latitude ($\\beta_{R06}=0.4\\arcdeg S,\\ 0.4\\arcdeg S,\\ 7.1\\arcdeg N$, $\\beta_{C06}=6\\arcdeg N$).  Nothing this large was reported by C06/C07.  Combining all results from the different surveys we place an overall 95\\%-confidence upper limit on the TNO surface density of: \\mbox{$\\sum N (m_R<35.0)$} = \\mbox{$2.8\\times 10^{9}$ deg$^{-2}$}.  A direct extrapolation of the observed CLF slope ($\\alpha\\approx$0.7) would yield a density of \\mbox{$\\sum N (m_R<35.0)$} = \\mbox{$1.9\\times 10^{8}$ deg$^{-2}$}.  Our work is the first complete and careful examination of observational data for KBO occultations and their proper analysis and interpretation."
    },
    "0801/0801.4474_arXiv.txt": {
        "abstract": "We report on \\namefull, a faint eclipsing system composed of two M dwarfs.  The variability of this system was originally discovered during a pilot study of the 2MASS Calibration Point Source Working Database.  Additional photometry from the Sloan Digital Sky Survey yields an 8--passband lightcurve, from which we derive an orbital period of $2.6390157 \\pm 0.0000016$ days.  Spectroscopic followup confirms our photometric classification of the system, which is likely composed of M0 and M1 dwarfs.  Radial velocity measurements allow us to derive the masses (M$_1 = 0.66 \\pm 0.03 \\msun$; M$_2 = 0.62 \\pm 0.03 \\msun$) and radii (R$_1 = 0.64 \\pm 0.08 \\rsun$; R$_2 = 0.61 \\pm 0.09 \\rsun$) of the components, which are consistent with empirical mass--radius relationships for low--mass stars in binary systems.  We perform Monte Carlo simulations of the lightcurves which allow us to uncover complicated degeneracies between the system parameters.  Both stars show evidence of H$\\alpha$ emission, something not common in early--type M dwarfs.  This suggests that binarity may influence the magnetic activity properties of low-mass stars; activity in the binary may persist long after the dynamos in their isolated counterparts have decayed, yielding a new potential foreground of flaring activity for next generation variability surveys. ",
        "introduction": "\\label{sec-intro}  Low--mass dwarfs ($0.07\\msun \\leq$ M$_{\\star} \\leq 0.7\\msun$) comprise $\\sim$ 75\\% of all stars in the Milky Way, making them the most common luminous objects in the Galaxy \\citep{1995AJ....110.1838R}. A significant amount of theoretical work has been devoted to constructing models that describe the physical processes in the interiors and atmospheres of these low--mass stars \\citep[e.g.][]{1993ApJ...406..158B,1998A&A...337..403B,1999ApJ...512..377H,2000ARA&A..38..337C}. Differences in the predictions of these models can be subtle, and distinguishing between them requires precise empirical constraints on fundamental stellar properties \\citep[mass, radius, luminosity, and effective temperature;][]{2006Ap&SS.304...89R}.  As low--mass stars are faint, small, and possess complex spectra dominated by strong molecular bands, measuring these parameters is challenging. Double--lined eclipsing binaries with detached, low--mass components of similar spectral type offer the best opportunity for accurate and precise measurements of these fundamental properties.  Although binary systems are more common than single stars at high stellar masses \\citep[$>$ 1 M$_{\\odot}$; eg.][]{Abt-83,Duquennoy-91}, binaries are not very common in low--mass stars \\citep{1997A&A...325..159L,1997AJ....113.2246R,2004ASPC..318..166D}. In fact, the combination of low binary fraction and the sheer dominance (by number) of low--mass stars suggests that most of the stars in the Galaxy are actually single stars \\citep{2006ApJ...640L..63L}.  Because of the low binary fractions and the faintness of M dwarfs, relatively few low--mass main sequence binaries have been studied \\citep[e.g.][and references therein]{Creevey-05,LopezM-05,Hebb-06,Ribas-03,2006Ap&SS.304...89R,Blake-07}. Adding to the census of double--lined eclipsing binary systems is critical because they provide highly accurate, direct measurements of the masses and radii of the components, nearly independent of any model assumptions.  While the known number of eclipsing low--mass binaries is small, next-generation variability and planet hunting surveys should increase the number of known systems.  Measurements of the masses and radii of the individual components of these low--mass systems have revealed potentially serious inadequacies in stellar evolution models \\citep{Ribas-03,2006Ap&SS.304...89R}, whereas measurements and models agree for stars over $ 1\\msun$.  These discrepancies apparently extend down into the brown dwarf regime as evidenced by the first L--dwarf binary \\citep{Stassun-06, 2007ApJ...664.1154S}.  In particular, the mass-radius, mass-temperature, and mass-luminosity relationships predicted by stellar theory are inconsistent with a large fraction of observed M dwarf components \\citep[e.g.][]{Hebb-06}.  Identifying and characterizing new low--mass eclipsing binaries will provide stronger constraints for theoretical models and help reveal the cause of the current discrepancies between observed and predicted properties of low--mass stars.  In this manuscript, we present the discovery of a new double-lined eclipsing binary system, \\namefull.  In \\S 2, we outline the photometric and spectroscopic observations of this binary system.  Our determination of physical parameters and spectral types are detailed in \\S 3.  Discussion of the discrepancies between empirical and theoretical mass-radius relations is outlined in \\S 4, followed by our conclusions in \\S 5. ",
        "conclusions": "\\label{sec-conclude} We report the discovery and characterization of the double--lined eclipsing binary system \\namefull.  Photometric and spectroscopic evidence suggests that both components are M dwarfs, and we adopt classifications of M0 and M1 as their subtypes.  We resolve splitting of the H$\\alpha$ emission line with spectroscopic observations using HIRES at Keck, leading to a radial velocity curve and estimates of the masses and radii of each star.  Simulated data sets created by adding noise to the best fit provides uncertainties on and covariances between the system parameters.  We emphasize that there exist complicated degeneracies between parameters in eclipsing systems that can only be fully explored with such detailed analyses.  Our analysis is consistent with previous studies of double--lined eclipsing M--dwarf systems.  An empirical study by \\cite{2006ApJ...651.1155B} yields a quadratic mass--radius relationship spanning the range of $0.2 - 0.8 \\msun$ and spectral types from late K to late M.  This empirical relation accurately predicts the radii of \\name's components.    We observe H$\\alpha$ emission from both components, an unlikely scenario given their early--M spectral types and their distance from the Galactic plane.  If magnetic activity is enhanced in M dwarfs in binary systems, the binary population including M--dwarfs components may present an additional foreground of stellar flares for next generation time domain surveys \\citep{2006ApJ...644L..63K,2004ApJ...611..418B}."
    },
    "0801/0801.4197_arXiv.txt": {
        "abstract": "We consider the impact of thermal inflation -- a short, secondary period of inflation that can arise in supersymmetric scenarios -- on the stochastic gravitational wave background.    We show that while the primordial inflationary gravitational wave background is essentially unchanged at CMB scales, it is massively diluted at solar system scales and would be unobservable by a BBO style experiment. Conversely, bubble collisions at the end of thermal inflation can generate a new stochastic background. We calculate the likely properties of the bubbles created during this phase transition, and show that the expected amplitude and frequency of this signal would fall within the BBO range. ",
        "introduction": "In the coming decades, it is very likely that gravitational waves will be observed by direct-detection experiments. These include the presently operating LIGO and VIRGO observatories, and proposed experiments such as LISA \\cite{LISA}, BBO \\cite{BBO} and DECIGO \\cite{Kawamura:2006up}. It is likely that the first sources to be detected will be localized, transient events, but many cosmological processes yield a stochastic background of gravitational waves. These backgrounds are a key target for the nascent field of gravitational astronomy, since they probe the primordial universe in  ways that are presently impossible.  The most widely discussed cosmological signal is the inflationary gravitational wave background, generated by quantum fluctuations of spacetime. The amplitude of this signal decreases with the inflationary energy scale, rendering it effectively invisible in models where inflation occurs well below the GUT scale.  A stochastic background of gravitational waves can also be generated whenever the universe undergoes a phase transition. Unlike the inflationary background, which is scale invariant, these backgrounds will be peaked at a wavelength correlated with the characteristic scale of the transition. In these scenarios, the gravitational waves are sourced  by the non-zero quadrupole (and higher) moments of the matter distribution \\cite{WeinbergGRBook}.  Recently, several groups have investigated the possibility that gravitational waves are produced during the preheating phase at the end of inflation \\cite{GarciaBellido:2007af,Easther:2006gt,Easther:2006vd,Khlebnikov:1997di,Dufaux:2007pt, GarciaBellido:2007dg,Easther:2007vj}.  The momentum modes of fields coupled to the inflaton undergo resonant pumping, which renders the universe highly inhomogeneous, and sources the emission of gravitational radiation.   Intriguingly, this signal is  {\\em easier\\/} to detect in models where the inflationary scale is substantially below the GUT scale, as first argued in \\cite{Easther:2006gt}, and confirmed in \\cite{Easther:2006vd}.   Moreover, it probes the mechanism that \\emph{ends\\/} inflation, which is normally assumed to be inaccessible to direct observational tests.  In addition, a second (and not entirely unrelated) mechanism could generate a significant background -- namely the collision of percolating bubbles formed during a first order phase transition   \\cite{Kosowsky:1991ua,Turner:1992tz,Kosowsky:1992rz,Kosowsky:1992vn, Kamionkowski:1993fg,Caprini:2007xq}. As we will see, it has the same scaling properties as the preheating signal -- namely the amplitude need not depend on the energy scale, but its (present-day) wavelength decreases linearly with this scale. Assuming an energy scale somewhere between the electroweak and GUT scales, the Hubble scale will lie between $10^{-5}\\eV$ and $10^{14}\\GeV$. We will see that if the characteristic bubble radius is within a few orders of magnitude of the Hubble scale at which they form,  the resulting gravitational wave background will have a present-day frequency in the range $10^{-3}$ to $10^{9} \\Hz$ \\cite{Easther:2006gt,Kamionkowski:1993fg}.  To predict the primordial background, one must know the transfer function, which is determined by the effective equation of state of the universe between the epoch when the signal is generated and the present day. A further complication is that there is no guarantee that inflation happens only once in  the early universe. One may imagine primary and secondary periods of inflation -- where the primary period is responsible for solving the usual cosmological initial conditions problems and laying down the perturbations that seed large scale structure, and the secondary period lasts for $\\leq \\mathcal{O}(10)$ $e$-foldings, and occurs at some energy scale below that of the first inflationary period and above (or at) the scale at which baryogenesis takes place.  If significantly more than  $\\mathcal{O}(10)$ $e$-folds occurs in the secondary phase, it will erase all relics of the previous phase, returning us to what is effectively a single phase model.  In addition to modifying the mapping between scales in the present epoch and the moment these (comoving) scales leave the horizon during the primary phase, a secondary phase of inflation also erases the primordial inflationary background of gravitational waves at wavelengths which have reentered the horizon before the secondary phase begins.  This leads to a situation where a high scale primary phase would generate a gravitational wave background visible in the B-mode of the CMB polarization, but completely \\emph{invisible\\/} to direct detection experiments operating at solar system scales, such as BBO.  A cosmological model with multiple inflationary phases may appear  baroque. However \\emph{thermal inflation} \\cite{Lyth:1995hj,Lyth:1995ka} is a secondary period of inflation that arises in  supersymmetric theories when some (almost) flat direction with negative mass-squared at the origin -- which we will call the flaton --  acquires a thermal contribution to its effective potential, giving it a positive effective mass.   The flaton is thus pinned at the origin, and the finite potential energy at this point drives a secondary period of inflation. Once thermal inflation sets in, the temperature falls rapidly and when it drops below the flaton's mass scale (usually assumed to be a typical soft supersymmetry breaking mass scale $\\sim \\TeV$) the flaton rolls aways from the origin, and thermal inflation naturally shuts off.   The total amount of inflation is typically small, about 10 $e$-folds or less.  Thermal inflation requires only a flat direction with negative mass-squared at the origin, both of which are natural in supersymmetric models.  From the model building perspective, its importance lies in its ability to dilute the moduli and gravitinos, which are generically produced after the primary period of inflation in supersymmetric theories. These particles either decay, wrecking the successful predictions of Big Bang Nucleosynthesis, or provide a dark matter density that massively overcloses the present universe \\cite{Coughlan:1983ci,Banks:1993en,de Carlos:1993jw,Khlopov:1984pf,Ellis:1984eq}.  These particles are produced before thermal inflation begins and are too weakly coupled to be produced after it ends, so their contribution is thus diluted to safe levels.  Thermal inflation is perhaps the best-motivated solution to the moduli problem, which is endemic to supersymmetric models of the early universe, including many of those derived from string theory. Thermal inflation has a relatively low reheating scale (typically $\\mathcal{O}(10) \\GeV$) but it can give rise to baryogenesis \\cite{Jeong:2004hy,Felder:2007iz}.  Unfortunately, as we will show in Section~\\ref{sect:forecasts}, thermal inflation wipes out any primordial inflationary gravitational wave signal in the range of frequencies accessible to BBO \\footnote{Indeed it also wipes out the primordial \\emph{density} spectrum at the same frequencies, although the scales affected are not important cosmologically and hence not constrained by large scale structure observations.}. Thus, even if the primary phase of inflation occurs at a  high enough scale to produce a detectable tensor background, thermal inflation erases it at short scales \\cite{Mendes:1998gr}. If this was the end of the story, one might regard it as an unfortunate side-effect of solving the moduli problem, but there is a final piece to the puzzle: thermal inflation ends via bubble nucleation, and the collisions between these bubbles as they percolate produces a significant stochastic gravitational wave background. We will see that this process may generate $\\Omega_\\mathrm{GW} h^2 (f \\approx 1 \\Hz) \\sim 10^{-17}$, and is potentially detectable by BBO.  Thermal inflation is a complicated scenario, at least when compared to simple toy models of inflation. However, it relies on generic ingredients found in many supersymmetric or stringy models of the early universe; it is the intrinsic richness of these models that generates their complicated phenomenology. In particular, the moduli problem is robust enough to have survived 25 years. Conversely, thermal inflation does not require any exotic physics, unnatural parameter values or initial conditions, and the complexity of this scenario may in fact be typical of realistic models of the early universe. Finally, while the primordial gravitational background is erased, it would need to have been near the top of its plausible range to have been observable even without dilution, and furthermore in supersymmetric and stringy models the bottom of this range is arguably more natural \\cite{Randall:1995dj,Kadota:2003fs,Kadota:2003tn}. The bubble collision signature is generic to thermal inflation and effectively expands the list of the models probed by BBO.  In Section 2 we summarize thermal inflation, and describe the process of bubble nucleation that occurs as it ends. In Section 3 we describe the production of gravitational waves  by bubble collisions and estimate the expected spectrum. In Section 4 we compute the suppression to the primordial inflationary background by a secondary period of inflation, combine this with the likely spectrum from bubble collisions, and assess the  prospects for detecting this signal in future observatories, the likelihood it can be distinguished from other stochastic astrophysical backgrounds, and the likely tensor signal  from the primary inflationary phase, which would be undiluted at CMB scales, or via pulsar timing experiments \\cite{Allen:1996vm}. Finally we sum up in Section 5. ",
        "conclusions": "This analysis draws on a  number of apparently disparate topics in theoretical cosmology. Firstly, it looks carefully at thermal inflation, considering both its expansion history and the properties of the bubbles created by the phase transition that marks its end. Historically, the primary motivation for considering thermal inflation is that it provides a natural solution to the moduli problem, which is endemic to supersymmetric models and arises from the production of a condensate of long-lived moduli particles which either disrupt nucleosynthesis as they decay or massively overclose the universe. Thermal inflation dilutes these unwanted particles to acceptable levels.   However, we show that thermal inflation also dilutes the {\\em primordial\\/} gravitational wave background at small scales, rendering it effectively unobservable by direct detection experiments such as BBO, but preserving any signal that is present in the CMB B-mode.  The detailed dependence of the BBO signal on the effective equation of state in the primordial universe has been discussed elsewhere (e.g.\\  \\cite{Boyle:2005se}) but the dilution factor implied by thermal inflation would effectively erase any primordial signal.  Conversely, thermal inflation ends via a first order phase transition, so the percolation of the resulting bubbles can generate a new background of gravitational waves.  This spectrum is confined to a few decades in frequency -- unlike the scale-free inflationary background -- and is potentially detectable by BBO style experiments.  While the detailed mechanism is different, this signal has much in common with the gravitational wave background that can be generated during preheating or parametric resonance \\cite{GarciaBellido:2007af,Easther:2006gt,Easther:2006vd,Khlebnikov:1997di, Dufaux:2007pt,GarciaBellido:2007dg,Easther:2007vj}.  Further study will be needed to see if the characteristic spectra can be easily distinguished from one another.  The analysis here only considers the background generated by the bubble collisions themselves -- a period of turbulent mixing after the bubble collisions would also source gravitational waves, and this signal has not yet been computed. It has the potential to significantly enhance the gravitational wave signal, since the bubble nucleation phase is followed by a period of flaton domination, mimicking a matter dominated universe, and diluting the background produced by colliding bubbles. If the turbulence lasts well into effective matter dominated phase, any gravitational waves sourced via the turbulent mixing will be less strongly diluted. Consequently, the spectrum computed here represents a lower bound on the total signal. Further, the background will be strongly inhomogenous on small scales due to the bubble collisions and other preheating mechanisms. It is unclear exactly how efficient this preheating is and so how much it changes the effective equation of state, hence the last factor in \\eq{eqn:energydensity} which may be significant. Thus, the suppression factor could again be reduced, leading to a stronger gravitational wave signal in the present epoch. These questions will most likely be settled via direct numerical simulations, and we will tackle these projects in future work."
    },
    "0801/0801.1298_arXiv.txt": {
        "abstract": "Over a timescale of a few years, an observed change in the optically thick radio continuum flux can indicate whether an unresolved \\HII region around a newly formed massive star is changing in size. In this Letter we report on a study of archival VLA observations of the hypercompact \\HII region G24.78+0.08 A1 that shows a decrease of $\\sim 45$ $\\%$ in the 6-cm flux over a five year period. Such a decrease indicates a contraction of $\\sim 25$ $\\%$ in the ionized radius and could be caused by an increase in the ionized gas density if the size of the \\HII region is determined by a balance between photoionization and recombination. This finding is not compatible with continuous expansion of the \\HII region after the end of accretion onto the ionizing star, but is consistent with the hypothesis of gravitational trapping and ionized accretion flows if the mass-accretion rate is not steady. ",
        "introduction": "\\label{intro}  The formation of massive stars ($M>8$ $\\Msun$) by accretion presents a number of theoretical difficulties; among them, that once a star attains a sufficient mass its surface temperature is high enough to produce a small \\HII region (see the reviews on ultracompact and hypercompact \\HII regions by Churchwell 2002; Kurtz 2005; and Hoare et al. 2007). The thermal pressure differential between the hot ($\\sim 10^4$ K) ionized gas and the cold ($\\sim 100$ K) molecular gas can potentially reverse the accretion flow of molecular gas and prevent the star from ever reaching a higher mass. However, recent models that include the effects of gravity \\citep[e.g.,][]{Keto07} have shown that if the \\HII region is small enough (ionized diameter $\\leq 1000$ AU for total stellar mass $\\geq 350$ $\\Msun$), the gravitational attraction of the star(s) dominates the thermal pressure and the molecular accretion flow can cross the ionization boundary and proceed toward the star as an ionized accretion flow within the \\HII region. In this stage the \\HII region is said to be gravitationally trapped by the star.   Previous studies have demonstrated a number of observational techniques bearing on the evolution of small \\HII regions around newly formed stars. Accretion flows onto and through \\HII regions can be directly observed by mapping molecular and radio recombination lines (RRLs) at very high angular resolution \\citep[e.g.,][]{Sol05, KW06}. If the \\HII region is too small to be spatially resolved, the frequency dependence of the  velocities and widths of RRLs can be used to infer steep density gradients and supersonic velocities within the \\HII regions \\citep{KZK07}, as expected from the presence of ionized accretion flows or bipolar outflows. Radio continuum observations at very high angular resolution, made at two epochs some years apart, can also be used to directly observe changes in size of small \\HII regions (Franco-Hernandez \\& Rodriguez 2004). Changes in size indicate whether the evolution of the \\HII region is consistent with pressure-driven expansion or gravitational trapping.  In this Letter we demonstrate yet another technique -- how a comparison of radio continuum observations at different epochs can be used to infer size changes in \\HII regions even if the observations do not spatially resolve them. At optically thick frequencies the radio continuum flux depends to first approximation only on the size of an \\HII region and is independent of its internal density structure. Therefore, even if the size is not known, a change in the flux over time still indicates a change in the size of the \\HII region. This technique is particularly valuable because it can be used on the smallest and youngest \\HII regions, and because lower angular resolution radio continuum observations generally require less observing time than spectral line and high angular resolution observations.  For the study presented in this Letter we selected the hypercompact (HC) \\HII region G24.78+0.08 A1 which lies at the center of a massive molecular accretion flow \\citep{Bel04, Bel06} and also has multi-epoch 6 cm radio continuum observations in the VLA archive. G24.78+0.08 was detected in the centimetric (cm) continuum by \\cite{Beck94}, and later resolved by \\cite{Cod97} into a compact (A) and an extended (B) component. A millimeter (mm) interferometric study \\citep[]{Bel04} revealed the presence of two massive rotating toroids centered in respective dust cores (A1 and A2). The compact cm emission comes from the mm component A1 (hereafter G24 A1), and has recently been resolved by \\cite{Bel07}. If G24 A1 is ionized by a single star, its spectral type should be earlier than O9 \\citep[]{Cod97}. Also, G24 A1 likely powers a massive CO outflow \\citep[]{Fur02}.  The infalling and rotating molecular gas and  bipolar outflow all suggest ongoing accretion. However, based on proper motions of H$_2$O masers around the \\HII region, \\cite{Bel07} and \\cite{Mos07} proposed that at the present time, the \\HII region is expanding into the accretion flow.  The suggested timescale for the expansion is short enough that we should be able to detect a corresponding increase in the optically thick radio continuum flux within a few years. ",
        "conclusions": "Our analysis of archival VLA observations of the \\HII region G24.78+0.08 A1 indicates a contraction of its radius between 1984.36 and 1989.48. This finding is consistent with the hypothesis that this HC \\HII is the inner, ionized part of the larger scale accretion flow seen in the molecular line observations of \\citet{Bel04,Bel06}. Future high-resolution RRL observations could unambiguously resolve the ionized-gas kinematics and confirm this hypothesis."
    },
    "0801/0801.3254_arXiv.txt": {
        "abstract": "We discover a faint shell-type radio and X-ray source, G353.6-0.7, associated with HESS J1731-347. G353.6-0.7 is likely an old supernova remnant (SNR), based on radio (0.8 GHz, 1.4 GHz and 5 GHz), infrared (8 $\\mu$m from the GLIMPSE Legacy Project and 21 $\\mu$m from the Midcourse Space Experiment), and X-ray (0.1 keV - 2.4 keV from the ROSAT survey and 5 - 20 keV from the INTEGRAL survey) data. The SNR, centered at ({\\sl l}, {\\sl b})=(353.55, -0.65) with a radius of $\\sim$ 0.25$^{\\circ}$, closely matches the outline of the recently discovered extended TeV source HESS J1731-347, which has no previously identified counterpart. A diffuse X-ray enhancement detected in the ROSAT all-sky survey is coincident with lower half shell of the SNR. Therefore the SNR is the best radio counterpart of both the HESS source and the diffuse X-ray enhancement. G353.6-0.7 has an age of $\\sim$ 27000 yrs.  Altogether, the new discovery provides the best case that an old SNR emits TeV $\\gamma$-rays. : ",
        "introduction": "About 52 Galactic TeV $\\gamma$-ray sources have been detected by ground-based high energy telescopes (e.g. the High Energy Stereoscopic System (HESS), Aharonian et al. 2006; the MAGIC telescopes, Albert et al. 2006) so far. Some of them do not have counterparts in other wavebands. Seeking their Galactic counterparts in other wavebands could play a key role for us to understand the origin of these sources and distinguish different acceleration mechanisms to emit TeV $\\gamma$-rays. Among identified counterparts, young shell-type supernova remnants (SNRs) are one of the four Galactic populations to generate very high energy (VHE) $\\gamma$-rays, i.e. pulsar wind nebulae (PWN), X-ray binaries, shell SNRs and a young stellar cloud Westerlund 2 (Aharonian et al. 2007, Landi et al. 2006, Torres et al. 2003, Enomoto et al. 2002). Recently, the old SNR W41 has been proposed to associate with the extended TeV source HESS J1834-087 based on observations (Tian et al. 2007; Leahy \\& Tian 2008), which is theoretically expected (Yamazaki et al. 2006, Fang \\& Zhang 2008). However, this is the only likely association of an old SNR with a TeV source up to now. We report here the discovery of a new faint radio and X-ray shell which closely matches the extended TeV source HESS J1731-347 without previously identified counterpart (Aharonian et al. 2008). We use new radio, infrared and X-ray data to identify the shell as an old SNR, therefore obtaining second association of a TeV source with an old SNR. ",
        "conclusions": "\\subsection{An old SNR} SNR G353.6-0.7 has a large angular size and a large physical size ($\\sim$ 28 pc in diameter for a likely distance of $\\sim$ 3.2 kpc), and faint radio and X-ray emissions. This suggests that G353.6-0.7 is quite old, because young SNRs ($\\le$ 2000 yrs) usually have a size $<$ 10 pc, and produce strong radio and X-ray emission. Applying the Sedov model (Cox 1972), we estimate its age based on the known radius R (14 pc) and ambient density $n_{0}$ assuming a typical explosion energy (10$^{51}$ erg). The likely absorption of X-rays from the western half of G353.6-0.7 hints that ambient density $n_{0}$ is not very low. The density $n_{0}$ can be estimated from the $\\gamma$ flux of HESS J1731-347 ($\\it{F_{\\gamma}}$$\\approx$1.2$\\times$10$^{-11}$ cm$^{-2}$ s$^{-1}$ which is given by $dF_{\\gamma}/dE(>E_{min}) = N_{0} \\left(\\frac{E}{1\\:{\\rm TeV}}\\right)^{-\\Gamma}$. Here $E_{min}$=0.5 TeV, $\\Gamma$=2.26$\\pm$0.10, $N_{0}$=6.1$\\pm$0.8 $\\times$10$^{-12}$ cm$^{-2}$ s$^{-1}$. See Aharonian et al. 2008). Using the formulae 1 and 2 of Aharonian et al. (2006), we obtain the $n_{0}$ of $\\sim$ 5 cm$^{-3}$ (assuming a conversion efficiency of 0.1 - 0.2 and explosion energy 10$^{51}$ erg). If G353.6-0.7 is still in the Sedov-Taylor phase, it has a Sedov age of $\\sim$ 25000 yrs.  However, the radius at which cooling affects the dynamics of an SNR is $\\sim$ 10 pc for $n_{0}$ $\\approx$ 5 cm$^{-3}$, so G353.6-0.7 has already entered the radiative phase. So we here obtain a radiative age of $\\sim$ 27000 yrs for G353.6-0.7. %  \\begin{figure} \\vspace{45mm} \\begin{picture}(80,80) \\put(322,-15){\\special{psfile=f3.ps hoffset=0 voffset=0 hscale=45 vscale=43 angle=90}} \\end{picture} \\caption[xx]{The 1420 MHz image of G353.6-0.7 with ROSAT image contours (red: 0.11, 0.19, 0.23, 0.8 in units of 10$^{-3} counts s^{-1} pixel^{-1}$) and HESS J1731-347 image contours (green) overlaid. The HESS J1731-347 image has same contours as Fig. 5 of Aharonian et al. (2008), i.e. starting at $4\\sigma $ in $1\\sigma $ steps, which is of $\\gamma$-ray excess counts smoothed with a Gaussian filter with standard deviation $0.1^\\circ $}. % \\end{figure}  \\subsection{Radio, X-rays and TeV $\\gamma$-rays from G353.6-0.7} The observations from radio, X-ray and $\\gamma$-ray all indicate that SNR G353.6-0.7 is the counterpart of the newly detected TeV $\\gamma$-ray source HESS J1731-347 (see Fig. 3). The TeV $\\gamma$-rays overlap the entire radio SNR, including the east half seen in X-rays and the west half not seen in X-rays. The $^{12}$CO column density increases toward the Galactic plane, consistent with absorbing the X-rays from the west half of the SNR, but allowing both radio and TeV $\\gamma$-rays to be observed. Old SNRs emitting TeV $\\gamma$-rays are rare. W41 has been considered as the only likely candidate so far (Tian et al. 2007). Young SNRs found to show $\\gamma$-ray emission have $F_{\\gamma(1-10TeV)}/F_{X(2-10 keV)}$ of typically below $\\sim$2, $F_{\\gamma(1-10TeV)}/F_{radio(0.01-100GHz)}$ of typically below $\\sim$10 (Yamazaki et al. 2006).  The flux density of G353.6-0.7 is S$_{1420MHz}$$\\approx$ 2.2 Jy, giving an integrated flux from 10$^{7}$ to 10$^{11}$ Hz of F$_{radio}$=5.2 $\\times$ 10$^{-13}$ ergs cm$^{-2}$ s$^{-1}$, assuming a radio spectral index of $\\sim$ 0.5. We calculate the $\\gamma$-ray flux of F$_{\\gamma(1-10TeV)}$ $\\approx$ 1.7$\\times$ 10$^{-11}$ ergs cm$^{-2}$ s$^{-1}$ in the 1-10 TeV band, so the ratio R = F$_{\\gamma}$/F$_{radio}$ is $\\sim$ 33. Also given a distance of 3.2 kpc to G353.6-0.7, the $\\gamma$-ray luminosity is $\\sim$ 2$\\times$10$^{34}$ ergs s$^{-1}$ in the 1 - 10 TeV range, which is larger than that of RX J1713.7-3946, one of the intrinsically brightest galactic TeV Gamma-ray sources discovered and thought to originate from pion decays. So, SNR G353.6-0.7 is a powerful cosmic ray accelerator.  The theoretically predicted $\\gamma$-rays to X-rays flux ratio for old SNRs is possibly $> \\sim$ 100 (Yamazaki et al. 2006), giving an upper limit to 2-10 keV X-ray emission from G353.6-0.7 of about 1.7$\\times$ 10$^{-11}$ ergs cm$^{-2}$ s$^{-1}$. We checked the INTEGRAL survey which covers G353.6-0.7 and found no detection within the SNR (JEM-X in energy range of 5 - 20 keV). Assuming an individual INTEGRAL observation for an integration time of 2 ks, the 5 $\\sigma$ source detection limit is 1.2$\\times$ 10$^{-10}$ ergs cm$^{-2}$ s$^{-1}$ on-axis in the energy range of 2 - 10 keV (Brandt et al. 2003). This is consistent with the derived X-ray upper limit flux from the old SNR.  The fact that the X-rays are visible in SNR G353.6-0.7 shows that SNR shock velocity is still fast. This means that G353.6-0.7 is not too old (i.e. $<$10$^{5}$ yrs), because an old SNR is expected to emit few thermal X-rays which originate from the high temperature gas heated by the fast-enough shock.  Because of poor statistics in the ROSAT X-ray spectrum, we can not exclude that part of the X-rays  are from synchrotron emission. Because no pulsar/PWN is detected within the remnant until now by radio and X-ray observations, it is unlikely that the TeV $\\gamma$-rays originate from particles accelerated to high energy at a termination shock formed in the pulsar wind. This strengthens our conclusion that the TeV $\\gamma$-rays originate from the old SNR.  TeV $\\gamma$-rays can originate from an old SNR or from a molecular cloud encountered by the SNR, via pion decay with the pions produced by proton-proton collisions (Yamazaki et al. 2006). CO observations can detect all CO cloud components in the direction of G353.6-0.7, therefore help to confirm and distinguish different acceleration mechanisms for emission of TeV $\\gamma$-rays. E.g. if there was a CO cloud overlapping the radio, X and $\\gamma$-ray emissions at the same distance as SNR G353.6-0.7, and if a broadened CO emission profile was detected, indicating a shock interaction, it would be a strong argument to support the SNR interacting with the cloud to emit TeV $\\gamma$-ray emission. However, the existing $^{12}$CO observations in the direction of G353.6-0.7 have too low resolution (1/8 degree) to give any constraint regarding SNR-CO cloud interaction. Further high-resolution CO and deep X-ray observations can be used to reveal the real radiative mechanism for TeV emission.  In summary, this new discovery of a radio SNR overlapping the TeV source HESS J1731-347 provides the best case that an old SNR emits TeV $\\gamma$-rays."
    },
    "0801/0801.3638_arXiv.txt": {
        "abstract": "Optical spectra obtained in 2006-07 of the nova-like cataclysmic variable QU~Car are studied for radial velocities, line profiles, and line identifications. We are not able to confirm the reported 10.9 hr orbital period from 1982, partly because our sampling is not ideal for this purpose and also, we suspect, because our radial velocities are distorted by line profile changes due to an erratic wind.  P-Cygni profiles are found in several of the emission lines, including those of C IV.  Carbon lines are abundant in the spectra, suggesting a carbon enrichment in the doner star.  The presence of [O III] 5007$\\AA$ and [N II] 6584$\\AA$ is likely due to a diffuse nebula in the vicinity of the system.  The wind signatures in the spectra and the presence of nebular lines are in agreement with the accretion wind evolution scenario that has been suggested to lead to SNeIa.  We argue that QU~Car is a member of the V~Sge subclass of CVs, and a possible SNeIa progenitor.  It is shown that the recent light curve of QU~Car has $\\sim$1 mag low states, similar to the light curve of V~Sge, strengthening the connection of QU~Car with V~Sge stars, supersoft x-ray sources, and SNeIa progenitors. ",
        "introduction": "Although identification of the progenitors of Supernovae Ia (SNeIa) remains controversial, it is accepted that they originate in binary systems in which at least one component is a white dwarf (WD); those systems are grouped under the wide umbrella of cataclysmic variables (CVs). CVs are semi-detached binaries in which a WD is accreting material from its lower main sequence or evolved companion (Warner 1995). When the magnetic field of the WD is low, a luminous accretion disk leads the gas onto the WD; if the magnetic field of the WD exceeds $\\sim$10MG, gas is guided onto the WD's magnetic poles via its magnetic field lines. Depending on the orbital period, the mass transfer rate ($\\dot{M}$), the white dwarf temperature and the donor star's nature, CVs present a rich group of members, each providing valuable information on accretion physics, stellar evolution, binary star interaction and cataclysmic explosions.  Current theories for SNeIa progenitors hold that, either via Roche lobe overflow of the companion or via a wind, the WD accumulates H or He-rich material which is then burned to C and O on the WD. Under suitable conditions (determined primarily by the stability of $\\dot{M}$ onto the WD) the WD mass reaches the Chandrasekhar limit initiating a series of thermonuclear reactions eventually leading to a SNeIa (Hillebrandt \\& Niemeyer 2000). Although outlining this scenario is rather straightforward, the specifics are far from being understood or defined. For example, the nature of the companion star is uncertain, as well as how $\\dot{M}$ remains low enough to prevent the expansion of the star's outer layers but high enough to allow stable H or He burning on the WD atmosphere without leading to a nova explosion. The WD mass necessary to initiate the thermonuclear reactions leading to the explosion is about 1.4M$_{\\sun}$ and recent theoretical models conclude that a fast rotating WD with initial mass as low as 0.7M$_{\\sun}$ can accrete efficiently and eventually reach the Chandrasekhar limit (Hachisu, Kato \\& Nomoto 2007).  Currently, candidate progenitors are either double-degenerate (WD+WD) systems, or binaries that harbor a post main sequence companion (the single degenerate binary, or SDB, scenario). Double-degenerate systems are less promising candidates due to their low overall mass or long orbital periods (Parthasarathy et al. 2007). Favorable candidates are the SDBs (Langer et al. 2000) for which stellar evolution codes predict the presence of substantial outflows; the outflows result from H- or He-rich material that is processed on the WD and expelled under appropriate conditions (e.g. Han \\& Podsiadlowski 2004). Badenes et al. (2007) predict that the presence of stable, high-velocity winds (v$_{outflow}$$\\ge$1000km/sec) should lead to the formation of cavities in the surrounding ISM. However, such cavities are not observed around existing SNeIa remnants. Badenes et al. (2007) suggest that if outflows exist from SNeIa progenitors, they should be irregular and variable in time and should not be able to affect the surrounding pre-SN environment.  Among the most promising SNeIa progenitors are some types of transient supersoft x-ray sources (SSS), recurrent novae, and V~Sge-type CVs. The latter is a small category of semi-detached binaries with high mass transfer rates  ($\\sim$10$^{-7}$M$_{\\sun}$yr$^{-1}$) allowing nuclear burning on the white dwarf as mass transfer proceeds; these are considered to be the galactic counterparts of extragalactic SSS sources. Hachisu $\\&$ Kato (2003; hereafter HK03) model the long-term optical light curve of the prototype of the class, V~Sge, and correlate the observed bright/faint states ($\\sim$1 mag in amplitude) of the system with the predictions of their accretion wind evolution scenario (AWE). In short, the mass of the WD increases until it reaches a critical limit at which time the WD burns the accumulated H and expels a large portion of the processed material (primarily C and O) in a fast wind. This wind sweeps away the disk surface increasing its luminosity (leading to an overall increase of the system brightness by up to 1 mag) and strips off the outer layers of the secondary star which becomes smaller than its Roche lobe; this halts mass transfer. The disk is then accreted onto the WD. Mass accretion eventually stops, leading to a cessation of the wind phase. Mass transfer starts again with the donor star regaining contact with its Roche lobe; this is also the beginning of a new cycle. At the same time, the white dwarf accumulates enough material to eventually reach the Chandrasekhar limit and finally explode as a SNeIa.  In this paper we present a spectroscopic study of QU~Car, an unusual yet bright and understudied CV. QU~Car was first reported in Stephenson et al. (1968) as an irregular variable HD310376, similar to the bright low-mass X-ray binary Sco X-1 and the old nova HR Lyr. QU~Car shares with Sco X-1 rapid and erratic photometric variations, as well as marked changes in the emission line spectrum.  However, in contrast to Sco X-1, QU~Car sometimes has He absorption features and relatively weak x-ray emission. (Schild 1969). In the optical, the system is non-eclipsing and exhibits flickering of 0.1-0.2 mag on timescales of 1-10 min (Schild 1969; Gilliland $\\&$ Phillips 1982, hereafter GP82). In our study, we present new optical spectra suggesting that QU~Car could be a member of the V~Sge category and a likely SNeIa progenitor. The following section presents our data and data reduction technique, followed by the results of our analysis in $\\S3$ and our discussion in $\\S4$. We summarize our findings in $\\S5$. ",
        "conclusions": "We have presented a new spectroscopic study of the bright nova-like QU~Car. The highlights of this work are:  \\begin{itemize}  \\item We explored the orbital period of the system using He II radial velocities from this study and from GP82.  We are unable to confirm the 10.9 hr period of GP82.  We argue that this situation is lilely due to rapid variations (compared to orbital) in the line profiles (P-Cygni) of the He II line due to an erratic wind, combined with our relatively short time series.  We use as support for this argument the fact that the wind variability time scales in NL CVs are in this range, and the presence of fast monotonic RV changes in our data as well as in GP82.  \\item Forbidden lines are detected in the QU~Car spectrum, consistent with substantial mass outflow.  The [O III] 5007 appears split, likely due to contributions from the front and back sides of an expanding flow.  \\item It is pointed out that the light curve of QU~Car has low states very similar to those in V~Sge.  In V~Sge these low states are taken as evidence for the operation of accretion wind evolution cycle that is thought to lead to SNeIa.  We propose that QU~Car is also a strong candidate for a SNeIa progenitor, based on this and other similarites to V~Sge systems. These similarites include high $\\dot{M}$ and evidence for a strong wind from P Cygni profiles and nebular lines.  \\end{itemize}  It is surprising that so little optical work has been done on this bright, important system.  It is important to determine/refine the orbital period of QU~Car.  Photometric monitoring is needed to define the characteristics of the photometric cycle and allow accurate comparisons with the predictions of the accretion wind model. Higher resolution spectral data will provide information on the various ionization regions and possible emission lines from the nebula.  High resolution corongraphic imaging should be employed to search for nebulosity. QU~Car appears to be the brightest known member of the V~Sge category, and a  likely SNeIa progenitor.  Further study can provide important information on the nature and the conditions that lead to mass accumulation onto the white dwarf during the wind accretion cycle."
    },
    "0801/0801.2122_arXiv.txt": {
        "abstract": "Self-consistent stellar models including all effects of atomic diffusion and radiative accelerations as well as mass loss are evolved from the pre main sequence for stars of 1.35-1.5\\,M$_{\\odot}$ at solar metallicity (Z=0.02). A mass loss rate similar to the solar mass loss rate is sufficient to reproduce observations of the star $\\tau$UMa. We discuss the effect of mass loss on the iron convection zone that naturally appears beneath the main hydrogen convection zone of these stars. We also find that the effects of mass loss can be distinguished from those caused by   turbulence, but are nevertheless able to explain the particularities of the AmFm phenomenon. ",
        "introduction": "\\label{intr}  Since 1970 it is generally agreed that atomic diffusion driven by radiative accelerations plays a role in creating the anomalous surface abundances of F, A and B stars (Michaud 1970). However, some 40 years later, many questions remain as to the exact behavior of many physical processes within the stable envelopes of these stars. In fact, for Fm stars, two competing scenarios which have each had their share of success are presently being confronted to observations. The  ``classical'' scenario Watson (1971) suggests that separation occurs below the H-He convection zone. In this framework, models which only consider atomic diffusion without extra mixing generate predicted anomalies that are 3-5 times larger than the ones observed (Turcotte et al. 1998), thus implying that there is at least one competing process that slows the effects of separation. This lead to more recent models (Richer et al. 2000, Richard et al. 2001) in which turbulence enforces mixing down to about 200\\,000\\,K. In these models, it is implied that separation occurs deeper in the star.  Like turbulence or rotation, mass loss is another macroscopic process that can reduce inhomogeneities on the surface of these stars. However, until now, only static stellar models have included the effects of mass loss (Michaud et al. 1983, Alecian 1996 for Ca and LeBlanc et Alecian 2007 for Sc) with the latter paper showing that the actual depth at which separation occurs is still uncertain.  With self-consistent models of Fm stars (6500\\,K\\,$\\leq$\\,\\teff\\,$\\leq$\\,7000\\,K) we will show that mass loss can reduce predicted abundances to the observed levels. The first aim is to see to what extent observations can constrain the importance of mass loss and if its effects can be deciphered from the ones encountered with turbulence. We will also discuss the implications of our models on the depth of chemical separation. ",
        "conclusions": "Our results seem to suggest that the scenario in which separation is to take place  at $T \\sim 2\\times 10^5$\\,K in Fm stars must be favored over the classical scenario. In this framework, a mass loss rate of the order of the solar mass loss rate is able to reduce predicted anomalies to the observed abundances of $\\tau$UMa. However, it is too early to eliminate the possibility of separation below the H convection zone. Abundance determinations between Al and Ar could help distinguish between these two regimes. It is also shown that turbulence and mass loss affect anomalies differently, though the discrepancy is slight in the models shown. Once again, more observations are required to further constrain these two mechanisms. More massive models in the \\teff\\,\\,range where observations are not as scarce are also needed. They are currently being calculated. In any case, it is shown that reasonable mass loss rates can effectively reduce the anomalies predicted by atomic diffusion models to the observed levels."
    },
    "0801/0801.0582_arXiv.txt": {
        "abstract": "The TRACER instrument (``Transition Radiation Array for Cosmic Energetic Radiation'') has been developed for direct measurements of the heavier primary cosmic-ray nuclei at high energies. The instrument had a successful long-duration balloon flight in Antarctica in 2003. The detector system and measurement process are described, details of the data analysis are discussed, and the individual energy spectra of the elements O, Ne, Mg, Si, S, Ar, Ca, and Fe (nuclear charge Z=8 to 26) are presented. The large geometric factor of TRACER and the use of a transition radiation detector make it possible to determine the spectra up to energies in excess of 10$^{14}$ eV per particle. A power-law fit to the individual energy spectra above 20 GeV per amu exhibits nearly the same spectral index ($\\sim$ 2.65 $\\pm$ 0.05) for all elements, without noticeable dependence on the elemental charge Z. ",
        "introduction": "\\label{sec:intro}  The energies of cosmic rays observed near earth extend over a very wide range, from below 10$^8$ eV to more than 10$^{20}$ eV per particle.  Up to about 10$^{10}$ eV per amu, the particle energies and intensities are significantly affected by solar modulation, but for the nearly ten remaining decades, the arriving particles are believed to represent the ambient cosmic-ray population in the local galaxy. Over this range, the total cosmic-ray intensity (the ``all particle'' differential energy spectrum, see, for instance, the compilation of \\citet{cronin97}) decreases monotonically by more than 30 orders of magnitude, initially as a power law $\\propto$ E$^{-2.7}$, steepening slightly to E$^{-3.0}$ at the ``knee'' above 10$^{15}$ eV per particle, and exhibiting small changes in slope again (a ``second knee'' and an ``ankle'') in the 10$^{18}$ -- 10$^{19}$ eV region. Besides these relatively minor changes, the overall spectrum is remarkably featureless, even though a variety of processes may contribute to the cosmic-ray flux. The general consensus is that the particles below the knee are generated inside our Galaxy and are contained by the galactic magnetic fields for millions of years, while at the highest energies ($\\geq$ 10$^{18}$ eV) the gyration radii become so large that galactic containment is no longer effective, and that the cosmic rays may then be of extragalactic origin.  The all-particle energy spectrum refers to the overall spectrum of cosmic rays without differentiating the individual components.  To obtain deeper insight, details of the composition of the cosmic-ray particles must be studied.  For instance, very accurate measurements of not only the elemental, but also the isotopic composition are now available at low energies ($\\leq$ 10$^{9}$ -- 10$^{10}$ eV per particle).  These have determined such important quantities as the average containment time of cosmic rays in the galaxy from observations of radioactive clock-nuclei ($\\tau \\approx 1.5 \\times 10^{7}$ y at these energies (first reported by \\citet{gmunoz75}; more recent work by \\citet{yanasak01}), and they have excluded fresh supernova ejecta as source material of cosmic rays from the abundances of Co- and Ni-isotopes \\citep{wiedenbeck99}.  At relativistic energies only the abundances of the elemental species, without isotopic detail, have been accessible to measurements (e.g. \\citet{engelmann90}), but even these encounter increasingly severe systematic and statistical uncertainties as the energy increases.  Very few direct observations (e.g. \\citet{muller91}) have provided composition detail with single-element resolution above ~10$^{13}$ eV per particle.  The work described in this paper represents an attempt to improve this situation, and to determine the energy spectra of individual elemental species up to much higher energies.  The current paradigm of the origin of galactic cosmic rays postulates first-order Fermi acceleration in interstellar shock fronts from supernova (SN) explosions as the most potent contributor to the cosmic-ray flux below the knee. This mechanism, first proposed by \\citet{bell78}, must be very efficient indeed, since as much as $\\sim$10\\% of the total kinetic energy released in supernova explosions is required to sustain the cosmic-ray flux. The SN shock acceleration mechanism predicts a source energy spectrum in the form of a power law $\\propto$ $E^{-\\Gamma}$, with a spectral index $\\Gamma \\approx 2.0$ for strong shocks.  Such a spectrum is much harder than the $E^{-2.7}$ spectrum that seems to be typical for most primary cosmic-ray nuclei below the knee.  This difference in spectral slope can be explained if the propagation and containment of cosmic rays in the galaxy are energy-dependent.  Such a behavior had indeed been inferred, even before the SN shock acceleration model was formulated, on the basis of measurements of the ``L/M ratio'', i.e., of the abundance of secondary, spallation-produced cosmic rays (such as the light elements Li, Be, B) relative to their primary parents (such as C and O), \\citep{juliusson72}. One concludes from such measurements that the average amount of interstellar matter encountered by cosmic-ray nuclei during their propagation through the galaxy decreases with energy. The effect is often parameterized by a propagation pathlength $\\Lambda$, which depends on energy as $\\Lambda \\propto E^{-0.6}$ for relativistic nuclei \\citep{engelmann90}. The primary cosmic-ray spectrum at the source would then be approximately $\\propto E^{-2.1}$ or $E^{-2.2}$, close to what the shock acceleration model predicts.  This fact provides strong but indirect evidence for the validity of the shock acceleration model.  However, it must be kept in mind that current measurements of the L/M ratio do not extend beyond 10$^{11}$ eV per amu (or about 10$^{12}$ eV per particle).  Hence, there is an extended region of energies (10$^{12}$ -- 10$^{15}$ eV per particle) where the SN shock acceleration model is assumed to be valid, albeit without the benefit of much observational support.  The SN shock acceleration process is expected to become inefficient at energies around $Z \\times$ 10$^{14}$ eV (where $Z$ is the atomic number of a cosmic-ray nucleus) \\citep{lagage83}, but this limit is not observationally confirmed. One might refer to the steepening of the all-particle spectrum at the \"knee\" as evidence, but even then the origin of cosmic rays with energies far beyond the knee remains a mystery.  The SN-shock acceleration model has recently found strong observational support through the detection of TeV gamma rays from shell-type SN-remnants \\citep{aharonian04}, although there remains a debate whether the parent particles of these gamma rays are electrons or, indeed, protons and nuclei \\citep{berezhko06}.  Clearly, the present observational evidence about high-energy cosmic rays is inadequate to provide answers to some of the most fundamental questions about their origin and galactic propagation, and there exists an undeniable need for improved measurements at higher energies. The TRACER program has been developed to provide direct measurements of the energy spectra of the heavier cosmic-ray nuclei through balloon flights above the atmosphere, up to energies between 10$^{14}$ and 10$^{15}$ eV per particle [The acronym TRACER stands for ``Transition Radiation Array for Cosmic Energetic Radiation'']. Higher energies are currently studied through indirect observations with ground-based air-shower installations. However, attempts to obtain composition details from air-shower observations are affected by a number of systematic uncertainties.  One of the goals of TRACER is to approach the energy region of air shower measurements, with the hope of providing some cross-calibration with the indirect measurements. ",
        "conclusions": "\\label{sec:conclusions}  The energy spectra of the cosmic-ray nuclei determined in this investigation extend to energies around $10^{14}$ eV per particle, or significantly higher for the more abundant nuclei. We believe that this data set represents the most comprehensive information about the spectra of the heavier primary nuclei currently available.  We have studied in some detail how these results provide constraints on our current understanding of the origin of galactic cosmic rays and of models of acceleration and propagation. We realize that the upper energy limit of the measurement is still below the region of the cosmic-ray ``knee'', so perhaps it is not surprising that no significant changes in cosmic-ray composition are yet observed. However, the great similarity between the individual energy spectra is remarkable. The result of our studies of these features is currently being prepared for a separate publication.  It is clear that an extension of the measurements to still higher energy is desirable. The dynamic range of the TRD used here is not exhausted with the current measurement. Hence, in order to reach higher energy, the TRACER detector would either have to be increased in size, which seems impractical as the instrument is already the largest balloon-borne cosmic-ray detector in existence, or would have to be subjected to additional long-duration balloon flights. Ideally, an instrument like TRACER should be exposed in space for several years.  An important objective for measurements beyond the current TRACER results would be an extension of the dynamic range of the detector, such that the light secondary elements (Li, Be, B) as well as carbon and nitrogen can be covered together with the heavier nuclei. If successful to sufficiently high energies ($\\sim 10^3$ GeV per amu), this measurement would lead to a determination of the L/M abundance ratio, i.e., to a determination of the propagation pathlength at high energy.  This would be an essential step towards a better understanding of the galactic propagation of cosmic rays. To make this measurement possible, the electronics of TRACER has recently been completely refurbished, and another successful long-duration balloon flight was performed in 2006. The analysis of these new data is currently in progress and results will be reported in due course."
    },
    "0801/0801.2342_arXiv.txt": {
        "abstract": "% {} {We present integral-field spectroscopy of the ionized gas in the central regions of four galaxies with a low surface brightness disk taken with the Visible Multi Object Spectrograph at the Very Large Telescope and aimed at testing the accuracy in the determination of the central logarithmic slope $\\alpha$ of the mass density radial profile $\\rho(r)\\propto r^\\alpha$ in this class of objects.} {For all the sample galaxies we subtracted from the observed velocity field the best-fit model of gas in circular motions and derived the residuals. Only ESO-LV~5340200 is characterized by a regular velocity field. We extracted the velocity curves of this galaxy along several position angles, in order to estimate the uncertainty in deriving the central gradient of the total mass density from long-slit spectroscopy.} {We report the detection of strong non-ordered motions of the ionized gas in three out of four sample galaxies. The deviations have velocity amplitudes and spatial scales that make not possible to disentangle between cuspy and core density radial profiles.} {} ",
        "introduction": "\\label{sec:introduction}  Low surface brightness (LSB) galaxies are defined as disk galaxies with a central face-on brightness fainter than 22.6 $B-$mag arcsec$^{-2}$. In comparison with high surface brightness (HSB) galaxies, the LSB galaxies have higher mass-to-light ratios and are dominated by the dark matter (DM) even in the central regions \\citep[e.g.,][]{Debl1996,Cour1999,Borr2001}. Therefore, the rotation curves of LSB galaxies are thought to represent an ideal test-bed to check the predictions of $N-$body simulations in a cold dark matter (CDM) universe, where the mass density radial profile of the DM halos is described by a steep power law $\\rho(r)\\sim r^\\alpha$ \\citep[i.e. $\\alpha\\lesssim-1$,][]{Nava1997,Nava2004,Moore1999,Diem2005}.  On the other hand, the observations seem to be inconsistent with the CDM scenario. The DM mass density distribution derived by H~{\\small I}\\ rotation curves of most of LSB galaxies is described by a radial profile with a constant core profile \\citep[i.e. $\\alpha\\simeq0$,][]{mcgaugh1998,salu2001}. This is true also for the ionized-gas rotation curves \\citep[see][]{Mcga2001,Debl2001Let,Mcga2003}. The spatial resolution of these data allows to detect cuspy DM density profiles when they are present. But, other authors do find an agreement of the radio \\citep{vdbosch2000} and optical \\citep{Swat2003} data with the CDM predictions. Despite this observational effort, a unique answer to the central DM density profile slope has not been found.  However, measuring the mass density profile using the gas as potential tracer in galactic centers may result problematic not only because of the resolution effects. Indeed, the gas can have an intrinsic velocity dispersion and does not necessarily move on perfectly ballistic orbits \\citep{Bert1995,Cinz1999}. Moreover, the emissivity distribution of the gas is often clumpy \\citep{Vand2001} and the presence of asymmetric and/or decoupled structures could remain undetected in long-slit observations \\citep[see][and references therein]{Cocc2005}. More mundanely, the imperfect centering and orientation of the slit in optical spectroscopy and the beam smearing in radio observations contribute to the total uncertainties on the measured kinematics.  To address all these issues, it is crucial to obtain high-resolution gas kinematics along different axes. Only recently, the two-dimensional velocity field of the ionized gas in LSB galaxies has been measured by means of integral-field units spectroscopy. \\citet{Simo2005} studied a sample of low mass spiral galaxies with a velocity field which can be explained by either a cored or a cuspy density radial profile. \\citet{Kuzi2006} observed a sample of LSB galaxies. They found that DM halos with a core of constant mass density provide a better fit to the data with respect to those with a density cusp.  In this paper we focus on the intrinsic limitation of the ionized-gas kinematics in tracing the galactic central potential and mass distribution. The work is based on the kinematics in the center of four galaxies with a LSB disk measured by integral-field spectroscopy. The observations and data reduction are described in Sect. \\ref{sec:observations}. The measurement and analysis of the kinematics and distribution of the ionized gas are discussed in Sect. \\ref{sec:analysis}. The results are given in Sect. \\ref{sec:results}. Finally, the conclusions are presented in Sect. \\ref{sec:conclusions}.   \\begin{figure*} \\includegraphics[width=.42\\textwidth]{ms6580_f1a.ps} \\includegraphics[width=.42\\textwidth]{ms6580_f1b.ps} \\includegraphics[width=.42\\textwidth]{ms6580_f1c.ps} \\hfill \\includegraphics[width=.42\\textwidth]{ms6580_f1d.ps}  \\caption{Maps of the distribution and kinematics of the ionized gas for the sample galaxies. East is up and North is right. The ranges are indicated at the right of each panel. {\\em Upper panel:} reconstructed image of the galaxy obtained by integrating the stellar continuum between 5728 and 5758 \\AA . {\\em Middle left panel:} \\ha\\ intensity map. {\\em Middle rigth panel:} \\niig\\ (or \\siip ) intensity map. {\\em Lower left panel:} heliocentric line-of-sight velocity without applying any correction for galaxy inclination. {\\em Lower right panel:} line-of-sight velocity dispersion corrected for instrumental velocity dispersion. The morphological classification, inclination, major-axis position angle, diameter of the 25 $B-$mag arcsec$^{-2}$ isophote, and total blue magnitude are taken from \\citet[][ESO-LV]{ESOLV}. The systemic velocities are from this paper. The distances are derived from the systemic velocities corrected to the CMB reference frame following \\citet{Fixsen1996} and assuming $H_0=75$ \\kmsmpc .} \\label{fig:results} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusions}  The results of the study of the sample galaxies can be summarized into two remarkable aspects, concerning the presence of non-ordered motions of the ionized gas and the uncertainties of recovering the logarithmic slope of the radial profile of the mass density in galaxies with a regular kinematics.  In three objects we found the evidence for both non-ordered motions and kinematically-decoupled components. ESO-LV~3520470 and ESO-LV~4000370 host a small ($\\sim1$ kpc) structure which is rotating around the galaxy minor axis. In ESO-LV~1860550 we found the evidence for a gaseous component infalling toward the galactic center. There is no strong correlation between the distribution and kinematics of the ionized-gas of these galaxies and their starlight distribution. Indeed, there is no hint for the presence of the decoupled components in broad-band images \\citep[see][]{Pizz2005..LSB}.  These kinematic disturbances seriously affect the determination of the radial profile of the mass density. In fact, the typical sizes of these features (0.5--1 kpc) are comparable to the radial range where the differences between the core and cuspy density radial profiles are expected to be observed. Their velocity amplitudes (30--50 \\kms) are also of the same order of both the observed rotation velocities and velocity differences inferred by the different mass density profiles. The majority of previous studies on the central mass distribution of LSB galaxies were based on long-slit spectroscopy, where the presence of kinematic irregularities could have been undetected. Non-circular motions are common phenomenon in inner regions of disk galaxies. Indeed, \\citet{Cocc2004} found that about $50\\%$ of the unbarred bright galaxies show strong off-plane and non-circular gas motions in their centers.  ESO-LV~5340200 was the only sample galaxy characterized by a regular velocity field. It was too distant to recover the mass density radial profile in the central kpc with the desired accuracy. However, it represented a good test-bed to estimate the uncertainties of the measurement of the central velocity gradient. This is crucial in deriving the logarithmic slope $\\alpha$ of the mass density radial profile in objects displaying regular kinematics. We found $\\alpha=1.2\\pm0.3$. The error was obtained by fitting a power law to the density distribution derived from the different rotation curves extracted along several position angles.  Therefore, it represents a good estimate of the scatter obtained by deriving $\\alpha$ using the gas kinematic measured from long-slit spectroscopy. On the other hand, the central velocity gradient of ESO-LV~5340200 is larger than that of typical LSB galaxies. We concluded the error we derived can be considered an upper limit to the typical uncertainty in the estimation of logarithmic slope of the mass density. Although it may be substantially reduced with the analysis of the full two-dimensional velocity field \\citep{Simo2005,Kuzi2006}, the presence of non-circular and non-ordered gas motion remains an unresolved issue. A way to circumvent it is using the stellar kinematics. The measurement of the stellar line-of-sight velocity distribution \\citep{Pizz2005..LSB} allows to build reliable mass models providing a valid alternative to the usual gas-based mass distribution determinations \\citep{Pizz2008, Magorrian2008}."
    },
    "0801/0801.2497_arXiv.txt": {
        "abstract": "All the Trojan asteroids orbit about the Sun at roughly the same heliocentric distance as Jupiter.  Differences in the observed visible reflection spectra range from neutral to red, with no ultra-red objects found so far. Given that the Trojan asteroids are collisionally evolved, a certain degree of variability is expected. Additionally, cosmic radiation and sublimation are important factors in modifying icy surfaces even at those large heliocentric distances.  We search for correlations between physical and dynamical properties, we explore relationships between the following four quantities; the normalised visible reflectivity indexes ($S'$), the absolute magnitudes, the observed albedos and the orbital stability of the Trojans.  We present here visible spectroscopic spectra of 25 Trojans. This new data increase by a factor of about 5 the size of the sample of visible spectra of Jupiter Trojans on unstable orbits. The observations were carried out at the ESO-NTT telescope (3.5m) at La Silla, Chile, the ING-WHT (4.2m) and NOT (2.5m) at Roque de los Muchachos observatory, La Palma, Spain.  We have found a correlation between the size distribution and the orbital stability. The absolute-magnitude distribution of the Trojans in stable orbits is found to be bimodal, while the one of the unstable orbits is unimodal, with a slope similar to that of the small stable Trojans. This supports the hypothesis that the unstable objects are mainly byproducts of physical collisions.  The values of $S'$ of both the stable and the unstable Trojans are uniformly distributed over a wide range, from $0\\ \\%/1000\\AA $ to about $15\\ \\%/1000\\AA$. The values for the stable Trojans tend to be slightly redder than the unstable ones, but no significant statistical difference is found. ",
        "introduction": "The main motivation of this work is to look for correlations between physical and dynamical properties.% We shall discuss ways in which the diversity in surface properties can develop, how relationships with orbital stability may originate and what this may imply.  The Trojan asteroids are believed to have been formed in the outer Solar System (see for example Marzari \\& Scholl 1998 or Morbidelli \\etal 2005) so, ices may be expected to be their main component. However, at the moment, no water or other more volatile substances have been detected on the surface of a Trojan (Yang \\& Jewitt 2007, Emery \\& Brown 2003, Emery \\& Brown 2004). Infrared observations indicate the presence of silicates (Cruikshank \\etal 2001, Emery \\etal 2006). In the case of Ennomos, an unusually hight-albedo object, the surface-content of water ice has been quantified to be below $10$\\% in mass (Yang \\& Jewitt 2007).  Cosmic Radiation modifies the spectroscopic properties of these types of compounds. The precise effect depends on the chemical composition of the asteroidal surface, that of the incoming cosmic radiation and its energy. Laboratory experiments of ion-radiation, of energies in the order of keV, gradually `flattens' the spectral slopes of organic-complexes such as asphaltite and kerite (Moroz et al. 2004), while it produces red and dark residuals upon ice-surfaces such as methanol, methane and benzene, turning the spectra to neutral for very high doses (Brunetto et al. 2006a). Micrometeorite bombardment also lowers the albedo and reddens the surface of silicates rich in olivine, pyroxene and serpentine, as indicated by experiments using UV-laser pulses which simulate this effect (Brunetto et al. 2006b); the corresponding timescale to modify the spectral slope on main belt asteroids are estimated to range between 10$^8$yr to 10$^{10}$yr. On the other hand, experiments of ion-irradiation on a sample of the meteorite Epinal indicate that this timescale for an S-type asteroid could be as short as 10$^4$yr to 10$^6$yr in the NEA population (Strazzulla et al. 2005). The timescale in which the surface spectral properties of asteroids will be modified in the Jupiter region is not simple to estimate as it depends on the modulation of the solar activity at these distances. But, since the density of the cosmic radiation originated in the Sun decreases as the heliocentric-distance squared, on may argue that, if any of these processes is relevant for the Trojans, these timescales should be increased by at least an order of magnitude.  The abundance of irradiated material decreases with depth, hence, if material originally from the interior is exposed in some objects, color differences would emerge between members of the same population. Collisions between objects is an obvious way of achieving this and, according to Dell'Oro \\etal (1998), collisions can be significant in the Trojan population. In the trans-Neptunian population, Gil-Hutton (2002) has shown that links between surface and orbital properties can be established in this way.  Although the Trojan asteroids are located at a large heliocentric distance, ice sublimation could also be an important factor affecting their surfaces. A thin insulator mantle can form in a very short timescale that ranges from $10^4$ yr to $10^6$ yr (Jewitt 2002).  There is a slow dynamical evaporation of bodies from the Trojan clouds due to gravitational perturbations, with the $5:2$ mean motion resonance with Saturn being very important (Levison \\etal 1997, Nesvorny \\& Dones 2002, Tsiganis \\etal 2005). In particular, Thersites (1868), has a high probability of leaving the Trojan swarm in less than $50 Myr$ (Tsiganis \\etal 2000). But the out-flux rate produced by collisions is much larger. Marzari \\etal (1995) have found that there is an out-flux of one Trojan asteroid greater than $1 km$ in size every $1000 yr$. The smaller byproducts of a collision are more numerous and tend to have greater post-encounter velocities and are thus less likely to remain in the most stable regions. Hence, collisions could create variability in surface properties and at the same time produce a correlation between physical properties and dynamical state.  Within the Trojan Clouds there exist a number of dynamical families, which are believed to have been formed collisionally. Fornasier \\etal (2004), Fornasier \\etal (2007)  and Dotto \\etal (2006) found that, for a given family, there is a similarity in the surface properties of the members, except for a few objects that are physically different and these have been called {\\it interlopers}. Unfortunately, the information regarding the surfaces of Trojan asteroids that belong to families is rather sparse to date.     It must also be noticed that interloper objects such as captured comets might exist in the Trojan clouds (Rabe 1972, Yoder 1979). Indeed, objects in Centaur-type orbits are known to evolve into temporary Trojans for time-spans of some 0.5 Myr (Horner \\etal 2005) and `transitional' objects with orbits that can evolve into short period comets are known to exist in the $1:1$ resonance with Jupiter (Karlsonn 2004).   Up to now, it was not possible to investigate correlations between dynamical stability and spectroscopic properties because the colors of the transitional Trojans were largely unknown. In order to rectify this, we have obtained new low-resolution visible spectra of an additional $24$ Trojans in unstable orbits.  This article is organised as follows. In section~\\ref{sec:os} we discuss the methods used to study the dynamical properties of the Trojan asteroids. % We present the data used in this study in section~\\ref{sec:vs}, both the observations carried out for the purpose of this investigation and also the data taken from other authors. Correlations between physical and dynamical properties are discussed in section~\\ref{sec:corr}. % Finally, in section~\\ref{sec:disc}, we discuss our results ",
        "conclusions": "\\label{sec:disc}   Before performing the observation presented here, we found only $4$ Trojan asteroids with determined visual spectra in orbits with $LCE>0.53 \\times 1/(10^5 yr)$ (see tables~\\ref{tab:Spout1}, \\ref{tab:Spout2} and \\ref{tab:Spout3}). Therefore, our data has increased by a factor of five the number of measured visible-spectra of Trojans on unstable orbits. Using this increased sample, we have studied the distribution of visible-spectroscopic observations of the Trojan asteroids and search for correlations with orbital stability, size and albedo.  Some physical characteristics, such as the size distribution, are expected to be related with orbital stability, due to energy equipartition after a physical collision. We find that, while the absolute-magnitude distribution of the Trojans in stable orbits is bimodal, the one corresponding to the unstable orbits is unimodal, with a slope similar to the one for the small stable Trojans. Given the location of the node for the bimodal distribution of the stable objects and the smallest absolute magnitude of the transitional ones, we deduce that fragments of collision become more abundant at $H$ $\\gtrsim$ $9.4$. On the other hand, we find no correlation between size and spectral slope. The resident and the transitional (smaller) objects, do not differ statistically in their surface properties. Therefore, we conclude that the collisional process that feeds the unstable orbits and creates smaller bodies is not altering noticeably the distribution of surface properties."
    },
    "0801/0801.0989_arXiv.txt": {
        "abstract": "We have used the \\rxte\\ and \\inte\\ satellites simultaneously to observe the High Mass X-ray binary \\igr. The spectra obtained in the 3--80 keV range have allowed us to perform a precise spectral analysis of the system along its binary orbit. The spectral evolution confirms the supergiant nature of the companion star and the neutron star nature of the compact object. Using a simple stellar wind model to describe the evolution of the photoelectric absorption, we were able to restrict the orbital inclination angle in the range 37--75 degrees. This analysis leads to a wind mass-loss rate from the companion star of $\\sim$ 10$^{-7}$ M$_{\\sun}$/year, consistent with its expected spectral type. We have detected a soft excess in at least three observations, for the first time for this source. Such soft excesses have been reported in several HMXBs in the past. We discuss the possible origin of this excess, and suggest, based on its spectral properties and occurrences prior to the superior conjunction, that it may be explained as the reprocessing of the X-ray emission originating from the neutron star by the surrounding ionised gas. ",
        "introduction": "High Mass X-ray Binaries (HMXBs) are binary systems consisting of a compact object orbiting a massive companion star. Prior to the launch of the INTErnational Gamma-ray Astrophysics Laboratory (\\inte) in 2002, most of the known HMXBs contained a Be companion. In Be-type HMXBs, the compact object emits strong X-ray flashes when it crosses the equatorial plane of the companion star, where a thick disk of matter originating from the stellar wind is present. O and B-type stars have more isotropic stellar winds which absorb the X-ray emission of the compact object, rendering them almost undetectable below a few keV. Since the launch of \\inte, however, many such systems have been found \\citep[see e.g.][]{IGR_1725}, raising the ratio of Be-HMXBs to O,B-HMXBs close to 0.5. The study of this new population may have strong implications on the scenarios of binary formation and evolution. In this perspective, the use of X-ray spectroscopy at different orbital phases makes it possible to probe the stellar wind, providing two-dimensional information on the density and ionization structure of the wind. For instance, the soft excess that is present in the soft X-ray spectra of many HMXBs, whose origin is still quite mysterious, is linked to the physics of the wind close to the compact object, especially the region where the fast moving stellar wind collides with the slow moving and highly ionized gas surrounding the compact object \\citep{soft}.   \\igr\\ was discovered by \\inte\\ by the satellite \\inte\\ in March 2003 \\citep{decouverte}. Prior to our study, it was identified an a HMXB with an orbital period of 13.552 days \\citep{periode}, composed probably of a neutron star \\citep{Jerome} orbiting a B0.5 I type supergiant companion star \\citep{Hanni}. Thus, it is a good candidate to study the properties of the obscured X-ray binaries. ",
        "conclusions": "The study of X-ray binaries is challenging since it is often difficult if not impossible to identify their visible and infrared counterparts. Even if an infrared counterpart were observed, the distance to some systems prohibits the measurement of the orbital characteristics. Our study shows that X-ray observations can overcome these limitations and produce very precise inferences. The \\rxte\\ and \\inte\\ observations of \\igr\\ have led to good measurements of the orbital period of the system and constraints on its inclination angle.  Moreover, we can use the compact object to probe the stellar wind of the companion. In the case of \\igr, we diagnosed the type of the companion (supergiant O or B), the wind density and its structure around the neutron star. More precise observations could lead to constraints on the mass and radius of the companion, and better constraints on the stellar wind. This allows the study of new \\inte\\ sources, either distant or highly absorbed, and ultimately the determination of new useful data for X-ray binary evolution scenarios."
    },
    "0801/0801.0492_arXiv.txt": {
        "abstract": " ",
        "introduction": "During the last two decades, the focus of star formation research has shifted from understanding the collapse of a single dense core into a star to studying the formation hundreds to thousands of stars in molecular clouds.  In this chapter, we overview recent observational and theoretical progress toward understanding star formation on the scale of molecular clouds and complexes, i.e the macrophysics of star formation \\citep{2007ARA&A..45..565M}.  We begin with an overview of recent surveys of young stellar objects (YSOs) in molecular clouds and embedded clusters, and we outline an emerging picture of cluster formation.  We then discuss the role of turbulence to both support clouds and create dense, gravitationally unstable structures, with an emphasis on the role of magnetic fields (in the case of distributed stars) and feedback (in the case of clusters) to slow turbulent decay and mediate the rate and density of star formation.  The discussion is followed by an overview of how gravity and turbulence may produce observed scaling laws for the properties of molecular clouds, stars and star clusters, and how the observed, low star formation rate may result from self regulated star formation. We end with some concluding remarks, including a number of questions to be addressed by future observations and simulations. ",
        "conclusions": ""
    },
    "0801/0801.0201_arXiv.txt": {
        "abstract": "We find that deviation from the virial equilibrium of the Abell Cluster A586 yields evidence \t\t\tof the interaction\tbetween dark matter and dark energy. We argue that this interaction might imply a violation of the Equivalence Principle. These evidences are found in the context of two different models of dark\tenergy-dark matter interaction. ",
        "introduction": "Current cosmological data strongly suggests that the existence dark energy (DE) and dark matter (DM) is crucial for a suitable description of universe's evolution.  Even though observations are consistent with the $\\Lambda$CDM model, a deeper insight into the nature of DE and DM might require more complex models -- in particular considering the interaction between these components. However, so far, no evidence of this putative interaction has been presented. In this work, we argue that data on the cluster A586 shows evidence that DE-DM interaction is an active dynamical factor. Furthermore, we argue that this suggests evidence of violation of the Equivalence Principle (EP).   In what follows, we set the general framework to treat the interaction between DE and DM and consider two very different, phenomenologically viable models: one based on \\textit{ad hoc} DE-DM interaction [\\cite{Amendola}], the other on the generalized Chaplygin gas (GCG) model with explicit identification of DE and DM [\\cite{Bento04}]. Our observational inferences are based on cluster A586, given its stationarity, spherical symmetry and wealth of available observations [\\cite{Cypriano:2005}]. We also compare our results with other cosmological observations [\\cite{Guo07}]. ",
        "conclusions": "Observations of cluster A586 [\\cite{Cypriano:2005}] suggest evidence of departure from virialization given that A586 is very spherical and relaxed (from its mass distribution and Gyrs without mergers). The generalized Layzer-Irvine equation allows to interpret this departure as due to interaction with DE. We can therefore link the observed virialization to interaction [\\cite{Berto07}a,b] for two different interacting DE-DM  models: an interacting quintessence with constant $\\omega_{DE}$ {\\tiny }[\\cite{Amendola}] and a Chaplygin gas with $\\omega_{DE}=-1$ {\\tiny }[\\cite{Bento04}].  Based on the evidence of interaction, we argue that the Equivalence Principle should be violated as seen through the bias parameter [\\cite{Berto07}a] and on Baryon/DM asymmetric collapse [\\cite{OBPedro07b}]. Our results are consistent with CMB, supernovae and Baryon acoustic oscillations as well as the simulation of the fall of the Sagittarius dwarf galaxy into our own. These results are quite encouraging and suggest that our method should be further employed into other cluster systems (see e.g. [\\cite{Abdalla:2007rd}])."
    },
    "0801/0801.2518_arXiv.txt": {
        "abstract": "% We present GOSSIP (Galaxy Observed-Simulated SED Interactive Program), a new tool developed to perform SED fitting in a simple, user friendly and efficient way. GOSSIP automatically builds-up the observed SED of an object (or a large sample of objects) combining magnitudes in different bands and eventually a spectrum; then it performs a $\\chi^2$ minimization fitting procedure versus a set of synthetic models. The fitting results are used to estimate a number of physical parameters like the Star Formation History, absolute magnitudes, stellar mass and their Probability Distribution Functions. User defined models can be used, but GOSSIP is also able to load models produced by the most commonly used synthesis population codes. GOSSIP can be used interactively with other visualization tools using the PLASTIC protocol for communications. Moreover, since it has been developed with large data sets applications in mind, it will be extended to operate within the Virtual Observatory framework. GOSSIP is distributed to the astronomical community from the PANDORA group web site (\\makeURL{http://cosmos.iasf-milano.inaf.it/pandora/gossip.html}). ",
        "introduction": "GOSSIP is a tool created to fit the electro-magnetic emission of an object (the SED, Spectral Energy Distribution) against synthetic models, to find the simulated one that best reproduces the observed data. It has been developed to perform this task in a simple, user friendly and efficient way.\\\\ GOSSIP was born within the frameworks of the VVDS (Le F\\'evre et al. 2005) and the zCOSMOS (Lilly et al. 2007) surveys; therefore it has been optimized to work on huge amounts of data as the ones provided by modern photometric and spectroscopic surveys.\\\\ GOSSIP has been written by the \\htmladdnormallinkfoot{PANDORA Group}{http://cosmos.iasf-milano.inaf.it/pandora} at INAF IASF-Milano, using the \\htmladdnormallinkfoot{Python}{http://www.python.org} language for the graphical part and the C language for the high performance computational tasks. ",
        "conclusions": "We have presented GOSSIP, a new tool dedicated to SED fitting. It has been developed in order to perform this task in a simple, user friendly and efficient way. GOSSIP builds-up observed SEDs, loads synthetic models, performs $\\chi^2$ minimization fitting between them and estimates physical parameters for huge samples of objects. The features implemented within this program make it a very useful tool for the analysis of the large data samples currently available to the astronomical community."
    },
    "0801/0801.1031_arXiv.txt": {
        "abstract": "We present a phase-coherent timing analysis of the intermittent accreting millisecond pulsar \\saxjfull. A new timing solution for the pulsar spin period and the Keplerian binary orbital parameters was achieved by phase connecting all episodes of intermittent pulsations visible during the 2001 outburst. We investigate the pulse profile shapes, their energy dependence and the possible influence of Type I X-ray bursts on the time of arrival and fractional amplitude of the pulsations. We find that the timing solution of \\saxjfull\\ shows an erratic behavior when selecting different subsets of data, that is related to substantial timing noise in the timing post-fit residuals. The pulse profiles are very sinusoidal and their fractional amplitude increases linearly with energy and no second harmonic is detected. The reason why this pulsar is intermittent is still unknown but we can rule out a one-to-one correspondence between Type I X-ray bursts and the appearance of the pulsations. ",
        "introduction": "Recently it has been shown (\\citealt{Kaaret}, \\citealt{Galloway07}, \\citealt{Casella}, \\citealt{Gavriil}, \\citealt{Altamirano}) that intermittent pulsations can be observed during the transient outbursts of some accreting millisecond pulsars (AMPs). It was emphasized (\\citealt{Casella}, \\citealt{Gavriil}, \\citealt{Altamirano}) that a new class of pulsators is emerging with a variety of phenomenology i.e., pulsations lasting only for the first two months of the outburst (HETE J1900.1-2455, \\citealt{Kaaret}, \\citealt{Galloway07}), switching on and off throughout the observations (SAX J1748.9-2021, \\citealt{Altamirano}, and HETE J1900.1-2455 \\citealt{Galloway07}), or even appearing for only ~150 s over 1.3 Ms of observations (Aql X-1, \\citealt{Casella}).  This is different from the behavior of other known AMPs, where the pulsations persist throughout the outburst until they drop below the sensitivity level of the instrumentation.   \\citet{Galloway07} and \\citet{Altamirano} pointed out the clear, albeit non-trivial, connection between the appearance of pulsations in HETE J1900.1-2455 and SAX J1748.9-2021 and the occurrence of thermonuclear bursts (Type I X-ray bursts). \\citet{Galloway07} demonstrated that increases of pulse amplitude in HETE J1900.0$-$2455 sometimes but not always follow the occurrence of Type I bursts nor do all bursts trigger an increase in pulse amplitude. \\citet{Altamirano} showed that in \\saxjfull\\ the pulsations can re-appear after an absence of pulsations of a few orbital periods or even a few hundred seconds. All the pulsing episodes occur close in time to the detection of a Type I X-ray burst, but again not all the Type I bursts are followed by pulsations so it is unclear if the pulsations are indeed triggered by Type I X-ray bursts.\\\\ The pulse energy dependence and the spectral state in which pulsations are detected also differ between these objects. As shown by \\citet{Cui}, \\citet{Galloway05}, \\citet{Gierlinski} and \\citet{Galloway07}, the fractional amplitude of the pulsations in AMPs monotonically decreases with increasing energies between 2 and $\\approx 20$keV. However in IGR J00291+5934, \\citet{Falanga} found a slight decrease in the pulse amplitude between 5 and 8 keV followed by an increase up to energies of 100 keV. For the single pulse episode of Aql X-1 the pulsation amplitude increased with energy between 2 and $\\approx 20$keV (\\citealt{Casella}). Interestingly in Aql X-1 the pulsations appeared during the soft state, contrary to all the other known AMPs which pulsate solely in their hard state (although many of these objects have not been observed in a soft state). Therefore it is interesting to investigate these issues for \\saxjfull\\ as well.\\\\ In order to study these various aspects of \\saxjfull\\ an improved timing solution is necessary. The spin frequency of the neutron star and the orbital parameters of \\saxjfull\\ were measured by \\citet{Altamirano} using a frequency-domain technique that measures the Doppler shift of the pulsar spin period due to the orbital motion of the neutron star around the center of mass of the binary. The source was found to be spinning at a frequency of 442.361 Hz in a binary with an orbit of ~8.7 hrs and a donor companion mass of $\\sim 0.1-1M_{\\odot}$. However, the timing solution obtained with this technique was not sufficiently precise to phase connect between epochs of visible pulsations. The intermittency of the pulsations creates many gaps between measurable pulse arrival times, complicating the use of standard techniques to phase connect the times of arrival (TOAs). The current work presents a follow-up study to the discovery paper of intermittent pulsations in \\saxjfull\\ (\\citealt{Altamirano}) and provides for the first time a timing solution obtained by phase connecting the pulsations observed during the 2001 outburst. In \\S 2 we present the observations and in \\S 3 we explain the technique employed to phase connect the pulsations. In \\S 4 we show our results and in \\S 5 we discuss them. In \\S 6 we outline our conclusions. ",
        "conclusions": "We have shown that it is possible to phase connect the intermittent pulsations seen in \\saxjfull\\ and we have found a coherent timing solution for the spin period of the neutron star and for the Keplerian orbital parameters of the binary.\\\\ We found strong correlated timing noise in the post-fit residuals and we discovered that this noise is strongest in low fractional amplitude pulses and is not related with the orbital phase. Higher-amplitude ($\\simgreat 1.8\\%$ rms) pulsations arrive systematically later than lower-amplitude ($\\simless 1.8\\%$ rms) ones, by on average 145$\\mu$s. The pulsations of \\saxjfull\\ are sinusoidal in the 5-17 keV band, with a fractional amplitude linearly increasing in the energy range considered. The pulsations appear when the source is in the soft state, similarly to what has been previously found in the intermittent pulsar Aql X-1. The origin of the intermittency is still unknown, but we can rule out a one-to-one correspondence between Type I X-ray bursts and the appearance of pulsations."
    },
    "0801/0801.3805_arXiv.txt": {
        "abstract": "This is a study of the scattering and absorption of planar gravitational waves by a Kerr black hole in vacuum. We apply the partial wave method to compute cross sections for the special case of radiation incident along the rotation axis. A catalogue of numerically-accurate cross sections is presented, for a range of incident wavelengths $M\\omega \\le 4$ and rotation rates $a \\le 0.999M$. Three effects are studied in detail: polarization, helicity-reversal and glory scattering. First, a new approximation to the polarization in the long-wavelength limit is derived. We show that black hole rotation distinguishes between co- and counter-rotating wave helicities, leading to a term in the cross section proportional to $a\\omega$. Second, we confirm that helicity is not conserved by the scattering process, and show that superradiance amplifies the effect. For certain wavelengths, the back-scattered flux is enhanced by as much as $\\sim 35$ times for a rapidly-rotating hole (e.g. for $a = 0.999M$ at $M\\omega = 0.945$). Third, we observe regular glory and spiral scattering peaks in the numerically-determined cross sections. We show that the angular width and intensity of the peaks may be estimated via a semi-classical approximation. We conclude with a discussion of the observable implications of our results. ",
        "introduction": "Introduction} Gravitational waves are propagating ripples in spacetime that are predicted by General Relativity (GR). There is strong indirect evidence for their existence, for example, from thirty-five years of pulsar timing measurements \\cite{Lorimer-1998}. Due to the tiny amplitude of waves reaching Earth (with a dimensionless strain $h \\sim 10^{-21}$), gravitational waves are yet to be detected directly. Many experimentalists are optimistic that first-light observations will occur within the decade. Ground-based detectors such as LIGO \\cite{Frey-2007} are reaching sufficient sensitivities such that even non-detection will constrain astrophysical theories (for example, on the origin of gamma ray bursts \\cite{LIGO-2007}). Plans for a space-based detector (LISA) are well advanced \\cite{Danzmann-2003}.  Gravitational waves are of interest to astronomers because they are generated by some of the most energetic astrophysical processes, such as binary mergers, supernovae and galaxy collisions. Electromagnetic radiation carries relatively little information about the most energetic regions of such processes, because photons are strongly scattered, absorbed and thermalized by intervening matter. On the other hand, gravitational waves are only weakly coupled to matter, and carry information about the dynamics at the heart of such processes.  Black holes are likely to be powerful sources of gravitational waves. For example, at the end of the inspiral phase, the merger of a pair of black holes creates a distinctive ``chirp'' signal. Rapid progress has been made in numerically modelling the merger process \\cite{Pretorius-2005}, providing template signals for experimentalists seeking to separate a weak gravitational-wave signal from a noisy background.  In addition to emitting radiation, black holes will scatter and absorb radiation that impinges upon them. To some extent, this is a secondary effect. However, scattering processes have long been of foundational interest to physicists, and this process is no exception. In this paper, we seek to answer the following question: how are gravitational plane waves scattered and absorbed by rotating (Kerr) black holes?  In our scenario, we assume a monochromatic wave produced by some far-distant source impinges on an isolated black hole. We assume the incident wave is essentially planar, long-lasting, and weak enough to be treated through the linearized approximation. Here, we consider only the special case of on-axis incidence. Our principle aim is to determine the absorption cross section $\\sigma_a$ and the differential scattering cross section $d \\sigma / d \\Omega$ as a function of scattering angle $\\theta$. Inevitably, the long-range $1/r$ ``Newtonian'' component of the potential leads to a divergent scattering cross section in the direction of incidence. In this respect, a black hole is no different to any generic massive source. More interestingly though, the signal scattered through large angles carries information about the near-horizon structure of the hole. In this paper we compute this signal and examine in detail the effect of black hole rotation.  Many studies of black hole scattering have been made over the last forty years. The most relevant for our purposes are studies of time-independent scattering, where it is assumed that the incident wave is monochromatic and long-lasting. Pioneering contributions include the works of Hildreth \\cite{Hildreth-1964}, Matzner and co-workers \\cite{Matzner-1968, Chrzanowski-1976, Matzner-1977, Matzner-1978, Handler-1980}, Sanchez \\cite{Sanchez-1976, Sanchez-1977, Sanchez-1978a, Sanchez-1978b} and Mashhoon \\cite{Mashhoon-1973,Mashhoon-1974, Mashhoon-1975}. Early studies focussed on the simplest setup: scalar (Klein-Gordon) waves incident on a non-rotating hole. Later works extend the analysis to higher spin: fermionic \\cite{Dolan-2006}, electromagnetic \\cite{Fabbri-1975} and gravitational \\cite{Futterman-1981} fields. A number of authors have considered the polarizing effect of a rotating hole on waves with spin \\cite{Mashhoon-1973, Mashhoon-1974, Mashhoon-1975, DeLogi-1977, Guadagnini-2002, Barbieri-2004, Barbieri-2005, Guadagnini-2008}. The physical origin of regular oscillations in cross-section-versus-scattering-angle plots was clarified by DeWitt-Morette \\cite{DeWitt-Morette-1984}, Zhang \\cite{Zhang-1984} and others \\cite{Matzner-1985, Anninos-1992}. Such oscillations are a diffraction effect due to interference of null rays passing close to the unstable photon orbit. In other words, the large-angle scattering cross section contains information about the near-horizon structure of the hole.  The efforts of Matzner and co-workers \\cite{Matzner-1968, Matzner-1977, Matzner-1978, Handler-1980, Futterman-1981} culminated in the publication of a monograph \\cite{Futterman-1988}, which set out in detail the partial-wave approach to computing black hole scattering cross sections. In this work, a range of plots of numerically-determined cross sections were presented. We believe that the plots for the rotating black hole are not accurate, due to numerical difficulties arising from the long-range nature of the gravitational field. Our belief is supported by more recent work by Glampedakis and Andersson \\cite{Andersson-1995} and \\cite{Glampedakis-2001}, who examined the scattering of scalar waves by a Kerr hole. A primary aim of this paper is to combine the numerical accuracy of Glampedakis and Andersson with the partial wave methods of Matzner \\emph{et al.} for the gravitational wave case.  A further objective is to examine in detail the coupling between the rotation of the hole and the helicity of the incident wave. The black hole is expected to have a polarizing effect \\cite{Mashhoon-1973}. A key result of this paper is that the polarizing effect is present even in the long-wavelength limit. A further observation is that parity-dependence of the scattering potential gives rise to a non-zero amplitude for helicity reversal. The helicity-reversed flux is maximal in the antipodal direction. We find that, close to the extremal rotation limit, the flux scattered in the backward direction can be significantly enhanced by the effects of superradiance.  This paper is structured as follows. In Section \\ref{sec-background} we outline basic results for black hole wave scattering and quote the key analytic results of this work. In Section \\ref{sec-partial-wave-method} we introduce the partial-wave method of Matzner \\emph{et al.} (\\ref{subsec-partialwave}), Teukolsky's equations (\\ref{subsec-teuk}), and the spin-weighted spheroidal harmonics (\\ref{subsec-spheroidal}). In Section \\ref{sec-long-wavelength} we apply the partial wave method to derive a new approximation for the scattering cross section and polarization in the long wavelength regime. In the process, we derive asymptotic results for the phase shifts (\\ref{subsec-mst}), the spheroidal harmonics (\\ref{subsec-spheroidal-expansion}) and the scattering amplitudes (\\ref{subsec-amplitudes}). Section \\ref{sec-nummeth} describes methods for the numerical calculation of cross sections. The numerical task splits into three parts. First (\\ref{subsec-num-spheroidal}) we compute the spin-weighted spheroidal harmonics using a spectral decomposition method. Second (\\ref{subsec-phaseshifts}) we calculate phase shifts using a Sasaki-Nakamura transformation. Third (\\ref{subsec-reduction}) we introduce a series reduction method for improving the convergence properties of partial wave series. In Section \\ref{sec-results} we present numerical results for the absorption and scattering cross sections and the polarization. We conclude in Section \\ref{sec-discussion} with a discussion of the physical implications of our results. Note that units $G = c = 1$ are employed throughout this paper. ",
        "conclusions": "} In this study we have shown that it is possible to compute accurate gravitational-wave cross sections using a partial-wave series approach. Although the key formulae (\\ref{f-def}--\\ref{T-trans}) were derived three decades ago by Matzner \\emph{et al.} \\cite{Matzner-1978, Handler-1980, Futterman-1988}, we believe this is the first time that accurate results for the rotating case have been presented.  Early studies \\cite{Matzner-1978, Handler-1980, Futterman-1988} encountered two obstacles to progress. Firstly, the Teukolsky equation has a long-ranged potential for perturbations with spin $|s| > 1/2$. This leads to so-called `peeling' behaviour in the asymptotic solutions which renders the determination of phase shifts numerically difficult. Secondly, the partial wave series (\\ref{f-def}) has an infinite number of terms which increase in magnitude with $l$, and it is divergent at $\\theta = 0$.  The first obstacle is overcome by transforming the radial equation. Matzner and Ryan \\cite{Matzner-1978} used Press and Teukolsky's transformation \\cite{Press-1973}; we used a Sasaki-Nakamura transformation \\cite{Sasaki-1982}. Finding the phase shifts is then relatively straightforward -- although of course we benefit from the rapid advance in computing power over the last thirty years.  It is overcoming the second obstacle -- approximating an infinite, divergent series -- that is the key to numerical progress. Regardless of computing power, the series must be truncated at some $l = l_{\\text{max}}$. Yet, the magnitude of the coefficients in the series (\\ref{f-def}) increases with $l$. Matzner \\emph{et al.} tried decomposing the phase shifts into a long-range `Newtonian' contribution (that varies logarithmically with $l$) and a short-range remainder (that decays with $l$). This approach works well for scalar waves \\cite{Glampedakis-2001}, but it does not seem to work for higher spins \\cite{Dolan-2006}. In truncating the series, Matzner \\emph{et al.} introduce numerical error with an angular frequency proportional to $1 / l_\\text{max}$. See for example Fig. 14 in \\cite{Handler-1980}. In this work, we circumvented the obstacle by employing a series reduction method to improve the convergence properties of the series \\cite{Yennie-1954}.  The results presented here shed new light on three intriguing scattering phenomena: polarization, helicity violation, and glory/spiral diffraction. Below we briefly recap the important points.  The coupling between the rotation of the black hole and the helicity of the incident wave means that a rotating black hole has a polarizing effect. In other words, since waves of co- and counter-rotating helicities are distinguished by a rotating black hole, a partial polarization is induced in an initially unpolarized beam. The polarization is a function of scattering angle. In the long-wavelength limit, the partial polarization is given by (\\ref{pol-final}). At shorter wavelengths the polarization it is a more complicated function of scattering angle (see e.g. Fig. \\ref{fig-csec-kerr-highE}), due to diffractive effects.  The parity-dependence of the phase shifts leads to a scattered component $M^{-2} |g(\\theta)|^2$ (Eq. \\ref{g-def}) which has the opposite helicity to the incident flux. The helicity-reversed flux is largest in the antipodal direction ($\\theta = 180^\\circ \\equiv \\pi$). Conversely, the helicity-conserved flux $|f(\\theta)|^2$ is zero in the backwards direction.  In the long-wavelength limit we showed that $M^{-2} |g(\\pi)|^2 = 1 + 4 a \\omega + \\mathcal{O}(\\omega^2)$.  In the non-rotating case, we find $|g(\\pi)|^2$ has a maximum of just $M^{-2} |g(\\pi)|^2 \\sim 1.2$ at $M \\omega = 0.3$ (Fig. \\ref{fig-csec-schw-approx}), and it is suppressed at shorter wavelengths (Fig. \\ref{fig-csec-schw2}). For a co-rotating incident helicity, the back-scattered flux is enhanced by superradiance in the $l = 2$, $m = 2$ mode, which occurs for frequencies $\\omega < \\omega_c = a / M r_+$. We find the backscattered flux may be enhanced by an order of magnitude or more as the extremal rotation limit is approached ($a \\rightarrow 1$). In particular, for $\\as = 0.999$ we observe a maximum enhancement of $M^{-2} |g(\\pi)|^2 \\sim 35$ at $M\\omega \\approx 0.945$. This contrasts with the scalar-wave ($s=0$) case, for which superradiance has a negligible effect \\cite{Glampedakis-2001}.  At short wavelengths, $\\lambda \\lesssim r_S$ ($M\\omega \\gtrsim 1$), we observe regular peaks and troughs in the scattering cross sections (Figs. \\ref{fig-csec-schw2}, \\ref{fig-csec-schw-spins} and \\ref{fig-csec-kerr-highE}). This is a diffraction effect, due to the interference of paths scattered through angles $\\theta, 2\\pi - \\theta, 2\\pi + \\theta$ \\cite{Anninos-1992}. The angular width of the oscillations is (approximately) proportional to the ratio of the inner unstable orbit radius to the incident wavelength. Hence, the angular width is roughly proportional to $M\\omega$, but also depends on rotation $\\as$ and incident helicity. Again, since co- and counter-rotating helicities are distinguished by a rotating black hole, the scattered wave acquires a partial polarization which depends on scattering angle.  Unlike light, gravitational radiation is able to `see through' the complicated and messy astrophysical region surrounding a black hole (e.g. accretion disk, jets, synchrotron radiation, etc) to probe the near-horizon geometry. Is there any prospect of detecting the gravitational radiation that is scattered by black holes? Since gravitational radiation has yet to be detected \\emph{at all}, the short answer is: not yet! The best hope in the short term is that experimentalists detect the strongest signals generated by highly dynamic, strong-field interactions, such as the final merger of black holes. Dynamic (i.e. time-dependent) perturbations will excite black hole `quasinormal modes' (QNMs) \\cite{Vishveshwara-1970b}. The frequencies and decay times of QNMs would provide a clear signature of the underlying black hole parameters $M$ and $\\as$.  Diffraction patterns like those shown in Figs. \\ref{fig-csec-schw2}, \\ref{fig-csec-schw-spins} and \\ref{fig-csec-kerr-highE} may provide an alternative signature for future detectors. Let us imagine a supermassive black hole (such as Sag.~A*) illuminated by a strongly-collimated and long-lasting beam of gravitational radiation. Let us assume the incident radiation has a short wavelength $\\lambda \\ll r_s$ and a wide-band spectrum $\\Delta \\omega_0 \\gg 1 / M$. Our angle of observation $\\theta$ is effectively fixed, but we are free to vary the observation frequency. Changing the observation frequency by $\\delta \\omega \\sim 1 / M$ would bring us from an interference maximum to a minimum \\cite{Futterman-1988}, and, if the BH is rotating, the polarization of the signal would be modified. Hence the characteristic signature of BH scattering would be a signal whose intensity and polarization oscillates rapidly over a frequency $\\delta \\omega$. %  Two extensions of this work suggest themselves. Firstly, one could treat waves approaching from an arbitary angle of incidence relative to the BH rotation axis. The relevant partial wave formulae, given in \\cite{Futterman-1988}, involve an additional sum over the azimuthal numbers $m$. We would expect to find maximal polarization for waves impinging along the axis of rotation, and zero polarization for waves approaching in the equatorial plane. Secondly, it would be relatively straightforward to apply the numerical method detailed here to the scattering of electromagnetic waves from a rotating hole, since the partial wave formulae are supplied in \\cite{Futterman-1988}, and an extension of the Sasaki-Nakamura to $|s| = 1$ is available in \\cite{Hughes-2000b}. In addition, it should also be straightforward to extend the long-wavelength approximation of section \\ref{sec-long-wavelength} to the electromagnetic case. We hope this would shed light on the physical origin of the factor of two difference between (\\ref{pol-final}) and the result (\\ref{pol-guadagnini}) of Barbieri and Guadagnini \\cite{Barbieri-2005} for scattering from classical rotating matter. We hope to undertake this calculation in the near future.  \\ack Thanks to Bahram Mashhoon for email correspondence, Lu\\'is Crispino and Ednilton Oliveira for interesting discussions, and Marc Casals for proof-reading the manuscript. Thanks also to Calvin Smith, Paul Watts and Adrian Ottewill for helpful comments. Financial support from the Irish Research Council for Science, Engineering and Technology (IRCSET) is gratefully acknowledged.  \\appendix"
    },
    "0801/0801.4108_arXiv.txt": {
        "abstract": "If the spontaneous breaking of Peccei-Quinn symmetry comes from soft supersymmetry breaking, the fermionic partners of the symmetry-breaking fields have mass of order the gravitino mass, and are called flatinos. The lightest flatino, called here the flaxino, is a CDM candidate if it is the lightest supersymmetric particle. We here explore flaxino dark matter assuming that the lightest ordinary supersymmetric particle is the stau, with gravity-mediated supersymmetry breaking. The decay of the stau to the flaxino is  fast enough not to spoil the standard predictions of Big Bang Nucleosynthesis, and its track and decay can be seen in future colliders. ",
        "introduction": "By observing the temperature anisotropy of Cosmic Microwave Background (CMB), the WMAP Collaboration reported that the density parameter of dark matter (DM)  or cold dark matter (CDM)  $\\Omega_{\\rm DM, obs}$ is \\begin{eqnarray} \\label{eq:OmegaWMAP3} \\Omega_{\\rm DM, obs}h^{2} = 0.104^{+0.007}_{-0.010}, \\end{eqnarray} at 68 \\% C.L.~\\cite{Spergel:2006hy} where $h$ is the normalized Hubble constant.  One of the most popular candidates for DM is the  lightest supersymmetric particle (LSP) in supersymmetric theories with R-parity.  Recent studies showed that Big Bang Nucleosynthesis (BBN) puts interesting limitations on the candidates of LSP DM.  A long-lived next LSP (NLSP) whose lifetime is longer than $0.1$ sec is dangerous for the successful BBN because the decaying NLSP produces a lot of (high-energy) daughter particles such as photons and hadrons during/after BBN, which can destroy $^{4}$He (and D) and non-thermally produce another light elements copiously~\\cite{radBBN,hadBBNold,hadBBNrec}.  In particular, a non-thermal hadron emission due to the decaying NLSP is severely constrained by the observational light-element abundances of D, $^{4}$He, $^6$Li, and $^7$Li~\\cite{hadBBNold,hadBBNrec}.  In the popular case of the neutralino (gaugino or higgsino) LSP, as is well-known, the unstable gravitino decay puts stringent upper bound on the reheat temperature after primordial inflation, which could be problematic for various cosmological considerations, e.g., of the inflation models. This problem can be avoided if the gravitino is the LSP. (See also~\\cite{Moroi:1993mb} for constraints on the reheat temperature in the gravitino LSP scenario).  However, the decays of ordinary supersymmetric particles to the gravitino turn out to be dangerous as well.  The case of neutralino NLSP, because the branching ratio into hadrons is close to the order of ${\\cal O}(1)$, is almost excluded~\\cite{Feng:2004mt,Steffen:2006hw} by considering BBN constraints~\\cite{hadBBNold,hadBBNrec}. Slepton NLSP scenarios appear much more favorable than neutralino NLSP scenarios in  cosmology.  Although the sneutrino NLSP scenario should be possible because of the smaller hadronic and radiative decaying modes, the primordial abundance of the sneutrino is  too small to satisfy the dark matter density for its daughter LSP particle~\\cite{Feng:2004zu,Kanzaki:2006hm}. Therefore, we would need another production mechanism of the LSP in cosmology, which usually requires a fine tuning.  Recently, the stau NLSP scenario has been discussed a lot as there are specifically severer/more interesting problems in BBN related to  the long-lived negatively charged particle~\\cite{Pospelov:2006sc,Kohri:2006cn,Kaplinghat:2006qr,Cyburt:2006uv}. See also long-term problems discussed in literature~\\cite{CargedDMold} and  topics in doubly-charged particles~\\cite{doubleC}.  If the lifetime of the negatively charged particle $C^{-}$ such as $\\tilde{\\tau}^{-}$ is longer than 10$^{2}$ sec, bound states with ambient light elements can be produced. In particular,   the bound state with $^{4}$He denoted by ($^{4}$He,$C^{-}$), which is produced after 10$^{3}$ sec, significantly enhances the reaction rate D+($^{4}$He,$C^{-}$) $\\to$ $^{6}$Li + $C^{-}$ + $\\gamma$~\\cite{Pospelov:2006sc} by a factor of $\\sim {\\cal O}(10^{7})$~\\cite{Hamaguchi:2007mp}. Then $^{6}$Li is over-produced by the catalyzed BBN, and this scenario is severely constrained by observations~\\cite{Kawasaki:2007xb}.  \\smallskip  Various attempts \\cite{Bird:2007ge,CBBNsolution,Buchmuller:2007ui,CBBNrecent} have been made to tackle this problem and realize a successful dark matter production mechanism.  In this paper, we suggest  the ``flaxino'' DM as a viable option to avoid the above-mentioned difficulties for the neutralino or slepton NLSP.  The flaxino is an axino-like fermion appearing in flat-direction axion models where the spontaneous breaking of Peccei-Quinn symmetry comes from soft supersymmetry breaking and thus all of the fermionic partners of the symmetry-breaking fields have mass of order the gravitino mass. The flaxino has a coupling to the ordinary superparticles suppressed by the axion scale $F_a$ and thus attractively leads to a fast decay of NLSP,  which should be compared with the gravitino coupling suppressed by the Planck scale. ",
        "conclusions": "The lightest supersymmetric particle (LSP) is a dark matter candidate. It is usually said that the LSP will be either the lightest superpartner of a Standard Model particle (so called LOSP), the gravitino or the axino.  For cosmological constraints and collider signals of the LSP, it matters what is the NLSP.  In this paper we revisited the axino-like LSP arising in the flat DFSZ axion models for the case of the stau NLSP.  Following \\cite{Chun:2000jx}, we assumed that the spontaneous breaking of Peccei-Quinn symmetry comes only from supersymmetry breaking. Then all of the PQ particles except the axion will have masses of order the gravitino mass. The scalar particles are called flatons and their superpartners are called flatinos.  In this scenario the saxion is a linear combination of flatons and the axino a linear combination of flatinos, but they are not mass eigenstates and have no special status. So the LSP candidate from the PQ sector is not the axino, but the lightest flatino which we have dubbed the flaxino. We have explored the simplest version of this scenario, which  has two flatinos and assumes  gravity mediated supersymmetry breaking.  We  have found that the stau NLSP decay to the flaxino LSP is fast enough to maintain the standard predictions of Big Bang Nucleosynthesis.  The big difference from the gravitino dark matter scenario is that the flaxino/axino coupling ($\\sim 1/F_a$) is much larger than the gravitino coupling ($\\sim 1/M_P$), which makes the stau NLSP decay much faster.  In the DFSZ axion model,  the stau-flaxino/axino coupling appears at tree level as a consequence of the $\\mu$ term mixing between the higgsino and flaxino/axino.  This makes the NLSP decay more efficient compared with the KSVZ axion models.  In a specific DFSZ model whose parameters are constrained by various cosmological considerations, we identified the region of the stop and stau masses which is compatible with the required dark matter density. Within the conventional bound on the axion scale $F_a \\lesssim 10^{12}$ GeV, the stau decay to the flaxino LSP never becomes dangerous for BBN. Furthermore, with the chosen parameter sets in the first case ($\\tilde{\\mu}=12$), the stau decay length turns out to be ${\\cal O}$(1) m so that a fair number of stau decays can be captured in colliders when the axion scale is close to its lower bound, $F_a \\sim 10^{10}$ GeV. In the second case ($\\tilde{\\mu}=34$), the flaxino LSP is naturally realized in the broad parameter space of the squark mass. Although the decay length is longer than ${\\cal O}$(10) m, the additional detectors proposed by \\cite{Hamaguchi:2004df,Feng:2004yi,De Roeck:2005bw} will catch the staus even in this case.  Let us finally remark that, in our scheme, the unstable gravitino decaying to ordinary superparticle or flaxino does not contradict with BBN as the primordial gravitino abundance will be diluted away by the entropy production after thermal inflation caused by the flat direction in the model.  A gravitino even lighter than the flaxino is also allowed since the flaxino decay to the gravitino and the axion causes no trouble with BBN."
    },
    "0801/0801.3458_arXiv.txt": {
        "abstract": "I demonstrate that precision timing of millisecond pulsars possess the capabilities of detecting the gravitational effects of intervening galactic substructure.  This analysis is applicable to all types of collapsed baryons including stars, planets, and MACHOs, as well as many types of dark matter, including primordial black holes, scalar miniclusters, and sufficiently dense clumps of cold dark matter.  The physical signal is quantified and decomposed into observable and unobservable components; templates for the observable signals are also presented.  Additionally, I calculate the expected changes in the observed period and period derivatives that will result from intervening matter.  I find that pulsar timing is potentially a very useful tool for probing the nature of dark matter and to learn more about the substructure present within our galaxy. ",
        "introduction": "Our modern understanding of cosmology has revolutionized our picture of our own galaxy.  What was once thought to be an island in the universe dominated by stars is now better modeled as a disk of mostly non-luminous baryons embedded in a much more massive halo composed primarily of dark matter \\citep{Cole:2005sx}.  However, the mass distribution and clumpiness of the both the disk and the halo, as well as the nature of the dark matter, are at present unknown.  Baryons collapse to form molecular clouds, stars, and planets, among other things \\citep{FP:04}.  The location and abundance of the baryonic components of the galaxy are known primarily from their optical effects, such as emission and absorption features.  This tactic does not work for dark matter, though, since it does not interact electromagnetically with any appreciable strength \\citep{CDMS:2005}.  However, both dark matter and baryons do interact gravitationally.  Thus, if a galactic probe that is highly sensitive to the gravitational influence of the components of the galaxy can be identified, it will not only have the potential to detect the baryons that optical searches have missed thus far, but can also be used to detect dark matter within our own galaxy.  Such a galactic probe highly sensitive to gravitational effects has been known since 1982: the millisecond pulsars \\citep{MSP:82}.  Over one hundred millisecond pulsars have already been identified in our galaxy, both in the disk and the halo \\citep{ATNF}.  Because of the stability and high accuracy of the timing measurements of millisecond pulsars, even small changes in the gravitational environment as the light from a pulsar travels to Earth can, in principle, be detected \\citep{LD:95,Fargion:97,Hosokawa:1999,Siegel:2007fz,Seto:2007}.  With this paper, I seek to quantify the physical effects and observable signatures of baryonic and dark matter on the timing measurements of millisecond pulsars, with a view towards using millisecond pulsars to detect heretofore invisible baryons and dark matter.  The layout of this paper is as follows: section 2 discusses the physical gravitational effects of matter on the light-travel-time from a pulsar to Earth.  Section 3 focuses on what component of the physical signal is uniquely ascribable to this gravitational effect, identifying points of confusion and their resolution.  In section 4, templates for identifying transiting matter are created and presented, along with probabilities for detecting dark matter via this method.  Finally, section 5 concludes this paper with a discussion of possible applications and future work as well as briefly addressing more complicated models of dark matter. ",
        "conclusions": "I have demonstrated the quantifiable, observable effects of transiting matter on pulsar timing measurements.  Perhaps most importantly, I have crafted templates detailing the observable signal in the timing residuals as well as predicting the corresponding changes in $P$ and $\\dot{P}$ which accompany such a transit.  These are effects which can lead to the indirect detection of either collapsed baryonic matter or dark matter in between Earth and a millisecond pulsar.  Current timing accuracies for millisecond pulsars are typically of order $\\sim 1 \\, \\mu$s, although there do exist a number of millisecond pulsars with accuracies of $\\sim 100 \\, $ns \\citep{Kaspi:94,vanStraten:01,Jacoby:05,Hotan:2006}, such as B1937+21, J1909-3744, J0437-4715, and J1713+0747.  Additionally, uncertainties in $P$ are typically $\\mathcal{O}(10^{-16} \\, \\mathrm{s})$ and uncertainties in $\\dot{P}$ are typically $\\mathcal{O}(10^{-25})$, although some are more accurate, for example, J1909-3744 has an uncertainty in $P$ of $2.0 \\times 10^{-18} \\, \\mathrm{s}$ and in $\\dot{P}$ of $2.0 \\times 10^{-25}$ \\citep{ATNF}.  While timing accuracies will have to improve by two orders of magnitude to detect clumps of mass $M_\\mathrm{dm} \\simeq 10^{-4} \\, M_\\odot$ with somewhat fortuitous transit parameters, measurements of the uncertainties in $P$ are nearly at the required level to give evidence of a transit today.  (I remind the reader that the $M_\\mathrm{dm}$ used in earlier sections are $0.01 \\, M_\\odot$.)  I recommend that everyone working with millisecond pulsars compare their timing residuals to these templates and also look for the corresponding systematic, coherent changes over time in the period and period derivatives of their pulsars.  While the probability of observing a transit sufficiently close to the LOS in observing any one pulsar is exceedingly small, there are four reasons for optimism.  First, the analysis presented is for one individual pulsar, but over 100 millisecond pulsars are presently known, with the total detectable galactic pulsar population estimated to be approximately 30,000 \\citep{Kiel:2007}, many of which will be discovered and timed to unprecedented accuracy with the Square Kilometer Array \\citep{SKA:2005}.  Second, this assumes a distance to the pulsar of interest of $1 \\, \\mathrm{kpc}$.  While this is a typical distance to the most accurate millisecond pulsars at present, millisecond pulsars have been detected as far away as the globular clusters, and new estimates are that millisecond pulsar timing is feasible for objects as far away as $\\sim 1.2 \\, \\mathrm{Mpc}$ \\citep{LOFAR:2007}, which would open up an extremely long baseline to millisecond pulsars in other galaxies (e.g., Andromeda), should this be correct.  Third, the gravitational effects should be easily separable from other effects (e.g., dust lanes, hot gas, etc.) because of the greyness of the gravitational effects.  While electromagnetic effects cause shifts of differing magnitudes as a function of frequency \\citep{Snez:2007}, the gravitational time delay is frequency independent.  Other gravitational sources, such as gravitational scattering of the pulsar with a star or planet, or a gravitational wave effect on timing, will have a different template signal and can always be discriminated against by continued timing measurements.  Finally, improvements in software \\citep{TEMPO:2006} combined with future improvements in instrumental timing accuracies \\citep{Shearer:08} may indeed lead to residuals knowable to $\\mathcal{O}( 1 \\, \\mathrm{ns})$, which would be a great boon to the range of detectability of transiting matter.  This method is primarily aimed at pulsars located in the galactic disk, where precision timing is most effective.  It is worth noting that there is discussion as well of precision timing and the possibility of observing the Shapiro effect for pulsars in the halo as well \\citep{HaloPulsars:2007}.  Timing of millisecond pulsars is sufficiently accurate that it should be able to determine whether or not matter of a given mass has transited close to the LOS to the pulsar over the timescales of observation.  Projects such as the Parkes Pulsar Timing Array \\citep{Manchester:07} are just coming online at present, and have the ability to detect these effects.  While this analysis is not applicable to extended sources that cross the LOS, it is applicable to any source that does not, as well as any point-like source that does.  Future work may attempt to address that problem, made more complex by the changing effective gravitational potential.  Such a template would be extremely useful for searching for various types of dark matter, such as Scalar Miniclusters \\citep{Zurek:2007} or WIMPy microhalos with small impact parameters.   I acknowledge Kathryn Zurek, first and foremost, for her contributions to this work during the early stages, without which this paper would not have been possible.  I also acknowledge Snezana Stanimirovic for useful discussions concerning pulsar astronomy and timing, and Jim Chisholm for discussions on Primordial Black Holes.  Many thanks as well to Craig Hogan, Jim Fry and Daniel Eisenstein for comments on drafts of this paper."
    },
    "0801/0801.3172_arXiv.txt": {
        "abstract": "Possible influence of primordial black hole (PBH) evaporations on cosmic microwave background (CMB) is investigated. The spectrum distortions of CMB from the black-body spectrum are described by the chemical potential $\\mu$ and the Compton parameter $y$. From COBE/FIRAS limits on $\\mu$ and $y$, the power law index $n$ of primordial density fluctuations and the mass fraction of PBHs $\\beta$ are constrained by employing the peak theory for the formation process of PBHs. Constraints set here are $n < 1.304$ and $n<1.333$ in the thresholds of peaks $\\zeta_{\\rm th} =0.7$ and $\\zeta_{\\rm th} =1.2$, respectively, for the PBH mass range between $2.7\\times 10^{11}$g and $1.6 \\times 10^{12}$ g, and $n < 1.312$ and $n<1.343$ in the thresholds of peaks $\\zeta_{\\rm th} =0.7$ and $\\zeta_{\\rm th} =1.2$, respectively, for the PBH mass range between $1.6 \\times 10^{12} ~{\\rm g}$ and $3.5\\times 10^{13}$ g, which correspond to the comoving scales between $3 \\times 10^{-18}$ Mpc and $ 4\\times 10^{-17}$ Mpc. The constraint on the PBH fraction, which is the direct probe of the amplitude of density fluctuations in these scales, stays at almost the same value as $\\beta<10^{-21}$ in these mass ranges. It is also found that, with these constraints, UV photons injected by PBH evaporations are unlikely to ionize the majority of hydrogen atoms. ",
        "introduction": "The properties of the primordial density fluctuations in large scales are revealed by Wilkinson Microwave Anisotropy Probe (WMAP) and galaxy redshift surveys \\cite{wmap,2df,sdss}. The results of these observations are essentially consistent with the predictions of the inflationary model, i.e., primordial density fluctuations with an almost scale-invariant spectrum and the random Gaussian statistics.  Taking the second order terms of the inflaton potential and density perturbations into account, however, we expect to have the departure from the scale invariant spectrum and the existence of the non-Gaussian components. Therefore, precise measurements of the spectrum and the non-Gaussian components are crucial for a detailed understanding of inflation. For example, the first year WMAP data together with galaxy redshift surveys and Ly-$\\alpha$ measurements preferred the non-zero running spectral index \\cite{wmap}, while the third year data with new Ly-$\\alpha$ measurements do not strongly favor the running \\cite{spergel}.  The Ly-$\\alpha$ measurements are sensitive to the power spectrum at $\\sim 1 ~{\\rm Mpc}$, while measurements of strong gravitational lensing have a sensitivity on much smaller scales such as $\\sim 10$-$100 ~{\\rm kpc}$. These methods are, however, limited in both scales and epochs. So far the best probe of the small scale density fluctuations is provided by the abundance of primordial black holes (PBHs)~\\cite{pbh}. The formation of PBHs took place during the radiation dominated epoch due to the gravitational collapse of the high density region at the horizon scale when the amplitude of over-density exceeded a critical threshold. Therefore the resultant mass spectrum and the abundance of PBHs depend on the amplitude of the power spectrum for the primordial density fluctuations at the horizon crossing epoch. It is known that PBHs eventually evaporate while emitting Hawking radiation~\\cite{hawkingrad}. The lifetime of a PBH is proportional to the cubic of its mass. Therefore, PBHs with mass less than $10^{15}$g should have evaporated away by the present epoch. When they evaporate, they emit black-body radiation and numerous kinds of particles such as neutrinos, electrons and protons.  Since PBHs with mass larger than $10^{15}$g survive in the present universe, we can set a constraint on their abundance from the fact that PBH density cannot exceed the average matter density observed at present. We can therefore introduce $\\beta(M)$ which is the fraction of the regions of mass $M$ collapsing into PBHs~\\cite{beta}. This can be related to the mass fraction of PBHs at the time of formation as $\\beta=\\rho_{{\\rm PBH},i}/\\rho_{r,i}$ where $\\rho_{{\\rm PBH},i}$ and $\\rho_{r,i}$ are the energy densities of PBHs and background radiation at the PBH formation epoch, respectively. The constraint on the density parameter of PBHs at present $\\Omega_{{\\rm PBH},0}<1$ implies $\\beta < 10^{-18} (M/10^{15}{\\rm g})^{1/2}$.  Another important constraint on the abundance of PBHs can be set by the phenomena of evaporation. PBHs emit many kinds of particles when they evaporate. Among these are diffuse gamma rays. The emitted photons from evaporated PBHs, after the recombination epoch, contribute to the diffuse gamma rays. The upper limit on diffuse gamma rays sets a stringent constraint on the abundance of PBHs~\\cite{gamma}. It is also known that big bang nucleosynthesis~\\cite{nucle} and entropy production in the early universe also provide constraints on PBH abundance for various mass ranges (see the review \\cite{abundance}).   In this paper, we investigate spectrum distortions of cosmic microwave background radiation (CMB) caused by PBH evaporation. Naselskii was the first to study the effect of PBHs on the recombination process and to estimate the allowed abundance of PBHs that evaporated around the recombination epoch~\\cite{distort}. Naselskii and Shevelev estimated distortions of CMB due to electrons and positrons from PBH evaporations \\cite{naselskii-shevelev}. Ricotti et al. have also investigated the effects of non-evaporating PBHs on CMB \\cite{ricotti}. X-rays emitted by gas accretion into non-evaporating PBHs were found to modify recombination and reionization. Therefore, by studying the effect of such X-rays on the $y$-distortion and the reionization, they obtained a PBH's constraint which is larger than $10^{15} M_\\odot$. Here, we adopt a modern analysis of CMB spectrum distortions, i.e., $\\mu$ and $y$ distortions on which COBE/FIRAS has set stringent observational limits, where $\\mu$ and $y$ are the chemical potential and the Compton $y$-parameter, respectively. The process we consider here is that photons which have evaporated from PBHs directly induce the $\\mu$-distortion in CMB and also hit electrons in the surrounding medium. Accordingly, these electrons scatter CMB photons and generate the $y$-distortion. While electrons directly emitted by PBHs also generate the $y$-distortion via inverse Compton scattering, resultant photons are too energetic to be observed as CMB distortions in COBE/FIRAS data. Therefore we ignore electron emission from PBHs in this paper.  Distortions are created when the energy injection into CMB arises between redshift $z\\sim 10^6$ and recombination $z\\sim 1000$. Before $z \\sim 10^6$, CMB can recover the black-body distribution via double Compton scattering or free-free emission. After recombination ($z < 1000$), there is almost no free electron to be a target of photons. Therefore we calculate CMB distortions caused by photons emitted by PBHs in this epoch. Consequently, we can obtain the constraint on the evaporated PBH abundance using COBE/FIRAS limits, $|\\mu|<9\\times 10^{-5}$ and $y<1.5 \\times 10^{-5}$ \\cite{firas}. Note that the mass range of PBHs constrained here is between $10^{11}$ and $10^{13}$g, since this mass range corresponds to the PBHs whose lifetimes are within $1000<z< 10^6$.    In order to estimate the number density of PBHs, we need the threshold value of density fluctuations required for the overdense regions to collapse into PBHs. The threshold value was traditionally calculated as the critical density contrast in the simplified cosmological model by the analytic method, where the value obtained is $\\delta_c = 1/3$ for the radiation dominant era \\cite{pbh}. There has been resent progress in estimating the threshold value by employing numerical simulations based on general relativity~\\cite{n-j, s-s}. In this paper, therefore, we adopt the threshold value of Shibata and Sasaki \\cite{s-s}.  This paper is organized as follows. In Sec. II we review the formula for evaluating the PBH mass function by using the peak theory. In Sec. III we calculate CMB distortions, $\\mu$ and $y$, caused by Hawking radiation from PBHs and we set constraints on the primordial power spectral index and the PBH abundance. Conclusions are presented in Sec. IV. In this paper we assume that $h=0.70 \\ (H_0=h \\times 100 ~{\\rm km/s \\cdot Mpc})$, $\\Omega _{\\rm b} h^2 =0.022$ and $\\Omega_{\\rm M} h^2 =0.11$ \\cite{spergel}. And we set $c=\\hbar=1$, where and $c$ and $\\hbar$ are the speed of light and Planck's constant over $2\\pi$, respectively. ",
        "conclusions": "In this paper, we investigated the possible influence of PBH evaporation on CMB. We set new and stringent constraints on the PBH mass fraction and the primordial power spectrum of density fluctuations from the observational upper bounds of $\\mu$ and $y$-distortions obtained by COBE/FIRAS. For this analysis, we employed the peak theory, with a threshold described as $\\zeta_{\\rm th}$.  It was shown that $\\mu$ and $y$-distortions set limits on different ranges of PBH masses. From the $\\mu$-distortion, we can test the mass range between $ 2.7\\times 10^{11} {\\rm g}$ and $ 1.6 \\times 10^{12}$g, which correspond to $M_{\\rm DC}$ and $M_{\\rm freeze}$, that is, the masses of PBHs which evaporate away at the decoupling epochs of the double Compton and the thermalization, respectively. We obtain constraints on the power law spectral index as $n<1.304$ and $n<1.333$ for $\\zeta_{\\rm th} =0.7$ and $\\zeta_{\\rm th} =1.2$, respectively. For the PBH abundance, we set the limit $\\beta<10^{-21}$ for the mass range we considered here.  On the other hand, from the $y$-distortion, we investigated the PBH mass range between $1.6 \\times 10^{12} {\\rm g}$ and $3.5\\times 10^{13}$g, which correspond to $M_{\\rm freeze}$ and $M_{\\rm RC}$. Here $M_{\\rm RC}$ corresponds to the mass of PBHs which evaporate away at the recombination epoch. We obtained constraints on the power law spectral index as $n<1.312$ and $n<1.343$ for $\\zeta_{\\rm th} =0.7$ and $\\zeta_{\\rm th} =1.2$, respectively. For the PBH abundance, we set the limit $\\beta<10^{-21}$ for the mass range considered here.    It turns out that our constraints on the spectral index are looser than those obtained from the observation of CMB temperature fluctuations by WMAP satellite and other large scale observations such as galaxy surveys of 2dF and SDSS, and surveys of Ly-$\\alpha$ forest, which imply a spectral index $n = 0.947 \\pm 0.015$ \\cite{wmap}. However, it is suggested by many authors that general inflation models produce not a simple power law spectrum but a spectrum with a running spectral index which has many branches. Accordingly it may be no use to constrain the power law index from observations on various scales. All we can do is to set constraints on the amplitude of each scale associated with the observation. From this point of view, the mass fraction of PBHs $\\beta$ can provide unique information for the fluctuation amplitude on very small scales, $M_{\\rm DC} < M < M_{\\rm RC}$ which corresponds to the comoving scales between $3 \\times 10^{-18}$Mpc and $ 4\\times 10^{-17}$Mpc. Since these small scales correspond to the Horizon scales right after the end of inflation, we can say that the last stage of inflation can by revealed by PBHs.     Finally, we would like to mention that PBHs can be formed not only by the primordial density perturbations but also by the collisions of bubbles of the broken symmetry phase \\cite{bubble} or by the collapse of cosmic strings \\cite{string}. Our constraint on $\\beta$ is even applicable for such PBHs."
    },
    "0801/0801.1494_arXiv.txt": {
        "abstract": "We report on interferometric imaging  of the CO J=1--0 and J=3--2 line emission from the  controversial QSO/galaxy pair HE\\,0450--2958.  {\\it The  detected CO  J=1--0 line  emission is  found associated  with the disturbed  companion galaxy not  the luminous  QSO,} and  implies $\\rm M_{gal}(H_2)\\sim (1-2)\\times 10^{10}\\,  M_{\\odot}$, which is $\\ga 30\\% $ of  the dynamical  mass in its  CO-luminous region.  Fueled  by this large gas  reservoir this galaxy is  the site of  an intense starburst with $\\rm  SFR\\sim 370\\, M_{\\odot  }\\, yr^{-1}$, placing it  firmly on the  upper  gas-rich/star-forming   end  of  Ultra  Luminous  Infrared Galaxies  (ULIRGs,  $\\rm L_{IR}>10^{12}\\,  L_{\\odot  }$).  This  makes HE\\,0450--2958 the  first case of  extreme starburst and  powerful QSO activity, intimately  linked (triggered  by a strong  interaction) but not coincident.  The lack of CO emission towards the QSO itself renews the controversy regarding its host  galaxy by making a gas-rich spiral (the  typical  host  of  Narrow  Line  Seyfert~1  AGNs)  less  likely. Finally,  given  that  HE\\,0450--2958  and similar  IR-warm  QSOs  are considered   typical   ULIRG$\\rightarrow   $(optically   bright   QSO) transition candidates, our results raise the possibility that some may simply  be  {\\it  gas-rich/gas-poor (e.g.   spiral/elliptical)  galaxy interactions} which ``activate'' an optically bright unobscured QSO in the gas-poor  galaxy, and a starburst  in the gas-rich  one.  We argue that such  interactions may  have gone largely  unnoticed even  in the local  Universe   because  the  combination  of   tools  necessary  to disentagle the progenitors (high  resolution and S/N optical {\\it and} CO imaging) became available only~recently. ",
        "introduction": "HE\\,0450-2958  is an optically  bright ($\\rm  M_v=-25.8$) IR-selected QSO at a redshift of z=0.286 (Low et al. 1988), and one of the few in the  local Universe  whose position  in  the (60$\\mu  $m, 25$\\mu  $m) versus (100$\\mu $m, 60$\\mu $m) far-IR color-color diagram lies in the area between Ultra Luminous  Infrared Galaxies (ULIRGs) and optically bright QSOs defined  by $\\rm -0.8 < \\alpha (100, 60)  < 0.4$ and $\\rm -2.0 <  \\alpha (60, 25)  < - 0.8$  (Canalizo \\& Stockton  2001).  Its complex  environment, marked  by tidal  interactions,  was identified early  with ground-based  (Hutchings,  \\& Neff  1988)  and {\\it  HST} imaging  (Boyce et  al.  1996).   These efforts  revealed  a strongly interacting system consisting of the QSO and a disturbed galaxy $\\sim 1.5''$ away, setting this QSO/galaxy  pair on par with typical ULIRGs in terms  of morphology, and  IR luminosity ($\\rm  L_{IR}\\sim 5\\times 10^{12}\\, L_{\\odot}$).  The  scarcity  of such  IR  {\\it  and}  optically luminous  QSOs  made HE\\,0450-2958   an  early   favorite  candidate   for   undergoing  a ULIRG$\\rightarrow $(optically  bright QSO) transition  (Hutchings, \\& Neff 1988; Canalizo \\& Stockton 2001) in a scenario first outlined by Sanders  et   al.   (1988).   The  latter  was   proposed  after  FIR color-color diagrams have  identified a population of dust-enshrouded QSOs via  their AGN-heated  dust component (de  Grijp, Miley,  \\& Lub 1987), and  it was motivated  by the similar  bolometric luminosities and  space densities  of optically  powerful QSOs  and ULIRGs  in the local  Universe.   It involves  two  gas-rich  galaxies whose  strong dynamical  interaction induces  a  starburst and  fuels  an AGN  {\\it within} heavily dust-obscured  environments.  Then, as star formation uses  and  disperses  their  large molecular  gas  reservoirs,  their initially cool  Spectral Energy  Distribution (SED) of  dust emission changes  towards one  dominated  by warm  AGN-heated  dust with  much warmer  IR   colors.   Eventually  an   unobscured  optically  bright QSO~emerges out of the original  ULIRG.  The large dust and molecular gas reservoirs found in {\\it optically} selected QSOs where they fuel intense starbursts  (e.g.  Alloin  et al.  1992;  Haas et  al.  2000; Evans  et al.  2001)  certainly support  an AGN-starburst  link, with dynamical  interactions/mergers  as   the  likely  trigger  for  both activities,  suggested also  by  optical studies  (e.g.  Canalizo  \\& Stockton   2001).   In  this   context  the   pivotal  role   of  the ULIRG$\\rightarrow $QSO  transition objects can then  be understood as the ``markers'' of a special  evolutionary stage during which the QSO emerges  from a  still ongoing  dusty starburst,  and lasting  only a fraction of already short  molecular gas consumption timescales ($\\rm \\la 10^8\\,yrs$).   HE\\,0450-2958  acquired further singular  significance when  Magain et al.   (2005), after  carefully subtracting  the AGN  emission  from an ACS/HST  image, failed  to  find  the QSO  host  galaxy expected  from well-established  and  tight   correlations  of  (host  galaxy)-quasar properties (e.g.   McLure \\& Dunlop  2002; Floyd et al.   2004).  This potentially  pivotal result  has been  upheld by  independent analysis (Kim  et  al.  2007),  and  follow-up  VLT spectroscopic  observations weakened the  hypothesis for a significantly  dust-enshrouded QSO host galaxy  (Letawe  et  al.    2007).   The  possibility  of  an  ejected (black-hole)+(fuel)  system as  responsible for  a ``naked''  QSO, its serious consequences and observable  signatures have been discussed in several papers.  Summarizing their  results here, such a configuration can be produced either by a Newtonian three-body ``kick'' exerted on a lighter  black hole  by a  black hole  binary residing  in  a gas-poor elliptical (Haehnelt, Davies, \\& Rees  2005; Hoffman \\& Loeb 2006), or by the recoil of a coalesced  binary of {\\it spinning} black holes due to an  asymmetric gravitational wave emission (Haehnelt  et al.  2005; Loeb  2007).  The  stage for  both  scenaria demands  a strong  galaxy interaction,  and both have  important consequences  for gravitational wave  detections by  future  detectors  such as  {\\it  LISA}, and  the hierarchical build-up  of supermassive  black holes in  galaxy centers (Haehnelt et al.  2005 and references~therein).   In the  present work we report  on the observations of  CO J=1--0, the prime molecular line  used to trace metal-rich H$_2$  gas in galaxies, using the  Australian Telescope Compact  Array (ATCA), and  the J=3--2 line  with the Smithsonian  Submillimeter Array  (SMA).  The  paper is organized  as  follows:  a)  we  present the  observations  and  their analysis (section 2),  b) estimate the molecular gas  and dust mass of the system and discuss the properties of the companion galaxy (Section 3),  c)  present the  implications  regarding  the  AGN and  its  host (section   4),  and   d)  discuss   possible  ramifications   for  the ULIRG$\\rightarrow  $QSO evolutionary  scheme (section  5).  Throughout this     work      we     adopt     a      cosmology     with     $\\rm H_{\\circ}=71\\,km\\,s^{-1}\\,Mpc^{-1}$,  $\\rm  \\Omega  _M=0.27$ and  $\\rm \\Omega  _{\\Lambda}=0.73$,   for  which  the   luminosity  distance  of HE\\,0450-2958  at  z=0.286  is   $\\rm  D_{L}=1460.4\\,  Mpc$  and  $\\rm 1''\\rightarrow 4.28\\,kpc$. ",
        "conclusions": "We performed  CO line emission  imaging observations of  the enigmatic QSO/galaxy pair HE\\,0450-2958  using the ATCA (CO J=1--0)  and the SMA (CO J=3--2), detecting  the J=1--0 line and placing  an upper limit on the J=3--2 line, our conclusions are as follows   \\noindent 1. The CO J=1--0 line emission is associated with the companion galaxy not  the QSO, it  corresponds to  $\\rm M_{gal}(H_2)  \\sim (1-2)\\times 10^{10}\\,  M_{\\odot}$ and  amounts to  $\\ga  30 \\%$  of the  enclosed dynamical mass  within $\\sim 8$\\,kpc. This large  gas reservoir fuels extreme  star formation at  a rate  of $\\rm  \\sim 370\\,  M_{\\odot }\\, yr^{-1}$,  efficiencies  of  $\\rm L_{IR}/M_{gal}(H_2)\\sim  (90-165)\\, M_{\\odot }/L_{\\odot }$, and places  the companion galaxy on the upper gas-rich and star  forming end of ULIRGs.  This  property, along with the  disturbed   nature  of  the  system,   solidifies  the  strongly interacting status of this controversial galaxy/QSO pair.   \\noindent 2. No  molecular gas  is detected  towards  the QSO  itself with  $\\rm M_{QSO-host}(H_2)/M_{gal}(H_2)\\leq 1/5$,  consistent with the little dust extinction or star  formation activity deduced for its putative host galaxy.   This weakens the case  for a host  galaxy typical for Narrow Line  Seyfert 1  AGNs since these  are gas-rich  spirals with luminous circumnuclear starbursts.   HE\\,0450-2958 is thus the first known  case of  a gas-rich  extreme  starburst and  a powerful  QSO, intimately linked (via a strong interaction), but not~coincident.   \\noindent 3. Most of the radio continuum in this system is due to low-brightness AGN-driven jets rather than  star formation activity.  Deeper radio continuum observations  are needed  to reveal their  morphology and extent, which in turn can yield more clues regarding its AGN.   \\noindent 4. The early  prominence of HE\\,0450-2958  as an IR-selected  QSO that exemplifies   the   ULIRG$\\rightarrow   $(optically  luminous   QSO) transition  in  a  popular  evolutionary scenario,  and  its  hereby revealed nature  as a  strong interaction between  a gas-rich  and a gas-poor, possibly elliptical, galaxy  opens up the possibility that some  of these  so-called  transition objects  are such  interacting pairs  instead.  Such   gas-poor/gas-rich  galaxy  interactions  can activate a  starburst in the  gas-rich progenitor, and  an optically luminous AGN  in the gas-poor  one.  Their composite  dust continuum would then  be typical of IR-warm  ULIRG$\\rightarrow $QSO transition objects.  Sensitive  optical and CO  imaging at high  resolution can now  reveal  and  study   such  dynamical  interactions,  and  these transition objects are excellent candidates."
    },
    "0801/0801.2424_arXiv.txt": {
        "abstract": "Foreground subtraction is the biggest challenge for  future redshifted $21 \\, {\\rm cm}$ observations  to probe reionization. We use a short  GMRT observation at $153 \\, {\\rm  MHz}$ to  characterize the statistical properties of the background radiation across  $\\sim1^{\\circ}$ to sub-arcminutes  angular scales, and across  a  frequency band of $5 \\, {\\rm MHz}$ with $62.5 \\, {\\rm kHz}$ resolution. The statistic we use is the visibility correlation function, or equivalently the angular power spectrum $C_l$. We present the results obtained from using relatively unsophisticated, conventional data calibration procedures. We find that even fairly simple minded calibration allows one to estimate the visibility correlation function at a given frequency $V_2(U,0)$.  From our observations we find that $V_2(U,0)$ is consistent with foreground model predictions at all angular scales except the largest ones probed by our observations where the the model predictions are somewhat in excess. On the other hand the  visibility correlation between  different frequencies $\\kappa(U,\\Delta \\nu)$, seems to be much more sensitive to calibration errors. We find a rapid decline in $\\kappa(U,\\Delta \\nu)$, in contrast with the prediction  of less than $1 \\%$ variation across  $2.5 \\, {\\rm MHz}$. In this case however, it seems likely that a substantial part of the discrepancy may be due to limitations of data reduction procedures. ",
        "introduction": "Observations of  redshifted $ 21 {\\rm cm}$ radiation from the large scale distribution of neutral  hydrogen (HI) are perceived as one of the most promising future probes of the Universe at high redshifts (see \\citealt{furla} for a recent review). Observational evidence from  quasar absorption spectra \\citep{becker,fan} and the CMBR  \\citep{spergel,page} together imply that the HI was reionized  over an extended period spanning the redshift range $6 \\le z \\le 15$ (for reviews see \\citealt{barkana,fck06,cf06}).  Determining how and when the Universe was reionized is one of the most important issues that will be addressed by future $21 {\\rm cm}$ observations. The Giant Meter Wave  Radio Telescope  (GMRT  \\footnote{http://www.gmrt.ncra.tifr.res.in}; \\citealt{swarup}), currently  functioning at several frequency bands in the  range $150$ to $ 1420 \\, {\\rm MHz}$ is very well suited for carrying out initial investigations towards detecting the reionization HI signal.  There are several upcoming    low-frequency instruments  such  as   LOFAR\\footnote{http://www.lofar.org/}, MWA\\footnote{http://www.haystack.mit.edu/arrays/MWA}, 21CMA\\footnote{http://web.phys.cmu.edu/$\\sim$past/} and   SKA\\footnote{http://www.skatelescope.org/} which are being built specifically with these observations in view.  It is currently perceived that a statistical analysis of the fluctuations in the redshifted $21 \\, {\\rm cm}$ signal holds the greatest potential for observing   HI at high redshifts \\citep{BS1,zald,morales,BA5,BP5}. Correlations among the visibilities measured in radio-interferometric observations directly probe the HI power spectrum  at the epoch where the radiation originated. The reionization visibility signal at the GMRT is expected to be $\\sim 1 \\,{\\rm mJy}$ and smaller (\\citealt {BA5}).  This HI signal is present as a minute component of the background in all low frequency observations, and it is buried in foreground radiation from other astrophysical sources whose contribution is $4$ to $5$ orders of magnitude larger. Extracting the HI signal from the foregrounds is a major challenge.  Individual sources can be  identified and removed from the image at a flux level which depends on the sensitivity. The contribution from the remaining discrete sources could be large enough to overwhelm the  HI signal \\citep{dmat1}.  The diffuse synchrotron emission from our Galaxy \\citep{shaver} is  another  important component. Foreground sources include  free-free  emission from ionizing halos  \\citep{oh}, faint   radio loud quasars \\citep{dmat1} and synchrotron emission from low redshift galaxy clusters \\citep{dmat}.  The foregrounds are expected to have a continuum spectra, and the contribution at two different frequencies separated by $\\Delta \\nu \\sim 1 \\,{\\rm MHz}$ are expected to be highly correlated. The HI signal is expected to be  uncorrelated at such a frequency separation and this holds the promise of allowing us to separate the signal from the foregrounds. A possible approach is to subtract a best fit continuum spectra for each line of sight \\citep{wang} and then use the residuals to determine the HI power spectrum. An alternate approach is to first determine the statistical properties of the total radiation and then subtract out the   smooth  $\\Delta \\nu$ dependent part to extract the HI signal \\citep{zald}. The issue of foreground removal has also been studied by \\citet{mor2} and \\citet{mcquinn}.  It is crucial to accurately characterize the foregrounds in order to be able to detect the HI signal  in future observations. In this paper we used  GMRT observations to characterize the foregrounds  at $153 \\, {\\rm MHz}$ which corresponds to an HI signal from $z=8.3$. To the best of our knowledge this is the first attempt to directly characterize the foregrounds at angular scales ($\\sim 1^{\\circ}$ to sub-arcminute) and frequency coverage  ($6 \\, {\\rm MHz}$ with $62.5 \\, {\\rm kHz}$ resolution) relevant for detecting the  reionization HI signal.  We next present a brief outline of the paper. In Section 2 we describe the observations and data reduction while in Section 3   we present ``visibility-correlations'' which we use to quantify the statistical properties of our radio-interferometric data.  Section 4 presents the predictions of existing foreground models, and in Section 5 we present our results and discuss their implications. ",
        "conclusions": "We have determined the observed value  $V_2(U,\\Delta \\nu)$ of the visibility correlation  estimator $\\hat{V}_2(U,\\Delta \\nu)$ for    {\\bf data I} and {\\bf data  R}  which are before and after source subtraction  respectively. Baselines in the range $20 \\le U \\le 2 \\times 10^4$, and frequency channels $21$ to $100$ were used for the analysis. Visibilities  $V(\\u+\\Delta \\u,\\nu + \\Delta \\nu)$ within the  disk  $\\mid \\Delta \\u \\mid \\le D=5$ were correlated with $V(\\u,\\nu)$. Here  $\\Delta \\nu$ was restricted to $\\mid \\Delta \\nu \\mid < 2.5 \\, {\\rm MHz}$ which corresponds to  a separation of $40$ channels. Note that the correlation of a visibility with itself was not included. The value of $D$ was chosen such that it is both less than $(\\pi \\theta_0)^{-1}=8$, and  also large enough that a reasonable number of visibility pairs that contribute to the correlation. Figure \\ref{fig:r1} shows $V_2(U,\\Delta \\nu)$ as a function of $U$ for  $\\Delta \\nu=0$.  Equivalently, we may also interpret this as   the multi-frequency angular  power spectrum $C_l(\\Delta \\nu)$  at $\\Delta \\nu=0$.  For both the data-sets  the real part of $V_2(U,0)$ is found to be considerably larger than the imaginary part. This  is consistent with  the discussion of  Section 3.1, and we expect the real  part to provide an estimate of the foreground contribution $V_2(U,0)$.  The $1-\\sigma$ error bars shown in the figure have been determined based on the error estimates discussed in Section 3.2. The uncertainty in $F_2(U,0)$ is mainly  due to  the limited number of independent estimates, the system noise makes a smaller contribution.  Though the results for {\\bf data I} over the range $200 \\le U \\le \\le 2 \\times 10^4$ looks like a power law $V_2(U,0)\\propto U^{-\\alpha}$ with a very small slope $0 \\le \\alpha \\le 0.25$,  we do not find a fit with an acceptable value of $\\chi^2$ per degree of freedom.  The real part of $V_2(U,0)$ falls to nearly one-fourth of its original value at most of the  $U$ bins when the directly detected sources are subtracted out. This indicates that a large part of the contribution to $V_2(U,0)$ in {\\bf data I}  is from these resolved sources, and we may interpret $V_2(U,0)$ as arising primarily from  these sources. {\\bf Data R} is expected to contain contributions from point sources below  the detection limit of our image, diffuse sources, system noise, limitations in our imaging and source subtraction procedure and residual RFI. We will assume for the moment that these effects can be ignored, but return to this issue later in this section.  \\begin{figure*} \\includegraphics[width=70mm,angle=-90]{vcorc.ps} \\caption{This shows the real (upper curve) and imaginary (lower curve) parts of the observed visibility correlation $V_2(U,0)$ as a function of $U$ for the two   data-sets indicated in the figure. As shown here, this may also be interpreted as $C_l(0)$ as a function of $l$.} \\label{fig:r1} \\end{figure*}  Figure \t\\ref{fig:r2} shows the observed $V_2(U,0)$ plotted against the predictions of the foreground models discussed in Section 4.  The brightest source in our image has  flux $890 \\, {\\rm mJy}$. Based on this we use  $S_c=900 \\, {\\rm mJy}$ for the point  source contribution to {\\bf data I}.  The clustering of point sources  dominates at baselines $U<150$ ($\\theta >0.7^{\\circ}$), while the Poisson fluctuations  of the point sources dominates at larger baselines.  The diffuse Galactic synchrotron radiation is much smaller than the point source contribution  at all baselines. The errors in the model prediction are quite large and are mainly  due to the Poisson fluctuations of the point sources. The model predictions are found to be  consistent with  the observed values of $V_2(U,0)$ except at the smallest $U$ value which corresponds to an angular scale of $\\sim 1.8^{\\circ}$.  At these baselines the convolution with the primary beam pattern (eq. (\\ref{eq:a4})) becomes important. We have not included this, and the  actual model predictions would  possibly be somewhat smaller if  this were  included. As noted in Section 4., the amplitude of the synchrotron contribution  is very sensitive to the value of the spectral index. The amplitude decreases by a factor of $\\sim 18$ if $\\bar{\\alpha}=2.5$ instead of the value $\\bar{\\alpha}=2.8$ used here. This changes the total foreground contribution only at small baselines ($U \\le 100$) where the model then becomes consistent with our observations.  The limiting flux for source subtraction is $\\sim 8 \\, {\\rm mJy}$, and hence we use  $S_c= 10 \\, {\\rm mJy}$  for {\\bf data R}. The model prediction is dominated by Galactic synchrotron radiation at $U < 150$,  point source clustering in the range $150 \\le U \\le 2 \\times 10^3$ and point source Poisson fluctuations at $U > 2 \\times 10^3$. The model predictions fall short of the observations at all baselines except the smallest $U$ value where it overshoots the observations.  Since the model prediction for $S_c=10 \\, {\\rm mJy}$ falls very much short of  the observations, we also consider $S_c=100 \\, {\\rm mJy}$ where the dominant contribution is Galactic synchrotron at $U < 60$,  point source clustering in the range $60 \\le U \\le 400$ and point source Poisson fluctuations at $U > 400$. We find that the observations are slightly above the $1-\\sigma$ error-bars at  baselines  $U > 100$, whereas they exceed the model predictions at baselines $U<100$. A point to  note is that at the smallest baseline the  prediction for the Galactic synchrotron  radiation exceeds the  observation. This may  be a consequence of the possibility  that the background radiation is relatively  low in the direction of our observation. Estimates from the  \\citet{haslam} map at $408 \\, {\\rm MHZ}$ show a relatively low brightness temperature of $\\sim 30 \\, {\\rm K}$ towards the direction of our observation.  \\begin{figure*} \\includegraphics[width=70mm,angle=-90]{vcormodel.ps} \\caption{The thick solid line shows the real  part of the observed visibility correlation $V_2(U,0)$ as a function of $U$ for the two data-sets indicated in the figure. As shown here, this may also be interpreted as   $C_l(0)$ as a function of $l$.  For {\\bf data I} the thin solid line  shows the total model prediction for $S_c=900  \\, {\\rm mJy}$. Also shown are  the contributions from point source Poisson (dash-dot), point source clustering (dot) and Galactic synchrotron (dash-dot-dot-dot).  For {\\bf data R} the  thin solid line shows the total model predictions for $S_c=100 \\, {\\rm mJy}$     and and the long dashed line for $10 \\, {\\rm mJy}$. The dash-dot-dot-dot curve shows the Galactic synchrotron contribution.} \\label{fig:r2} \\end{figure*}   We quantify the $\\Delta \\nu$ dependence of $V_2(U,\\Delta \\nu)$ using  $\\kappa(U,\\Delta \\nu)$ which is defined as \\begin{equation} \\kappa(U,\\Delta \\nu)=\\frac{V_2(U,\\Delta \\nu)}{V_2(U,0)}\\,. \\end{equation}  We expect the visibilities $V(\\u,\\nu)$  and $V(\\u,\\nu+\\Delta \\nu)$ to get decorrelated as $\\Delta \\nu$ is increased, and hence we expect $0 \\le \\mid \\kappa(U,\\Delta \\nu) \\mid \\le 1$. Figure \\ref{fig:r3} shows $\\kappa(U,\\Delta \\nu)$ for different values of $U$. The foreground models predict a smooth $\\Delta \\nu$ dependence for $\\kappa(U,\\Delta \\nu)$.  The departure from $\\kappa(U,\\Delta  \\nu)=1$ is predicted to be less than $1 \\%$  for $\\Delta \\nu < \\, 2.5 \\,{\\rm MHz}$.  The observed behavior of $\\kappa(U,\\Delta \\nu)$ is quite different from the model predictions.  At the small baselines $U < 1000$ we find that $\\kappa(U,\\Delta \\nu)$ falls sharply within  the first three  channels. In the $U=47$ bin $\\kappa(U,\\Delta \\nu)$ fluctuates at large $\\Delta \\nu$ whereas it remains roughly constant at $U=360$. In both cases this  value of $\\kappa(U,\\Delta \\nu)$  is smaller for {\\bf  data R} as  compared to {\\bf data I}.  At $U=2200$, for {\\bf data I} $\\kappa(U,\\Delta \\nu)$  falls gradually  with increasing $\\Delta \\nu$, and the  visibilities are uncorrelated ($\\kappa(U,\\Delta \\nu)\\sim 0$) by $\\Delta \\nu \\sim 2.5 \\, {\\rm MHz}$. Interestingly, for {\\bf data R} we find that $\\kappa(U,\\Delta \\nu)$ shows a sudden increase to $\\kappa(U,\\Delta \\nu)>1$ at very small $\\Delta \\nu$ ($< 0.5 \\, {\\rm MHz}$), after which  $\\kappa(U,\\Delta \\nu)$ falls and becomes negative by $\\Delta \\nu \\sim < 2 \\, {\\rm  MHz}$. It appears that in this $U$ bin our source subtraction procedure has introduced excess correlations between the visibilities at small $\\Delta \\nu$ and introduces anti correlations at large $\\Delta \\nu$. At $U=4200$, for {\\bf data I}  the value of $\\kappa(U,\\Delta \\nu)$ oscillates with increasing $\\Delta \\nu$. At large $\\Delta \\nu$ {\\bf data  R} also shows  a similar behavior except that the  $\\kappa(U,\\Delta \\nu)$ values are smaller. The behavior of {\\bf data R} is quite different from that of {\\bf data I} at  very small $\\Delta \\nu$  where there are two small oscillations that cross $\\kappa(U,\\Delta \\nu)=1 $.  \\begin{figure*} \\includegraphics[width=100mm,angle=-90]{delnuc.ps} \\caption{This shows $\\kappa(U,\\Delta \\nu)$ as a function of $\\Delta \\nu$ for the different $U$ values shown in the figure. The upper curve (at large $\\Delta \\nu$)  shows {\\bf data I} while the lower shows {\\bf data R}.} \\label{fig:r3} \\end{figure*}  The first point that emerges from our results is that the observed visibility correlations $V_2(U, 0)$ is consistent with the predictions of the  existing foreground models at all baselines except the smallest one  which probes angular scales $\\sim 1^{\\circ}$. The observations are in excess of the model prediction at the smallest baseline. The second point is that $V_2(U,\\Delta \\nu)$ shows considerable $\\Delta \\nu$ dependence,  there being changes of order unity within $\\Delta \\nu= 2.5 \\, {\\rm MHz}$.  This rapid change in the visibilities $V_2(U,\\nu)$ across frequency channels is contrary to the foreground models which predict changes less than $1 \\%$.  It is well appreciated that accurate subtraction of the foreground emission requires very exacting calibration. In contrast,  we have followed fairly standard calibration procedures. As such it seems likely that the discrepancy between our observations and existing predictions is probably not genuine; indeed there are a several purely instrument related  possibilities that may account for the discrepancies between our observational findings and existing models for the foreground emission. We take up first  the issue of calibration error which will introduce phase and amplitude errors in the visibilities. The fact that the values of $\\kappa(U,\\Delta \\nu)$ are generally smaller for {\\bf data R} as compared to {\\bf data I} may be interpreted as indicating that the visibilities $V(U,\\nu)$  are a combination of two parts, a correlated part which arises from for e.g. the effect of calibration errors on discrete sources, and another whose contribution to  different channels is uncorrelated. The ``halos'' that we see around the bright sources is a clear indication that calibration problems exist in our data. Phase errors  which vary with channel would cause decorrelation of the visibilities across  different frequencies. Further, one would expect that the phase errors increase with increasing baseline length, which is qualitatively consistent with what we see in Fig.~\\ref{fig:r3}. In contrast to the situation for $\\kappa(U,\\Delta \\nu)$, the contribution from the source subtraction residuals  to $V_2(U,0)$ (Fig.~\\ref{fig:r2}) can be estimated to be small as follows. There are only $\\sim 100$ imaging artifacts with absolute value of flux  $> 20$~mJy ({\\bf Data R}, Figure \\ref{fig:a3}) , while about 10,000 such sources would be needed to  produce the observed visibility correlation of $\\sim 4\\,$Jy$^{2}$ ({\\bf Data R}, Figure \\ref{fig:r2}).    The 2D Fourier relation between  the sky brightness and the visibilities assumed in Section~3 is not strictly valid for GMRT's large field of view $(\\theta_{\\rm FWHM}=3.8^{\\circ}$). In addition to $u-v$ which are the components of the baseline in the plane normal to the direction of observation, it is also necessary to consider $w$ the component along the observing direction. This is a possible source of error in our visibility correlation analysis. To asses the impact of the $w$ term we have repeated the analysis using only a limited range of baselines for which $w\\le 100$. We find that limiting the maximum $w$ value  does not make  any  qualitative change in our results. The conclusions are unchanged even if we impose $w \\le 50$.    Residual RFI is another  possibility. The visibilities were clipped at $12 \\, {\\rm Jy}$ (Section 2.) and this is expected to remove the strong RFI, but weak RFI contributions will persist in the data. The RFI electric fields  at  any two antennas  is correlated with a time delay $\\tau$ which  depends on  position of the RFI source relative to the antennas and the direction of observation. The RFI contribution behaves like the system noise if $\\tau$ is greater than $\\tau_c$ the coherence time of the RFI signal. In this case the RFI effectively increases $\\sigma$ the rms. fluctuations of the visibilities. This only changes the error estimates,  and does not affect the expected visibility correlations.  RFI sources for which  $\\tau <\\tau_c$ are expected to affect the visibility correlations. This contribution will depend on the distribution of the time delays $\\tau$s and the frequency spectrum of the  RFI sources. The analysis of this is beyond the scope of this paper. Work is currently underway at the GMRT to implement more sophisticated real time as well as offline RFI mitigation schemes.  Future  observations will help assess the improvement that these schemes as well as better calibration procedures make on the problem of foreground subtraction. Polarization leakage is another important  issue that we plan to take up in future work."
    },
    "0801/0801.3614_arXiv.txt": {
        "abstract": "We present a systematic study of the HNCO, C$^{18}$O, $^{13}$CS, and C$^{34}$S emission towards 13 selected molecular clouds in the Galactic center region \\footnote{Based on observations carried out with the IRAM 30-meter telescope. IRAM is supported by INSU/CNRS (France), MPG (Germany) and IGN (Spain).}. The molecular emission in these positions are used as templates of the different physical and chemical processes claimed to be dominant in the circumnuclear molecular gas of galaxies. The relative abundance of HNCO shows a variation of more than a factor of 20 among the observed sources. The HNCO/$^{13}$CS abundance ratio is highly contrasted (up to a factor of 30) between the shielded molecular clouds mostly affected by shocks, where HNCO is released to gas-phase from grain mantles, and those pervaded by an intense UV radiation field, where HNCO is photo-dissociated and CS production favored via ion reactions. We propose the relative HNCO to CS abundance ratio as a highly contrasted diagnostic tool to distinguish between the influence of shocks and/or the radiation field in the nuclear regions of galaxies and their relation to the evolutionary state of their nuclear star formation bursts. ",
        "introduction": "The central few hundred parsecs of many galaxies, including our own, harbor large amounts of molecular gas and are the sites of past, present or future star formation. The interaction of the gas with the gravitational potential in such areas, with the radiation and outflows from recently formed and evolved stars, and with the central objects makes that the physical processes influencing this molecular gas are much more complex than in disk clouds. These physical processes not only dominate the energy balance in the circumnuclear environment of galaxies, but they also determine if stars could be formed within such clouds. In the central 300 pc of our own galaxy we find giant molecular clouds (GMC) harboring regions of massive star formation and heavily affected by shocks \\citep{Pintado01}; large regions where molecular gas is pervaded by intense UV photodissociating radiation (PDRs) stemming from massive star clusters \\citep{Krabbe95,Figer99,Rodriguez01}; strong X-ray emission \\citep[][and references therein]{Sidoli01} as well as the ultimate massive black hole candidate \\citep{Eisen05,Ghez05}. This variety of energetic processes not only modifies the distribution, kinematics, density and temperature of the gas but also drives different types of chemistry in the molecular gas in the Central Molecular Zone \\citep[CMZ, ][]{Morris96}.  Comparing the chemical composition between Galactic regions with known dominant heating mechanisms can be used to trace the heating mechanisms in the central region of galaxies. Unbiased spectral line surveys of selected regions within our Galaxy allow to identify the different chemical signatures which trace the different dominant heating mechanisms affecting the nuclear interstellar medium (ISM) in galaxies with different type of activity and/or in different evolutionary states \\citep{Martin06b}.  One of the molecules whose abundance appears to be a particularly unambiguous tracer of shock heating is HNCO. This asymmetric rotor has been observed in high mass molecular hot cores with fractional abundances relative to H$_2$ of $\\sim10^{-9}$ \\citep{Zinchen00} as well as in translucent and dark clouds with similar abundances \\citep{Turner99}. The observations of HNCO toward the photodissociation regions in the Orion bar shows abundances of $\\lesssim$1.6$\\times$10$^{-11}$ \\citep{Jansen95}, indicating  this molecule is easily destroyed by the UV radiation. Within the Galactic center region, the low resolution C$^{18}$O (which is thought to be a tracer of the total H$_2$ column density) and HNCO maps by \\citet{Dahmen97} as well as the close up view towards the Sgr\\,A region by \\citet{Lindqvist95} clearly show differences in their distribution which can not be ascribed to only differences in excitation but to real differentiation in the chemical properties of the different molecular complexes in the Milky Way. \\cite{Minh06} have mapped the Sgr\\,B2 region and they identified an expanding ring-like structure in the HNCO emission, clearly different from that of $^{13}$CO and HCO$^+$. This suggests that HNCO is likely enhanced by the shock interaction between the expanding ring and the main cloud. Similar chemical differences in the distribution of the HNCO emission have also been observed at larger scales towards external galaxies \\citep{Meier05}. Therefore HNCO seems to be a potential good tracer of different heating mechanisms in the nuclei of galaxies. However, there is not a systematic study of the physical conditions and the relative abundances of HNCO in a sample of molecular clouds in the CMZ affected by different type of heating mechanisms.  In this paper we present systematic multitransition observations of HNCO toward 13 molecular clouds in the Galactic center which are thought to be affected by different processes such as shocks, UV radiation, and X-rays. We explore the potential of HNCO to be used as a chemical discriminator of the different physical processes affecting the chemistry of the ISM. We show that the HNCO abundance ratio with respect to C$^{18}$O, $^{13}$CS, and C$^{34}$S is an excellent discriminator between the molecular clouds affected by the UV radiation and moderate shocks. The proposed CS vs HNCO diagnostic diagram can be used to trace these type of activities in the extragalactic ISM. ",
        "conclusions": "We have studied the HNCO, $^{13}$CS, C$^{34}$S, and C$^{18}$O emission towards 13 positions in the Central Molecular Zone in the Galactic center region. Column densities and excitation temperatures have been derived for each position. These set of sources represent a sample of the main different chemical processes affecting the molecular material.  With variations of up to a factor of 20 in abundance, the molecule of HNCO is shown to be one of the most contrasted species among the observed sources. The derived HNCO/$^{13}$CS abundance ratio shows particularly high contrast of up to a factor of 30 between sources. This ratio allows the grouping of the observed sources in three categories: 1) Giant Molecular Clouds showing a high HNCO/$^{13}$CS ratio, 2) Photo dissociated regions and high velocity shocks interactions with the lowest observed ratios, and 3) hot molecular cores with intermediate HNCO/$^{13}$CS ratios but with significantly higher H$_2$ column densities.  The negative correlation between the CS and HNCO abundances relative to H$_2$ observed in these groups of sources support the idea of HNCO being ejected from grain mantles by low velocity shocks while CS formation is favored in PDRs environments via ion reactions.  We have used the results from the Galactic center sources as templates for the observed chemistry in extragalactic nuclei. Thus, we compared the observed abundances in the Galactic center with the available data on five starburst galaxies. A similarly large variation in the HNCO abundance is observed with sources such as M\\,82 and NGC\\,253 showing abundances similar to the galactic PDRs while, on the other end, Maffei\\,2 appears to be more similar to the Galactic center GMCs. The CS vs HNCO abundance diagram is suggested as a powerful tool to study the chemical differentiation in the nuclei of galaxies."
    },
    "0801/0801.4943_arXiv.txt": {
        "abstract": "% We highlight the importance of transit observations on understanding the physics of planetary atmospheres and interiors.  Transmission spectra and emission spectra of atmospheres allow us to characterize these exotic atmospheres, which possess TiO, VO, H$_2$O, CO, Na, and K, as principal absorbers.  We calculate mass-radius relations for water-rock-iron and gas giant planets and examine these relations in light of current and future transit observations.  A brief review is given of mechanisms that could lead to the large radii observed for some transiting planets. ",
        "introduction": "% The blanket term ``hot Jupiter\" or even the additional term ``very hot Jupiter'' belies the diversity of these highly irradiated planets.  Each planet likely has its own unique atmosphere, interior structure, and accretion history.  The relative amounts of refractory and volatile compounds in a planet will reflect the parent star abundances, nebula temperature, total disk mass, location of the planet's formation within the disk, duration of its formation, and its subsequent migration (if any).  This accretion history will give rise to differences in core masses, total heavy elements abundances, and atmospheric abundance ratios.  Through transit observations we are able to probe nearly \\emph{20 orders of magnitude} in pressure in these planets, as seen in \\mbox{Figure~\\ref{moth}}, from the outer reaches of the escaping atmosphere to the high pressures of dense central cores.  Since atmospheric and interior physics are both quite large topics, here we will focus specifically on the atmospheres of close-in giant planets at photospheric pressures, which \\emph{Spitzer} data is allowing us access to, as well as the structure and mass-radius relations of Earth-like to gas giant planets.  At least 21 gas giants are now known, as well as a possible Neptune-class planet, with even smaller planets on the horizon. \\begin{figure}[!ht] \\begin{center} \\includegraphics[scale=0.62]{f1.eps} \\end{center} \\caption{Exosphere \\cp[adapted from][]{Yelle04} , atmosphere \\citep{Fortney05}, and interior \\cp{Fortney07a} \\emph{P-T} profile for a generic HD 209458b-like planet.}\\label{moth} \\end{figure} \\begin{figure}[!ht] \\begin{center} \\includegraphics[scale=0.53]{f2.eps} \\end{center} \\caption{Flux incident upon a collection of hot Jupiter planets.  At left is incident flux as a function of planet mass, and at right as a function of planet surface gravity.  In both figures the labeled dotted lines indicate the distance from the Sun that a planet would have to be to intercept this same flux.  Diamonds indicate the transiting planets while triangles indicate non-transiting systems  (with minimum masses plotted but unknown surface gravities).  The error bars for HD 147506 and HD17156 indicate the variation in flux that each planet receives over their eccentric orbits.  Flux levels for pM Class and pL Class planets are shown, with the shaded region around $\\sim$0.04-0.05 AU indicating a possible transition region.  ``Hot Neptune'' GJ 436b experiences less intense insolation, off the bottom of this plot at 3.2$\\times$10$^7$ erg s$^{-1}$ cm$^{-2}$.}\\label{flux} \\end{figure} ",
        "conclusions": "The future will surely see additional discoveries in atmospheric and interior physics.  In particular, a large amount of \\emph{Spitzer} data on atmospheric thermal emission is just beginning to be published.  Thermal emission light curves, as a function of orbital phase, will take on increased importance.  Mass/radius determinations for less irradiated giant planets will test theories for planetary evolution.  Certainly the next frontier is smaller masses, now that GJ 436b in known to transit \\cp{Gillon07}.  Paced by discoveries from space missions (such as \\emph{COROT} and \\emph{Kepler}) and ground-based surveys, as well as precise followup characterization observations via \\emph{Spitzer}, \\emph{Hubble}, \\emph{MOST}, and the ground, the future is even more exciting than the present."
    },
    "0801/0801.4491_arXiv.txt": {
        "abstract": "{Evolutionary models for massive stars, accounting for rotational mixing effects, do not predict any core-processed material at the surface of B dwarfs with low rotational velocities. Contrary to theoretical expectations, we present a detailed and fully-homogeneous, NLTE abundance analysis of 20 early B-type dwarfs and (sub)giants that reveals the existence of a population of nitrogen-rich and boron-depleted, yet intrinsically slowly-rotating objects. The low-rotation rate of several of these stars is firmly established, either from the occurrence of phase-locked UV wind line-profile variations, which can be ascribed to rotational modulation, or from theoretical modelling in the pulsating variables. The observational data presently available suggest a higher incidence of chemical peculiarities in stars with a (weak) detected magnetic field. This opens the possibility that magnetic phenomena are important in altering the photospheric abundances of early B dwarfs, even for surface field strengths at the one hundred Gauss level. However, further spectropolarimetric observations are needed to assess the validity of this hypothesis.} ",
        "introduction": "\\label{sect_intro} Fast rotation is one of the most distinctive features of massive stars and may give rise to large-scale fluid motions in the stellar interior which dredge up a significant amount of CNO-processed material from the stellar core to the surface (Heger \\& Langer \\cite{heger_langer}; Meynet \\& Maeder \\cite{meynet_maeder00}). On the other hand, such mixing processes can transport the boron surviving close to the stellar surface, after the star has evolved off the zero age main sequence (ZAMS), to deeper layers where this fragile element will quickly be destroyed. Therefore, the B and CNO surface abundances of OB stars are powerful probes of rotation-related mixing phenomena throughout the entire stellar interior and at different evolutionary phases. As such, they can be used to constrain theoretical models including rotational effects, such as meridional currents or shear instabilities, and to ultimately calibrate some poorly-known quantities, such as the diffusion coefficients for the transport of the chemical elements. Of particular relevance for the present work is the recent recognition that dynamo-generated magnetic fields could also significantly alter the photospheric CNO content (Maeder \\& Meynet \\cite{maeder_meynet}; Heger \\etal \\cite{heger}).  The evolutionary models mentioned above do not predict any detectable nitrogen overabundance (i.e. above 0.1--0.2 dex) in B dwarfs with rotation rates below $\\sim$100 km s$^{-1}$. However, it has now been known for a long time (but often overlooked) that some apparently slowly-rotating objects actually present relatively large quantities of N-enriched material at their surface, with overabundances reaching up to a factor three (e.g. Kilian \\cite{kilian92}; Gies \\& Lambert \\cite{gies_lambert}). Such an extra amount of deep mixing in OB dwarfs would clearly have consequences on their post main-sequence evolution and other related, important issues (e.g. stellar populations in external galaxies).  We have recently revealed the same peculiarity in a number of $\\beta$ Cephei stars, which have been inferred from asteroseismic or line-profile variation studies to be {\\em intrinsically} very slow rotators. As an illustration, the B1.5--B2 subgiant \\dcet \\ has a true rotation rate of at most 28 km s$^{-1}$, as deduced from the detection of a rotationally-split multiplet in {\\em MOST} data (Aerts \\cite{aerts06}), yet exhibits an unexpected N excess reaching a factor about four (Morel \\etal \\cite{morel}; hereafter M06). This is puzzling when considering, for instance, that no contamination from CN-processed material has been reported in the literature for some fast-rotating analogs in low-metallicity environments where, in addition, the abundance changes should be more easily detectable (e.g. several B1.5 subgiants in the LMC with equatorial rotational velocities exceeding $\\sim$130 km s$^{-1}$; Korn \\etal \\cite{korn}). To shed more light on this issue, here we extend the M06 study and present an NLTE abundance analysis of a sample of 20 dwarfs/(sub)giants comprising stars not established as pulsating variables, with special emphasis on the possible role played by magnetic fields to shape the CNO surface abundance patterns of early B-type stars during the early phases of their evolution. ",
        "conclusions": "\\label{sect_conclusion} \\subsection{A population of slowly-rotating, N-rich Galactic B dwarfs}\\label{sect_population} Our NLTE abundance analysis confirms the results of previous studies (e.g. Kilian \\cite{kilian92}; Gies \\& Lambert \\cite{gies_lambert}) and supports the existence of a population of B stars which, despite having low rotation rates, displays the observational signature of deep mixing during the early phases of their evolution. Of the 36 (non-supergiant) B stars observed by Gies \\& Lambert (\\cite{gies_lambert}), about one third exhibit a [N/C] ratio above about --0.4 dex. Such a high incidence of N enriched stars in a primarily magnitude limited (with an additional constraint on the projected rotational velocity) sample suggests that such objects make up a significant fraction of all B dwarfs in the solar neighbourhood. A sizeable, similar population has also been shown to exist in the Magellanic Clouds (Hunter \\etal \\cite{hunter07a},b; Trundle \\etal \\cite{trundle07}).  Mass transfer processes in close interacting binaries may also dramatically alter the CNO surface abundances of the two components (e.g. Vanbeveren \\cite{vanbeveren}) and may be responsible for the existence of the so-called OBN stars whose unusual chemical properties are believed to be connected to binarity (e.g. Levato \\etal \\cite{levato}). However, one expects in that case a much higher N overabundance and C depletion than observed in our sample (Sch\\\"onberner \\etal \\cite{schonberner}). A good example of such a star with a more extreme CNO abundance pattern is the blue straggler \\object{$\\theta$ Car} (Hubrig et al., in prep.). On the other hand, \\bcep \\  is the only N-rich star known to be member of a binary system. The orbit is very wide (${\\cal P}_{\\rm orb}$$\\sim$90 yr; Pigulski \\& Boratyn \\cite{pigulski_boratyn}) and the existence of a past episode of mass transfer is debatable. In the case of \\piori, the strong boron depletion but roughly normal nitrogen abundance have been claimed to be the unambiguous signature of rotational mixing (Fliegner \\etal \\cite{fliegner}). Furthermore, three other single-lined binary systems (\\gpeg, \\ \\bcru \\ and \\toph) have normal abundances. Summarizing, mass exchange processes do not appear as a viable explanation for all the N overabundances observed.  With the caveat that not all the N-rich stars in our sample have been the subject of intensive monitoring campaigns, one can note that at least two stars in this subsample (\\gori \\ and \\ecma) are classified as poor or rejected $\\beta$ Cephei candidates in the recent catalogue of Stankov \\& Handler (\\cite{stankov_handler}). As can be seen in Fig.\\ref{fig_hr}, two N-rich stars (\\icma \\ and \\tsco) also fall well outside the instability domains for pressure and gravity modes.\\footnote{A blue loop evolution can be safely ruled out for the evolved star \\icma \\ according to evolutionary models (both with and without rotation included). The observed abundance anomalies are indeed much less severe than expected if the star had gone through the first dredge-up (e.g. Venn \\cite{venn99}).} On the other hand, some well-known pulsators are N-normal (e.g. \\gpeg, \\bcma). This argues against a purely pulsational origin for the nitrogen enrichment.  \\begin{figure*} \\centering \\includegraphics[width=16cm]{ms8999_fig4.ps} \\caption{Position of the programme stars in the Hertzsprung-Russell diagram: ({\\em a}) \\gpeg, ({\\em b}) \\zcas, ({\\em c}) \\dcet, ({\\em d}) \\neri, ({\\em e}) \\piori, ({\\em f}) \\gori, ({\\em g}) \\object{HD 36591}, ({\\em h}) \\bcma, ({\\em i}) \\xcma, ({\\em j}) \\15cma, ({\\em k}) \\icma, ({\\em l}) \\ecma, ({\\em m}) \\object{HD 85953}, ({\\em n}) \\bcru, ({\\em o}) \\vcen, ({\\em p}) \\tsco, ({\\em q}) \\toph, ({\\em r}) \\voph, ({\\em s}) \\bcep \\ and ({\\em t}) \\twelvelac. {\\em Open circles}: Group I (B- and N-normal stars), {\\em open squares}: Group II (N-normal, but B-depleted stars), {\\em filled circles}: Group III (B-depleted and N-rich stars). See Fig.\\ref{fig_boron} for a definition of the three subgroups. In the online version of this journal, green and magenta symbols refer to stars with and without an N excess, respectively. The stars lacking boron data (see right-hand side of Fig.\\ref{fig_boron}) are included in this figure assuming that the 3 N-normal and 2 N-rich stars belong to Groups I and III, respectively. The size of the symbols is linearly proportional to the $\\Omega R \\sin i$ values (see Table~\\ref{tab_parameters}). The luminosities were computed using {\\em Hipparcos} parallaxes and the bolometric corrections of Flower (\\cite{flower}). The extinction in the $V$ band, $A_V$, was derived from the theoretical ($B-V$) colour indices of Bessell \\etal (\\cite{bessell}). The evolutionary tracks for solar metallicity and without rotation are taken from Schaller \\etal (\\cite{schaller}). The initial masses are indicated on the right-hand side of this figure and the ZAMS is shown as a dotted line. The theoretical instability strips for $\\beta$ Cephei and SPB stars are shown as long dashed and dotted-dashed lines, respectively (Miglio \\etal \\cite{miglio}). The models have been computed assuming a metallicity $Z$=0.01 (value typical of the $\\beta$ Cephei stars in our sample; see M06), updated OP opacities, the solar mixture of Asplund \\etal (\\cite{asplund}) and the mean Ne abundance of Cunha \\etal (\\cite{cunha}).} \\label{fig_hr} \\end{figure*}  On the other hand, except perhaps in the case of \\icma, the stellar winds are too weak to remove the outermost stellar layers and expose boron-poor material. As an illustration, mass-loss rates $\\dot{M}$$\\sim$10$^{-7}$ M$_{\\sun}$ yr$^{-1}$ are required (e.g. Venn \\etal \\cite{venn02}), whereas two of the most evolved stars in our sample fall short by at least one order of magnitude: $\\dot{M}$$\\sim$6 and 9 $\\times$ 10$^{-9}$ M$_{\\sun}$ yr$^{-1}$ for \\bcma \\ and \\ecma, respectively (Cassinelli \\etal \\cite{cassinelli}; Drew \\etal \\cite{drew}). The same conclusion holds for the main-sequence star \\tsco, with $\\dot{M}$$\\sim$2 $\\times$ 10$^{-8}$ M$_{\\sun}$ yr$^{-1}$ (Donati \\etal \\cite{donati06b}).  It can be argued that some stars in our sample were rapid rotators on the ZAMS, hence experiencing strong rotational mixing and surface N enrichment, but that substantial loss of angular momentum ensued, spinning down the stars to the observed levels. This hypothesis is not supported by evolutionary models, which only predict a slow decline of the rotation rate during core-hydrogen burning (e.g. Meynet \\& Maeder \\cite{meynet_maeder03}), as confirmed by observations of early-B dwarfs in young open clusters (e.g. Huang \\& Gies \\cite{huang_gies}; Wolff \\etal \\cite{wolff}). However, the case of the magnetic stars is more controversial. For instance, a magnetically-confined wind could constitute an effective way of dissipating angular momentum (as in the case of the O star \\object{HD 191612}; Donati \\etal \\cite{donati06a}). The influence of a large-scale magnetic field with a dipole configuration on a radiatively-driven outflow can be parametrized by a dimensionless confinement parameter, \\begin{equation} \\eta=B^2_{\\rm pol}R^2_{\\star}/\\dot{M}v_{\\infty} \\end{equation} where $B_{\\rm pol}$ is the polar field strength, $R_{\\star}$ is the stellar radius, $\\dot{M}$ is the mass-loss rate and $v_{\\infty}$ is the terminal wind velocity (ud-Doula \\& Owocki \\cite{ud_doula}). The emergent flow is channeled along closed magnetic loops  and deflected towards the magnetic equatorial plane for $\\eta$ $\\gtrsim$ 1. Assuming an oblique rotator model for \\zcas, \\voph \\ and \\bcep \\ leads to a polar field of about 300 G (Neiner \\etal \\cite{neiner03a},b; Henrichs \\etal \\cite{henrichs}). The field geometry is more complex in \\tsco, but the average field emerging at the surface typically reaches similar values (Donati \\etal \\cite{donati06b}). If we adopt, as representative values, $B_{\\rm pol}$=300 G, $R_{\\star}$=5 R$_{\\odot}$, $\\dot{M}$=10$^{-8}$ M$_{\\sun}$ yr$^{-1}$ (see Sect.~\\ref{sect_population}) and $v_{\\infty}$=1500 km s$^{-1}$, we obtain $\\eta$$\\sim$30. It is thus likely that the stellar outflow would be magnetically confined within the Alfv\\'en radius, $R_A$, in our magnetic targets. Following Donati \\etal (\\cite{donati01}), the magnetic braking timescale, $\\tau$, can be roughly estimated through the following expression \\begin{equation} \\tau=k\\frac{M_{\\star}}{\\dot{m}}\\left(\\frac{R_{\\star}}{R_A}\\right)^2 \\end{equation} where $k$ is the fractional gyration radius, $M_{\\star}$ is the stellar mass and $\\dot{m}$ is the effective mass-loss rate (i.e. the amount of wind material not trapped within the magnetosphere). Characteristic values for $\\dot{m}$ and $R_A$ have been obtained in the case of \\bcep \\ and \\tsco \\ by Donati \\etal (\\cite{donati01}, \\cite{donati06b}) using the magnetically-confined wind-shock model of Babel \\& Montmerle (\\cite{babel_montmerle}). This led to magnetic braking timescales that are much longer than the estimated stellar ages (110 Myrs and 5 Gyrs, respectively). There is considerable uncertainty surrounding $\\dot{m}$ and the loss of angular momentum through magnetic phenomena is not fully understood (for instance, some chemically peculiar Bp stars undergo unexpectedly strong, sometimes abrupt, increases in their rotational period: Mikul\\'{a}\\v{s}ek \\etal \\cite{mikulasek}; Trigilio \\etal \\cite{trigilio}), but it appears unlikely that the slow rotation rate of the N-rich stars results from magnetic braking. Another argument against this interpretation is the fact that some non-magnetic stars in our sample are also intrinsically slow rotators (Table~\\ref{tab_parameters}). This suggests that another mechanism could also be at work, a possibility being that this property is inherited from the pre-main sequence phase (e.g. St\\c{e}pie\\'n \\cite{stepien}; Wolff \\etal \\cite{wolff}).  \\subsection{A link with magnetic fields?}\\label{sect_fields} Rotational mixing effects are expected to be more conspicuous at lower metallicities and for increasing stellar mass, age and rotation rate (e.g. Meynet \\& Maeder \\cite{meynet_maeder00}). There are no trends between the abundance data and $\\Omega R\\sin i$ (Fig.\\ref{fig_trends_CNO}) or $\\Omega R$, but this is somewhat expected considering the limited range in rotation rate sampled by our sample (and the noise introduced by the unknown inclination angles in the case of the projected rotational velocities). More telling is the lack of evidence for the stars with an N excess to be on average older and/or more massive than the stars with normal N abundances (Fig.\\ref{fig_hr}). As can also be seen in Fig.\\ref{fig_boron}, all the N-rich stars reach very similar (``saturation''?) levels of N enhancement ($\\sim$0.5 dex) despite having different physical properties. The inadequacy of the theoretical models to reproduce our observations is well-illustrated by the unexpected detection of a factor of 3 nitrogen excess in \\tsco \\ and \\voph. First, these stars are very slow rotators (especially \\tsco \\ with $\\Omega R$$\\sim$6 km s$^{-1}$; Donati \\etal \\cite{donati06b}) and rotational mixing should be particularly inefficient in that case. This problem could be lessened if one assumes a much steeper internal rotation law than currently assumed, leading to an increased transport of the chemical elements through shear instabilities. A decline of the rotational velocity by a factor 3--5 from the core to the surface has been inferred from asteroseismological studies in  \\neri \\ (Pamyatnykh \\etal \\cite{pamyatnykh}) and \\vcen \\ (Dupret \\etal \\cite{dupret}). Although this information is unfortunately lacking for the N-rich stars, this appears in accordance with the predictions of evolutionary models (Meynet \\& Maeder \\cite{meynet_maeder00}). Second, these stars are also probably very young (see Fig.\\ref{fig_hr}) and should not have had time to build up a significant N overabundance at the surface. Taken together, this suggests that an additional physical process to rotational mixing might play a key role in the appearance of N-enriched material at the stellar surface.  In this respect, it is remarkable that all observational data collected to date point to a higher incidence of a magnetic field in stars with peculiar abundances: all observations of the N-rich stars led to a magnetic field detection, whereas all stars with solar abundances remained undetected.\\footnote{We do not discuss here the case of the N-rich, apparently slowly-rotating $\\beta$ Cephei star \\object{HD 180642} (Morel \\& Aerts \\cite{morel_aerts}), as the magnetic field detection is not completely secure and needs to be confirmed: although several Zeeman features are seen in Stokes {\\em V} spectra obtained with FORS1/VLT, none of the two spectropolarimetric measurements formally exceeds the 3$\\sigma$ detection threshold (Hubrig et al., in prep.).} The only two exceptions to this rule are \\gpeg \\ and \\object{HD 85953} (Fig.\\ref{fig_boron}). The $\\beta$ Cephei star \\gpeg \\ has been claimed to possess a very weak magnetic field with a longitudinal component varying from --10 to 30 G (Butkovskaya \\& Plachinda \\cite{butkovskaya}). On the other hand, a magnetic field has been detected in the Slowly Pulsating B (SPB) star \\object{HD 85953} at the 3-$\\sigma$ level in 3 out of 5 measurements spread over 3.5 yrs with FORS1/VLT (Hubrig \\etal \\cite{hubrig06}; Hubrig et al., in prep.). The latter example illustrates the paucity of the magnetic fields measurements in B-type stars and that caution should be exercised when using these data. The field detections reported in Table~\\ref{tab_magnetic_fields} are for the most part at the limit of the instrumental capabilities (especially with the older generation of spectropolarimeters such as MuSiCoS) and are based on snapshot observations with a poor temporal sampling. Only in the cases of \\zcas, \\tsco, \\voph \\ and \\bcep \\ have sufficient time-resolved measurements been collected to provide detailed information on the field strength and global morphology (Neiner \\etal \\cite{neiner03a},b; Donati \\etal \\cite{donati06b}; Henrichs \\etal \\cite{henrichs}). Considering that the null detections are also often based on a handful of measurements and might be refuted when more intensive and precise observations are undertaken, the separation between magnetic and non-magnetic stars should only be regarded as arbitrary at this stage. It is also conceivable that the N enrichment took place when the star was very young and going through a magnetic phase (similarly to \\object{$\\theta^1$ Ori C}), and that the field strongly decayed afterwards (e.g. Hubrig \\etal \\cite{hubrig07}) to a point where it is no longer detectable with current instrumental facilities. In this context, it is interesting to note that an N excess is already present in the young magnetic stars \\tsco \\ and \\voph \\ (see Fig.\\ref{fig_hr}). Although highly speculative, this might explain why a slowly-rotating, N-rich dwarf such as \\object{1 Cas} ([N/C]=--0.07$\\pm$0.20 dex; Gies \\& Lambert \\cite{gies_lambert}) remained undetected during the intensive MuSiCoS observations of Schnerr (\\cite{schnerr07}). A correlation between the amount of CN-cycle processed material at the surface and the field strength would strongly support a connection between deep mixing and magnetic phenomena. There is no trend in our data between [N/C] and the rms longitudinal field $\\overline{\\left< B_l \\right>}$ (Table~\\ref{tab_magnetic_fields}), but it should be kept in mind that the rotational period is poorly sampled in most cases and that the magnetic observations are restricted to longitudinal field measurements, which are aspect dependent, and to fields larger than $\\sim$40--50 Gauss. Despite these limitations, however, forthcoming magnetic observations of OB stars conducted using the new generation of spectropolarimetric instruments (such as FORS1/VLT, ESPaDOnS/CFHT or Narval/TBL) will clearly prove essential to establish the existence of a link between a nitrogen enrichment and magnetic fields. Such a relationship could in principle help to select the most promising targets for magnetic field surveys of massive stars (indeed, an N excess was already one of the selection criteria used by Schnerr \\cite{schnerr07}).  Strong, ordered kG-strength magnetic fields have been detected in two slowly-rotating O-type stars, namely in the O6.5f?pe--O8fp star \\object{HD 191612} (Donati \\etal \\cite{donati06a}) and in the young, peculiar O7 star \\object{$\\theta^1$ Ori C} (Wade \\etal \\cite{wade}). To our knowledge, however, no quantitative CNO abundance analysis has been performed for either of these two objects. Only in the former has a nitrogen enhancement been suggested based on the observed line strengths (Walborn \\etal \\cite{walborn}). Unfortunately, a similar lack of CNO abundance data also applies to classical, magnetic Bp stars lying in a similar temperature domain as our programme stars. Two strongly magnetic, He-rich dwarfs (\\object{HD 66522} and \\object{HD 96446}) have been analyzed under the assumption of LTE by Zboril \\& North (\\cite{zboril}). The CNO abundances do not clearly differ from those of the non-magnetic stars in their sample, but no firm conclusion can be drawn from such a limited number of objects. Systematic studies have been conducted for Ap stars, both for boron (e.g. Leckrone \\cite{leckrone}) and CNO (e.g. Roby \\& Lambert \\cite{roby}). However, these stars are much cooler than our targets and diffusion effects have long been known to play a key role in the dramatic surface abundance anomalies observed in this temperature range.  Nevertheless, radiatively-driven microscopic diffusion has also been proposed (Bourge \\etal \\cite{bourge07}) as a possible explanation for the appearance of core-processed material at the surface of slowly-rotating $\\beta$ Cephei variables (we recall that 5 N-rich targets are confirmed members of this class; see Table~\\ref{tab_parameters}). Support for the occurrence of diffusion comes from the fact that it may also account for the accumulation of iron in the driving zone and the excitation of the pulsation modes through the $\\kappa$ mechanism (e.g. Bourge \\etal \\cite{bourge06}). Magnetic fields were not taken into account in these models, but they are expected to amplify the transport of the ionized species along the field lines and to lead to a large-scale, inhomogeneous distribution of some elements at the stellar surface. However, the Mg, Al, Si, S and Fe abundances do not reveal any clear deviation from a scaled solar pattern (see Fig.\\ref{fig_trends_He_Fe}) and do not provide evidence for the existence  of diffusion effects capable of significantly altering the surface abundances. This is supported by a comparison of the abundance properties of the two subsamples of N-normal and N-rich stars, which reveals no significant differences for chemical elements other than nitrogen (Table~\\ref{tab_comparison}). Two stars apparently have a low helium content (\\object{HD 36591} and \\icma), but this is at odds with the expected/observed position of the He weak stars in the $T_{\\rm eff}$-$\\log g$ plane (Babel \\cite{babel}; Bychkov \\etal \\cite{bychkov}) and may be regarded with some suspicion. Slightly subsolar He abundances are in general derived within our sample (Fig.\\ref{fig_trends_He_Fe}). We suspect that the adopted microturbulent velocities derived from the analysis of the \\ion{O}{ii} features are not adequate for properly modelling the \\ion{He}{i} lines, and that lower values should ideally be employed instead (as also concluded for \\ion{N}{ii}; Sect.~\\ref{sect_sensitivity}).  \\begin{table} \\centering \\caption{Comparison between the chemical properties of the two subsamples of N-normal and N-rich stars.} \\label{tab_comparison} \\begin{tabular}{lcc} \\hline\\hline & N-normal sample & N-rich sample\\\\ & (10 stars) & (10 stars)\\\\ \\hline $$He/H              & 0.071$\\pm$0.007 & 0.087$\\pm$0.020\\\\ $\\log \\epsilon$(C)  & 8.21$\\pm$0.09   & 8.11$\\pm$0.10\\\\ $\\log \\epsilon$(N)  & 7.67$\\pm$0.11   & 7.98$\\pm$0.09\\\\ $\\log \\epsilon$(O)  & 8.51$\\pm$0.09   & 8.44$\\pm$0.16\\\\ $\\log \\epsilon$(Mg) & 7.39$\\pm$0.11   & 7.38$\\pm$0.12\\\\ $\\log \\epsilon$(Al) & 6.13$\\pm$0.08   & 6.10$\\pm$0.12\\\\ $\\log \\epsilon$(Si) & 7.18$\\pm$0.07   & 7.19$\\pm$0.10\\\\ $\\log \\epsilon$(S)  & 7.20$\\pm$0.11   & 7.15$\\pm$0.11\\\\ $\\log \\epsilon$(Fe) & 7.28$\\pm$0.10   & 7.26$\\pm$0.08\\\\ \\hline \\end{tabular} \\end{table}  Recent codes attempting to model rotational mixing in the presence of magnetic fields (Maeder \\& Meynet \\cite{maeder_meynet}; Heger \\etal \\cite{heger}) also face difficulties for accounting for the data presented in this paper. For instance, they predict a noticeable helium enrichment, which is not observed (the only exception is \\voph, but the slight He excess is more likely to be due to magnetic phenomena; see Neiner \\etal \\cite{neiner03b} and M06). Furthermore, these models assume a magnetic field created via dynamo action in the stellar envelope (Spruit \\cite{spruit}), whereas at least three N-rich stars in our sample (\\zcas, \\voph \\ and  \\bcep) have been claimed to possess a global dipole field most likely of fossil origin (Neiner \\etal \\cite{neiner03a},b; Henrichs \\etal \\cite{henrichs}). A remnant field from the star formation phase is also more likely in \\tsco, although the field topology is probably much more complex (Donati \\etal \\cite{donati06b}). In all cases, the periods found in the magnetic data are compatible with those modulating the UV wind line profiles, thus supporting an oblique rotator model. Clearly much progress remains to be done before the production of magnetic fields in massive stars and the mixing processes taking place in their interiors are satisfactorily understood (see, e.g. Zahn \\etal \\cite{zahn})."
    },
    "0801/0801.3564_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:1}  The exploration of the nature of the phase space that we live in, is a naturally attractive endeavour. On astronomical scales, this translates to an exercise in understanding the state of the phase space in the neighbourhood of the Sun. Of course, for a long time, this was impeded by the dearth of sufficiently large samples of proper motion data of the nearby stars. However, with the information available in Woolley's catalogue \\cite{woolley}, Agris Kalnajs concluded the local velocity space to be basically bimodal \\cite{agris} and attributed this behaviour to the rough proximity of the Sun to the Outer Lindblad Resonance of the central bar in the Galaxy ($OLR_b$). The observational domain expanded with the availability of the transverse velocity information of stars in the vicinity of the Sun, as measured by {\\it Hipparcos}. Since then many workers have attempted to chart the local phase space distribution $DF$ \\cite{jameswalter, fux_aa, skuljan} and investigate the origin of the structure in the same \\cite{fux_aa, walterolr, tremainewu, quillen, famaey05}.  Here we pursue the hypothesis that the observed state of the local phase space owes to the dynamical effect of Galactic features such as the central bar or the spiral pattern. However, dynamical modelling of the Solar neighbourhood is incomplete without the inclusion of the effects of {\\it both} the bar and the spiral pattern; after all, a major resonance of the bar occurs in the vicinity of the Sun, (as suggested by the Kalnajs Mechanism) and the inter-arm separation in the Galactic spiral pattern falls short of the average epicyclic excursion of a star at the Solar radius \\cite{chakrabarty_aa}.  However, this joint perturbation has been considered only occasionally \\cite{chakrabarty_aa, quillen}. That too, it is contentious if the Galactic spiral arms should attach themselves to the ends of the bar or rotate with a pattern speed ($\\Omega_s$) that is distinct from that of the bar ($\\Omega_b$).  The former scenario was invoked to produce self-consistent models of the galaxies NGC 3992, NGC 1073, and NGC 1398 \\cite{kaufmann96} while the best-fit model for NGC 3359 was reported to manifest distinct pattern speeds for the bar and the spiral \\cite{boon05}. Such dissimilar pattern speeds were considered for the first time by \\cite{chakrabarty_aa} in her modelling of the Solar neighbourhood. The effect of the relative separation between the locations of the resonances due to the Galactic bar and the spiral pattern is expected to bear interesting dynamical consequences.  In this article, we report the results of modelling of the local phase space, undertaken with the aim of disentangling the mystery about the origin of the observed phase space structure. In particular, we address the question of the imperativeness of the presence of chaos, in order to explain this structure. ",
        "conclusions": "\\noindent In this contribution, we have attempted to understand the state of the local phase space via dynamical modelling that includes the influence of the bar and the spiral pattern in the Galaxy. The clumpy structure of the 2-D local velocity distribution was attributed to the major and minor resonances of these perturbations, as well as the emergence of chaoticity triggered primarily by the spiral potential. Only the models that take the gravitational potentials of both bar and spiral are found to offer viable representations of the Solar neighbourhood dynamics; overlap of the results from the successful models indicate a fast bar, with a pattern speed of 57.4$^{+3.0}_{-2.5}$\\kmskpc and a bar angle that lies in the range [0$^\\circ$, 30$^\\circ$].  The currently used $p$-value formalism is not satisfactory on account of the lack of automativeness and inability to rank or grade the models. In fact, we are hoping to improve upon this technique, with a Bayesian measure of the quality of a simulation. This is to aid the expansion of the relevant parameter space, particularly, in the inclusion of the effects of the halo and progression to a full six dimensional phase space."
    },
    "0801/0801.3752_arXiv.txt": {
        "abstract": "{ We present simultaneous high-temporal and high-spectral resolution observations at optical and soft X-ray wavelengths of the nearby flare star  CN Leo.  During our observing campaign a major flare occurred, raising the star's instantaneous energy output by almost three orders of magnitude.  The flare shows the often observed {\\it impulsive behavior}, with a rapid rise and slow decay in the optical and a broad soft X-ray maximum about 200 seconds after the optical flare peak. However, in addition to this usually encountered flare phenomenology we find an extremely short ($\\tau _{dec}$ $\\approx$ 2 sec) soft X-ray peak, which is very likely of thermal, rather than non-thermal nature and temporally coincides with the optical flare peak.  While at hard X-ray energies non-thermal bursts are routinely observed on the Sun at flare onset, thermal soft X-ray bursts on time scales of seconds have never been observed in a solar nor stellar context.   Time-dependent, one-dimensional hydrodynamic modeling of this event requires an extremely short energy deposition time scale $\\tau _{dep}$ of a few seconds to reconcile theory with observations, thus suggesting that we are witnessing the results of a coronal explosion on CN~Leo. Thus the flare on CN Leo provides the opportunity to observationally study the physics of the long-sought \"micro-flares\" thought to be responsible for coronal heating. } ",
        "introduction": "The basic energy release in solar and stellar flares is thought to occur in the coronal regions of the underlying star, by conversion of non-potential magnetic energy through magnetic reconnection \\citep{priest02}. Particle acceleration takes place and the accelerated particles and/or plasma waves move along the magnetic field lines and reach and heat cooler atmospheric layers \\citep{syrat72,brown73}.  Because of their small cooling times the heated photospheric layers start radiating immediately, serving as a proxy indicator for non-thermal particles, while the heated dense chromospheric layers \"evaporate\" leading to a soft X-ray flare \\citep{brown73,neupert68,peres00}. The heating and cooling time scales and the dynamic response of the various atmospheric layers are important for our understanding of solar and stellar flares and of coronal heating \\citep{kostiuk74,somov81,tamres86,fisher87}; a popular hypothesis \\citep{klimchuk06,parker88} even attributes all of the required energy input to the quiescent solar corona to (nano)-flare input.  Solar and stellar flares are observed to occur over vastly different time scales ranging from a few seconds \\citep{vilmer94,schmitt93} to more than a week \\citep{kuerster96}, and the response of the heated plasma sensitively depends on the mode of energy input. Since the flare process involves plasma differing in temperature and density by more than four orders of magnitude, varying on short spatial and temporal scales, simultaneous multiwavelength data with sufficient spectral coverage and sufficient spectral and temporal resolution are required for any sensible observational diagnostics.  Since the primary energy release and energy loss process of a flare is coronal, space-based X-ray observations are required. Using ESA's XMM-Newton X-ray observatory and ESO's Ultraviolet-Visual Echelle Spectrograph (UVES) on Kueyen on May 19$^{th}$ 2006  we observed a spectacular flare on the nearby (d~=~2.39~pc) flare star CN~Leo (spectral type: M5.5-M6.0, $T_{eff}$ = 2800-K~2900~K).  Despite its low apparent rotational velocity of $v\\ \\sin i < 3.0$~km~s$^{-1}$ \\citep{fuhr04}, CN~Leo shows all the attributes of a magnetically active star, including chromospheric and coronal emission as well as flaring in the optical and in the X-ray bands \\citep{fuhr07}. ",
        "conclusions": "Heating by so-called \"microflares\" and \"nanoflares\" \\citep{cargill04,klimchuk06} is a popular hypothesis to explain the heating of the corona of the Sun and that of other stars.  Their nanoflare model assumes that any coronal loop is composed of many unresolved strands of magnetic flux and that these strands  are \"heated impulsively by a small burst of energy (which, for convenience, are called a nanoflare, although the range of energies is completely arbitrary\". Further, the strands are thermally isolated from each other and can be viewed as elemental magnetic flux tubes with a small value for the plasma beta \\citep{cargill04}. An individual heating event on such a strand leads to a dynamic evolution of the affected plasma located on this particular strand.  In a typical observation of a star and even in a typical observation of the solar corona, a multitude of heating events occur more or less simultaneously on different strands located in the field of view or are integrated over during a typical exposure time. Therefore in this picture any solar or stellar coronal observation represents a mean of the heating and cooling history of possibly a very large number of individual energy releases.  For a computation of the global properties of solar and stellar coronae in terms of the distributions of plasma densities, temperatures and flow velocities, a knowledge and modeling of the relevant heating and cooling processes is mandatory. Such models specifically assume an impulsive heating of the strands of magnetic flux, which can be recognized only in data with very high time resolution ($\\sim$ 1 sec), and then cool, first, by conduction and then by radiation \\citep{cargill04}.  This is precisely what appears to have happened on CN~Leo:  An energy release on a time scale of less than 10 sec in a single loop has to our knowledge never been observed in any star nor on the Sun. Therefore, with our high time-resolution soft X-ray and optical observations of the impulsive flare on CN~Leo -- by its energetics definitely not a nanoflare -- we may have isolated the relevant physics for nanoflare heating of the corona of the Sun and that of the stars."
    },
    "0801/0801.4957_arXiv.txt": {
        "abstract": "s{ The Alpha Magnetic Spectrometer (AMS) is a particle detector, designed to search for cosmic antimatter and dark matter and to study the elemental and isotopic composition of primary cosmic rays, that will be installed on the International Space Station (ISS) in 2008 to operate for at least three years. The detector will be equipped with a ring imaging \\CK\\ detector (RICH) enabling measurements of particle electric charge and velocity with unprecedented accuracy. Physics prospects and test beam results are shortly presented.} \\vspace{-1.2cm} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.2019.txt": {
        "abstract": "The \\cesam code is a consistent set of programs and routines which perform calculations of 1D quasi-hydrostatic stellar evolution including microscopic diffusion of chemical species and diffusion of angular momentum. The solution of the quasi-static equilibrium is performed by a collocation method based on piecewise polynomials approximations projected on a B-spline basis; that allows stable and robust calculations, and the exact restitution of the solution, not only at grid points, even for the discontinuous variables. Other advantages are the monitoring by only one parameter of the accuracy and its improvement by super-convergence. An automatic mesh refinement has been designed for adjusting the localisations of grid points according to the changes of unknowns. For standard models, the evolution of the chemical composition is solved by stiffly stable schemes of orders up to four; in the convection zones mixing and evolution of chemical are simultaneous. The solution of the diffusion equation employs the Galerkin finite elements scheme; the mixing of chemicals is then performed by a strong turbulent diffusion. A precise restoration of the atmosphere is allowed for. ",
        "introduction": "Within the limitations due to electronic degeneracy, \\cesam allows the computation of the quasi-static evolution of stellar models as long as the assumption of quasi-static equilibrium remains valid, that is to say, until the exhaustion of oxygen in the core. The modular structure of \\cesam facilitates the choices among several physical formalisms, for equation of state (hereafter {\\small EOS}), convection, opacities, diffusion coefficients, \\etc Many nuclear networks and initial mixtures are available which allow to optimise the physical description according to the kind of model and evolutionary phase of interest. Mass loss and infall of planetoids are also implemented.\\\\  \\paragraph{The available packages.} Earlier versions, {\\sc cesam2-3-4}, were programmed in F77 and {\\sc cesam5} in F90. Though obsolete, {\\sc cesam4} and {\\sc cesam5} are available at:  \\centerline{\\tt\\small http://www.obs-nice.fr/morel/CESAM}  \\noindent In the early 2000's, \\cesam was re-programmed in F95 and named {\\sc cesam2k}. Three versions are now available: \\begin{itemize} \\item The ``{\\em fixed version}'' and the ``{\\em fixed COROT version}'', limited to 3$\\alpha$ burning, available respectively at:  \\centerline{\\tt\\small http://www.obs-nice.fr/cesam/}  \\centerline{\\tt\\small http://perso.obspm.fr/$\\sim$lebreton/Modeles/CESAM.html}  \\item The ``{\\em $\\beta$ version}''  more complete, not fixed, still in development (not free of bugs) is hereafter signalled by the flag $(\\beta)$. It includes evolution up to oxygen burning, diffusion of the angular momentum and other developments of minor importance. It is available at:  \\centerline{\\tt\\small http://www.obs-nice.fr/morel/CESAM} \\end{itemize} Each package contains five directories and two files: \\begin{itemize} \\item ``{\\tt SOURCE}'', contains the {\\sc fortran} sources. \\item ``{\\tt EXPLOIT}'', contains programs to exploit the models and examples of input files. \\item ``{\\tt SUN\\_STAR\\_DATA}'', contains physical data and programs for their implementation. \\item ``{\\tt TESTS}'', contains programs performing various checks. \\item ``{\\tt SCRIPTS}'', contains scripts to be used for the implantation and operating, and a {\\tt MAKEFILE}. \\item ``{\\tt aide\\_mem2k.ps}'', a short guide of directions for use, the ``{\\em aide-m\\'emoire}'' (hereafter Paper 2). \\item ``{\\tt cesam2k.ps}'', a complete description of the numerical aspects and physics implemented, the ``{\\em notice}'' (hereafter Paper 3). \\end{itemize} The source is structured in 14 {\\em modules}: numerical routines, opacities, convection, etc. All the F95 routines of the source are compiled once. The requirements for the calculations are read in external files to be supplied by the user. The run is interactive. Messages displayed in French, or in English, allow the control of calculations. The extent of convection zones, a H--R diagram, and the profiles of temperature, pressure, luminosity and abundances are displayed\\footnote{Use of the {\\tt PGPLOT} package.} on line. \\cesam has been especially designed to facilitate the implementation of various physical constants, opacities, {\\small EOS}, atmosphere, nuclear networks etc. So, its overall structure is separated in two spaces: \\begin{enumerate} \\item A {\\em``physical space''} where the coefficients of the differential equations are written in a form close to their physical formalism. \\item A {\\em ``numerical space''} where the differential equations are formally solved. \\end{enumerate} Therefore \\cesam allows to implement physical processes, and physical data, without any knowledge of numerical methods involved for the solution of the equations. For the physics, the use of generic routines makes the reading of the algorithms easier.  \\paragraph{Units and values of physical constants.} \\cesam uses {\\tt cgs} units except for the mass, radius and luminosity expressed in solar units. Two sets of fundamental constants are implemented which correspond to widely used values \\citep{c68,c88,li94,c20}. For each calculation, one of these sets is chosen as the unique source of fundamental constants. Other constants are initialised locally, \\eg the mass excesses, in the routine performing the calculations of thermonuclear reaction rates.  \\paragraph{Input files.} Very often only the {\\em ``input data file''} (hereafter IDF) is needed. It is read at the onset of the run and collects all the requirements needed for the calculations: \\begin{itemize} \\item physical parameters: mass, chemical composition, mi\\-xing-length parameter, etc. \\item numerical parameters: maximum number of shells, kind of precision, etc. \\item criteria for halting the computations: age to be reached, value of the hydrogen abundance at centre, etc. \\item names and locations of the external data files containing the data of the tabulated {\\small EOS} and opacity, names of the physical routines to be used, name of the model, of the set of units to be used, etc. \\end{itemize} The other input files have only specific functions, among them the most useful are: \\begin{itemize} \\item ``{\\tt mixture}'' allows the use of an initial mixture not implemented in \\cesamp \\item``{\\tt modif\\_mix}'' allows to modify the abundances of some species in a mixture already implemented. \\item ``{\\tt reglages}'' allows to personalise the kind of precision to be used (see below). \\item ``{\\tt planet}'' contains the characteristics of infall of planetoids. \\item ``{\\tt rap\\_iso}'' allows to modify the isotopic ratios. \\item ``{\\tt vent}'' defines the chemical composition of the wind when it differs from the atmosphere composition. \\item ``{\\tt zoom}'' allows to fit the resolution of the display for the plot on line. \\item ``{\\tt langue}'' allows to have the messages in English. \\end{itemize} The meaning of all items are explained in Paper~2. Examples of files are given in the directory {\\tt EXPLOIT}.  \\paragraph{Output files.} At the end of each time step, a {\\em ``return binary file''} (hereafter RBF) is created. It contains all the data needed to initialise or pursue a computation. There is the possibility, either to save all RBF, or to keep only the last RBF created. On request, {\\em ``output data files''} (hereafter ODF) are created. Three ODF are designed for adiabatic, non adiabatic and inversion asteroseismic investigations. These ODF also serve as input for some programs of the directory {\\tt EXPLOIT}. An ODF concerns the diffusion of the angular momentum ($\\beta$). All items of output  files are detailed in Paper~2. It is also possible to create personalised ODF.  \\paragraph{The kinds of precision.}\\label{sec:preci} To optimise the calculations, sets of parameters, named {\\em ``r\\'eglages''}, are fixed according to the kind of models to be calculated and to their subsequent use. The most useful {\\em r\\'eglages} are: \\begin{itemize} \\item ``{\\tt realistic precision}'' for standard evolutions. \\item ``{\\tt super precision}'' used when a high level accuracy is needed. \\item ``{\\tt solar accuracy}'', close to {\\tt super precision}, but especially designed for seismological investigations; the number of shells of the last model is increased up to its maximal value. \\item ``{\\tt corot}'', close to {\\tt super precision}, but especially designed for investigations connected to CoRot. \\item ``{\\tt advanced}'', for the computation of early type star models evolved up to the oxygen burning. \\item ``{\\tt normal precision}'', for exploratory work. \\item ``{\\tt low mass}'', for the computation of late type star models. \\item ``{\\tt reglages}'', in that case, the parameters, designed by the user, are read on an input file named ``{\\tt reglages}''. \\end{itemize} The larger the expected accuracy, the larger the computational expense.  \\paragraph{Operating data and programs.} The directory {\\tt EXPLOIT} contains: \\begin{itemize} \\item examples of input files as IDF, ``{\\tt reglages}'', ``{\\tt planet}'', ``{\\tt modif\\_mix}'', etc. \\item ASCII files of preliminary models for the initialisations of PMS and ZAMS models. \\item miscellaneous programs to make plots of the chemical composition profile, or an extension of the grid on a given set of radius, or create the IDF for solar calibration, etc. \\end{itemize} The flow chart of \\cesam is described in \\Sect{sec:flow}. As the numerical features are detailed in the appendix of \\citet{m97}\\footnote{Only available on electronic form.} (hereafter Paper I) they are only succinctly recalled in \\Sect{sec:num}, except for the automatic allocation of mesh points described in \\Sect{sec:mvgrid}. The restitution of the atmosphere is outlined in \\Sect{sec:atm}. The algorithms performing the temporal evolution are described in \\Sect{sec:evol}. The nuclear network is detailed in \\Sect{sec:chem} and the implementation of the rotation is described in \\Sect{sec:rot}. The various formalisms of convection implanted in \\cesam are described in \\Sect{sec:conv}. Mass loss formalisms and infall of planetoids are described in \\Sect{sec:loss}. {\\small EOS} and opacities data available are listed in \\Sect{sec:eos}. ",
        "conclusions": ""
    },
    "0801/0801.0603_arXiv.txt": {
        "abstract": "We show that, within some modified gravity theories, such as the Palatini models, the non-linear nature of the field equations implies that the usual na\\\"{i}ve averaging procedure (replacing the microscopic energy-momentum  by its cosmological average) could be invalid. As a consequence, the relative motion of particles in Palatini theories is actually indistinguishable from that predicted by General Relativity. Moreover, there is no WEP violation. Our new and most important result is that the cosmology and astrophysics, or put more generally, the behaviours on macroscopic scales, predicted by these two theories are the same, and as a result the naturalness problems associated with the cosmological constant are not alleviated. Palatini gravity does however predict alterations to the internal structure of particles and the particle physics laws, \\emph{e.g.}, corrections to the hydrogen energy levels. Measurements of which place strong constraints on the properties of viable Palatini gravities. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.2740_arXiv.txt": {
        "abstract": "{We present first results of a systematic investigation of the incomplete Paschen-Back effect in magnetic Ap stars. A short overview of the theory is followed by a demonstration of how level splittings and component strengths change with magnetic field strength for some lines of special astrophysical interest. Requirements are set out for a code which allows the calculation of full Stokes spectra in the Paschen-Back regime and the behaviour of Stokes $I$ and $V$ profiles of transitions in the multiplet 74 of Fe\\,{\\sc ii} is discussed in some detail. It is shown that the incomplete Paschen-Back effect can lead to noticeable line shifts which strongly depend on total multiplet strength, magnetic field strength and field direction. Ghost components (which violate the normal selection rule on $J$) show up in strong magnetic fields but are probably unobservable. Finally it is shown that measurements of the integrated magnetic field modulus $H_s$ are not adversely affected by the Paschen-Back effect, and that there is a potential problem in (magnetic) Doppler mapping if lines in the Paschen-Back regime are treated in the Zeeman approximation.} ",
        "introduction": "\\label{sec:intro}  Ever since Pieter Zeeman (1897) discovered the splitting of spectral lines in a magnetic field, astrophysicists have tried to take advantage of the diagnostic capabilities of this effect. Thanks to the Zeeman effect, G.E. Hale (1908) was able to demonstrate the presence of strong magnetic fields in sunspots and over the years countless Zeeman observations (of ever increasing precision) of lines originating in all parts of the solar atmosphere have provided deep insights into the physics of the outer layers of the sun. The discovery by Friedrich Paschen and Ernst Back (1921) that in strong fields some transitions change from the anomalous Zeeman effect to the normal Zeeman effect, has not resulted in any revision of the traditional interpretation of the observations, because even in the strongest solar magnetic fields (about 0.4\\,T) only very few lines are noticeably affected. Not until almost 50 years later was the question of the Paschen-Back (PB) effect in the solar spectrum first addressed by Engvold et al. (1970). Similarly, after the discovery by H.W. Babcock, starting 1947, of very strong magnetic fields in some upper-main sequence chemically peculiar stars (Ap stars), no serious thought was given to the Paschen-Back effect. Although it is clear that in a 3.4\\,T field like that of HD\\,215441 (Babcock 1960), many spectral lines would exhibit the transition from the Zeeman to the PB regime (which we shall call the incomplete or partial PB effect), the first papers on the Paschen-Back effect in a stellar context appeared only much later (Kemic 1975; Stift 1977).  Recently, interest in the Paschen-Back effect has revived, mainly in the solar context and for molecular lines. As to atomic transitions, let us note that the He\\,{\\sc i}\\,10830\\,{\\AA} triplet has received special attention (Socas Navarro et al. 2004), as has for example the Mn\\,{\\sc i}\\,15262.7\\,{\\AA} line (Asensio Ramos et al. 2007). In the stellar context, the work of Mathys (1990) constitutes a particularly detailed modelling attempt and discussion of the incomplete PB effect in the Fe\\,{\\sc ii} lines at $\\lambda$\\,6147.7 and $\\lambda$\\,6149.2. The PB effect on hyperfine structure and its observational consequences have been investigated by Landolfi et al. (2001). It should be kept in mind that these results have generally been obtained in the Milne-Eddington approximation, and that heavy blending as often found in Ap stars could not be taken into account.  Our interest in the Paschen-Back effect has been stimulated by the difficulties we encountered when trying to match high resolution spectral observations of strongly magnetic Ap stars with synthetic spectra (see Fig.\\,\\ref{fig:shift}). The Zeeman doublet of Fe\\,{\\sc ii} at 6149\\,{\\AA} in particular not only displays non-symmetric relative intensities of the components -- already noted by Mathys (1990) -- but it also proves impossible to model the velocity shift (relative to the reference RV determined from magnetic null lines of iron) within the Zeeman regime. From Fig.\\,\\ref{fig:shift} it emerges that even at the very moderate field strength of 0.36\\,T this shift is already observable.  Assuming that the Paschen-Back effect is responsible for both asymmetries and shifts, the question naturally arises whether the measurement of the magnetic field modulus $H_s$ (the absolute field strength integrated over the visible hemisphere), which is based on the interpretation of the observed splitting of the $\\lambda\\,6149$ line in terms of simple Zeeman splitting, may be systematically affected by the PB effect. This has not been studied so far, and we are also not aware of any discussion in the literature of the observed shift of the $\\lambda\\,6149$ line.  In the past, neither computers nor codes were powerful enough to embark on any ambitious program of modelling stellar Stokes profiles in the incomplete PB regime. Now that multi-processor and/or multi-core architectures have become affordable, COSSAM (Stift 2000) -- a state-of-the-art  object-oriented and fully parallel Stokes code -- has been modified in such a way as to allow the calculation for realistic stellar atmospheres of a multiplet in the incomplete PB regime, with full blending from the remaining spectral lines which are assumed to display classical anomalous Zeeman patterns. These tools make it possible to systematically explore how the observed Stokes profiles are affected by the partial PB effect, to have a close look at the diagnostic content of Paschen-Back lines, and to model in detail selected spectral intervals in Ap stars with very strong fields. First important results are presented in this paper.   \\begin{figure} \\includegraphics[width=84mm]{HD66318_l6149.eps} \\includegraphics[width=84mm]{HD188041_l6149.eps} \\caption{ RV-corrected Stokes\\,$I$ spectra in the vicinity of the Fe\\,{\\sc ii} lines at 6147 and 6149\\,{\\AA} of the stars HD\\,66318 (top) and HD\\,188041 (bottom). The magnetic field modulus is 1.5\\,T for the former star, 0.36\\,T for the latter. The dashed red lines give the zero field positions of the lines of Cr, Fe and Nd. The full red lines are plotted for the only purpose of illustrating the line shift in $\\lambda\\,6149$; they result from a straightforward spectral synthesis, assuming normal Zeeman splitting. The RV has been determined from the magnetic null lines ($g_{\\rm eff} = 0.0$) of Fe\\,{\\sc i} at $\\lambda\\,5434.52$ and $\\lambda\\,5576.09$.} \\label{fig:shift} \\end{figure} ",
        "conclusions": "\\label{seq:conclusion}  High resolution spectral observations of magnetic Ap stars make it clear that the Paschen-Back effect cannot be neglected in the interpretation and modelling of a number of spectral lines, whether it be the Li\\,{\\sc i} $\\lambda\\,6707$ resonance line or the Fe\\,{\\sc ii}\\,$\\lambda\\,6149$ Zeeman doublet. Because of its simple structure, the latter is particularly useful for the determination of the magnetic field modulus $H_s$. We have calculated in detail the splittings of the spectroscopic terms involved in these transitions and the relative intensities of the subcomponents and found, somewhat to our surprise, that some lines can enter the incomplete Paschen-Back regime already at field strengths of a bare 0.05\\,T as is the case for the lithium resonance lines. Looking further to other chemical elements, we have discovered that in extreme cases (Cr\\,{\\sc i} is an example), up to an estimated 18\\% of the lines may exhibit the transition to the partial Paschen-Back regime in Ap stars with strong magnetic fields of the order of 2\\,T.  The development of CossamPaschen, a modified version of COSSAM -- the polarised line synthesis code established by Stift (2000) -- has made it possible over the last few years to explore the effect of the partial Paschen-Back regime on Stokes spectra. CossamPaschen allows the detailed and realistic modelling of a PB multiplet, blended with lines which are treated in the Zeeman approximation. First ever detailed Stokes profiles in realistic stellar atmospheres have been presented at the CP\\#AP Workshop in Vienna (Stift 2007) and provide valuable insight into the sometimes spectacular changes of selected spectral lines. These calculations show that the centres-of-gravity of the Fe\\,{\\sc ii}\\,$\\lambda\\,6149$ line and of its sister line at $\\lambda\\,6147$ are shifted in magnetic fields, either towards the red or towards the blue, depending on line strength, magnetic field strength, and field direction. A comparison between the observations plotted in Fig.\\,\\ref{fig:shift} and the theoretical results displayed in Figs.\\,\\ref{fig:6238_I} and \\,\\ref{fig:6238_V} shows gratifying agreement.  Another interesting finding is the confirmation of the existence of normally forbidden ``ghost components'' in strong magnetic fields. The intensity of these components is a very strong function of magnetic field strength, and in Ap stars some of them may actually significantly contribute to the {\\em local} spectra in places with very strong magnetic fields (of the order of 4\\,T). However, the strong dependency of intensity and position on field strength will probably lead to the {\\em global}  signature of the ``ghost components'' begin washed out.  An important result of our investigation is the confirmation that $H_s$ measurements obtained from the splitting of the Fe\\,{\\sc ii}\\,$\\lambda\\,6149$ line and interpreted in terms of classical Zeeman splitting do not need to be corrected for the PB effect. For all practical purposes, even in the strongest fields encountered in Ap stars, the PB splitting is not different from the Zeeman splitting. As for maps derived by means of (magnetic) Doppler mapping, we cannot provide such a reassuring confirmation and there is a real danger that the inclusion of lines subject to the partial PB regime may lead to spurious results. What can be a serious nuisance may also provide improved diagnostic capabilities and so we are looking forward to improved imaging codes incorporating the full treatment of the partial Paschen-Back effect."
    },
    "0801/0801.3485_arXiv.txt": {
        "abstract": "We measure the two-point spatial correlation function for clusters selected from the photometric MaxBCG galaxy cluster catalog for the Sloan Digital Sky Survey (SDSS). We evaluate the correlation function for several cluster samples using different cuts in cluster richness. Fitting the results to power-laws, $\\xi_{cc}(r) = (r/R_0)^{-\\gamma}$, the estimated correlation length $R_0$ as a function of richness is broadly consistent with previous cluster observations and with expectations from N-body simulations. We study how the linear bias parameter scales with richness and compare our results to theoretical predictions. Since these measurements extend to very large scales, we also compare them to models that include the baryon acoustic oscillation feature and that account for the smoothing effects induced by errors in the cluster photometric redshift estimates. For the largest cluster sample, corresponding to a richness threshold of $\\Ng\\ge 10$, we find only weak evidence, of about $1.4-1.7\\sigma$ significance, for the baryonic acoustic oscillation signature in the cluster correlation function. ",
        "introduction": "Galaxy clusters have long been recognized as powerful cosmological probes \\citep{cluster_cosmo,RozoEtal2007cosmo}. In particular, measurement of the cluster mass function vs. redshift constrains cosmological parameters, including those associated with dark energy \\citep{cosmo_DEprobe_wang,cosmo_DEprobe_haiman,RCSconstraints,Rozo:2007yt,mantz07}. This has motivated the design of new wide-area cluster surveys in the optical \\citep{DES}, X-ray \\citep{ConX}, and using the Sunyaev-Zel'dovich effect (SZE) \\citep{SPT_1,SPT_2}. The utility of this probe hinges on limiting the uncertainty in the relation between cluster mass and whatever observable (e.g., optical richness, X-ray luminosity, or SZE flux decrement) is used as a proxy for it. Measurement of the two-point correlation function of clusters can help calibrate such mass-observable relations and thereby improve the resulting cosmological constraints \\citep{maj, lima}.  Since clusters are the largest virialized mass concentrations in the Universe, measurement of their spatial clustering also provides insight into models of large-scale structure formation and tests theoretical frameworks, such as the Halo Model, that describe the relation between the galaxy, cluster, and dark matter distributions, i.e., the bias \\citep{RDF,bahcall_sim,SSS2006,Schulz2006,Smith2007}.  On very large scales, $r \\sim 100\\Mpc$, the two-point correlation function or power spectrum of clusters should show evidence of baryon acoustic oscillations (BAO) \\citep{Angulo2005}. In concert with measurements of the cosmic microwave background anisotropy, the BAO scale provides an estimate of cosmic distance and thereby a geometric probe of dark energy \\citep{Seo2003,hu03}. A possible detection of BAO in the power spectrum of clusters from the Abell/ACO catalog was reported in  \\cite{mnb}, while earlier studies had claimed evidence for a feature in the cluster correlation function at $r \\sim 125\\Mpc$ \\citep{kopylov,mo,einasto1,einasto2,einasto3}.  The BAO feature was detected at $3.4\\sigma$ significance in the two-point correlation function of $\\sim 47,000$ luminous red galaxies (LRGs) with spectroscopic redshifts in the range $0.16<z<0.47$ in the Sloan Digital Sky Survey (SDSS) \\citep{Eisenstein2005}. Using slightly larger samples, the BAO feature was also detected in the LRG power spectrum \\citep{Huetsi2006,Tegmark2006,Percival2006}.  The BAO feature has also been inferred from the large-scale clustering of $\\sim 600,000$ LRGs identified in the deeper SDSS photometric survey \\citep{Blake2007,Padmanabhan2007}. Although the photometric catalog covers a larger volume and contains many more galaxies than the spectroscopic sample, the galaxy redshifts in the former must be estimated photometrically. With a photometric redshift uncertainty of $\\sigma_z \\sim 0.03$ for this sample, much of the information in the radial modes of the power spectrum is lost, and the BAO feature was detected at less than $3\\sigma$ significance. Based on theoretical considerations presented in \\cite{Seo2003} and \\cite{Blake2005}, \\cite{Blake2007} presented a direct measurement of the 3D power spectrum, with a correction for the damping of power in the radial direction due to photometric redshift uncertainties. By contrast, \\cite{Padmanabhan2007} measured the angular power spectrum in photometric redshift slices and used it to reconstruct the three-dimensional power spectrum.  In this paper, we present measurements of the two-point correlation function for optically selected galaxy clusters in the Sloan Digital Sky Survey (SDSS) and study the possible detection of baryonic oscillations. The cluster samples are derived from the MaxBCG catalog \\citep{maxBCGsample,Koester2007}, in which clusters are identified as concentrations of red-sequence galaxies; the colors of these early-type galaxies are used to estimate photometric redshifts for the clusters. The two point correlation function for galaxy clusters in SDSS was studied in \\cite{Basilakos} using an earlier cluster catalog.  We focus on measurement of the 3D cluster correlation function. In comparing the observations to models of large-scale structure, the model predictions are corrected for the effects of photometric redshift errors. We use two different methods for estimating this correction which are in good agreement for the photometric redshift error $\\sigma_z \\sim 0.01$ characteristic of the MaxBCG catalog.  The cluster correlation function measurements provide weak evidence ($\\sim 1.4-1.7\\sigma$) for the presence of BAO in the cluster spatial distribution. An independent measurement of the power spectrum for the same cluster sample has been presented in \\citep{Huetsi2007}, also indicating weak evidence ($\\sim 2\\sigma$) for acoustic features in the power spectrum.  The paper is organized as follows. In section \\ref{sec:maxbcg_cf} we describe the MaxBCG cluster catalog, the samples we derive from it, and our measurements of the cluster correlation function. In section \\ref{sec:model} we introduce our model for the cluster correlation function and discuss the corrections due to photometric redshift errors, presenting two different correction methods. In section \\ref{sec:covariance} two different estimates of the correlation function covariance matrix are described and compared. In section \\ref{sec:results} we compare the model to the data, extracting estimates of the cluster bias as a function of richness and mass, and comparing the goodness of fit for models with and without BAO. We present our conclusions in section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}    We have measured the large-scale 3D correlation function for four richness samples in the MaxBCG catalog of optically selected galaxy clusters from the SDSS, currently the largest cluster catalog available. Since the cluster redshifts in this sample were estimated photometrically, our modeling of the observed correlation function includes a careful treatment of the impact of photo-$z$ errors. The geometric approach to photo-$z$ errors we have introduced should be of broad utility in the analysis of future large photometric galaxy surveys.  On scales $r=20-60 \\Mpc$, the cluster correlation function is well fit by a power law in separation, with a correlation scale $R_0$ that increases with cluster richness. We have determined the relation between the correlation scale $R_0$ and the mean cluster separation $d$, finding qualitative agreement with the compilation of previous cluster measurements presented in \\cite{bahcall_data}. The scaling of $R_0$ with $d$ is also consistent with that predicted in N-body simulations of $\\Lambda$CDM \\citep{bahcall_sim}, but with a slightly ($10-15$\\%) higher value of $R_0$ at fixed $d$.  We have modeled the large-scale correlation function on scales $r=20-195 \\Mpc$ using a non-linear model of the $\\Lambda$CDM power spectrum that includes the effects of non-linear damping of the baryon acoustic peak. Non-linear damping, coupled with the estimated photo-$z$ errors, imply that we do not expect a robust detection of the BAO feature in these samples. Indeed, we find that the data set does not yield a clear detection of baryon acoustic features in the correlation function: for the largest sample considered, $\\Ng\\ge 10$, the significance for the best-fit BAO model with respect to a featureless model is about $1.4-1.7\\sigma$, depending on whether the covariance matrix is determined using the jackknife procedure or linear perturbation theory.  Comparison of the clustering on scales less than and greater than $r \\sim 60 \\Mpc$ provides weak evidence that the clustering amplitude for the $\\Ng\\ge 10$ and 11 samples is larger on large scales than on small scales, relative to a $\\Lambda$CDM power spectrum. Such a suggestion of extra large-scale power has also been seen in the clustering of the SDSS photometric LRG sample \\citep{Blake2007,Padmanabhan2007}, but statistical confirmation will require samples covering larger volumes.  Finally, we have combined the clustering measurements with weak lensing calibration of the mass-richness relation \\citep{Johnston2007} to directly infer the bias as a function of halo mass, a fundamental quantity in studies of structure formation. Again, the trend of increasing bias with halo mass is qualitatively consistent with the predictions of $\\Lambda$CDM, but the amplitude of the bias is $\\sim 15-20$\\% higher than the model prediction. This disagreement is reduced for lower values of the power spectrum amplitude $\\sigma_8$. Given the large intrinsic scatter in the relation between halo mass and cluster optical richness, however, as well as uncertainty in the photo-$z$ error dispersion $\\sigma_z$, it is not yet clear whether this is a significant discrepancy. Moreover, an elevated cluster bias could be a sign of assembly bias \\citep{wu}."
    },
    "0801/0801.4570_arXiv.txt": {
        "abstract": "We carry out a comprehensive analysis of the behavior of the magnetorotational instability (MRI) in viscous, resistive plasmas. We find exact, non-linear solutions of the non-ideal magnetohydrodynamic (MHD) equations describing the local dynamics of an incompressible, differentially rotating background threaded by a vertical magnetic field when disturbances with wavenumbers perpendicular to the shear are considered. We provide a geometrical description of these viscous, resistive MRI modes and show how their physical structure is modified as a function of the Reynolds and magnetic Reynolds numbers. We demonstrate that when finite dissipative effects are considered, velocity and magnetic field disturbances are no longer orthogonal (as it is the case in the ideal MHD limit) unless the magnetic Prandtl number is unity.  We generalize previous results found in the ideal limit and show that a series of key properties of the mean Reynolds and Maxwell stresses also hold for the viscous, resistive MRI. In particular, we show that the Reynolds stress is always positive and the Maxwell stress is always negative.  Therefore, even in the presence of viscosity and resistivity, the total mean angular momentum transport is always directed outwards.  We also find that, for any combination of the Reynolds and magnetic Reynolds numbers, magnetic disturbances dominate both the energetics and the transport of angular momentum and that the total mean energy density is an upper bound for the total mean stress responsible for angular momentum transport. The ratios between the Maxwell and Reynolds stresses and between magnetic and kinetic energy densities increase with decreasing Reynolds numbers for any magnetic Reynolds number; the lowest limit of both ratios is reached in the ideal MHD regime. ",
        "introduction": "\\label{sec:intro}  The magnetorotational instability \\citep[MRI,][]{BH91, BH98} has been widely studied in the inviscid and perfectly conducting, magnetohydrodynamic (MHD) limit.  The departures from this idealized situation are usually parametrized according to the Reynolds ${\\rm Re} = vl/\\nu$ and magnetic Reynolds ${\\rm Rm} = vl/\\eta$ numbers, where $v$ and $l$ stand for the relevant characteristic velocity and lengthscale and $\\nu$ and $\\eta$ stand for the kinematic viscosity and resistivity. The ideal MHD regime is then formally identified with the limit ${\\rm Re}~,~{\\rm Rm}~\\rightarrow~\\infty$. There are many situations of interest in which the effects of dissipation need to be considered.  From the astrophysical point of view, accretion disks around young stellar objects constitute one of the most compelling reasons for investigating the MRI beyond the ideal limit. In particular, there is great interest in understanding to what extent can MHD turbulence driven by the MRI enable efficient angular momentum transport in cool, poorly conducting, protoplanetary disks \\citep[see, e.g.][]{BB94, Jin96, Gammie96, SM99, SW05}.  Most of the studies addressing the effects of dissipation in non-ideal MRI have usually focused in inviscid, resistive plasmas. However, accretion disks are characterized by a wide range of magnetic Prandtl numbers, with ${\\rm Pm}=\\nu/\\eta$ varying by several orders of magnitude across the entire disk \\citep[see, e.g.,][]{BH08}. In order to understand the behavior of the MRI under these conditions it is necessary to relax the assumption of an inviscid plasma.  A large fraction of the shearing box simulations addressing the non-linear regime of the MRI have been carried out in the ideal MHD limit, i.e., without including explicit dissipation in the codes \\citep[see, e.g.,][]{HGB95,Brandenburg95,Sanoetal04}.  However, even in the absence of explicit viscosity and resistivity, finite difference discretization leads to numerical diffusion/dissipation. Therefore, even in this type of simulations, it is necessary to understand the impact of these numerical artifacts that lead to departures from the ideal MHD regime and how similar they are when compared with physical (resolved) dissipation.  A handful of numerical studies with explicit resistivity but zero physical viscosity have been carried out in order to understand the effects of ohmic dissipation in the saturation of MRI-driven turbulence \\citep[see, e.g.,][]{SIM98, SI01, FSH00, Sanoetal04, TSD07}. In particular, \\citet{SS03} have shown that the saturation level of the stresses increases with increasing magnetic Reynolds number and seem to converge to an asymptotic value for magnetic Reynolds numbers larger than unity.  Recent work has pointed out problems with convergence in zero-net-flux numerical simulations of ideal MHD driven by the MRI \\citep{PCP07, FPI07} implying the necessity of incorporating explicit dissipation in the codes.  Numerical studies with both resistivity and viscosity, in the presence of a mean vertical magnetic field \\citep{LL07} and in the case of zero net flux \\citep{FPII07}, have begun to uncover how the characteristics of fully developed MRI-driven turbulence depends on the Reynolds and magnetic Reynolds numbers.  Even though the ranges in Reynolds and magnetic Reynolds numbers that can be currently addressed is still limited, the results obtained from the simulations suggest that the magnetic and kinetic energies contained in turbulent motions in the saturated regime depend on the values of the microphysical viscosity and resistivity. In particular, the mean angular momentum transport in the turbulent state increases with increasing magnetic Prandtl number.  From the experimental perspective, understanding the effects of non-vanishing resistivity and viscosity in the behavior of the MRI seems imperative, since the physical conditions achievable in the laboratory depart significantly from the ideal MHD regime \\citep{JGK01, GJ02, Sisanetal04, LGJ06, SSW03}.  Liquid metals (such as sodium, gallium, and mercury) are often characterized by rather low magnetic Prandtl numbers (${\\rm Pm}\\simeq 10^{-5}$--$10^{-7}$). Although the regime of Reynolds numbers involved is still orders of magnitude smaller than any astrophysical system with similar magnetic Prandtl numbers, MRI experiments offer one of the few prospects of studying anything close to MHD astrophysical processes in the laboratory.  A number of analyses addressing some aspects of the impact of viscosity and resistivity on the MRI in various dissipative limits appear scattered throughout the literature on theoretical, numerical, and experimental MRI.  More recently, \\citet{LB07} have found particular solutions of the viscous, resistive MHD equations (including even a cooling term) in the shearing box approximation. However, we are unaware of any comprehensive, systematic study addressing how the MRI behaves in viscous, resistive, differentially rotating magnetized plasmas for arbitrary combinations of the Reynolds and magnetic Reynolds numbers.  The aim of this work is to carry out this analysis in detail.  The rest of paper is organized as follows.  In \\S~\\ref{sec:assumptions}, we state our assumptions. In \\S~\\ref{sec:solution}, we solve the eigenvalue problem defined by the MRI for arbitrary Reynolds and magnetic Reynolds numbers.  We provide closed analytical expressions for the eigenfrequencies and the associated eigenvectors. In \\S~\\ref{sec:mri_modes}, we address the unexplored physical structure of MRI modes for finite Reynolds and magnetic Reynolds numbers and derive simple analytical expressions that describe these modes in various asymptotic regimes.  In \\S~\\ref{sec:stresses_energies}, we calculate the correlations between magnetic and velocity MRI-driven perturbations that are related to angular momentum transport and energy densities. We find that some key results previously shown to hold in the ideal MHD limit \\citep{PCP06} are also valid in the non-ideal regime. In particular, we show that even though the effectiveness with which the MRI disrupts the laminar flow depends on the Reynolds and magnetic Reynolds numbers, the instability always transports angular momentum outwards.  We also find that magnetic perturbations dominate both the energetics and the transport of angular momentum for any combination of the Reynolds and magnetic Reynolds numbers. In \\S~\\ref{sec:discussion} we summarize our findings and discuss the implications of our study. ",
        "conclusions": "\\label{sec:discussion}    We investigated the effects of viscosity and resistivity on the stability of differentially rotating plasmas threaded by a magnetic field perpendicular to the shear. We have shown that the most powerful incompressible MRI modes are exact solutions of the MHD equations for arbitrary combinations of the Reynolds and magnetic Reynolds numbers. We have derived analytical expressions for the eigenfrequencies as well as for the eigenmodes describing the MRI in viscous, resistive media and provided a detailed description of the physical properties of these modes in various dissipative regimes.  We have shown that the scalings derived for the marginally stable mode, the most unstable wavenumber, and the maximum growth rate with magnetic Reynolds number, $k_{\\rm c}, k_{\\rm max}, \\gamma_{\\rm max} \\propto {\\rm Rm}$, valid for resistive, inviscid plasmas, see equations~(\\ref{eq:k_c_lim_reggrm}),~(\\ref{eq:k_max_lim_reggrm})~and~(\\ref{eq:gamma_max_lim_reggrm}), as well as \\citealt{SM99}, also hold when finite Reynolds numbers are involved.  This is true as long as the magnetic Prandtl number is of order unity or smaller, as it is usually the case in many astrophysical systems (such as accretion disks around cataclysmic variables and young stellar objects, as well as the Sun) and also in MRI laboratory experiments.  Furthermore, we have addressed in detail, for the first time to our knowledge, the physical properties of the MRI in highly viscous, slightly resistive media.  These conditions are expected to be found in the hot, diffuse gas in galaxies and galaxy clusters. In this case, we found that the critical wavenumber for the onset of the MRI, the most unstable wavenumber, and the maximum growth rate scale with the Reynolds and magnetic Reynolds numbers according to $k_{\\rm c}\\propto ({\\rm Re\\,Rm})^{1/3}$ and $k_{\\rm max}, \\gamma_{\\rm max} \\propto {\\rm Re}^{1/2}$, see equations~(\\ref{eq:k_c_lim_rellrm}),~(\\ref{eq:k_max_lim_rellrm})~and~(\\ref{eq:gamma_max_lim_rellrm}).  We have provided a thorough geometrical description of the viscous, resistive MRI modes in terms of the angles that define the planes containing the velocity and magnetic field perturbations.  In the ideal MHD limit, these planes are orthogonal, with the plane containing the velocity disturbances laying at $45\\degr$ with respect to the radial direction.  We have shown that velocity and magnetic field perturbations are still orthogonal if the magnetic Prandtl number is unity, but that the planes containing them tend to be aligned with the radial and azimuthal directions, respectively, when the Reynolds number increases. In the regime of large Reynolds numbers and small magnetic Reynolds numbers, magnetic and velocity field perturbations tend to be orthogonal and aligned with the azimuthal and radial directions, respectively. On the other hand, in the regime of small Reynolds numbers and large magnetic Reynolds numbers, the perturbed magnetic field tends to be aligned with the azimuthal direction but the velocity field perturbations do not tend to be aligned with the radial direction. The angle between both fields is determined entirely by the epicyclic frequency $\\kappa$. It would be very interesting to understand to what extent this geometrical dependence of MRI modes on the Reynolds and magnetic Reynolds numbers influences the physical properties of kinetic and magnetic cells in fully developed viscous, resistive MHD turbulence.  \\begin{figure*}[t] \\includegraphics[width=0.675\\columnwidth,trim=0 0 0 0]{f12a.eps} \\includegraphics[width=0.675\\columnwidth,trim=0 0 0 0]{f12b.eps} \\includegraphics[width=0.675\\columnwidth,trim=0 0 0 0]{f12c.eps} \\caption{Mean total stress responsible for angular momentum transport, $\\bar{T}_{r\\phi}=\\bar{R}_{r\\phi}-\\bar{M}_{r\\phi}$, calculated according to equations~(\\ref{eq:mean_Rrphi})~and~(\\ref{eq:mean_Mrphi}), in various dissipative regimes for Keplerian rotation.  \\emph{Left}: Mean total stress $\\bar{T}_{r\\phi}$ as a function of the magnetic Reynolds number for different values of the Reynolds number.  The thick solid line denotes the inviscid limit, i.e., ${\\rm Re}\\rightarrow \\infty$.  The thin solid lines, in decreasing order, correspond to ${\\rm Re} =10, 1, \\ldots, 10^{-4}$.  For magnetic Reynolds numbers larger than unity, this ratio is independent of ${\\rm Rm}$ regardless of the value of ${\\rm Re}$. Moreover, the asymptotic value of this ratio for ${\\rm Rm}\\ll 1$ is independent of the Reynolds number.  \\emph{Middle}: Mean total stress $\\bar{T}_{r\\phi}$ as a function of the magnetic Prandtl number for different values of the Reynolds number.  From left to right, the curves correspond to ${\\rm Pm}=10^3, 10^2, \\ldots, 1$ (thick solid line), $\\ldots, 10^{-6}$.  \\emph{Right}: Mean total stress $\\bar{T}_{r\\phi}$ as a function of the Reynolds number for different values of the magnetic Reynolds number.  The thick solid line corresponds to the ideal conductor limit, i.e., ${\\rm Rm}\\rightarrow \\infty$. The thin solid lines, in decreasing order, correspond to ${\\rm Rm} =10, 1, 0.1$.  Note that the stress scale in these plots is arbitrary.} \\label{fig:Txy_RmPm} \\end{figure*}     In the ideal MHD limit, the exact (primary) MRI modes are known to be unstable to parasitic (secondary) instabilities \\citep{GX94}. These parasitic modes have long been suspected to enable the mechanisms that disrupt the primary modes providing an avenue toward saturation.  The solutions derived in this work describe the dynamics of primary MRI modes in viscous, resistive media enabling the study of parasitic instabilities for arbitrary combinations of Reynolds and magnetic Reynolds numbers.  The modifications in the relative directions of the velocity and magnetic field perturbations characterizing the primary viscous, resistive MRI modes described above can have an important impact on the development and evolution of parasitic instabilities in the presence of dissipation.   We have shown that, for any combination of the Reynolds and magnetic Reynolds numbers, the mean Reynolds stress, $\\bar{R}_{r\\phi}=\\langle \\delta v_r(z,t) \\, \\delta v_\\phi(z,t) \\rangle$, is always positive and the mean Maxwell stress, $\\bar{M}_{r\\phi}=\\langle \\delta b_r(z,t) \\, \\delta b_\\phi(z,t) \\rangle$, is always negative. This implies that the mean total stress, $\\bar{T}_{r\\phi}=\\bar{R}_{r\\phi}-\\bar{M}_{r\\phi}$ is always positive, leading always to an outward transport of angular momentum.  We have also demonstrated that both the ratio between magnetic and kinetic stresses, $-\\bar{M}_{r\\phi}/\\bar{R}_{r\\phi}$, and the ratio between magnetic and kinetic energy densities, $\\bar{E}_M/\\bar{E}_K$, are always dominated by the magnetic contribution.  These last two statements, support a somewhat unexpected result since it is tempting to think that velocity perturbations would dominate both the transport of angular momentum and the energy density in highly resistive, inviscid plasmas.  It would be very interesting to understand if and how the value of these ratios in the saturated turbulent state vary as a function of the Reynolds and magnetic Reynolds numbers.  Sano and collaborators have studied the linear \\citep{SM99} and non-linear \\citep{SIM98, SI01, SS03, Sanoetal04} evolution of the MRI for inviscid, resistive MHD.  The simulations in \\citet{SS03} show that for small magnetic Reynolds numbers the stresses at saturation increase rapidly with increasing ${\\rm Rm}$ and that there exists a critical magnetic Reynolds number, of order unity, beyond which turbulent stresses are rather insensitive to ${\\rm Rm}$. The fact that this same behavior is indeed seen when the stresses are due to viscous, resistive MRI modes (see Figure~\\ref{fig:Txy_RmPm}) rises the question of how strong is the influence of long-lived, channel-like modes on the fully developed turbulent state reached in shearing box simulations with net magnetic flux through the vertical boundaries.  Systematic numerical studies of viscous, resistive MHD shearing flows have begun to uncover the dependencies of microphysical dissipation on the mean transport properties of MRI-driven turbulence. Numerical simulations with both zero \\citep{FPII07} and non-zero net magnetic fluxes \\citep{LL07} lead to the conclusion that angular momentum transport increases with increasing magnetic Prandtl number when the Reynolds number is held constant. This behavior can also be identified when examining the stresses due to viscous, resistive MRI modes, see~Figure~\\ref{fig:Txy_RmPm}.  The effects of varying the Reynolds number at fixed magnetic Prandtl number on the saturation of MRI-driven turbulence are currently rather uncertain \\citep[see, in particular, the discussion in][]{LL07}.  The global trends exhibited by the available simulations suggest that the stresses at saturation increase with increasing Reynolds number for magnetic Prandtl numbers smaller than unity while the opposite behavior is observed for magnetic Prandtl numbers larger than unity. If confirmed, these results suggest that the mechanisms leading to saturation might operate differently depending on whether the magnetic Prandtl number is larger or smaller than unity.  In any case, having obtained a better understanding of the behavior of the most unstable MRI modes as a function the magnetic Prandtl number it would be very interesting to follow the evolution of the viscous, resistive MRI from the linear to the non-linear regime. By performing numerical simulations with the same Prandtl number (both larger and smaller than unity) and different Reynolds numbers we could see whether there is an inversion of the trends observed in the linear regime (i.e., higher stresses at higher Reynolds numbers for fixed magnetic Prandtl numbers) after the exact solutions break down.  The comparison between the late time behavior of the viscous, resistive MRI modes and fully developed MHD turbulence with dissipation will shed light into the mechanisms that lead to the saturation of the MRI in different dissipative regimes.  Finally, most current numerical algorithms employ finite difference methods (with constrained transport for the evolution of the magnetic field).  The leading order errors in first-order upwind methods behave like diffusion, however, in second-order central difference methods, the leading order errors are dispersive (artificial viscosities are usually employed to damp unphysical oscillations near shocks). The comparison between numerical solutions from different algorithms can provide estimates of these errors.  However, it is difficult to quantify these numerical artifacts based on analytical studies of ideal MHD. The analytical solutions derived in this paper, on the other hand, describe the effects of arbitrary combinations of viscosity and resistivity [see also \\citet{LB07} who derived results to leading order in $(\\eta-\\nu)/k$]. It should now be possible to better measure the numerical viscosity and resistivity for a wide range of Reynolds and magnetic Reynolds numbers by comparing numerical solutions and analytical solutions of non-ideal MRI.  The results presented in this paper provide ideal benchmarks to the study numerical artifacts generated by different algorithms in various dissipative regimes.  \\appendix"
    },
    "0801/0801.1213_arXiv.txt": {
        "abstract": "We combine VLT/ISAAC NIR spectroscopy with archival HST/WFPC2 and HST/NICMOS imaging to study the central 20\\arcsec$\\times$20\\arcsec\\ of M83. Our NIR indices for clusters in the circumnuclear star-burst region are inconsistent with simple instantaneous burst models. However, models of a single burst dispersed over a duration of 6~Myrs fit the data well and provide the clearest evidence yet of an age gradient along the star forming arc, with the youngest clusters nearest the north-east dust lane. The long slit kinematics show no evidence to support previous claims of a second hidden mass concentration, although we do observe changes in molecular gas velocity consistent with the presence of a shock at the edge of the dust lane. ",
        "introduction": "M83 (NGC 5236) is a nearby (4.5 Mpc; \\citealt{Thim03}) grand design spiral galaxy SAB(s)c \\citep{deVauc91} in a group containing several bright galaxies including Centaurus A (NGC 5128) and NGC 4945. Within the central 20\\arcsec\\ lies the photometric peak (hereafter referred to as the nucleus) offset from the centre of symmetry of the outer bulge isophotes and the global gas kinematics \\citep{Wolst88,Gallais91,Thatte00,Sakamoto04} and surrounded by a semicircular starburst region, some 3\\arcsec\\ - 8\\arcsec\\ from the nucleus. M83 has a close dynamical companion, NGC 5253 \\citep{Rog74} which also contains recent star formation, albeit much more compact than that of M83. The closest passage of NGC 5253 occurred 1-2 Gyrs ago \\citep{Rog74} and so the activity in these galaxies is likely driven by internal inflow of gas via the bar rather than from direct interaction as the clusters are mostly 1-10 Myrs in age \\citep[hereafter H01]{Harris01}.   \\subsection{Central star formation}\\label{sec:csf} ",
        "conclusions": "\\label{sec:conc}  We have shown that the NIR indices \\wpa\\ and \\wco\\ can be used to age-date young star clusters. Furthermore, \\wpa\\ is much less affected by extinction than \\wha, which requires significant correction. The young clusters in the circumnuclear arc are inconsistent with simple stellar populations produced by an instantaneous burst; the observations require a finite episode of star formation lasting around 6 Myrs. Such models together with our observations provide the clearest evidence yet of an age gradient along the star forming arc, as first found in H01. However, our NIR observations further show that the circumnuclear arc does not stop at the edge of the dust lane, but continues through and along its north-west edge.  Our stellar kinematics show no evidence for a second obscured nucleus at either the location of D06b or T00 and we conclude that the velocity gradients in the gas kinematics (ionised and molecular) are a result of a shock front on the south-eastern edge of the north-eastern dust lane. We therefore explore alternative explanations for the formation of the circumnuclear arc, the most plausible being the $x_2$ orbital motion of the clusters from their place of birth in the dust lane.  Future higher-resolution integral field spectroscopy in the NIR has the potential to build on this study and find the youngest clusters in the arc (suspected to be on the north-west edge of the northern dust lane) and also address previous claims of an age gradient \\emph{perpendicular} to the arc (H01)."
    },
    "0801/0801.0215_arXiv.txt": {
        "abstract": "We construct a new model of Type Ia Supernovae (SNe Ia), based on the single degenerate scenario, taking account of the metallicity dependences of the white dwarf (WD) wind and the mass-stripping effect on the binary companion star. Our model naturally predicts that the SN Ia lifetime distribution spans a range of $0.1-20$ Gyr with the double peaks at $\\sim 0.1$ and $1$ Gyr. While the present SN Ia rate in elliptical galaxies can be reproduced with the old population of the red-giants+WD systems, the large SN Ia rate in radio galaxies could be explained with the young population of the main-sequence+WD systems. Because of the metallicity effect, i.e., because of the lack of winds from WDs in the binary systems, the SN Ia rate in the systems with [Fe/H] $\\ltsim -1$, e.g., high-z spiral galaxies, is supposed to be very small. Our SN Ia model can give better reproduction of the [($\\alpha$, Mn, Zn)/Fe]-[Fe/H] relations in the solar neighborhood than other models such as the double-degenerate scenario. The metallicity effect is more strongly required in the presence of the young population of SNe Ia. We also succeed in reproducing the galactic supernova rates with their dependence on the morphological type of galaxies, and the cosmic SN Ia rate history with a peak at $z \\sim 1$. At $z \\gtsim 1$, the predicted SN Ia rate decreases toward higher redshifts and SNe Ia will be observed only in the systems that have evolved with a short timescale of chemical enrichment. This suggests that the evolution effect in the supernova cosmology can be small. ",
        "introduction": "The star formation histories of galaxies are imprinted on stellar populations of the galaxies. Because of the age-metallicity degeneracy for stellar populations, it is difficult to determine ``age'' from their colors and spectra alone. However, the elemental abundance ratios can provide independent information on ``age'', because different heavy elements are produced from different supernovae with different timescales \\citep[e.g.,][]{pagel1997,matteucci2001}. To derive the information of ``ages'' from elemental abundance patterns as a cosmic clock, the most important uncertainty is the lifetime of Type Ia Supernovae (SNe Ia), i.e., the evolutionary time (delay-time) of the progenitor from the main-sequence through the explosion, which is mainly determined by the lifetime of the relatively low-mass ($\\ltsim 8 M_\\odot$) companion star of the exploding white dwarf (WD).  There exist two distinct types of supernova explosions \\citep[e.g.,][]{arnett1996,alex1997}; One is Type II supernovae (SNe II), which are the core collapse-induced explosions of massive stars ($\\gtsim \\, 8M_\\odot$) with short lifetimes of $10^{6-7}$ yrs, and produce more $\\alpha$-elements (O, Mg, Si, S, Ca, and Ti) relative to Fe with respect to the solar ratios (i.e., [$\\alpha$/Fe] $>0$). Type Ib and Ic Supernovae (SNe Ib and Ic) are also the core-collapse supernovae, but have lost the envelope material before the explosions. Recently, it is found that bright core-collapse supernovae, hypernovae (HNe), produce an important amount of Fe \\citep[e.g.,][]{nom06}. In chemical evolution models, the contributions of SNe Ib, Ic, and HNe can be included in that of SNe II \\citep{kob06}. The other is SNe Ia, which are the thermonuclear explosions of accreting WDs in close binaries and produce mostly Fe and little $\\alpha$-elements \\citep[e.g.,][]{nom94,hil00}. The progenitors of the majority of SNe Ia are most likely the Chandrasekhar (Ch) mass WDs (e.g., Nomoto, Iwamoto \\& Kishimoto 1997 for a review; see also Nomoto et al. 2007), although the sub-Ch mass models might correspond to some peculiar subluminous SNe Ia. The early time spectra of the majority of SNe Ia are in excellent agreement with the synthetic spectra of the Ch mass models, while the spectra of the sub-Ch mass models are too blue to be compatible with observations \\citep{hof96,nug97}.  For the evolution of accreting C+O WDs toward the Ch mass, two scenarios have been proposed; One is a double-degenerate (DD) scenario, i.e., merging of double C+O WDs with a combined mass surpassing the Ch mass limit \\citep{ibe84,web84}. From the long-term observational search (SPY project) of the WD binaries for the DD systems, more than thousand WDs are found. However, their combined masses are smaller than the Ch mass, except for a few candidates near the Ch mass \\citep{nap07,gei07,tov07}. Theoretically, it has been suggested that the merging of WDs leads to accretion-induced collapse rather than SNe Ia \\citep{sai85,sai98}. The other is a single-degenerate (SD) scenario, i.e., the WD mass grows by accretion of hydrogen-rich matter via mass transfer from a binary companion \\citep[e.g., ][for reviews]{nom00a,liv01}. Two progenitor systems have been found for red-giant (RG) and near main-sequence (MS) companions \\citep[e.g., ][]{hac99u}. The SD scenario has predicted the detection of hydrogen-rich circumstellar matter from the companion star for a certain class of SN Ia binary systems, and such circumstellar matter has been actually detected for some SNe Ia \\citep[e.g., ][]{ham03,pat07}.  The lifetime of SNe Ia has been estimated from different aspects as follows, and these estimates conflict each other; (1) {\\it Population synthesis} --- The lifetime distribution of SNe Ia has been predicted with the population synthesis models for the evolution of binary systems \\citep[e.g.,][]{yun05,bel05,tut07}. The predicted rate of SD systems in those models is much smaller than DD systems, and the majority of SNe Ia have the lifetime as short as $\\sim 0.1$ Gyr \\citep[e.g.,][]{rui95,yun98}. However, the optically thick winds from the WD, which were not taken into account in the above models, have been shown to play an essential role in the evolution of the accreting WDs \\citep{hac96,hac99,hac99u}. Furthermore, the stripping effect may increase the parameter space for the SD systems \\citep{hac08}. If these effects are taken into account in the population synthesis models, the DD rate could be smaller, and the SD rate could be larger.  (2) {\\it Chemical evolution} --- The typical lifetime of SNe Ia could have been constrained from the chemical evolution of the Milky Way Galaxy, most notably in the [$\\alpha$/Fe]-[Fe/H] relations. Metal-poor stars with [Fe/H] $\\ltsim -1$ have [$\\alpha$/Fe] $\\sim 0.5$ for Mg, Si, and Ca on the average \\cite[e.g.,][]{mcw95,cay04}, while disk stars with [Fe/H] $\\gtsim -1$ show a decrease in [$\\alpha$/Fe] with increasing metallicity \\cite[e.g.,][]{edv93,ben04}. Such an evolutionary change in [$\\alpha$/Fe] against [Fe/H] has been explained with the early heavy-element production by SNe II and the delayed enrichment of Fe by SNe Ia \\citep{mat86}. Conversely, chemical evolution models can constrain the nature of progenitor systems of SNe Ia; for example, Yoshii, Tsujimoto, \\& Nomoto (1996) estimated the lifetime of SN Ia progenitors to be as long as $0.5-3$ Gyr.  (3) {\\it Supernova rates} --- The redshift evolution of the SN Ia rate can give a rough estimate of the typical lifetime of SNe Ia (i.e., delay-time) for the assumed star formation history. The decrease of the SN Ia rate seen at $z \\sim 1$ corresponds to the delay-time of $\\sim 3$ Gyr \\citep{str04}. However, the observed cosmic star formation rates (SFRs) involve uncertainties from dust extinction and completeness. The environmental dependence on galaxy evolution must be important if the SNe Ia rate depends on metallicity. SN Ia rates depending on the galaxy type can give stronger constraint on the SN Ia progenitor models \\citep{cal06}. \\citet{del05} suggested that the short lifetime of $\\sim 0.1$ Gyr is required to explain the high SN Ia rate in radio-loud elliptical galaxies, and \\citet{sul06} suggested the similar lifetime distribution from the relation between the SN Ia rate and the specific SFR. Therefore, the bimodal distribution of the lifetime of SNe Ia, so-called ``A+B'' model, has been proposed \\citep{sca05,man06}.  In addition, Kobayashi et al. (1998, hereafter K98) have found the metallicity effect on the WD winds, consequently, on the occurrence of SNe Ia, based on the simulations of the progenitor evolution by Hachisu et al. (1999ab). If the iron abundance of the progenitors is as low as [Fe/H] $\\ltsim -1.1$, then the wind is too weak for SNe Ia to occur. Thus the SN Ia rate in the system with [Fe/H] $\\ltsim -1.1$ is supposed to be ``very small''. This prediction has not been observationally confirmed yet, and several SNe Ia in metal-poor environment have been reported \\citep[e.g.,][]{pri08}. Statistical studies of a large sample are required, since there must be a spread in the metallicity distribution function of stars in a galaxy and also the [O/Fe] is uncertain (see \\S \\ref{sec:zeffect} for the details).  If the lifetime distribution function of SNe Ia depends on metallicity, the detailed calculation of chemical evolution is even more important to understand the observations of supernova rates and abundance ratios. The redshift evolution of the cosmic SN Ia rate is affected by chemical enrichment histories in various types of galaxies. The estimates of the star formation histories of galaxies from their abundance ratios are affected by the metallicity effect on the SN Ia rate. In addition, Fe can be produced by the four sources \\citep{kob06}: i) SNe Ia, ii) relatively large Fe production from low-mass SNe II with $M=13-15M_\\odot$, iii) large Fe production from HNe, or iv) the SN I.5 explosion of some AGB stars. Therefore, not only [$\\alpha$/Fe] but also other elements should be considered to discuss the contribution of SNe Ia.  This may be the case for dwarf spheroidal (dSph) galaxies; the observed low [$\\alpha$/Fe] at low [Fe/H] in dSphs is often suggested to be due to the significant contribution of SNe Ia at such low metallicity \\cite[e.g.,][]{ven04}. However, there is a large variety among galaxies, and a decreasing trend in [Mg/Fe] at [Fe/H] $< -1$ may be seen in Draco \\citep{she01} and Sculper \\citep{she03}, but not for other $\\alpha$ elements such as Si and Ca. For more metal-poor stars, some stars show high [$\\alpha$/Fe] \\citep{koc08,fre09}, but others show low [$\\alpha$/Fe] \\citep{aok09,fre09b}. Since [Mn/Fe] in dSph galaxies is as low as in Milky Way halo stars, the abundance pattern can be better explained with low-mass SNe II than SNe Ia \\citep[][hereafter K06]{tol03,kob03,kob06}. This is also supported by the low [Ba/Fe] in \\cite{aok09}. Chemodynamical evolution of dSph galaxies are complicated because of the low SFR and large dark matter content, and can be largely affected only by a few supernovae. At least, the low [$\\alpha$/Fe] in low metallicity stars does not necessarily mean the contribution of SNe Ia, and certainly more detailed studies are required under the reasonable model of SN Ia progenitors.  Similarly, it is argued that damped Lyman $\\alpha$ (DLA) systems show low [$\\alpha$/Fe] at low [Fe/H] \\citep[e.g., ][]{pet99}. \\citet{pro02}, however, concluded that [Si/Fe] $\\sim 0.3$ at [Fe/H] $< -1.5$, and also similar values were presented in \\citet{des07}. The elemental abundance ratios in DLA systems are affected by the dust depletion. The dust depletion has been often estimated using [Zn/Fe] ratio, but there is no assurance that [Zn/Fe] $\\sim 0$ for low-metallicity. In fact, an increasing trend of [Zn/Fe] has been observed in the Milky Way Galaxy \\citep{pri00,nis07}. \\citet{nis07} showed that [S/Zn], which is free from the dust effect, of halo stars is similar to that in DLA systems. Such [S/Zn] and [Si/Fe] could be explained by hypernovae. Since the observation is for neutral gas in galaxies, it can be easily affected by dilution of hydrogen in chemodynamical simulations \\citep{kob07}. Again, more detailed studies are required.  The modeling of SNe Ia in the chemical evolution was first made by Greggio \\& Renzini (1983), where the SN Ia rate is calculated from the distribution of binaries based on Whelan \\& Iben (1973), and then extended by Greggio (1996) including the efficiency of accretion of the envelope of the secondary to the WD. This model has recently been updated in \\citet{gre05} and \\citet{mat06}. An alternative model was proposed by K98, fully based on the SD scenario, where the metallicity effect is essential to explain the [O/Fe]-[Fe/H] relation in the solar neighborhood. Kobayashi, Tsujimoto, \\& Nomoto (2000, hereafter K00) applied this model for different types of galaxies, and succeeded in reproducing the present SN II and Ia rates in different types of galaxies. Matteucci \\& Recchi (2001, hereafter MR01), however, argued that they could not reproduce K98 results and that the lack of SNe Ia at low metallicities produces results at variance with the observed [O/Fe]-[Fe/H] relation. However, this is due to their misunderstanding of the parameterization of the K98 models, as described in detail in \\S 5.  In this paper, we first show the lifetime distribution functions derived from different SN Ia models. We then introduce our new SN Ia model by taking into account the stripping effect \\citep{hac08} and show the lifetime distribution function as a function of metallicity (\\S 2). We then evaluate the SN Ia models from the comparison with observations. The evolution of the elemental abundance ratios in the solar neighborhood (\\S 3) and the galactic and cosmic supernova rates (\\S 4) can be better reproduced with our SN Ia model including the metallicity effect. In \\S 5, we discuss the two formulations in K98 and MR01 in detail, and reproduce both results using the alternative formulation. In \\S 6, we summarize our conclusions. The detailed discussion on the elemental abundance ratios in other galaxies will be in the second paper. ",
        "conclusions": "We construct a new SN Ia model, based on the SD scenario for the SN Ia progenitors, taking account of the metallicity dependences of the WD wind (K98) and the mass-stripping effect on the binary companion star (HKN08). The lifetime distribution of SNe Ia is determined from the main-sequence mass range of the companion stars in the WD binary systems, and is given as a function of metallicity (Fig. \\ref{fig:life3}). The SN Ia rate in the systems with [Fe/H] $\\ltsim -1$, e.g., $r \\gtsim 10r_{\\rm e}$ in intermediate-mass galaxies, $r \\gtsim 50$ kpc in dwarf galaxies, and high-z spiral galaxies, is supposed to be very small (\\S \\ref{sec:zeffect}).  Our model naturally predicts that the SN Ia lifetime distribution spans a range of $t_{\\rm Ia} \\sim 0.1-20$ Gyr with the double peaks at $\\sim 0.1$ and $1$ Gyr reflecting the two systems of the companion stars: the MS+WD and RG+WD systems, respectively. With the chemical enrichment history in the Galactic disk, the fraction of the young population of SNe Ia is $10\\%$ and $50\\%$ for $t_{\\rm Ia}<0.1$ and $1$ Gyr, respectively (Fig. \\ref{fig:snlife}). In the galaxies with a initial star burst, the SN Ia rate shows a strong peak at $t \\sim 1.5$ Gyr, which may generate the SN Ia driven galactic winds.  We make a comparison of the lifetime distribution functions derived from different SN Ia models, and evaluate the results from the observational constraints as summarized in Table \\ref{tab:summary}.  \\noindent (1) Contrary to the argument in Matteucci \\& Recchi (2001, MR01), our SN Ia models give better reproduction of the [$\\alpha$/Fe]-[Fe/H] relation in the solar neighborhood. In the DD and MR01-like models, the typical lifetimes of SNe Ia are $\\sim 0.1$ and $0.3$ Gyr, respectively, which results in the too early decrease in [$\\alpha$/Fe] at [Fe/H] $\\sim -2$. If we do not include the metallicity effect in our model, [$\\alpha$/Fe] decreases too early because of the shortest lifetime, $\\sim 0.1$ Gyr, of the MS+WD systems. In other words, the metallicity effect is more strongly required in the presence of the young population of SNe Ia, to be consistent with the chemical evolution of the solar neighborhood.  \\noindent (2) At the same [Fe/H], [Mn/Fe] starts to increase toward higher metallicity, because SNe Ia produce [Mn/Fe] $>0$. This evolutionary change is started from the same [Fe/H] as the decreasing trends of [$\\alpha$/Fe]. [Zn/Fe] evolves differently for the different SNe Ia models, because of the combination of the different lifetime of SNe Ia and the metallicity effect on Zn yields of SNe II. Further observations of these elements at [Fe/H] $\\gtsim -1.5$ is critically important in order to identify the SN Ia progenitors and to clarify the metallicity dependence on the nucleosynthesis yields. The lack of the similar [Mn/Fe] increase in dSphs suggests that those are not enriched by SNe Ia.  \\noindent (3) To explain the observed SN Ia in ellipticals at $z=0$, the lifetimes of a sufficient fraction of SNe Ia should be longer than $\\sim 10$ Gyr, which is satisfied in our SD model and the MR01-like model, but not in the DD model. Our SN Ia models are also successful in reproducing the SN II, Ibc, and Ia rates with their dependence on the galaxy type \\citep{man06}; This cannot be reproduced with the DD model. The large SN Ia rate in radio galaxies could be explained with the young population of the MS+WD systems in our model.  \\noindent (4) We also provide the cosmic supernova rate history as a composite of those in spirals and ellipticals. The cosmic SN Ia rate in comoving density shows a peak at $z\\sim1$ and decreases toward higher redshift in our model. In contrast, this rate increases toward higher redshifts in the DD and MR01-like models. Because of the metallicity effect, i.e., because of the lack of winds from WDs in the binary systems at [Fe/H] $< -1.1$, the SN Ia rate in the low-metallicity systems, e.g., high-z spiral galaxies, is supposed to be very small in our model. This metallicity effect will appear only around [Fe/H] $\\sim -1$ (Fig. \\ref{fig:zeffect}). In contrast, the SN Ia rate in such systems is as high as present in the DD and MR01-like models. In our models, at $z \\gtsim 1$, SNe Ia will be observed only in the systems that have evolved with a short timescale of chemical enrichment. This suggests that the evolution effect in the supernova cosmology can be small."
    },
    "0801/0801.0486_arXiv.txt": {
        "abstract": "We develop an unconditionally stable numerical method for solving the coupling between two fluids (frictional forces/heatings, ionization, and recombination), and investigate the dynamical condensation process of thermally unstable gas that is provided by the shock waves in a weakly ionized and magnetized interstellar medium by using two-dimensional two-fluid magnetohydrodynamical simulations. If we neglect the effect of magnetic field, it is known that condensation driven by thermal instability can generate high density clouds whose physical condition corresponds to molecular clouds (precursor of molecular clouds). In this paper, we study the effect of magnetic field on the evolution of supersonic converging HI flows and focus on the case in which the orientation of magnetic field to converging flows is orthogonal. We show that the magnetic pressure gradient parallel to the flows prevents the formation of high density and high column density clouds, but instead generates fragmented, filamentary HI clouds. With this restricted geometry, magnetic field drastically diminishes the opportunity of fast molecular cloud formation directly from the warm neutral medium, in contrast to the case without magnetic field. ",
        "introduction": "It is widely known that radiative cooling and heating in the interstellar medium (ISM) leads to a stable cold dense medium and a warm diffuse medium (Field, Goldsmith \\& Harbing 1969; Wolfire et al. 1995, 2003). These stable phases are called the cold neutral medium (CNM), and the warm neutral medium (WNM). The CNM is observed as HI clouds ($n \\sim 10^{1-2}$ cm$^{-3}$, $T\\sim 10^{2}$ K), and cold neutral clouds that contain substantial amounts of H$_{2}$ are called molecular clouds ($n \\sim 10^{2-3}$ cm$^{-3}$, $T\\sim 10$ K). Since molecular clouds are the sites of star formation, the knowledge of how molecular clouds are formed is crucial to understanding the initial conditions of star formation.  Thermal instability (TI) is the most promising formation mechanism of the CNM. Many authors studied the dynamical condensation process of the ISM driven by TI. Koyama \\& Inutsuka (2000, 2002), and Inutsuka, Koyama, \\& Inoue (2005) studied the CNM formation as a result of TI in the layer compressed by shock propagation. Hennebelle \\& P\\'erault (1999), Audit \\& Hennebelle (2005), Hennebelle \\& Audit (2007), Heitsch et al. (2005, 2006), and Vazquez-Semadeni et al. (2006, 2007) studied an analogous process in supersonic converging flows, which essentially include two shocks. These studies showed that TI can generate molecular clouds, by piling up the WNM.  However, the above-mentioned studies do not include the effect of magnetic field. The dynamical formation process of the CNM with the effect of magnetic field is studied by Hennebelle \\& P\\'erault (2000) using one-dimensional, ideal MHD simulations. They showed that magnetic pressure prevents TI, if the initial angle between the magnetic field and the fluid velocity is larger than $20^\\circ$-$30^\\circ$ with a few microgauss initial field strength. Recently, Inoue, Inutsuka, \\& Koyama (2007) have shown that if the effect of ambipolar diffusion is taken into account, TI can create the CNM even with a microgauss magnetic field which is orthogonal to the orientation of the condensation.  In Inoue, Inutsuka, \\& Koyama (2007), we used a thermally unstable equilibrium as an initial condition and performed one-dimensional numerical simulations. However, thermally unstable gas is likely to be supplied by the shock compression of the ISM, and in multiple dimensions, growth of TI along the magnetic field seems to be faster than the mode that is perpendicular to the magnetic field. In this paper, to investigate how the CNM is formed under more realistic situations, we study the dynamical condensation process of thermally unstable gas that is supplied by the shock compression of the WNM employing two-dimensional two-fluid magnetohydrodynamical simulations.  This paper is organized as follows. In \\S 2 we provide the basic equations of the weakly ionized interstellar medium and assumptions used in this paper. In \\S 3 we describe a method for solving the basic equations numerically, in which an unconditionally stable scheme that solves the coupling between two fluids is proposed and tested. The results of two-fluid simulations of TIs in weakly ionized media are shown in \\S 4. Finally, in \\S 5 we summarize our results and discuss their implications. ",
        "conclusions": "We have developed an unconditionally stable numerical method for solving the coupling between two fluids, which can adequately solve the evolution of a partially ionized medium from weak-coupling to strong coupling regimes. By using two-dimensional two-fluid MHD simulations based on this method, we have investigated the dynamical formation process of the CNM in both the magnetized and unmagnetized media. Our findings are as follows.  In the unmagnetized case, the converging flows of the WNM create a complex slab bound by quasi-standing shocks in which turbulent cloudlets of the CNM are generated as a consequence of TI. Because the thermal pressure in the slab has to sustain the ram pressure of the flows, the CNM whose pressure is higher than the average pressure of the ISM and whose density is $n\\simeq 10^{3}$ cm$^{-3}$, which correspond to the observational characteristics of molecular clouds, is inevitably formed. Thus, as long as the converging flows continue, the column density of the CNM increases, and a molecular cloud would be generated in the slab.  In the magnetized case, magnetic pressure inside the shocked region can sustain the ram pressure of the flows. This leads to two important effects on the evolution of the shocked region. One is that, because the thermal pressure does not have to be high in order to oppose the ram pressure, it decreases due to cooling. This makes the physical conditions of the resulting CNM similar to those of HI clouds. Another one is that the shocks continue to propagate, supplying thermally unstable gas to an extended region. This makes it difficult to generate high-column density clouds in a compact region like molecular clouds. This is because ambipolar diffusion is not effective globally, and TI cannot pile up material in the direction parallel to the converging flows.  In this paper, as a first step, we examined the situation in which the initial angle between the magnetic field and the velocity of the colliding flows is $90^\\circ$, which maximizes the effect of magnetic field. If the initial angle is small enough, it seems possible to generate the CNM whose physical condition corresponds to molecular clouds, as in the unmagnetized case. How large is the critical angle that changes the character of the resulting CNM? In a one-dimensional case, according to Hennebelle \\& P\\'erault (2000), magnetic pressure prevents condensation, if the initial angle is larger than $20^\\circ$-$30^\\circ$ with a few microgauss initial field strength (see Fig. 10 of their paper). In more realistic two-dimensional situation, as we have shown in this paper, the mode of TI along the magnetic field always generates the CNM that corresponds to a HI cloud. Thus, their conclusion needs to be modified. However, their result tells that if the initial angle is smaller than their critical angle, magnetic pressure cannot be strong enough to support the ram pressure. In such a case, the resulting shocked slab would have a physical condition that can generate molecular clouds (precursor of molecular clouds), provided the flows last long.  However, if this conjecture is right, smallness of the critical angle will drastically diminish the opportunity of fast molecular cloud formation directly from the WNM whose time scale is on the order of growth time scale of TI, or cooling time, ($\\sim1$-$10$ Myr). Thus, there is a possibility that most of the molecular clouds are formed from HI clouds, as suggested from observations of clouds in Local Group galaxies (Fukui 2007; Blitz et al. 2007), via more complex processes (e.g., experiences of TI more than once due to successive shocks of supernovae and possibly through the effect of potential well of stellar spiral waves or self-gravity) rather than the direct pile-up of WNM as a consequence of TI. The diffuculty of the fast formation of molecular clouds in the presence of magnetic field perpendicular to the inflow is also discussed from another point of view in Hartmann et al. (2001).  In order to confirm this possibility, the following studies are necessary: (1) similar calculations of converging flows varying the angles between converging flows and magnetic field to show the smallness of the critical angle even in multiple dimensions; and (2) a global study of the magnetized ISM to learn about the typical angle between the mean magnetic field and converging flows or interstellar shocks. In subsequent papers, we will present studies relevant to these points.  Note that the discussion shown in this section is based on one-/two-dimensional simulations. However, especially in MHD, dynamical behavior in 3D can be different from that in 2D (see, e.g., Stone \\& Gardiner 2007 for Rayleigh-Taylor instability in a magnetized medium). Thus, we should bear in mind the possibility that the difference due to the dimensional restriction might change the outcome."
    },
    "0801/0801.2160_arXiv.txt": {
        "abstract": "We  explore the  role of  AGN in  establishing and/or  maintaining the bimodal   colour   distribution  of   galaxies   by  quenching   their star-formation and  hence, causing their  transition from the  blue to the  red cloud.   Important tests  for this  scenario include  (i) the X-ray properties  of galaxies in  the transition zone between  the two clouds and (ii) the incidence of AGN in post-starbursts, i.e.  systems observed   shortly  after   ($<1$\\,Gyr)  the   termination   of  their star-formation.   We  perform  these  tests  by  combining  deep  {\\it Chandra} observations with multiwavelength data from the AEGIS survey. Stacking the X-ray photons  at the positions of galaxies ($0.4<z<0.9$) not individually  detected at X-ray wavelengths  suggests a population of obscured  AGN among sources in  the transition zone and  in the red cloud.   Their mean X-ray  and mid-IR  properties are  consistent with moderately obscured low-luminosity AGN, Compton thick sources or a mix of both.  Morphologies  show that major mergers are  unlikely to drive the evolution  of this  population but minor  interactions may  play a role.  The  incidence of  obscured AGN in  the red cloud  (both direct detections and  stacking results) suggests that  BH accretion outlives the termination of the star-formation.   This is also supported by our finding that  post-starburst galaxies at  $z \\approx 0.8$ and  AGN are associated,  in agreement  with recent  results at  low-$z$.   A large fraction of  post-starbursts and red cloud galaxies  show evidence for at least moderate levels of AGN obscuration.  This implies that if AGN outflows  cause  the  colour  transformation of  galaxies,  then  some nuclear gas and  dust clouds either remain unaffected  or relax to the central galaxy regions after quenching their star-formation. ",
        "introduction": "A major recent development in extragalactic astronomy is the discovery that  most  of   the  spheroids  in  the  local   Universe  contain  a super-massive black hole  (BH; e.g. Magorrian et al.   1998). The masses of  these  monsters  are tightly  correlated  to  the  stellar  velocity dispersion of  the host  galaxy bulges (e.g.   Ferrarese et  al.  2000; Gebhardt et al. 2000), suggesting  that the formation and evolution of spheroids and the  build-up of the super-massive BHs  at their centres are interconnected.  This interplay may hold the key for understanding some of the observed trends in galaxy evolution.   One  of  the  fundamental  properties  of galaxies  is  their  bimodal distribution  in rest-frame  colour.  This bimodality  is observed  to beyond redshift $z\\approx1$ and is believed to hold important clues on galaxy assembly (e.g. Strateva et al.  2001; Baldry et al.  2004; Bell et al.   2004; Weiner et al.   2005; Cirasuolo et  al. 2007).  Evolved spheroidal   galaxies  are   found   in  the   ``red-cloud''  of   the colour-magnitude diagram (CMD), star-forming systems define the ``blue cloud'',  while the  region  between these  overdensities is  sparsely populated.   This bimodal distribution  is likely  to be  the combined result of evolution effects  and external factors such as interactions and mergers (e.g.  Bell et al.  2004, 2006a, 2006b; Blanton 2006).  In the  simplest interpretation,  the CMD  is consistent  with  a picture where  the galaxy  star-formation is  truncated, by  a  mechanism that remains  to  be identified,  leading  to  the  ageing of  the  stellar population and  the rapid transition from  the blue to  the red cloud. AGN are  proposed to play a  major part in this  process by regulating the star-formation  in galaxies causing their transition  in colour or maintaining them in the red cloud.  Modelling  work indeed  suggests that  the energy  released by  AGN is sufficient to  either heat up  or blow away  the cold gas  of galaxies (e.g.   Silk  \\& Rees  1998;  Fabian  1999;  King 2003),  irreversibly altering their evolution.   Two main routes are proposed  by which AGN feedback affects the  host galaxy.  The first one  operates during the luminous  and  high  accretion  rate  stage of  BH  growth,  which  is suggested to occur during major  mergers (Hopkins et al.  2005, 2006a, b).  These  catastrophic events  also produce nuclear  starbursts that obscure the central engine for  most of its active lifetime.  When the AGN becomes sufficiently luminous it drives outflows, which sweep away the  nuclear  gas  and  dust  clouds, thereby  quenching  the  nuclear star-formation.  Following this stage, the AGN also declines as the BH runs out of accreting material and eventually switches off.  In addition  to the  cold gas accretion  mode above,  simulations also include a second AGN feedback recipe, which is invoked to suppress the accretion  of hot gas  onto the  central galaxy  of large  dark matter haloes (e.g.  Croton  et al.  2006; Okamoto et  al.  2007; Cattaneo et al. 2007).  This  mode operates in systems where  the supermassive BHs have  already been  built  through mergers  and  are in  place at  the central  galaxy  regions.  In  this  picture,  as  galaxies enter  the massive  dark matter  haloes of  groups  or clusters,  their cold  gas supply is cut off leading  to passive evolution and their migration to the red cloud (Dekel \\&  Birnboim 2006; Croton et al. 2006).  However, cooling flows in these dense environments could produce a reservoir of cold gas in the central galaxy regions, the raw material for starburst and  QSO activity.   Low-level BH  accretion is  therefore  invoked to counterbalance cooling flows by producing low-luminosity AGN that heat the bulk  of the  cooling gas and  thereby suppressing  any subsequent star-formation in these galaxies.   In  the two  AGN feedback  prescriptions above,  the  outflow scenario during  the luminous  and  high  accretion rate  phase  of BH  growth, assigns  AGN  a  central  role  in  establishing  the  bimodal  colour distribution of galaxies by directly causing their transition from the blue to the red cloud.  The  distribution of X-ray selected AGN in the CMD  appears to  be broadly  consistent  with this  picture.  AGN  are preferentially associated with  galaxies in the red cloud,  at the red limit of the  blue cloud and in the transition  zone in between. These associations   are  suggestive   of  AGN-driven   truncation   of  the star-formation (Nandra et al.  2007).  However, more detailed study of the  properties of  AGN reveal  a number  of inconsistencies  with the merger scenario.  Obscured AGN are predominantly hosted by galaxies in the red cloud (Nandra et  al.  2007; Rovilos \\& Georgantopoulos 2007), with prominent  bulges and  little morphological evidence  for ongoing major mergers  (Grogin et  al.  2003; Pierce  et al.  2007).   The red optical colours of these systems  are most likely because of old stars and not  dusty star-formation.  Although there are  examples of deeply buried AGN  in star-forming  galaxies (e.g. Genzel  et al.   1998; Cid Fernandes et al.  2001; Franceschini  et al.  2003; Georgakakis et al. 2004; Alexander  et al.  2005),  the frequency of these  systems among the obscured  X-ray population at  $z\\approx1$ is not yet  clear.  For example, Rovilos  et al.  (2007) argue that  a non-negligible fraction of X-ray  selected AGN at $z  \\approx 1$ have  $\\mu$Jy radio continuum emission (1.4\\,GHz) that is consistent with star-formation in the host galaxy.  These  authors however,  do not find  a strong  trend between X-ray obscuration and the incidence of faint radio emission, likely to be associated with starbursts.  A plausible interpretation of the evidence above is that current X-ray surveys are  biased against the early  phase of BH  evolution. At this initial  stage of  their formation,  AGN  may be  deeply buried  under star-forming clouds and/or low luminosity because the BH mass is still small, although growing rapidly.  If there is such a population of obscured  and/or intrinsically  faint  young AGN  below the  detection threshold of  current surveys, it may  be possible to  find them using stacking  analysis.  Alternatively,  AGN  may not  be responsible  for quenching the star-formation in  galaxies and producing the bimodality of the  CMD. A key test  of this scenario  is the incidence of  AGN in post-starburst  galaxies,  i.e.  systems  observed  shortly after  the termination  of  their   star-formation  ($<1$\\,Gyr).   In  the  local Universe  ($z  \\lesssim  0.1$),  a  number  of  studies  point  to  an association between  {\\it optically} selected  AGN and post-starbursts (e.g.  Kauffmann  et al. 2004;  Goto 2006; Yan  et al.  2006).   At $z \\approx  1$, close  to the  peak of  the AGN  density in  the Universe (e.g. Barger et  al. 2005; Hasinger et al. 2005),  there is still very limited  information  on  the  link  between X-ray  selected  AGN  and post-starbursts.  We  address  the  issues  above  using data  from  the  All-wavelength Extended Groth strip International Survey (AEGIS; Davis et al.  2007). X-ray  stacking  is  employed  to  search  for  deeply  buried  and/or low-luminosity AGN  among optically selected galaxies  in the redshift interval  $0.4<z<0.9$   to  explore  the  role  of   BH  accretion  in establishing  the bimodal  colour distribution  of these  systems.  We also study the X-ray  properties of post-starbursts as identified from Keck spectra at $z\\approx 0.8$  to investigate the link between recent star-formation events and active BHs.  We  adopt $\\rm H_{0} = 70 \\, km \\,   s^{-1}  \\,   Mpc^{-1}$,   $\\rm  \\Omega_{M}   =   0.3$  and   $\\rm \\Omega_{\\Lambda} = 0.7$.   \\begin{figure} \\begin{center} \\rotatebox{0}{\\includegraphics[height=0.9\\columnwidth]{cmd.eps}} \\end{center} \\caption{ Rest-frame $U -  B$ colour against B-band absolute magnitude for DEEP2  galaxies (small  blue dots) in  the range  $0.4 < z  < 0.9$ estimated by Willmer et  al. (2006). Post-starburst galaxies are shown with the (red)  triangles.  The continuous line is  defined by Willmer et   al.  (2006)   to  separate   the  red   from  the   blue  clouds. }\\label{fig_cmd} \\end{figure}   \\begin{figure*} \\begin{center} \\rotatebox{0}{\\includegraphics[height=0.7\\columnwidth]{red_24.eps} \\includegraphics[height=0.7\\columnwidth]{blue_24.eps} \\includegraphics[height=0.7\\columnwidth]{sed_24.eps}} \\end{center} \\caption{Mid-IR  colour-colour  plot.  For  clarity we  use  different panels.  Left: red  cloud galaxies  in the  redshift  range $0.4-0.9$. Middle: blue cloud  galaxies in the same redshift  interval. In both panels we only  plot sources detected in all  4 IRAC/MIPS bands used to  estimate colours.  Right:  expected tracks  of different  galaxy types in the range $z=0-1$ using observed mid-IR SEDs from the SINGS described by  Dale et al.  (2007).   NGC\\,1097 Sy1:  long-dashed dotted; NGC\\,5408 Im:  short-dashed dotted; Sc spiral  galaxy NGC\\,628: dotted;  Sab spiral  NGC\\,4450: continuous; elliptical galaxy NGC\\,584: long  dashed. The dots mark the position of $z=0$ and the crosses correspond to $z=0.4$. }\\label{fig_lacy24} \\end{figure*} ",
        "conclusions": "We  explore  the  role  of  AGN in  establishing  the  bimodal  colour distribution  of  galaxies by  quenching  the  star-formation of  blue star-forming  systems  causing their  transformation  to red,  evolved galaxies.  The main conclusions from this study are summarised below  1. There  is  evidence  for  AGN  activity  among  galaxies  in  the transition zone between  the red and blue clouds of  the CMD.  A large fraction of these  accreting BHs are obscured, suggesting  that if AGN outflows  are related  to the  colour transformation  of  galaxies, at least some nuclear gas and dust  gas clouds are either not affected or can efficiently reform after the truncation of the star-formation.  2. Morphological analysis suggests that  major mergers do not dominate the  evolution of  this population.   Minor interactions  however, may play a role.  3. We find an association  between BH accretion and post-starbursts at $z \\approx  0.8$, in agreement with  studies on the  properties of AGN hosts at low redshift, $z \\approx 0.1$.  4.  AGN  activity outlives  the termination  of the  star-formation. A large fraction of  active BHs are present in  red cloud galaxies. This is in contrast  to optically selected AGN, which  lie predominantly in the valley between the two clouds."
    },
    "0801/0801.2356_arXiv.txt": {
        "abstract": "We present some preliminary results from new {\\it Chandra} and {\\it XMM-Newton} X-ray observations of the nearby barred spiral galaxy NGC~1672. It shows dramatic nuclear and extra-nuclear star formation activity, including starburst regions located near each end of its strong bar, both of which host ultraluminous X-ray sources (ULXs). With the new high-spatial-resolution {\\it Chandra} imaging, we show for the first time that NGC~1672 possesses a faint ($L_X\\sim10^{39}$\\,erg s$^{-1}$), hard central X-ray source surrounded by an X-ray bright circumnuclear starburst ring that dominates the X-ray emission in the region. The central source may represent low-level AGN activity, or alternatively the emission from X-ray binaries associated with star-formation in the nucleus. ",
        "introduction": "The presence of a bar in a galaxy plays an important role in its evolution.  Approximately two-thirds of all spiral galaxies are barred \\citep{sheth07}, so any complete picture of galaxies must include the impact of bars on their properties.   For example, there is believed to be a causal connection between the existence of a bar and a galaxy's circumnuclear properies.  Bars can drive gas from the disc to the inner regions as gas clouds lose angular momentum as they encounter the bar and sink towards the galaxy's centre.  Indeed, a substantial fraction of galaxies with bars do show enhanced star-formation activity in their central regions, some of which are also known to have circumnuclear rings ($\\sim$20\\%, \\citealt{knapen05}).  A question that we are addressing with this study is: can the material driven to the centre of a galaxy by the bar also fuel an active galactic nucleus (AGN)?  At present, optical studies show no significant correlation between Seyfert activity and bars (e.g. \\citealt{ho97}; \\citealt{knapen05}).  However, an X-ray study by \\cite{maiolino99} shows that there is a strong correlation between absorbing column density toward Type-2 Seyfert nuclei and the presence of a strong bar: {\\it more than 80\\% of Compton thick Seyfert 2s are barred}.  This suggests that even if low-luminosity AGN activity is present, it may be obscured from view at optical wavelengths and/or diluted by strong star-formation activity, which could lead to a misclassification as a pure star-forming nucleus.  The topic of the influence of bars on galaxy evolution is not well studied at X-ray energies, but we are now able to utilize the excellent complementary imaging and spectral capabilities of {\\it Chandra} and {\\it XMM-Newton} to make detailed investigations of the X-ray characteristics of these systems.  X-ray studies offer a powerful means of probing star-formation activity in galaxies, i.e. from X-ray binaries (XRBs), supernova remnants (SNRs) and hot diffuse gas.  X-ray observations are also ideal for searching for obscured AGN activity in the form of a hard central point source that may not be visible at other wavelengths.  To address the nature of the nuclear and off-nuclear X-ray sources in barred galaxies, we have obtained new X-ray observations of NGC~1672 with both {\\it Chandra} (40\\,ks) and {\\it XMM-Newton} (50\\,ks).   Here we show some preliminary results from these observations; a full analysis will be presented in Jenkins et al. (2008) {\\it in prep}.    \\begin{figure*} \\begin{center} \\scalebox{0.8}{\\includegraphics{fig1_compressed.eps}} \\caption{{\\it Chandra} ACIS-S (left) and {\\it XMM-Newton} EPIC (right) 3-color images of NGC~1672. Red=0.3--1\\,keV, Green=1--2\\,keV \\& Blue=2--10 keV emission. North is up, east to the left.} \\label{fig:rgb} \\end{center} \\end{figure*} ",
        "conclusions": "With high-spatial-resolution {\\it Chandra} imaging we have shown, for the first time, that NGC~1672 possesses a hard central X-ray source, with a low 2--10\\,keV luminosity of $\\sim10^{39}$ erg s$^{-1}$.   This in turn is surrounded by an X-ray bright circumnuclear star-forming ring, which dominates the X-ray emission in the central region of the galaxy.  If the faint, hard X-ray source denotes low-level AGN activity, this could be an example of a much broader trend.  With similar studies on large samples of nearby galaxies, we will address the important question of the real nature of nuclear activity in barred versus non-barred systems.  For example, if detected, what range of X-ray luminosites and absorbing columns do these sources display?  What effect does a circumnuclear ring have on the luminosity, i.e. are they all low-luminosity because the ring may be stopping the bar-driven material from reaching the nucleus (as suggested by \\citealt{ho97})?  This information will be a crucial element in assessing the overall demography of massive black holes in the nearby Universe."
    },
    "0801/0801.4371_arXiv.txt": {
        "abstract": "Both core accretion and disk instability appear to be required as formation mechanisms in order to explain the entire range of giant planets found in extrasolar planetary systems. Disk instability is based on the formation of clumps in a marginally-gravitationally unstable protoplanetary disk. These clumps can only be expected to contract and survive to become protoplanets if they are able to lose thermal energy through a combination of convection and radiative cooling. Here we present several new three dimensional, radiative hydrodynamics models of self-gravitating protoplanetary disks, where radiative transfer is handled in the flux-limited diffusion approximation. We show that while the flux-limited models lead to higher midplane temperatures than in a diffusion approximation model without the flux-limiter, the difference in temperatures does not appear to be sufficiently high to have any significant effect on the formation of self-gravitating clumps. Self-gravitating clumps form rapidly in the models both with and without the flux-limiter. These models suggest that the reason for the different outcomes of numerical models of disk instability by different groups cannot be attributed solely to the handling of radiative transfer, but rather appears to be caused by a range of numerical effects and assumptions. Given the observational imperative to have disk instability form at least some extrasolar planets, these models imply that disk instability remains as a viable giant planet formation mechanism. ",
        "introduction": "Observations of protoplanetary disks around T Tauri stars traditionally imply disk masses in the range of 0.01 to 0.1 $M_\\odot$ (Kitamura et al. 2002). However, these disk masses may well be underestimated by as much as a factor of 10 (Andrews \\& Williams 2007). In addition, young stellar objects are likely to have had even higher disk masses at ages younger than typical T Tauri stars (a few Myr), as their protostellar disks transitioned into protoplanetary disks. Combined with the need for planet formation theorists to prefer increasingly higher disk masses in order to account for the timely formation of gas giants by core accretion (e.g., Inaba et al. 2003 suggest a 0.08 $M_\\odot$ disk, while Alibert et al. 2005 considered disks as massive as 0.1 $M_\\odot$), it is becoming clear that at least some protoplanetary disks are likely to have experienced a phase of gravitational instability, which might have led to the rapid formation of gas giant planets by the disk instability mechanism (e.g., Boss 1997; Mayer et al. 2002). The absence of IR excesses in $\\sim$ 65\\% of the youngest stars observed by {\\it Spitzer} suggests that the majority of protoplanetary disks dissipate on time scales of $\\sim$ 1 Myr or less (Cieza et al. 2007). While core accretion models can be constructed that permit giant planet formation times less than 1 Myr (Chambers 2006), other assumptions can require formation times of several Myr (e.g., Inaba et al. 2003; Alibert et al. 2005).  Considering that estimates of the frequency of gas giant planets around G dwarfs with orbits inside $\\sim$ 20 AU range from $\\sim$ 20\\% to $\\sim$ 40\\%, there is a need for at least one robust formation mechanism for gas giant planets. It is important to note that {\\it both} core accretion and disk instability appear to be needed to explain the range of extrasolar planets detected to date. Core accretion would seem to be the preferred mechanism to form giant planets with very large inferred core masses. E.g., HD 149026b has been inferred to have a core mass of $\\sim$ 70 $M_\\oplus$ with a gaseous envelope of $\\sim$ 40 $M_\\oplus$ (Sato et al. 2005), though the formation of this planet is hard to explain even by core accretion (Ikoma et al. 2006). Disk instability would seem to be the preferred mechanism for forming gas giants in very low metallicity systems (e.g., the M4 pulsar planet, where the metallicity [Fe/H] = -1.5 [Sigurdsson et al. 2003], and perhaps the giant planets orbiting HD 155358 and HD 47536, both of which have [Fe/H] = -0.68 [Cochran et al. 2007]). While disk instability is somewhat insensitive to metallicity (Boss 2002), recent models have suggested that higher metallicity could aid in the formation of giant planets by disk instability (Mayer et al. 2007). Disk instability also appears to be needed to form gas giant planets around M dwarf stars (Boss 2006b), though the situation regarding formation by core accretion is unclear at present (Laughlin et al. 2004; cf., Kornet et al. 2006).  Based on all of these observations and detections, then, the main theoretical questions would seem to be whether there is a formation mechanism that can account for the {\\it majority} of extrasolar planets, and if so, which mechanism it is. In order to answer these questions, theorists have been busily examining core accretion and disk instability in increasingly greater detail. Core accretion has been subjected to considerably greater scrutiny than disk instability, given that it has been the generally accepted mechanism for giant planet formation for almost three decades, dating back to Mizuno (1980), whereas disk instability is only a decade-old as a serious alternative to core accretion (Boss 1997), and is just now beginning to be investigated sufficiently to discover its strengths and weaknesses.  Theorists studying disk instability are divided into two distinct camps, those whose numerical models support the possibility of forming giant planets by this means (e.g., Boss 1997, 2000, 2001, 2002, 2004, 2005, 2006a,b, 2007; Mayer et al. 2002, 2004, 2007), and those whose numerical models (e.g., Pickett et al. 2000; Cai et al. 2006; Boley et al. 2006, 2007a,b) or analytical arguments (e.g., Rafikov 2007) do not. The reason for this basic difference in outcomes is presently unclear, and may be a combination of many effects (Nelson 2006; Boss 2007), such as numerical spatial resolution, gravitational potential solver accuracy, use of artificial viscosity, degree of stellar irradiation, detailed radiative transfer effects, and spurious numerical heating.  Recently attention has been focused on the role of radiative losses from the surface of the disk. A disk instability is likely to be stifled if the optically-thick clumps that form are unable to lose at least some of the thermal energy produced by compressional heating during contraction to protoplanetary densities. While vertical convection appears to be able to cool the disk midplane (Boss 2004; Boley et al. 2006; Mayer et al. 2007; Rafikov 2007), this thermal energy must eventually be radiated away at the disk's surface. Models employing the flux-limited diffusion approximation have been presented by Boley et al. (2006, 2007b) and Mayer et al. (2007), reaching opposite conclusions regarding the possibility of disk instability forming protoplanets. The treatment of radiative boundary conditions for the disk differs for each group. Boley et al. (2006) fit an atmosphere to the flux originating from the interior of the disk. Mayer et al. (2007) assume blackbody emission at the disk surface (for particles defined as being on the surface), while the present models use an envelope bath with a fixed temperature, typically 50 K.  With the exception of a single test model mentioned in passing by Boss (2001), all of the author's disk instability models since Boss (2001) have employed diffusion approximation radiative transfer {\\it without} a flux-limiter, for reasons of computational performance. We present here three new models that explore in some detail the effects of including a flux-limiter in disk instability models, in the hopes of helping to decide if this particular numerical choice is responsible for the distinct disparity in outcomes of disk instability models. ",
        "conclusions": "The results presented here confirm the statement made by Boss (2001) that the inclusion of a flux-limiter in these calculations does not lead to significantly different outcomes for the progress of a disk instability calculation. Even with the steeper vertical temperature gradient near the disk surface when a flux-limiter was employed (Figures 7 and 8), the corresponding midplane temperature increased by no more than 2\\%. Similarly, Boss (2007) investigated the effects of several other changes in the treatment of radiative transfer in these models, finding that the numerical assumption that had the largest effect was the relaxation of the monotonically declining vertical temperature profile, which resulted in clumps that were no more than a factor of 2 times less dense than when monotonicity was enforced. For comparison, for models FL1 and TE in Figure 7 and 8, the maximum clump densities differed by only 25\\%, implying even less of a difference between models FL1 and TE and the two models (H and TZ) from Boss (2007).  Evidently disk instability is tolerant of a range of treatments of the radiative transfer, at least up to a point. If there is a means for a clump to cool enough to contract, the clump will find this means to allow its survival. In this context it is of interest to note that analytical evaluations of disk instability (e.g., Rafikov 2007) have been restricted to considering plane parallel (one dimensional) disk models, where the entire disk midplane must be cooled, in order to cool the disk midplane anywhere at all. In a more realistic three dimensional disk model, of the sort depicted in the present numerical models, only the limited midplane region inside the dense clump needs to lose thermal energy, in any direction, in order for the clump to continue to contract and possibly survive to become a gas giant planet. This is a considerably relaxed criterion for cooling and ultimate clump survival compared to the cooling of an entire slab of midplane gas and dust. Similarly, a hot spot on the disk surface above a contracting clump will find it easier to radiate away its thermal energy than if the entire disk surface has the same vertical thermal profile as that under the hot spot.  Given the apparent observational need for disk instability to be able to form gas giant planets in some protostellar environments, if flux-limiters and other radiative transfer effects are not the main reason for the discrepant outcomes in models of disk instability, then there must be other reasons, or combinations of reasons, for these differences, as examined and discussed in some detail by Nelson (2006) and Boss (2007). Spurious heating of the inner disk associated with numerical oscillations is one possible source of these discrepancies that deserves further scrutiny (Boley et al. 2006, 2007b), as this leads to gravitational stability in the same region of the disk where clumps form in other disk instability models (Boss 2007).  Because of the unsatisfactory nature of the theoretical understanding of disk instabilities at present, it is important to continue to undertake code tests. The present models have shown that the use of a flux-limiter has relatively little effect on the evolution of an instability during the phase when the disk is already dynamically unstable. However, it is also important to investigate the role of a flux-limiter during earlier phases of evolution, before the disk becomes unstable, in order to learn if the flux-limiter can affect clump formation if applied from the very beginning of the evolution. A new model is underway that investigates this possibility, and the results will be presented in a future paper. Other code tests should also be sought, similar to the radiative transfer tests advanced by Boley et al. (2006, 2007b), except for spherical geometry instead of slab geometry, so that the present code can be tested in a similar manner."
    },
    "0801/0801.1624_arXiv.txt": {
        "abstract": "Present and planned dark matter detection experiments search for WIMP-induced nuclear recoils in poorly known background conditions.  In this environment, the maximum gap statistical method provides a way of setting more sensitive cross section upper limits by incorporating known signal information. We give a recipe for the numerical calculation of upper limits for planned directional dark matter detection experiments, that will measure both recoil energy and angle, based on the gaps between events in two-dimensional phase space. ",
        "introduction": "Introduction} Dark matter comprises approximately 25\\% of the mass of the universe~\\cite{r:spergel2003}.  The generic dark matter candidate is a weakly interacting massive particle (WIMP).  If WIMPs are supersymmetric particles, the predicted mass is in the range of 10 to 10$^4$ GeV/c$^2$, and the expected cross section lies in the range of $10^{-42}$ to $10^{-48}$ cm$^2$~\\cite{r:ellis2005}. Many experiments seek to detect dark matter particles via their elastic scattering interactions with detector nuclei~\\cite{r:gaitskell2004}.  Recent measurements~\\cite{r:xenon2007,r:cdms2005} limit the cross section to be less than approximately $5\\times10^{-44}$ cm$^2$. Given the small size of the expected WIMP cross section, discrimination against backgrounds is of paramount importance in direct detection experiments.  For the same reason, there is much to gain by optimizing the statistical methods used to interpret experimental data as upper limits on the WIMP interaction cross section~\\cite{r:yellin}.  In this paper, we develop a new statistic, the maximum patch, for setting limits on the WIMP-nucleus interaction cross section.  This method is motivated by directional dark matter experiments, which seek to measure both the nuclear recoil energy, and the recoil angle of the struck nucleon in WIMP-nucleon interactions.  In section \\ref{sec:overview} we introduce the theoretical distributions used for setting limits, and in sections \\ref{sec:1dstatistics} and \\ref{sec:2dstatistics} we discuss limit setting techniques within the context of one- and two-dimensional WIMP detection experiments.  In section \\ref{sec:comparison} we compare results in the cases of (i) observing a signal with no background, (ii) observing some signal and some background, and (iii) observing only background. ",
        "conclusions": "Conclusions}  In this paper, we have developed a new method for setting upper limits in two dimensions. The motivation for our maximum patch method is directional dark matter detection, but it is generally applicable to any two-dimensional data sets for which the distribution of the signal is known.  The approach is an extension of the one-dimensional maximum gap method~\\cite{r:yellin}, a statistic often used to set limits on the WIMP-nucleon interaction cross section in direct detection dark matter experiments.  To directly detect dark matter requires unprecedented control and understanding of backgrounds.  The great advantage of the maximum gap and patch methods is that they require no knowlege of the background distribution to set conservative upper limits on the WIMP nucleon scattering cross section.  The scattering kinematics of a dark matter signal in one and two-dimensional direct detection experiments are relatively simple.  This information is included in a straightforward way in the maximum gap and patch methods.  In particular, the maximum patch method incorporates information about the large expected angular anisotropy of dark matter scattering into the limit setting procedure.  We demonstrate that for simplistic background assumptions, the maximum patch method and two-dimensional dark matter detection yield a large gain in sensitivity, especially at low WIMP masses, over one-dimensional dark matter detection."
    },
    "0801/0801.0764_arXiv.txt": {
        "abstract": "We use Vector Spectromagnetograph (VSM) chromospheric full-disk magnetograms, from the Synoptic Optical Long-term Investigations of the Sun (SOLIS) project, to study the distribution of magnetic field flux concentrations within the polar caps. We find that magnetic flux elements preferentially appear toward lower latitudes within the polar caps away from the poles. This has implications on numerous solar phenomena such as the formation and evolution of fine polar coronal structures (i.e., polar plumes). Our results also have implications for the processes carrying the magnetic flux from low to high latitudes (e.g., meridional circulation). ",
        "introduction": "Near minima of the activity cycle of the Sun, solar poles are characterized by unipolar magnetic dominated areas, in particular at photospheric heights (Babcock \\& Livingston 1958; Babcock 1959; Howard 1972; Timothy, Krieger, \\& Vaiana 1975; Harvey, Harvey, \\& Sheeley 1982; Varsik, Wilson, \\& Li 1999). The so-called ``polar caps'' are quite prominent features due to their large sizes and their continuous presence through the solar cycle. They are the counterpart of the coronal polar holes where the fast solar wind streams (Schwenn, Muelhaeuser, \\& Rosenbauer 1979) and plasma heating (Schatzman 1949; Osterbrock 1961) originate. The physical processes laying behind these phenomena are still far from being completely understood. A significant portion of the mostly unipolar magnetic flux opens increasingly with height and fills most the heliosphere that contributes to the formation of the heliospheric magnetic field.  The solar polar caps are not sufficiently studied for different reasons (observational and modeling). In addition to the important projection effects due to the solar geometry and the lack of off-ecliptic-plane observational facilities of the magnetic field, the latter is intrinsically weak close to the solar poles (Grigor'ev 1988; Bogart, Hoeksema, \\& Scherrer 1992; Varsik, Wilson, \\& Li 1999; Varsik et al. 2002). In addition, the solar poles are only partially observable from the ecliptic plane due to $B_0$. However, instrumental limitations in matter of sensitivity  are the major problem for accurately measuring and studying the polar fields.  Complex dynamics in the solar upper convection zone together with other dynamical phenomena (e.g., differential rotation), magnetic fields of decaying active regions drift monotonously to form the large-scale field observed on the solar surface, in particular the polar caps (see Leighton 1964; DeVore \\& Sheeley 1987; Sheeley, Nash, \\& Wang 1987; Bilenko 2002). The same processes contributes significantly to the so-called polar magnetic reversal (Babcock \\& Babcock 1955) that is the change of the dominant magnetic polarity to the opposite from a solar cycle to the following one (Bilenko 2002). Under the effect of solar differential rotation, supergranular diffusion and meridional circulation, the poleward migration of the active region trailing polarities (in relative excess due to the diffusion of the leading polarity of emerging bipolar regions near solar minimum across the solar equator as suggested by Wang, Nash, \\& Sheeley 1989) has been suggested to be responsible of the reversal of the magnetic polarity of the solar poles (e.g., Babcock \\& Babcock 1955; Babcock 1961; Leighton 1964). Meridional flows has been observed and identified as the main process of flux transport from low to upper latitudes (Howard 1974; Duvall 1979; Cameron \\& Hopkins 1998; Snodgrass \\& Dailey 1996; Durrant, Turner, \\& Wilson 2004) through helioseismology and magnetic tracers (Duvall 1979; LaBonte \\& Howard 1982; Ulrich et al. 1988; Howard \\& LaBonte 1981; Topka et al. 1982; Wang, Nash, \\& Sheeley 1989). The weak flow, of few times 10~m~s$^{-1}$ at best, can be measured only to mid solar latitudes since the techniques used fails above roughly $45^\\circ-50^\\circ$. However, it is of capital importance to obtain additional constraints on these flows up to high latitudes close to the solar poles. This would be of great use for models dealing with flux transport and the solar dynamo.  Raouafi, Harvey, \\& Solanki (2006a,b; 2007) studied the plasma dynamic properties in the coronal polar plumes by means of forward modeling. They found that modeled spectroscopic emissions from polar coronal plumes rooted close to the solar poles do not match the observed coronal spectra by the UltraViolet Coronagraph Spectrometer (UVCS; Kohl et al. 1995) on board the Solar and Heliospheric Observatory (SOHO; Domingo, Fleck, \\& Poland 1995). They speculated that polar plumes would preferentially be based within the polar holes about 10$^\\circ$~away from the solar poles. It is well known from previous studies that polar plumes are the product of magnetic reconnection of a relatively large unipolar flux element with an emerging opposite polarity. Thus, the study of the polar magnetic flux distribution would yield insights and clues about the distribution of the coronal structures, such as polar plumes. In the present paper, we confirm Raouafi, Harvey, \\& Solanki's results through the study of the polar flux distribution as a function of latitude and suggest the probable effects on solar phenomena that are at the origin of such a distribution.   Saito (1958) used eclipse observations and found a similar distribution of the so-called ``polar rays''. ",
        "conclusions": "We studied the distribution of magnetic flux elements as a function of latitude near solar minimum of activity utilizing LOS-magnetograms from the SOLIS-VSM magnetograph. Chromospheric magnetograms are preferred to photospheric ones because of the better signal in the LOS magnetic component due the canopy structure of the magnetic field. Additional effects, such as the seething in the horizontal photospheric magnetic fields, are avoided by using chromospheric magnetograms. The magnetic data is recorded almost daily from September through December 2006. The north pole data was chosen for two reasons: \\begin{itemize} \\item[$\\bullet$] the north polar cap was better developed than the southern one during this period of time; \\item[$\\bullet$] the north pole was more visible than the southern one ($B_0>0$). \\end{itemize}  A recognition method of closed structures has been used to select magnetic flux elements. Thresholds on size and shape of the elements to be selected are tunable parameters of the method. The average spherical coordinates (latitude and longitude) of the selected elements are computed and their daily latitudinal distributions are determined. In order to remove any effect of the surface area of the polar cap, the histograms are normalized with respect to the area corresponding to each latitude bin taking into account the different solar geometrical parameters for the each magnetogram. Since daily distributions do not reflect the real density of magnetic elements for statistical reasons, monthly averages of the normalized histograms are then computed.  The density of the unsigned magnetic flux elements within the polar magnetic cap (normalized by the surface area of the polar cap) is found to be strongly dependent on solar latitude. Flux elements are found relatively uniformly dense up to latitudes of about $75^\\circ-80^\\circ$ followed by a significant decrease in their latitudinal distribution at higher latitudes close to the solar pole. Although, a difference in the distribution of relatively small and large flux elements is found, the density decreasing behavior toward the pole is found in all the data considered (for details see Raouafi, Harvey, \\& Henney 2007). We believe that the mechanisms responsible for the transport of the magnetic flux of low latitude decaying active regions toward high polar latitudes are less efficient close to the poles. This means that the meridional circulation responsible for the flux transport slows before reaching the poles. Such a result would have a significant impact on the theories and models dealing with flux transport. These results also put important constraints on solar phenomena that are inaccurately (if at all) determined by other means, such as helioseismic studies that become inefficient at high latitudes."
    },
    "0801/0801.0287_arXiv.txt": {
        "abstract": "This work tabulates measured and derived values of coefficients for Lorentz and CPT violation in the Standard-Model Extension. Summary tables are extracted listing maximal attained sensitivities in the matter, photon, and gravity sectors. Tables presenting definitions and properties are also compiled. ",
        "introduction": "\\label{Introduction}  Recent years have seen a renewed interest in experimental tests of Lorentz and CPT symmetry. Observable signals of Lorentz and CPT violation can be described in a model-independent way using effective field theory \\cite{kp}.  The general realistic effective field theory for Lorentz violation is called the Standard-Model Extension (SME) \\cite{dcak,akgravity}. It includes the Standard Model coupled to General Relativity along with all possible operators for Lorentz violation. Both global \\cite{dcak} and local \\cite{akgravity} Lorentz violation are incorporated. Since CPT violation in realistic field theories is accompanied by Lorentz violation \\cite{owg}, the SME also describes general CPT violation. Reviews of the SME can be found in Refs.\\ \\cite{cptprocs,rb}.  Each Lorentz-violating term in the Lagrange density of the SME is constructed as the coordinate-independent product of a coefficient for Lorentz violation with a Lorentz-violating operator. The Lorentz-violating physics associated with any operator is therefore controlled by the corresponding coefficient, and so any experimental signal for Lorentz violation can be expressed in terms of one or more of these coefficients.  The Lorentz-violating operators in the SME are systematically classified according to their mass dimension, and operators of arbitrarily large dimension can appear. At any fixed dimension, the operators are finite in number and can in principle be enumerated. A limiting case of particular interest is the minimal SME, which can be viewed as the restriction of the SME to include only Lorentz-violating operators of mass dimension four or less. The corresponding coefficients for Lorentz violation are dimensionless or have positive mass dimension.  The results summarized here concern primarily but not exclusively the coefficients for Lorentz violation in the minimal SME. We compile data tables for these SME coefficients, including both existing experimental measurements and some theory-derived limits, and we provide tables listing some relevant definitions and properties. We also extract summary tables listing our best estimates for the maximal attained sensitivities in three sectors: ordinary matter (electrons, protons, and neutrons), photons, and gravity. The tables include results available from the literature up to \\papersversdate. More recent papers can be found online \\cite{akweb}.  The order of the tables is as follows. Table \\ref{contents} contains a list of all tables. The three summary tables are presented next, Tables \\ref{sum:matter}, \\ref{sum:photon}, and \\ref{sum:gravity}. These are followed by the data tables, Tables \\ref{dat:electr} to \\ref{dat:nonmin}. The properties tables appear last, Tables \\ref{QEDlagrangian} to \\ref{spherical_harm}.  A description of the summary tables is given in Sec.\\ II. Information about the format and content of the data tables is presented in Sec.\\ III, while Sec.\\ IV provides an overview of the properties tables. The bibliography for the text and all the tables follows Sec.\\ IV. ",
        "conclusions": ""
    },
    "0801/0801.1142_arXiv.txt": {
        "abstract": "We present a detailed analysis of archival \\textit{Hubble Space Telescope} data that we use to measure the proper motion of the Crab pulsar, with the primary goal of comparing the direction of its proper motion with the projected axis of its pulsar wind nebula (the projected spin axis of the pulsar).  Combining data from 47 observations spanning $>10\\,$yr with two different instruments, and using the best available measurement techniques and latest distortion models, we are able to demonstrate that our measurement is robust and has an uncertainty of only $\\pm0.4\\,\\masyr$ on each component of the proper motion.  However, we then consider the various uncertainties that arise from the need to correct the proper motion that we measure to the local standard of rest at the position of the pulsar and find $\\mu_\\alpha = -11.8 \\pm 0.4 \\pm 0.5\\,\\masyr$ and $\\mu_\\delta = \\mbox{+4.4} \\pm 0.4 \\pm 0.5\\,\\masyr$ relative to the pulsar's standard of rest, where the two uncertainties are from the measurement and the reference frame, respectively.  If we then wish to compare this proper motion to the symmetry axis of the pulsar wind nebula, we must consider the unknown velocity of the pulsar's progenitor (assumed to be $\\sim 10\\,\\kms$), and hence add an additional uncertainty of $\\pm2\\,\\masyr$ to each component of the proper motion, although this could be a factor of 10 larger if the pulsar's progenitor had an anomalously high velocity ($>100\\,\\kms$).  This implies a projected misalignment with the nebular axis of $14\\degr\\pm2\\degr\\pm9\\degr$, consistent with a broad range of values including perfect alignment. We use our proper motion to derive an independent estimate for the site of the supernova explosion with an accuracy that is 2--3 times better than previous estimates.  We conclude that the precision of individual measurements which compare the direction of motion of a neutron star to a fixed axis will often be limited by fundamental uncertainties regarding reference frames and progenitor properties. The question of spin-kick (mis)alignment, and its implications for asymmetries and other processes during supernova core-collapse, is best approached by considering a statistical ensemble of such measurements, rather than detailed studies of individual sources. ",
        "introduction": "The high velocity nature of the neutron star population has been apparent almost as long as the existence of neutron stars has been recognized \\citep{go70}. Recent statistical studies of the radio pulsar velocity distribution \\citep[e.g.,][]{acc02,bfg+03,hllk05,fgk06} yield mean three-dimensional velocities of 300--500~\\kms, with a high velocity tail extending beyond 1000~\\kms.  A variety of physical mechanisms have been proposed as the origin of high velocities.  Perhaps the first was disruption of binaries through mass loss in supernovae (\\citealt{b61}; \\citealt*{ggo70}; \\citealt{it96}), although it is difficult for binary disruption alone to account for some of the highest observed velocities (\\citealt*{hla93}; \\citealt{cvb+05}). The most natural source of such high velocities appears to be asymmetries in the birth supernovae of pulsars \\citep{s70,vdhvp97,pzvdh99}, although exactly how an asymmetry in the core collapse process in a supernova explosion is converted to a birth kick imparted to a nascent neutron star remains unclear \\citep*{lcc01}.  While hydrodynamic or convective instabilities are the most plausible route \\citep[e.g.,][]{bh96,jm96,lg00,spj+04,jsk+05,skjm06,bld+06}, more exotic mechanisms such as asymmetric neutrino emission in the presence of strong magnetic fields \\citep{al99b} or some combination of the above \\citep{sbhf05} cannot be ruled out.  Of these kick mechanisms, many predict kicks vectors that relate to the spin axis orientation of the nascent neutron star: the alignment (or lack thereof) of the natal kick with the neutron star spin axis could provide a specific discriminant between various mechanisms (\\citealt*{bhf95}; \\citealt{sp98,cowsik98,lcc01,romani05}).  Even such parameters such as the number and timescale of kick components, coupled with the initial spin period of the neutron star, can be constrained through observations of an ensemble of sources (\\citealt*{drr99}; \\citealt{jhv+06,wlh06,wlh07,nr07,rankin07}).   \\subsection{The Crab Pulsar and its Nebula} The Crab pulsar (\\object[PSR B0531+21]{PSR~B0531+21}) and nebula have been observed by virtually every telescope capable of pointing at the system, and the Crab may be the most studied system in all of astronomy.  The Crab nebula possesses a general symmetry axis, visible in images at most wavelengths.  Recent observations have delineated this axis with striking clarity \\citep[e.g.,][]{hmb+02,nr04}.  The X-ray jet, in particular, allows us to trace the symmetry axis to the neutron star location, and provides a natural association with the rotation axis of the pulsar itself (since  every other vector would be rotation averaged).  The symmetry axis is roughly aligned with the proper motion \\citep[e.g.,][]{cm99}, but the observational uncertainties have made the alignment hard to quantify.  A precise proper motion vector for the Crab pulsar could be quantitatively compared to the jet direction to establish whether the natal kick is aligned with the spin axis, as many theories predict.   Given its prominent place in our understanding of neutron stars, it is perhaps surprising that the proper motion and distance of the Crab pulsar are not better known.  Compared to many fainter objects, the precision of our measurements is lacking.  While there were a number of early attempts to measure the proper motion of the Crab pulsar, these were generally inconsistent with each other \\citep{mink70}.  Perhaps the first reliable measurement was that of \\citet[][hereafter \\citetalias{wm77}]{wm77}, who found\\footnote{The analysis of \\citetalias{wm77} was based on the B1950 frame, but at this level of precision the precession between that and the J2000 frame does not change the proper motion significantly.}  $(\\mu_\\alpha,\\mu_\\delta) = (-13\\pm2,7\\pm3)\\,\\masyr$ from photographic plates spanning epochs from 1899 to 1976.  There have not been any direct (i.e.\\ geometric) distance measurements of the pulsar itself, but the distance was estimated based on various lines of evidence to lie between 1.4 and 2.7$\\,$kpc \\citep{t73}. In spite of the wealth of observations since then, including a treasure trove of \\textit{Hubble Space Telescope} (\\hst) images with a time baseline of $>10$ years, these estimates had not significantly improved until recently. For example, \\citet[][hereafter \\citetalias{cm99}]{cm99} estimate a proper motion $(\\mu_\\alpha,\\mu_\\delta) = (-17\\pm3,7\\pm3)\\,\\masyr$, which is consistent with the earlier estimate, but does not improve upon its accuracy.  The main obstacle to more accurate measurements is not (as in many cases) the limitations of faint objects, but rather the fact that the Crab nebula is too bright for most interferometric radio observations (it raises the system temperature too much), that there are no suitable interferometric calibrators nearby, and that the rotational stability is not sufficient (due to glitches) for precise pulse timing over a long time baseline.  Because of these reasons, high angular resolution optical observations are so far the only way to measure the astrometric parameters (proper motion and parallax) of the Crab pulsar, and with the current generation of instruments we are limited to data from \\hst.   \\citet[][hereafter \\citetalias{nr06}]{nr06} attempted a significantly more detailed astrometric analysis of the Crab pulsar compared to \\citetalias{cm99}, taking advantage of new \\hst\\ data spanning 7 years and trying to account for many sources of uncertainty not addressed by \\citetalias{cm99}.  \\citetalias{nr06} found a result that is discrepant with that of \\citetalias{cm99}: $(\\mu_\\alpha,\\mu_\\delta) =(-15.0\\pm0.8,1.3\\pm0.8)\\,\\masyr$.  This shows a significant misalignment with the projected spin-axis of the pulsar ($26\\degr\\pm3\\degr$) and as such reverses previously held notions of spin-kick alignment, but as we discuss below (and as \\citetalias{nr06} acknowledge) even this analysis still is not as accurate as possible. Additionally, a large number of new observations with the Advanced Camera for Surveys (ACS) have become publicly available.  These were taken primarily for studying the dynamics and polarization of the Crab's pulsar wind nebula \\citep[e.g.,][]{hmb+02}, and as such they were not ideal for astrometry (the exposure times were long enough that the pulsar saturated, the dithering strategy was not optimal, and they used a limited range of roll angle), but nonetheless they are an important resource.  Motivated by the importance of the Crab pulsar in our understanding of neutron stars and supernova remnants in general, by the large amount of data available on it, and by the limitations of previous analyses, we have attempted to re-measure the proper motion of the Crab pulsar as well as assess the possibility of a parallax measurement with future data (although given the recent failure of the ACS instrument, future observations may not be possible until the installation of the upcoming Wide Field Camera 3).    The organization of our paper is as follows: in \\S~\\ref{sec:an} we describe our analysis, noting departures from previous analyses, although the majority of the fitting techniques are similar to those we used in \\citet*{kvka07}, and we refer readers there for more details.  After refining the proper motion measurement, we discuss in \\S~\\ref{sec:ref} the limitations on our knowledge of the proper motion imposed by the unknown velocity of the pulsar's progenitor, as well as uncertainties in the corrections to the pulsar's local standard of rest.  These transformations and their associated uncertainties limit the accuracy of the comparison between the proper motion and the projected spin-axis of the pulsar.  We then give our conclusions in \\S~\\ref{sec:conc}.  Finally, we include a discussion of the prospects for a parallax distance for the Crab pulsar in Appendix~\\ref{sec:par}. In what follows, we define our proper motions in Right Ascension and Declination $(\\mu_\\alpha,\\mu_\\delta)$ so that the scales are the same and no $\\cos\\delta$ term is necessary.  All uncertainties are 1-$\\sigma$ unless otherwise stated, and all position angles are measured east of north. ",
        "conclusions": "\\label{sec:conc}  While we assumed a velocity of $\\pm10\\,\\kms$ for the Crab pulsar's progenitor, it is possible that the progenitor itself had a much larger space velocity.  This was proposed early on by \\citet{mink70}, who considered the high proper motion of the pulsar to be a relic of the progenitor's velocity and that the progenitor itself was a runaway star (e.g., \\citealt{b61}) from the Gem~OB1 association. However, \\citet{mink70} later dismissed this hypothesis since he did not believe that the supernova was a type~II explosion.  \\citet{ggo70} have a similar idea, where they consider the Crab pulsar and the nearby pulsar B0525+21 (two of the first pulsars to be discovered; \\citealt{sr68}) as former binary companions ejected from the Gem~OB1 association; \\citet{hla93} attempted to measure the proper motion of PSR~B0525+21 but did not find a statistically significant result. Later analyses, such as \\citet{pols94} and \\citet{md01}, have revived the hypothesis that the progenitor had a large ($>100\\,\\kms$) space velocity, with \\citet{pols94} arguing that the large height below the Galactic plane (200$\\,$pc for a distance of 2$\\,$kpc, which is larger than the scale-height of OB stars; \\citealt{reed00}; \\citealt*{ecca06}) and the presumed evolutionary state of the progenitor (inferred from the elemental and velocity structure of the Crab Nebula) suggest that the Crab pulsar was formed from the second explosion in a binary system, and that it had a large space velocity from the first explosion.  Currently we cannot say definitively whether or not the progenitor had a significant space velocity and must therefore treat this as an overall uncertainty on our whole analysis.   \\subsection{Location of the Explosion Center} A number of authors have estimated the ``divergent point'' for the Crab Nebula: the point from which all of the filaments seem to traveling outwards, which is presumed to be the center of the explosion.  Among the more reliable measurements are those of \\citet{trimble68}, \\citetalias{wm77}, and \\citet[][who also review the situation]{nugent98}, which we list in Table~\\ref{tab:exp}. Most of these determinations trace the filaments back and find best-fit dates for the explosion of $\\approx 1130$~CE instead of the commonly accepted 1054~CE \\citep{sg02}, with the difference caused by unmodeled acceleration of the filaments.   \\begin{deluxetable*}{l c c l c c} \\tablecaption{Comparison Between Our Proper Motion and Divergent Points\\label{tab:exp}} \\tablewidth{0pt} \\tabletypesize{\\footnotesize} \\tablehead{ \\colhead{Reference} & \\mc{2}{c}{Divergent Point (J2000)\\tablenotemark{a}} & \\colhead{Divg.\\ Date\\tablenotemark{b}} & \\colhead{$\\Delta r_{\\rm min}$\\tablenotemark{c}} & \\colhead{$\\Delta r$(1054 CE)\\tablenotemark{d}}\\\\ \\cline{2-3} & \\colhead{$\\alpha$} & \\colhead{$\\delta$} & \\colhead{(CE)} & \\colhead{(arcsec)} &\\colhead{(arcsec)} \\\\ } \\startdata \\citet{trimble68} & $05^{\\rm h}34^{\\rm m}32\\fs72\\pm0\\fs12$ & $+22\\degr00\\arcmin47\\farcs5\\pm1\\farcs4$ & $1067\\pm138$ & $0.5$ & 0.6\\\\ \\citetalias{wm77} & $05^{\\rm h}34^{\\rm m}32\\fs67\\pm0\\fs06$ & $+22\\degr00\\arcmin47\\farcs6\\pm0\\farcs9$ & $1114\\pm78$ & 0.7 & 1.0\\\\ \\citet{nugent98} & $05^{\\rm h}34^{\\rm m}32\\fs84\\pm0\\fs12$ & $+22\\degr00\\arcmin48\\farcs0\\pm1\\farcs3$ & \\phn$947\\pm138$ & 0.6 & 1.5\\\\ \\tableline This work\\tablenotemark{e}  &$05^{\\rm h}34^{\\rm m}32\\fs74\\pm0\\fs03$ & $+22\\degr00\\arcmin47\\farcs9\\pm0\\farcs4$ & 1054 & \\nodata & \\nodata\\\\ \\enddata \\tablenotetext{a}{The positions of the divergent point according to \\citet{trimble68} and \\citet{nugent98} were computed from the offsets given in \\citet{nugent98} between the divergent point/explosion center and the star $5\\arcsec$ to the north-east of the pulsar, whose position we take to be: $\\alpha=05^{\\rm h}34^{\\rm m}32\\fs17$, $\\delta=+22\\degr00\\arcmin56\\farcs0$ from 2MASS (the star is 2MASS~J05343217+2200560).  For \\citetalias{wm77}, we take the divergent point directly from their paper (Eqn.~17).} \\tablenotetext{b}{Date of closest approach between our proper motion vector projected backwards and the divergent point.  The uncertainties are only measurement uncertainties --- no reference frame uncertainties are included.} \\tablenotetext{c}{Closest approach between our proper motion projected backward and the estimated divergent point of the filaments.} \\tablenotetext{d}{Distance between our proper motion projected backward and the divergent point for 1054~CE.} \\tablenotetext{e}{Not truly a divergent point, but rather the location of the pulsar projected back to 1054~CE.} \\end{deluxetable*}   With the proper motion that we derive, the explosion centers from the literature, and the nominal explosion date of 1054~CE, we can perform three tests: we have a time, a displacement, and a velocity, and we can use any two of those to estimate the third.  First, we can measure how close our proper motion comes to the various explosion centers for the nominal explosion date, and we give these values in the last column (``$\\Delta r$(1054~CE)'') of Table~\\ref{tab:exp}.  Second, we can compute the dates of closest approach between our projected proper motion vectors and the explosion centers, which serve as our own estimates of the explosion dates. The uncertainties on those values are dominated by the uncertainties of the divergent point measurements ($\\sim 1\\arcsec$), and the values are given in the ``Divg.~Date'' column of Table~\\ref{tab:exp}, with the approach distances for those dates given in the ``$\\Delta r_{\\rm min}$'' column.  Finally, we can compute our own estimate for the explosion position, taking our proper motion and assuming the date of 1054~CE, and we give this in the last row of Table~\\ref{tab:exp}.  In general, all of the values --- the approach distances for 1054~CE, the explosion dates, and our inferred explosion center --- are consistent at better than 1-$\\sigma$.  The first two elements are largely consistency checks: this shows that our proper motion is consistent with the independent divergent point estimates, and that our reference frame corrections are consistent (although they were not explicitly the same).  As discussed in \\citetalias{wm77}, the location of the divergent point depends on the choice of reference frame, so we cannot address any of the larger reference frame uncertainties.  We note, though, that in contrast to our proper motion that approaches the divergent points with distances of $<1\\arcsec$, the proper motion of \\citetalias{nr06} does approximately a factor of 3 worse.  The final element that we have computed, our own estimate for the explosion center, is more precise than previous estimates by a factor of $\\sim$2--3 in each axis.  This may serve to help constrain future measurements of the filament motions and acceleration, as its independence from the filaments themselves should improve the reliability of the measurements.   \\subsection{Spin-Kick MisAlignment} \\citet{nr04} fit a model to the torus seen in the \\hst\\ data, and find a best-fit torus symmetry axis of $304.0\\degr\\pm0.1\\degr$, which is the projection of the spin axis on the plane of the sky.  This implies a projected misalignment of $14\\degr\\pm2\\degr\\pm9\\degr$, as seen Figure~\\ref{fig:wfc}.  This projected misalignment is less than that found in \\citetalias{nr06}, and is significant if one only considers the measurement uncertainty: with all of the uncertainties, the misalignment is consistent with a broad range of values, including zero.  Perhaps the next best case of a pulsar wind nebula giving the projected spin axis is that of the \\object[PSR B0833-45]{Vela pulsar}, where the proper motion (corrected for Galactic rotation and solar motion) is $45\\pm1.3\\,\\masyr$ at a position angle of $301\\degr\\pm 2\\degr$ \\citep{dlrm03}.  The symmetry axis of the torus is at a position angle of $310.6\\degr\\pm0.1\\degr$ \\citep{nr04}, giving the projected misalignment of $10\\degr\\pm2\\degr$ \\citep{nr07}.  For this system, the proper motion reference frame and corrections should be better defined than those for the Crab: the proper motion is measured in the radio, so the reference sources are at infinite distance; and the Vela pulsar is reasonably close (with a distance measured through geometric parallax) and not located at the Galactic anti-center, so the effects of Galactic rotation are much better understood.  However, \\citet{nr04} still fail to include any allowance for the unknown velocity of the progenitor: a $\\pm 10\\,\\kms$ velocity at a distance of $287\\,$pc is $\\pm7\\,\\masyr$ (or an angular uncertainty of $\\sim 9\\degr$), so again this completely dominates the measurement uncertainty and makes the degree of misalignment consistent with zero.  The two examples considered here, the Crab and Vela pulsars, while the two best X-ray tori \\citep{nr04}, may not be good cases for computing misalignment.  This is because they are both moving at smaller velocities than the average pulsar population ($120\\,\\kms$ and $61\\,\\kms$, vs.\\ $\\sim400\\,\\kms$; e.g., \\citealt{hllk05,fgk06}), so the unknown velocity of the progenitor has a correspondingly greater contribution.  In fact, there is a  bias in favor of tori being found around relatively slow pulsars.  This is because the transition from a ``bubble'' pulsar wind nebula (PWN) with a torus to a bow-shock PWN \\citep[see][]{gs06} occurs when the pulsar has traveled roughly 68\\% of the distance from the center of the of the supernova remnant to its edge \\citep*{vdsdk04}, largely independent of the pulsar's velocity.  So slower pulsars will spend longer in the bubble/torus phase.  For faster moving pulsars, whose proper motions and projected spin axes are not as well determined (in general they are further away), the uncertainty will be closer to $1\\degr$ and will not dominate over the measurement uncertainties.  In addition, if the model of \\citet{nr07} is correct, we would expect the faster moving pulsars to be intrinsically closer to alignment.  We also note that, for pulsars moving at $>200\\,\\kms$, all of the reference frame uncertainties discussed here will lead to uncertainties of $<10$\\% on the space velocity: this will typically be less than the uncertainty on the distance.  Therefore, in studying the magnitude of pulsar velocities (or of the pulsar population) the reference frame uncertainties will not be significant.  However, we can also look at the situation from the other side.  If the spin and kick axes were perfectly aligned in the reference frame of the progenitor's motion, we could still (erroneously) infer a misalignment because of a high progenitor space velocity.  In practice, though, it is difficult to disentangle the effects of intrinsic and apparent misalignments for single objects.  A further complication comes from the fact that all alignments are examined only in projection: inclinations of nebulae can be estimated (although not directly measured), but radial velocities of pulsars are entirely unknown, and without the third dimension any observed alignments could still be coincidences.   A number of other pulsar wind nebulae also have symmetry axes \\citep[e.g.,][]{pksg01,hgh01,nr04}, although the lower fluxes and larger distances make most of these hard to observe in detail.  There are other situations where alignments are inferred but not measured directly: for instance, using an offset from the center of a supernova remnant to derive a kick direction (this approach can introduce substantial systematic errors of its own; see \\citealt{gcs+06}), and deriving a rotation axis from fitting radio polarization data \\citep{drr99,lcc01,rn03,wlh06,rankin07}.  For those systems with constrained proper motions and rotations axes, \\citet{nr04} find projected misalignments of $\\sim 10\\degr$, although most are consistent with zero.  \\citet{wlh06} and \\citet{nr07} find similar misalignments for a larger sample using polarization data (also see \\citealt{drr99}), but especially if we restrict the sample to the younger pulsars (where Galactic acceleration should not have modified the initial velocity) then the conclusions are similar to \\citet{nr04} (also see \\citealt{jhv+05}).  This suggests that, at least statistically, there still is a considerable degree of alignment between projected spin axes and proper motions.  From this, \\citet{wlh07} and \\citet{nr07} argue that the asymmetries experienced by proto-neutron stars following core-collapse simultaneously can impart these stars with both kick and spin, and that these asymmetries consist of a stochastic ensemble of thrusts, each long enough to result in rotational averaging of the resultant linear momentum vector. The uncertainties discussed in this paper illustrate the difficulties in measuring precision alignments (or lack thereof) in any individual object. Further progress in these studies, for example, via detailed analyses of how the degree of alignment depends on parameters such as space velocity and surface magnetic field strength, is best achieved by adding to the total number of pulsars with information on the orientations of both spin and kick."
    },
    "0801/0801.1232_arXiv.txt": {
        "abstract": "We derive new constraints on the mass of the Milky Way's dark matter halo, based on a set of halo stars from SDSS as kinematic tracers. Our sample comprises $2401$ rigorously selected Blue Horizontal-Branch (BHB) halo stars at $\\rm |z| \\ge 4$ kpc, and with distances from the Galactic center up to $\\sim 60$ kpc, with photometry and spectra drawn from SDSS DR-6. With distances accurate to $\\sim 10\\%$, this sample enables construction of the full line-of-sight velocity distribution at different Galactocentric radii. To interpret these distributions, we compare them to matched mock observations drawn from two different cosmological galaxy formation simulations designed to resemble the Milky Way, which we presume to have an appropriate orbital distribution of halo stars. Specifically, we select simulated halo stars in the same volume as the observations, and derive the distributions $\\rm P(V_{los}/V_{cir})$ of their line-of-sight velocities for different radii, normalized by the simulation's local circular velocity. We then determine which value of $\\rm V_{cir}(r)$ brings the observed distribution into agreement with the corresponding distributions from the simulations. These values are then adopted as observational estimates for $\\rm V_{cir}(r)$, after a small Jeans Equation correction is made to account for slight data/simulation differences in the radial density distribution. This procedure results in an estimate of the Milky Way's circular velocity curve to $\\sim 60$~kpc, which is found to be slightly falling from the adopted value of $\\rm 220~km~s^{-1}$ at the Sun's location, and implies M$(<60 \\rm~ kpc) = 4.0\\pm 0.7\\times 10^{11}$M$_\\odot$. The radial dependence of $\\rm V_{cir}(r)$, derived in statistically independent bins, is found to be consistent with the expectations from an NFW dark matter halo with the established stellar mass components at its center. If we assume an NFW halo profile of characteristic concentration holds, we can use the observations to estimate the virial mass of the Milky Way's dark matter halo, M$_{\\rm vir}=1.0^{+0.3}_{-0.2} \\times 10^{12}$M$_\\odot$, which is lower than many previous estimates. We have checked that the particulars of the cosmological simulations are unlikely to introduce systematics larger than the statistical uncertainties.  This estimate implies that nearly 40\\% of the baryons within the virial radius of the Milky Way's dark matter halo reside in the stellar components of our Galaxy. A value for M$_{\\rm vir}$ of only $\\sim 1\\times10^{12}$M$_\\odot$ also (re-)opens the question of whether all of the Milky Way's satellite galaxies are on bound orbits. ",
        "introduction": "The visible parts of galaxies are, in the current paradigm for galaxy formation, concentrations of baryons at the center of much larger dark matter halos, which have assembled through hierarchical merging and gas cooling. Understanding the properties of these dark matter host halos, their virial masses, concentrations, and radial mass profiles, {\\it vis-a-vis} the luminous properties of the main galaxy at their center, is crucial for modeling the dynamics of the central galaxy, for connecting observations of galaxies to large-scale cosmological dark matter simulations, and for understanding what fraction of baryons in the halo ended up as stars in the central galaxy. In turn, the extended stellar distributions of galaxies, in particular the stellar halos of nearby galaxies, offer some of the best probes to test generic predictions about the nature of dark matter mass profiles (Navarro et al. 1996).  The Milky Way and its surrounding halo are of particular interest, as our internal position permits the placement of unique constraints on the Galaxy's stellar mass content, on its dark matter profile at large radii, and on the 3-D shape of its dark matter halo. Yet, our location within the Galaxy also complicates some measurements, such as the extended rotation curve of gas in its disk. As a result, the dark mass profile for the Milky Way between $\\sim 10$ kpc and $\\sim 100$ kpc and the halo's virial mass have not been previously constrained to better than a factor of 2 to 3. In practice, it has proven useful to quantify the halo mass profile by either a circular velocity curve, $\\rm V_{cir}(r)$, or by the escape velocity curve, $\\rm V_{esc}(r)$.  In previous work, the most common tools used to estimate the Milky Way halo mass are the escape velocity and the velocity dispersion profile of the tracer populations (i.e., halo stars, or the Milky Way's satellite galaxies and globular clusters). The escape velocity provides constraints on the gravitational potential at the relevant position (Little \\& Tremaine 1987; Zaritsky et al.  1989; Kulessa \\& Lynden-Bell 1992; Kochanek 1996). Recent work has shown that the total mass of the halo is around $2 \\times10^{12} ~ {\\rm M}_\\odot$. Wilkinson \\& Evans (1999) used the velocities of 27 satellite galaxies and globular clusters to find a halo mass of $\\sim 1.9^{+3.6}_{-1.7} \\times 10^{12}~{\\rm M}_\\odot$ by adopting a truncated, flat rotation curve halo model.  Sakamoto, Chiba \\& Beers (2003) used a sample including 11 satellite galaxies, 137 globular clusters, and 413 solar neighborhood field horizontal-branch stars, along with a flat rotation curve model, to obtain a total halo mass of $2.5^{+0.5}_{-1.0} \\times 10^{12}~{\\rm M}_\\odot$ or $1.8^{+0.4}_{-0.7} \\times 10^{12}~{\\rm M}_\\odot$, depending on the inclusion (or not) of Leo I. Recently, Smith et al. (2007) estimated a halo mass of $\\sim 1.42^{+1.14}_{-0.54} \\times 10^{12}~{\\rm M}_\\odot$, based on a sample of high-velocity stars from the RAVE survey and two published databases, and an adiabatically contracted NFW halo model. Battaglia et al. (2005, 2006) used a derived velocity dispersion profile to determine a total mass of $0.5 \\sim 1.5 \\times 10^{12}~{\\rm M}_\\odot$, with some dependence on the model adopted for the halo profile (see also Dehnen, McLaughlin \\& Sachania 2006).  In order to improve the precision of the mass estimate of the halo, the fundamental first step is to begin with high-quality data (tracers with accurate distances and radial velocities), augmented with an efficient method of analysis capable of extracting the maximum amount of information. For example, Wilkinson \\& Evans (1999) comment that a data set of $\\sim 200$ radial velocities of BHB stars reduces the uncertainty in the mass estimate of the halo to $\\sim 20$\\%.  This goal has proven elusive, even with data samples twice this size (e.g. Sakamoto et al. 2003), due to the use of a relatively nearby sample of tracers.  However, we also expect a large and distant data set to constrain the gravitational potential at different positions, which is ultimately a more reliable probe of the dark matter halo. The planned or already-underway kinematically unbiased surveys, such as the European Space Agency's astrometric mission Gaia (Turon et al. 2005), the Space Interferometry Mission (SIM; Unwin et al. 2000), the RAdial Velocity Experiment (RAVE; Steinmetz et al. 2006), the Sloan Digital Sky Survey (SDSS; York et al. 2000), and the Sloan Extension for Galactic Understanding and Exploration (SEGUE; Newberg et al. 2003, Beers et al.  2004), will finally make it feasible to obtain the required large sample of tracers with well-measured parameters.  Blue Horizontal-Branch (BHB) stars are excellent tracers of Galactic halo dynamics, because they are luminous and have a nearly constant absolute magnitude within a restricted color range (see, e.g., Sirko et al. 2004 for BHB stars in SDSS). SDSS and SEGUE specifically targeted BHB stars for spectroscopy; we employ these stars (and other serendipitously discovered BHB stars with available spectra, e.g., mis-identified QSOs) to derive precision constraints on the mass of the Milky Way's dark matter halo.  In \\S 2 we describe the assembly of a particularly conservative (i.e., low contamination) selection of BHB stars from the spectra of candidate BHBs in SDSS DR-6. Two different cosmological simulations, one of a galaxy resembling the Milky Way with a $\\sim 2 \\times 10^{12}~{\\rm M}_\\sun $ halo, and one of a Milky-Way-like galaxy simulation with a $\\sim 1 \\times 10^{12}~{\\rm M}_\\sun $ halo, are presented in \\S 3. This section also describes the estimation of the circular velocity curve as a function of Galactocentric radius, $\\rm V_{cir}(r)$, obtained by comparing the observed BHB radial velocities to the kinematics of the simulations. In \\S 4 we present the resulting estimates of the Milky Way's halo circular velocity curve, and discuss implications for the virial mass of the Milky Way's halo. Our conclusions are summarized in \\S 5. ",
        "conclusions": "We have constrained the mass distribution of the Milky Way's dark matter halo, by analyzing the kinematics of nearly $2400$ BHB stars drawn from SDSS DR-6, which reach to Galactocentric distances of $\\sim \\rm 60$ kpc. To obtain a ``clean sample'' of BHB stars, we have re-analyzed all candidate BHB spectra, following the prescription by Sirko et al. 2004, which should result in a contamination fraction (mostly by BS stars) of well below 10\\%. The metallicity distribution, centered on [Fe/H] $ \\sim -2$~dex, confirms that most sample members must be halo stars. The distances to the BHB stars are known to $\\sim 10\\%$.  To account for the complex survey geometry and for plausible orbital distributions of our sample of BHB stars, we have compared the observed radial velocities to analogous quantities drawn from the ``star particles'' in galaxy formation simulations that resemble the Milky Way. In particular, we have placed a virtual ``observer'' at 8.0 kpc from the simulation centers, looking in the actual SDSS directions and sampling radial velocities for stars that are more than 4 kpc above and below the disk plane. We then explored to what mass scale, or $\\rm V_{cir}(r)$, the simulations need to be scaled to in order to match the observed line-of-sight velocity distribution in a set of radial bins. In this analysis, we adjusted this scaling using the Jeans Equation, to account for the slight difference in the radial profile of observed halo stars ($\\rm \\rho \\sim r^{-3.5}$) and simulated halo stars ($\\rm \\rho \\sim r^{-2.9}$) over this radial range.  This procedure results in direct estimates of $\\rm V_{cir}(r)$ from $\\sim 10$ kpc to $\\sim 60$ kpc, the best such estimates to date over this range. The circular velocity estimate varies slightly with radius, dropping from $\\sim 220~ \\rm km~s^{-1}$ at 10~kpc to $\\sim 170~ \\rm km~s^{-1}$ in the most distant two bins. Applying this procedure to two independent cosmological simulations (Simulation I and Simulation II, respectively) results in consistent estimates of $\\rm V_{cir}(r)$. The mass enclosed within 60~kpc, constrained quite directly by the data, is found to be $4.0\\pm 0.7\\times 10^{11}$M$_\\odot$. As a result of the much more extensive data set provided by SDSS, the uncertainties on this estimate are substantially lower than those obtained by previous comparable work. A simple, Jeans- Equation based modeling approach, assuming (an-)isotropies of either $\\rm \\beta~ = ~0.37 ~or~ \\beta~ = ~0$ (as found for the halo stars in the cosmological simulations) yields results that are consistent with these values.  Although each of the $\\rm V_{cir}(r)$ points were estimated independently, the implied overall profile is consistent with both the mass profile in the simulations and with a parameterized mass model that combines a fixed disk and bulge model with an NFW dark matter halo, whose concentration c corresponds to the expected mean value for its virial mass. We have also explored halo fits with the concentration c as a free parameter, as well as halo profiles that have been modified by adiabatic contraction. We find that our data cannot discriminate well whether adiabatic contraction occurs or not -- an uncontracted halo of higher than average concentration and a contracted halo of (initially) low concentration fit comparably well. The resulting virial masses, $\\rm M_{vir}=1.0^{+0.3}_{-0.2} \\times 10^{12}$M$_\\odot$, are consistent for both fitting approaches. We have also checked that these results are quite robust with respect to distance errors, modest sample contamination ($\\leq 10\\% $), and the choice of a different, independent galaxy formation simulation.  The estimate of $\\rm M_{vir}$, which does imply an extrapolation from r$_{\\rm max}=60$ kpc to the virial radius of $\\sim 250$ kpc, is consistent with a recent estimate from a much smaller sample of halo stars (Battaglia et al. 2006), but it is lower than previous estimates that also rely on the kinematics of satellite galaxies (e.g., Kochanek et al. 1996). However, recent results on the LMC (Kallivayalil et al. 2006; Besla et al. 2007) indicate that not even the Magellanic Clouds may have been bound to the Milky Way for long, posing a potential conceptual problems for the use of satellite dynamics. It should be interesting to explore how our new constraint modifies the Local Group timing argument and its implication for M31's halo mass (Li \\& White 2008).  The estimate of $\\rm M_{vir}\\sim 10^{12}$M$_\\odot$, together with an estimated total cold baryonic mass of $6.5\\times 10^{10}$M$_\\odot$ and a global baryon mass fraction of $0.17$, implies that nearly 40\\% of all baryons within the virial radius have cooled to form the Milky Way's stars and (cold) gas, consistent with recent estimates for galaxies of that mass scale, based on statistical arguments (van den Bosh et al. 2007; Gnedin et al. 2007).  We note that parts of our analysis have been performed under the assumption that the stellar halo of the Galaxy is considered as a single relaxed population, or one that matches the simulations. Our data on the overall dynamics are consistent with that hypothesis. Recent evidence from Carollo et al. (2007) and Miceli et al. (2007) suggest that the halo may well be more complex, and comprise two distinct populations of inner- and outer-halo stars, with (slightly) different net rotational velocities and spatial profiles. We defer all analysis of kinematic and spatial sub-structure in this particular sample to a future paper.  The SDSS and SEGUE surveys have shown in this context that they can provide unparalleled sets of kinematic tracers for the Milky Way, enabling a direct ``circular velocity curve'' estimate of the Milky Way extending to 60 kpc. Once the full set of spectroscopy from SEGUE is available, a much larger set of stars for such an analysis should be available.  The proposed extension of SDSS, known as SDSS-III (and using a more sensitive, 1000 fiber spectrograph), will provide the opportunity to obtain higher quality spectra for fainter, more distant BHB stars, thus extending the reach of our analysis to over 100 kpc."
    },
    "0801/0801.3237_arXiv.txt": {
        "abstract": "We compare the performance of two alternative algorithms which aim to construct a force-free magnetic field given suitable boundary conditions. For this comparison, we have implemented both algorithms on the same finite element grid which uses Whitney forms to describe the fields within the grid cells. The additional use of conjugate gradient and multigrid iterations result in quite effective codes.  The Grad-Rubin and Wheatland-Sturrock-Roumeliotis algorithms both perform well for the reconstruction of a known analytic force-free field. For more arbitrary boundary conditions the Wheatland-Sturrock-Roumeliotis approach has some difficulties because it requires overdetermined boundary information which may include inconsistencies. The Grad-Rubin code on the other hand loses convergence for strong current densities. For the example we have investigated, however, the maximum possible current density seems to be not far from the limit beyond which a force free field cannot exist anymore for a given normal magnetic field intensity on the boundary. ",
        "introduction": "\\label{sec:Intro}  With the advent of vector magnetographs which measure the line-of-sight component of the photospheric magnetic field and, except for a 180$^{\\circ}$ ambiguity, also its component normal to the line-of-sight, the interest in extrapolating these measurements into the corona have grown enormously.  The line-of-sight component of the photospheric magnetic field has been observed for decades now but these measurements alone supply only boundary information at best sufficient for a Laplace field model of the coronal magnetic field. The vector magnetograph observations now available considerably constrain the photospheric horizontal field and allow estimates also of the coronal current density. As a consequence much more realistic coronal field models can be based on these observations.  Since the magnetic field, at least in the lower corona, completely dominates the plasma forces, a ``force-free'' approximation of the field for stationary situations seems to be a tolerable assumption. ``Force'' in this context means the magnetic Lorentz force $\\vec{j}\\crss\\vec{B}$. All other MHD forces like gravity, pressure, etc. are neglected because they are about three orders of magnitude smaller than $jB$. Hence, the current and magnetic field vectors should be aligned to better than half a degree.  These simplifications accepted, the magnetic field in some domain $V$ of the corona may be described by \\begin{equation} \\div{\\vec{B}}=0 \\;;\\quad \\crl{\\vec{B}}=\\vec{j} \\;;\\quad \\vec{j}\\crss\\vec{B}=0\\,. \\label{FF}\\end{equation} The alignment of current and magnetic field causes the problem to be nonlinear, hence the questions which boundary information is to be supplied and how to solve (\\ref{FF}) are by no means trivial.  Boundary conditions which seem necessary and sufficient for (\\ref{FF}) are \\citep{Boulmezaoud:Amari:2000} \\begin{equation} \\vec{n}\\sdot\\vec{B} \\mathtext{on} \\partial V \\;;\\quad \\vec{n}\\sdot\\crl{\\vec{B}} \\mathtext{on} (\\partial V)^- \\mathtext{or} (\\partial V)^+\\,, \\label{BC}\\end{equation} where $(\\partial V)^{\\pm}$ is that part of the surface of $V$ where $\\vec{n}\\sdot\\vec{B}$ is either $>$ 0 or $<$ 0. But not all the boundary values which comply with (\\ref{BC}) are allowed. The field line twist which can be stationarily maintained, i.e., $\\vec{n}\\sdot\\crl{\\vec{B}}$ on $\\partial V$, is limited by the field energy sustained from the exterior, i.e., by $\\vec{n}\\sdot\\vec{B}$ (see section \\ref{results:TwistedLoop}).  Quite some effort has gone into attempts to solve (\\ref{FF}) for given boundary values either in the form (\\ref{BC}) or differently. Among the most promising schemes are those suggested a long time ago by Grad and Rubin (\\citeyear{Grad:Rubin:1958}, abbreviated GR) and more recently by Wheatland, Sturrock, and Roumeliotis (\\citeyear{Wheatland:etal:2000}, abbreviated WSR). The GR code has first been implemented by Sakurai (\\citeyear{Sakurai:1981}) and further developed by Amari et al. (\\citeyear{Amari:etal:1999a}), R\\'egnier et al. (\\citeyear{Regnier:etal:2002}) and Wheatland (\\citeyear{Wheatland:2004}). The WSR scheme has been extended by Wiegelmann (\\citeyear{Wiegelmann:2004}) and Wiegelmann and Inhester (\\citeyear{Wiegelmann:Inhester:2003}). The development still is an active area of research. A comparison and description of codes in use and other alternative approaches can be found in Schrijver et al. (\\citeyear{Schrijver:etal:2006}).  A difficulty for a thorough comparison of the various schemes, however, is the fact that they are often implemented very heterogeneously and also include many different details which could easily speed up or slow down their performance and sometimes may even obscure the basic advantages or disadvantages of an approach.  In order the make the two schemes comparable, we apply them to the same problem defined on exactly identical grids. Error norms to measure the performance are exactly the same. The special grid which we use is explained in section~\\ref{sec:Grid}. We are convinced that it has many advantages for electromagnetic problems like (\\ref{FF}).  The WSR and the GR algorithms have been extensively described elsewhere, so that in sections \\ref{sec:Wheatland} and \\ref{sec:GradRubin} we restrict ourselves to the basics and rather emphasize some or the details of our implementation. The numerical results obtained for two different examples are presented and discussed in the final parts \\ref{sec:Results} and \\ref{sec:Discussion}. ",
        "conclusions": "\\label{sec:Discussion}  Force-free extrapolation codes available nowadays are capable to compute a field model in boxes as big as 256$\\times$256$\\times$256 \\citep[e.g.,][]{Schrijver:etal:2006}. Compared to the resolution of modern vector magnetograms with pixel arrays of 1000$\\times$1000, the calculated extrapolation models are most often quite limited in their resolution. Keeping in mind that they have to oversample the observation in order to limit the discretization error, the observed data typically has to be smoothed before it can serve as boundary condition for an extrapolation \\citep{Wiegelmann:etal:2005b}. There is therefore a definite need for fast and effective force-free reconstruction codes which can handle bigger models in order to make proper use of the resolution which the observations provide.  In this paper we have implemented and tested two alternative schemes for the extrapolation of a force-free field from boundary data. The algorithms differ considerably and in order to compare these different approaches unobscured by differences in the numerical coding, we implemented them as far as possible in a similar way. The discretization on a finite element grid by means of Whitney forms results in codes with a very efficient performance. In addition, we have improved the WSR code considerably using an efficient conjugate gradient iteration.  The two schemes differ not only in their approach towards a solution, but they also differ in the boundary information which has to be supplied. The WSR code requires boundary information which clearly overdetermines the problem unless parts of the boundary fields are left to be varied. But the full magnetic field vector on even part of the boundary is a very strong constraint and there is a great danger to impose inconsistent boundary values. This may be one reason why the WSR code converges well for known solutions like the Low and Lou model \\citep{Low:Lou:1990} where precise and consistent boundary values are supplied, while a little less reliable boundary values like those we retrieved from the result of a GR extrapolation slow down the WSR convergence speed markedly.  The WSR code has as a free parameter the weight between the Lorentz force and the divergence term in (\\ref{L_B}). In our implementation it is hidden in the constant of the weight function $w$. We found const $\\sim$ 1 a good value for the problems which we have delt with here. However, since the (\\ref{FF}) is nonlinear, the performance and also the optimal parameters may vary with the problem studied. As an example, note the difference in the iteration history for the Low and Lou model for const = 2 and 5 in Figure~\\ref{fig:PltHstp}. Here, the convergence changes drastically if the constant is modified by only a small amount.  Typical computation times on a $n$ = 64 grid with our code are about 20 minutes on an ordinary home computer. At present, the codes still include some checkout overhead, which could be dispensed with. Major improvements can probably be obtained by spreading the calculation onto multiple grids. We have made first tests with the WSR code for an adaptive enhancement of the grid resolution. Here, the solution on the lower grid $n$ was used as initial iterate on the next finer grid $2n$. We estimate that this way the computation time can be reduced by about a factor 1/3. Even more can probably be gained if the a true multigrid scheme is used."
    },
    "0801/0801.2007_arXiv.txt": {
        "abstract": "{ We use observations of Sunyaev-Zel'dovich effect and X-ray surface brightness to reconstruct the radial profiles of gas temperature and density under the assumption of a spherically symmetric distribution of the gas. The method of reconstruction, first raised by Silk \\& White, depend directly on the observations of the Sunyaev-Zel'dovich effect and the X-ray surface brightness, without involving additional assumptions such as the equation of state of the gas or the conditions of hydrostatic equilibrium. We applied this method to the cluster RX J1347.5--1145, which has both the Sunyaev-Zel'dovich effect and X-ray observations with relative high precision. It is shown that it will be an effective method to obtain the gas distribution in galaxy clusters. Statistical errors of the derived temperature and density profiles of gas were estimated according to the observational uncertainties. ",
        "introduction": "Galaxy cluster is known to be the largest virial gravitational bound system in the universe. The strong gravitational potential heats the gas (Hydrogen and Helium) to be fully ionized with a very high temperature $\\sim10^7-10^8$K. The thermal electrons collide with ions and emit bremsstrahlung radiation, which is mainly in X-ray band, making galaxy clusters strong X-ray sources. The intensity of thermal bremsstrahlung radiation is approximately proportional to $T_e^{1/2}n_e^2$ (\\cite{1999PhR...310...97B}). So by measuring the X-ray surface brightness, we can obtain information about the gas temperature and density distributions of the galaxy cluster. One of the most popular models is the so-called isothermal $\\beta$ model (\\cite{1976A&A....49..137C}), in which the temperature distribution as a function of radius is assumed to be constant and the density profile takes the form $n_e(R)=n_{e0}(1+R^2/R_c^2)^{-3\\beta/2}$. This simple model gives good approximations to many galaxy clusters. However, more and more observations indicate that the temperature distribution may not be a constant over the whole galaxy. Hughes et al. (1988) introduced a revised model with a constant temperature in a range $r<r_{\\rm iso}$ and a decrease at large radius, while the electron density has a truncation at a radius $r_{\\rm lim}$. This form provides an adequate description for some of the clusters. Further modification on the $\\beta$ model has been proposed to deal with the existence of cooling flow in the center of the cluster (\\cite{1984Natur.310..733F}).   In order to get rid of the degeneracy between $T_e$ and $n_e$, further assumptions like the hydrostatic equilibrium condition or the polytropic equation of state (EOS) of the electron gas (\\cite{1987ApJ...317..593C}; \\cite{2000A&A...360L..43X}; \\cite{2001ApJ...547...82W}), or the spectrometric analysis (\\cite{2006ChJAA...6..181J}) are needed. However, the results are still model dependent. A realistic reconstruction is expected to come directly from the observations.   The inverse Compton scattering between the electrons in clusters and the cosmic microwave background (CMB) photons, namely the Sunyaev-Zel'dovich (SZ) effect (\\cite{1969Ap&SS...4..301Z}; \\cite{1970Ap&SS...7....3S}, b), provides a possible way to achieve this goal. The distortion of the CMB spectrum is related to the electron temperature and density, and will provide another equation of $T_e$ and $n_e$. Combining the SZ effect and the X-ray observations, one can obtain the electron temperature and density profiles. This method was first suggested by \\cite{1978ApJ...226L.103S}, but has not yet been put into practice due to the limited instrumental sensitivity and resolution in the detection of the temperature fluctuation of CMB. Up to the 1990s some preliminary observational results of the radial temperature distribution were obtained for a few clusters (\\cite{1991ApJ...379..466B}; \\cite{1994ApJ...420...33B}), but it was still difficult to obtain precise reconstructions of the temperature and density profiles because the data points were too sparse and uncertain (\\cite{2003ASPC..301..365W}).   In this paper, we try with this method to acquire some preliminary information about the gas temperature and density distribution. For simplicity, we assume a spherically symmetric distribution of the gas in the clusters. We adopt the $\\Lambda$CDM model with $\\Omega_M=0.27$, $\\Omega_{\\Lambda}=0.73$, and Hubble constant $H_0=71$ km s$^{-1}$ Mpc$^{-1}$ (\\cite{2003ApJS..148..175S}). This paper is organized as follows: we describe the method of reconstruction in Section~\\ref{method}. In Section~\\ref{apply} we apply this method to the galaxy cluster RX J1347.5--1145 to reconstruct its temperature and density profiles, together with an estimate of the uncertainties. The conslusions and discussions are presented in Section~\\ref{cd}. ",
        "conclusions": "\\label{cd} Using the observational data of the SZ effect and X-ray surface brightness, we applied the method suggested by \\cite{1978ApJ...226L.103S} to the galaxy cluster RX J1347.5--1145, to reconstruct its gas temperature and density profiles. This is a direct way to obtain the temperature and density profiles of galaxy clusters, without additional theoretical assumptions. However, the quality of the observational data strongly affect the reconstruction results. Our attempt on cluster RX J1347.5--1145, which has both the X-ray and SZ measurements with relative high precision, demonstrates the effectiveness of this method. The derived results show similar behaviors as the previous studies. It indicates that there is a cooling flow in the center of the cluster, but the central value we derived is lower than the others. Poor quality and possible systematic uncertainties of the SZ effect data might be responsible for this discrepancy. It should be noted that the model parameters of the X-ray surface brightness and SZ effect were derived from finite area images around the center of the cluster, then extrapolated to large radii. This may result in additional uncertainty (see the discussion in Section~\\ref{apply}). When high quality measurements of the SZ effect become available in the future, it will be a powerful tool to study the intrinsic gas distribution independent of any theoretical arguments."
    },
    "0801/0801.2694_arXiv.txt": {
        "abstract": "Nowadays, compact sources like surfaces of nearby stars, circumstellar environments of stars from early stages to the most evolved ones and surroundings of active galactic nuclei can be investigated at milli-arcsecond scales only with the VLT in its interferometric mode. We propose a spectro-imager, named VSI (VLTI spectro-imager), which is capable to probe these sources both over spatial and spectral scales in the near-infrared domain. This instrument will provide information complementary to what is obtained at the same time with ALMA at different wavelengths and the extreme large telescopes. ",
        "introduction": "\\label{sect:overview}  The VLTI Spectro Imager will provide the ESO community with spectrally-resolved near-infrared images at angular resolutions down to 1.1 milliarcsecond and spectral resolutions up to $R=12000$. Targets as faint as $K=13$ will be imaged without requiring a brighter nearby reference object; fainter targets can be accessed if a suitable off-axis reference is available. This unique combination of high-dynamic-range imaging at high angular resolution and high spectral resolution for a wide range of targets enables a scientific programme which will serve a broad user community within ESO and at the same time provide the opportunity for breakthroughs in many areas at the forefront of astrophysics.  A great advantage of VSI is that it will provide these new capabilities while using technologies which have extensively been tested in the past and while requiring little in terms of new infrastructure on the VLTI. At the same time, VSI will be capable to make maximum use of the new infrastructure as it becomes available. VSI provides the VLTI with an instrument capable of combining up to 8 telescopes, enabling rapid imaging through measurement of up to 28 visibilities in hundreds of wavelength channels within a few minutes. Operations with less than 8 telescopes is the scope of the first phases of VSI.  Three development phases are foreseen: VSI4 combining 4 telescopes (UTs or ATs), VSI6 combining 6 telescopes (4UTS+2ATs or 4ATs+2UTS and eventually 6ATs), and perhaps ultimately, in the long-run, VSI8 combining 8 telescopes (4UTs+4ATs or eventually 8ATs). \\textbf{The current studies were focused on a 4-telescope version with an upgrade to a 6-telescope one}. The instrument contains its own fringe tracker and wavefront control in order to reduce the constraints on the VLTI infrastructure and maximize the scientific return. ",
        "conclusions": ""
    },
    "0801/0801.0975_arXiv.txt": {
        "abstract": "The Warm-Hot Intergalactic Medium (WHIM) arises from shock-heated gas collapsing in large-scale filaments and probably harbours a substantial fraction of the baryons in the local Universe. Absorption-line measurements in the ultraviolet (UV) and in the X-ray band currently represent the best method to study the WHIM at low redshifts. We here describe the physical properties of the WHIM and the concepts behind WHIM absorption line measurements of \\ion{H}{i} and high ions such as \\ion{O}{vi}, \\ion{O}{vii}, and \\ion{O}{viii} in the far-ultraviolet and X-ray band. We review results of recent WHIM absorption line studies carried out with UV and X-ray satellites such as FUSE, HST, Chandra, and XMM-Newton and discuss their implications for our knowledge of the WHIM. ",
        "introduction": "\\label{Introduction}  As recent cosmological simulations imply, the temperature of the intergalactic medium (IGM) undergoes a significant change from high to low redshifts parallel to the proceeding of large-scale structure formation in the Universe (e.g., \\citealt{cen1999}; \\citealt{dave2001}). As a result, a substantial fraction of the baryonic matter in the local Universe is expected to reside in the so-called Warm-Hot Intergalactic Medium (WHIM). The WHIM represents a low-density ($n_{\\rm H}\\sim10^{-6}-10^{-4}$ cm$^{-3}$), high-temperature ($T\\sim 10^5-10^7$ K) plasma that primarily is made of protons, electrons, \\ion{He}{ii}, and \\ion{He}{iii}, together with traces of some highly-ionised heavy elements. The WHIM is believed to emerge from intergalactic gas that is shock-heated to high temperatures as the medium is collapsing under the action of gravity in large-scale filaments (e.g., \\citealt{valageas2002}). In this scenario, part of the warm (photoionised) intergalactic medium that gives rise to the Ly\\,$\\alpha$ forest in the spectra of distant quasars (QSO) is falling in to the potential wells of the increasingly pronounced filaments, gains energy (through gravity), and is heated to high temperatures by shocks that run through the plasma.  Because of the low density and the high degree of ionisation, direct observations of the shock-heated and collisionally ionised WHIM are challenging with current instrumentation (in contrast to the photoionised IGM, which is easily observable through the Ly\\,$\\alpha$ forest). Diffuse emission from the WHIM plasma must have a very low surface brightness and its detection awaits UV and X-ray observatories more sensitive than currently available (see, e.g., \\citealt{fang2005}; \\citealt{kawahara2006}). The most promising approach to study the WHIM with observations at low redshift is to search for absorption features from the WHIM in FUV and in the X-ray regime in the spectra of quasars, active galactic nuclei (AGN) and other suited extragalactic background sources. As the WHIM represents a highly-ionised plasma, the most important WHIM absorption lines are those originating from the electronic transitions of high-ionisation state ions (hereafter referred to as \"high ions\") of abundant heavy elements such as oxygen and carbon. Among these, five-times ionised oxygen (\\ion{O}{vi}) is the most valuable high ion to trace the WHIM at temperatures of $T\\sim 3\\times 10^5$ K in the FUV regime. In the X-ray band, the \\ion{O}{vii} and \\ion{O}{viii} transitions represent the key observables to trace the WHIM at higher temperatures in the range $3\\times 10^5 < T < 10^7$. In addition to the spectral signatures of high ions of heavy elements the search for broad and shallow Ly\\,$\\alpha$ absorption from the tiny fraction of neutral hydrogen in the WHIM represents another possibility to identify and study the most massive WHIM filaments in the intergalactic medium with FUV absorption spectroscopy. Finally, for the interpretation of the observed WHIM absorption features in UV and X-ray spectra the comparison between real data and artificial spectra generated by numerical simulations that include realistic gas physics is of great importance to identify possible pitfalls related to technical and physical issues such as limited signal-to-noise ratios and spectral resolution, line-broadening mechanisms, non-equilibrium conditions, and others.  In this chapter, we review the physics and methodology of the UV and X-ray absorption measurements of warm-hot intergalactic gas at low redshift and summarise the results of recent observations obtained with space-based observatories. The outline of this chapter is the following. The ionisation conditions of the WHIM and the most important absorption signatures of this gas in the UV and X-ray band are presented in Sect.\\,2. Recent UV absorption measurements of the WHIM at low redshift are discussed in Sect.\\,3. Similarly, measurements of the WHIM in the X-ray are presented in Sect.\\,4. In Sect.\\,5 we compare the results from WHIM observations with predictions from numerical simulations and give an overview of WHIM measurements at high redshift. Finally, some concluding remarks are given in Sect.\\,6. ",
        "conclusions": ""
    },
    "0801/0801.2141_arXiv.txt": {
        "abstract": "\\re{Photometric maps, obtained with {\\it Spitzer}'s Infrared Array Camera (IRAC), can provide a valuable probe of warm molecular hydrogen within the interstellar medium.}  IRAC maps of the supernova remnant IC443, extracted from the {\\it Spitzer} archive, are strikingly similar to spectral line maps of the H$_2$ pure rotational transitions that we obtained with the Infrared Spectrograph (IRS) instrument on {\\it Spitzer}.   IRS spectroscopy indicates that IRAC Bands 3 and 4 are indeed dominated by the H$_2$ $v=0$--0 S(5) and S(7) transitions, respectively.  Modeling of the H$_2$ excitation suggests that Bands 1 and 2 are dominated by H$_2$ $v = 1-0$ O(5) and $v=0$--0 S(9).   Large maps of the H$_2$ emission in IC433, obtained with IRAC, show band ratios that are inconsistent with \\re{the presence of gas at a single temperature.}  The relative strengths of IRAC Bands 2, 3, and 4 are consistent with pure H$_2$ emission from shocked material with a power-law distribution of gas temperatures.  CO vibrational emissions \\ree{do not} contribute significantly to the observed Band 2 intensity. Assuming that the column density of H$_2$ at temperatures $T$ to $T+dT$ is proportional to $T^{-b}$ for temperatures up to 4000~K, we obtained a typical estimate of 4.5 for $b$.  The power-law index, $b$, shows variations over the range $\\sim 3 - 6$ within the set of different sight-lines probed by the maps, with the majority of sight-lines showing $b$ in the range $4-5$.  The observed power-law index is consistent with the predictions of simple models for paraboloidal bow shocks. ",
        "introduction": "Warm molecular hydrogen, at temperatures of several hundred to several thousand K, has been observed widely in molecular clouds that have been heated by interstellar shock waves (e.g. Gautier et al.\\ 1976, Treffers et al.\\ 1976, Combes \\& Pineau des Forets 2000.)  Such clouds are frequently associated with protostellar outflows and supernova remnants (SNR), and emit a rich spectrum of rovibrational and pure rotational H$_2$ emissions.  While the near-infrared rovibrational transitions have been observed extensively with ground-based observatories, observations of the mid-infrared pure rotational transitions are significantly hampered by atmospheric absorption; thus ground-based observations can provide only a fragmentary picture of the pure rotational spectrum of warm H$_2$, and space-based observations from the {\\it Infrared Space Observatory (ISO)} or the {\\it Spitzer Space Telescope} have been needed to provide a complete spectrum.  In particular, recent observations with the Infrared Spectrograph (IRS) on {\\it Spitzer} have provided sensitive spectral line maps for the S(0) through S(7) pure rotational transitions of H$_2$, allowing the gas temperature and ortho-to-para ratio of the emitting H$_2$ to be determined.  Such observations have revealed the presence of an admixture of multiple gas temperatures along every sight-line that has been observed, with non-equilibrium ortho-to-para ratios readily apparent in the cooler gas components, and have provided important constraints upon theoretical models for interstellar shock waves (Neufeld et al.\\ 2006a, 2007).  While warm molecular hydrogen has been studied primarily through spectroscopic observations with the IRS on {\\it Spitzer}, photometric observations with the Infrared Array Camera (IRAC) are also possible.  As noted recently by Reach et al.\\ (2006; see their Fig.\\ 1), the $v=1-0$ O(5) and $v=0-0$ S(4) - S(13) transitions of H$_2$ fall within the IRAC bandpass and can contribute significantly to the intensities observed with IRAC; in some sources, the IRAC intensities are potentially dominated by H$_2$ line emissions.  In this paper, we investigate further the utility of IRAC observations in mapping the distribution of warm molecular hydrogen.  In \\S 2, we discuss IRAC and IRS observations carried out toward the supernova remnant IC443.  In \\S 3, we compare the spatial distribution of the H$_2$ line emissions, observed with IRS, with the photometric maps obtained with IRAC.  In \\S 4, we discuss how the relative intensities within the four IRAC bands can be used to constrain the physical conditions within the warm molecular gas. ",
        "conclusions": "The excitation of molecular hydrogen behind interstellar shock waves has been the subject of considerable discussion over the past two decades.  In the late 1980's, with the advent of spectral line mapping at near-IR wavelengths, two key features of the H$_2$ emission were noted: (1) the relative line strengths observed toward shocked regions are \\ree{typically} inconsistent with a single excitation temperature (Burton et al.\\ 1989); and (2) even line ratios that are strongly dependent upon the gas temperature often show a remarkable constancy throughout a given source (Brand et al.\\ 1989).  \\ree{These behaviors have subsequently been observed in a wide variety of sources, and are particularly evident in {\\it Spitzer} observations where the S(1) -- S(7) pure rotational lines are mapped over a common region (e.g.\\ in both the Herbig Haro objects -- HH54 and HH7 -- observed by Neufeld et al.\\ 2006a; and in all the supernova remnants -- IC443, W28, W44, 3C391 -- observed by Neufeld et al.\\ 2007).}  It was recognized (Smith \\& Brand 1990) that these behaviors are unexpected for single planar shocks that are of ``C-type'' in the designation of Draine (1980).  In such shocks, ion-neutral drift is a critical feature of the physics (Mullan 1971), allowing the velocities of the ionized and neutral species to change continuously (hence the designation ``C-type'' for ``continuous'') in a warm region of roughly uniform temperature: as a result (a) the rotational diagrams for the H$_2$ level populations behind such shocks show little curvature (e.g. Neufeld et al.\\ 2006a, Appendix B); and (b) temperature behind the shock front is a strong function of the shock velocity\\footnote{ While planar C-type shocks were subsequently shown to be unstable (Wardle 1991), numerical simulations of the non-linear development of the instability (Neufeld \\& Stone 1997; MacLow \\& Smith 1997) suggested that both results (a) and (b) still apply.}. The former is inconsistent with observational feature (1) above, and the latter is inconsistent with (2) unless the shock velocity is remarkably constant throughout the mapped region.  Two separate explanations were offered to account for the observations, i.e.\\ the key features (1) and (2) described above.  First, Burton et al.\\ (1989) noted that the cooling region behind a ``J-type'' shock might successfully explain the H$_2$ line ratios.  In a ``J-type'' shock, where the ion and neutral species are well-coupled and show velocities that jump discontinuously at the shock front (hence the designation ``J-type'' for ``jump''), the gas is initially heated to a high temperature and then cools radiatively, resulting in an admixture of gas temperatures.  Provided the shock is not fast enough to destroy molecular hydrogen, J-type shocks yield a rotational diagram with significant curvature but which shows a relatively weak dependence upon the shock velocity.  As a second and alternative explanation, Smith, Brand \\& Moorhouse (1991; hereafter SBM) suggested bow shocks as the source of the observed H$_2$ emissions.  Such shocks result when a supersonic flow encounters a clump of material or when a supersonic clump travels through an ambient medium.  In SBM's model, multiple unresolved bow shocks are assumed to be present in each telescope beam.  At the head of each bow shock, where the gas moves perpendicular to the shock front, the effective shock velocity is greatest; here the shock may become dissociative and of J-type.  On the sides of the bow shock, where the material enters the shock front at an oblique angle, the shock velocity is smaller and the shock is non-dissociative and of C-type.  Assuming a parabolic shape for the shock surface, and decomposing a bow shock into a set of planar shocks of varying shock velocity, SBM considered the resultant H$_2$ emission. The H$_2$ rotational diagram is indeed expected to show significant curvature, consistent with the observations.  Furthermore, provided that the shock velocity at the head of the bow shock is sufficient to dissociate molecules, SBM showed that the H$_2$ line ratios are independent of the flow velocity.  Bow shocks offer several advantages over J-type shocks in explaining the observations.  First, given the typical ionization level and magnetic field strengths in molecular clouds, non-dissociative shock waves are expected to be of C-type, not J-type.  Second, high-resolution observations of shocked H$_2$ (e.g. Schultz et al.\\ 1999), carried out after the explanations of Burton et al.\\ (1989) and SBM were offered, have indeed revealed the presence of bow shocks that were unresolved in previous observations.  Finally, recent observations of the S(0) -- S(7) rotational transitions of H$_2$ (Neufeld et al.\\ 2006a), carried out with {\\it Spitzer}/IRS toward six separate sources, have shown that the H$_2$ ortho-to-para ratio is smaller than its equilibrium value, {\\it with the cooler components along every sight-line showing the greatest departures from equilibrium.}  This behavior is simply inconsistent with the J-type shock model, in which the warmer gas components evolve to become the cooler components.  To the contrary, the cooler components must arise in separate shocked regions.  We can compare the power-law index, $b$, derived in \\S 4 above with the predictions of bow shock models.  For a curved shock front with a paraboloidal shape, material crossing the shock front with a perpendicular velocity between $v$ to $v+dv$ is associated with a shock surface area $dA \\propto v^{-4} dv$ (Smith \\& Brand 1990); this proportionality applies for all velocities less than the maximum shock velocity, $v_{\\rm max}$, attained at the head of the bow shock.  Based upon a fit to the planar C-type shock models of Kaufman \\& Neufeld (1996), Neufeld et al.\\ (2006a, Appendix B) found that the warm shocked material could be adequately approximated by a slab of constant temperature $\\propto v^{1.35}$ and column density $N ({\\rm H}_2) \\propto v^{-0.75}$.   These proportionalities yield a power-law temperature distribution, with the mass at temperature between $T$ and $T+dT$ obeying $$dM \\propto N({\\rm H}_2)  dA \\propto v^{-4.75} dv \\propto T^{-(3.75+1.35)/1.35} dT \\propto T^{-3.78} dT \\eqno(3).$$  For bow shocks with a maximum shock velocity sufficient to dissociate H$_2$ at the head of the bow, this leads to a power-law index $b \\sim 3.8$ extending to the maximum possible H$_2$ temperature, $T_{max}$.  For slower bow shocks, the temperature distribution extends to some temperature lower than $T_{max}$.  Thus an admixture of bow shocks, with a range of velocities some of which insufficient to achieve temperature $T_{\\rm max}$ at the head of the bow shock, could lead to an overall temperature distribution which is steeper than $dN \\propto T^{-3.8} dT$ (i.e.\\ to values of $b$ greater than 3.8).  These simple predictions are broadly consistent with the observations, which indicate values of $b < 3.8$ to be quite rare.  In this scenario, the observed spatial variations in $b$ could result from variations in the distribution of velocities at the heads of a set of unresolved bow shocks.  We note, however, that our assumption of parabolic bow shocks in steady-state is a highly idealized one, and that numerical simulations would be needed to provide a more realistic prediction for the shock structure and resultant temperature distribution in the presence of possible instabilities (including the Wardle instability)."
    },
    "0801/0801.2155_arXiv.txt": {
        "abstract": "Mass estimates of K~giants are generally very uncertain. Traditionally, stellar masses of single field stars are determined by comparing their location in the Hertzsprung-Russell diagram with stellar evolutionary models. Applying an additional method to determine the mass is therefore of significant interest for understanding stellar evolution. We present the time series analysis of 11 K~giants recently observed with the \\wire\\ satellite. With this comprehensive sample, we report the first confirmation that the characteristic acoustic frequency, \\numax, can be predicted for K~giants by scaling from the solar acoustic cut-off frequency. We are further able to utilize our measurements of \\numax\\ to determine an asteroseismic mass for each star with a lower uncertainty compared to the traditional method, for most stars in our sample. This indicates good prospects for the application of our method on the vast amounts of data that will soon come from the COROT and Kepler space missions. ",
        "introduction": "According to theoretical calculations of stellar evolution, stars of a large mass range, corresponding to main-sequence spectral classes from late B to K, all end up in roughly the same part of the Hertzsprung-Russell (H-R) diagram when they evolve to become red giants with spectral classes from late G to M. Here, the evolution tracks are closely spaced, which makes it difficult to estimate the mass, and hence to determine the progenitor of a red giant star from its position in the H-R diagram.  The investigation of stellar oscillations (asteroseismology), in particular solar-like oscillations, provides a unique tool to probe the interior and hence also the mass of stars. Much excitement therefore followed the first clear evidence of solar-like oscillations in a red giant star \\citep{Frandsen02}. % Subsequently, researchers have seen evidence of % oscillations in about a dozen K and late G type giants, both in nearby field stars \\citep{Ridder06,Barban07,Tarrant07} and members of the open cluster \\objectname[M67]{M67} \\citep{Stello07}.  However, % it remains uncertain if accurate mode frequencies are possible to obtain, since observational indications are that mode lifetimes could be too short \\citep{Stello06}, in disagreement with theory \\citep{HoudekGough02}. Further, these stars might pulsate predominantly in radial mode overtones as suggested by \\citet{Dalsgaard04} or indeed show additional non-radial pulsations as indicated by \\citet{Hekker06} and \\citet{Kallinger07}. Without a full understanding of which modes we observe, we cannot exploit the full potential of the asteroseismic analysis.  In this paper we aim at utilizing the characteristic acoustic frequency, denoted \\numax. This asteroseismic quantity is relatively insensitive to the mode lifetime and the number of excited modes, but still enables us to obtain information about the stellar interiors. In particular, we will use \\numax\\  to estimate the masses of a sample of K~giants observed with the star tracker on the {\\it Wide-Field Infrared Explorer} (\\wire) satellite. ",
        "conclusions": "We have analyzed photometric time series of 11 nearby K~giants % obtained with the \\wire\\ satellite. The Fourier transforms show clear evidence that oscillation power shifts to higher frequencies for less luminous stars, as anticipated from scaling the solar frequency of maximum power, \\numax$_{,\\odot}$. We were able to measure \\numax\\  and made simulations of each star to obtain a realistic uncertainty of the measurement. Using a simple scaling relation, which relates this frequency to the stellar parameters $T_{\\mathrm{eff}}$, $L/\\mathrm{L}_{\\odot}$, and $M/$M$_{\\odot}$, we estimated an asteroseismic mass. For several stars this approach provides a significantly lower uncertainty of the mass relative to the classical mass estimate based purely on comparing stellar evolution tracks with the location in the H-R diagram.  These results show exciting prospects for the current COROT mission \\citep{Baglin98} and the upcoming Kepler satellite \\citep{Dalsgaard07}, which will both provide much more extended times series than \\wire. With longer time series we can potentially acquire lower uncertainties in the \\numax\\ measurements, and hence more precise mass estimates. This will indeed be possible with Kepler, which will obtain parallaxes, and hence luminosities, of its target stars. Our approach % could be particularly valuable in cases where the oscillation power does not allow accurate detection of the large frequency separation, $\\Delta\\nu$, due to an insufficient number of modes with high signal-to-noise. This might, in fact, include most faint solar-like pulsators, as well as bright giant stars if indeed the mode lifetime does not increase with increasing oscillation periods \\citep{Stello06}."
    },
    "0801/0801.0366_arXiv.txt": {
        "abstract": "We propose a theory to explain the formation of both spirals and rings in barred galaxies using a common dynamical framework. It is based on the orbital motion driven by the unstable equilibrium points of the rotating bar potential. Thus, spirals, rings and pseudo-rings are related to the invariant manifolds associated to the periodic orbits around these equilibrium points. We examine the parameter space of three barred galaxy models and discuss the formation of the different morphological structures according to the properties of the bar model. We also study the influence of the shape of the rotation curve in the outer parts, by making families of models with rising, flat or falling rotation curves in the outer parts. The differences between spiral and ringed structures arise from differences in the dynamical parameters of the host galaxies. ",
        "introduction": "Bars are a very common feature of disc galaxies. In a sample of $186$ spirals drawn from the Ohio State University Bright Spiral Galaxy Survey, Eskridge {\\it et al.} (\\cite{esk00}) find that $56\\%$ of the galaxies in the near infrared are strongly barred, while an additional $6\\%$ are weakly barred. A large fraction of barred galaxies show two clearly defined spiral arms (e.g. Elmegreen \\& Elmegreen \\cite{elm82}), often departing from the end of the bar at nearly right angles. This is the case for instance in NGC~1300, NGC~1365 and NGC~7552. Deep exposures show that these arms wind around the bar structure and extend to large distances from the centre (see for instance Sandage \\& Bedke \\cite{san94}). Almost all researchers agree that spiral arms and rings are driven by the gravitational field of the galaxy (see Toomre \\cite{too77} and Athanassoula \\cite{ath84}, for reviews). In particular, spirals are believed to be density waves in a disc galaxy (Lindblad \\cite{lind63}). Toomre (\\cite{too69}) found that the spiral waves propagate towards the principal Lindblad resonances of the galaxy, where they damp down, and thus concludes that long-lived spirals need some replenishment. There are essentially three different possibilities for a spiral wave to be replenished. First, it can be driven by a companion or satellite galaxy. A direct, relatively slow, and close passage of another galaxy can form trailing shapes (e.g. Toomre \\cite{too69}; Toomre \\& Toomre \\cite{too72}; Goldreich \\& Tremaine \\cite{gol78}, \\cite{gol79}; Toomre \\cite{too81} and references therein). They can also be excited by the presence of a bar. Several studies have shown that a rotating bar or oval can drive spirals (e.g. Lindblad \\cite{lind60}; Toomre \\cite{too69}; Sanders \\& Huntley \\cite{san76}; Schwarz \\cite{sch79}, \\cite{sch81}; Huntley \\cite{hun80}). The third alternative, proposed by Toomre (\\cite{too81}), is the swing amplification feedback cycle. This starts with a leading wave propagating from the centre towards corotation. In doing so, it unwinds and then winds in the trailing sense, while being very strongly amplified. This trailing wave will propagate towards the centre, while a further trailing wave is emitted at corotation and propagates outwards, where it is dissipated at the Outer Lindblad Resonance. The inwards propagating trailing wave, when reaching the centre will reflect into a leading spiral, which will propagate outwards towards corotation, thus closing the feedback cycle. Danby (\\cite{dan65}) argued that orbits in the gravitational potential of a bar play an important role in the formation of arms. He noted that orbits departing from the vicinity of the equilibrium points located at the ends of the bar describe loci with the shape of spiral arms and can be responsible for the transport of stars from within to outside corotation, and vice versa. Unfortunately, he did not set his work in a rigorous theoretical context, so that it remained purely phenomenological. He also investigated whether orbits can be responsible for ring-like structures, but in this case, he did not consider orbits departing from the ends of the bar as he previously did when accounting for the spiral arms.  Strongly barred galaxies can also show prominent and spectacular rings or partial rings. The origin of such morphologies has been studied by Schwarz (\\cite{sch81}, \\cite{sch84}, \\cite{sch85}), who followed the response of a gaseous disc galaxy to a bar perturbation. He proposed that ring-like patterns are associated to the principal orbital resonances, namely ILR, CR, and OLR. There are different types of outer rings. Buta (\\cite{but95}) classified them according to the relative orientation of the ring and bar major axes. If these two axes are perpendicular, the outer ring is classified as $R_1$. If the two axes are parallel, the outer ring is classified as $R_2$. Finally, if both types of rings are present in the galaxy, the outer ring is classified as $R_1R_2$.  In Romero-G\\'omez {\\it et al.} (\\cite{rom06,rom07}), we propose that rings and spiral arms are the result of the orbital motion driven by the invariant manifolds associated to periodic orbits around unstable equilibrium points. In Romero-G\\'omez {\\it et al.} (\\cite{rom06}), we fix a barred galaxy potential and we study the dynamics around the unstable equilibrium points. We give a detailed definition of the invariant manifold associated to a periodic orbit. For the model considered, the invariant manifolds delineate well the loci of an $rR_1$ ring structure. In Romero-G\\'omez {\\it et al.} (\\cite{rom07}), we construct families of models based on simple, yet realistic, barred galaxy potentials. In each family, we vary one of the free parameters of the potential and keep the remaining fixed. For each model, we numerically compute the orbital structure associated to the invariant manifolds. In this way, we are able to study the influence of each model parameter on the global morphologies delineated by the invariant manifolds.  In Sect. \\ref{sec:mod}, we first present the equations of motion and the galactic models used in the computations. In Sect. \\ref{sec:inv}, we give a brief description of the dynamics around the unstable equilibrium points. In Sect. \\ref{sec:res}, we show the different morphologies that result from the computations. In Sect. \\ref{sec:dis}, we compare our results with some observational features and conclude. ",
        "conclusions": "\\label{sec:dis}  The family of Lyapunov orbits is unstable and becomes stable only at high energy levels (Skokos {\\it et al.} (\\cite{sko02}). We compute the invariant manifolds of Lyapunov orbits in the range of energies in which they are unstable. In the right panel of Fig. \\ref{myfig3}, we plot the invariant manifolds for two different energy levels of a given model. We find that the locus of the invariant manifolds is independent of the energy. As the energy increases, however, the size of the Lyapunov orbits also increases and thus the size of the invariant manifolds. Nevertheless, as we consider more energy levels, we find that the density in the central part is higher (see left panel of Fig. \\ref{myfig3}). Therefore, the thickness of the ring observed is smaller than the thickness of the invariant manifold of higher energies. We also compute the radial and tangential velocities on the galactic plane and in a non-rotating reference frame along the ring and we observe that they are small perturbations of the circular velocity. The maximum deviation from the typical circular velocity of $200 \\rm{km}\\,\\rm{s}^{-1}$ is $\\pm 20 \\rm{km}\\,\\rm{s}^{-1}$ (Athanassoula {\\it et al.} \\cite{ath07}).  In the case of the spiral arms, we compute the density profile on two different angles, namely one near the beginning of the arm and one at the end. We find that the first cut has a narrow and high density profile, while the second cut, has a wide and low density profile. This is the typical behaviour of the grand design spiral arms, namely they are denser and brighter near the bar ends and then they become more diffuse (Athanassoula {\\it et al.} \\cite{ath07}).  \\begin{figure}[h] \\centering \\hspace{1.cm} \\includegraphics[scale=0.3,angle=-90.]{romerogomez_png_fig4.ps}\\hspace{0.55cm} \\includegraphics[scale=0.3,angle=-90.]{romerogomez_png_fig5.ps} \\caption{{\\it Left panel:} Density profile on a cut across the ring. {\\it Right panel:} Two unstable invariant manifolds for different values of $E_J$. Note how similar the regions they delineate are.} \\label{myfig3} \\end{figure}  To summarise, our results show that invariant manifolds describe well the loci of the different types of rings and spiral arms. They are formed by a bundle of trajectories linked to the unstable regions around the $L_1/L_2$ equilibrium points. The study of the influence of one model parameter on the shape of the invariant manifolds in the outer parts reveals that only the pattern speed and the bar strength affect the galaxy morphology. The study also shows that the different ring types and spirals are obtained when we vary the model parameters.  We have compared our results with some observational data. Regarding the photometry, the density profiles across radial cuts in rings and spiral arms agree with the ones obtained from observations. The velocities along the ring also show that these are only a small perturbation of the circular velocity."
    },
    "0801/0801.2649_arXiv.txt": {
        "abstract": "We investigate the extinction curves of young galaxies in which dust is supplied from Type II supernovae (SNe II) and/or pair instability supernovae (PISNe). Since at high redshift ($z>5$), low-mass stars cannot be dominant sources for dust grains, SNe II and PISNe, whose progenitors are massive stars with short lifetimes, should govern the dust production. Here, we theoretically investigate the extinction curves of dust produced by SNe II and PISNe, taking into account reverse shock destruction induced by collision with ambient interstellar medium. We find that the extinction curve is sensitive to the ambient gas density around a SN, since the efficiency of reverse shock destruction strongly depends on it. The destruction is particularly efficient for small-sized grains, leading to a flat extinction curve in the optical and ultraviolet wavelengths. Such a large ambient density as $n_{\\rm H}\\ga 1$ cm$^{-3}$ produces too flat an extinction curve to be consistent with the observed extinction curve for \\sdss\\ at $z=6.2$. Although the extinction curve is highly sensitive to the ambient density, the hypothesis that the  dust is predominantly formed by SNe at $z\\sim 6$ is still allowed by the current observational constraints. For further quantification, the ambient density should be obtained by some other methods. Finally we also discuss the importance of our results for observations of high-$z$ galaxies, stressing a possibility of flat extinction curves. ",
        "introduction": "Dust grains play an important role in the formation and evolution of galaxies. Dust grains control the energy balance in the interstellar medium (ISM) by absorbing stellar light and reemitting it in far infrared (FIR). Also, the surface of dust grains is a site for an efficient formation of H$_2$ molecules \\citep[e.g.][]{cazaux04}, which act as an effective coolant in metal-poor ISM. Those effects of dust turn on even at $\\sim 1$\\% of the solar metallicity according to the calculation by \\citet{hirashita02}, who argue that the star formation rate is enhanced because of the first dust enrichment in the history of galaxy evolution. The first sources of dust in the Universe are Type II (core-collapse) supernovae (SNe II) or pair instability supernovae (PISNe), since the lifetimes of their progenitors are short ($\\sim 10^6$ yr). In the local Universe, dust grains are also produced by evolved low mass stars \\citep{gehrz89}, but this production mechanism requires much longer ($\\ga 1$ Gyr) timescales. The first dust supplied by SNe II or PISNe may trigger the formation of low-mass stars via dust cooling \\citep{schneider03,omukai05}.  To quantify the above effects of dust in the early stages of galaxy evolution, it is crucial to know how much dust and what species of grains form in supernovae (SNe). It has been suggested by some observations of nearby SNe that dust is indeed produced in SNe, although the quantity of formed dust is still debated \\citep[e.g.][]{moseley89,dunne03,morgan03,hines04,sugerman06,meikle07,rho07}. By treating the nucleation and accretion in SNe, the dust composition and size distribution are theoretically calculated \\citep*{kozasa89,kozasa91}. Recently, in order to examine the effects of dust in Population III (Pop III) objects, the formation of dust in SNe II and PISNe is extensively examined (\\citealt{todini01}; \\citealt[hereafter N03]{nozawa03}, \\citealt{schneider04}). The motivation for considering PISNe comes from some evidence indicating that the stars formed from metal-free gas, Population III (PopIII) stars, are very massive with a characteristic mass of a few hundred solar masses \\citep[e.g.][]{nakamura01,bromm04}. Such massive stars are considered to begin pair creation of electron and positron after the helium burning phase, and finally an explosive nuclear reaction disrupts the whole stars \\citep{fryer01,heger02}. This explosion is called PISN.  Recently, the hypothesis that the dust formation is dominated by SNe II and/or PISNe at high redshift ($z$) is observationally examined for some objects. We expect that at $z>5$, when the cosmic age is less than 1 Gyr, the main sources of dust grains are SNe II and PISNe \\citep[e.g.][]{dwek07}. Extinction curves can be used to investigate the dust properties \\citep[e.g.][]{mathis90}. By using a sample of broad absorption line (BAL) quasars, \\citet{maiolino04a} show that the extinction properties of the low-$z$ ($z<4$) sample is different from those of the high-$z$ ($z>4.9$) sample. This result is suggestive of a change in the dust production mechanism in the course of galaxy evolution. The highest-$z$ BAL quasar in their sample, SDSS J104845.05+463718.3 (hereafter \\sdss) at $z=6.2$ shows an extinction curve flat at wavelength $\\lambda \\ga 1700$ \\AA\\ and rising at $\\lambda\\la 1700$ \\AA. \\citet{maiolino04b} show that the extinction curve of \\sdss\\  is in excellent agreement with the SN II dust models by \\citet{todini01}. \\citet{bianchi07} consider dust destruction by reverse shock in SNe, suggesting that 2--20\\% of the initial dust mass survives, and that the extinction curve after the destruction is still consistent with that of \\sdss. More recently, \\citet{stratta07} show that the dust extinction in the host galaxy of GRB 050904 at $z=6.3$ can be explained by the extinction curve of \\sdss, further supporting that the SNe II are the main sources of dust at $z>6$. Also, \\citet{willott07} find a similar extinction property for CFHQS J1509-1749 ($z=6.12$) to that of \\sdss.  There are other series of theoretical papers on the extinction curves of high-$z$ objects. \\citet[][hereafter H05]{hirashita05} calculate the extinction curve based on the dust production calculation by N03. They also reproduce the extinction curve of \\sdss\\ by using the the dust production in SNe II, although the dust composition and size distribution are different from those of \\citet{todini01}. Recently, \\citet[][hereafter N07]{nozawa07} have treated the dust destruction by the reverse shock as done by \\citet{bianchi07}, but considering the motion of dust relative to gas caused by the drag force and the destruction of dust in the radiative phase as well as in the non-radiative phase of supernova remnants. Then, they show that the size distribution of grains supplied in the ISM is strongly modified by the reverse shock. Grains smaller than $\\sim 0.02~\\mu$m are efficiently destroyed if the ambient hydrogen number density is larger than 0.1 cm$^{-3}$. Thus, it is important to reexamine the consistency between the observed extinction curve and the reverse shock destruction.  In this paper, we calculate the extinction curves based on the dust properties calculated by N07, who have focused on the effect of reverse shock destruction in SNe. Then, we compare the results with observed extinction curves at high $z$. This paper is organized as follows. First, we describe our theoretical treatment to calculate the extinction curves of SN II and PISN dust in Section \\ref{sec:model}. We show and examine our results in Section \\ref{sec:results}. We discuss our results from the observational viewpoint in Section \\ref{sec:obs}, and finally give the conclusion of this paper in Section \\ref{sec:sum}. ",
        "conclusions": "\\label{sec:sum}  We have theoretically investigated the extinction curves of grains produced in SNe II and PISNe. Since at high $z(>5)$, low-mass stars cannot be dominant sources for dust grains, SNe II and PISNe, whose progenitors are massive stars with short lifetimes, can govern the dust production. While our previous works (N03 and H05) did not consider the reverse shock destruction induced by a collision with ambient interstellar medium, we adopt the composition and size distribution of grains of N07, who take into account the reverse shock destruction.  We have found that the extinction curve is sensitive to the ambient gas density around SNe, since the efficiency of reverse shock destruction is largely dependent on it. The destruction is significant for small-sized grains, leading to a flat extinction curve in the optical and UV. Such a large ambient density as $n_{\\rm H}\\ga 1$ cm$^{-3}$ produces too flat an extinction curve to be consistent with the observed extinction curve for \\sdss\\ at $z=6.2$. Although the extinction curve is highly sensitive to the ambient density, the hypothesis that the  dust is predominantly formed by SNe at $z\\sim 6$ is still allowed if $n_{\\rm H}$ is smaller than 1 cm$^{-3}$. For further quantification, the ambient density should be constrained by some other methods.  It is worth noting that a flat extinction curve produced by effective reverse shock destruction may explain the absence of reddening of systems in which dust is known to be present by far-infrared/sub-mm emission or by depletion of heavy elements."
    },
    "0801/0801.2539_arXiv.txt": {
        "abstract": "Searches for supersymmetry are among the most exciting physics goals at Run~II of the Tevatron. In particular, in supersymmetric models with light charginos, neutralinos and sleptons, associated chargino--neutralino production can potentially be observed as multi--lepton events with missing energy. We discuss how, in the generic TeV--scale MSSM, the prospects for these chargino--neutralino searches are impacted by cosmological considerations, namely the neutralino relic abundance and direct detection limits. We also discuss what an observation of chargino--neutralino production at the Tevatron would imply for the prospects of future direct dark matter searches without assuming specific patterns of supersymmetry breaking. ",
        "introduction": "A consensus has formed within the astrophysics community in support of the conclusion that the majority of our universe's mass takes the form of cold, collisionless dark matter~\\cite{dmreview}. Despite the very large body of evidence in favor of dark matter's existence, the nature of this elusive substance remains unknown. Of the many dark matter candidates to have been proposed, one of the most compelling and most often studied is the lightest neutralino in $R$-parity conserving models of supersymmetry~\\cite{jungman}.\\bigskip  Among the most prominent missions of the Tevatron's Run~II are its searches for supersymmetry. Results from Tevatron searches for squarks and gluinos~\\cite{tevsquarksgluinos}, neutralinos and charginos~\\cite{tevneucha}, stops and sbottoms~\\cite{stopsbottom}, and the Higgs bosons of the Minimal Supersymmetric Standard Model (MSSM)~\\cite{tevhiggs} have each recently been published. While no evidence for supersymmetry has yet been found, in many cases these results represent the strongest limits to date. Although the Tevatron is not well suited to directly place limits on the properties of the lightest neutralino, the results of these other searches can have considerable implications for the nature of this dark matter agent.  One of the prime channels for observing supersymmetry at the Tevatron is associated neutralino--chargino production~\\cite{trilep_early,trilep_newer}.  These particles can decay to the lightest neutralino and leptons through the exchange of either sleptons or gauge bosons, resulting in events featuring three leptons and missing energy. The results of these searches are somewhat model dependent, but the current results from CDF and D0 can be used to exclude charginos as heavy as approximately 150~GeV in some models, well beyond LEP's chargino mass limit of 104~GeV. By the end of Run II, the Tevatron is expected to exclude selected models with charginos not far below 200~GeV.\\bigskip  Neutralino dark matter can be detected through its elastic scattering with nuclei. Experimental efforts designed to observe such events are known as direct detection. The prospects for this class of techniques depends on the composition of the lightest neutralino, as well as on the masses and couplings of the exchanged squarks and Higgs bosons. Generally speaking, information from collider searches for supersymmetry, whether detections or constraints, can be used to better estimate the prospects for the detection of neutralino dark matter. The relationship between Tevatron and LHC searches for heavy MSSM Higgs bosons and direct searches for neutralino dark matter has been studied in detail elsewhere~\\cite{marcela}. Here, we return to this theme, but focussing on searches for trilepton events from associated neutralino--chargino production at the Tevatron (see also Ref.~\\cite{trilep_dm}).  In the past, it has been possible to link the Tevatron's trilepton signature to other signatures for new physics, for example the decay $B_s \\to \\mu \\mu$~\\cite{bmumu}. Such links rely for example on a correlation between light sleptons and small values of $\\tan\\beta$ for the chargino and neutralino decays on one hand and the pseudoscalar Higgs boson mass and large values of $\\tan\\beta$ in flavor physics on the other. Naively, similar correlations should be present when the dark matter candidates annihilate mainly through an $s$-channel Higgs boson and the trilepton signature requires relatively light supersymmetric scalars. Moreover, one could imagine correlations between these signals in the co-annihilation region, if the lightest slepton is mass degenerate with the lightest neutralino, limiting the visibility of the trilepton channel. However, these fairly obvious correlations rely on a series of assumptions. First, the different MSSM scalar masses have to be correlated. Secondly, the light scalar masses should in some way be linked to the lightest neutralino mass and to the mass difference between the light chargino and neutralino. Last but not least, the dark matter particle should annihilate dominantly through one channel. The aim of this analysis is to determine how much of a correlation between dark matter and Tevatron searches survives if we assume only a TeV--scale MSSM spectrum with no specific patterns of supersymmetry breaking.\\bigskip  This article is structured as follows: In Section~\\ref{sec:tri}, we discuss the searches for associated chargino-neutralino production at the Tevatron. In Section~\\ref{sec:dm}, we turn our attention to the thermal relic abundance of neutralinos, focussing on those models within the reach of the Tevatron and the correlation between Tevatron measurements and the neutralino's relic density. In Section~\\ref{sec:direct}, we discuss direct detection prospects for such models and the correlations between those and Tevatron observations.  Finally, we summarize our results and conclusions in Section~\\ref{sec:conclusion}. ",
        "conclusions": ""
    },
    "0801/0801.0974.txt": {
        "abstract": "We investigate the effects of ambipolar diffusion and the Hall effect on the stability of weakly-ionized, magnetized planar shear flows. Employing a local approach similar to the shearing-sheet approximation, we solve for the evolution of linear perturbations in both streamwise-symmetric and non-streamwise-symmetric geometries using WKB techniques and/or numerical methods. We find that instability arises from the combination of shear and non-ideal magnetohydrodynamic processes, and is a result of the ability of these processes to influence the free energy path between the perturbations and the shear. They turn what would be simple linear-in-time growth due to current and vortex stretching from shear into exponentially-growing instabilities. Our results aid in understanding previous work on the behaviour of weakly-ionized accretion discs. In particular, the recent finding that the Hall effect and ambipolar diffusion destabilize both positive and negative angular velocity gradients acquires a natural explanation in the more general context of this paper. We construct a simple toy model for these instabilities based upon transformation operators (shears, rotations, and projections) that captures both their qualitative and, in certain cases, exact quantitative behaviour. ",
        "introduction": "The importance of understanding the physics of shear flows has been appreciated for well over a century, starting with the pioneering work of \\citet{helmholtz1868} and \\citet{kelvin1871}. The stability of these flows in a wide variety of situations has been thoroughly studied and reviewed, most notably in Chandrasekhar's (1961) classic text. Since that time, observations have revealed that shear flows are commonplace in astrophysical systems, playing likely roles in the stability and collimation of jets \\citep{ftz81,fj84,bbr84}, the solar corona \\citep{kads93}, bipolar outflows from young stellar objects \\citep{pringle89,bachiller96}, rotating stars \\citep{cowling51,gs67}, accretion discs \\citep{pl95,bh98}, and even the Earth's magnetopause \\citep{mckenzie70}. In nearly all of these systems, magnetic fields play a vital role in determining the structure and evolution. There are, however, systems where the importance of magnetic fields remain unclear, due to poor ionization and therefore weak coupling between the dominant ion and neutral species and the magnetic field. These include molecular clouds and their cores \\citep{mc99}, galactic molecular discs \\citep{bb94}, protostellar accretion discs \\citep{gammie96,hs98,ftb02}, protostellar outflows and disc winds \\citep{wk93}, dwarf nova disks \\citep{gm98,ss03}, shock waves in dense molecular clouds \\citep{wardle91,dmk93,rc07}, and protoplanetary discs \\citep{sm99,smun00,sw05,cmc07}. In these systems, various non-ideal magnetohydrodynamic (MHD) effects may come into play. For example, the neutral particles may drift relative to the ions in a process referred to as ambipolar diffusion. If, on the other hand, the ions drift relative to the electrons (and thereby the magnetic field), then the Hall effect occurs.  Given the high occurrence of shear flows and low levels of ionization in various astrophysical environments, it is not surprising that examples can be found when the two coincide. Perhaps the most obvious examples are outflows from star-forming regions, where shear instabilities can occur at the interface between the jet and the ambient material \\citep{wzhc04}. The resulting turbulent boundary layer can transfer linear momentum from the jet to the ambient medium, suggesting a possible mechanism for entrainment. Even at the launching point of these outflows, there may be differing velocity profiles in the ion and neutral fluids (see fig. 2 of Wardle \\& K\\\"{o}nigl 1993), conditions ripe for both ambipolar diffusion and shear. Another notable example concerns the interface between the magnetically-active and -inactive (`dead') regions in protoplanetary discs. It has been suggested that the well-coupled outer layers of protoplanetary discs will be subject to the magnetorotational instability (MRI; Balbus \\& Hawley 1991), while the shielded midplane layers will remain dormant \\citep{gammie96}. In this situation, angular momentum can be effectively transported in the outer layers but not in the midplane, setting up a velocity profile in the vertical (away from the midplane) direction \\citep{fs03,fn06}.  Here we undertake a study of a magnetized, planar shear layer in the presence of ambipolar diffusion and the Hall effect, and find that instability arises when shear and non-ideal MHD effects act in concert. The instabilities are similar to those found by \\citet{wardle99}, \\citet{bt01}, \\citet{kb04}, and \\citet{desch04} in the context of the MRI. Here we show that they are actually much more general. Our approach not only provides us with a clearer physical picture of the instabilities, free from the complications of rotation, but also brings qualitatively new results. In the case of ambipolar diffusion, instability arises from the combination of shear and the anisotropic nature of its wave damping. When the Hall effect is present, epicyclic-like motions set up by electromagnetic (whistler) waves couple to the background shear. When the handedness of these waves is opposite to that of the shear, instability occurs. It is notable that neither of these processes relies on rotational kinematics. In fact, these instabilities are neither versions of the Kelvin-Helmholtz instability \\citep{wzhc04}, nor versions of the MRI \\citep{bh91,bh92b}, despite their reliance on the presence of both shear and magnetic fields. The unstable modes are a result of the ability of non-ideal MHD processes to open new pathways for the fluid to tap into the free energy of shear.  An outline of the paper is as follows. In \\S \\ref{S_equations}, we discuss the formulation of the problem and give the basic equations to be solved. We then consider the evolution of Eulerian perturbations to these equations in comoving, local Lagrangian coordinates \\citep{glb65}. After deriving two coupled equations for the evolution of the relevant magnetic field eigenvectors in \\S \\ref{S_shearingsheet}, we solve them both analytically and numerically in several different situations. In \\S \\ref{S_AD}, we first restrict our attention to the effects of ambipolar diffusion on the stability of the system. After a detailed analysis and discussion, we then examine the effects of Hall electromotive forces (HEMFs) in Section \\ref{S_hall}. Contact with prior work is emphasized. In \\S \\ref{S_discussion}, we construct and analyse a toy model that captures all the salient features of these instabilities. Section \\ref{S_summary} summarizes our findings and conclusions. ",
        "conclusions": "\\label{S_discussion}  Despite the impression one may get from the abundance of mathematical manipulations in the preceding sections, the instabilities themselves are actually quite simple. They are the result of combinations of shears, rotations (in the case of the Hall effect), and projections (in the case of ambipolar diffusion). Much in the way that the MRI may be understood by two orbiting masses connected by a spring \\citep{bh92a}, there exists an equally intuitive toy model that captures the essence of the shear instabilities examined in this paper. Through simple matrix multiplication, several main results can be recovered without recourse to the lengthy mathematical manipulations undertaken in the preceding sections. All we require is some linear algebra.  Consider the following abstract problem, in which a position vector $|r_0\\rangle = |x_0,y_0\\rangle$ (in Dirac notation), having the coordinates $[x_0,y_0]$ in a two-dimensional $x$-$y$ Cartesian coordinate system, is subjected to various shear, rotation, and projection operators given by \\begin{equation} \\msb{S} \\equiv \\left[\\begin{array}{cc} 1 & 0 \\\\ \\varepsilon &  1 \\\\ \\end{array}\\right]\\,, \\qquad \\msb{R} \\equiv \\left[\\begin{array}{rr} \\cos\\theta & -\\sin\\theta \\\\ \\sin\\theta &  \\cos\\theta \\\\ \\end{array}\\right]\\,, \\qquad{\\rm and}\\qquad \\msb{P} \\equiv \\left[\\begin{array}{cc} \\sin^2\\phi & -\\sin\\phi\\cos\\phi \\\\ -\\sin\\phi\\cos\\phi & \\cos^2\\phi \\\\ \\end{array}\\right]\\,, \\end{equation} respectively. We denote by the column vector $|r_n\\rangle = |x_n,y_n\\rangle$ the position vector $|r\\rangle$ after $n$ transformations have been applied to $|r_0\\rangle$. To be precise, the combination $\\msb{S}|r\\rangle$ results in a shearing of $|r\\rangle$ along the $y$-axis by a distance $\\varepsilon x$; $\\msb{R}|r\\rangle$ results in a counter-clockwise rotation of $|r\\rangle$ through an angle $\\theta$; and, $\\msb{P}|r\\rangle$ projects $|r\\rangle$ onto the unit vector $|-\\sin\\phi,\\cos\\phi\\rangle$. These will be put in a physical context below.  We first turn our attention to the shear and projection operators. Applying the combination $\\msb{S}\\msb{P}$ (a projection followed by shear) to the initial state vector $|r_0\\rangle$ advances it to $|r_1\\rangle$: \\footnote{This requires $|r_0\\rangle\\ne|\\cos\\phi,\\sin\\phi\\rangle$, which corresponds to a zero eigenvalue of the operator $\\msb{SP}$. In this case, the projection gives the null vector, and there is nothing left to shear. For this particular initial state vector, the combination $\\msb{PS}$ works fine.} \\begin{equation} | r_1\\rangle = \\msb{SP}|r_0\\rangle = (x_0\\sin\\phi - y_0\\cos\\phi)\\left[\\begin{array}{c} \\sin\\phi\\\\ \\varepsilon\\sin\\phi-\\cos\\phi\\\\ \\end{array}\\right]\\,. \\end{equation} A graphical depiction of this process is given in Fig. \\ref{F_ambipicture}. While $|r_0\\rangle$ is clearly not an eigenvector of the operator $\\msb{SP}$, it is straightforward to show that $|r_1\\rangle$ {\\em is} an eigenvector, with eigenvalue $(1-\\varepsilon\\sin\\phi\\cos\\phi)$. It is then trivial to write down the result for general $n \\ge 1$: \\begin{equation}\\label{E_ambirecursive} | r_n\\rangle = (1-\\varepsilon\\sin\\phi\\cos\\phi)\\,|r_{n-1}\\rangle = (1-\\varepsilon\\sin\\phi\\cos\\phi)^{n-1}\\,| r_1\\rangle \\,. \\end{equation} Put differently, after this transformation has been performed just once, all subsequent applications simply stretch (or contract) the vector along $|r_1\\rangle$ by a factor equal to the associated eigenvalue. Note that in order for any evolution to occur, $\\phi\\ne m\\pi/2$ with $m$ an integer. For growth, we require \\begin{equation}\\label{E_projectioncriterion} \\varepsilon\\sin\\phi\\cos\\phi < 0\\,. \\end{equation} We now place these results in the context of the ambipolar-diffusion--shear instability. With $\\varepsilon = 2A\\Delta t$ and $|r\\rangle=\\delta\\bb{B}$, our shear operator corresponds physically to the production of $\\delta B_y$ due to the shearing of $\\delta B_x$ in a time $\\Delta t$. When $\\phi$ is taken to be the angle between a background magnetic field and the $x$-axis, the projection operator embodies ambipolar diffusion acting on the vector $|r\\rangle=\\delta\\bb{B}$. Their combination leads precisely to the behaviour described in \\S \\ref{S_behaviour}. In fact, it embodies the {\\em exact} solution for the case $\\bb{k}=k_z\\ez$ and $\\bb{B} = B\\cos\\phi\\ex + B\\sin\\phi\\ey$, as can be readily verified by taking the limit $\\Delta t\\rightarrow 0$ in Equation (\\ref{E_ambirecursive}) to arrive at the differential equation: \\begin{equation} \\frac{dr}{dt} = -A\\sin(2\\phi)\\,r\\,. \\end{equation} If $A\\sin(2\\phi)<0$, we have exponential growth. The growth rate is $A\\sin(2\\phi)$, which has its maximum at precisely the Oort $A$ value. That the maximum growth rate occurs at $\\phi=\\pi/4$ is in accord with the well-known ``$\\sin\\,2l$\" law (e.g., Mihalas \\& Binney 1981), a result of particular significance; we refer the reader to \\S 2.4 of \\citet{bh92a} for a full discussion of this topic.  Next we turn our attention to the shear and rotation operators. Applying the combination $\\msb{S}\\msb{R}$ (a rotation followed by shear) to the initial state vector $|r_0\\rangle$ advances it to $|r_1\\rangle$: \\begin{equation} | r_1\\rangle = \\msb{SR}|r_0\\rangle = (x_0\\sin\\theta - y_0\\cos\\theta)\\left[\\begin{array}{c} 1\\\\ \\varepsilon\\\\ \\end{array}\\right] + (x_0\\sin\\theta + y_0\\cos\\theta)\\left[\\begin{array}{c} 0\\\\ 1\\\\ \\end{array}\\right]\\,. \\end{equation} A graphical depiction of this process is given in Fig. \\ref{F_hallpicture}. For $n \\ge 2$, it is possible to show that a recursion relation exists between successive state vectors: \\begin{equation}\\label{E_hallrecursive} |r_n\\rangle - (2\\cos\\theta-\\varepsilon\\sin\\theta)|r_{n-1}\\rangle + |r_{n-2}\\rangle = 0\\,. \\end{equation} This equation has the general solution: \\begin{equation} |r_n\\rangle = \\left(\\frac{\\lambda^n_+ - \\lambda^n_-}{\\lambda_+ - \\lambda_-}\\right)\\,|r_1\\rangle - \\left(\\frac{\\lambda^{n-1}_+ - \\lambda^{n-1}_-}{\\lambda_+ - \\lambda_-}\\right)\\,|r_0\\rangle\\,, \\end{equation} where \\begin{equation} \\lambda_\\pm = \\left(\\cos\\theta - \\frac{\\varepsilon}{2}\\sin\\theta\\right) \\pm \\left[\\left(\\cos\\theta - \\frac{\\varepsilon}{2}\\sin\\theta\\right)^2 - 1\\right]^{1/2} \\end{equation} are the characteristic roots of Equation (\\ref{E_hallrecursive}). With $\\theta = \\omega\\Delta t$ and $|r\\rangle=\\delta\\bb{B}$, this transformation corresponds to the physical behaviour described in \\S \\ref{S_hall} concerning the evolution circularly-polarized electromagnetic waves (with frequency $\\omega$) in the presence of shear; namely, in a time $\\Delta t$, the perturbed magnetic field vector is rotated by $\\msb{R}$ through an angle $\\omega\\Delta t$ and sheared by $\\msb{S}$ along the $y$-axis. Taking the $\\Delta t \\rightarrow 0$ limit of Equation (\\ref{E_hallrecursive}) gives the differential equation \\begin{equation} \\frac{d^2 r}{dt^2} = -\\omega(\\omega+2A)\\,r\\,, \\end{equation} whose solutions are exponentially growing if \\begin{equation} \\frac{2A}{\\omega} < -1\\,. \\end{equation} The similarity between this instability criterion and Equation (\\ref{E_hallcrit}) is striking. When $|A/\\omega|$ equates to unity, the growth rate attains its maximum value: precisely the Oort $A$ constant. This simple model captures all the salient features of the Hall--shear instability.  \\begin{figure} \\begin{center} \\includegraphics[angle=-90,width=6.9in,clip]{kunz07_fig7.eps} \\end{center} \\caption{Evolution of the initial state vector $|r_0\\rangle$ (solid arrow) under the transformation $\\msb{SP}$. (a) A projection is applied to $|r_0\\rangle$ that retains only its component lying along the unit normal $|-\\sin\\phi,\\cos\\phi\\rangle$. The result is $\\msb{P}|r_0\\rangle$ (grey arrow). (b) This vector is then sheared along the $y$-axis into $|r_1\\rangle$ (open arrow). This process represents the ambipolar-diffusion--shear instability (see text for details). Figure is drawn to scale for comparison with Fig. \\ref{F_hallpicture}.}\\label{F_ambipicture} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[angle=-90,width=6.9in,clip]{kunz07_fig8.eps} \\end{center} \\caption{Evolution of the initial state vector $|r_0\\rangle$ (solid arrow) under the transformation $\\msb{SR}$. (a) A counter-clockwise rotation by $\\theta$ is applied to $|r_0\\rangle$, taking it into $\\msb{R}|r_0\\rangle$ (grey arrow). (b) This vector is then sheared along the $y$-axis into $|r_1\\rangle$ (open arrow). The process represents the Hall--shear instabiliy (see text for details). Figure is drawn to scale for comparison with Fig. \\ref{F_ambipicture}.}\\label{F_hallpicture} \\end{figure}    In this paper, we have investigated the stability of weakly-ionized, magnetized planar shear flows to linear disturbances. Employing a local approach similar to the shearing-sheet approximation of \\citet{glb65}, we have derived two coupled differential equations governing the evolution of magnetic field perturbations in the presence of either ambipolar diffusion or the Hall effect. Solutions are found by WKB methods and by direct numerical integration. We find that instability arises from the combination of shear and non-ideal MHD processes, and is a result of the ability of these processes to open new pathways for the fluid to feed off the free energy of shear. They turn what would be simple linear-in-time growth due to current and vortex stretching from shear into exponential instabilities. We have also constructed a simple toy model based on transformation operators that not only captures all the qualitative results of this paper, but also matches the exact {\\em quantitative} solution in some specific instances.  In the case of ambipolar diffusion, anisotropic damping leads to the generation of magnetic field perturbations perpendicular to the background magnetic field. What ensues is a competition between ambipolar diffusion and shear, which is trying desperately to stretch the magnetic field along the stream-wise direction. In the end they both win, and the perturbations grow in a direction somewhere between the two that depends upon the geometry of the magnetic field and the ratio of the neutral-ion collision time-scale to the shearing time-scale. The resulting growth rates are on the order of $0.1\\,|2A|$ (in the case of a time-independent background). It is notable that, in the general case of a time-dependent background, the growth rates increase in time without bound since highly-trailing (or -leading) shearing waves are trivially unstable. If the exponential growth phase lasts sufficiently long, growth rates may become comparable to or even exceed the shearing rate.  In the case of the Hall effect, it is not current damping but rather current generation that gives rise to a magnetic field component perpendicular to the shear. Instability arises from the influence of shear on the propagation of circularly-polarized electromagnetic (whistler) waves, and is present so long as the shearing frequency $|2A|$ is larger than the ion cyclotron frequency (times a factor proportional to the degree of ionization). Growth rates are on the order of $0.5\\,|2A|$. In contrast to the work of \\citet{wardle99} and \\citet{bt01}, we find that instability depends not on $\\bb{\\Omega}\\bcdot\\bb{B}$ or $(\\bb{k}\\bcdot\\bb{B})(\\bb{k}\\bcdot\\bb{\\Omega})$, respectively, but rather on $(\\bb{k}\\bcdot\\bb{B})(\\bb{k}\\bcdot\\bb{\\Gamma})$, where $\\Gamma=\\del\\btimes\\bb{v}$ is the vorticity of the background. This should be negative for destabilization. Provided a given wavevector begins its evolution unstable (stable), it will remain unstable (stable) until non-linear processes intercede. The physical reason that typical growth rates for the Hall--shear instability are so much greater than the maximum growth rate for the ambipolar-diffusion--shear instability is that the Hall--shear instability employs a conservative process (cyclotron gyrations) rather than a dissipative process (ambipolar diffusion). The difference in these rates is related to the rate at which ion-neutral friction heats the gas.  In both cases, unstable wavenumbers can be found for any sign of the velocity gradient. This explains why these processes were found to destabilize both inwardly- and outwardly-decreasing angular velocity gradients in accretion disks \\citep{bt01,kb04}. While the analysis is complicated somewhat by the time-dependence of the shearing background, we find that the evolution unfolds as a series of time-independent problems, similar to studies of nonaxisymmetric instabilities in discs (e.g., Balbus \\& Hawley 1992b, Balbus \\& Terquem 2001). In the limit of small horizontal ($y$) wavenumber, $k_y/k_z \\ll 1$, the instability criterion may be written \\begin{equation} \\kdotva + k^2\\bb{\\eta}\\,\\bb{:}\\,\\del\\,\\bb{v} < 0\\,, \\end{equation} with the resistivity tensor $\\bb{\\eta}$ being given by either Equation (\\ref{E_eta}) (for ambipolar diffusion) or Equation (\\ref{E_eta2}) (for the Hall effect). The final term here is the key to instability; in formal Cartesian index notation $(i,j,k)$, it is $k^2\\eta_{ij}\\partial v_j/\\partial x_i$. The maximum growth rate for a given $\\bb{\\eta}$ is \\begin{equation} \\sigma_{\\rm max} = \\left|\\frac{\\bb{\\eta}\\,\\bb{:}\\,\\del\\,\\bb{v}}{2\\det^{1/2}(\\bb{\\eta}) +\\,{\\rm tr}(\\bb{\\eta})}\\right|\\,. \\end{equation} For both non-ideal MHD effects, this is independent of the degree of ionization. Off-diagonal elements in the resistivity tensor are essential. No shear instabilities are present for isotropic damping processes, such as Ohmic dissipation.  The impact of these results on astrophysical systems is unclear at this point. Our neglect of rotation precludes a straightforward application to accretion discs. There is a fundamental difference between planar shear flows and disc systems: in a planar shear flow there is only one characteristic gradient, $dv/dx$, whereas in a disk system there are two, one for the angular velocity, and one for the angular momentum. In astrophysical discs, the omitted Coriolis force generally dominates the shear dynamics. There is no asymptotic domain for either linear or non-linear perturbations in which the governing dynamical equations behave locally like Cartesian shear. It seems that the best we can do here is to claim linear instability in constant-specific-angular-momentum discs. In addition, the well-known fact that inviscid planar shear layers are hosts to a plethora of non-linear instabilities suggests that the non-ideal MHD effects investigated in this paper may play only secondary roles.  This may be undue pessimism, however. The instabilities investigated in this paper {\\em do} have rotational counterparts, even in Keplerian systems. The real utility of this calculation, therefore, is that by removing rotation from the problem we obtain a clearer physical picture of what is going on in actual discs. One important simplification is that the MRI is absent. This is expected, of course, since the MRI is not {\\em just} a shear instability. Rather, the MRI plays the crucial role of redistributing angular momentum and thereby opening paths to lower energy states in differentially-rotating discs. On the other hand, the ambipolar-diffusion-- and Hall--shear instabilities have no preference for the source of the shear, whether it be in linear or angular velocity. The extent to which they lead to enhanced angular momentum transport is not governed by magnetic stresses, as in the case of the MRI, but rather whether or not the bulk fluid is responsive enough to utilize an increasingly radial magnetic field."
    },
    "0801/0801.3709_arXiv.txt": {
        "abstract": "We outline the steps needed in order to incorporate the evolution of single and binary stars into a particular Monte Carlo code for the dynamical evolution of a star cluster. We calibrate the results against $N$-body simulations, and present models for the evolution of the old open cluster M67 (which has been studied thoroughly in the literature with $N$-body techniques). The calibration is done by choosing appropriate free code parameters. We describe in particular the evolution of the binary, white dwarf and blue straggler populations, though not all channels for blue straggler formation are represented yet in our simulations. Calibrated Monte Carlo runs show  good agreement with results of $N$-body simulations not only for global cluster parameters, but also for e.g. binary fraction, luminosity function and surface brightness. Comparison of Monte Carlo simulations with observational data for M67 shows that is possible to get reasonably good agreement between them. Unfortunately, because of the large statistical fluctuations of the numerical data and uncertainties in the observational data the inferred conclusions about the cluster initial conditions are not firm. ",
        "introduction": "The modelling of individual globular clusters has a long history (see \\cite{MH1997}, especially \\S\\S 3 and 7.7).  Much of the focus of this work is on static models such as the King model and its variants.  In this kind of modelling the dynamical history of the cluster is almost irrelevant, except for the general assumption that the cluster is almost relaxed.  By contrast, there have been a small number of studies based on techniques which can follow the dynamical evolution of a cluster.  Most of this work has been performed with a Fokker-Planck scheme using finite differences, but also there are examples of the use of fluid and Monte Carlo methods  (Tab. \\ref{table:history}).  \\begin{table*} \\begin{minipage}{120mm} \\caption{Dynamical evolutionary models of individual globular clusters} \\begin{tabular}{lll} \\hline Cluster&Method&Reference\\\\ \\hline 47 Tuc&Fokker-Planck&\\citet{behleretal2003}, \\citet{murphyetal1998}\\\\ M15&Fokker-Planck&\\citet{murphyetal2003,murphyetal1997,murphyetal1994}, \\citet{dulletal1997}, \\citet{grabhornetal1992}\\\\ M30&Fokker-Planck&\\citet{howelletal2000}\\\\ N6397&Fokker-Planck&\\citet{dulletal1994}, \\citet{drukier1993, drukier1992}\\\\ N6624&Fokker-Planck&\\citet{grabhornetal1992}\\\\ $\\omega$ Cen&Monte Carlo&\\citet{gh2003}\\\\ M3&Moment equations&\\citet{angelettietal1980} \\end{tabular}\\label{table:history} \\end{minipage} \\end{table*}   In the present paper we develop the Monte Carlo technique further, and apply it to a new object.  The dynamical ingredients of the Monte Carlo code are essentially the same as those described in \\citet{giersz2006}, whose code embodies several features introduced by \\citet{stod1986}, whose code was in turn based on that originally devised by \\citet{henon1971}.  Three main features distinguish the code which is described in the present paper from that used by \\citet{gh2003} in their work on $\\omega$ Cen: (i) it now incorporates dynamical interactions between binary and single stars, between pairs of binaries, and interactions of three single stars resulting in the creation of new binaries, all using cross sections; (ii) it replaces the skeletal approach to stellar evolution taken from \\citet{cw1990} by the algorithms of \\citet{hurleyetal2000} for the evolution of single stars, supplemented by the methods of \\citet{hurleyetal2002} for the internal evolution of binary stars; (iii) a better treatment of the escape process in the presence of a static tidal field according to the theory proposed by \\citet{baum2001}.   There are several factors which motivate this work. Star clusters are the focus of several intensive observational campaigns \\citep[e.g.][]{bedinetal2001,bedinetal2003,grindlayetal2001,piottoetal2002, kaliarietal2003,kafkaetal2004,richeretal2004,andersonetal2006}, which are now turning to an examination of the parameters of their populations of binaries and blue stragglers (BS).  Dynamical models are needed for the design and interpretation of observational programmes: how is the period distribution and the spatial distribution of binaries affected by dynamical evolution?  Another problem is the abundance and spatial distribution of blue stragglers, which can only be answered by a technique which follows simultaneously both their dynamics and internal evolution.   While $N$-body techniques may ultimately be the method of choice for such studies, systems with the size of a globular cluster are likely to remain beyond reach for some years, simply because of the number of stars and the population of binaries.  After all, it is only recently that the ``hardest\" open clusters have been modelled at the necessary level of sophistication, and even then the typical simulation takes one month \\citep{hurleyetal2005}.  These authors focused on the old open cluster M67, which has been chosen by the MODEST international collaboration (MOdelling DEnse STellar systems) \\citep{sillsetal2003} as a target cluster for comparison between observations and various techniques of numerical simulation. We also focus on this cluster, partly for the purpose of refining our calibration of the Monte Carlo method.  This paper begins in Sec.2 with a summary of the features which have been added to the Monte Carlo scheme.  We also show there how we calibrate the Monte Carlo technique with $N$-body simulations.  Next (Sec. 3) we apply the technique to construct a dynamical evolutionary model of the old open cluster M67, and compare our results with observations. We give predictions for the initial parameters of the old open cluster M67. The final section summarises our conclusions, and discusses some of the main limitations of our models. ",
        "conclusions": "In this paper we have presented an advanced Monte Carlo code for the evolution of rich star clusters, including most aspects of dynamical interactions involving binary and single stars, and the internal evolution of single and binary stars. It was shown that the free parameters of the Monte Carlo code can be successfully calibrated against results of small $N$-body simulations and simulations of the old open cluster M67.  Sec.\\ref{sec:refinement} described our best models of M67. The results  show that an equally good fit to the observational data can be achieved by models which differ substantially in some initial parameters, such as the slope of the IMF. By contrast it seems that the other parameters are better constrained, at least within the ensemble of models which we studied. Actually, none of the models can successfully fit all the observational properties of M67 that we have studied, but we have argued that the remaining mismatches can be understood in terms of known characteristics of the Monte Carlo method, or the observational problem of subtracting the background.  These difficulties are especially pronounced at the bright and faint ends of the luminosity function. The most satisfactory models are characterised by the following initial parameters: $N$ about 30000, $r_t$ about 33 pc, $f_b$ about 50\\% and $r_t/r_h$ about 10. The word ``about'' is used deliberately, because of the large effect of statistical fluctuations in such small systems.  It is worth noting that a satisfactory fit can be achieved for a large range of values of the power-law index of the IMF for low mass stars, though it seems that values of $\\alpha_{IMF}$ in the range 0.5 -- 0.7 give slightly better agreement with observations.  Finally, these Monte Carlo simulations clearly show the strong influence of statistical fluctuations on the observational properties of a cluster. In consequence it seems that recovery of the initial cluster parameters from a comparison between numerical models and observations may be very difficult or even impossible, particularly for models with low or moderate $N$. More observational data is needed to constrain the models better, but the data on such properties as the number of particular kinds of binaries, pulsars or blue stragglers seem to provide only very weak constraints.  Despite these successes in fitting $N$-body models and the open cluster M67, the code has some known shortcomings, which we summarise here.  \\begin{enumerate} \\item {\\sl Cross sections}:  it has become apparent \\citep{fr2007} that the use of cross sections can lead to some systematic errors in the evolution of core parameters and other quantities.  Within the context of the Monte Carlo code we are working on, it is known how to replace these by explicit numerical calculation of the interactions \\citep{gs2003}, and this will be our next improvement. One side effect of the current absence of explicit interactions is that we do not yet model collisions which occur during them;  therefore one channel for the formation of blue stragglers is missing from these simulations. \\item{\\sl Higher-order multiples}:  It is widely argued that primordial triples and higher multiples should be incorporated into simulations along with primordial binaries. In any case, hierarchical triples form abundantly in binary-binary interactions \\citep{mikkola1984}. Such higher-order multiples are ignored in the present Monte-Carlo code, as cross sections for interactions with other objects have not yet been devised. Hierarchical triples and higher-order multiples can be introduced as new species when explicit calculation of interactions has been incorporated. \\item{\\sl Escape}: the Monte Carlo code described here incorporates a tidal cutoff, and a simple modification based on the theory devised by \\citet{baum2001}. Other treatments are possible and worth trying. \\item{\\sl Rotation}: the Monte Carlo code is based on spherical symmetry, and would require rather fundamental and very difficult reconstruction in order to cope with cluster rotation. As was pointed out by \\citet{kimetal2004} the rotation only somewhat accelerates the rate of core collapse. \\item{\\sl Static tide}: the effect of tidal shocks have been extensively studied (e.g. \\citet{ko1995}) and it would be possible to add the effects as another process altering the energies and angular momenta of the stars in the simulations. The addition of tidal shocks will be more important when modelling Galactic globular clusters than open clusters, which usually are confined inside the Galactic disk. \\end{enumerate}  Despite these limitations, some of which are difficult to cure, the Monte Carlo model presented in this paper shows its potential power in simulations of star clusters, from open clusters to rich globular clusters. Monte Carlo models are feasible in a reasonable time (a week or so) for globular clusters, which are too large for direct $N$-body models, and future papers in this series will present results on  M4 and several other globular clusters. The data provided by Monte Carlo simulations are as detailed as those  provided by an $N$-body code. No available simulation methods, except Monte Carlo and $N$-body methods, can really provide that kind of comprehensive information. Even when $N$-body simulations eventually become possible, Monte Carlo models will remain as a quicker way of exploring the parameter space for the large scale $N$-body simulations."
    },
    "0801/0801.4403_arXiv.txt": {
        "abstract": "We use the \\emph{Spitzer} Space Telescope to search for infrared excess at 24, 70, and 160\\,$\\mu$m due to debris disks around a sample of 45 FGK-type members of the Hyades cluster. We supplement our observations with archival 24 and 70\\,$\\mu$m \\emph{Spitzer} data of an additional 22 FGK-type and 11 A-type Hyades members in order to provide robust statistics on the incidence of debris disks at 625\\,Myr of age, an era corresponding to the late heavy bombardment in the Solar System. We find that none of the 67 FGK-type stars in our sample show evidence for a debris disk, while 2 out of the 11 A-type stars do so. This difference in debris disk detection rate is likely to be due to a sensitivity bias in favor of early-type stars. The fractional disk luminosity, $L_{DUST}/L_*$, of the disks around the two A-type stars is $\\sim$4$\\times$10$^{-5}$, a level that is below the sensitivity of our observations toward the FGK-type stars. However, our sensitivity limits for FGK-type stars are able to exclude, at the 2-$\\sigma$ level, frequencies higher than 12\\% and 5\\% of disks with  $L_{DUST}/L_* > $1$\\times$10$^{-4}$  and $L_{DUST}/L_* > $5$\\times$10$^{-4}$, respectively. We also use our sensitivity limits and debris disk models to constrain the maximum mass of dust, as a function of distance from the stars, that could remain undetected around our targets. ",
        "introduction": "Soon after IRAS discovered cold circumstellar disks around main-sequence (MS) stars \\citep{AuBeGi84}, it was realized that these disks could not be made of primordial material. Followup CO observations \\citep[e.g.][]{YaHaOm93} showed that molecular gas was highly depleted around these disks. Since, in the absence of gas, the survival time of dust due to dissipation processes such as radiation/wind pressure and the Poynting-Robertson effect is much shorter than the ages of MS stars, these systems are believed to be debris disks where dust is continuously replenished by collisions between planetesimals, the building blocks of planets. Because of their probable connection with the formation of planetary systems, debris disks  rapidly became  the subject of many studies. However, IRAS was only sensitive enough to study bright nearby objects and most of the pre-\\emph{Spitzer} statistics come from surveys performed by IRAS's successor, the ISO satellite. \\citet{HaDoJM01} studied 84 nearby ($\\mathrm{d} < 25$\\,pc) A,F,G, and K stars for which ISO was sensitive to photospheric fluxes and detected 60\\,$\\mu$m excess in $\\sim$50\\% of the stars younger than 400\\,Myr and in 10\\% of the stars older than 400\\,Myr. They suggest that this sudden decrease in the fraction of stars with disks around 400\\,Myr is related to the lifetime of planetesimals that replenish the dust. \\citet{SpSaSi01} observed $\\sim$150 pre-main sequence stars and young main sequence stars and detected 60\\,$\\mu$m excess in $\\sim$ 25\\% of the objects. Their observations do not confirm a sudden decrease in the disk fraction around 400\\,Myr but rather suggest a power law relationship (index $\\sim$ --2) between the age of the star and the fractional dust luminosity, $L_{DUST}/L_*$. They argue that such a power law naturally arises in collisionally replenished debris disks.  The discrepancies in the results from these ISO surveys can probably be traced back to the different target selection criteria and observing strategies and to small number statistics.  However, since conclusions from \\citet{HaDoJM01} and  \\citet{SpSaSi01} were not totally consistent, ISO was unable to provide a clear picture for the evolution of debris disks.  Fortunately, \\emph{Spitzer}'s unprecedented sensitivity has recently allowed many studies of large samples of pre-MS and MS stars in the mid- and far-IR. These studies are rapidly providing important clues on the evolution of debris disks. Rieke et al. (2005) studied over 250 A-type stars and concluded that, even though the incidence of 24 $\\mu$m excess clearly decreases with increasing stellar age, a very large dispersion on the magnitude of the IR-excesses is seen at every age. Su et al. (2006) finds that A-type stars show a similar qualitative behavior at 70 $\\mu$m. However, at this wavelength, the IR excess declines more slowly with time than at 24 $\\mu$m.  The decay timescale is estimated to be $\\sim$150 Myr for the 24 $\\mu$m  excess and $\\sim$850 Myr for the 70 $\\mu$m excess. Trilling et al. (2007) studied field FGK-type stars and found that $\\sim$4$\\%$ and $\\sim$16$\\%$ of them  show detectable IR excesses at 24 and 70$\\mu$m, respectively. These disk frequencies are roughly a factor of two lower than those found for A-type stars at the same wavelengths. They also find that solar-type stars have an almost constant disk frequency beyond 1 Gyr and hence conclude that the their debris disk decay timescale is significantly larger than for A-type star.  Gautier et al. (2007) studied a sample of nearby M-type stars but detected no significant far-IR excesses. However, given the sensitivity of their survey, they conclude that the average MIPS excesses, measured in photospheric flux units, of the M-type stars are at least a factor of four lower than those of solar-type stars  The \\emph{Spitzer} study of the Hyades presented in this paper is intended to provide  additional clues by giving robust statistics on the frequency of debris disks at 625\\,Myr of age for a homogeneous sample of MS stars. At 46\\,pc, the Hyades is the nearest star cluster to the Sun, and represent a sample of stars formed at the same epoch with the same heavy element abundance ($[\\rm{Fe/H}] = 0.13 \\pm 0.01$) \\citep{PaSnCo03}.  The 625\\,Myr age of the Hyades places them at an extremely interesting era in the evolution of planetary systems. This age corresponds almost exactly to the era of the late heavy bombardment (LHB) in our Solar System about 3.9\\,Gyrs ago \\citep{TePaWa73}. The cratering record of the Moon, Mars, and Mercury all indicate that the inner planets experienced intense bombardment by large bodies at that time. There is still intense debate as to whether the LHB represented merely the end of an exponential decrease in the impact rate from the formation of the terrestrial planets \\citep{We75,We77,NeIv94}, or was instead a short intense spike in the bombardment rate when the Solar System was about 600\\,Myr old \\citep{Ry90,Cohen02}. In either case, the LHB of our Solar System clearly indicates that at an age of $\\sim$600 My there was still a major debris disk present that was undergoing a rapid evolution.  Large, asteroid-size bodies had been built up during the early planet building era, but not all of these bodies had been incorporated into the planets. At the time of the LHB,  these bodies were undergoing an era of significant collisions with the inner planets, and presumably with each other as well. These collisions would have generated large amounts of smaller particles, ranging all of the way down to dust particles. If most planetary systems go through a LHB-type event at a similar age, it is quite reasonable to assume that other stellar systems might have also similar remnant debris disks at $\\sim$600 Myr of age. However, it has been recently argued that events such that of the LHB can be triggered by sudden dynamical interactions between planets after a long quiescent  period of time (Gomes  et al. 2005).  In that case, the incidence of debris disks on Hyades members is not expected to be significantly different from that of stars that are somewhat older (or younger) than them.  Here, we analyze deep MIPS  24 and 70 $\\mu$m observations for a sample of 78 Hyades stars, enough to provide robust statistics on the status of debris disks at 625\\,Myr of age. In Section ~\\ref{oba_data},  we describe the sample of Hyades stars, our observations, and the data reduction procedures. In Section ~\\ref{results}, we establish our disk identification criteria and present our detection statistics. In Section ~\\ref{Discussion}, we compare our detection statistics to recent \\emph{Spitzer} results, use our sensitivity limits and debris disk models to constrain the maximum mass of dust that could remain undetected around our targets, and discuss the implications of our results for debris disk evolution models and the late heavy bombardment in the Solar System. ",
        "conclusions": "\\label{Discussion}  \\subsection{Comparison to Recent Spitzer Results}  Recent \\emph{Spitzer} surveys have provided robust statistics on the debris disk frequencies around nearby stars against which our results can be compared. In order to make more meaningful comparisons, we divide the sample into FGK-type stars (67 objects) and A-type stars (11 objects). There are two motivations for doing so. First, most of the previous studies are restricted to either one of these groups. Second, given the strong luminosity dependence on spectral type, the 70\\,$\\mu$m observations are sensitive to much smaller 70\\,$\\mu$m excesses (in units of photospheric fluxes) and  $L_{DUST}/L_*$ values for A-type stars than for FGK-type stars.  \\subsubsection{FGK-type vs A-type stars}\\label{A-type}  To estimate the sensitivity difference between the 70\\,$\\mu$m observations of A-type stars and FGK-type stars, we calculate the ratio of 5 times the flux uncertainties to the estimated photospheric values (from Table~\\ref{tab:MIPS}). A cumulative histogram of this ratio is shown in Figure~\\ref{fig:70micronCDF} for FGK and A-type stars. For A-type stars, the 70\\,$\\mu$m observations can detect fluxes that are $\\sim$1-2$\\times$ those of the expected photospheres. However, for most of the FGK-type stars, the 70\\,$\\mu$m observations are only sensitive enough to detect fluxes that are $\\sim$15$\\times$ the expected photospheric values. As discussed in Section 3.2, this sensitivity limitation is mostly due to the fact that the stellar photospheres of solar-type stars, at the distance of the Hyades, fall below the 70 $\\mu$m extragalactic confusion limit for MIPS (1-$\\sigma$ $\\sim$2 mJy, Bryden et al. 2006).  The difference in 70\\,$\\mu$m sensitivity is even larger when it is calculated in terms of minimum detectable disk luminosity, $L_{DUST}/L_*$. Following Bryden et al. (2006), we calculate minimum disk luminosity as a function of 70\\,$\\mu$m excess flux by setting the emission peak at 70\\,$\\mu$m (T$_{DUST} = 52.5 K)$, according to:  \\begin{equation}\\label{FDL} \\frac{L_{DUST}}{L_*}(minimum) = 10^{-5}\\left(\\frac{5600 K}{T_{*}}\\right)^3\\frac{F_{DUST,70}}{F_{*,70}} \\end{equation} where $F_{DUST,70}$ is the flux of the dust and $F_{*,70}$ is the flux of the stars, both at 70\\,$\\mu$m. By setting $F_{DUST,70} = 5\\sigma_{70}-F_{*,70}$, we calculate the minimum $L_{DUST}/L_*$ values that are detectable for A-type and FGK-type stars. The results are shown in Figure~\\ref{fig:70micronCDF2}, which demostrates that the 70\\,$\\mu$m observations  of A-type stars are sensitive enough to detect disk with $L_{DUST}/L_*$ values in the 10$^{-6}$-10$^{-5}$ range. However for most of the FGK-type stars, the 70\\,$\\mu$m observations are only sensitive enough to detect disks with $L_{DUST}/L_{STAR} \\gtrsim$1-2$\\times$10$^{-4}$. We also use equation ~\\ref{FDL} to estimate $L_{DUST}/{L_*}$ values of 3.4$\\times$10$^{-5}$ and 4.7$\\times$10$^{-5}$ for the debris disks around HD28226 and HD28355, respectively. Since there are only 2 FGK-type objects for which the 70\\,$\\mu$m observations are sensitive enough to detect disks fainter than $L_{DUST}/{L_*} \\sim 4.0 \\times$10$^{-5}$, we conclude that the difference in the detection rate of debris disks around A-type stars and FGK-type stars is the result of a sensitivity bias rather than a real effect.  \\subsubsection{Comparison to FGK-type field stars}  \\citet{BrBeTr06} present 24 and 70\\,$\\mu$m observations for 69~FGK nearby (distance $\\sim$10-30\\,pc) field stars with a median age of $\\sim$4\\,Gyrs. They find 24\\,$\\mu$m excess around only one of their targets. This is consistent with the 24\\,$\\mu$m excess rate of 0$\\%$ we find for our 67 FGK-type stars. At 70\\,$\\mu$m, they identify 7 debris disks. This excess rate ($\\sim$10$\\%$), if taken at face value, seems inconsistent with our results. However, given the smaller distances to the stars involved, their survey was more sensitive than ours to faint disks. Since the stellar photospheres of their targets are above the  70 $\\mu$m confusion limit of MIPS, they were able to detect the stellar photosphere of most of them and identify very faint disks ($L_{DUST}/L_* \\geq 10^{-5}$). They also find that the disk frequency increases from 2$\\%\\pm2\\%$ for disks with $L_{DUST}/L_*\\geq 10^{-4}$ to $12\\%\\pm5\\%$ for disks with $L_{DUST}/L_* \\geq 10^{-5}$. Recent results by Trilling et al. (2008) confirm that the frequency of debris disks with $L_{DUST}/L_* \\geq 10^{-4}$ around solar type stars is $\\sim$2$\\%$.  Figure~\\ref{fig:70micronCDF2} shows that there are only 22 objects for which our 70\\,$\\mu$m observations are sensitive enough to detect a debris disk with $L_{DUST}/L_* \\geq 10^{-4}$. We use binomial statistics to show that if the incidence of disks brighter than $L_{DUST}/L_* = 10^{-4}$ is in fact 2\\% as found by \\citet{BrBeTr06} and Trilling et al. (2007), there was a 26$\\%$ probability that our survey would find zero disks. Also, given the cumulative distribution of sensitivities shown in Figure~\\ref{fig:70micronCDF2}, we use binomial statistics to calculate the minimum  disk frequencies, as a function of $L_{DUST}/L_*$, that would be \\emph{excluded} at the 1- and 2-$\\sigma$ level (i.e., the disk frequencies that would give our survey a 32\\% and 5\\% chance to result in zero detections). These  disk frequencies are tabulated in Table~\\ref{tab:disklimits}.  Based on the statistics for disks with $L_{DUST}/L_* \\geq 10^{-4}$, Table~\\ref{tab:disklimits} suggests that a debris disk fraction in the Hyades $\\sim$2.5 and $\\sim$6 times larger than in the field can be excluded at the 1-$\\sigma$ and 2-$\\sigma$ level, respectively.  Thus, we conclude that the debris disk fraction of the FGK-type Hyades stars (age $\\sim$625\\,Myr) is consistent with that in the field (age $\\sim$~4~Gyrs), but that $\\sim$6$\\times$ higher values cannot excluded from the currently available data.  \\subsection{Comparison to debris disk models}  In this section, we use the debris disk model developed by Augereau et al. (1999) to constrain the location and mass of the circumstellar material around HD 28266 and HD 28355, the two A-type stars identified in section ~\\ref{res_70} as having real 70 $\\mu$m excesses. We also use this model to estimate the maximum encompassed mass of dust, as a function of distance from the stars, that could remain undetected around the A-type and FGK-type stars in our sample that do not show significant IR excesses.  \\subsubsection{The debris disks around HD 28266 and HD 28355}  We limit the exploration of the parameter space to the disk parameters that affect most the global shape of an SED, namely the minimum grain size $a_{\\rm min}$, the peak surface density position $r_0$ and the total dust mass $M_{\\rm dust}$ (or, equivalently, the surface density at $r_0$). We adopted a differential grain size distribution proportional to $a^{-3.5}$ between $a_{\\rm min}$ and $a_{\\rm max}$, with $a_{\\rm max} = 1300\\,\\mu$m, a value sufficiently large to not affect the SED fitting in the wavelength range we consider. Following Augereau et al. (1999), the disk surface density $\\Sigma(r)$ is  parametrized by a two power-law radial profile $\\Sigma(r) = \\Sigma(r_0) \\sqrt{2} \\left(x^{-2\\alpha_{\\rm in}}+x^{-2\\alpha_{\\rm out}} \\right)^{-1/2}$ with $x = r/r_0$, and where $\\alpha_{\\rm in} = 10$ and $\\alpha_{\\rm out} = -3$ to simulate a disk peaked around $r_0$, with a sharp inner edge, and a density profile decreasing smoothly with the distance from the star beyond $r_0$. The optical properties of the grains were calculated for astronomical silicates (optical constants from Weingartner $\\&$ Draine (2001)), and with the Mie theory valid for hard spheres. The grain temperatures were obtained by assuming the dust particles are in thermal equilibrium with the central star. NextGen model atmosphere spectra (Hauschildt et al. 1999) scaled to the observed K-band magnitudes, were used to model the stellar photospheres.  For each of the stars with 70 $\\mu$m excess,  HD 28266 and HD 28355, we calculated $15000$ SEDs ($0.3\\,\\mu$m$\\leq \\lambda \\leq 950\\,\\mu$m), for $75$, logarithmically-spaced values of $a_{\\rm min}$ between $0.05\\,\\mu$m and $100\\,\\mu$m, and for $200$ values of $r_0$, logarithmically-spaced between $10$\\,AU and $500$\\,AU. For each model, the dust mass was adjusted by a least-squares method, assuming purely photospheric emission in the 2MASS bands and by fitting the measured MIPS $24\\,\\mu$m and $70\\,\\mu$m flux densities. The results are summarized in Table 5, and the SEDs are displayed in Figure~\\ref{SEDs}. Results in Table 7 are listed for two different regimes of minimal grain sizes, namely, $a_{\\rm min}$ $>$ 10 $\\mu$m and $a_{\\rm min}$ $<$ 0.5 $\\mu$m in order to illustrate the strong dependence of $r_0$ and M$_{dust}$ on the assumed grain size distribution.  Even though neither the position of the peak surface density $r_0$, nor the minimum grain size $a_{\\rm min}$, can be uniquely determined with so few observational constraints, some models can be eliminated. In particular, given the large luminosity of A-type stars, the small or nonexistent excess at 24 $\\mu$m implies that the disks of HD 28226 and HD 28355 are significantly dust-depleted within \\emph{at least} $\\sim$40 AU from the star. However, ``inner-holes'' larger than 300 AU can not be excluded if very small grains are present (i.e., a$_{\\rm min}$ $<$ 0.5 $\\mu$m). Similarly, even though the best-fit disk models imply dust masses of the order 10$^{-2}$M$_{\\oplus}$, disk masses $\\sim$0.1M$_{\\oplus}$ could be accomodated for both stars.  \\subsubsection{Limits on dust masses as a function of radius}  In section ~\\ref{A-type}, we found that the 70 $\\mu$m observations of the Hyades were sensitive to significantly lower fractional disk luminosities for A-type stars than for FGK-type stars. However, since the 70 $\\mu$m observations probe larger radii around A-type stars than around FGK-stars, this implies that the mass of the grains needed to produce a given fractional luminosity, as calculated in section~\\ref{A-type}, is \\emph{larger} around A-type stars than it is around FGK-type stars (see Gautier et al. 2007 for a discussion of origin of the dependence of dust mass sensitivity on stellar luminosity).  Therefore, the degree to which the MIPS observations are sensitive to smaller amounts of dust around A-type stars than around FGK-type stars is not immediately obvious. In order to explore this last point, we estimate the maximum encompassed mass of dust, as a function of distance from the stars, that could remain undetected around the A-type and FGK-type stars in our sample that do not show significant IR excesses.  Following Cieza et al. (2007), we use the optically thin disk models discussed above to constrain the maximum amount of dust that could be present within 300 AU of the A-type and FGK-type stars in our sample. Using the 70 $\\mu$m 5-$\\sigma$ upper limits (and the 160 $\\mu$m limits, when available), we calculated $15000$ models analogous to those calculated for HD 28266 and HD 28355 for each of the stars in our sample. For each model, we calculated the mass encompassed within a radius $r$, as a function of this radius. With this approach, we estimate the maximum dust mass in the circumstellar regions of the Hyades stars with no detectable emission in excess to the photospheric emission. The results for A-type and FGK-type stars, for two different minimal grain size regimes (a$_{\\rm min}$ $>$ 10 $\\mu$m in blue and a$_{\\rm min}$ $<$ 0.5 $\\mu$m in red), are shown in Figure \\,\\ref{MvsR}. This figure shows that the absence of IR excess at 24 and 70 $\\mu$m constrains the total mass of dust within 10 AU of the solar type stars to be $<$ 10$^{-4}$ M$_{Earth}$. This limit is an order of magnitude lower for A-type stars. The range of encompassed dust masses as a function of radius for the best-fit disk models of HD 28355 corresponding to the case where a$_{\\rm min}$ $>$ 10 $\\mu$m is shown for comparison (green region). The fact that, for FGK-type stars, the green region lies below the blue region (i.e., the case corresponding to the same grain size distribution), implies that disks similar to that found around HD 28355 would not be detectable by our observations if they were also present around the FGK-type Hyades members. Given the similarities of the inferred disk properties of HD 28355 and HD 28266 (see  Table 5), the same conclusion applies for HD 28266. Thus, Figure \\,\\ref{MvsR} strengthen our conclusion from section ~\\ref{A-type} stating that the difference in detection rate of debris disks around A-type and FGK-type stars is due to a sensitivity bias rather than to a real difference in the incidence or properties of debris disks around stars of different spectral types.   \\subsection{Debris disk evolution and the Late Heavy Bombardment}  \\subsubsection{Steady State vs. Stochastic Evolution}  \\citet{RiSuSt05} studied a sample of 266 A-type stars with \\emph{Spitzer}, ISO, or IRAS 24 $\\mu$m/25 $\\mu$m data. They used this very large sample to establish statistically significant trends of IR-excess with age. They find that: (1) at all ages, the population is dominated by stars with little or no IR excess, (2) stars with a wide range of excesses are seen at every age, and (3) both the frequency and the magnitude of the IR excess decreases with time. In particular, they find that the upper envelope of the evolution of the excess ratio with time can be fitted by $t_o/t$, with $t_o\\sim$150\\,Myr. Similar trends are seen in the 70\\,$\\mu$m excesses of A-type stars \\citep{SuRiSt06}, with the difference that the decay time seems to be considerably larger, $t_o \\gtrsim$400\\,Myr, suggesting inside-out disk clearing.  Based on these results, \\citet{RiSuSt05}  argue that the evolution of debris disks is the convolution of a stochastic and a steady component. They suggest that, at any given age, the debris disks detected are those that have experienced large planetesimal collisions in the recent past. This stochastic evolution is on top of steady decrease in the number of parent bodies in the belts  of planetesimals where the dust is produced, which would explain the overall decrease of IR excess with age. However, it has also been argued that a stochastic component in the evolution of debris disk is not necessary to explain the diversity of disk properties observed at a given age. \\citet{WySmSu07} construct a simple collisional model, where the mass of planetesimals is constant until the largest ones reach collisional equilibrium, at which point mass falls as 1/time. They propose that the large spread in IR properties observed at any given age can be explained in terms of the initial distributions of masses and temperatures of the planetesimal belts producing the dust. They argue that their simple model can account for the 24 and 70\\,$\\mu$m statistics presented by \\citet{RiSuSt05} and \\citet{SuRiSt06} using realistic belt parameters, and thus that transient events are not \\emph{required} to explained the observations. Given the limited observational constraints available, the models presented by \\citet{WySmSu07} do not rule out the possibility that stochasticity plays an important role in the evolution of most debris disks. Our results could provide additional constrains  to these kinds of models because, unlike the studies by \\citet{RiSuSt05} and \\citet{SuRiSt06}, our study  provides robust statistics for the debris disks at a single, well defined age.  \\subsubsection{Implications for the Late Heavy Bombardment in the Solar System}\\label{LHB}  The 625\\,Myr age of the Hyades corresponds almost exactly to the era of the late heavy bombardment \\citep[LHB,][]{TePaWa73,GoLeTs05}. Thus, \\emph{if} the Solar-type (FGK) Hyades stars resemble the Sun at 625\\,Myr of age, our statistics could provide valuable clues on this important event in the history of the Solar System.  The cause and the duration of the LHB is still a matter of debate. Proposed causes for a intense spike in the impact rate include the formation of Uranus and Neptune \\citep{LeDoCh01}, the presence of a fifth terrestrial planet in a low-eccentricity orbit which became dynamically unstable at an age of about 600\\,Myr \\citep{ChLi02}, or impacts by bodies left over from planetary accretion \\citep{MoPeGl01}. More recently, \\citet{GoLeTs05} propose that the LHB was triggered by the sudden migration of the giant planets that occurred after a long quiescent period of time.  In their model, soon after the dissipation of the solar nebula, the orbits of Jupiter and Saturn started to slowly diverge  due to the interaction with the massive disk of planetesimals that was still present.  They argue that $\\sim$700\\,Myr later, when Jupiter and Saturn crossed their 1:2 mean motion resonance, their orbits became eccentric and temporally destabilized those of Uranus and Neptune.  The reconfiguration of the orbits of the giant planets resulted in the perturbation and massive delivery of planetesimals to the inner Solar System, which according to their models, lasted between 10-150\\,Myr.  The observational signatures of a LHB-type event as seen from a distance of 46 pc are not known, but it as been suggested that they could be those of a family of rare Solar-type stars  characterized by the presence of a bright ``hot disk'' around an object hundreds of million years old. These objects present excess IR emission, originating in the terrestrial planet regions, with $F_{DUST}/F_* > 10^{-4}$, a level that is $>1000$ times larger than steady state evolution models can explain \\citep{WySmGr07}. There are currently only 5 known objects that fall into this category of ``hot transient disks'': BD~+20~307 \\citep[age $\\sim$300\\,Myr,][]{SoZuWe05}, HD72905 \\citep[age $\\sim$400\\,Myr,][]{BeBrSt06}), $\\eta$ Corvi \\citep[age $\\sim$1\\,Gyr,][]{WyGrDe05}, HD69830 \\citep[age $\\sim$2\\,Gyr,][]{BeBrRi05}, and $\\tau$ Ceti (age $\\sim$10 Gyr, Di Folco et al. 2007) which represent $\\sim$2$\\%$ of all the Solar-type stars surveyed. \\emph{If} this group of objects corresponds to those that are currently experiencing events similar to the LHB and the Hyades stars resemble the Sun at 625\\,Myr of age, then  the fact that none of the 67 Solar-type stars in our Hyades sample has  a ``hot transient disk'' implies one of two possibilities: (1) the likelihood of an  event similar to the LHB is not significantly higher at $\\sim$625\\,Myr than it is at any other age, or (2) events like the LHB are very short spikes with a duration much closer to the lower limit of 10\\,Myr suggested  by \\citet{GoLeTs05} than to their 150\\,Myr upper limit.  If the likelihood of a LHB-type event is approximately constant with time, then a 2$\\%$ incidence in the Solar neirghborwood (median age $\\sim$4000\\,Myr) would imply a total duration of $\\sim$80\\,Myr. However, if such an event is more likely to occur around an age of $\\sim$625\\,Myr, then our non detections would only be consistent with a much shorter duration. Thus, the implication of our results on the LHB could depend on the age distribution of these ``hot transient disks'', which still remains largely unconstrained since only 5 of such examples are currently known.  Fortunately, as more \\emph{Spitzer} observations are reported, this distribution will become better constrained.  Understanding the debris disk phenomenon has been a high priority of the \\emph{Spitzer}'s mission. As a result, the number of debris disk studies has increased dramatically over the last few years. Each of these studies is providing new clues and constraints, from which it will eventually emerge a much clearer picture of the evolution of debris disks and its connection to the history of the Solar System."
    },
    "0801/0801.4897_arXiv.txt": {
        "abstract": "%  We are using ``broadband'' (4.6 to 43 GHz) multi-frequency VLBA polarization observations of compact AGN to investigate the 3-D structure of their jet magnetic ({\\bf B}) fields. Observing at several frequencies, separated by short and long intervals, enables reliable determination of the distribution of Faraday Rotation, and thereby the intrinsic {\\bf B} field structure. Transverse Rotation Measure (RM) gradients were detected in the jets of 0954+658 and 1418+546, providing evidence for the presence of a helical {\\bf B} field surrounding the jet. The RM in the core regions of 2200+420 (BL Lac), 0954+658 and 1418+546 display different signs in different frequency-intervals (on different spatial scales); we suggest an explanation for this in terms of modest bends in a helical {\\bf B} field surrounding their jets. In future, polarization observations with a combination of VSOP-2 at 8, 22 and 43 GHz and ground arrays at frequencies with corresponding resolution will help map out the distributions of Faraday rotation, spectral index and the 3-D {\\bf B} field structure both across the jet and closer to the central engine, providing strong constraints for any jet {\\bf B} field models. ",
        "introduction": "%  AGN jets emit synchrotron radiation that is often highly linearly polarized, which provides important information on the degree of order and orientation of the magnetic ({\\bf B}) field in these jets. ``Blazars'' have jets pointed close to our line of sight (LoS) and often exhibit strong variability in total flux and linear polarization over a broad range of frequencies from $\\gamma$-ray to radio.  Recent MHD simulations have provided an almost complete explanation of how these jets are launched, accelerated and collimated close to the black hole, see \\cite{MeierJapan} and references therein. The dominant {\\bf B} field structure on these scales is helical. It is possible that this structure changes when the flow gets disrupted by shocks after a few hundred Schwarzschild radii, but observational evidence of helical fields on scales larger than this \\citep{Asada2002, GabuzdaMurray2004, Mahmud2008} suggests that remnants of the earlier {\\bf B} field structure remain or that a current driven helical kink instability is generated \\citep{NakamuraJapan, Carey2008}.   Faraday Rotation manifests itself as a linear dependence of the observed polarization angle ($\\chi_{obs}$) on the wavelength ($\\lambda$) squared described by the formula $\\chi_{obs}=\\chi_0+RM\\lambda^2$, where the Rotation Measure (RM) is proportional to the integral of the electron density and the dot product of the {\\bf B} field along the LoS with the path length $\\vec{dl}$. Hence, a positive/negative RM tells you that the LoS {\\bf B} field is pointing towards/away from the observer and combined with the corrected polarization orientation we are provided with a 3-D view of the {\\bf B} field structure. However, if the rotating medium is mixed in with the emitting plasma then internal Faraday rotation occurs and the above formula does not generally apply. Theoretical models \\citep{Burn1966, Cioffi1980} of internal FR in a spherical or cylindrical region have shown that if a rotation of greater than $45\\degr$ is observed then the Faraday rotation must be external.    \\begin{figure} \\begin{minipage}[t]{5.0cm} \\begin{center} \\includegraphics[width=5.0cm,clip]{osullivanFig1.eps} \\end{center} \\end{minipage} \\hfill \\begin{minipage}[t]{7.0cm} \\begin{center} \\includegraphics[width=6.5cm,clip]{osullivanFig2.eps} \\end{center} \\end{minipage} \\caption{Left: Plots of $\\chi$ vs. $\\lambda^2$ for the whole observed frequency range of the core regions of 1156+295 and 0954+658 clearly showing a transition in both cases where physically different regions are being probed. Right: For a ``head-on'' ($\\theta<1/\\Gamma$) or ``tail-on'' ($\\theta>1/\\Gamma$) view of a helical {\\bf B} field, different LoS components will dominate (shown by solid circles surrounding the dot or cross which indicate the direction of the LoS component) producing RMs with different signs in an unresolved jet.} \\end{figure} ",
        "conclusions": "Polarization observations with VSOP-2 have the potential to greatly improve our understanding of the compact inner regions of AGN jets. The excellent resolution provided by VSOP-2, combined with an array of ground telescopes, will enable multi-frequency Faraday rotation analysis much closer to the central engine. The higher resolution of VSOP-2 will also reduce the effect of beam depolarization so that new polarization could be detected in previously unpolarized observations of AGN cores.  Observations using the full power of VSOP-2 with dual polarization at 8, 22 and 43 GHz combined with matched-resolution ground array observations at multiple bands at 15, 22, 43 and 86 GHz will facilitate unparalleled analysis of jet physics in these compact regions. Information on the 3-D {\\bf B} field structure, electron density, 2-D internal structure and spectral index much closer to where these jets are being launched, accelerated and collimated will test current theoretical models of jet production and propagation through space."
    },
    "0801/0801.1156_arXiv.txt": {
        "abstract": "We present a high resolution dark matter reconstruction of the $z=0.165$ Abell 901/902 supercluster from a weak lensing analysis of the {\\it HST} STAGES survey. We detect the four main structures of the supercluster at high significance, resolving substructure within and between the clusters.  We find that the distribution of dark matter is well traced by the cluster galaxies, with the brightest cluster galaxies marking out the strongest peaks in the dark matter distribution. We also find a significant extension of the dark matter distribution of Abell 901a in the direction of an infalling X-ray group Abell $901\\alpha$.  We present mass, mass-to-light and mass-to-stellar mass ratio measurements of the structures and substructures that we detect.  We find no evidence for variation of the mass-to-light and mass-to-stellar mass ratio between the different clusters.  We compare our space-based lensing analysis with an earlier ground-based lensing analysis of the supercluster to demonstrate the importance of space-based imaging for future weak lensing dark matter `observations'. ",
        "introduction": "Observations and theory both point to the importance of environment on the properties of galaxies. Early-type galaxies are typically found in more dense regions compared to late-type galaxies \\citep{Dressler}, galaxy colour and luminosity are found to be closely related to galaxy density \\citep{Blanton05} and the fraction of star-forming galaxies also shows a strong sensitivity to the density on small $<1$ Mpc scales \\citep{Balogh04,Blanton06}. Theoretically there are a number of physical mechanisms that could cause these effects in dense environments.  These processes can change the star formation history, gas content and/or morphology of a galaxy through, for example, ram-pressure stripping \\citep{GG72,Larson,Balogh00} and/or the tidal effects of nearby galaxies \\citep[galaxy harassment,][]{Moore96} and/or the tidal effects of the dark matter potential \\citep{Bekki,Moore98}. They depend differently on cluster gas, galaxy density and the dark matter potential however, with ram-pressure stripping dependent on the gas distribution compared to tidal effects which are dependent on the overall potential. A key difficulty in disentangling these effects observationally, is that typically the tidal potential is only constrained in a global sense through the measured velocity dispersion of a cluster, or a richness or total luminosity estimate.  This results in an assumed spherical tidal potential model that is smoothed over the small scales that are relevant for tidal stripping and harassment studies.  In this paper we study the complex Abell 901/902 supercluster, hereafter A901/2, in the first of a series of papers from the STAGES\\footnote{`Space Telescope A901/902 Galaxy Evolution Survey' ({\\it HST} GO-10395, PI M. E. Gray), www.nottingham.ac.uk/$\\sim$ppzmeg/stages } collaboration. From a rich multi-wavelength dataset the A901/2 supercluster permits a thorough investigation of the relationships between galaxy morphology \\citep[from Hubble Space Telescope ({\\it HST}) and ground-based imaging,][] {GrayPeng,Lane}, luminosity, stellar mass and colour \\citep[from the COMBO-17 survey with 17-band optical imaging,][]{Wolf03, Borch}, star formation rates \\citep[from 24$\\mu$m {\\it Spitzer} data,][]{Bell07}, galaxy density \\citep{wgm,Gray04}, and the hot intra-cluster medium \\citep[from XMM observations,][]{Gilmour,Gray07}. One of the key reasons to obtain {\\it HST} imaging of this supercluster was to construct a high resolution, reliable and accurate map of the projected total mass density distribution. Using weak gravitational lensing techniques we are able to reconstruct the distribution of both dark and luminous matter and quantify the significance of the structures that are seen, updating the previous ground-based weak lensing analysis of \\citet{meg02}. This extra dimension to the multi-wavelength view of A901/2 will be a key ingredient in future studies where we hope to be able to separate the effects of tidal and gas-dynamical influence on galaxy formation and evolution.  Weak gravitational lensing is now a well established method for studying the distribution of dark matter.  Light from distant galaxies is deflected by the gravitational effect of the intervening structures, inducing a weakly coherent distortion in the shapes of galaxy images. The strength of this lensing effect is directly related to the projected mass along the line of sight, and it can therefore be used to map dark matter in dense regions \\citep[see for example][]{meg02,Gavazzi04,Dietrich,Clowe06,Mahdavi}.  The first weak lensing analysis of A901/2 by \\citet{meg02} used deep ground-based $R$-band observations from the COMBO-17 survey \\citep{Wolf03}. This analysis revealed three significant peaks in the dark matter distribution at the locations of the A901a, A901b and A902 clusters, in addition to a low significance south west peak co-incident with a galaxy group, hereafter referred to as the SW group. This analysis also showed a filamentary extension between the A901a and A901b clusters.  As this filament was located across the CCD chip boundary in the mosaic image, however, \\citet{meg02} could not rule out the possibility of this structure originating from residual uncorrected distortions from the point spread function (PSF) of the telescope and detector. The \\citet{meg02} ground-based analysis also reported a candidate giant arc. The STAGES {\\it HST} imaging can rule out this candidate arc as a co-incidental alignment of objects.  STAGES does however resolve several other candidate arcs around supercluster galaxies, which will be presented in \\citet{Aragon} and \\citet{GrayPeng}.  Using the accurate photometric redshift information from the A901/2 17-band observations of the COMBO-17 survey, where the photometric redshift error $\\sigma_z \\sim 0.02/(1+z)$ for $R<24$, \\citet{A9013D} extended the \\citet{meg02} analysis, by creating a three-dimensional reconstruction of the A901/2 dark matter distribution.  This analysis revealed a previously unknown higher redshift cluster located behind A902 that is at $z=0.46$.  This cluster was named, and hereafter referred to as, CB1 by \\citet{A9013D}. We have updated the redshifts of both CB1 and A901/2 in this analysis based on an improved photometric redshift catalogue and the addition of some spectroscopic redshifts.  In this analysis we revisit the dark matter distribution in A901/2 using deep {\\it HST} observations.  The dominant source of noise in the weak lensing analysis of clusters is the Gaussian noise introduced from the random intrinsic ellipticities of galaxies.  Weak lensing maps of dark matter on small scales therefore benefit greatly from the high resolution that {\\it HST} has to offer.  {\\it HST} triples the number density of resolved galaxies from which the lensing signal can be measured, reducing the intrinsic ellipticity noise on small scales. In addition, the high resolution space-based data permits higher signal-to-noise shape measurements and a narrower PSF, thus implying a more accurate PSF correction.  This paper is organised as follows.  In section~\\ref{sec:theory} we describe the weak lensing theory that is related to this analysis, and the maximum likelihood method that we use to reconstruct the dark matter distribution. We describe the data and weak lensing measurement method in section~\\ref{sec:data}.  We present our results in section~\\ref{sec:res}, including NFW profile mass measurements in section~\\ref{sec:nfwmass} and the dark matter reconstruction and ground-based comparison in section~\\ref{sec:maps}.  A first comparison of the dark matter, galaxy light and stellar mass distribution is presented in section~\\ref{sec:masslight} along with mass, mass-to-light  and mass-to-stellar mass ratio measurements. A more detailed comparison of the mass, gas and galaxies of A901/2 will appear in a forthcoming analysis. We investigate the significance of the supercluster substructure that is resolved in our dark matter reconstruction in section~\\ref{sec:peaksig} and discuss our findings and conclude in section~\\ref{sec:conc}.  Throughout this paper we assume a $\\Lambda$CDM cosmology with $\\Omega_m =0.3$, $\\Omega_\\lambda =0.7$, and ${\\rm H}_0 = 100 {\\rm h}\\, {\\rm km \\, s^{-1} Mpc^{-1}}$.  All magnitudes are given in the Vega system. ",
        "conclusions": ""
    },
    "0801/0801.3479_arXiv.txt": {
        "abstract": "Hypervelocity stars (HVS) traverse the Galaxy from the central black hole to the outer halo. We show that the Galactic potential within 200 pc acts as a high pass filter preventing low velocity HVS from reaching the halo. To trace the orbits of HVS throughout the Galaxy, we construct two forms of the potential which reasonably represent the observations in the range 5--10$^5$ pc, a simple spherically symmetric model and a bulge-disk-halo model. We use the Hills mechanism (disruption of binaries by the tidal field of the central black hole) to inject HVS into the Galaxy and to compute the observable spatial and velocity distributions of HVS with masses in the range 0.6--4 M$_\\odot$. These distributions reflect the mass function in the \\GC, properties of binaries in the \\GC, and aspects of stellar evolution and the injection mechanism. For 0.6--4 \\msun\\ main sequence stars, the fraction of unbound HVS and the asymmetry of the velocity distribution for their bound counterparts increases with stellar mass. The density profiles for unbound HVS decline with distance from the \\GC\\ approximately as $r^{-2}$ (but are steeper for the most massive stars which evolve off the main sequence during their travel time from the \\GC); the density profiles for the bound ejecta decline with distance approximately as $r^{-3}$. In a survey with limiting magnitude $V \\lesssim$ 23, the detectability of HVS (unbound or bound) increases with stellar mass. ",
        "introduction": "The components of the Milky Way sample the diverse forms of matter in the universe from the black hole at the dynamical center to the dark matter component of the halo. Hypervelocity stars (HVSs), a class of recently discovered objects, are a probe linking the central black hole and its history to the outer halo. HVSs probably originate from interactions with the black hole; they travel out into the halo and beyond, serving as tracers of the Galactic potential on scales from 5 pc to 10$^5$ pc.  To make testable predictions of the relationships between the observable properties of HVSs and aspects of the Galactic center and halo, we adopt the massive black hole (MBH)-binary star interaction first proposed by Hills \\citep{hil88} as the injection mechanism for HVSs. We then explore the impact of the Galactic potential, the properties of binaries in the Galactic Center, the mass function in the Galactic Center, and stellar evolution on the observable spatial and velocity distributions of HVSs as a function of stellar mass.   \\subsection {History and Challenges}  \\citet {hil88} showed that HVSs are inevitable when sufficiently tight binary stars are disrupted by the tidal field of the central MBH. Much later, \\citet{yu03}  analyzed the production of HVSs and argued that the Hills production rates were excessive.  They considered other mechanisms for the production of hypervelocity stars including interactions of single stars with a binary central black hole.  In 2005, \\citet{bro05} reported the first discovery of a HVS.  The star, SDSS J090745.0+024507, is leaving the Galaxy with a heliocentric radial velocity of 853$\\pm$12 km~s$^{-1}$ and has the largest velocity ever observed in the Milky Way halo. Subsequent photometry revealed that the object is a slowly pulsating B main-sequence star \\citep{fue06}. Only interaction with a massive black hole can plausibly accelerate a 3 M$_\\odot$ main-sequence B star to such an extreme velocity.  The \\citet{bro05} discovery inspired further observations along with a wealth of theoretical models. \\citet{ede05} discovered a hypervelocity main sequence B star plausibly ejected from the LMC \\citep[see also][] {bon07} and \\citet{hir05} reported a helium-rich subluminous O star probably ejected from Sgr A$^*$. These three first discoveries were serendipitous.  \\citet{bro06a} then carried out a targeted search for hypervelocity stars by using the SDSS to select faint B-type stars with lifetimes consistent with travel times from the Galactic center but that are not a normally expected stellar halo population. So far this extensive survey has yielded another 7 hypervelocity stars \\citep{bro06a, bro06b, bro07a, bro07b}. The velocity of each of these stars exceeds the escape velocity from the Milky Way.  The observational discoveries have inspired theorists to reconsider production mechanisms and ejection rates of hypervelocity stars. Ejection mechanisms include the original Hills binary encounter with a MBH \\citep{hil88, yu03, bkg06}, encounters between single stars and binary black holes \\citep{yu03, ses06, ses07a, ses07b, merritt06}, encounters of single stars with an intermediate mass black hole (IMBH) inspiraling toward a MBH \\citep{han03, gua05, levin06, bau06}, and single star encounters with a cluster of stellar mass black holes (SMBHs) around a MBH \\citep{ole08}. Recently \\citet{lu07} predicted hypervelocity ejection of tight binary stars in interaction with a central binary MBH. Discovery of a single tight binary HVS would be a strong indication of a central binary MBH.  Absent the discovery of a binary HVS, distinguishing among the ejection mechanisms requires predictions of the resulting spatial and velocity distribution of the HVSs along with computation of expected ejection rates (e.g. Sesana et al 2006, 2007a). For example, the \\citet{hil88} mechanism may produce the largest velocities \\citep{ses07a} and a binary black hole may produce an anisotropic distribution of ejecta \\citep{ses06}.  Dynamical and evolutionary considerations for stellar populations near the \\GC\\ also constrain HVS ejection mechanisms. For example, observations of the current population of B stars at the \\GC\\ favor the \\citet{hil88} mechanism over models involving an inspiraling IMBH or a cluster of SMBHs \\citep{per07b}.  Ejection rate estimates depend on an understanding of the way stars are scattered onto orbits intersecting the MBHs ``loss cone'' and on the specific ejection mechanism. Perets et al (2007a, 2007b) argue that scattering by giant molecular clouds gives HVS production rates from binary disruption by a single MBH that are consistent with the observations.  In addition to the HVSs, all of the ejection mechanisms produce a bound population of ejecta \\citep[e.g.,][]{bkg06}. \\citet{bro07a, bro07b} demonstrate the very probable presence of a bound population of ejected B stars. These stars have velocities between 275 \\kms\\ and 450 \\kms. In this velocity range the Brown et al surveys detect 18 outgoing stars and only 1 incoming star, consistent with the $\\sim$ 200--300~Myr lifetime of these main sequence stars and a Galactic center origin \\citep[see also][]{sve08}.  Because HVSs travel across the Galaxy, they are powerful tracers of the Galactic potential \\citep{gne05, yu07}. \\citet{yu07} emphasize that the population of bound and unbound ejecta constrain the anisotropy of the halo independent of the ejection mechanism.  The power of HVSs as probes of the Galactic center and as tracers of the Galactic potential provides strong motivation for acquisition of larger samples of these objects. Brown et al. focus on B stars because these objects are not an indigenous population in the halo and they are observable to large distances where the contrast between the density of HVSs and any indigenous population is largest. \\citet{koll07} suggest searches for possibly more numerous old-population, lower mass HVSs. Detection of lower mass HVSs could provide strong tests of models for their origin.  \\subsection {From the Galactic Center to the Halo}  In the \\citet{hil88} mechanism, HVSs naturally sample both the stellar mass function and the properties of the binary population in the Galactic Center. The ejecta then traverse the Galaxy and become observable in the halo. We show that the observable velocity and spatial distributions of HVSs are sensitive to the Galactic potential, particularly on scales of 5--200 pc from the \\GC. These distributions also reflect stellar evolutionary timescales.  To include as complete a picture of the consequences of the Hills ejection mechanism as possible, we consider a broad range of issues in this paper. In \\S2.1 we simulate the MBH-binary interaction following the procedures of  \\citet{bkg06}. Motivated by \\citet{koll07}, we extend the simulations to cover lower mass stars and unequal mass binaries. In \\S2.2 we derive forms for the Galactic potential which are a good representation of the observations from 5--10$^5$ pc. Our results show that most of the deceleration of HVSs occurs within the central 200 pc of the Galaxy; thus, a proper fit to the observations in this region is a crucial step in understanding HVSs.  In \\S2.3, \\S2.4, and \\S2.5, respectively,  we describe our integration of HVS orbits through the Galaxy, the procedure we use to construct mock catalogs of HVSs, and the impact of azimuthal deflections on HVS orbits.  We continue with more detailed discussions of the observable properties of HVSs in a simple spherically symmetric potential (\\S3) and in a three-component Galaxy (\\S4). One of the most striking general results of this investigation is that a central potential which matches observations in the central 200 pc of the Galaxy acts as a high pass filter preventing lower velocity HVSs from reaching the halo.  In \\S3 and \\S4 we predict the spatial and velocity distributions of HVSs as a function of stellar mass and we address the underlying physical explanations for the behavior of these distributions. In \\S5 we connect the predictions with the observations by calculating relative space densities as a function of stellar mass (\\S5.1). This discussion emphasizes that the relative numbers of bound stars are sensitive to the mass function in the Galactic Center and relative stellar lifetimes; relative numbers of unbound stars are nearly independent of stellar lifetime.  In \\S5.2 and \\S5.3 we analyze some sample HVSs search approaches.  We argue that searches become increasingly difficult with decreasing stellar mass (\\S5.2). We also demonstrate that short lifetimes make the discovery of post main sequence HVS unlikely. We conclude in \\S6. ",
        "conclusions": "HVSs are a fascinating newly discovered class of objects because they connect the Galactic Center with the outer halo of the Milky Way. We explore these connections by using the \\citet{hil88} model to inject stars into the Galactic potential. We track the journeys of these ejected stars across the Galaxy and derive simulated catalogs of observable HVSs.  The foundation for our model includes the construction of forms for the Galactic potential which fit observations over the range 5--10$^5$ pc. We demonstrate that potentials which match the observations within the central 200 pc of the Galaxy are crucial for understanding HVSs. Our approximations to the potential may be useful for other astrophysical problems which connect the central regions of the Galaxy to its outer reaches.  We show that important aspects of the median speeds and shapes of the observable velocity distributions of HVSs are set by the Milky Way potential at r $\\lesssim$ 200 pc. For the potentials we construct to match the observations in this central region, median speeds are larger and velocity histograms are broader. Thus, low velocity ejected stars have less penetration into the outer halo at every stellar mass. These results indicate that HVSs might be useful probes of the triaxial potential of the central Galaxy.  The models predict the spatial and velocity distributions of observable HVSs. They also provide a physical understanding of the origin of the dependences of these distributions on stellar mass for $m =$ 0.6--4 \\msun\\ and on distance in the halo for $r =$ 10--200 kpc. Here, we enumerate the main predictions of the model. We concentrate on the predictions resulting from our model potentials which match observations in the \\GC. For all of these issues, we consider the subtle effects on unequal mass binaries in the text (\\S3); here, we focus on the results for equal mass binaries.  \\begin{itemize}  \\item {\\it Stellar evolution} affects the observability of HVSs. It removes them at large radii where the travel time from the \\GC\\ exceeds the stellar lifetime.  Because the lifetimes of low mass stars are long, the observable properties of low mass HVSs depend only on the potential and the properties of binaries at the Galactic Center.  \\item For 3--4 M$_\\odot$ ejected stars, the fraction of unbound HVSs increases dramatically with $r$ in the Galactic halo (Fig. \\ref{fig:fracunbound}). Because the travel time from the \\GC\\ is a significant fraction of the main sequence lifetime for these stars, the median {\\it stellar age} of 3--4 \\msun\\ HVSs increases monotonically with $r$.  In contrast, bound stars dominate the population of long-lived, low mass HVSs. Thus, the median stellar ages are essentially independent of $r$.  \\item The {\\it velocity distribution} of 3--4 M$_\\odot$ HVSs is asymmetric with a long tail toward positive velocities. The shape of the velocity distribution again reflects coincidence of stellar lifetimes and travel times.  The velocity distributions of HVSs are increasingly symmetric with decreasing stellar mass.  \\item  As emphasized by Bromley et al. (2006), the short lifetimes of 3--4 M$_\\odot$ HVS require that they be injected with large velocities to reach the outer halo. Thus the {\\it median speed} of massive HVSs increases with $r$.  For low mass stars, the median speed depends only weakly on $r$. For all stars, the median speed increases with stellar mass at fixed $r$.  \\item The {\\it density profiles} of unbound HVSs are approximately $\\rho \\propto r^{-n}$ with $n =$ 2--2.5.  For the most massive stars, their finite lifetime removes stars at large $r$, steepening the density profile ($n \\gtrsim$ 3 at $r \\gtrsim$ 80 kpc).  The density profile for bound, mostly low mass HVS is roughly $\\rho \\propto r^{-3}$.  We compute the detailed behavior of these profiles (Fig. \\ref{fig:densallhex}).  \\item {\\it Scaled density profiles} show that the relative numbers of observable unbound HVSs as a function of stellar mass are relatively independent of the stellar lifetime. They depend mostly on the mass function at the Galactic Center at the time of ejection.  In contrast, the relative numbers of bound HVSs are a function of the stellar lifetimes and the mass function.  \\item We predict the {\\it relative observable numbers} as a function of stellar mass for the range 0.6--4 M$_\\odot$. In a magnitude limited survey, the main factors that set detectability are (1) the accessible volume, (2) the fraction of unbound HVSs, (3) the asymmetry of the velocity distribution, and (4) contamination by indigenous stellar populations. Detection of HVSs is increasingly difficult with decreasing mass because all of these issues become less and less favorable.  We also argue that post main sequence stars are poor targets because their detectability is subject to yet another limit, the very short lifetimes of these phases.  \\end{itemize}  Samples of HVSs which are sufficiently large to explore these predictions provide a strong test of the model for injection of the stars into the Galactic potential. Coupled with observations of the stellar population at the Galactic Center, observations of HVSs provide promising probes of the binary population, the stellar mass function, and the central potential of the Milky Way \\citep[see also][]{per07b}.  So far, the observations by \\citet[][and references therein]{bro07b} indicate that there are $\\sim$ 100 detectable HVS with $m \\approx$ 3--4 M$_\\odot$. Dectection of comparable lower mass populations may be feasible. Samples of hundreds of HVSs promise strong tests of models like the one we construct."
    },
    "0801/0801.1854.txt": {
        "abstract": "{Stochastic inflation describes the global structure of the inflationary universe by modeling the super-Hubble dynamics as a system of matter fields coupled to gravity where the sub-Hubble field fluctuations induce a stochastic force into the equations of motion. The super-Hubble dynamics are ultralocal, allowing us to neglect spatial derivatives and treat each Hubble patch as a separate universe. This provides a natural framework in which to discuss probabilities on the space of solutions and initial conditions. In this article we derive an evolution equation for this probability for an arbitrary class of matter systems, including DBI and k-inflationary models, and discover equilibrium solutions that satisfy detailed balance. Our results are more general than those derived assuming slow roll or a quasi-de Sitter geometry, and so are directly applicable to models that do not satisfy the usual slow roll conditions. We discuss in general terms the conditions for eternal inflation to set in, and we give explicit numerical solutions of highly stochastic, quasi-stationary trajectories in the relativistic DBI regime. Finally, we show that the probability for stochastic/thermal tunneling can be significantly enhanced relative to the Hawking-Moss instanton result due to relativistic DBI effects. }  \\preprint{PI-COSMO-58} \\begin{document} ",
        "introduction": "The great success of the inflationary paradigm has given impetus to each of the areas of physics upon which it touches. Its essential character, though, is captured by the dynamics of a scalar degree of freedom that sources gravity in such a way as to generate an accelerating expansion rate for the Universe. The basic predictions of inflation require only an understanding of the classical dynamics of a scalar field in the overdamped limit. Though simple, meeting this requirement has launched a thousand models and remains an active area of research. Theoretical discrimination among models is difficult to achieve, particularly when some models possess exotic features like non-standard kinetic terms. Even within models, it is difficult to say with certainty whether particular inflationary trajectories are natural or contrived.  It is no easy task to find an approach general enough to encompass the wide variety of inflationary models that still provides useful insight into inflationary dynamics. One way forward comes from taking a statistical approach, a technique that goes under the name of stochastic inflation. Pioneered by Starobinsky \\cite{starobinsky}, stochastic inflation treats sub-Hubble scalar field fluctuations as a stochastic source of noise for the mean value of the field coarse-grained over a Hubble volume. The noise crosses the horizon during the inflationary expansion, allowing it to influence the super-Hubble behavior of the field. Studying this interplay allows us to determine the evolution of probability density for particular solutions to the field equations, a necessary first step towards understanding which inflationary trajectories are likely and which are not.  There has been a great deal of work done applying Starobinsky's methods to a variety of inflationary scenarios and in understanding other stochastic effects for inflationary fields \\cite{Goncharov:1987ir, Rey:1987,Sasaki:1988a,Hosoya:1988b,Graziani,Linde:1993xx, Prokopec:2007ak}. Readers interested in a review of stochastic techniques and applications are referred to \\cite{Linde:2005ht}. Some recent progress includes finding an analog of Starobinsky's celebrated Fokker-Planck equation for Dirac-Born-Infeld (DBI) inflation \\cite{Silverstein:2003hf} in Ref.~\\cite{Chen:2006hs}; also, probabilities and stationary solutions in k-inflation \\cite{ArmendarizPicon:1999rj} were studied in Ref.~\\cite{Helmer:2006tz}. In this work, we describe general techniques that capture the stochastic behavior of inflating fields without relying on slow roll.  This allows us to treat either exotic DBI or k-essence fields and canonical fields within a unified framework.  To make clear what reductions of the full dynamics we make, we shall initially formulate the stochastic system on the full phase space. There are two ways in which the coarse-graining between super-Hubble and sub-Hubble modes can be achieved, which we refer to as the one- and two-noise approaches. In the one-noise approach, the stochastic noise is viewed as an additional force in the classical equations of motion, so that the momentum receives kicks, but the time derivative of the field is related to the momentum according to its classical equation. In the two-noise approach both the field and momentum receive stochastic kicks. The two approaches are qualitatively (although not identically) the same. The two-noise approach has the advantage that if the sub-Hubble perturbations are assumed to be adiabatic, then it is equivalent  to the stochastic Hamilton-Jacobi approach considered in Ref.~\\cite{Chen:2006hs} after a phase space transformation. A disadvantage to this approach is that when applied to models such as DBI inflation in which the scalar has a speed limit, stochastic jumps will typically violate this limit, even though on average it is satisfied. By contrast, the speed limit continues to be satisfied even with stochastic fluctuations included in the one-noise approach.  In the two-noise case we can find a one parameter family of equilibrium solutions which satisfy detailed balance. This additional parameter is a reflection of the fact that the solutions retain a partial memory of their initial state. This is the same parameter that arises as the integration constant in solving the Hamilton-Jacobi equation, so it is fixed by the initial conditions for the field and momentum. One of the most important reasons to understand whether equilibrium solutions exist in any model of inflation is to understand stochastic tunneling, i.e. tunneling described by the Hawking-Moss instanton \\cite{Hawking:1981fz}. If the inflationary potential has a variety of metastable minima for which the tunneling rate to other vacua is low, there is usually sufficient time for a partial equilibrium to set in before the field tunnels to the true vacuum. The equilibrium solutions we find through our stochastic approach allow us to derive the transition rates for thermally-activated/stochastic tunneling, thus generalizing the Hawking-Moss instanton result without making use of difficult to interpret Euclidean gravity techniques. Qualitatively, what we find is a transition amplitude similar to the usual one when $c_s$, the effective sound speed of the perturbations, is everywhere close to unity; however, in regions of phase space for which $c_s \\ll 1$, we find a significantly enhanced tunneling rate. This is of significant interest in understanding tunneling in models of the so-called landscape of string theory (cf. \\cite{HenryTye:2006tg,Podolsky:2007vg,Sarangi:2007jb,Copeland:2007qf}). Furthermore, the tunneling rate, like the equilibrium solution, retains a partial memory of the initial probability state.  One practical way to make use of our framework is to use it to give explicit solutions for stochastic inflationary trajectories. We did this for DBI inflation. In agreement with earlier results in \\cite{Chen:2006hs}, we find that in a phenomenologically viable DBI model with only an $m^2 \\phi^2$ potential, eternal inflation is not possible, at least in the relativistic DBI regime. However, by adding higher order polynomial corrections to the potential, which are expected to kick in at large field values, we can find solutions that possess  both a phenomenologically viable regime at small values of $\\phi$ and a stochastic DBI regime at large values of $\\phi$. Intriguingly, for certain initial conditions the evolution of $\\phi$ in this stochastic regime deviates strongly from its classical solution. The field undergoes an apparently metastable random walk, yet with $c_s$ on average well below unity. This situation is quite different from the stochastic regimes encountered in minimally coupled models in that the relativistic ``speed limit\", rather than Hubble damping, is setting the field's dynamics. Since brane inflation and other uses of the DBI action are still under active development, it will be important to determine whether regimes that exhibit these dynamics exist in fully worked out constructions. If they do, it will be necessary for stochastic dynamics to be taken into account when determining what inflationary trajectories are available and likely.  Although our generalized stochastic approach does not provide any unambiguous measure for inflation, it does allow us to state clearly the assumptions behind recent proposals for a measure. For instance, the phase space measure proposed by Gibbons and Turok \\cite{GT} assumes that stochastic diffusion is absent from the scalar field evolution, and so including the stochastic noise tells us precisely how this measure will be adjusted when fluctuations are included. We find that applying their classical arguments to DBI inflation gives precisely the same $e^{-3N}$ suppression. Alternatively, one can incorporate volume weighting into the Fokker-Planck equation, although this procedure is not manifestly gauge invariant.  Recent proposals have attempted to improve on this (see e.g. \\cite{Hartle:2007gi,Linde:2007nm}).  We shall give a somewhat pedagogical discussion, emphasizing a number of standard results not often found together in the literature. For instance, we demonstrate that the stochastic approach requires neither a nearly de Sitter geometry nor a slowly rolling field, but only that the modes are exiting the horizon and the growing mode is dominant, and can be formulated in way that consistently incorporates gravitational backreaction. We begin in Sec.~\\ref{PSFP} with derivation methods for both the Langevin and Fokker-Planck equations. We emphasize the crucial role of the consistent ultralocal truncation of GR coupled to matter fields with its resulting scaling symmetry -- which provides a natural definition of time through e-folds of volume increase. We apply  these methods to DBI inflation in Sec.~\\ref{DBIinflation}. We then discuss several of the properties and consequences of the stochastic dynamics in Sec.~\\ref{properties}, including a general prescription for when quantum corrections come to dominate over classical ones -- a necessary prerequisite for eternal inflation. In Sec.~\\ref{numerics} we numerically calculate some explicit stochastic trajectories for DBI inflation. We focus on trajectories exhibiting strongly stochastic behaviour in a highly relativistic regime. Most intriguing is our discovery that apparently metastable ultra-relativistic random walks exist for super-Planckian field values. For these fields, the value of $\\phi$ does not systematically roll in any particular direction, yet the field's mean sound speed, $c_s$, is well below unity. Finally, in Sec. \\ref{tunneling}, we calculate thermally-activated tunneling rates using our stochastic approach and demonstrate the above-mentioned enhancement in tunneling rate due to $c_s < 1$ effects. In Appendix A we demonstrate that the Hubble parameter is the generator of e-folding time-translations for ultralocal gravity and in Appendix B we give a formal path integral translation between the Langevin (``equation of motion\") and Fokker-Planck (``conservation of probability density\") approaches to stochastic dynamics. ",
        "conclusions": "Recent progress toward building phenomenologically viable inflationary models in string theory has many consequences, chief among them the need to grapple with the landscape of vacua. Together with the rise in popularity of other non-standard inflationary models, the time is ripe to reconsider some of the standard techniques used in exploring the full implications of inflation. In this article we review some the salient features of the stochastic inflationary framework, taking a very general, model independent approach. Using our formalism, we specialize to non-minimal kinetic models such as DBI inflation that appear naturally in a stringy framework. We use this as an example of the concrete uses of our formalism, and were able to discover novel stochastic behavior for DBI inflation.  By solving the stochastic dynamics numerically, we have found that metastable equilibrium states can exist in which the field is trapped by quantum corrections at some fixed value. We also find that the rate for Hawking-Moss type thermal tunneling can be significantly enhanced in regions of phase space where relativistic DBI effects are important."
    },
    "0801/0801.0293_arXiv.txt": {
        "abstract": "We summarize the main highlights from \\textit{ISO} observations towards Sgr~B2 and Orion~KL  in the far--IR domain ($\\sim$43 to 197\\,$\\mu$m). Both Star--Forming Regions are among the best sources to construct a  \\textit{template} for more distant and unresolved regions (e.g., extragalactic). We  stress some peculiarities in the interpretation (excitation and radiative transfer) of far--IR spectral lines and dust continuum emission. ",
        "introduction": "Sgr~B2 ($d$\\,$\\simeq$8.5\\,kpc) and Orion~KL ($d$\\,$\\simeq$450\\,pc) can be considered to be the two most remarkable giant molecular clouds for Astrochemistry and Star-Forming Region (SFR) studies. They are also very appropriate sources to construct a  \\textit{template} for more distant (e.g., fainter) and unresolved regions (e.g., extragalactic).  Sgr~B2, represents the most active burst of high--mass star formation in the Galactic Center region (Lis \\& Goldsmith 1989). Its geometrical properties  (clumped extended envelope, centrally condensed hot cores and H\\,{\\sc ii} regions), its physical conditions (widespread warm gas, enhanced turbulence and  UV-- and X--ray radiation fields) as well as its  chemical complexity (extended emission of refractory and organic species) mimic a miniature ``Galactic Center\". Therefore, it provides the closest \\textit{guide} for studies of  starbursts vs. active galactic nuclei.  Orion is the nearest and best studied high--mass SFR (Genzel \\& Stutzki~1989). Due to its proximity, high spectral and angular resolution observations allow us to separate the  very different physical components associated with high--mass SFRs (hot cores, outflows, shocks, PDR-like interfaces and ambient gas). In particular, Orion has been traditionally used to \\textit{test} our understanding of the physics and chemistry in hot cores (e.g.,~IRc2) and extended shocked regions (e.g.,~Peaks 1/2).   The cores of both SFRs are the most  prolific sites of molecular line emission/absorption in the Galaxy (in terms of density and intensity/depth of lines). Due to the large column density of warm material  (Sgr~B2) or due to its proximity (Orion), both complexes are among the brightest far--IR sources in the entire sky. For these reasons, Sgr~B2 and Orion~KL were fully surveyed between $\\sim$43 and 197\\,$\\mu$m at the maximum spectral resolution provided by the ISO/LWS Fabry--P\\'erot ($\\lambda/\\Delta \\lambda$\\,$\\sim$10,000 or $\\sim$30\\,km\\,s$^{-1}$). These data cover an extremely important region of the electromagnetic spectrum that can not be accessed from the ground. The resulting spectra remain unique examples of the diagnostic power of this domain, and archetypes of what future far--IR space missions will routinely observe. In this constribution  we review the main highlights of ISO data and we stress some peculiarities in the interpretation of far--IR line and continuum emission. ",
        "conclusions": ""
    },
    "0801/0801.1406_arXiv.txt": {
        "abstract": "We calculate the wavelength dependence of the ratio of the linear polarization degree to extinction (polarizing efficiency) $P(\\lambda)/A(\\lambda)$ from the ultraviolet to near-infrared. The prolate and oblate particles with aspect ratios from $a/b=1.1$ up to 10 are assumed to be rotating and partially aligned with the  mechanism of paramagnetic relaxation (Davis--Greenstein). Size/shape/orientation effects are analyzed. It is found that the wavelength dependence of $P(\\lambda)/A(\\lambda)$ is mainly determined by the particle composition and size whereas the values of $P(\\lambda)/A(\\lambda)$ depend on the particle shape, degree and direction of alignment. ",
        "introduction": "Modelling of the wavelength dependencies of interstellar extinction $A(\\l)$ and  linear polarization $P(\\l)$ allows one to obtain  information about such properties of interstellar grains as size,  composition, shape, etc. Interstellar polarization also tells us about the structure of magnetic fields because it arises due to dichroic extinction of non-spherical grains aligned in large-scale Galactic magnetic fields \\cite{gre68}, \\cite{ha04}. A correlation between observed interstellar extinction and polarization shows that the same particles are responsible for both phenomena.  Very often interstellar  extinction and  polarization are modelled separately (e.g., \\cite{weid01},  \\cite{km95}). The modelling of these phenomena usually includes consideration of normalized extinction $E(\\lambda-{\\rm V})/E\\rm (B-V)$ and normalized polarization given by Serkowski's curve ($P_{\\rm max}$ is the maximum degree of polarization  and $\\lambda_{\\rm max}$ the wavelength corresponding to it, $K$ is the coefficient) $P(\\lambda)/P_{\\rm max} = \\exp [-K \\ln^2 (\\lambda_{\\rm max}/\\lambda)]$ (see, e.g., \\cite{wcm93} and discussion in \\cite{w03}, \\cite{nvv2}). In these cases, the important observational data (absolute values of extinction and polarization) are ignored and normalized curves can be fitted using particles of different composition and slightly different sizes (see, e.g., discussion in \\cite{gre68}, \\cite{nvv2}). For example, the consideration of absolute extinction gives the possibility to use the cosmic abundances to constrain grain models (see \\cite{vihd06}). In the case of interstellar polarization it is better to interpret the wavelength dependence of polarizing efficiency per unit extinction $P(\\lambda)/A(\\lambda)$ (ratio of the linear polarization degree to extinction). This allows one to estimate the degree and direction of grain alignment and particle shape. Note that  previous consideration of polarizing efficiency was restricted by one wavelength band only (like $P({\\rm V})/A({\\rm V})$ or $P({\\rm K})/A({\\rm K})$).  Here, we analyze the extinction and polarization produced by homogeneous spheroidal particles with different aspect ratios $a/b$ ($a$ and $b$ are major and minor axes, respectively). Spheroidal particles have the important applications in various fields of science (spheroidal antennas, bacteria and microweeds, raining drops, etc.). By changing $a/b$ the particles with a shape varying from close to as sphere ($a/b \\simeq 1$) up to needles (prolate spheroids) or disks (oblate spheroids) can be studied. The particles are assumed to be rotating and partially aligned with the Davis--Greenstein mechanism (mechanism of paramagnetic relaxation) \\cite{dg51}. We focus on the behaviour of polarizing efficiencies and consider the normalized extinction $A(\\lambda)/A_{\\rm V}$ (i.e., ignore the cosmic abundances which will be taken into account in next paper). The theory is compared with observations of stars seen apparently through one interstellar cloud. ",
        "conclusions": "In the frame of the model of partially aligned spheroidal grains we studied the wavelength dependence of the ratio of the linear polarization degree to extinction (polarizing efficiency) $P(\\lambda)/A(\\lambda)$.  The main results  can be summarized as follows.  1. It is found that the wavelength dependence of $P(\\lambda)/A(\\lambda)$ is mainly determined by the particle composition and size.  2. The values of $P(\\lambda)/A(\\lambda)$ depend on the particle shape, degree and direction of alignment.  3. The  modelling of the only wavelength dependence of polarizing efficiency does not allow one to determine all parameters dust ensemble. Therefore, interpretation of the observations must include consideration of extinction.   \\noindent{\\bf Acknowledgements}  We are thankful to Vladimir Il'in for careful reading of manuscript. The work was partly supported by grants NSh 8542.2006.2, RNP 2.1.1.2152 and RFBR 07-02-00831 of the Russian Federation."
    },
    "0801/0801.4015_arXiv.txt": {
        "abstract": " ",
        "introduction": "Dark matter is frequently assumed to consist of massive weakly interacting particles which are stable (or have a very long lifetime) because their decay is forbidden by some (approximate) symmetry. However, it is also well-known that dark matter may originate in a hidden sector which is coupled to the standard model only via higher-dimension operators, ensuring that dark matter does not decay (see e.g.~\\cite{Kolb:1985bf,smallnumbers}).  We demonstrate that the latter scenario is realized under fairly general assumptions in the type IIB string theory landscape. The main starting point is the well-known fact that type IIB flux compactifications generically contain many strongly warped regions or throats~\\cite{Denef:2004ze}. This statement can be made quantitative~\\cite{Hebecker:2006bn} under the assumption that the fine-tuning of the cosmological constant requires manifolds with many 3-cycles and that, in many cases, such 3-cycles produce Klebanov-Strassler throats~\\cite{KS} if stabilized at small volume. We base our analysis on the fact that, after inflation ends, the standard model sector is heated to a certain temperature. Moreover, we assume that the throats have received no energy from the reheating process.\\footnote{ Alternatively, if this assumption is not fulfilled, i.e.~if the throats are heated directly by the reheating mechanism, a lower reheating temperature would be sufficient for our scenario. } Even under such minimal assumptions, a certain amount of energy is deposited in the various throats by energy transfer from the heated standard model sector. There exist certain optimal throat lengths (i.e. optimal warp factors) for which a given throat contains a maximal amount of cold dark matter in the late universe. Requiring furthermore that such an `optimal' throat is responsible for the dark matter observed today, we determine the reheating temperature to be between $ 10^{10}$~GeV and $10^{11}$~GeV. The IR scale of the relevant throat (i.e. the mass of dark matter particles) is either around $10^5$~GeV or around $10^{10}$~GeV. Moreover, we observe that the decay rate of throat dark matter to the standard model sector is not negligible. It may, for a certain range of parameters, lead to the discovery of this variant of dark matter via the observation of the diffuse $\\gamma$-ray background or the spectrum of very energetic $\\gamma$-rays.  The fact that dark matter can come from a hidden (or more precisely conformally sequestered) sector realized by a Klebanov-Strassler throat has already been emphasized in~\\cite{CT}. This paper focuses on scenarios where a Klebanov-Strassler throat is heated by the annihilation of a brane with an antibrane at the end of inflation. Subsequently, energy is transferred from this throat to other throats which may be present in the compact manifold. Throat-localized Kaluza-Klein (KK) modes, which are produced in this way, can be the observed dark matter if they are sufficiently long-lived and if certain parameters which determine their relic density are tuned. More precisely, three types of dark matter candidates are discussed in \\cite{CT}: Kaluza-Klein modes which are localized in the throat where inflation took place and those localized in the throat where the standard model lives have to carry an (approximately) conserved angular momentum in the throat in order to be sufficiently stable. Kaluza-Klein modes which are localized in other throats may, by contrast, be long-lived also without such an angular momentum.\\footnote{ In addition, \\cite{CT} also discusses particles on D-branes in these throats as a dark matter candidate.}  We approach the possibility of throat dark matter from a different and, to a certain extent, more general perspective: We do not assume reheating to be due to brane-antibrane annihilation in a warped region. Instead, we only rely on the fact that the standard model (which is realized in the unwarped part of the manifold) has a certain reheating temperature after inflation ends. In our approach, throats are not a `model building feature' introduced to realize inflation, uplifting, etc. Instead, we view the presence of (a potentially large number of) throats as a prediction of the type IIB string theory landscape. Accordingly, we consider scenarios in which various throats are present and perform a quantitative analysis of the energy density in these throats in the late universe. We attempt to avoid any other specific assumptions, relying only on the known fact that the standard model has a certain initial temperature in early cosmology. The resulting dark matter can then be viewed as a generic prediction of the type IIB landscape. It turns out that, given a reasonably large number of throats, the reheating temperature is the only parameter which has to be tuned to account for the observed dark matter density. We show quantitatively, using results from~\\cite{Hebecker:2006bn}, that sufficiently many throats are available in an important fraction of vacua of the type IIB landscape.  We focus in particular on the phenomenological importance of fermionic KK modes in the throat. To the best of our knowledge, this point has so far not received sufficient attention in the literature. The Klebanov-Strassler throat has $\\mathcal{N}=1$ supersymmetry. Generically, this supersymmetry is weakly broken by the compact space to which the throat is attached and in which supersymmetry has to be broken for phenomenological reasons. Therefore, light fermionic KK modes are guaranteed to be present in the spectrum. They turn out to play a central role in the phenomenology of sequestered dark matter. Furthermore, we use our new values for the energy transfer rates and the decay rates between throats~\\cite{hotthroats}. For temperatures and Kaluza-Klein masses smaller than the compactification scale, our rates differ from the values used in~\\cite{CT}. We restrict our ana\\-lysis to this case, which can also be characterized as the case of `overlapping potentials' in the quantum mechanical language of tunneling~\\cite{hotthroats}. Our reasons for focusing on this parameter range are as follows: On the one hand, a large compactification scale is required to make the KK modes sufficiently stable against decay to the standard model or other throats.\\footnote{ An exception is the decay rate of fermionic KK modes to the standard model sector which can be made small even for small compactification scales (cf.~Sect.~\\ref{DecaytoSM}).} On the other hand, temperatures above the compactification scale would imply that also the unwarped part of the manifold is heated up. The resulting gas of KK modes in the compact space may then destabilize the volume modulus, making the analysis of the early cosmological evolution much more difficult.  Our paper is organized as follows: Sect.~\\ref{thermalproduction} describes the thermal production of sequestered dark matter. Using the energy transfer rates from~\\cite{hotthroats}, we determine the energy density deposited in a throat by the heated standard model sector. We show that the KK modes which are produced in that way thermalize for a certain range of parameters. Whether this happens or not influences the further time evolution of the energy density significantly. Taking this into account, we calculate the late-time abundance of KK modes as a function of the reheating temperature and the IR scale of the throat. In Sect.~\\ref{candidates}, we discuss various decay channels of the KK modes. We show that they decay very quickly to a scalar state and its fermionic superpartner. These lightest KK modes can then decay to the standard model sector. We determine the corresponding decay rates which turn out to be different for the scalar and the fermion. Cosmological scenarios are discussed in Sect.~\\ref{scenarios}. First, we analyse setups with a single throat. We find that a moderately long throat gives a promising dark matter candidate which may allow for a discovery by upcoming $\\gamma$-ray experiments. Then, we consider scenarios with a large number of throats, using results on the distribution of multi-throat configurations from~\\cite{Hebecker:2006bn}. We find that a throat of the required length is in many cases present. Some issues concerning supersymmetry breaking in the throat sector are discussed in Sect.~\\ref{extra}. Finally, in Sect.~\\ref{summary}, we give a summary of our results. ",
        "conclusions": "\\label{summary} Strongly warped regions or throats are a common feature of the type IIB string theory landscape. KK modes whose wavefunctions are localized in a throat have redshifted masses, allowing for their production after reheating in the standard model sector, even if the reheating temperature is not very high. In addition, these KK modes are only very weakly coupled to the rest of the compact manifold, potentially resulting in a very long lifetime. These properties make the KK modes an interesting dark matter candidate.  We have considered a setup in which the standard model lives in the unwarped part of a compact manifold, which in addition has a certain number of throats. To be conservative, we have assumed that reheating only takes place in the standard model sector. Even under this minimal assumption, the throats are heated up by energy transfer from the hot standard model plasma. In Sect.~\\ref{et}, we have determined the energy density deposited in a throat, using our result for the corresponding energy transfer rate from a previous paper~\\cite{hotthroats}.  Throats have a dual description as a strongly coupled gauge theory with a large number of colours. From this gauge theory point of view, the energy density in a throat is initially in the form of gauge theory states with energy of the order of the reheating temperature. These states subsequently settle into a certain number of lighter glueballs with some distribution of kinetic energies. The knowledge of this distribution is important since it determines whether the glueballs thermalize and whether (and for how long) the energy density scales like radiation or like matter. Since we are at present unable to determine this distribution of kinetic energies, we have considered two extremal cases in Sect.~\\ref{te}. In the first case, the initial gauge theory state settles into a large number of glueballs with kinetic energies of the order of their mass. In the second case, the decay products of the initial state are an $\\mathcal{O}(1)$ number of glueballs which accordingly have kinetic energies of the order of the reheating temperature.  It turns out that the gauge theory thermalizes in both extremal cases if the energy density in this sector is above the critical energy density for deconfinement. The energy density thus scales like radiation with the expansion of the universe until the confinement phase transition takes place. Afterwards, it scales like matter. Taking this scaling into account, we have determined the contribution of the throat sector to the total energy density of the universe at our epoch. We have also determined this contribution for the case that the initial energy density in the throat sector is below the critical energy density and found that it differs considerably between the two extremal cases. We expect the true behaviour in this region of parameter space to be in between the two extremal cases.  After the confinement phase transition, different types of glueballs with mass of the order of the confinement scale $m_\\IR$ are formed. Similarly, if the gauge theory does not thermalize, the initial gauge theory states created at reheating settle into a certain number of light glueballs. As we have shown in Sect.~\\ref{processes}, these glueballs quickly decay to a lightest state and its superpartner. On the basis of a number of papers devoted to the spectrum of the KS gauge theory~\\cite{Krasnitz:2000ir,Caceres:2000qe, Amador:2004pz,Firouzjahi:2005qs,Noguchi:2005ws,Dymarsky:2006hn,Berg:2006xy, Dymarsky:2007zs,Benna:2007mb}, we expect these lightest states to be a scalar glueball and its spin-$\\smash{\\frac{1}{2}}$ superpartner. The glueballs couple (very weakly) to the standard model and other throats and can thus decay to these sectors. In Sect.~\\ref{DecaytoSM}, we have applied our results for this decay~\\cite{hotthroats} to scalar glueballs. A crucial difference with respect to our previous analysis is the fact that the bulk fields which mediate the decay are massive in a flux background. Depending on the mass of the decaying glueballs, this can lead to a significant suppression of the corresponding decay rate. In addition, scalar glueballs can decay to two gravitons. We have determined the corresponding rate in Sect.~\\ref{processes}.  Similarly, spin-$\\smash{\\frac{1}{2}}$ glueballs can in principle decay to a graviton and a gravitino. To get a stable dark matter candidate, we are mainly interested in setups where such decays are kinematically forbidden due to a heavy gravitino. This requires that the supersymmetry breaking scale is larger than the mass of the glueball. It may also mean that the superpartners of standard model particles are heavier than the glueballs. If R-parity is conserved, most decay channels of spin-$\\smash{\\frac{1}{2}}$ glueballs to the standard model sector involve such a superpartner and are therefore kinematically forbidden. In Sect.~\\ref{DecaytoSM}, we have identified an operator which does not involve a superpartner and which would allow the decay of spin-$\\smash{\\frac{1}{2}}$ glueballs to a Higgs and a lepton. If present, this operator would give the dominant decay channel even for maximally broken R-parity. Using the corresponding vertex, we have determined the decay rate of spin-$\\smash{\\frac{1}{2}}$ glueballs to the standard model.   In Fig.~\\ref{fig:plot}, we have plotted the contribution of a throat sector to the total energy density of the universe for fixed reheating temperature $T_\\RH$ and as a function of the IR scale $m_\\IR$. We expect this function to have two maxima at IR scales $ m_\\IRc$ and $T_\\RH$. For a throat with IR scale $m_\\IRc$, the dual gauge theory thermalizes precisely to the phase transition temperature and therefore the energy density scales like matter immediately after reheating. KK modes in throats with IR scale $T_\\RH$, on the other hand, are so massive that they become nonrelativistic immediately after reheating and the energy density again scales like matter afterwards.  One of the main underlying ideas of our analysis is the generic presence of throats in the type IIB landscape. It is then probable to have throats with these optimal IR scales in a given compact manifold. For simplicity, we have first discussed a scenario with a single throat in Sect.~\\ref{SingleThroat}. We have found that, in many cases, KK modes in throats with IR scale $T_\\RH$ have a decay rate which is too high for them to be a good dark matter candidate. However, if the gravitino is very heavy (high-scale SUSY breaking) and certain operators connecting the standard model and the moduli sector are suppressed, the fermionic glueballs may nevertheless survive and play the role of dark matter.  The more promising case is that of throats with IR scale $m_\\IRc$, to which we now turn. In this case, it follows immediately from our results of Sect.~\\ref{te} that a reheating temperature of $10^{10}$~GeV to $10^{11}$~GeV leads to the right amount of glueballs to account for the observed dark matter. The critical IR scale $m_\\IRc$ is a function of $T_\\RH$. After having fixed $T_\\RH$, we find a mass for our dark matter candidate which is between $\\smash{10^5}$~GeV and $10^6$~GeV. In Sect.~\\ref{ManyThroats}, we have used results from Ref.~\\cite{Hebecker:2006bn} on the distribution of throats in the landscape to consider a scenario with a large number of throats. We have found that there are setups in which the probability of having a throat with the required IR scale is large. The only free parameter which has then to be fixed is the reheating temperature.  Our dark matter scenario may lead to some interesting observable signatures. The dark matter glueballs decay to the standard model with a very low, but non-negligible rate. The decays produce photons to which experiments like GLAST or HESS may be sensitive. It turns out that the decay rates depend on two parameters: The 10d Planck mass $M_{10}$ enters via the flux stabilized mass of fields which mediate the decay. Moreover, the decay rate of spin-$\\smash{\\frac{1}{2}}$ glueballs depends on the coupling strength $\\lambda$ of the aforementioned operator which allows their decay to a lepton and a Higgs. In Sect.~\\ref{SingleThroat}, we have identified two interesting scenarios:  If $\\lambda$ is of the order 1, $M_{10}$ has to be very large in order to get a sufficiently stable dark matter candidate. Namely, for a lower 10d Planck scale, the spin-$\\smash{\\frac{1}{2}}$ glueballs decay to the standard model with a rate which is in conflict with measurements of the diffuse $\\gamma$-radiation. On the other hand, the 10d Planck scale cannot be larger than the 4d Planck scale. This makes it more probable that the lifetime of spin-$\\smash{\\frac{1}{2}}$ glueballs is in a range that can be probed with new, more sensitive experiments like GLAST. If the scenario with $\\lambda =\\mathcal{O}(1)$ is realized in nature, one can hope to detect a signal from the decaying glueballs in the near future.  If $\\lambda$ is much smaller than 1, a lower $M_{10}$ still leads to sufficiently stable spin-$\\smash{\\frac{1}{2}}$ glueballs. For a low 10d Planck scale, also the decay of scalar glueballs can become relevant for detection. In contrast to fermionic glueballs, scalar glueballs can decay directly to two photons. This decay channel leads to a sharp $\\gamma$-line at an energy of $10^5$~GeV to $10^6$~GeV which could be detected with experiments like HESS. In addition, the scalar glueball decays also produce a continuous $\\gamma$-ray spectrum to which e.g.~GLAST may again be sensitive. If the scalar glueballs make up an $\\mathcal{O}(1)$ fraction of the dark matter at our epoch, a 10d Planck scale of the order of $10^{13}$~GeV would allow for a detection of both signals in the near future. Such a 10d Planck scale corresponds to a compactification radius of the order of just $ 50\\,l_{\\bf string}$, which is not extremely large.  Another interesting effect is the decay of the scalar glueballs into two gravitons which may happen before our epoch. Since we consider a setup in which the corresponding decay of the spin-$\\smash{\\frac{1}{2}}$ glueballs into a graviton and a gravitino is kinematically forbidden, we still have a sufficiently stable dark matter candidate. For reasons that we have explained in Sect.~\\ref{processes}, we expect the scalar glueballs to have a lower abundance than the spin-$\\smash{\\frac{1}{2}}$ glueballs. The total dark matter abundance would then only change by a small factor when the scalar glueballs decay. It is possible, however, that the scalar glueballs have a comparable or even higher abundance than the spin-$\\smash{\\frac{1}{2}}$ glueballs. The change in the dark matter abundance could then be enormous. It would be interesting to consider the implications of this scenario for cosmology.  In a setup with a large number of throats one also expects throats with IR scales larger than $10^5$~GeV to $10^6$~GeV (the IR scale of the throat which provides the dark matter). The glueballs from these throats have shorter lifetimes and may decay already during nucleosynthesis. Therefore, successful nucleosynthesis may impose further constraints on our scenario. For a particular example, we have checked in Sect.~\\ref{ManyThroats} that the glueball decays do not affect nucleosynthesis due to a low branching ratio to the standard model. Examples are conceivable, however, in which bounds from nucleosynthesis are only marginally fulfilled. It would then be interesting to look for traces of glueball decays in the abundances of light elements. Unfortunately, we are at present unable to determine the contribution of throats with the aforementioned range of IR scales to the total energy density of the universe with sufficient precision. Progress in this direction requires a detailed understanding of hadronization in a strongly coupled gauge theory, which we therefore view as a further interesting topic for future research.  \\vspace*{.4cm} \\noindent {\\bf Acknowledgements}:\\hspace*{.5cm}We would like to thank F.~Br\\\"ummer, W.~Buchm\\\"uller, T.~Dent, J.~Ellis, W.~Hofmann, K.~Sigurdson and M.~Trapletti for helpful discussions."
    },
    "0801/0801.3897_arXiv.txt": {
        "abstract": "The Be/X-ray binary 1A 0535+262 was discovered in 1975 during a giant outburst. Afterwards it has shown periods of quiescence (flux below 10 mCrab), normal outbursts (10 mCrab-1Crab) and occasionally giant outbursts (several Crab). Ending 11 years of quiescence, the last giant outburst took place in May/June 2005, but the source was too close to the Sun to be observed by most satellites. A subsequent normal outburst took place in August 2005, which was observed by \\inte\\ and \\xte\\ TOO observations. Based on \\inte\\ data, we present results on the long term pulse period history of the source, on their energy dependent pulse profiles and on phase resolved spectra. ",
        "introduction": "Since its discovery in 1975 \\cite{rosenberg75}, the Be/X-ray binary 1A 0535+262 has been intensely studied. Further details may be found in a review by \\cite{giovannelli92}. The source has shown giant outbursts in April/May 1975 \\cite{rosenberg75}, October 1980 \\cite{nagase82}, March/April 1989 \\cite{sunyaev82}, February 1994 \\cite{finger94}, and in May/June 2005 \\cite{tueller05}. Then the source showed a normal outburst in August/September 2005, observed by \\inte and \\xte. During this outbusrt, the averge flux was 300 mCrab in the energy range 5-100 keV \\cite{kretsch05}. The last normal outburst took place in December 2005 \\cite{finger05}. Fig.~\\ref{fig:asm_lc} shows the \\xte \\asm long term light curve of the source during the last three outbursts. In this paper, we present preliminary results on the timing and spectral behaviour of the source based on \\inte \\ibis data from the August/September 2005 outburst. All data were reduced and analysed using \\inte OSA v5.1. Pulse phase resolved spectroscopy was performed using additionally the software provided by IASF/INAF Palermo. \\begin{figure} \\centering \\includegraphics[width=0.9\\linewidth]{asm_lc_a0535_integral.eps} \\caption{\\xte \\asm long term light curve of the last three outbursts. The dotted lines indicate the time of our \\inte observation, the inset gives a blow-up. The vertical lines indicate the times of periastron passage.\\label{fig:asm_lc}} \\end{figure} ",
        "conclusions": "In this poster we show some first results of our analysis of the \\inte 1A 0535+262 data from the August/September 2005 outburst. The source is found to pulsate up to 120 keV and energy dependent pulse profiles are observed. We also detect the presence of two phase dependent cyclotron lines at $\\sim$~45~keV and $\\sim$~100~keV in the hard X-ray spectrum of the source. This therefore confirms previous results by Kendziorra et al. \\cite{kend94} based on \\ttm\\ data taken during the March/April 1989 giant outburst. We wish to outline however that observation from  \\osse\\  on \\cgro\\ during the February 1994 giant outburst clearly detected only the 110 keV line. The presence of the fundamental line was not clear, but it was concluded that if the line was present, its optical depth should be significantly smaller than that of the 110 keV line \\cite{grove95}. Based on our results we can estimate from the fundamental line at $\\sim45$~keV the magnetic field of the source to be $\\sim4$~x$10^{12}$~G. Further detailed analysis of this outburst is ongoing. The next step will be to try different models for fitting the broad band continuum making use of observational data from \\jemx\\ and \\spi."
    },
    "0801/0801.1540_arXiv.txt": {
        "abstract": "We present a geometrical methodology to interpret the periodical light curves of Soft Gamma Repeaters based on the magnetar model and the numerical arithmetic of the three-dimensional magnetosphere model for the young pulsars. The hot plasma released by the star quake is trapped in the magnetosphere and photons are emitted tangent to the local magnetic field lines. The variety of radiation morphologies in the burst tails and the persistent stages could be well explained by the trapped fireballs on different sites inside the closed field lines. Furthermore, our numerical results suggests that the pulse profile evolution of SGR 1806-20 during the 27 December 2004 giant flare is due to a lateral drift of the emitting region in the magnetosphere. ",
        "introduction": "Soft Gamma Repeaters (SGRs) seemed weird since the first discovery on 1979 March 5 \\citep{mazets79} for their mysterious characteristics such as the large energy release, the repetitive emission of bursts in hard X-rays or soft gamma-rays bands, and the pulsed periodical emissions after the bursts and in the quiescent stages whose morphologies are both energy-dependent and time-dependent (see \\cite{woods04} for a recent review). So far, the catalogue\\footnote{http://www.physics.mcgill.ca/$\\sim$ pulsar/magnetar/main.html} has four SGRs confirmed plus one candidate. SGRs are found to be associated with young ($\\sim 10^4$ yr) supernova remnants (SNRs), and their spin periods are $5\\sim 8\\rm{s}$ and at a large spin down rate about $10^{-11}\\rm{s~s^{-1}}$, which give an inferred ultra-strong dipolar magnetic field at the order of $10^{14} \\sim 10^{15}\\rm{G}$.  A variety of models were proposed to understand the physics of SGRs, and the magnetar model is now widely accepted to ascribe SGRs as neutron stars with magnetic field of $10^{14}\\sim 10^{15}$G \\citep{td95,td96}. Unlike most of the pulsars in the neutron star family powered by their spin-down, the high-luminosity bursts and the persistent X-ray or soft $\\gamma$-ray pulsations of magnetars come from the decay of their ultra-strong magnetic fields. The spectrum of the persistent X-ray emission could be fitted by a superposition of a blackbody component and a power law, which suggests that besides the radiation from the neutron star surface, there is another component coming from the magnetosphere.  The high-luminous burst has now been successfully interpreted by the dissipation of magnetic energy. However, the persistent long-period pulsations in the quiescent X-ray emission of SGRs are still not well understood for their complicated and astonishing morphologies. \\cite{tlk02} assumed that the angular pattern of the X-ray flux are modified by the resonant cyclotron scattering at a distance about 50-100 km above the neutron star, and multi-pulses are generated by some effects associated with the twisted magnetosphere (e.g., the optical depth is thin near two magnetic poles, and thick at the magnetic equator). However, their numerical results from the Monte-Carlo simulation has a bias in favor of the orthogonal dipole (i.e., the magnetic axis is perpendicular to the spin axis), and the viewing angel has also to be nearly $90^0$ from the spin axis, for the photons escaping from optically thin poles \\citep{thompson07}. In addition, the separation of the strong peaks in their calculations is always $\\pi$ (half phase cycle), which differs from the observation. Thompson (private communication, 2005) agreed that the change in the persistent pulse profile reflects a re-distribution of persistent currents on closed field lines. However, he ascribed such re-distribution as the gradual change in the magnetic twist at a distance of 30-100 km, while our understanding for that is due to the gradual migration of the emitting regions (crustal platelet motion or hot plasma drifting in the magnetosphere, we will discuss the possibility of each candidate in \\S\\ref{app}).  In this paper, we present some simulated pulsed profiles by assuming the radiation region located in the closed field lines, and make attempts to simulate the radiation morphology evolution in one particular event, i.e., SGR 1806-20 burst on 27 December 2004. We first have a general review on the timing properties of SGRs in \\S\\ref{obs}. In \\S\\ref{model}, we present the motivation to reproduce the light curves in the closed field lines and the resultant profiles. In \\S\\ref{app}, we applied the geometrical model to the well-known December 2004 burst of SGR 1806-20 and calculate the radiation morphology by a three-dimensional magnetosphere simulation. We also try to explain the change in the persistent pulse profile. A brief discussion is given in \\S\\ref{dis}. ",
        "conclusions": "}  We have investigated the timing properties of SGRs in the persistent state and outburst tails, and ascribed the variety of the pulsed radiation morphologies as the emissions coming from the closed field line regions inside the neutron star magnetosphere. For the weak relativistic effects in the magnetosphere closer to the stellar surface, the totally different features of the light curve of SGRs from those of Crab and Vela could be explained. Furthermore, by assuming the emitting region drift in the magnetosphere, we are able to simulate the pulse profile evolution of SGR 1806-20 during the giant flare on 27 December 2004. In addition, when we take the emissions from both magnetic poles into account, we are also able to explain the occurrence of the third peak with the phase roughly consistent with the observation, which increases in strength relative to the two major ones (e.g. Figure 2 in \\cite{palmer05}) at the end of the tails after the outburst. We believe that when the emission mechanism becomes clear, the emissivity on the field lines should not be uniform. In that case we may be able to predict a stronger third peak.  \\cite{feroci01} made simple analytic fits to the flare light curves of SGR 1900+14 on 27 August 1998, and concluded that a fireball with contracting surface (rather than a cooling surface of fixed area) could provide a reasonable explanation to the decay tail of the outburst. They also proposed that the four-peaked profile was produced by several X-ray jets tied to the neutron star surface. However, the size of the radiation region in our fitting of SGR 1806-20 is unchanging, and the smooth decay in the tail phase of the giant flare might be due to the cooling of the plasma. It was believed that the phase stability of the pulses in light curve is due to the fixed location of the emitting region on the stellar surface. But our results show that a small migration could not break the stability in phase.  We have speculated some possible mechanisms, which cause the radiation region to migrate. It seems very clear that the physical motion of magnetic field lines, where the charged particles are trapped, must be very slow due to the extremely large drag force by the electrons in the core of neutron star. On the other hand, the $\\bf{E}\\times \\bf{B}$ drift seems more possible. However, how this residual electric field survives from the screen of electron/positron pairs is not clear. We argue that one possible way is that some platelet is in the open field lines. It is unclear if this situation always occurs.  The 3-100 keV phase-averaged spectrum of the pulsed tail during the 2004 burst is fitted by a blackbody function at the temperature of 5.1 keV plus a power law \\citep{hurley05}. We didn't give the calculated spectrum in this paper, for the radiation mechanism in the closed field lines needs further work. However, we argue that the power law component results from the synchrotron radiation by the charged particles gyrating along the magnetic field lines. We need more information in the higher energy band and the phase-resolved spectra to provide more constrains and modification on our geometrical model."
    },
    "0801/0801.1851_arXiv.txt": {
        "abstract": "We explore the dependence of the central logarithmic slope of dark matter halo density profiles $\\alpha$ on the spectral index $n$ of the linear matter power spectrum $P(k)$ using cosmological $N$-body simulations of scale-free models (i.e. $P(k) \\propto k^n$). These simulations are based on a set of clear, reproducible and physically motivated criteria that fix the appropriate starting and stopping times for runs, and allow one to compare haloes across models with different spectral indices and mass resolutions. For each of our simulations we identify samples of well resolved haloes in dynamical equilibrium and we analyse their mass profiles. By parameterising the mass profile using a ``generalised'' Navarro, Frenk \\& White profile in which the central logarithmic slope $\\alpha$ is allowed to vary while preserving the $r^{-3}$ asymptotic form at large radii, we obtain preferred central slopes for haloes in each of our models. There is a strong correlation between $\\alpha$ and $n$, such that $\\alpha$ becomes shallower as $n$ becomes steeper. However, if we normalise our mass profiles by $r_{-2}$, the radius at which the logarithmic slope of the density profile is $-2$, we find that these differences are no longer present. This is apparent if we plot the maximum slope $\\gamma=3(1-\\rho/\\bar{\\rho})$ as a function of $r/r_{-2}$ -- we find that the profiles are similar for haloes forming in different $n$ models. This reflects the importance of concentration, and reveals that the concentrations of haloes forming in steep-$n$ cosmologies tend to be smaller than those of haloes forming in shallow-$n$ cosmologies. We conclude that there is no evidence for convergence to a unique central asymptotic slope, at least on the scales that we can resolve. ",
        "introduction": "The currently favoured model of cosmological structure formation asserts that the Universe is dominated by some form of non-baryonic Cold Dark Matter (hereafter CDM), and it makes a number of generic predictions about how the dark matter clusters on small scales. Arguably the defining prediction of the CDM model is that the radial mass density profile of dark matter haloes is divergent at small radii, which leads to a central density ``cusp''. Characterising the form of this mass density profile has been one of the most active research problems in computational cosmology over the last decade.  The seminal study of Navarro, Frenk \\& White (1996,1997; hereafter NFW) introduced the concept of a ``universal'' mass density profile for dark matter haloes that form in hierarchical clustering cosmologies. This NFW profile is written as  \\begin{equation} \\label{eq:NFWspherical} \\rho(r) = \\frac{\\rho_c}{\\left( r/r_s \\right)(1+r/r_s)^{2}} \\ , \\end{equation}  \\noindent where $r_s$ is a scale radius and $\\rho_c$ is a characteristic density, which can be related to $r_s$ once the virial mass of the halo is fixed; equation~\\ref{eq:NFWspherical} describes a one-parameter family of curves. It is convenient to rewrite equation~\\ref{eq:NFWspherical} in the more general form  \\begin{equation} \\label{eq:NFWgeneral} \\rho(r) = \\frac{\\rho_c}{\\left( r/r_s \\right)^{\\alpha}(1+r/r_s)^{3-\\alpha}} \\ , \\end{equation}  \\noindent where $\\alpha$ is the central asymptotic logarithmic slope. The NFW profile is ``universal'' in the sense that it describes the ensemble averaged mass profile of dark matter haloes in dynamical equilibrium, independent of virial mass, cosmological parameters and initial power spectrum.  During the last decade numerous studies have investigated whether or not the NFW profile does indeed provide an adequate description of the mass profile of dark matter haloes, and to understand the physical mechanisms that shape the functional form of the profile. NFW recognised that $r_s$ correlates with the virial mass of the halo, increasing with decreasing virial mass. They argued that this virial mass--scale radius relation is an imprint of the hierarchical assembly of haloes; low-mass haloes tend to collapse before high-mass haloes, when the mean density of the Universe is higher, and $r_s$ reflects the mean density of the Universe at the time of collapse.  The halo sample that formed the basis of \\citet{1997ApJ...490..493N} contained of order $10^4$ particles (and hence resolving the profile down to about 10\\% of the virial radius, according to the convergence criteria of \\citet{2003MNRAS.338...14P}). Subsequent studies drew upon haloes containing about two orders of magnitude more particles within the virial radius and confirmed the basic finding of NFW that CDM haloes are cuspy (e.g. \\citealt{1997ApJ...477L...9F}; \\citealt{1998ApJ...499L...5M}; \\citealt{2000ApJ...529L..69J}; \\citealt{2000ApJ...535...30J}; \\citealt{2001ApJ...557..533F}; \\citealt{2002ApJ...574..538J}). These studies led to debate about the exact value of the logarithmic inner slope, which ranged from $\\alpha \\sim 1$~to~$1.5$ (e.g., \\citealt{1997ApJ...490..493N}, \\citealt{1999MNRAS.310.1147M}). The study of \\citet{2003MNRAS.338...14P} provided convergence criteria that allowed simulators to understand the impact of numerical artifacts on the central structure of haloes, and to identify the innermost radius at which the mass profile could be considered reliably resolved. More recently, a great deal of effort has gone into even higher resolution simulations that have revealed density profiles of CDM haloes to well within $1\\%$ of the virial radius (\\citealt{2004ApJ...606..625F}; \\citealt{2004ApJ...607..125T}; \\citealt{2004MNRAS.349.1039N}; \\citealt{2005MNRAS.357...82R}; \\citealt{2005MNRAS.364..665D}). The highest resolved simulation to-date reached an effective mass resolution of about 130 million particles in a cluster sized dark matter halo still supporting the evidence for a central cusp with a logarithmic inner slope of about $\\alpha \\sim 1.2$ \\citep{2005MNRAS.364..665D}.\\\\  Cosmological simulations allow us to characterise the functional form of the density profile and to explore the important physical processes (e.g.  merging, smooth accretion) that drive this form \\citep[cf., ][]{2006MNRAS.368.1931L, 2007ApJ...666..181S}, but a theoretical understanding of the origin of the mass profile is essential. Is the form of the profile set by non-linear processes during the virialisation of the halo, or is there an imprint of the primordial power spectrum $P(k)$?  One must make strong assumptions about the connection between density and velocity dispersion to obtain analytical predictions about halo structure. For example, under the assumption that the phase-space density is a power law in radius as suggested by \\citet{2001ApJ...563..483T} one can solve the spherical Jeans equation to obtain the density profile (\\citealt{2004MNRAS.352L..41H}; \\citealt{2005MNRAS.363.1057D}; \\citealt{2005ApJ...634..756A}). These studies, guided by the results of the numerical simulations, confirm the central logarithmic slope to be in the range $\\alpha \\sim 1$ to $2$, although \\citet{2006JCAP...05..014H} have claimed that equilibrated haloes have central slopes of $\\alpha \\sim 0.8$.  However, it is interesting to ask whether or not there is a dependence of halo structure on the primordial power spectrum $P(k)$. In particular, does the form of the profile depend on the (effective) spectral index $n$ of $P(k)$? One of the key predictions of the NFW papers is that the shape of the ``universal'' profile should hold in any hierarchical cosmology, including scale-free ones. This was confirmed by subsequent numerical simulations \\citep[e.g.][]{1996MNRAS.281..716C}. However, a number of analytical studies have claimed that the density profile should depend on spectral index \\citep[e.g., ][]{1985ApJ...297...16H, 1998MNRAS.293..337S, 2000ApJ...538..528S, 2003MNRAS.340.1199H, 2007ApJ...666..181S}, and these have been supported by numerical simulations providing similar evidence \\citep[e.g., ][]{1994ApJ...434..402C, 2001ApJ...554..114E, 2003MNRAS.344.1237R, 2004astro.ph..3352C, 2005MNRAS.357...82R, 2007ApJ...663L..53R}.  Within this context it is interesting to consider the findings of \\citet{2003MNRAS.344.1237R} and \\citet{2004astro.ph..3352C}. Both studies analysed high-resolution cosmological simulations and each argued that dark matter haloes at high redshifts have shallow central logarithmic slopes, with $\\alpha \\approx 0.2-0.4$. In particular, \\citet{2003MNRAS.344.1237R} claimed that the central logarithmic slope depends explicitly on the effective spectral index $n$, in agreement with the predictions of \\citet{2000ApJ...538..528S}. We note that \\citet[][]{2004ApJ...612...50C} performed and analysed simulations comparable to those presented in \\citet[][]{2003MNRAS.344.1237R} and found no evidence for the shallow ``cores'' ($\\alpha \\sim 0.2$); rather they found that their data favoured ``cusps'' ($\\alpha \\approx 1$). However, \\citet{2007ApJ...663L..53R} have recently revisited this topic using simulations with higher mass and force resolution, and they continue to argue in favour of shallow cores, in agreement with their earlier study.\\\\  \\citet{2003MNRAS.344.1237R} and \\citet{2007ApJ...663L..53R} sought to isolate the effect of spectral index $n$ on halo structure by using a series of simulations of the $\\Lambda$CDM model with different box-sizes and evolved for different numbers of expansion factors, to mimic a single $n$. In this short paper we revisit this topic and the claims of \\citet{2003MNRAS.344.1237R} and \\citet{2007ApJ...663L..53R} using high resolution cosmological $N$-body simulations of scale-free models in which we vary systematically the spectral index $n$. This ``clean'' approach allows us to study whether the density profiles of dark matter haloes indeed show a dependence on the spectral index $n$.  In \\S~\\ref{sec:numerical_simulations} we give a detailed description of how we have set up and run our scale-free simulations. There are some important issues to be considered when deciding on \\emph{when to start} a scale-free simulation and \\emph{when to stop} the simulation to compare halo properties. We encapsulate our findings in a set of well defined and physically motivated criteria. In \\S~\\ref{sec:analysis} we present the results of our analysis of the halo mass profiles, and we summarise our results in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  The primary motivation for this paper has been to establish what dependence, if any, the logarithmic slope $\\alpha$ of the inner mass density profile of dark matter haloes has on the spectral index $n$ of the linear matter power spectrum. For this purpose we have run a series of high resolution cosmological $N$-body simulations of scale-free models (i.e. $P(k) \\sim Ak^n$) with values of the spectral index $n$ varying between $-0.5 \\geqslant n \\geqslant -2.75$. By using scale-free models and fixing $n$ in this manner, the problem of identifying correlations between $\\alpha$ and $n$ becomes a relatively straightforward one. Relatively, because there is some freedom in the choice of criteria one can use to determine when to start and when to finish scale-free simulations.  To address this, we have derived clear, reproducible and physically motivated criterion that allows us to set up scale-free simulations and that is presented in \\S~\\ref{sec:starting}. In short, we start the simulations with the same integral power in the $k$-range defined by the the total number of particles $N$; it is especially important to have an objective criterion for starting simulations when studying mass profiles. To determine when to stop our simulations, we follow \\citet{1997ApJ...490..493N} and use the evolution of the typical collapsing mass $M_{*}$. We do so once the mass in a typical $M_{*}$ halo corresponds to approximately $42000$ particles.  Having established a set of well-defined criteria to set up and run cosmological simulations of scale-free models, we performed a sequence of high resolution runs ($1 h^{-1} \\rm Mpc$ boxes, $512^3$ particles) that we used to investigate the dependence of the central logarithmic slope of the dark matter halo density profile $\\alpha$ on the spectral index $n$. We varied $n$ between $-0.5$ and $-2.75$ and selected samples of well resolved haloes ($N_{\\rm vir} \\geqslant 3.15 \\times 10^4$) in dynamical equilibrium in each run (ranging from $\\sim 270$ haloes in the $n=-0.5$ run to $\\sim 40$ haloes in the $n=-2.75$ run). Using $\\chi^2$ fits to a generalised NFW profile, we identified preferred values of the inner slope $\\alpha$ and indeed, we found a trend for the inner slope to become shallower with steeper spectral index -- from $\\alpha \\simeq 1.6$ for $n=-0.5$ to $\\alpha \\simeq 1.2$ for $n=-2.75$.  However, we argue that it is not the central slope $\\alpha$ that depends on spectral index; rather, it is the scale radius of the halo, $r_s$, or our preferred measure $r_{-2}$, the radius at which the differential mass profile $\\rho\\,r^2$ reaches its maximum value. We have shown that haloes in different $n$ models have similar radial profiles of the maximum slope $\\gamma=3(1-\\rho/\\bar{\\rho})$ when normalised to $r_{-2}$. However, as already shown by~\\citet{1997ApJ...490..493N}, haloes that form in models with steeper $n$ tend to be less centrally concentrated than haloes forming in models with shallower $n$, and so their mass profiles can be resolved to smaller fractions of $r_{-2}$, where the flattening of the profile is more apparent. Haloes in these models would then appear to have shallower central profiles.\\\\  We noted in the introduction that using scale-free simulations to study the effect of the spectral index of the power spectrum on the central structure of dark matter haloes was preferable to the approach taken by \\citet{2003MNRAS.344.1237R}, \\citet{2004astro.ph..3352C}, \\citet{2004ApJ...612...50C}, and \\citet{2007ApJ...663L..53R}, in which the box size and analysis redshift were varied to capture the behaviour of different spectral indices. The claims of \\citet{2007ApJ...663L..53R} are of particular interest; these authors argue that halo concentration is a universal constant and that dwarf galaxies identified at $z$=10 have logarithmic slopes shallower than $0.5$.  Our results strongly disagree with these claims. Although the central slope $\\alpha$ that we obtain by fitting a generalised NFW profile does vary with $n$, this does not imply that the shape of the profile is sensitive to $n$, for the reasons presented above. Moreover, as we show in Figure \\ref{fig:maxslope_rr2}, density profiles normalised by $r_{-2}$, which is equivalent to concentration, have similar shapes on average, independent of spectral index; this implies that it is $r_{-2}$ and consequently concentration that depends on $n$. It is possible that our profiles for $\\gamma(r)$ may roll over and approach different asymptotic value at smaller radii than we can resolve with our simulations, but over the range of radii that we can reliably resolve we find no evidence for such behaviour. It should be noted that we checked for this by producing figure~\\ref{fig:maxslope_rr2} for haloes in the high-mass sample (cf. Table~\\ref{tbl:halo_select}) but we did not find any evidence for convergence.\\\\  We briefly consider the dependence of $r_{-2}$ on $n$. We show in figure~\\ref{fig:concentration} that haloes forming in cosmologies with steeper spectral indices tend to be less centrally concentrated than haloes forming in cosmologies with shallower spectral indices, confirming and extending results presented by~\\citet{1997ApJ...490..493N}. This is interesting because we expect that haloes forming in shallow-$n$ models do so predominantly by accretion, whereas haloes in steep-$n$ models form via merging.  If we examine the mass concentration relation in shallow-$n$ models (upper panel in figure~\\ref{fig:concentration}), we find that more massive haloes tend to be less centrally concentrated than their less massive counterparts. In contrast, we find that this relation is far less pronounced in steep-$n$ models.  Prescriptions for halo concentration, such as \\citet{2001ApJ...554..114E}, assert that concentration is determined by the halo's collapse redshift -- haloes that collapse earlier do so when the mean density of the Universe is higher, and so they tend to be more centrally concentrated. It has been argued that this reflects the rate at which the proto-halo accretes mass \\citep[cf.][]{2003ApJ...597L...9Z}. Haloes are said to assemble their mass in two distinct phases -- fast-accretion phase during which the depth of the potential well is set, and a slow-accretion phase during which mass is added to the outer parts of the halo \\citep[e.g][]{2006MNRAS.368.1931L}. This has important implications for the concentration. For example, \\citet{2002ApJ...568...52W} argued that during the fast accretion phase the central density is set by the background density, but as the accretion rate slows and the halo enters its slow accretion phase, the central density stays approximately constant and the halo's concentration grows in step with the virial radius.  Merging also plays an important role; for example, \\citet{2003ApJ...593...26M} looked at the effect of major mergers on concentration and found that violent relaxation drives the remnant's density profile towards a form that is essentially identical to the one it would have formed through pure accretion.  What can we learn about the physical origin of the density profile from our simulations? As we noted above, haloes in shallow-$n$ models grow predominantly by accretion, whereas those in steep-$n$ models grow via merging. This link between spectral index and accretion rate can be made explicit \\citep[see, for example,][]{2006MNRAS.368.1931L}, which makes the scale-free simulation a very interesting tool with which to study the origin of the profile. The results -- that concentrations tend to be higher in shallow-$n$ models, and that concentration decreases with increasing virial mass in these models (assuming that merging plays a role in the growth of the most massive halos) -- indicate that ``two-phase accretion'' may be very important. However, to properly address this question a more detailed study is required. We are currently undertaking such a study, an account of which will be presented in a future paper."
    },
    "0801/0801.4473_arXiv.txt": {
        "abstract": "We consider the influence of microlensing on different spectral bands of lensed QSOs. We assumed that the emitting X-ray, UV and optical regions are different in size, but that the continuum emission in these spectral bands is originating from an accretion disc. Estimations of the time scales for microlensing and flux amplification in different bands are given. We found that the microlensing duration should be shorter in the X-ray (several months) than in UV/optical emitting region (several years). This result indicates that monitoring of the X-ray variations in lensed QSOs that show a 'flux anomaly' can clarify the source of this anomaly. ",
        "introduction": "Recent observational and theoretical studies suggest that gravitational microlensing can induce  variability not only in optical light, but also in the X-ray emission of lensed QSOs \\citep{Chart02a,Chart04,Dai03,Dai04,pj06,Pop01,Pop03a,Pop03b,Pop06a,Pop06b}. Variability studies of QSOs indicate  that the size of the X-ray emitting region is significantly smaller ($\\sim$ several light hours), than the optical and UV emitting regions ($\\sim$ several light days).  Gravitational lensing is achromatic (the deflection angle of a light ray does not depend on its wavelength), but it is clear that if the geometries of the emitting regions at different wavelengths are different then chromatic effects could occur. For example, if the microlens is a binary star or if the microlensed source is extended \\citep{Griest_Hu92,Griest_Hu93,Bog_Cher95a,Bog_Cher95b,Zakh97,Zakh_Sazh98,Pc04} different amplifications in different spectral bands can be present. Studies aiming to determine the influence of microlensing on the spectra of lensed QSOs need to take into account the complex structure of the QSO central emitting region \\citep{Pc04}. Since the sizes of the emitting regions are wavelength dependent, microlensing by stars in the lens galaxy may lead to a wavelength dependent magnification. For example, Blackburne et al. (2006) reported such a 'flux anomaly' in quadruply imaged quasar 1RXS J1131-1231. In particular, they found discrepancies between the X-ray and optical flux ratio anomalies. Such anomalies in the different spectral band flux ratios can be attributed to micro- or milli-lensing in the massive lensing halo. In the case of milli-lensing one can infer the nature of substructure in the lensing galaxy, which can be connected to Cold Dark Matter (CDM) structures (see e.g. Dobler \\& Keeton 2006). Besides microlensing, there are several mechanisms which can produce flux anomalies, such as extinction and intrinsic variability. These anomalies were discussed in \\citet{Pc04} where the authors gave a method that can aid in distinguishing between variations produced by microlensing from ones resulting from other effects. In this paper we discuss consequences of variations due to gravitational microlensing, in which the different geometries and dimensions of emitting regions of different spectral bands are considered.  The influence of microlensing on QSO spectra emitted from their accretion discs in the range from the X-ray to the optical spectral band is analyzed. Moreover, assuming different sizes of the emitting regions, we investigate the microlensing time scales for those regions, as well as a time dependent response in amplification of different spectral bands due to microlensing. Also, we give the estimates of microlensing time scales for a sample of lensed QSOs.  In Section 2, we describe our model of the quasar emitting regions and a model of the micro-lens. In Section 3 we discuss the time scales of microlensing and in Section 4 we present our results. Finally, in Section 5, we draw our conclusions. ",
        "conclusions": "In this paper we calculated microlensing time scales of different emitting regions. Using a model of an accretion disc (in the center of lensed QSOs) that emits in the X-ray and UV/optical spectral bands, we calculated the variations in the continuum flux caused by a straight-fold caustic crossing an accretion disc. We also simulated crossings of accretion discs over microlensing magnification patterns for the case of image A of Q2237+0305 and for a \"typical\" lens system. From these simulations we concluded the following:  (i) one can expect that the X-ray radiation is more amplified than UV/optical radiation due to microlensing which can induce the so called 'flux anomaly' of lensed QSOs.  (ii) the typical microlensing time scales for the X-ray band are on order of several months, while for the UV/optical they are on order of several years (although the time scales obtained from microlensing magnification pattern simulations are longer in comparison to those obtained from caustic simulations).  (iii) monitoring of the X-ray emission of lensed QSOs can reveal the nature of 'flux anomaly' observed in some lensed QSOs.  All results obtained in this work indicate that monitoring the X-ray emission of lensed QSOs is useful not only to discuss the nature of the 'flux anomaly', but also can be used for constraining the size of the emitting region."
    },
    "0801/0801.4968_arXiv.txt": {
        "abstract": "The Canada-France-Hawaii Telescope Supernova Legacy Survey (SNLS) has created a large homogeneous database of intermediate redshift ($0.2 < z < 1.0$) type Ia supernovae (SNe\\,Ia). The SNLS team has shown that correlations exist between SN Ia rates, properties, and host galaxy star formation rates. The SNLS SN\\,Ia database has now been combined with a photometric redshift galaxy catalog and an optical galaxy cluster catalog to investigate the possible influence of galaxy clustering on the SN\\,Ia rate, over and above the expected effect due to the dependence of SFR on clustering through the morphology-density relation. We identify three cluster SNe\\,Ia, plus three additional possible cluster SNe\\,Ia, and find the SN\\,Ia rate per unit mass in clusters at intermediate redshifts is consistent with the rate per unit mass in field early-type galaxies and the SN\\,Ia cluster rate from low redshift cluster targeted surveys. We also find the number of SNe\\,Ia in cluster environments to be within a factor of two of expectations from the two component SN\\,Ia rate model. ",
        "introduction": "\\label{S1Intro}  The importance of understanding Type Ia supernovae (SNe\\,Ia) has risen greatly in recent years because of their pivotal role in demonstrating the accelerated expansion of the universe \\citep{R98,P99,A06,W-V07}. SNe\\,Ia are generally agreed to be carbon-oxygen white dwarfs which have accreted sufficient mass from their companion star to initiate a thermonuclear explosion. Despite this consensus, several models exist for the companion type, the accretion process, the delay time between star formation and SN\\,Ia explosion, and the explosion mechanism \\citep{Hill00,Hoflich03}. Photometric and spectroscopic observations, and explosion simulations, are making inroads towards a coherent picture of SN\\,Ia events \\citep{Mazzali07}. Another (albeit indirect) path of investigation is studying the SN\\,Ia rate in different stellar populations.  Compilation of the Cappellaro et al. (1999) supernova catalog with infrared data led Mannucci et al. (2005) to conclude that SN\\,II, SN\\,Ib/c, and SN\\,Ia rates per unit stellar mass are directly correlated with host morphology and (B-K) colour, and to infer a correlation with host star formation activity. Based on this, Scannapieco \\& Bildsten (2005) expressed the SN\\,Ia rate per unit stellar mass as the sum of a ``delayed'' component from old and a ``prompt'' component from young stellar populations. They parametrized the two components as A and B, proportional to the mass and star formation rate (SFR) of a galaxy respectively, and also found the ``prompt'' component to account for the the discrepancy between low observed rates of cluster SNe\\,Ia and the high cluster iron abundance \\citep{Maoz04}. This ``A+B'' two component model has since been confirmed at intermediate redshifts with the large SNLS catalog \\citep{S06a}; the observations can also be matched with progenitor populations that have distributions of delay times as described by Mannucci et al. (2006) and Pritchet et al. (2007).  From the two component model, the SN\\,Ia rate in galaxy clusters is expected to be lower than in the field due to the morphology-density relation \\citep{Postman84}: cluster galaxies are predominantly of early-type with little or no star formation. However, since clusters contain only a small fraction of the stellar mass of the Universe, it is conceivable that some hitherto undetected influence in such exotic environments could affect the SN\\,Ia rate. For example, the fraction of binary stars, or the rate of mass accretion onto the white dwarf, could be enhanced. The recent discovery of an enhanced nova rate in the core of elliptical galaxy M87 is evidence for the latter \\citep{M07}. A second example is the detected SN\\,Ia rate enhancement in radio-loud early-type galaxies \\citep{DV05}; such galaxies tend to be very luminous ellipticals in the centers of large clusters.  The Wise Observatory Optical Transient Survey (WOOTS) targeting 140 Abell clusters confirmed the SNe\\,Ia rate per unit mass in low redshift galaxy clusters to be consistent with the rate in early-type galaxies \\citep{KS07}. Recently, Mannucci et al. (2007) analyzed the Cappellaro et al. (1999) sample of 136 low redshift SNe ($\\rm z<0.04$) and found the SN\\,Ia rate in cluster early-type galaxies is enhanced by a factor of $\\gtrsim 3$ over field early-type galaxies. The large database of intermediate to high redshift SNe\\,Ia compiled by the Canada-France Hawaii Telescope Supernova Legacy Survey (CFHT SNLS) \\citep{A06}, and the publicly available photometric redshift galaxy and cluster catalogs for SNLS fields \\citep{I06,LFO} are ideal for extending these investigations to higher redshifts. As SNLS is not a cluster-targeted survey, further advantages over Sharon et al. (2007) include using a flexible parametrization of galaxy clustering in the environments of SNe\\,Ia, and determining simultaneous field SN\\,Ia rates for comparison to cluster rates.  \\S~\\ref{S2Obs} describes the SN\\,Ia, galaxy, and cluster catalogs used in this experiment. \\S~\\ref{S3SigParam} and \\S~\\ref{S4LFO} present two independent approaches to statistically evaluate the effects of clustering on the SN\\,Ia rate per unit mass. \\S~\\ref{S3SigParam} uses a continuous clustering strength estimator to compare the SN\\,Ia rate per unit mass inside and outside clustered environments. \\S~\\ref{S4LFO} identifies six SN\\,Ia in Olsen et al. (2007) galaxy clusters, considers the probability of these observations based on expectations from the two component rate model, and calculates the SN\\,Ia rate per unit mass in clusters. \\S~\\ref{S5Conc} reviews and discusses the papers main findings.  A flat cosmology with $\\rm H_0=70 \\ km \\ s^{-1} \\ Mpc^{-1}$, $\\Omega_{Lambda}=0.7$, and $\\Omega_M=0.3$ is used. ",
        "conclusions": "\\label{S5Conc}  Type Ia supernova, galaxy, and cluster catalogs generated from the first four years of CFHTLS Deep survey data were combined to search for an influence of clustering on the SN\\,Ia rate per unit mass. To avoid cluster-specific detection efficiencies and the inclusion of interloping galaxies as cluster members, we also considered subsets of regular and BCG-like early-type galaxies only. Results are dominated by the statistical uncertainties in identifying only three probable and three possible cluster SNe\\,Ia, and consistent with the results of low redshift cluster SNe\\,Ia rate studies. To capitalize on SNLS's inclusion of both field and cluster environments, we used the continuous clustering strength parameter ``significance'' to compare the SN\\,Ia rate per unit mass in and out of clustered environments, but did not find evidence of a clustering influence on the SN\\,Ia rate per unit mass. Future work includes incorporating existing radio and infrared catalogs to investigate the rates of SNLS SNe\\,Ia in radio loud and infrared luminous galaxies, and the VIRMOS spectroscopic redshifts for Deep field galaxies to explore the SNe\\,Ia rate in small groups."
    },
    "0801/0801.4535_arXiv.txt": {
        "abstract": "We derive the SZ effect arising in radio-galaxy lobes that are filled with high-energy, non-thermal electrons. We provide here quantitative estimates for SZ effect expected from the radio galaxy lobes by normalizing it to the Inverse-Compton light, observed in the X-ray band, as produced by the extrapolation to low energies of the radio emitting electron spectrum in these radio lobes. We compute the spectral and spatial characteristics of the SZ effect associated to the radio lobes of two distant radio galaxies (3C294 and 3C432) recently observed by Chandra, and we further discuss its detectability with the next generation microwave and sub-mm experiments with arcsec and $\\sim \\mu$K sensitivity. We finally highlight the potential use of the SZE from radio-galaxy lobes in the astrophysical and cosmological context. ",
        "introduction": "The Sunyaev-Zel'dovich effect (hereafter SZE, Sunyaev \\& Zel'dovich 1972, 1980) is a powerful probe of the physical conditions of electronic plasmas in astrophysical and cosmological context (see Birkinshaw 1999 for a review). A SZE can be produced by thermal (hot and/or warm) electrons in the atmospheres of galaxies and galaxy clusters (Birkinshaw 1999, Itoh et al. 1998, Colafrancesco et al. 2003), by non-thermal (and relativistic) electrons in clusters and in the cavities produced by AGN radio lobes (Ensslin \\& Kaiser 2000, Colafrancesco et al. 2003, Colafrancesco 2005) and by the secondary electrons produced by Dark Matter annihilation in cosmic structures (Colafrancesco 2004), beyond the kinematic SZE related to the bulk motion of such plasmas.\\\\ A SZE from the lobes of radio galaxies is also expected (Ensslin \\& Kaiser 2000, Colafrancesco et al. 2003, Blundell et al. 2006, Colafrancesco 2005, 2007) but it has not been detected yet. Nonetheless, several indirect limits can be set to the SZE from radio-galaxy lobes using both the available radio and X-ray information of these structures.  In fact, detailed studies of radio-galaxy lobes (e.g., Brunetti et al. 2002, Harris \\& Krawczynski 2002, Kataoka et al. 2003, Croston et al. 2005, Blundell et al. 2006 end references therein) have shown that these extended structures contain electrons that have passed through the inner jets and the hotspots and are diffusing on larger-scale lobes where they are currently available to produce inverse-Compton X-ray emission and the relativistic SZ effect which we want to discuss here, beyond the low-frequency synchrotron radio emission.\\\\ In this context, Erlund et al. (2006) found that the extended X-ray emission observed by Chandra along the lobes of three powerful, high-redshift radio galaxies 3C 432, 3C 294 and 3C 191 is likely due to Inverse Compton scattering (ICS) of Cosmic Microwave Background (CMB) photons by the relativistic electrons confined in the lobes. This is consistent with previous findings of X-ray emission from the lobes of FR-II radio galaxies (like 3C223 and 3C284, see Croston et al. 2004) attributed to the same ICS of CMB photons.\\\\ Several studies have been already performed on inverse-Compton scattering in the lobes of radio galaxies and lobe-dominated quasars at low redshift (see e.g. Hardcastle et al. 2002; Comastri et al. 2003; Croston et al. 2004; Bondi et al. 2004; Kataoka \\& Stawarz 2005; Croston et al. 2005; Hardcastle \\& Croston 2005). These studies have shown that the X-ray emission from the lobes of radio galaxies is two-sided and is not dominated by the jet inverse-Compton X-ray emission, as happens instead in powerful, core-dominated quasars (see, e.g., the prototypical case of  PKS 0637-752,  Tavecchio et al. 2000). Only if relativistic beaming effects are important, the jets will dominate the synchrotron and/or inverse-Compton emission on kpc scales. On the contrary, it is clear that relativistic beaming is never important in extended radio lobes.  Detailed studies of the ICS X-ray emission provide constraints to the energetics of the electron population in the lobes. Specifically, a lower limit on the electron energy in the lobe, which is of order of a few times $10^{59}$ erg, has been derived for relativistic electrons that have an energy cutoff $\\gamma_{min} \\sim 10^3$ (see, e.g., Erlund et al. 2006), where $E= \\gamma m_e c^2$ for relativistic electrons. Such a value of the lobe energy has to be considered, however, as a lower limit since the total energy estimate depends on the shape and on the extension of the electron spectrum (see eqs.\\ref{eq.espectrum} and \\ref{eq.energy} below) and might increase by even a large factor for lower values of $\\gamma_{min}$ and steep electron spectra.  Under the assumption that the X-ray emission from radio-galaxy lobes is due to ICS of CMB photons, a simple energy argument yields a value of the minimum electron energy $E_{min} \\approx 0.35 GeV (E_X/keV)^{1/2}$ that corresponds to $\\gamma_{min} \\approx 484$ in order to be detected by Chandra in the energy range $0.5 - 3.0$ keV.\\\\ There are also various arguments indicating that realistic values of $\\gamma_{min}$ are likely in the range $ \\sim 1 - 10^2$ (see e.g. Hardcastle et al. 2002; Comastri et al. 2003; Croston et al. 2004; Bondi et al. 2004; Kataoka \\& Stawarz 2005; Croston et al. 2005; Hardcastle \\& Croston 2005). In this respect, direct constraints on the value of $\\gamma_{min}$ of the order of $\\gamma_{min} \\sim {\\rm a \\, few} \\cdot 10^2 - 10^3$ can be derived in radio galaxy hotspots (see, e.g., Hardcastle 2001). Since we expect that there is adiabatic expansion of the relativistic plasma out of the hotspots, the value of $\\gamma_{min}$ in the radio galaxy lobes should then be substantially lower than the previous estimate, and very likely $< 10^2$ (see, e.g., also the discussion in Blundell et al. 2006).\\\\  Assuming that a unique electron spectrum is responsible for the radio (i.e. synchrotron) and X-ray (i.e. IC) emission features of the lobes, the same electron population that produces X-rays by ICS inevitably produces also an associated SZE which has a specific spectral shape that depends on the shape and on the energy extention of the relative electronic spectrum. Such SZE emission is also expected to be co-spatial with the relative ICS X-ray emission.\\\\ We present here predictions for the relativistic, non-thermal SZE produced by the electrons diffusing within powerful radio-galaxy lobes. A relevant aspect that motivated our study is that more reliable estimates of the lobe SZE are obtained by normalizing the energy of the electronic population residing in the lobes to the available X-ray observations of two specific sources, 3C432 (at $z=1.785$) and 3C294 (at $z=1.779$). We study the amplitude and the spatial distribution of the SZE from the lobes of these sources and provide estimates of detectability with the next generation SZE instruments with arcsec and $\\mu$K sensitivity (like ALMA and/or SPT). We finally highlight the potential astrophysical and cosmological relevance of combining SZE and X-ray observations of radio-galaxy lobes. For a direct comparison of our results with the X-ray observations of Erlund et al. (2006) we use the same cosmological model with $H_0=71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_0=1$ and $\\Omega_{\\Lambda} = 0.73$. ",
        "conclusions": "We have shown in this paper that detailed predictions for the SZE from radio-galaxy lobes can be obtained taking advantage of the constraints set on the energy/pressure by the available X-ray observations. The Chandra data for the lobes of 3C294 indicate that the associated SZE should be visible by the next generation SZ experiments with $\\sim$ arcsec and $\\sim \\mu$K sensitivity, if its electron spectrum extends to quite low energies $E_{min} \\simlt 1$ MeV. Radio galaxy lobes with much flatter spectra, like those of 3C432, result instead swamped by the large contamination of the synchrotron emission in the same radio lobes.\\\\ Steep spectra radio lobes whose IC emission is detected in the X-ray band by Chandra are the best candidates to detect the associated SZE.  We have also shown that the spectral information on the SZE from lobes has a relevant astrophysical impact. The SZE provides, in fact, information on the energetics/pressure and on the optical depth (geometry of the lobe): for the simple case of a power-law spectrum, it yields information on the minimum energy of the electrons and on the minimum projected size of the lobe. The case of more complex spectra can be treated in a similar way. This information cannot be provided by the IC X-ray emission limited in a finite energy band ( e.g., $\\sim 1-10$ keV), and hence to rather high values of $\\gamma_{min} \\sim 10^3$.\\\\ Moreover, since the zero $x_o(P_{e})$ of the SZE is moved to high frequencies depending only on the total electron pressure/energy density (see Figs.\\ref{fig.sz_3c432_dt} and \\ref{fig.sz_3c294_dt}), its detection provides direct information on the total pressure/energy density of the electron population in the lobe.\\\\ We outline various potential uses of the non-thermal SZE detectable in radio lobes in the astrophysical and cosmological contexts: i) constraints on the evolution of the electronic pressure in the radio lobes and the transition to the larger-scale cavities of the intra-cluster medium; ii) estimates of magnetic fields in lobes from a combination of radio and SZE  measurements, possibly obtainable with single experiments with large spectral coverage, like ALMA; iii) a full 3-D tomography of radio galaxy lobes obtainable by means of the next generation microwave polarization experiments; iv) a test of the evolution of the CMB temperature by using the ratio ${(\\Delta T)_{lobe} \\over F_{IC,lobe}}$. These applications are still preliminary but also potentially important with the advent of the next technology.  Most of the non-thermal SZE signal is predicted to be detectable in the outer regions of the radio-galaxy lobes, where the radio synchrotron contamination dims out. Therefore, SZE detectability is likely  possible with instruments with arcsec spatial resolution and $\\mu$K sensitivity (like ALMA) in distant sources, and with a somewhat lower resolution instruments (like SPT) in nearby radio galaxies.\\\\ All the arguments presented in this paper show that the study of the SZE from radio-galaxy lobes are highly complementary to the study of their IC X-ray emission and are also crucial for their astrophysical and cosmological relevance."
    },
    "0801/0801.3315_arXiv.txt": {
        "abstract": "The black-hole binary Cygnus X-1 was observed for 17 ks with the Suzaku X-ray observatory in 2005 October, while it was in a low/hard state with a  0.7--300  keV luminosity  of $4.6\\times 10^{37}$ erg s$^{-1}$. The  XIS and HXD spectra, spanning 0.7--400 keV, were reproduced successfully incorporating a cool accretion disk and a hot Comptonizing corona. The  corona is characterized by an electron temperature of  $ \\sim 100$ keV, and two optical depths of $\\sim 0.4$ and $\\sim 1.5$ which account for the harder and softer continua, respectively. The disk has the innermost temperature of $\\sim 0.2$ keV, and is though  to protrude  half way into the corona. The disk not only provides seed photons to the Compton cloud, but also produces a soft spectral excess, a mild reflection hump, and a weakly broadened iron line. A comparison with the Suzaku data on GRO~J1655$-$40 reveals several interesting spectral differences, which can mostly be attributed to inclination effects assuming that the disk has a flat geometry while the corona is grossly spherical. An intensity-sorted spectroscopy indicates that the continuum becomes less Comptonized when the source flares up on times scales of 1--200 s, while the underlying disk  remains unchanged. ",
        "introduction": "\\label{sec:intro}  Luminous soft X-ray radiation of accreting stellar-mass black holes (BHs) has generally been explained as  thermal emission from optically-thick (in particular ``standard'') accretion disks (Shakura \\& Sunyaev 1973; \\cite{Makishima1986,Dotani1997,RM06}), which are expected to form around them under rather high accretion rates. In contrast, their hard X-ray production process is much less understood, even though intense hard X-ray emission  characterizes black-hoe binaries (BHBs) among various types of  compact X-ray sources  in the Milky Way and Magellanic clouds.  Indeed,  BHBs  often emit a major fraction of their radiative luminosity in the hard X-ray band, in the form of spectral hard-tail  component if  they are in so-called high/soft state, or  as the entire power-law (hereafter PL) like continua if they are in so-called low/hard state (hereafter LHS) which appears under relatively low accretion rates. Furthermore, the hard X-ray emission (particularly in the LHS) involves another interesting aspect, namely the long-known aperiodic variation over a wide frequency range (e.g., \\cite{Oda1971,Oda1977,Noloan1981,Miyamoto1991,Pottschmidt2003,RM06}). These spectral and timing studies are not limited to stellar-mass BHs, since a fair fraction of active galactic nuclei (AGNs) are also considered to be in a state analogous to the LHS of BHBs, emitting luminous hard X-rays (e.g., \\cite{MDK1993}).  In hard X-rays, a  BHB in the LHS emits a roughly PL shaped continuum with a photon index $\\Gamma \\sim 1.7$ (e.g., \\cite{RM06}), but the spectrum gradually steepens toward higher energies. This has been explained  in terms of unsaturated Comptonization of some soft seed photons, in a cloud of hot Maxwellian electrons, or a ``corona\", with a  temperature of $T_{\\rm e}=50-100$ keV and an optical depth of $\\tau \\sim 1$ \\citep{SLE1976,ST1979,ST1980,G1997,DiSalvo2001,Frontera2001,ZG2004,Ibragimov2005}. As a result, the Comtpon $y$-parameter, $y \\equiv 4\\tau kT_{\\rm e}/m_{\\rm e} c^2$, also take a value of $\\sim 1$, where $k$ is the Boltzmann constant, $m_{\\rm e}$ is the electron mass and $c$ is the light velocity. The electrons will be  cooled by transferring  their energies to the  photons, while heated via Coulombic collisions by ions that acquire nearly a free-fall temperature (e.g., \\cite{ZLS1999}). The Comptonizing cloud may be identified with radiatively inefficient inner accretion flow (e.g., \\cite{Ichimaru1977, ADAF95}).  Superposed on the dominant continuum characterizing the LHS, we often observe a spectral hump at energies of  20--40 keV due to  ``reflection'' of the primary X-rays by some cool materials  (e.g., \\cite{LW1988,Inoue1989}), a fluorescent Fe-K emission line (\\cite{Basko1978,Kitamoto1990}), and an apparently ``smeared'' Fe-K edge (\\cite{Ebisawa1994}). These features can be explained by assuming that a fraction of the  primary hard X-rays are reprocessed by  cool materials located near the BH, presumably an optically-thick part of the  accretion disk (\\cite{Done1992,Marshall1993,Ebisawa1996,G1997}). This part of the disk may  at the same time produce an X-ray excess which we often observe at the softest end of the spectrum (e.g., \\cite{BH1991}), and furthermore, may provide the Comptonizing corona with the seed photons  (e.g., \\cite{SLE1976}).  In spite of general success of this currently popular interpretation of the LHS in terms of Comptonizing hot clouds and cool materials, details of the scenario still need to be elucidated. It is  yet to be confirmed that the seed  photons are provided by a cool  accretion disk (e.g., \\cite{ZLS1999}), rather than by thermal cyclotron emission in the Compton corona itself (e.g., \\cite{Kan2001}). The geometry and size of the Comptonizing cloud is not well known, except for loose constraints, including; that the cloud should be large enough for the compactness parameter to remain within a reasonable range (\\cite{LZ1987}), but should not be too large to be able to produce the hard X-ray variation on a $\\sim 1$ s time scale, and that it should have a sufficient vertical height so as to efficiently illuminate the reprocessing matter. Still unsettled is the relation between  the hot corona and the putative cool disk; whether the cool disk is  truncated at a large radius \\citep{Z1998, DZ1999,Z2004}, or it penetrates deep into the hot corona \\citep{Young2001,Miller2006}. Yet another  fundamental issue is how to understand  the violent hard X-ray variability in the context of thermal Comptonization scenario (e.g., \\cite{Miyamoto1987}).  To address these issues associated with the LHS of accreting BHs, a number of spectral studies have been conducted so far, by fitting observed broad-band spectra with theoretical spectral models (e.g., \\cite{G1997,Z1998,Dove1998,Frontera2001}). Similarly, extensive studies in the frequency domain have been conducted to characterize the random variations (e.g., \\cite{MK1989,NKM1994,Nowak1999b,Gilfanov2000,Revnivtsev2000,RM06}). However, the former approach is limited in that different modelings often degenerate. The latter approach has also a limitation, in that it may not  constrain details of the Comptonization scenario without invoking particular models \\citep{KHT1997,HKT1997}.  We analyze broad-band (0.7--400 keV) Suzaku data of Cygnus X-1 (Cyg X-1; Oda 1977; Lian and Nolan 1984) acquired in 2005 October. With the wide-band sensitivity, Suzaku is expected to provide  novel diagnostics of this  BHB, the first celestial object  suspected  to be a black hole  \\citep{Oda1971}. To overcome the above difficulties, we  compare the  Suzaku spectra of Cyg X-1 with those of  GRO~1655$-$40, searching the two spectra for differences attributable to their different system inclination, $i \\sim 45^\\circ$ of Cyg X-1 (\\cite{A2004}) and  $i \\sim 70^\\circ$  of GRO~1655$-$40  (\\cite{G2001}). We hence refer  to the Suzaku results on GRO~1655$-$40 by \\citet{Takahashi1655} as Paper I. We also analyze spectral changes associated with the random variations on a time scale of $\\sim 1$ s (section~\\ref{sec:ana_varspec}), to see how the Compton parameters are changing; this will be  complementary to the study of long-term spectral changes (e.g.,\\cite{Ibragimov2005}). We employ  a distance of  2.5 kpc \\citep{Margon1973,Bregman1973} to Cyg X-1. The errors refer to statistical 90\\% confidence limits. ",
        "conclusions": "Through the analyses in the preceding three subsections, the spectra changes associated with  the fast (1 to 200  s) intensity variations have been quantified in the following manner.   \\begin{enumerate}  \\item The intensity increase from LP to HP is explained by an increase in the seed-photon luminosity, $L_{\\rm seed}$.  \\item The spectral softening  from LP to  HP can be attributed to weakening in the Comptonization of \\texttt{compPSh} (either small $\\tau$, or lower $T_{\\rm e}$, or both).  \\item Unlike  $L_{\\rm seed}$, the directly visible disk luminosity, $L_{\\rm raw}$, does not increase significantly in HP, and could even decrease.   \\item As the source varies, the innermost temperature  of the cool disk is kept approximately constant at $T_{\\rm in}=0.19$ keV.   \\item On these  time scales, the Fe-K line and the reflection hump follow the continuum variations, because $\\Omega/2\\pi$ and the Fe-K line EW are both consistent with being constant.  \\end{enumerate}     \\subsection{Summary of the Data Analysis \\label{subsec:summary_of_analysis}}  The Suzaku observation of Cyg X-1, conducted for an exposure of 17 ks at a 0.7--300 keV luminosity of $4.6 \\times 10^{37}$ erg s$^{-1}$, has provided one of the highest-quality LHS spectra of this prototypical BHB. The XIS, HXD-PIN, and HXD-GSO spectra, altogether covering an extremely broad energy band of 0.7--400 keV, have been explained simultaneously and consistently, by invoking  two basic  constituents of the X-ray emission. One is the hot Comptonizing corona with $T_{\\rm e} \\sim 100$ keV, which manifests itself  in  the high-energy  spectral roll over detected with a high significance by  HXD-GSO. The other is the optically-thick cool disk with $T_{\\rm in}\\sim 0.2$ keV, which produces the soft excess and the Fe-K line in the XIS band, as well as the reflection hump in the HXD-PIN band with $\\Omega/2\\pi \\sim 0.4$.  The data are fully consistent with the cool disk supplying seed photons to the Compton cloud, although other possibilities may not be necessarily excluded. In the present modeling, the soft photons from the cool disk are considered to reach us either directly as the soft excess (the \\texttt{diskbb} component), or through weaker Comptonization with $y \\sim 0.3$ (\\texttt{compPSs}), or after experiencing  stronger Comptonization with $y \\sim 1.15$ (\\texttt{compPSh}).  Assembling together the two Compton components (\\texttt{compPSh, compPSs}) and the three  cool-disk related ones (\\texttt{diskbb}, reflection, and gaussian), we have successfully reproduced the overall continuum shape, including the complicated Fe-K line and soft-excess regions. The Fe-K  line is inferred to be broadened only weakly to $\\sim 1$ keV (in Gaussian $\\sigma$), without any evidence for relativistically broadened wings of which the detection is claimed from some BHBs  \\citep{Miller_diskline}. As seen in figure~\\ref{fig:bestfitmodels}, the inferred incident spectrum shows too complicated a shape, all over the energy range, to be approximated  by a single power-law.   Analyzing in section~\\ref{sec:ana_varspec} the characteristic flare-up behavior of Cyg X-1 on a time scale of $\\sim 1$ s, we found that the \\texttt{wabs*(diskbb+compPS+compPS+gau)} modeling applies also to the HP, LP, and the difference spectra. The flare-up behavior can be understood primarily as an increase in the seed-photon  supply, accompanied by a decreased Comptonization. While the seed-disk luminosity thus increases in HP, the luminosity of the directly visible \\texttt{diskbb} does not increase significantly. As the source varies, $T_{\\rm in}$ of the cool disk is kept  approximately constant.   \\subsection{Comparison with the BeppoSAX results \\label{subsec:bepposax}} Applying the double-\\texttt{compPS} modeling to the BeppoSAX data of Cyg X-1 in the LHS covering up to $\\sim 200$ keV (whereas Suzaku reached 400 keV), \\citet{Frontera2001} derived $T_{\\rm e} = 59 \\pm 5$ keV, $y=0.89 \\pm 0.01$, and $\\Omega/2\\pi = 0.25 \\pm 0.04$ for \\texttt{compPSh} (their  Compton component 1), and $T_{\\rm e} = 42 \\pm 19$ keV and $y=0.15 \\pm 0.01$ for \\texttt{compPSs}  (their \\texttt{copmPS} component 2). Except for some differences in the model parameters, the Suzaku and  BeppoSAX  results are hence consistent in that the  broad-band spectra can be reproduced by a pair of Comptonized components. The two results grossly agree on the  soft-excess parameters as well, although \\citet{Frontera2001} used a blackbody to model the soft excess and the seed-photon source.  The BeppoSAX observation was conducted at  a 0.5--200 keV unabsorbed flux of $4.2 \\times 10^{-8}$ erg cm$^{-2}$ s$^{-1}$, or a luminosity of $3.1 \\times 10^{37}$ erg s$^{-1}$ (in the same band at a distance of 3.2 kpc). The corresponding values from the present Suzaku observation are higher, $5.9 \\times 10^{-8}$ erg cm$^{-2}$ s$^{-2}$ ($4.4 \\times 10^{37}$ erg s$^{-1}$). Then, we would expect Suzaku to observe a lower $T_{\\rm e}$, based on the  general understanding that $T_{\\rm e}$ decreases toward higher luminosities due to enhanced Compton cooling (e.g., \\cite{Esin1998,Yamaoka2005,ZG2004}). Nevertheless, we in fact measured a 1.6 times {\\em higher} $T_{\\rm e}$ than BeppoSAX. One possible explanation to this behavior is that Cyg X-1 may not have fully returned to the LHS at the time of the BeppoSAX observation, which was conducted soon after the source transition from the high/soft state.   Prior to the present work, the double-Compton model was applied successfully  to wide-band LHS spectra of Cyg X-1 (\\cite{G1997,Frontera2001,Ibragimov2005}; though some explicitly invoking multiple values of $T_{\\rm e}$) and GRO~J1655$-$40 (Paper~I, subsection~\\ref{subsec:1655_cyg}). Therefore, we  regard this model as a promising description of the LHS spectra of BHBs. As discussed in Paper~I, the mixture of two Compton optical depths (plus the direct \\texttt{diskbb} component) allows either a ``spatial'' or a ``time-domain'' interpretation. In the former case, some fraction of the cool  disk is considered directly visible, while the remaining  part is covered by a hot electron cloud that has thinner and thicker portions. The latter (time domain) alternative assumes that the Compton cloud is rapidly varying among three typical conditions, represented by $\\tau=0$, $\\tau \\sim 0.4$, and $\\tau \\sim 1.5$. The actual condition may be a combination of these two simplified cases. Of course, as pointed out previously (e.g., Paper~I), the double-Compton modeling may be an approximation to more complex conditions involving more than two  Compton optical depths, or even a continuous distribution in $\\tau$ as suggested by some models (e.g., \\cite{Liu2002}). With the present data, it is virtually impossible to distnguish which is more likely, a continuous distributin in $\\tau$ or just two optical depths.  \\subsection{Comparison with GRO~J1655$-$40 \\label{subsec:1655_cyg}} A comparison between  Cyg X-1 (with $i \\sim 45^\\circ$) and  GRO~J1655$-$40 (with $i \\sim 70^\\circ$) would help elucidating  the source geometry  (section~\\ref{sec:intro}). After a brief attempt made in Paper~1 using the BeppoSAX results on Cyg X-1 \\citep{Frontera2001}, here we perform a more accurate comparison between the two objects using the Suzaku data on both. For this purpose, figure~\\ref{fig:bestfitmodels}b reproduces the best-fit  \\texttt{wabs*(diskbb+compPS+compPS+gau)} model for GRO~J1655$-$40, deduced in Paper~I using  the Suzaku data. On that occasion (2005 September 22 and 23), GRO~J1655$-$40  was in the LHS with a 0.7--300 keV  luminosity of $4.9 \\times 10^{36}$ erg s$^{-1}$ (Paper~I), which  is an order of magnutude lower than that of  Cyg X-1 in the present observation (both assuming isotropic emission). However, the difference becomes a factor of 3--5 when the luminosity is normalized to the Eddington value, because the BH in Cyg X-1  has a mass of $12-20~M_\\odot$, while that in  GRO~J1655$-$40 has $\\sim 6.5~M_\\odot$  (Paper~I).  \\begin{figure}[thb] \\begin{center} \\FigureFile(8cm,){cygx1_figure12.ps} \\end{center} \\caption{Background-subtracted  XIS (black), HXD-PIN (red), and HXD-GSO (green) spectra of GRO J1655-40 (Paper~I), divided by those of Cyg X-1 obtained in the present study.} \\label{fig:1655_cyg_specratio} \\end{figure}   Although the two  models in figure~\\ref{fig:bestfitmodels} generally look  alike, they differ in details. To grasp such  differences in a model-independent manner, we divided the XI2, HXD-PIN, and HXD-GSO spectra of  GRO~J1655$-$40 by the corresponding spectra of Cyg X-1, and obtained the results shown in figure~\\ref{fig:1655_cyg_specratio}. Below, we attempt to identify basic features in this  spectral ratio, and  to interpret them in terms of the spectral decomposition in figure~\\ref{fig:bestfitmodels} and table~\\ref{tbl:parameters}. Like in figure~\\ref{fig:HiLo_spec}b, the behavior of the ratio at the highest XIS energy range should be ignored, where the GRO~J1655$-$40 data may suffer from photon pile-up effects (Paper~I).  The  spectral ratio in figure~\\ref{fig:1655_cyg_specratio} is roughly constant over  $\\sim 3$ to $\\sim 30$ keV, but decreases significantly below $\\sim 2$ keV. This effect is too large to be explained by their difference in absorption, $N_{\\rm H} = 6.6 \\times 10^{21}$ cm$^{-2}$ in Cyg X-1 and $7.4\\times 10^{21}$ cm$^{-2}$ in GRO~J1655$-$40. Instead, it can be attributed to the stronger soft excess of Cyg X-1 (figure~\\ref{fig:bestfitmodels}), as the \\texttt{diskbb} component of Cyg X-1 is an order of magnitude more prominent  than that of GRO~J1655$-$40 when normalized to the overall continuum (while $T_{\\rm in}$ is nearly the same). This in turn  may be ascribed primarily to their difference in the $\\cos i$ factor, $\\cos(45^\\circ)/\\cos(70^\\circ)=2.1$, because we expect the direct disk emission to scale as $\\propto \\cos i$, while the Comptonized emission to be approximately isotropic. Strictly speaking,  this  estimate needs some caution. In subsection~\\ref{subsec:HiLo_compPSfits}, we introduced two quantities, $L_{\\rm raw}$ and $L_{\\rm seed}$. Because $R_{\\rm in}$  used to calculate $L_{\\rm raw}$ has already been corrected for the $\\cos i$ factor while $R_{\\rm in}^{\\rm seed}$ is not (subsection~\\ref{subsec:comppsfits}, table~\\ref{tbl:parameters}), we expect the ratio $L_{\\rm raw}/L_{\\rm seed}$ to be approximately inclination independent. Nevertheless, the ratio was measured to be $\\sim 0.44$ in Cyg X-1, while $< 0.19$ in GRO~J1655$-$40 (table~\\ref{tbl:parameters}). That is, the soft excess in GRO~J1655$-$40 is even weaker than would be predicted by the assumed $i=70^\\circ$. This is discussed later in subsection~\\ref{subsec:variation}.  In figure~\\ref{fig:1655_cyg_specratio}, we observe a small dip at $\\sim 6.5$ keV. Evidently, this feature arises because the Cyg X-1 spectrum exhibits a stronger Fe-K line (with an equivalent width of  $\\sim 300$ eV) than that of GRO~J1655$-$40 ($\\sim 80$ eV). This property is again ascribed, at least qualitatively, to the inclination effects: the cool flat disk will intercept a fraction of the Comptonized photons, and produce the fluorescent Fe-K photons, of which  the visibility must be proportional to $ \\cos i$. Quantitatively,  however, the difference in the equivalent width between the two sources amounts to a factor of $\\sim 4$, which exceeds the factor $\\sim 2$ difference  expected from $i=70^\\circ$. This is also discussed  in subsection~\\ref{subsec:variation}.   Over the 30--200 keV range, the ratio in figure~\\ref{fig:1655_cyg_specratio} increases with energy the  $E$ roughly as  $\\propto E^{0.3}$: GRO~J1655$-$40 has a harder continuum. This is consistent with  the PL approximation to their  hard X-ray spectra, which gave $\\Gamma=1.75$ for  Cyg X-1 (section~\\ref{subsec:convfits}) and $\\Gamma=1.39$ for GRO~J1655$-$40 (Paper~I). In terms of the Comptonization scenario, this slope difference can be attributed to the higher $y$-parameter of \\texttt{compPSh} in GRO~J1655$-$40  ($1.31^{+0.11}_{-0.08}$) than in Cyg X-1 ($1.15^{+0.05}_{-0.03}$). This in turn can be explained rather securely by the higher $T_{\\rm e}$ of GRO~J1655$-$40, unlike the case of the HP vs. LP comparison of Cyg X-1 in which the origin of the change in $y$ remained ambiguous. As discussed in Paper~I and subsection~\\ref{subsec:bepposax}, the higher $T_{\\rm e}$ of GRO~J1655$-$40 may be due to its lower luminosity.  In the PIN energy range, the ratio of figure~\\ref{fig:1655_cyg_specratio} exhibits a  slightly concave curvature. Because the continuta themselves in this energy range should be  simple power-laws dominated by \\texttt{compPSh}, this effect must be caused by  differences in the reflection features. In fact, figure~\\ref{fig:bestfitmodels} reveals two effects. One is that the reflection hump of Cyg X-1, relative to the Compton continuum, is  inferred to be stronger than that in GRO~J1655$-$40. This is  again considered an   inclination effect, because the reflected component is expected to be proportional to $\\cos i$ while the primary continuum is thought to be  rather isotropic. Actually, the values of $\\Omega/2\\pi$ in table~\\ref{tbl:parameters}, which is calculated  {\\it after correction} for $\\cos i$, becomes comparable between the two objects. The other effect is that the reflection hump in Cyg X-1, modeled with  $i = 45^\\circ$, emerges  at 30--50 keV, whereas that of  GRO~J1655$-$40, modeled with $i = 70^\\circ$, is peaked at higher energies. This is because a larger $i$ makes the reflection more forward directed, and hence the reflected photons suffer less Compton degradation \\citep{MZ1995}. These two effects in combination are considered to produce the slightly concave-shaped spectral ratio in the PIN band. Incidentally, the comparable values of $\\Omega/2\\pi$ between the two objects disfavors the idea that the hot Comptonizing corona may have a semi-relativistic outflow velocity away from the disk plane  \\citep{Malzac2001}, because in that case we would observe a smaller value of $\\Omega/2\\pi$ from Cyg X-1 due to the enhanced continuum.   To summarize, the spectrum of Cyg X-1 and GRO~J1655$-$40 are both dominated by the Comptonized continua, but Cyg X-1 exhibits more prominently the flat-disk-related spectral features presumably due to its lower inclination. Another inclination effect manifests itself in the  subtle differences in the 10--70 keV continuum shape, presumably due to inclination-dependent changes in the strength and peak energy of the reflection hump. These results are consistent with the view that the Comptonizing hot cloud is geometrically thick, and hence its emission  depends much less on the inclination than  the flat-disk related components.   Analyzing  broad-band BeppoSAX spectra of XTE~J1118+480, \\citet{Frontera2001b} discovered that this transient BHB has a very weak soft excess and a weak reflection hump. Like the present discussion on GRO~J1655$-$40, these properties of XTE~J1118+480 may  also be ascribed to the inclination effects, since this object  is also reported  to  have a rather high inclination as  $i=(68 \\pm 2)^\\circ$ \\citep{Gelino2006}.   \\subsection{The disk parameters in the low/hard state} \\label{subsec:diskpara}   In the high/soft state of Cyg X-1, a \\texttt{diskbb} fit to the ASCA data showed the optically-thick disk to have $r_{\\rm in}= 71$~km \\citep{Dotani1997}, which translate via equation~(\\ref{eq:rin2Rin}) to $R_{\\rm in}=90$~km for $i=30^\\circ$ \\citep{Makishima2000} or $R_{\\rm in}= 100$~km for $i=45^\\circ$. The disk is securely considered to reach the last stable orbit \\citep{Dotani1997,Makishima2000}. In the LHS, a similar cool disk is  considered to be present as well, but its conditions become significantly more difficult to estimate, because its emission is overwhelmed and strongly modified by the Compton process. As a result, there is considerable controversy over the innermost disk radius in the LHS.  Theoretically, the Advection Dominanted Accretion Flows  \\citep{ADAF95} and other related hot flow models, meant to explain the LHS,  all assume that the cool disk is truncated at some radius  larger than the last stable orbit. (One important difference is that the original ADAF model assumes synchrotron photons as the soft seed photons to be Compton up-scatterd into hard X-rays, whereas the present scenario assumes the cool disk to supply  the seed soft photons.) A gradual inward shift in this truncation radius will give rise to the correlated spectral softening, increased reflection, and increased variation power frequencies, as observed  (e.g., \\cite{Gilfanov1999,ZLS1999,Z2004,ZG2004,Done2007}). When the disk reaches the last stable orbit, a  transition to the high/soft state will take place. However, this view has been challenged by some LHS observations which suggest the disk to extend down to the last stable orbit \\citep{Miller2006,MHM2006,Rykoff2007}. Here we use the broad-band Suzaku data to infer the disk parameters in the LHS.   In  subsection~\\ref{subsec:convfits}, the conventional  \\texttt{diskbb+cutoffpl} fit to the present LHS data tentatively gave $r_{\\rm in} ~\\sim 13$ km, and hence $R_{\\rm in}= 18$~km for $i=45^\\circ$. This apparently implies that the cool disk extends down to a much smaller  radius than in the high/soft state. However,  as pointed out before \\citep{Ebisawa1996,DiSalvo2001,Ibragimov2005} and revealed by the present data (figure~\\ref{fig:bestfitmodels}), the LHS spectra of Cyg X-1 and other BHBs are significantly more complex than is described by a single power-law. In fact, our final solution invoking double \\texttt{compPS} has given an order of magnitude larger value, $R_{\\rm in} = 250$ km (table~\\ref{tbl:parameters}), which is consistent with the disk inner radius receding by a factor 2.5 compared to the high/soft state. The large difference in $R_{\\rm in}$  between the two modelings for the same data can be explained in the following manner. In our final model, the flux in the 1--2 keV range is mostly explained by  \\texttt{compPSs} (figure~\\ref{fig:bestfitmodels}a), with \\texttt{diskbb} carrying softer residuals. If, instead, a single straight continuum is employed, the \\texttt{diskbb} component becomes unusually hot trying to explain the 1--2 keV excess \\citep{DiSalvo2001,Ibragimov2005}, and hence the inner disk radius becomes apparently too small. This gives a lesson that disentangling the disk parameters, where the disk is not the dominant component in the spectrum, is subject to large uncertainties in the underlying continuum shape. Evidently, the case of GRO~1655$-$40 is subject to larger difficulties because of the weaker soft excess.   The  Comptonization scenario  involves another important effect which artificially reduces the disk radius: the disk photons are ``lost''  from the \\texttt{diskbb} component into the dominant Comptonized continua. In estimating the true disk radius, to be denoted $R_{\\rm in}^{\\rm tot}$, we must  take this effect into account, knowing that Comptonization conserves the photon number \\citep{KME2001,Aya+Max2004,Aya+Chris2004}. Practically, $R_{\\rm in}^{\\rm tot}$ can be derived by quadratically summing $R_{\\rm in}$ of \\texttt{diskbb}, $R_{\\rm in}^{\\rm seed}$ of \\texttt{compPSh}, and $R_{\\rm in}^{\\rm seed}$ of \\texttt{compPSs} as \\begin{equation} \\left(R_{\\rm in}^{\\rm tot} \\right)^2 = R_{\\rm in}^2 + \\left(R_{\\rm in-h}^{\\rm seed}\\right)^2 +\\left(R_{\\rm in-s}^{\\rm seed}\\right)^2~~. \\label{eq:Rtot} \\end{equation} Then, equations~ (\\ref{eq:Lseed}),  (\\ref{eq:Lraw}), and (\\ref{eq:Rtot}) readily yield the total luminosity $L_{\\rm disk}$ of the cool disk  as \\begin{equation} L_{\\rm disk} = L_{\\rm raw}+L_{\\rm seed} =4\\pi \\left({R_{\\rm in}^{\\rm tot}} \\right)^2 \\sigma_0 T_{\\rm in}^4~. \\label{eq:Ldisk} \\end{equation} From the values  in table~\\ref{tbl:parameters}  for the time-averaged spectrum, we  obtain $R_{\\rm in}^{\\rm tot} \\sim 330$ km, which is equivalent to  $\\sim 15R_{\\rm g}$ for $M_{\\rm BH}=15~M_\\odot$. The implications are that roughly half the disk area is directly visible, and the disk inner radius is  factor 3 larger than in the high/soft state.  Although the estimated value of $R_{\\rm in}^{\\rm tot}$ may still be relatively small compared with the original idea of disk truncation, it  can be subject to yet additional uncertainties which  take part in when transforming the raw \\texttt{diskbb} normalization  into a physical radius. The stress-free inner boundary condition is probably appropriate for a thin disk extending down to the last stable orbit. In contrast, a truncated disk may have continuously rising stress to its inner radius. Similarly, the environment of the disk is different between the two states: in the LHS, the disk is likely to be more strongly heated by irradiation and conduction from the hot cloud. Therefore, the color hardening factor in the LHS  may well be significantly larger than the value of 1.7 employed in equation~(\\ref{eq:rin2Rin}); then, the true effective temperature would be lower, and hence $R_{\\rm in}$ would become still larger.   Three independent pieces of evidence further argue against the small disk radius in the LHS. One is  the relatively small values of reflection, $\\Omega/2\\pi \\sim 0.4$, which requires the cool seed disk not to intrude too deeply into the hot corona. To reconcile the small values of $\\Omega/2\\pi$ with  a configuration wherein  the seed disk is nearly completely surrounded by the corona, we would have to invoke, e.g., mild relativistic coronal outflows (e.g., \\cite{Malzac2001}), but we already argued against it (subsection~\\ref{subsec:1655_cyg}). Second, the intrinsic width of the Fe-K line, when interpreted as relativistic broadening in the disk near the BH, can constrain the innermost locations of the cool disk. For this purpose, we re-fitted the average spectra with the final model, but with the Gaussian replaced by so-called \\texttt{diskline} model \\citep{Fabian89}. When fixing the inclination  at $45^\\circ$, and assuming the line emissivity to  scale as $r^{-3}$, the fit constrained the rest-frame line center energy as $6.3_{-0.1}^{+0.2}$  keV, and the innermost disk radius for the Fe-K line production as $13_{-7}^{+6} \\; R_{\\rm g}$  which is fully consistent with the value of $R_{\\rm in}^{\\rm tot}$ derived above. Finally,  as quoted in subsection~\\ref{subsec:diffspec}, \\citet{Revnivtsev1999} studied fast variations of Cyg X-1 in the Fe-K line energy region, and argued  the cool reprocessor to be located farther than $\\sim 100 R_{\\rm g}$ for $M_{\\rm BH}=15~M_\\odot$. Although the implied radius is much larger than considered here, this should be taken as another evidence supporting  the disk truncation. (Incidentally, the present results on the fast variability in the Fe-K lines and the Compton humps constrain their production sites to be closer than $ 1.3 \\times 10^4~R_{\\rm g}$.)   Based on these examinations, we estimate that the cool-disk radius of Cyg X-1 during the present LHS observation is $R_{\\rm in}^{\\rm tot} \\sim 15R_{\\rm g} \\sim 330$ km nominally, but is subject to rather large uncertainties, and is likely to be considerably larger than this.   All the considerations presented here will also apply to previous LHS studies, where the small inner disk radius derived from simple continuum fits was used to challenge the truncated disk model \\citep{Miller2006,MHM2006,Rykoff2007}. In short, the small disk radii claimed for some BHBs in the LHS would need a careful revision.   \\subsection{A possible interpretation of the fast X-ray variation} \\label{subsec:variation}  The remaining issue is how to interpret our results on the time variation, summarized in subsection~\\ref{subsec:sec4summary} and subsection~\\ref{subsec:summary_of_analysis}. In a standard accretion disk which we believe is approximately correct in the present case, the disk parameters are mutually related as $T_{\\rm in } \\propto (R_{\\rm in}^{\\rm tot})^{-3/4} \\dot{M}^{1/4}$ (e.g., \\cite{Makishima2000}), where $\\dot{M}$ is the mass accretion rate. Since  $T_{\\rm in}$ has been confirmed to change little as  Cyg X-1 varies, the variation must also keep  $R_{\\rm in}^{\\rm tot}$ and $\\dot{M}$ approximately constant. (An alternative case, namely, positively correlated changes in $R_{\\rm in}^{\\rm tot}$ and $\\dot{M}$, would be unlikely, because an increase in $\\dot {M}$ would enhance radiative cooling, and hence {\\em reduce} $R_{\\rm in}^{\\rm tot}$; \\cite{ZLS1999,ZG2004,Done2007}.) As  a result, the total disk luminosity $L_{\\rm disk}$, defined by equation~(\\ref{eq:Ldisk}), should also be kept rather constant. However, $L_{\\rm seed}$  has been found to  increase from LP to HP (subsection~\\ref{subsec:HiLo_compPSfits}). Therefore,  $L_{\\rm raw}$ must decrease in HP by 20--30\\%. As  confirmed in subsections \\ref{subsec:HiLo_compPSfits} and \\ref{subsec:diffspec}, the data permit (though do not require) such a decrease in $L_{\\rm raw}$ from LP to  HP.  Based on the above arguments, we presume that Cyg X-1 gets brighter as a larger fraction of the disk photons are intercepted by the Comptonizing corona, while the underlying cool disk itself is kept essentially unchanged. This will cause an increase in $L_{\\rm seed}$, and the associated  decrease in $L_{\\rm raw}$. Although this inference does not necessarily specify how the corona is actually changing, we may consider one possible scenario below.  Let us, for example, assume that the Compton cloud has a rather spherical shape, or a thick torus geometry, with its outer radius  located at $R_{cc} \\sim 100~R_{\\rm g} \\sim 2 \\times10^3$ km where the free-fall time is $\\sim 0.1$ s. (The actual in-fall time scale of the corona may well be longer, depending on the viscosity.) Based on our discussion made in subsection~\\ref{subsec:diskpara}, the cool disk is considered to  partially protrude into the corona, but truncated at, say, $R_{\\rm in}^{\\rm tot}  \\sim 50 ~R_{\\rm g}$. The Comptonization will then take place both in the overlapping region (between $R_{\\rm in}^{\\rm tot}$ and $R_{\\rm cc}$), and inside $R_{\\rm in}^{\\rm tot}$, of the corona. The corona is assumed to be rather turbulent with strong density inhomogeneities, so that we need two optical depths, $\\tau_{\\rm s} \\sim 0.4$ and $\\tau_{\\rm s} \\sim 1.5$, to express its effects. Furthermore, we expect $T_{\\rm e} $ to increase as the corona falls toward the BH, for the following two reasons. One is because, compared to the corona inside $R_{\\rm in}$, that in the overlapping region must be more strongly cooled by the disk photons, and less efficiently heated by protons due to lower densities. The other is because,  at about $R_{cc}$, the disk surface layer must evaporate into the corona, so it will take a fraction of a second (calculated at $R_{\\rm cc}$; \\cite{ei_coupling}) for the evaporated electrons to be heated by protons via Coulomb interaction. Then, the  observed $T_{\\rm e}\\sim 100$ keV should be considered as a radial average over the corona.  Under the configuration as assumed above, we further speculate that, for some unspecified reasons, the corona has a number of  ``holes'' (at least in the overlapping region between the disk and corona), through which the disk is directly visible (subsection~\\ref{subsec:bepposax}). With these, the  spectral composition of Cyg X-1 (and of GRO~J1655$-$40) can be  explained at least semi-quantitatively. Furthermore, this scenario can explain why the directly visible disk and the Fe-K lines of GRO~J1655$-$40 are weaker than those of Cyg X-1 beyond what is expected by the inclination difference, because the holes would act as a vertical collimator to direct these disk-related emission components toward pole-on directions. The disk reflection would also be affected, but thier intrinsically broad shape would remain rather intact.  In order to explain the fast variability, we may assume that the opening fraction of the coronal holes vary with time, due to  some fluctuations in the disk evaporation rate. As the hole area diminishes by an increased evaporation, we expect $L_{\\rm seed}$ to increase and $L_{\\rm raw}$ to decrease, as observed in the changes from LP to HP. Since the incremented portion of the corona must initially have a lower electron temperature due to the finite ion vs. electron coupling time, the HP spectrum is expected to be less Comptonized than in LP, again as observed. Electrons in that portion of the corona will be gradually heated up, on time scales of a fraction of a second, as they fall inwards across $R_{\\rm in}^{\\rm tot}$. This will  explain, at least qualitatively, the XIS to PIN phase lag which  would be too large to be explained by  time-of-flight delays of photons in the Compton process.      Through the detailed analysis of the Suzaku data of Cyg X-1, we  have obtained the following firm results. \\begin{enumerate} \\item The  0.7--400 keV Suzaku spectra are reproduced by a sum of a cool disk emission, two Compton continua both with reflection ($\\Omega/2\\pi \\sim 0.4$), and a mildly broad  Fe-K line.  \\item The  two Compton continua, with $\\tau \\sim 0.4$ and $\\sim 1.5$, can be described by a common temperature of  $T_{\\rm e} \\sim 100$ keV, although the data do not exclude them having different temperatures.  \\item The disk is characterized by $T_{\\rm in} \\sim 0.2$ keV and  $R_{\\rm in} \\gtrsim 330$ km ( $\\gtrsim 15~R_{\\rm g}$). Roughly half the  disk emission gets  Comptonized, while the rest forms the spectral soft excess.  \\item When Cyg X-1 brightens up on a time scale of 1--200 s, a larger fraction of disk photons becomes Comptonized, although the disk parameters remain unchanged. The spectrum softens, as a result of a decrease in either $T_{\\rm e}$ or $\\tau$.  \\item A comparison between Cyg X-1 (with $i \\sim 45^\\circ$) and GRO~J1655$-$40 (with $i \\sim 70^\\circ$) supports the view that the disk has a flat geometry and the Compton cloud is more spherical. \\end{enumerate}  Combining these results with some theoretical considerations, we propose the following speculative scenario on the LHS of Cyg X-1 (and possibly of other BHBs). That is, the Componizing corona, with a roughly spherical (or thick torus) geometry, extends up to, say, $\\sim 100 R_{\\rm g}$, and the cool disk protrudes  half way into it. The corona is highly inhomogeneous (as required by the double Compton $\\tau$), and has openings (as required by the soft excess) to allow a direct view of part of the disk. The hard X-ray intensity increases when the opening area of the corona decreases and hence a larger number of disk photons become Comptonized.     \\bigskip  {\\it Acknowledgement} \\ The authors would like to express their hearty thanks to all the members of the Suzaku Science Working Group, for their help in the spacecraft operation, instrumental calibration, and data processing. This work was partially supported by the Grant-in-Aid for Scientific Research on Priority Areas (Grant No. 14079101).  \\clearpage"
    },
    "0801/0801.2313_arXiv.txt": {
        "abstract": "Surface effects in strange-quark matter play an important role for certain observables which have been proposed in order to identify strange stars, and color superconductivity can strongly modify these effects. We study the surface of color superconducting strange-quark matter by solving the Hartree-Fock-Bogoliubov equations for finite systems (``strangelets'') within the MIT bag model, supplemented with a pairing interaction. Due to the bag-model boundary condition, the strange-quark density is suppressed at the surface. This leads to a positive surface charge, concentrated in a layer of $\\sim$ 1 fm below the surface, even in the color-flavor locked (CFL) phase. However, since in the CFL phase all quarks are paired, this positive charge is compensated by a negative charge, which turns out to be situated in a layer of a few tens of fm below the surface, and the total charge of CFL strangelets is zero. We also study the surface and curvature contributions to the total energy. Due to the strong pairing, the energy as a function of the mass number is very well reproduced by a liquid-drop type formula with curvature term. ",
        "introduction": "From rather general arguments it is expected that at low temperatures and high densities quark matter is in a color superconducting state~\\cite{colorsup}. More recently~\\cite{Alford97,Rapp97} it has been suggested that the diquark pairing gaps for quark matter at densities of several times nuclear matter saturation density could be of the order of $\\sim 100$ MeV. Since this could have important phenomenological consequences in particular for the interior of compact stars, this has triggered much work on color superconductivity in dense quark matter (for reviews, see, e.g., Ref.~\\cite{reviews}). These investigations of the QCD phase diagram have revealed a very rich phase structure with many different possible pairing patterns, depending on external conditions such as, for instance, electrical neutrality or quark masses. The largest diquark pairing gaps arise from scalar condensates, leading either to the two-flavor color superconducting (2SC) phase or to the color-flavor-locked (CFL) phase~\\cite{CFL,weakCFL}. The latter pairing pattern involves strange ($s$) quarks, in addition to the two light quark flavors, up ($u$) and down ($d$).  If color superconducting quark matter exists in nature, the most likely place to find it is the interior of compact stars because matter is compressed there to densities much higher than nuclear matter saturation density. However, it has been argued that strange-quark matter (SQM) might be absolutely stable~\\cite{Witten}. Under this hypothesis, even pure strange stars should exist~\\cite{strangestars}, i.e., stars entirely composed of SQM. Also small lumps of SQM, called ``strangelets,'' might be stable. Because of their low charge to baryon number ratio $Z/A$, strangelets have been proposed to populate ultra-high energy cosmic rays~\\cite{cosmicrays}.  In SQM without pairing, the density of strange quarks is supposed to be smaller than that of light quarks because of their higher mass. Consequently, SQM and strangelets are positively charged and the charge neutrality of strange stars has to be achieved via the presence of electrons. At the surface an atmosphere of electrons forms~\\cite{strangestars} which can potentially be detected~\\cite{Usov,UsovPage} via the emission of electron-positron pairs from an extremely strong electric field at the surface.  Recently another possible picture of the surface of a strange star has been proposed~\\cite{nuggets}: there could be a ``crust'' composed of strangelets immersed in an electron gas. Similar to an ordinary neutron star, there could be an interface between the crust and the interior in form of the famous ``pasta phases.'' Within this scenario the electric field at the surface would be strongly reduced. Obviously, surface effects for the strangelets play an important role for the description of this scenario. For instance, there is a critical surface tension deciding whether a homogeneous phase or the droplet phase is favored~\\cite{stablenuggets}.  Another question for which surface effects should be considered is the formation of a strange star in a supernova explosion. Before the explosion the original star contains hadronic matter. During the formation of the star, nucleation of strangelets sets in, leading then to a conversion of the entire star to SQM. For the nucleation process the properties of small strangelets are important.  Pairing tends to reduce the differences in density of different quark species.  For bulk quark matter in the CFL state, requiring color neutrality, all quarks are paired. The densities are thus equal and CFL quark matter is electrically neutral on its own, i.e., without any electrons~\\cite{RW01}. This would suggest dramatic changes in the properties of strangelets and SQM inside compact stars. For instance, the electrosphere at the surface of a strange star could completely disappear. But, the presence of the surface can modify this picture since it can lead to a non-zero surface charge which remains even for large objects. For example, the boundary condition of the MIT bag model suppresses the density of the massive strange quarks at the surface, resulting in a positive surface charge~\\cite{Madsen01}. Within this scenario, the total charge of a strangelet, following roughly $Z \\approx 0.3 A^{2/3}$, is drastically reduced with respect to ``normal'' strangelets. For strange stars, this requires the presence of electrons~\\cite{Usov04}. However, pairing has not been treated self-consistently in previous work (see, e.g., Ref.~\\cite{Madsen01}). In this paper we will therefore reinvestigate finite-size strangelets with pairing by considering quark matter in a color superconducting spherical bag, solving the Hartree-Fock-Bogoliubov (HFB) equations. We will show, in particular, that there exist CFL type solutions where all quarks are paired and the total charge of the strangelet strictly vanishes.  The outline of the paper is as follows. In Section~\\ref{sec:model} we will present our model for treating color superconducting quark matter in a finite volume. In Section~\\ref{sec:results} we will show numerical results. In Section~\\ref{sec:solutions} we discuss the possibility of qualitatively different configurations. In Section~\\ref{sec:chdens} we concentrate on the charge-density distributions of the CFL like solutions. In Section~\\ref{sec:liquiddrop} we discuss a liquid-drop like mass formua for the CFL-like solutions and calculate the surface tension. Finally, in Section~\\ref{sec:sum} we will summarize our results. ",
        "conclusions": "\\label{sec:sum} In this paper we have investigated finite lumps of color superconducting SQM. To that end we have treated the MIT bag model, supplemented with a pairing interaction, in the framework of HFB theory. This allows us to correctly include finite size effects for pairing, too. The calculation is numerically rather involved, since in addition to solving self-consistently the HFB equations, we have to determine the bag radius and the fractions of the different quark species by minimizing the total energy of the system.  As expected from previous MIT bag-model studies, we find a suppression of the strange-quark densities at the surface, resulting in a positive surface charge. Our main result is that, in spite of this surface charge, the total charge of the CFL type solution is zero due to pairing, as in bulk matter. Most of the positive surface charge is compensated in a negatively charged layer situated at about 1-3 fm below the surface. The origin of this concentration of the negative charge is pairing: Since all quarks are paired, the positive surface charge must be compensated on a length scale corresponding to the coherence length. The remaining negative charge, which is necessary to compensate all of the positive surface charge, is situated below this layer. With increasing distance from the surface, the charge density decreases on a length scale of $\\sim 8$ fm, corresponding to the Debye screening length. This number will probably be strongly decreased if the gluons are included in a perturbative way similar to the photon. In any case, in the biggest part of a large object, such as a strange star, one finds vanishing charge density if one goes more than a few tens of fm away from the surface. It remains to be investigated in which way our results change the traditional picture of the surface of a strange star and the detectability of smaller strangelets in current experiments such as AMS-02 or LSSS~\\cite{Madsen06}.  We have also compared our results for the energy per baryon of finite strangelets with a liquid-drop like formula. We obtain a surface tension of the order of 12-14 MeV, in reasonable agreement with previous studies where color superconductivity was not considered, and a strong curvature term which is crucial to reproduce the correct energies up to baryon numbers of several thousands. An interesting result is that, in the presence of color superconductivity, the liquid-drop formula describes very accurately the total energies even for $A\\lesssim 100$, at least for strangelets with even baryon number. The reason is that, since the gap $\\Delta$ is much larger than the spacing between the energy levels, shell effects are strongly suppressed."
    },
    "0801/0801.2255_arXiv.txt": {
        "abstract": "% We summarize the main properties of the extended UV (XUV) emission found in roughly 30\\% of the nearby spiral galaxies observed by the GALEX satellite. Two different classes of XUV disks are identified, the Type~1 XUV disks where significant, structured UV-bright features are found beyond the {\\it classical} azimuthally-averaged star-formation threshold, and the Type~2 XUV disks, which are characterized by very extended (seven times the area where most of the stellar mass is found), blue [(FUV$-$$K$)$<$5\\,mag] outer disks. These latter disks are extreme examples of galaxies growing inside-out. The few XUV disks studied in detail to date are rich in HI but relatively poor in molecular gas, have stellar populations with luminosity-weighted ages of $\\sim$1\\,Gyr, and ionized-gas metal abundances of $\\sim$Z$_{\\odot}$/10. As part of the CAHA-XUV project we are in the process of obtaining deep multi-wavelength imaging and spectroscopy of 65 XUV-disk galaxies so to determine whether or not these properties are common among XUV disks. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.3968_arXiv.txt": {
        "abstract": "We describe Monte Carlo models for the dynamical evolution of the nearby globular cluster M4.  The code includes treatments of two-body relaxation, three- and four-body interactions involving primordial binaries and those formed dynamically, the Galactic tide, and the internal evolution of both single and binary stars.  We arrive at a set of initial parameters for the cluster which, after 12Gyr of evolution, gives a model with a satisfactory match to the surface brightness profile, the velocity dispersion profile, and the luminosity function in two fields.  We describe in particular the evolution of the core, and find that M4 (which has a classic King profile) is actually a {\\sl post-collapse} cluster, its core radius being sustained by binary burning.  We also consider the distribution of its binaries, including those which would be observed as photometric binaries and as radial-velocity binaries.  We also consider the populations of white dwarfs, neutron stars, black holes and blue stragglers, though not all channels for blue straggler formation are represented yet in our simulations. ",
        "introduction": "The present paper opens up a new road in the study of the dynamical evolution of globular clusters.  We adopt the Monte Carlo method of Giersz \\citep{Gi1998,Gi2001,Gi2006}, which in recent years has been enhanced to deal quite realistically with the stellar evolution of single and binary stars, to study the dynamical history of the nearby globular cluster M4.  An earlier version of the code had already been used to study the dynamical history of $\\omega$ Cen \\citep{GH2003}, but at that time the treatment of stellar evolution was primitive and there were no binaries.  The new code has been thoroughly tested on smaller systems, by comparison with $N$-body simulations and observations of the old open star cluster M67 \\citep{GH2008}.  There we showed that the Monte Carlo code could produce data of a similar level of detail and realism as the best $N$-body codes.  Now for the first time we consider much richer systems, with about half a million stars initially, which are at present beyond the reach of  $N$-body methods.  This paper has a place within a long tradition of the modelling of globular star clusters, but the place is a distinctive one.  First, we are not concerned with a static model of a star cluster at the present day, like a King model.  We are concerned with issues where the dynamical history of the star cluster is important, where static models are uninformative.  Secondly, our aim is to construct a model of a specific star cluster, rather than trying to understand the general properties of the evolution of a population of star clusters. This has been done before, and a brief history is outlined in \\citet{GH2008}, but the present work takes these efforts onto a new level of realism, in terms of the description of stellar evolution, and dynamical interactions involving binary stars.  This problem is not easy.  Not only is it necessary to use an elaborate technique for simulating the relevant astrophysical processes, but it is necessary also to search for initial conditions which, after about 12Gyr of evolution, lead to an object resembling a given star cluster.  By ``resembling'' we do not simply mean matching the overall mass, radius and binary fraction of a cluster, for example, for two reasons:\\begin{enumerate} \\item  We have found that values found for data in the literature are highly uncertain, and different sources are contradictory.  These data are usually derived, in some model-dependent way, from such data as surface-brightness profiles and velocity dispersion profiles, and we prefer to compare our models directly with this data, and not with inferred global parameters. \\item  We have found that, even if one achieves a satisfactory fit to these profiles, the model may give a very poor comparison with the luminosity function. \\end{enumerate} From these considerations we conclude that a model which aims to fit only the mass and radius of a star cluster (say) may be very far from the truth.  Tackling this difficult problem is not just interesting, however.  We have been motivated by a number of pressing astrophysical problems. For example, the two nearby star clusters M4 (the subject of this study) and NGC 6397 (which we shall consider in our next paper in this series) have rather similar mass and radius, and yet one has a classic King profile, while the other is a well-studied example of a cluster with a ``collapsed core'' \\citep{Tr1995}.  Among possible explanations one may consider differences in the population of binaries, which are known to affect core properties, or in tidal effects.  Indeed the present paper will show that these two clusters may be much more similar than one would suppose from the surface brightness profiles alone.  A second motivation for our work is our involvement in observational programmes aimed at characterising the binary populations in globular clusters.  What differences (e.g. in the distributions of periods and abundances) should one expect to find between the core and the halo?  These issues are important in the planning of observations, and in their interpretation.  The cluster M4 is the focus of much of this effort because it is nearby, making it a relatively  easy target for deep observational study.  It was the first globular cluster to yield a deep sequence of white dwarfs \\citep{Ri1995}.  More recently it has been subjected to an intensive observational programme by the Padova group \\citep{bedinetal2001,bedinetal2003,andersonetal2006}, which includes searches for radial-velocity binaries in the upper main sequence \\citep{Somm}. It also turns out to be a cluster which (we conclude) started with only about half a million stars, which facilitates the modelling.  Along with the open cluster M67, M4 was chosen by the international MODEST consortium, at a meeting in Hamilton in 2005, as the focus for joint effort by theorists and observers, to cast light on its binary population and dynamical properties.  M67 has been modelled very successfully by \\citet{Hu2005}, using $N$-body techniques, and this paper represent the first theoretical step in a similar study of M4.  The paper is organised as follows.  First, we summarise features of the code and the models, the data we used, and our approach to the problem of finding initial conditions for M4.  Then we present data for our best models: surface brightness and velocity dispersion profiles, luminosity functions, the properties of the binary population, white dwarfs and other degenerate remnants, and the inferred dynamical state of the cluster.  The final section summarises our findings and discusses them in the context of work on other clusters, including objects to which we will turn in future papers. ",
        "conclusions": "In this paper we have presented a Monte Carlo model for the nearby globular cluster M4.  This model includes the effects of two-body relaxation, evaporation across the tidal boundary, dynamical interactions involving primordial and three-body binaries, and the internal evolution of single stars and binaries.  By adjustment of the initial parameters (total mass, tidal radius, initial mass function, binary fraction) we have found a model which, after 12Gyr of evolution, leads to a model with a surface brightness profile, velocity dispersion profile, and luminosity functions at two radii, all of which are in tolerable agreement with observational data.  It also leads to a number of photometric binaries roughly consistent with observation.  This model has a current mass of $4.6\\times10^4\\msun$, a $M/L_V$ ratio of 1.8, a binary fraction of almost 6\\%, and an almost flat lower mass function (Table \\ref{tab:mc_king}).  Almost 40\\% of the mass is in white dwarfs, and about 7\\% in neutron stars, though our model assumes a 100\\% retention rate.  They are strongly mass-segregated to the centre of the cluster.  There are no stellar-mass black holes left in the cluster.  The binaries have experience significant dynamical evolution, almost all the soft pairs having been destroyed.  The binary population as a whole is only slightly segregated towards the centre, but there is more evident segregation of bright binaries (such as those that would be more readily observed in a radial velocity search.)  The most significant new result from our model is the implication that M4 is a core collapse cluster, despite the uncollapsed appearance of the surface brightness profile.  Now we discuss a number of shortcomings of our model and other issues related to this study. \\begin{enumerate} \\item {\\sl Uniqueness:}  Though we have arrived at a broadly satisfactory model, it is not at all clear how unique it is. Certainly we were unable to construct models with much larger numbers of primordial binaries, or with significantly larger initial radius, as such models produced an insufficiently concentrated surface brightness profile.  Furthermore, the fact that the lower slope of the initial mass function is less steep than the canonical value of 1.3 should not be taken to imply that we have inferred the initial value. \\item {\\sl Fluctuations}: we have noticed that runs differing only in the initial seed of the random number generator can give surprisingly different surface brightness profiles.  In broad terms this confirms the important finding of \\citet{Hu2007}, though we have not yet established (as he did) that the presence or absence of black hole binaries is the underlying mechanism.  We shall return to this issue at appropriate length in the next paper in this series, on the cluster NGC 6397. \\item {\\sl The Initial Model:}  The initial model is astonishingly dense, the central density being of order $10^6\\msun$pc$^{-3}$.  It also has to be realised that we are imposing initial conditions at a time when the residual gas from the birth of the cluster has already dispersed.  Various properties of a star cluster, including its binary population and mass function, may change significantly during the phase of gas expulsion, which further undermines the power of our model to establish the initial conditions. Finally, we have assumed no initial mass segregation, even though recent work suggests that this should be present already before the cluster has assembled into a roughly spherical object \\citep{Ve2007}.  This is an aspect of the modelling that could be readily improved. \\item {\\sl The early evolution of the model:}  It is worth drawing attention here to the high initial escape velocity from the centre of our model, in the first few tens of millions of years.  Clearly this is dependent on the small initial radius.  The initial high density is also responsible for the very rapid initial evolution of the core. \\item {\\sl Imperfections of the modelling:} Several important improvements need to be made to our technique. \\begin{enumerate} \\item   At present few-body interactions are handled with cross sections.  The problem is not simply that these are not well known, especially in the case of unequal masses; it also means that we are unable to determine if a collision occurs during a long-lived interaction.  For this reason, we have said almost nothing about collision products (e.g. blue stragglers) in this paper.  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. \\item Long-lived triples  are neglected in the model at present.  These are 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}. \\item We assume that the tide is modelled as a tidal cutoff.  We have taken some care to ensure that the tidal radius is adjusted so that the model loses mass through escape at the same rate as an $N$-body model would, if immersed in a tidal {\\sl field} with the same tidal radius.  This means, however, that the effective tidal radius of the model is somewhat too small.  A better treatment of a steady tide may be possible, without this drawback. \\item We assume that the tide is steady, something which is not true for M4.    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.   On the other hand \\citet{Ag1988} found, on the basis of a simple model, that tidal evaporation was the dominant mechanism in the evolution of the mass of M4 at the present day, and that other factors (such as disk and bulge shocking) contributed at a level less than 1\\%. \\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 The search for initial conditions is still very laborious, and we are constantly seeking ways of expediting this.  Our most fruitful technique at present is the use of small-scale models which relax at the same rate as a full-sized model.  Not all aspects of the dynamics are faithfully rendered by a scaled model, but the technique appears to be successful as long as the fraction of primordial binaries is not too large. \\end{enumerate} \\end{enumerate}  Despite these caveats and shortcomings, it is clear that the Monte Carlo code we have been developing is now an extraordinarily useful tool for assessing the dynamics of rich star clusters.  For the cluster M4 it has led to the conclusion that M4 can be explained as a post-collapse clusters.  It also provides a wealth of information on the distribution of the binary population, which will be important for the planning and interpretation of searches for radial velocity binaries, now under way.  $N$-body models will eventually supplant the Monte Carlo technique, but at present are incapable of providing a star-by-star model of even such a small cluster as M4.  They can of course provide a great deal of general  guidance on the dynamical evolution of rich star clusters, and they underpin the Monte Carlo by providing benchmarks for small models.  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."
    },
    "0801/0801.4586_arXiv.txt": {
        "abstract": "White dwarfs (WDs) with carbon absorption features in their optical spectra are known as DQ WDs.  The subclass of peculiar DQ WDs are cool objects ($T_{\\rm eff}\\lesssim 6000$~K) which show molecular absorption bands that have centroid wavelengths $\\sim$100--300\\,\\AA\\ shortward of the bandheads of the \\ct\\ Swan bands.  These ``peculiar DQ bands\" have been attributed to a hydrocarbon such as \\cth.  We point out that \\cth\\ does not show strong absorption bands with wavelengths matching those of the peculiar DQ bands and neither does any other simple molecule or ion likely to be present in a cool WD atmosphere.  \\ct\\ seems to be the only reasonable candidate for producing the peculiar DQ bands.  Many characteristics of those bands can be understood if they are pressure-shifted Swan bands. While current models of WD atmospheres suggest that, in general, peculiar DQ WDs do not have higher photospheric pressures than normal DQ WDs do, that finding requires confirmation by improved models of WD atmospheres and of the behavior of \\ct\\ at high pressures and temperatures. If it is eventually shown that the peculiar DQ bands cannot be explained as pressure-shifted Swan bands, the only explanation remaining would seem to be that they arise from highly rotationally excited \\ct\\ ($J_{peak}\\gtrsim 45$). In either case, the absorption band profiles can in principle be used to constrain the pressure and the rotational temperature of \\ct\\ in the line-forming regions of normal and peculiar DQ WD atmospheres, which will be useful for comparison with models.  Finally, we note that progress in understanding magnetic DQ WDs may require models which simultaneously consider magnetic fields, high pressures and rotational excitation of \\ct. ",
        "introduction": "\\label{intro}  Carbon white dwarfs (DQ WDs) with $T_{\\rm eff}\\lesssim 11000$~K can show absorption from the Swan bands of \\ct. Peculiar carbon white dwarfs (DQ PEC or DQp WDs) are relatively cool WDs ($T_{\\rm eff}\\lesssim 6000$~K) which exhibit bands reminiscent of the Swan bands, but with more rounded profiles and with centroid wavelengths located $\\sim$100--300\\,\\AA\\ (700$\\pm$100~cm$^{-1}$) shortward of the bandheads of the Swan bands (Table \\ref{t1}; see also Figure 3 of {Schmidt}, {Bergeron}, \\& {Fegley} 1995 and Figure 30 of {Bergeron}, {Ruiz}, \\& {Leggett} 1997). The origin of these ``peculiar DQ bands\" or DQp bands has been a matter of some debate. {Liebert} \\& {Dahn} (1983) first suggested they were pressure-shifted Swan bands, but {Bergeron} {et~al.} (1994) cast doubt on that interpretation (see \\S~\\ref{nkT}). Strong magnetic fields can and do distort the Swan bands in some carbon WDs, but the DQp WDs are not all strongly magnetic ({Schmidt} {et~al.} 1995, 1999), so another explanation is required for their bands. One well-received explanation was put forward by Schmidt et al. (1995), who attributed the DQp bands to a hydrocarbon such as \\cth\\ rather than to \\ct. As a result, DQp WDs have also been referred to as \\cth\\ WDs. However, until the origin of the bands in such objects is conclusively settled, we prefer to use the nomenclature of DQp WDs and DQp bands.  In this contribution we report an unsuccessful literature search for known transitions from \\cth\\ and other simple molecules that correspond to the DQp bands (\\S~2).  We then revisit the issue of how \\ct\\ might produce the DQp bands (\\S~3).  We summarize our conclusions in \\S~4. ",
        "conclusions": "\\label{end}  We have shown that the molecular absorption bands in peculiar DQ WDs are very unlikely to be produced in any molecule other than \\ct\\ (\\S~2).  We have considered how the DQp bands can be explained as modified Swan bands (\\S~3) and conclude that they could either be pressure-shifted Swan bands (as originally suggested by {Liebert} \\& {Dahn} 1983) or Swan band absorption from highly rotationally excited levels, or both. Regardless, a continuum of increasingly shifted and reshaped bands is expected to be observed in DQp WDs, in agreement with recent observations (\\S~\\ref{rovib}).  One implication of our result is that the putative absence of WDs with $T_{\\rm eff}\\lesssim 6000$~K which exhibit molecular carbon absorption ({Dufour} {et~al.} 2005) is not real --- the DQp WDs are such objects. Another implication is that understanding the spectra of magnetic DQ WDs may require models which simultaneously account for the effects of pressure, rotational excitation and magnetic fields on \\ct\\ absorption.  Theoretical and experimental investigations of the properties of \\ct\\ in high-pressure, high-temperature helium are needed to determine if pressure-shifted Swan bands resemble the DQp bands.  If they do not, then the only remaining explanation for the DQp bands seems to be that they arise in rotationally excited \\ct\\ ($J_{peak}\\gtrsim 45$). That scenario would require a higher \\trot\\ in DQp WDs than in DQ WDs, perhaps due to \\ct\\ absorption arising deeper in the photosphere in DQp WDs ({Bues} 1999).  Even if pressure-shifted Swan bands are found to resemble the DQp bands, existing models of WD atmospheres suggest that DQp WDs do not have higher photospheric pressures than DQ WDs and thus that the peculiar DQ bands are not pressure-shifted Swan bands. Nonetheless, since there are known shortcomings in existing models of cool, He-rich WD stellar atmospheres, that tentative conclusion needs to be revisited when improved models become available.\\footnote{More speculatively, it would also be worth investigating the properties of carbon-atmosphere WDs ({Dufour} {et~al.} 2007b) at low \\teff, to see if they have high photospheric pressures. Although {Dufour} {et~al.} (2007b) state that no carbon-atmosphere WD is known at \\teff$<15,000$~K, carbon-atmosphere WDs were not thought to exist at all before their recent discovery. We cannot be confident that all carbon-atmosphere WDs develop a helium atmosphere as they cool until the origin of such objects is better understood.  If predictions can be made of the spectra of carbon-atmosphere WDs at low \\teff, the SDSS provides a large database in which to search for objects with such spectra.} If DQp WDs do have higher photospheric pressures than DQ WDs, the reason may be that low metal abundances reduce the He$^-$ opacity by restricting the availability of donor electrons.  Whichever explanation of the DQp bands is eventually confirmed, with accurate theoretical models it should be possible to use \\ct\\ absorption band profiles as a sort of barometer to constrain both the pressure and the rotational temperature of \\ct\\ in the line-forming regions of DQ and DQp WD atmospheres.  For example, examination of Swan band fits to DQ WD absorption profiles shows that the observed profiles are often more rounded than predicted ({Dufour} {et~al.} 2005; {Koester} \\& {Knist} 2006). Thus, the observed profiles may provide evidence for an increasing $J_{peak}$ or an increasing pressure shift with decreasing $T_{\\rm eff}$ even among normal DQ WDs.  However, these discrepancies could simply be due to imprecise knowledge of Swan band absorption parameters or to inexact treatment of the WD atmospheres or of \\ct\\ itself.  For example, {Dufour} {et~al.} (2005) calculate \\ct\\ line profiles using the impact approximation, which is known to break down in dense DQ WD atmospheres where the mean distance between particles is comparable to the particle size ({Koester}, {Weidemann}, \\& {Zeidler} 1982).  Observationally, more and better parallax values for DQ and DQp WDs would provide information on the range of masses seen in each type of WD and help determine what contribution different values of $g$ make to the DQp phenomenon. It might also be useful to search for and study other \\ct\\ bands in normal and peculiar DQ WDs. Whatever physical effect transforms the Swan bands into the DQp bands, its effect on other \\ct\\ energy states, and thus other \\ct\\ bands, may be different and may serve as a clue to its origin."
    },
    "0801/0801.1640_arXiv.txt": {
        "abstract": "We detect 353 X-ray point sources, mostly low-mass X-ray binaries (LMXBs), in four {\\em Chandra} observations of Centaurus~A (NGC 5128), the nearest giant early-type galaxy, and correlate this point source population with the largest available ensemble of confirmed and likely globular clusters associated with this galaxy. Of the X-ray sources, 31 are coincident with 30 globular clusters that are confirmed members of the galaxy by radial velocity measurement (2 X-ray sources match one globular cluster within our search radius), while 1 X-ray source coincides with a globular cluster resolved by HST images. Another 36 X-ray point sources match probable, but spectroscopically unconfirmed, globular cluster candidates. The color distribution of globular clusters and cluster candidates in Cen A is bimodal, and the probability that a red, metal rich GC candidate contains an LMXB is at least 1.7 times that of a blue, metal poor one. If we consider only spectroscopically confirmed GCs, this ratio increases to $\\sim$3. We find that LMXBs appear preferentially in more luminous (massive) GCs. These two effects are independent, and the latter is likely a consequence of enhanced dynamical encounter rates in more massive clusters which have on average denser cores. The X-ray luminosity functions of the LMXBs found in GCs and of those that are unmatched with GCs reveal similar underlying populations, though there is some indication that fewer X-ray faint LMXBs are found in globular clusters than X-ray bright ones. Our results agree with previous observations of the connection of GCs and LMXBs in early-type galaxies and extend previous work on Centaurus~A. ",
        "introduction": "\\label{sec:intro}  Observations of early-type galaxies with the {\\em Chandra} X-ray Observatory have shown that a significant fraction of low-mass X-ray binaries (LMXBs) in these galaxies are associated with globular clusters (GCs) \\citep[see][for a review]{fab06}. The fraction of LMXBs identified with known GCs around massive early-type galaxies varies from at least 20\\% in NGC~4697 \\citep{sarazin01} to a remarkable 70\\% in NGC~1399, the central giant elliptical of the Fornax cluster \\citep{angelini01}.  In the Milky Way $\\sim$10\\% of all bright LMXBs are found in GCs, even though the GCs account for $<10^{-3}$ of the stellar mass of the galaxy \\citep[e.g.][]{katz75,grindlay93}.  The overabundance of LMXBs in GCs is thought to be due to their high central densities, which result in greatly enhanced rates of dynamical interactions with respect to the field. These interactions can lead to the creation of LMXB progenitors via dissipative encounters between neutron stars and ordinary stars \\citep[e.g.][]{fpr75}, or encounters in which a compact object replaces a member of a binary in a three-body interaction \\citep[e.g.][]{hills76}. This picture is supported by observations in the Milky Way and nearby early-type galaxies \\citep{pooley03, jordan04, siv07, jordan07} which show that estimates of GC dynamical collision rates are good predictors for the presence of LMXBs in GCs.  It has even been suggested \\citep{white02} that almost all LMXBs are formed in the cores of GCs and either are later ejected from their host clusters due to the evaporation of the GCs or escape the cluster due to dynamical processes or kick velocities  imparted at birth. Recent analyses suggest that this is not the case though, as observations seem to require the presence of a bona fide population of field LMXBs besides a component whose formation is connected to GCs \\citep{irwin05,juett05}.  Moreover, \\cite{fabbiano07} have recently shown that GCs in NGC 3379 lack low luminosity LMXBs when compared to the field population, implying that the latter cannot have originated from GCs \\citep[see also ][]{voss07}. To date, large populations of LMXBs have been found in X-ray observations of E/S0 galaxies over a wide range of luminosities \\citep[e.g.][]{siv07,kundu07} and in large disk galaxies such as M31, NGC~3115, and NGC~4594 \\citep{fan05,kundu03,distefano03}.  In this paper, we present results from a correlation of the LMXB population, detected as point sources in Chandra/ACIS observations, with the optically detected GC population in the galaxy Centaurus~A (Cen~A, NGC 5128). A previous study of the connection between GCs and LMXBs in Cen~A, using smaller optical and X-ray catalogs, is presented in \\citet{m04}. Cen~A is the nearest massive early-type galaxy ($M_B\\! =\\! -21.1$) \\citep{dufour79}, and has been widely studied across the electromagnetic spectrum \\citep[for a comprehensive review, see][]{israel98}.  X-ray observations of Cen~A \\citep{feig81,turner97,kra00,kra02,kra03a,kra03b,hardcastle07} reveal several distinct emission components, including a bright, compact and variable nucleus, an X-ray jet perpendicular to the dust lane and aligned with radio features, and diffuse emission representing the hot phase of the interstellar medium.  In addition, several hundred bright point sources are seen in X-rays \\citep{kra01,vg06}, mostly representing the LMXB population.  Studying Cen~A has major advantages over studies of the LMXB/GC connection in other early-type galaxies. First, Cen~A is $\\sim$5 times closer than the often-studied galaxies in the Virgo and Fornax clusters, which allows for a much deeper census and better characterization of its X-ray binary population.  Second, at the distance of Cen~A, a scale of $1^{\\prime\\prime}$ corresponds to a linear distance of just 18~pc, which is comparable to the physical diameters of GCs.  Furthermore, Cen~A has a moderately large GC population \\citep[$\\simeq 1500$ clusters; see][]{harris06}, with a bimodal distribution in color, equally split between red and blue.  An intriguing aspect of the observations linking GCs with X-ray point sources has been that LMXBs in early-type galaxies are more likely to be found in the redder (higher-metallicity $Z$) GCs.  This trend already can be seen in the small Milky Way sample of LMXBs, but becomes much more evident in larger galaxies \\citep[e.g.][]{angelini01, kundu02,distefano03,jordan04,kim06,siv07,poss07,kundu07}. The physical mechanisms underlying this observational result are still unknown, although several possibilities have been put forward, including metallicity dependent variations in the initial mass function \\citep{grindlay87}, the link between metallicity and the outer convective zones of stars \\citep{ivanova06} and the effects of metallicity on radiation induced stellar winds \\citep{maccarone04}.  A recent spectroscopically based survey of the cluster ages ($\\sim 150$ GCs with Lick index measurements) by \\citet{beasley07} indicates that the great majority of the GCs are old, with ages $\\tau \\gtrsim 10$ Gyr.  However, perhaps 20\\% of them (almost all on the metal-rich side) scatter to younger ages in the 6 to 8 Gyr range.  For comparison, \\citet{rej05} show that the field halo stars in Cen~A have a mean inferred age of $8^{+3}_{-3.5}$ Gyr.  In the Milky Way, the metal-rich clusters, including the LMXB hosts, are mostly old systems \\citep{rosenberg99}. The two giant galaxies M87 and M49 in Virgo are found to have mostly old GC systems \\citep{cohen98,jordan02,puzia99,beasley00} and also show a preference for metal-rich GCs to host LMXBs \\citep{siv07}. Thus the preference of X-ray binaries to be in higher metallicity clusters does not seem to be associated with age and must be primarily due to metallicity \\citep[see also][]{kundu03}.  Recent work by \\cite{jordan07} exploits the distance to Cen~A in order to explore the detailed dependence of the incidence of LMXBs in GCs on the structural parameters of GCs such as their half-light radii and central densities.  \\cite{jordan07} 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.  Various studies indicate that the half-light radii of metal-rich GCs in large E galaxies are $\\sim 20\\%$ smaller than the metal-poor clusters (see \\cite{kundu01}, \\cite{jordan05}, and \\cite{gw07} among others).  \\cite{jor04} shows that this could be a result of mass segregation and the longer lifetimes of lower metallicity stars for a given mass, under the assumption that the half-mass radii distribution does not depend on metallicity. In this case, we expect dynamical effects to be largely decoupled from the metallicity of GCs, but if the half-light radii reflect indeed a difference in half-{\\it mass} radii at least part of the preference of LMXBs for metal-rich GCs could be due to their smaller average size and therefore denser cores. It is easy to show that dynamical processes cannot be responsible for all of the observed effect of metallicity on the incidence of LMXBs, even if there was a size difference consistent with that observed in other early-type galaxies. Assuming encounter rates scale with half-light radii $r_h$ as $r_h^{-2.5}$ \\citep[see ][]{siv07} metal-rich GCs would need to be $\\sim 50\\%$ smaller than their metal-poor counterparts in order to fully account for a threefold enhancement of the number of LMXBs in metal-rich GCs. This level of size difference is ruled out by the observations.  The metallicity itself seems to be an important factor in producing LMXBs.  This paper is organized as follows.  We introduce the observations and discuss data reduction and point source detection in \\S\\ref{s:obs}.  In \\S\\ref{s:matching}, we identify GCs from optical observations which are associated with the X-ray point sources, and discuss their properties in \\S\\ref{s:gcmatches}.  In \\S\\ref{s:xraypts}, we examine the X-ray properties of the point sources that are associated with GCs and compare with those not found in GCs.  Finally, results are summarized in \\S\\ref{s:remarks}. In this analysis, we adopt a distance of $3.8 \\pm 0.2$ Mpc for Cen~A \\citep[an average of five precise standard candles including TRGB, PNLF, SBF, LVPs, and Cepheids; see][and references therein]{mcl07, ferrarese07} and an integrated optical luminosity $M_V^t = -22.0$ \\citep{dufour79}. ",
        "conclusions": "\\label{s:remarks}  We have presented a study of the connection between GCs and LMXBs using a ground-based catalog of GCs and an X-ray point source list constructed using four Chandra/ACIS X-ray observations of the nearby early-type galaxy Cen~A, obtained over a four-year period, each with an exposure time of 37--50~ks.  Of the 353 X-ray point sources detected on one or more of these Chandra fields, 67 are found to be hosted by known GCs or GC candidates. Thirty of these matches correspond to GCs confirmed by radial velocity, and 1 match corresponds to a resolved GC by HST imaging. Even in the central heavily obscured dust lane, we find matches between the two populations, although the optical samples are significantly incomplete in that region.  As in many other early-type galaxies \\citep[e.g. ][]{peng06}, the GCs in Cen~A have a large range in metallicity and a bimodal color distribution.  A three-to-one majority of the matched X-ray point sources are seen to lie in the redder, more metal-rich GCs: specifically, 156/335 ($44\\%$) of the confirmed GCs are in  the metal-rich subpopulation, while 23/29 ($79\\%$) of the X-ray point sources matched to confirmed GCs lie in these clusters.  We also note a possible double X-ray point source detection in the radially velocity confirmed GC0233.  Also in agreement with previous studies, we find that luminous GCs are most likely to host LMXBs. Quantitatively, the number of LMXBs per unit GC mass is roughly constant.  The lack of correlation between GC size and mass implies that denser GCs are more likely to host LMXBs, as expected if most GC LMXB are formed dynamically by tidal capture or exchange processes in GC cores.  Our expanded sample of GC-LMXB identifications in Cen~A gives further support for a consistent picture which has been emerging from combined optical and {\\it Chandra} observations of early-type galaxies. Important factors determining the presence of LMXBs in GCs which can be easily determined from ground-based observations are their mass, which is a reflection of a more fundamental dependence on dynamical formation rates \\citep{jordan07}, and their metallicity. The latter is a bona fide effect which seems not to follow from any possible trends of GC structural properties or age with metallicity.  The X-ray luminosity function of the LMXBs that are found in GCs and of those that are not matched with GCs are found to have the same slope, revealing similar underlying populations. In fact, if one samples the areas of the galaxy that have more complete coverage for both the LMXB and GC populations (away from the dust lanes and within 10\\arcmin\\ of the galaxy center), half of the LMXBs are found within GCs.  However, there is some indication that, towards the faint end of the luminosity function, the fraction of LMXBs found in GCs is far lower than at the more luminous end, but it is not clear whether this is due to incompleteness effects at the faint end. Cen~A offers a unique opportunity to study this question because of its proximity, and we will probe the LMXB/GC connection with our upcoming deep observations to X-ray luminosities unattainable in more distant early-type galaxies. In addition, we plan to use high-resolution optical imaging material, from both ground based telescopes and from the HST/ACS, to probe the structural properties of the GCs with and without LMXBs, as well as the {\\sl Chandra} spectral properties of the X-ray sources."
    },
    "0801/0801.1689_arXiv.txt": {
        "abstract": "{V838 Mon erupted at the beginning of 2002 becoming an extremely luminous star with L=10$^6$~L$_{\\sun}$. The outburst was followed by the most spectacular light echo that revealed that the star is immersed in a diffuse and dusty medium, plausibly interstellar in nature. Low angular resolution observations of the star and its closest vicinity in the lowest CO rotational transitions revealed a molecular emission from the direction of V838~Mon. The origin of this CO emission has not been established.} {The main aim of this paper is to better constrain the nature of the CO emission. In particular, we investigate the idea that the molecular emission originates in the material responsible for the optical light echo.} {We performed observations of 13 positions within the light echo in the two lowest rotational transitions of $^{12}$CO using the IRAM 30~m telescope.} {Emission in CO $J=\\,$1--$\\,$0 and $J=\\,$2--1 was detected in three positions. In three other positions only weak $J=\\,$1$-\\,$0 lines were found. The lines appear at two different velocities $V_{\\rm LSR}=\\,$53.3~km~s$^{-1}$ and $V_{\\rm LSR}=\\,$48.5~km~s$^{-1}$, and both components are very narrow with {\\small FWHM}$\\,\\approx\\,$1~km~s$^{-1}$.} {The molecular emission from the direction of V838~Mon is extended and has a very complex distribution. We identify the emission as arising from diffuse interstellar clouds. A rough estimate of the mass of the molecular matter in those regions gives a few tens of solar masses. The radial velocity of the emission at 53.3~km~s$^{-1}$ suggest that the CO-bearing gas and the echoing dust are collocated in the same interstellar cloud.} ",
        "introduction": "}  V838~Mon is a star that brightened significantly at the beginning of 2002 becoming a very luminous object with a luminosity of 10$^6$~L$_{\\sun}$. The outburst lasted for $\\sim$3~months and was characterized by a rather complex light curve with three distinct maxima seen in the optical. Several broad P-Cygni profiles observed during the brightening imply that the event was accompanied by an outflow with terminal velocities of several hundred~km~s$^{-1}$ \\citep[e.g.][]{kip04}. The end of the eruption was characterized by a very steep decline with a total drop in brightness in the V band of about 8~mag within one month. Unlike novae, the object evolved from a hot stage with an F-type spectrum to progressively lower and lower temperatures and eventually it turned into a supergiant cooler than M10 \\citep{eva03}. Named the first L supergiant, the star has exhibited an extraordinary molecular spectrum studied extensively in optical \\cite[e.g.][]{pav06} and infrared \\citep[e.g.][]{lyn04}. Even five years after the eruption, the object continues to evolve - after four years of a relatively constant photometric brightness, an eclipse-like event has been observed in the B and V light curves \\citep{mun07}. Detailed descriptions of the spectral and photometric behavior of V838~Mon can be found in: \\cite{mun02}, \\cite{kim02}, \\cite{wis03}, \\cite{cra03,cra05}, \\cite{kip04}, and \\cite{tyl05}.  Although there is great interest in V838~Mon among astrophysicists and many efforts have been made to gain a deeper understanding of the object, its nature remains unclear. A wide range of mechanisms have been proposed to explain the outburst including {\\it (i)} thermonuclear models: a nova-like \\citep{mun02} or a He-shell flash \\citep{law05, munhan05}, and {\\it (ii)} merging events: a stellar merger of stars with masses of 8~M$_{\\sun}$ and 0.3~M$_{\\sun}$ \\citep{soktyl03,soktylconf}, and a giant swallowing planets \\citep{ret03,ret06}. It seems that the most promising is the stellar merger scenario \\citep{tylsok06,soktylconf}.  Since February 2002 a light echo of V838~Mon has been observed \\citep{han02}. Its spectacular evolution, especially well documented by the series of the Hubble Space Telescope (HST) images \\citep{bond03,bondconf}, revealed that V838~Mon is immersed in a diffuse and dusty medium with a very complex spatial distribution. There is no agreement about the origin of this matter. An analysis of the light echo images performed by \\cite{tyl04} shows that the dusty region has a rather complex distribution  and, most probably, is of interstellar origin \\citep[see also][]{cra05,tyl05b}. On the other hand, there are suggestions that the matter might have been ejected by the object itself in the past \\citep{bond03,bondconf}.  A light echo analogue was discovered in the infrared with the Spitzer Space Telescope \\citep{baner06}. The analysis of the infrared emission gives a few tens to a few hundreds of solar masses as an order of magnitude estimate for the mass of the gas associated with the emitting dust. This result is a strong argument supporting the idea of an interstellar origin of the echoing matter.  \\cite{kam07a,kam07b} discovered emission in the $^{12}$CO (2--1) and (3--2) transitions at the position of V838~Mon. The emission was detected inside the large KOSMA-telescope beams with half power widths (HPBWs) of 130$\\arcsec$ and 82$\\arcsec$. The detected lines are very narrow and appear at LSR velocity of 53.3~km~s$^{-1}$. The data analysis performed in \\cite{kam07b} suggest that the emission is extended, but on the basis of the low-resolution observations it was not certain. Moreover, in observations performed in three epochs Kami\\'{n}ski et al. found small but possibly real variations in the intensity of the CO (2--1) line. Due to the unknown contribution to the intensity changes from antenna-pointing errors, the finding is also very uncertain.  \\cite{degu07} observed the region around V838~Mon in the CO (1--$\\,$0) transition and detected a narrow emission line 30$\\arcsec$ north from the star position. This detection together with the observations reported in \\cite{kam07b} show that some form of molecular matter must exist close to the direction of V838~Mon. Its nature has not been sufficiently established, but it is tempting to link the CO emitting gas with the dusty medium illuminated by the eruption of V838~Mon.  In the current paper, we present follow-up observations of the field around V838~Mon in the CO (1--$\\,$0) and (2--1) rotational transitions with a much better angular resolution than the observations reported in \\cite{kam07a,kam07b}. The observational details are provided in Sect.~\\ref{obs}, while the data are described in Sect.~\\ref{res}. In Sect.~\\ref{dis}, we discuss the results, in particular, we investigate the possible origin of the CO emission found in the field around V838~Mon. Sect.~\\ref{sum} contains final conclusions and emphasizes the need for future observations of the echo region in molecular transitions. ",
        "conclusions": "}  We present observations towards 13 positions in the field of the light echo of V838 Mon in the CO (1--$\\,$0) and (2--1) transitions. The measurements reveal an extended molecular region around the star at two distinct radial velocities.  The CO emitting region is elongated in the north-south direction and exhibits a very complex distribution on the plane of the sky. We identify the CO lines as emerging from diffuse interstellar clouds. No molecular emission that can be associated with the star itself was detected.  The possible association of the molecular emission with the light echo material has been investigated. Although the CO emission appears in the field of the light echo, its detailed spatial distribution correlates only weakly with the light echo image. On the other hand, the velocity of the CO emission agrees very well with the velocity of the SiO maser discovered from the direction of V838~Mon, making the collocation of the dust and CO-bearing gas probable.  To more deeply investigate the origin of the molecular emission in the field of the echo, more extended map of the region in different molecular transitions is needed. A fully sampled map of the echoing region would be helpful to draw more conclusive statements about the suggested connection with the light echo material."
    },
    "0801/0801.4779_arXiv.txt": {
        "abstract": "We report on variations in important X-ray emission lines in a series of \\chandra\\ grating spectra of the supermassive colliding wind binary star \\ec, including key phases around the X-ray minimum/periastron passage in 2003.5. The X-rays arise from the collision of the slow, dense wind of \\ec\\ with the fast, low-density wind of an otherwise hidden companion star. The X-ray emission lines provide the only direct measure of the flow dynamics of the companion's wind along the wind-wind collision zone.  We concentrate here on the silicon and sulfur lines, which are the strongest and best resolved lines in the X-ray spectra. Most of the line profiles can be adequately fit with symmetric Gaussians with little significant skewness. Both the silicon and sulfur lines show significant velocity shifts and correlated increases in line widths through the observations. The $\\R=$ forbidden-to-intercombination ratio from the \\ion{Si}{13} and \\ion{S}{15} triplets is near or above the low-density limit in all observations, suggesting that the line-forming region is $>1.6$ stellar radii from the companion star. We show that simple geometrical models cannot simultaneously fit both the observed centroid variations and changes in line width as a function of phase. We show that the observed profiles can be fitted with synthetic profiles with a reasonable model of the emissivity along the wind-wind collision boundary. We use this analysis to help constrain the line formation region as a function of orbital phase, and the orbital geometry. ",
        "introduction": "\\label{sec:Introduction}  The supermassive star \\ec\\ \\citep{davidson97a} is notorious for its extraordinarily large luminosity and its implicitly large mass \\citep[$L>4\\times10^{6}L_{\\odot}$ and $M \\sim 100 M_{\\odot}$,][]{hillier01}, the beautiful bipolar ``Homunculus'' nebula which shrouds it \\citep{gaviola50}, its wild instability (most notably the ``Great Eruption'' of 1843 which created the Homunculus) and its continued broad-band variations \\citep{sterken96,davidson99b}.  Understanding \\ec\\ is important for a wide variety of astrophysical topics regarding the formation and evolution of extremely massive stars, the processes by which such stars lose mass and angular momentum, and the ways in which they interact with their surroundings.  \\ec\\ exhibits variability over a wide range of wavelengths, from radio \\citep{duncan03}, through infrared \\citep{whitelock94,whitelock04,damineli96,damineli97,damineli00,davidson00}, optical \\citep{steiner04}, and ultraviolet \\citep{smith04a} to X-rays \\citep{ishibashi99,corcoran05b}. All these variations have a characteristic cycle of almost exactly 2024 days, which strongly suggests that \\ec\\ is a long period ($P=2024$ day) binary \\citep{damineli96,damineli97}. The observed variability is believed to result (directly or indirectly) from the interaction of the fast wind ($\\vc \\sim 3000~\\kmps$) and ionizing radiation from the companion with the dense, slow wind of the Luminous Blue Variable (LBV) primary ($\\veta \\sim 500~\\kmps$). In this scheme, the X-rays are produced by the collision of the two stars' winds, which causes the companion's fast wind to be shock-heated to tens of MK \\citep{pittard98a,pittard02a}. The high temperature of the shocked wind of the companion explains the hard X-rays ($kT \\ga 4~\\kev$) first directly associated with the star by \\einstein\\ \\citep{seward79}.  Similar hard X-ray emission is seen from WR~140, the ``canonical'' long period eccentric massive colliding wind binary \\citep{pollock05}.  Our understanding of the system has become more sophisticated due in part to dense multiwavelength monitoring near the times of the X-ray minima in 1998 and 2003.5.  These observations showed that, at the same time that the X-ray brightness of the source reaches minimum, the ionization state of the circumstellar medium rapidly decreases \\citep{1995ApJ...441L..73D,nielsen07}, the infrared \\citep{whitelock04} and millimeter-wave \\citep{2005A&A...437..977A} brightness of the source also drops, absorption components in excited \\ion{He}{1} P-Cygni emission lines undergo rapid blue-to-red velocity shifts \\citep{nielsen07}, \\ion{He}{2} 4686-\\angstrom\\ emission \\citep{steiner04,2006ApJ...640..474M} appears, shows a similar blue-to-red centroid shift, then disappears, and the far UV flux from \\ec\\ drops rapidly \\citep{2005ApJ...633L..37I}. In X-rays, the hottest electron temperature stayed the same, but the ionization balance of Fe ions changed remarkably \\citep{hamaguchi07}. In all colliding wind models these changes (which last only about 90 days of the 2024-day cycle) occur near periastron passage, and require a high eccentricity ($e\\sim0.9$).  However, important details regarding the nature of the wind-wind collision are still not well constrained; there is still debate concerning, for example, whether the X-ray minimum occurs near inferior conjunction (when the companion is in front of the LBV primary) or superior conjunction; the magnitude of the companion's wind velocity; and the mass loss rates from either star. These uncertainties limit our understanding of how the companion star affects the system, and, ultimately limit our knowledge of the evolutionary state of the system.  The  detailed analysis of excited \\ion{He}{1} P-Cygni absorption lines in spatially resolved spectra by \\citet{nielsen07}  showed radial velocity variations which mimic the orbital radial velocity variations expected in an eccentric ($e\\approx0.9$) binary system with the semi-major axis pointed towards the observer (longitude of periastron $\\omega\\sim 270^{\\circ}$) and an assumed inclination $i\\approx41^{\\circ}$. These spectral variations suggest that the ionized helium zone in the wind of the cool, massive primary star approaches the observer prior to periastron passage.  They also showed that the velocity amplitude of the \\ion{He}{1} P-Cygni absorption components ($\\sim$140~\\kmps) was larger than expected if the absorption arises in the dense wind of the more massive star. They concluded that the velocity variations are probably strongly influenced by ionization effects due to the interaction of the companion star's photospheric UV radiation with the wind of the cool primary star.  They also suggested that some of the \\ion{He}{1} emission might originate within or near the wind-wind collision and thus could be a diagnostic of that collision.  However the complex influence of the companion's radiation with the primary wind makes interpretation of such diagnostics far from straightforward. (For an alternative explanation of the \\ion{He}{1} observations, in which the lines are assumed to form in the acceleration zone of the secondary, see \\citealp{kashi07}.)  X-ray line profiles provide the most direct probe of the dynamics of the wind of the unseen companion after it is shock-heated in the wind-wind interaction, since these lines originate in the high temperature plasma near the wind-wind shock interface. X-ray lines directly reflect the dynamic properties of this hot shocked gas.  In this paper we present our analysis of the high resolution X-ray grating spectra of \\ec\\ obtained by the High Energy Transmission Grating Spectrometer \\citep[HETGS;][]{markert94} on the \\chandra\\ X-ray Observatory \\citep{weisskopf02} obtained as part of a large observing campaign around the time of the 2003.5 X-ray minimum.  A preliminary  analysis of these data has appeared in \\citet{henley05b}.  In this paper we discuss our analysis of  spectra in the energy range near 2~\\kev\\ obtained by the Medium and High Energy Gratings (MEG and HEG).   This energy range is dominated by line emission from Si and S hydrogen-like and helium-like ions.  These lines form in the cooler regions of the shocked gas farther along the wind-wind collision zone, and thus provide a better measure of the flow dynamics of the shock-heated wind of the companion along the colliding wind interface than the iron lines, which originate near the hottest part of the shock near the stagnation point where flow velocities are low. In this energy range the HETGS first order spectra has sufficient resolution to resolve the component lines of the He-like triplets providing useful density and temperature diagnostics. Unfortunately, potentially crucial line emission from C, N, and O (which could be used to measure abundances of the shocked companion's wind and help constrain the evolutionary state of the companion) are not observable in the central source due to the heavy absorption by the cold gas and dust in the Homunculus.  This paper is organized as follows.  The observations and the data reduction are described in \\S\\ref{sec:Observations}, and the HETGS silicon and sulfur emission lines are discussed in \\S\\ref{sec:Spectra}. In \\S\\ref{sec:GeometricalModel} we apply a simple geometrical model of the wind-wind collision to the variations in line centroids and widths. In \\S\\ref{sec:SyntheticProfiles} we apply synthetic colliding wind line profiles to the observed HETGS silicon and sulfur profiles.  We discuss the results of the emission line analysis in \\S\\ref{sec:Discussion}, and our conclusions are presented in \\S\\ref{sec:Conclusions}. Throughout this paper we quote $1\\sigma$ errors. ",
        "conclusions": "\\label{sec:Conclusions}  We have presented our analysis of resolved silicon and sulfur X-ray emission lines from a series of HETGS observations of \\ec\\ at key orbital phases.  These lines originate in the wind-wind collision zone where the slow, dense wind of \\ec\\ interacts with the fast, low-density wind of a massive companion star.  We have shown that line profile variations around the orbit are not consistent with simple geometrical models of the line forming region.  A more physically realistic model which takes into account the detailed geometry of the contact discontinuity and allows for variations in the emissivity distribution along the shock boundary can produce both the variations in the line centroids and the observed changes in line width. This analysis allows us to probe directly both the temperature distribution along the shock boundary and also the flow of the shocked wind of the companion away from the stagnation point at the apex of the shock.  The \\SiXIV\\ and \\SXVI\\ emission appear to originate from similar regions, which is as expected given the range of temperatures at which they are produced. However, away from periastron they originate closer to the stagnation point than is expected from hydrodynamical simulations. This may be because the wind-wind collision is out of equilibrium, and the line emission is originating from a region which would be too hot if the wind-wind collision were in equilibrium (although it should be noted that we did not find unambiguous evidence of non-equilibrium conditions from the line flux ratios). Just before periastron the size of the \\SiXIV-emitting region is closer to that which is expected from a hydrodynamical simulation, suggesting that the shocked gas has equilibrated at the time of CXO$_{030616}$. We find that the flow speed along the wind-wind collision surface is $\\sim$3$\\times$ the flow speed given by the analytic model of \\citet{canto96}. This larger flow speed approaches a large fraction of the terminal velocity of the companion star's wind far from the stagnation point, which is in fair agreement with detailed numerical hydrodynamic models of the flow, and suggests that the mixing of the cold postshock primary wind into the hot postshock primary wind is relatively unimportant, perhaps due to the presence of magnetic fields.  We can obtain a good fit to the profiles with an inclination and longitude of periastron similar to those which have previously been assumed ($i \\approx 50\\degr$ and $\\omega \\approx 270\\degr$), although this is only true if we include the effects of absorption by the unshocked wind of the companion. Given the simplifying assumptions inherent in the model, the uncertainty of the various input model parameters, and the fact that we did not search for a global best fit, we cannot rule out other possible orbital orientations. Nevertheless, an important result of this analysis is that it shows that colliding wind models can fit the detailed flow dynamics as shown by the variations in X-ray line profiles, without recourse to additional flows in the system.  These results must be considered preliminary since the observed line profile variations need to be confirmed as dependent on orbital phase rather than simply secular changes in the wind. \\chandra\\ HETGS observations are scheduled for Fall 2008 and late 2008, and ideally more will be carried out around the time of the next X-ray minimum, expected in 2009 January. These observations will complement ongoing X-ray monitoring with \\rxte\\ and ground-based optical and radio observations. Meanwhile, the line profile model can be improved by the inclusion of more realistic absorption from the wind of \\ec\\ and from the wind of the companion star, which might give further insight into the mass loss rate of the wind from the companion, and by more detailed numerical models including turbulence to directly determine the dependence of the theoretical line profiles on the orientation and geometry of the colliding wind region."
    },
    "0801/0801.4123_arXiv.txt": {
        "abstract": "{} {To constrain the equation of state of super-nuclear density matter and  probe the interior composition of the X-ray pulsar in SAX J1808.4-3658. In our estimation, we consider both its persistent 2.49 ms X-ray pulsations discovered by Wijnands and van der Klis from using the Rossi X-ray Timing Explorer, which is interpreted to come from an accreting-powered millisecond X-ray pulsar in the low mass X-ray binaries, and the corresponding mass-radius data analyzed of the light curves of SAX J1808.4-3685 during its 1998 and 2005 outbursts by Leahy et al. from assuming a hot spot model where the X-rays are originated from the surface of the neutron star. } {The interior composition of the X-ray pulsar in SAX J1808.4-3658 is obtained from comparing the mass-radius relations, Keplerian motions  and gravitational wave radiation instabilities of the predictions under different theoretical models with the analytically mass-radius district and rotation frequency of observational data.} {We show that SAX J1808.4-3658 may include the exotic matter such as hyperon or quark matter, but we can't distinguish the star with hyperon matter from quark matter. } {} ",
        "introduction": "The properties of nuclear matter under extreme conditions are correlative to the in-depth experiment and theoretical research of the heavy nuclear physics and nuclear astrophysics intimately. Neutron stars provide us a unique spherical environment to study the properties of matter with low temperature and high density. Since the discovery of the first pulsar by Hewish et al. (1968) and then the confirmation of it to be a fast rotational neutron star by Gold and Pacini (1968), the compositions and properties of the interiors of neutron stars have attracted much attention (\\cite{Pandharipande, Glendenning1, Glendenning2, Sahu, Thorsson, Alford1}). If we have made this significant problem clear, it would modify our understanding on the structure of the matter, strengthen the cognition about the formation of the universe and in all contribute most to the research of astrophysical physics, nuclear physics and particle physics.  However, in fact due to the uncertain nuclear physics, the equation of state for neutron star at super-nuclear density is still an indeterminacy, on which the maximum mass, radius and rotation frequency depend strongly. Many investigators spontaneously expect to probe the matter under these conditions and constrain or rule out some equations of state through astrophysical observations (\\cite{Glendenning3, Lattimer, Lackey, Lavagetto, Klahn, Lattimer1, Pan}). Although researchers universally impose the inferred masses and radii or rotation frequencies for pulsars on the equations of state, their common treatments cannot uniquely examine and distinguish all the classes of  compositions inside neutron stars, i.e. it could rule out neither traditional neutron stars nor hyperon stars or quark stars determinately. This problem should be carefully considered in other effective ways for the further estimation. ",
        "conclusions": "The composition  of matter in the core of neutron stars has attracted much attention owing to its important significance. Up to now, SAX J1808.4-3685 is one of seven known accreting ms X-ray pulsar. Its persistent 2.49 ms X-ray pulsations was discovered by Wijnands and van der Klis and interpreted as due to the ms rotation of the center neutron star. And Leahy et al. have analyzed its corresponding mass-radius relation. In our estimation, we try to probe the inner components of it in our own way by comparing the mass-radius relations and genuine rotation frequencies under different theoretical models with the observational data. We finally come to a conclusion that the pulsar in SAX J1808.4-3685 is a star containing exotic matter, which just bases on these two observational properties, but we can't distinguish it with hyperon matter from quark matter at the present time, which must depend on more observational information about it, such as the thermal emission data.   \\begin{figure} \\centering \\includegraphics[width=0.51\\textwidth]{fig2.eps} \\caption{Limit rotation frequencies for various equations of state of compact stars. These dashed lines represent the r-mode instability limits, and dotted lines the Keplerian limits. The corresponding thick solid lines are the genuine upper frequencies. The light yellow district refers to the mass prediction of the X-ray pulsar in SAX J1808.4-3658 at 99.7\\% confidence lever analyzed by Leahy et al. (2007), and the green dash-dot-dot horizontal line the corresponding 2.49 ms rotation frequency discovered by Wijnands and van der Klis (1998). } \\label{mf} \\end{figure}"
    },
    "0801/0801.4912_arXiv.txt": {
        "abstract": "We analyze integral-field ionized gas and stellar line-of-sight kinematics in the context of determining the stellar velocity ellipsoid for spiral galaxies observed by the Disk-Mass Survey. Our new methodology enables us to measure, for the first time, a radial gradient in the ellipsoid ratio $\\sigma_z/\\sigma_R$. Random errors in this decomposition are 15\\% at two disk scale-lengths. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.4315_arXiv.txt": {
        "abstract": "X-ray emission due to inverse-Compton scattering of microwave background photons by electrons in the lobes of powerful radio galaxies has now been seen in a large number of objects.  Combining an inverse-Compton model for the lobe X-ray emission with information obtained from radio synchrotron emission provides a method of constraining the electron population and magnetic field energy density, which cannot be accomplished using the radio data alone. Using six frequencies of new and archival radio data and new {\\it XMM-Newton} observations of the Fanaroff \\& Riley class II radio galaxy 3C\\,353, we show that inverse-Compton emission is detected in the radio lobes of this source at a level consistent with what is seen in other objects.  We argue that variations in the X-ray/radio ratio in the brighter eastern lobe require positionally varying magnetic field strength.  We also examine the X-ray nucleus and the cluster, Zw\\,1819.1-0108, spatially and spectrally. ",
        "introduction": "Extended X-ray emission from the lobes of powerful, FR\\,II \\citep{fr74jlg} radio galaxies and quasars is thought to be produced by inverse-Compton scattering of the cosmic microwave background (CMB) \\citep{feigelson95jlg,tashiro98jlg,tashiro01jlg,isobe02jlg,hardcastle02jlg} and of infrared photons from the core \\citep*{brunetti97jlg}. Although the synchrotron self-Compton emission from the lobes can also be modelled, it is negligible compared to the inverse-Compton emission from other populations except in the very smallest lobes \\citep{hardcastle02jlg}.  Recent studies of the integrated X-ray properties of radio lobes have been carried out for large samples \\citep{croston05jlg, kataoka05jlg}, showing that the CMB is the dominant photon population in most cases.  Combining an inverse-Compton/CMB model for the lobe X-ray emission with information obtained from radio synchrotron emission provides a method of constraining the electron population and magnetic field energy density, which cannot be accomplished using radio data alone. Unlike the jets, the lobe material is not moving relativistically so there are no beaming effects to consider during this analysis \\citep[e.g.][]{mackay73jlg}.  Understanding the contributions of the electron densities of the synchrotron and inverse-Compton electron populations and the magnetic field allows us to estimate directly the total energy in the radio source, and thus the amount of energy that can be transferred to the environment, without assuming equipartition, and also allows us to investigate the distribution of internal energy within the source.  \\citet{hardcastle05jlg} recently carried out a spatially resolved X-ray inverse-Compton study of the bright X-ray and radio source Pictor A.  They used the variation of the X-ray/radio ratio across the lobes to investigate the variation in electron density and magnetic field strength throughout the source.  They found that variations in either component alone could not explain the observed X-ray and radio properties of the lobe.  Similar spatially resolved X-ray studies of other bright FR\\,II radio galaxies are needed to follow up these results.  Here we report on observations of the radio galaxy 3C\\,353.  3C\\,353 is a FR\\,II radio galaxy associated with the cluster Zw\\,1718.1-0108.  Although it is one of the brightest extragalactic sources in the sky at low frequencies, it is relatively poorly studied, presumably due to its low declination and low Galactic latitude.  \\citet{swain96jlg} observed 3C\\,353 at four radio frequencies with the NRAO Very Large Array (VLA) and published a study of the polarisation variation across the jets, which favoured a spine-sheath model \\citep*{swain98jlg}. Their VLA observations revealed filamentary structure within the lobes making 3C\\,353 an excellent target for a spatially resolved study of electron distribution and magnetic fields.  X-ray emission from the cluster associated with 3C\\,353 was first detected by \\citet{iwasawa00jlg} in {\\it ASCA} and {\\it ROSAT} images.  The radio source resides in a giant elliptical galaxy on the edge of the cluster, which Iwasawa et al. found to be bright in the X-ray with a luminous point source coinciding with the radio galaxy's core.  Iwasawa et al. determined global cluster temperatures and used optical observations to identify the cluster's members.  They confirmed a redshift of $z=0.0304$ for 3C\\,353, a redshift of $z=0.028$ for three bright member galaxies, and used the velocity distribution to confirm 3C\\,353 as a member of the same system.  In this paper we use new {\\it XMM-Newton} and radio observations to investigate the nature of the electron distribution and the magnetic fields in the lobes and hotspots of 3C\\,353.  We also determine the cluster temperature, density and pressure profiles and examine the cluster's interaction with the lobes of 3C\\,353. Throughout the paper we use a cosmology in which $H_{0} = 70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{m} = 0.3$, and $\\Omega_{\\Lambda} = 0.7$.  The angular scale is 1\\,arcsec = 0.61\\,kpc.  We define the spectral index in the sense that $S_{\\nu} \\propto \\nu^{-\\alpha}$.  The photon index $\\Gamma = 1 + \\alpha$. ",
        "conclusions": "Our results can be summarized as follows:  \\begin{itemize}  \\item By fitting an inverse-Compton model to the lobes, we found the East lobe to be electron dominated and the West lobe to be consistent (within the large errors) with equipartition.  \\item We determined that a variation in the electron spectron cannot account for the varying X-ray and radio emission alone, but that a change in the magnetic field strength across the lobes is required.  \\item We have obtained a good X-ray spectrum of the nucleus of 3C\\,353.  Both the X-ray and optical properties of this source are ambiguous but it appears to lie in a region of parameter space in between those normally occupied by narrow-line and low-excitation radio galaxies.  \\item We have detected an X-ray counterpart for the East hotspot offset by $4.0 \\pm 0.5$\\,kpc.  \\item The northern and southern sub-clusters were found to be isothermal with a temperature difference of $\\sim 1$\\,keV supporting a model in which they are two originally separate components undergoing a merger with no evidence for a violent interaction.  \\end{itemize}"
    },
    "0801/0801.2657_arXiv.txt": {
        "abstract": "We investigate a wide range of possible evolutionary histories for the recently discovered Bootes dwarf spheroidal galaxy, a Milky Way satellite.  By means of $N$-body simulations we follow the evolution of possible progenitor galaxies of Bootes for a variety of orbits in the gravitational potential of the Milky Way.  The progenitors considered cover the range from dark-matter-free star clusters to massive, dark-matter dominated outcomes of cosmological simulations. For each type of progenitor and orbit we compare the observable properties of the remnant after 10~Gyr with those of Bootes observed today.  Our study suggests that the progenitor of Bootes must have been, and remains now, dark matter dominated.  In general our models are unable to reproduce the observed high velocity dispersion in Bootes without dark matter.  Our models do not support time-dependent tidal effects as a mechanism able to inflate significantly the internal velocity dispersion.  As none of our initially spherical models is able to reproduce the elongation of Bootes, our results suggest that the progenitor of Bootes may have had some intrinsic flattening.  Although the focus of the present paper is the Bootes dwarf spheroidal, these models may be of general relevance to understanding the structure, stability and dark matter content of all dwarf spheroidal galaxies. ",
        "introduction": "\\label{sec:intro}  The dwarf satellite galaxies of the Milky Way have attracted considerable interest in recent years due to their high apparent mass to light ratios, which suggest the presence of considerable quantities of dark matter~\\citep[DM:][]{aaronson83,mateo98,kleyna01}.  Given the apparent absence of dynamically significant DM in globular star clusters, the dwarf spheroidal galaxies (dSphs) are particularly valuable for DM studies as they are the smallest stellar systems known to contain DM~\\citep{gilmore07}.  $N$-body simulations of galaxy formation in a $\\Lambda$CDM cosmology predict that a galaxy like the Milky Way should be surrounded by several hundred low-mass satellite halos~\\citep[e.g.][]{moore99}.  The small number of known dSphs in the vicinity of the Milky Way has often been cited as a problem for $\\Lambda$CDM and many attempts have been presented in the literature to explain why the numbers of satellites which formed stars might be much lower than the total number of substructures.  The past two years have witnessed the discovery of a plethora of new low-luminosity Milky Way satellites~\\citep{willman05,be06,zu06,be07,walsh07}.  These objects probe a previously unexplored regime for galaxies, extending the faint end of the galaxy luminosity function by almost four magnitudes.  As \\citet{gilmore07} discuss, these new objects also extend to fainter magnitudes the apparent bi-modality in the size distribution of low-luminosity stellar systems, with dSphs exhibiting core radii which are always more than a factor of four larger than the half-light radii of the most extended star clusters.  Given that all the dSphs brighter than $M_{\\rm V}=-8$ show evidence of DM, while star clusters appear to be purely stellar systems, understanding the physical origin of this size difference may have implications for our knowledge of DM.  It is therefore of particular importance to determine whether or not the newly discovered satellites display evidence of DM.  The goal of this paper is to investigate whether the current observational data for one of these newly identified satellites can be used to constrain its DM content.  \\citet{be06} reported the discovery of a faint dwarf galaxy in the constellation of Bootes.  The object was discovered during a systematic search for over-densities of stars in the magnitude range $16 \\leq r \\leq 22$ in the Sloan Digital Sky Survey~\\citep[SDSS][]{yo00}.  The morphology of the stellar iso-density contours based on the SDSS data suggested that the system was quite irregular with some hints of the presence of internal substructure, possibly indicating that the object was in the process of tidal disruption. The aim of this paper is to test this hypothesis and constrain the properties of potential progenitors of this system.  The investigation of the evolution of the dSph galaxies around the Milky Way by means of $N$-body simulations has a long history.  Many authors \\citep[e.g.][]{jo99,jo02,ma02,read06} try to model the initial cosmological halos with or without the luminous component inside, follow their evolution and compare the final results with the population of dSph galaxies of the Milky Way.  Mostly these studies focused on the more luminous dwarfs, which were the only known satellites prior to the recent discoveries.  In this paper we follow a different approach.  We use $N$-body simulations to investigate the evolution of possible progenitors of Bootes (henceforth Boo).  Our goal is to understand the present-day properties of Boo and to determine whether these can be used to constrain the DM content of this satellite and of its progenitor.  In \\S~\\ref{sec:obsprop} we describe the observed properties of Boo.  \\S~\\ref{sec:setup} summarises the initial conditions for our models. \\S~\\ref{sec:results} presents the results of our simulations for a range of possible progenitors and Galactocentric orbits.  Finally, in \\S~\\ref{sec:conclus} we assess the relative merits of the various models and describe the follow-up observations which are required to distinguish between them. ",
        "conclusions": "\\label{sec:conclus}  In this paper we have investigated a wide range of possible progenitors and evolutionary histories for the recently discovered Bootes dwarf spheroidal satellite galaxy of the Milky Way.  As we discussed in \\S~\\ref{sec:inimod}, although the outer iso-density contours of the stellar distribution are elongated, an observation sometimes interpreted as suggestive of tidal disturbance, this interpretation is not consistent with the observed velocity dispersion of $\\approx 7$~km\\,s$^{-1}$.  Our initial attempts to reproduce both the distorted outer morphology and high velocity dispersion using models which were close to, or slightly beyond, the point of complete disruption were unsuccessful.  We therefore focused on bound models either with or without the presence of significant quantities of dark matter.  Our model (A) is an extended 'star cluster' model and reproduces the surface brightness profile and distorted outer morphology of Boo reasonably well.  As expected for a purely stellar system from our argument in \\S~\\ref{sec:inimod}, however, the final velocity dispersion is significantly smaller than that observed in Boo.  It is therefore very unlikely that a purely stellar model could be the origin of Boo.  Our dark matter-dominated models models (cases (B) - (D)) can reproduce both the surface brightness profile and the velocity dispersion, but not the morphology of the outer stellar distribution. This morphological limitation is because all our progenitors are initially spherical, and their dark matter halos protect them from significant tidal disturbance.  Even if the orbit of Boo has a very small perigalacticon (case (D)), the remnant does not exhibit tidal tails on scales as small as those observed in Boo.  For these models therefore, if the elongation of the outer contours of Boo is real and not an artifact of the sparse photometric data, then it must result from intrinsic elongation of the progenitor system.  As stated in \\S~\\ref{sec:inimod} the putative S-shape of the contours in these cases do not represent the on-set of tidal tails but might be caused by tidal torques acting on an intrinsically flattened system which is at least partly rotationally supported but remains deeply embedded in its DM halo.  Given the very low stellar mass of a faint satellite like Boo, a dwarf disc galaxy progenitor (such as was used by \\citet{ma02} to model the Carina dSph) would either have to be orders of magnitude lower in baryonic mass than typically observed stellar discs, or would have to lose almost all of its initial mass (both dark and baryonic). The latter scenario would either require an extremely low star formation efficiency in the original disk (i.e.\\ most of the baryonic component is still in the form of gas, which gets easily stripped by ram pressure stripping and the cosmic ultraviolet background at high redshift as suggested by \\citet{may07}), or very strong tidal disturbance. Again we emphasise that to reproduce the high velocity dispersion of Boo, we require a dark matter mass in the inner regions which would preclude this level of tidally-induced mass loss from the central parts - in our study it was not possible to reproduce the data of Boo with a disrupted object of any kind.   On the basis of our study, we conclude that it is unlikely that the progenitor of Boo contained no dark matter.  In fact, we found no viable models in which dark matter was not always dominant.  External tidal effects on the evolution of the progenitor of Boo do not significantly affect the observed stellar velocity dispersion in the inner regions, or the structure of the galaxy.  Our most likely models have final DM masses in the range $1-7 \\times 10^{7}$M$_\\odot$.  By changing the possible orbit of Boo we found that the closer the orbit gets to the Galactic centre (i.e.\\ smaller perigalacticon and higher eccentricity) the larger is the mass-to-light ratio exhibited in the undisturbed central part of the final model.  In order to understand properly the details of the specific Bootes galaxy progenitor, however, further data are required.  First, deeper photometric data extending to larger radii around the satellite are essential in order to establish the full extent of the remnant and whether there is any evidence that the elongation is tidal in origin (e.g. whether the elongation links to larger scale tidal arms). Secondly, a larger kinematic data set is required to study the variation of the velocity dispersion and mean velocity with position in Boo.  This would yield further information about the mass distribution within the remnant.  Both these studies are currently underway. \\\\   \\noindent {\\bf Acknowledgements}: MF, VB and DBZ acknowledge funding by STFC.  MIW acknowledges support from a Royal Society University Research Fellowship.  We thank Rainer Spurzem and Walter Dehnen for useful comments."
    },
    "0801/0801.4906_arXiv.txt": {
        "abstract": "{We have constructed a grid of photoionization models of spherical, elliptical and bipolar planetary nebulae. Assuming different velocity fields, we have computed line profiles corresponding to different orientations, slit sizes and positions. The atlas is meant both for didactic purposes and for the interpretation of data on real nebulae. As an application, we have shown that line profiles are often degenerate, and that recovering the geometry and velocity field from observations requires lines from ions with different masses and different ionization potentials. We have also shown that the empirical way to measure mass-weighted expansion velocities from observed line widths is reasonably accurate if considering the HWHM. For distant nebulae, entirely covered by the slit, the unknown geometry and orientation do not alter the measured velocities statistically. The atlas is freely accessible from internet. The Cloudy\\_3D suite and the associated VISNEB tool are available on request. } \\addkeyword{methods: numerical} \\addkeyword{line: profiles} \\addkeyword{turbulence} \\addkeyword{planetary nebulae}  \\begin{document} ",
        "introduction": "\\label{sec:Intro} As is known, nebular line profiles bear information on the motions of the emitting gas. The study of line profiles thus allows one to measure the expansion velocity of nebulae and even to determine their internal velocity field. The velocity field is an important key for understanding the dynamics and the genesis of planetary nebulae. The expansion velocity allows one to derive the expansion age, and from this, such parameters as the masses of the central stars \\citep{1997A&A...318..256G}. However, the determination of these characteristics is a very difficult task in general, since  only the velocity field perpendicular to the plane of the sky can be observed, through the study of line profiles. Even if one could, for a series of ions with different ionization potentials, obtain line profiles in every pixel of a high-resolution image of a PN, one would still need to make assumptions about the projection of the velocity vectors in the plane of the sky, in order to reach a full description of the velocity field. Fortunately, most PNe exhibit some kind of overall regularity in their morphology, and some degree of symmetry (apparently circular, ellipsoidal, bipolar...). So, if an image of the nebula is available, it is possible, with the help of a few assumptions,  to determine the entire nebular geometry and velocity field. Most of the time, however, only limited data are available: line profiles obtained through one slit,  or even, if the object has very small angular dimensions,  line profiles corresponding to the emission of the entire object. If the object is faint, line profiles can be obtained only for the brightest lines, the extreme case being when they are available only for \\oiii/. Finally, morphological information on the object may be completely lacking, as is the case for PNe in distant galaxies. In such cases, what is the value of the information that one can derive from line profiles?  In practice, the analysis of line profiles of planetary nebulae has followed several paths: i) A simple determination of the expansion velocities from observed line widths or from line splitting \\citep[most of the references quoted in][]{1989A&AS...78..301W}; ii) Fitting of line profiles with spherical models \\citep{GAZ03}; iii) Modeling emission line profiles and Position-Velocity diagrams \\citep{1984A&AS...58..273S,1968ApJ...153...49W} and more recently \\citet{2005AJ....130.2303M} and \\citet{2006RMxAA..42...99S,Stef_Canaries07} which  their powerful tool SHAPE; iv) Fitting of line profiles and other nebular characteristics with a 3D photoionization model \\citep{2000ApJ...533..931M,2000ApJ...537..853M,1984A&AS...58..273S}; v) Tomography \\citep[][and references therein]{2006A&A...451..937S}.  The question of the robustness of the results is difficult and is almost never addressed in such studies. The atlas of line profiles that we present in this paper is meant as a tool to visualize the variety of possible line profiles, depending on the morphology of the nebula, the velocity field and the observing conditions. The atlas is to be installed in a virtual observatory environment, that will collect in a large data base the models and line profiles computed by any user with the tools we have developed \\citep{2006IAUS..234..467M}.  The description of the atlas is given in Sec.~\\ref{sec:Descri}. In the second part of this paper, we use the atlas, in its present stat, for two applications: a search for models reproducing a given profile (Sec.~\\ref{sec:appli1}) and an estimation of the robustness of the determination of expansion velocities based on profiles (Sec.~\\ref{sec:appli2}). ",
        "conclusions": "\\label{sec:conclusion}  Using the pseudo-3D photoionization code Cloudy\\_3D, we have constructed a grid of photoionization models of planetary nebulae with different morphologies, including elliptical and bipolar geometries. We have then assumed different expansion laws, and computed line profiles corresponding to different orientation of the models considered as virtual nebulae, and to different slit sizes and positions. The resulting atlas is available at \\url{http://132.248.1.102/Atlas_profiles/}. In its present version, it is already useful to visualize the variety of profiles that can be obtained according to the physical properties of the nebulae and to the observing conditions, and to help in the interpretation of observational data.  We have shown that in many circumstances, line profiles are degenerate, and the recovery of the true geometry and velocity field from observations requires lines from ions with different masses and different ionization potentials.  We have also shown that the empirical way to measure mass-weighted expansion velocities from the HWHM of observed line profiles is accurate within about 20\\% in the  range of geometries and velocity fields considered in the present version of the atlas. In the case of observations of distant nebulae which are entirely covered by the slit, the unknown geometry and orientation of the nebula do not alter the measured velocities, when considered statistically.  We plan to integrate our atlas in a virtual observatory environment. For the time being, the Cloudy\\_3D suite and the associated VISNEB tool are available on request from C.M. and may be used as a help to interpret observed expansion profiles.  One application of the atlas will be to compare the synthetic line profiles and PV-diagrams with observed data such as the SPM Kinematic Catalogue of Planetary Nebulae \\citep{2006RMxAC..26..161L}."
    },
    "0801/0801.3698_arXiv.txt": {
        "abstract": "We present numerical simulations of primordial supernovae in cosmological minihalos at $z \\sim$ 20.  We consider Type II supernovae, hypernovae, and pair instability supernovae (PISN) in halos from 6.9 $\\times$ 10$^5$ - 1.2 $\\times$ 10$^7$ $\\Ms$, those in which Population III stars are expected to form via H$_2$ cooling.  Our simulations are novel in that they are the first to follow the evolution of the blast from a free expansion on spatial scales of 10$^{-4}$ pc until its approach to pressure equilibrium in the relic \\ion{H}{2} region of the progenitor, $\\sim$ 1000 pc.  Our models include nine-species primordial chemistry together with all atomic H and He cooling processes, inverse Compton cooling and free-free emission.  The supernovae evolve along two evolutionary paths according to whether they explode in \\ion{H}{2} regions or neutral halos. Those in \\ion{H}{2} regions first expand adiabatically and then radiate strongly upon collision with baryons ejected from the halo during its photoevaporation by the progenitor.  Explosions in neutral halos promptly emit most of their kinetic energy as x-rays, but retain enough momentum to seriously disrupt the halo.  We find that the least energetic of the supernovae are capable of destroying halos $\\lesssim$ 10$^7$ $\\Ms$, while a single PISN can destroy even more massive halos.  Blasts in \\ion{H}{2} regions disperse heavy elements into the IGM, but neutral halos confine the explosion and its metals.  Primordial supernova remnants develop dynamical instabilities at early times capable of enriching up to 10$^6$ $\\Ms$ of baryons with metals to levels $\\gtrsim$ 0.01 $\\Zs$, well above that required for low-mass star formation.  In \\ion{H}{2} regions, a prompt second generation of stars may form in the remnant at radii of 100 - 200 pc in the halo. Explosions confined by large halos instead recollapse, with infall rates in excess of 10$^{-2}$ $\\Ms$ yr$^{-1}$ that heavily contaminate their interior.  This fallback may either fuel massive black hole growth at very high redshifts or create the first globular cluster with a radius of 10 - 20 pc at the center of the halo.  Our findings allow the possibility that the first primitive galaxies formed sooner, with greater numbers of stars and distinct chemical abundance patterns, than in current models. ",
        "introduction": "Numerical studies indicate that primordial stars form in the first pregalactic objects to reach masses of $\\sim$ 5 $\\times$ 10$^{5}$ $M_{\\odot}$ at $z \\sim$ 20 - 50 \\citep{bcl99,abn00,abn02,gao07}. The models suggest that Population III stars form in isolation (one per halo) and that they are likely very massive, from 100 to 500 $M_{\\odot}$.  More recent work surveying a larger sample of halos \\citep{oshea07b} extends the lower mass limit of primordial stars down to 15 $M_{\\odot}$.  Very massive stars easily photoionize their halos, typically in a few hundred kyr with final \\ion{H}{2} region radii of 2500 - 5000 pc while \\ion{H}{2} regions of less massive stars are confined to very small radii ($<$ 10$^{-3}$ pc) by the high central densities deep in the halo \\citep{wan04,ket04,abs06,awb07,jet07,yet07,wa07a,wa07b}. In the latter case, high recombination rates and pressure from overlying matter due to the depth of the halo dark matter potential well halts the expansion of the heated bubble \\citep{ftb90}.  Unlike massive stars in the Galaxy today, little mass loss occurs in Population III stars because there are no line-driven winds from their pristine atmospheres \\citep{bhw01,kud00,vkl01}, although mixing in rapidly rotating stars may cause some mass loss later in their lives \\citep{met07}.  Consequently, winds do not clear gas from cosmological halos in the first generation of stars as they do from molecular cloud cores today.  However, if the ionization front (I-front) of the star breaks out of the halo it will sweep more than half of the baryons interior to the virial radius into a thick shell $\\sim$ 100 pc in radius by the end of the life of the star \\citep{wan04,ket04}.  The diffuse gas enclosed by the shell is uniform and at densities far lower than in the original halo, 0.1 - 1 cm$^{-3}$ instead of more than 10$^{10}$ cm$^{-3}$.  The ultimate fate of a Population III star depends on its mass at the end of its main sequence lifetime (MSL).  Those lying between 10 $M_{\\odot}$ and 40 $M_{\\odot}$ die in Type II supernova explosions, ejecting roughly 20\\% of their mass into the intergalactic medium (IGM) as metals and leaving behind a compact remnant.  Stars between 140 $M_{\\odot}$ and 260 $M_{\\odot}$ explode in extremely energetic pair instability supernovae (PISN) hundreds of times more powerful \\citep{hw02}. As much as 50\\% of the mass of the star is dispersed as heavy elements into the IGM in such events, with no compact remnant.  However, these predictions are based on one-dimensional non-rotating stellar evolution models that make a range of assumptions (e.~g. mixing parameters) and should be considered qualitative, rather than quantitative, in nature.  How far elements are expelled from the star depends strongly on the state of the halo prior to the explosion.  Supernova (SN) remnants in halos completely ionized by the star typically expand to half the diameter of the H II region.  The shock front comes into pressure equilibrium with the hot recombining gas more quickly than if the blast propagates in a neutral medium of comparable density \\citep{get07}.  Supernovae in neutral halos, which result when low-mass Population III stars fail to ionize them, evolve quite differently.  Because no winds evacuate gas from the center of the halo, the blast encounters circumstellar densities several orders of magnitude greater than in OB associations in the Galaxy, with much higher energy losses at very early times.  Since the blast energy is much greater than the binding energy of the gas in the halo, it is often supposed that PISN easily destroyed halos.  Here, destruction of the halo means complete expulsion of its gas, not dispersal of the dark matter itself, which remains mostly undisturbed.  However, as \\citet{ky05} point out, this argument ignores radiative losses from the shock in the large densities deep in the halo. They find that bremsstrahlung and inverse Compton cooling (IC) radiate away the explosion's energy before it even sweeps up its own mass in ambient gas, quenching the blast and preventing any dissemination of metals into the halo.  Explosions with hundreds of times the binding energy may be necessary to disperse more massive halos under these conditions.  However, a difficulty with this and other studies of primordial supernovae in cosmological halos is how the explosion itself is initialized.  This is typically done by depositing the energy of the blast as thermal energy on the mesh and allowing pressure gradients to launch shocked flows into their surroundings \\citep{byh03,ky05,get07}.  This is partly justified for explosions in \\ion{H}{2} regions whose central densities are quite low and in which early radiative losses from the remnant are minimal.  Too much energy loss from the early remnant would be ignored by this approach for SN in neutral halos, which cool strongly from very early times due to high ambient densities.  In fact, explosions initialized by thermal pulses in neutral halos cool before any of their energy can be converted into motion \\citep{ky05}.  When the energy of the blast is deposited as heat to central mesh zones, their temperatures skyrocket to tens of billions of degrees K, driving extreme bremsstrahlung and IC losses before pressure gradients can accelerate any flows.  These initial conditions do not properly represent the free expansion that actually erupts through the stellar envelope and whose momentum cannot be radiated away.  The early stages of the explosion in a nearly undisturbed halo are much better modeled by a cold free expansion in which the energy of the blast is kinetic rather than thermal.  Furthermore, supernovae modeled in cosmological density fields with smoothed particle hydrodynamics methods \\citep{byh03,get07} also lack the numerical resolution to capture the early hydrodynamic behavior of the blast wave.  Dynamical instabilities might manifest that lead to early mixing of heavy elements in the ejecta with ambient halo material on much earlier time scales than in current simulations.  How easily neutral halos are destroyed by primordial supernovae and how metals mix with the early IGM thus remains unclear.  One key issue is how much gas is displaced in the halo by the momentum of the ejecta after its energy is lost to cooling processes in the expanding flow at early times. In \\ion{H}{2} regions the shock later strongly interacts with the dense shell of ionized gas expanding within the \\ion{H}{2} region.  How well do metals in the remnant mix with the dense shell, and if dynamical instabilities arise that fragment the shock into enriched clumps, are they liable to collapse into new stars?  Numerical simulations excluding these phenomenon indicate that metals do not significantly mix with the IGM until gas falls back into the dark matter potential after about a Hubble time \\citep{get07}.  Capturing the true evolution of the blast presents numerical challenges due to the spatial scales at play.  To accurately represent the initial ejecta as a free expansion, its hydrodynamical profile cannot extend beyond radii that enclose more surrounding gas than ejecta mass.  In the high central densities of neutral halos such radii are 5 $\\times$ 10$^{-4}$ - 1 $\\times$ 10$^{-3}$ pc, but the blast must then be evolved to tens or hundreds of parsecs to evaluate its true impact on the halo.  Furthermore, multiple sites of intense luminosity will erupt in the flow over all these scales from the earliest times, so the chemistry of the flow must be solved consistently with hydrodynamics to determine where gas is collisionally ionized and therefore strongly radiating.  The attendant time scales can restrict integration times to much lower values than Courant times.  To address these questions we have performed one-dimensional models of Type II supernovae, hypernovae, and PISN in neutral and ionized cosmological halos at redshift 20.  Our initial conditions are spherically-averaged halos derived from cosmological initial conditions with the Enzo adaptive mesh refinement (AMR) code \\citep{enzom} that are photoionized by the star in the ZEUS-MP reactive flow radiative transfer code \\citep{wn06}.  The aim of our survey is to determine the explosion energies required to destroy cosmological halos that are thought to form primordial stars by H$_2$ cooling.  We also evaluate their observational signatures and how their metals will mix with the ambient medium as a prelude to three-dimensional AMR calculations now in progress.  Our radiation hydrodynamics scheme and initial profiles for the free expansion are presented in $\\S$ 2.  We also describe the expanding grids used in this study and the time-dependent boundary conditions that present the ambient medium to the expanding flow.  \\ion{H}{2} region profiles for the explosions and the grid of supernova and halo models selected for this survey are outlined in $\\S$ 3.  The evolution of blasts in both neutral halos and \\ion{H}{2} regions are examined together with dominant energy loss mechanisms for each case in $\\S$ 4.  We discuss observational measures for both kinds of explosion in $\\S$ 5, chemical enrichment of baryons in the halos in $\\S$ 6 and conclude in $\\S$ 7. ",
        "conclusions": "Primordial supernovae can be more damaging to cosmological minihalos than previous studies suggest when they are properly initialized with kinetic rather than thermal energy.  This is primarily because the momentum of the ejecta cannot be radiated away, even when radiative cooling is highly efficient, as in partially ionized or neutral clouds.  Even blasts that are trapped by massive halos severely disrupt them, mixing their interiors with metals and creating massive fallback onto any black hole that remains.  Population III stars of lower mass that are capable of explosions (15 - 40 $\\Ms$) easily destroy halos $\\lesssim$ 10$^7$ $\\Ms$ but not more massive ones, in part because they cannot preionize them.  On the other hand, primordial stars from 140 - 260 $\\Ms$ will destroy halos even more massive than 10$^{7}$ $\\Ms$ in pair instability supernova blasts.  To a good approximation, any supernova will destroy any halo less massive than 10$^{7}$ $\\Ms$ when photoevaporation of the halo prior to the blast is taken into account.  Primordial supernovae evolve in one of two ways.  Blasts in \\ion{H}{2} regions remain free expansions out to several pc and are decelerated first by adiabatic losses and then later by radiative losses.  In addition to the light curve of the radioactive decay of their ejecta, these remnants radiate strongly upon collision with the dense shell of the \\ion{H}{2} region, usually by 10$^4$ - 10$^5$ yr.  While most of their luminosity is in atomic H and He lines, particularly powerful blasts in \\ion{H}{2} regions can channel up to 20\\% of their energy into the CMB.  Explosions in ionized halos evacuate gas from their virial radii to densities below 10$^{-4}$ cm$^{3}$.  Our models are in qualitative agreement with earlier one-dimensional Lagrangian and three-dimensional smoothed particle hydrodynamics (SPH) calculations of Type II and pair instability SN in primordial \\ion{H}{2} regions.  Our results are also consistent with early one-dimensional calculations of supernova blasts in uniform diffuse media \\citep{chev74,ciof88}:  they too manifest reverse shocks that separate from forward shocks at relatively small radii and exhibit multiple shock reflections throughout the interior.  On the other hand, blasts in nearly neutral halos promptly radiate most of their kinetic energy as x-rays but retain sufficient momentum for the ejecta to reach radii of 10 - 20 pc.  These contained explosions leave little observational footprint but radically alter the evolution of gas within the halo, first by contamination with heavy elements and then by successive cycles of fallback and rebound.  This is a departure from earlier supernovae models in neutral halos initialized with thermal rather than kinetic energy that is radiated away before launching any flow out into the halo.  Again, in agreement with earlier work by \\citet{chev74} and \\citet{ciof88}, we find that neither explosion pathway is well represented by the canonical phases of idealized remnants in uniform media (Sedov-Taylor, PDS, or MCS).  In all cases inclusion of the dark matter potential is crucial; models would otherwise fail to capture fallback in trapped explosions or erroneously predict expansion of the remnant until it achieves pressure equilibrium with the IGM.  Primordial supernovae occurred in neutral halos only if the halo was $\\gtrsim$ 10$^7 \\Ms$ and the progenitor was not very massive.  Given that stars normally form by H$_2$ cooling well before the halo reaches this mass, is it likely that such blasts ever happened?  The answer is yes if one considers halos at slightly later redshifts in rising Lyman-Werner backgrounds, whose effect is to postpone star formation in halos until they reach larger masses and become self-shielded to photodissociating photons \\citep{mac01,oshea07a,wa07c}.  Numerical surveys indicate that when stars form in more massive halos whose virial temperatures approach the threshold for atomic line cooling, they still do so by molecular hydrogen cooling and are subject to the range in masses examined in this paper \\citep{oshea07b}.  It is probable that explosions and fallback in more massive halos were more common at somewhat lower redshifts (15 - 20) than those of the very first stars.  We also point out that there are a greater range of spin parameters for star-forming halos than are sampled in current studies.  Halos with higher angular momenta exhibit markedly lower central infall rates \\citep{oet05} that lead to much less massive stars.  Finally, since the processes that halt accretion onto primordial stars remains poorly understood, low-mass supernovae in relatively massive halos cannot be ruled out.  The recent discovery of the first potential pair instability supernova ever observed \\citep{ns07} lends hope that this sort of explosion may have occurred in the high redshift universe.  Population III PISN might be seen at moderate redshift ($z \\sim$ 6) with current technology \\citep{scan05}, or out to significantly higher redshift with the \\textit{James Webb Space Telescope} \\citep{wl05}.  Scannapieco et al. use the KEPLER code to create mock supernova light curves and demonstrate that the typical luminosities of pair instability supernovae are not much brighter than standard Type Ia supernovae.  However, they remain brighter for much longer periods of time and have distinctive spectral features, making them easy to discriminate from standard Type Ia or Type II supernovae.  Unfortunately, time dilation may make the light curves of extremely high redshift supernovae difficult to detect, especially those powered by the pair instability, since they have long evolutionary times in the rest frame.  Their observation would require very different survey strategies than those employed in low redshift supernova searches.  Furthermore, the predicted Population III supernova rate varies widely depending on a range of basic assumptions \\citep{wl05,wa05,mes06}.  If Population III supernovae are ever observed, they may be useful probes of their immediate environments.  Peculiar features in a supernova light curve can be exploited to infer the properties of circumstellar gas \\citep{sm07}.  If Population III stars are rapidly rotating, they may explode as gamma ray bursts \\citep{bl06} and could be used to probe the interstellar medium of the halos in which they form, as has been done in the $z <$ 4 universe \\citep{cet05,pet06,pet07}.  The ionization front of the progenitor itself would also be prone to violent dynamical instabilities in three dimensions due to H$_2$ cooling, forming gas clumps in the relic \\ion{H}{2} region with which the blast would interact \\citep{wn07a,wn07b}.  They would be most prominent in \\ion{H}{2} regions trapped in massive halos, partly because D-type fronts better shield H$_2$ from Lyman-Werner photons from the star and partly because R-type I-fronts that break free of halos would tend to photoevaporate any clumps that may have formed, although numerical simulations indicate that some would survive \\citep{wn07a,wn07b}. I-front instabilities may also drive turbulence in the surrounding medium.  Both fragmentation and turbulence will enhance the mixing of metals in halos and should be included for self-consistency in future three dimensional models.  Molecular hydrogen and HD could also cause the remnant to cool and fragment, a process not addressed in these models.  Explosions in ionized halos would sweep up H$_2$ and HD that are rapidly forming in the relic \\ion{H}{2} region.  However, molecular cooling would not occur until late times in our models because the remnant remains fully ionized and hot for several Myr, as shown in Figure \\ref{fig:260Ms_halo1_2}f.  At 7.93 Myr, H$_2$ might form in the forward shock as it partially recombines and cools to 10000 K, forming a dense shell that might be prone to breakup.  In all likelihood, prior instabilities, such as Rayleigh-Taylor instabilities in the shock, would pre-empt those due to H$_2$ cooling, but the radiative losses might still be relevant to the kinematics of the remnant.  Molecular hydrogen cooling might also be viable in trapped explosions.  Figures \\ref{fig:40Ms_halo3_2}b and f show that the leading shock in the contained blast recombines and cools at much earlier times, at about the same time instabilities would erupt in the contact discontinuity.  Rayleigh-Taylor instabilities in the reverse shock and Vishniac instabilities in the forward shock mediated by H$_2$ cooling could evolve together, destroying the shell from within and without.  In the absence of magnetic fields in the vicinity of the progenitor, how would the free expansion interact with the \\ion{H}{2} region or neutral halo?  Simple collisional cross section arguments imply that the ejecta particles would travel to great distances with little deflection by the surrounding gas unless magnetic fields mediate hydromagnetic interactions between the blast and its environment.  If no primordial fields were present, the ejecta might interact differently with the halo than predicted by the fluid equations, which are predicated upon mean free paths in the flow being far shorter than characteristic lengths in the flow itself.  Nevertheless, for now it is reasonable to suppose that magnetic fields arise due to the Weibel instability \\citep[e.~g. ][]{md07}, effectively coupling the two plasma streams.  This mechanism has been invoked to explain the origin of GRB afterglows \\citep{mlm07}.  An important phenomenon not noted in earlier surveys of primordial supernovae are the large central infall rates that can result from fallback when remnants fail to escape massive halos, in excess of 10$^{-2}$ $\\Ms$ yr$^{-1}$ and lasting 1 - 2 Myr in some instances.  Such infall, although less coordinated, would certainly appear in three dimensional simulations and lead to rapid growth of the central black hole. The growth would be episodic since bounceback would periodically reverse the inflow, and three dimensional radiation hydrodynamical accretion simulations would be required to ascertain how much gas at the center of the halo would fall into the black hole. Nevertheless, this mechanism may provide the means whereby black holes of several thousand solar masses could form in pregalactic structures at $z \\sim$ 20, which numerical studies suggest could be the origin of the supermassive black holes occupying the centers of most massive galaxies today \\citep{liet07}.  Periodic accretion onto the black hole would emit bursts of radiation from the halo that might over time contribute to an x-ray background at high redshifts and strongly affect the thermal and chemical properties of the halo.  One alternative to this scenario is that gas in the trapped explosion instead cools to tens of Kelvin by metal-line cooling and then fragments, forming many low mass stars whose peculiar motions are confined to the dark matter potential of the halo.  Several tens of thousands of solar masses of gas could fuel low mass star formation in this manner.  If large numbers of small stars can truly form in the halo the result could be a globular cluster 15 - 30 pc in diameter, the radii to which gas in the halo is disturbed by the blast.  High dynamical range adaptive mesh refinement (AMR) simulations are necessary to determine if gas in the halo can fracture and collapse in this fashion.  Finally, our supernova models may hint at an alternative pathway for the assembly of the first primitive galaxies.  \\citet{get07} find that the first supernovae preferentially expel metals into the voids interspersed among cosmological filaments.  Since densities are relatively low there, the metals do not promote immediate star formation, and must generally await merger timescales to be taken up into a new generation of stars.  If supernovae ejecta instead mix efficiently with the dense shells of relic H II regions and cause them to cool and fragment, a second generation of stars with different chemical abundance patterns might promptly form.  Likewise, explosions trapped by massive halos could also create small populations of new stars before mergers with other halos.  If so, the first primitive galaxies were assembled sooner, with more stars and different chemical abundances than in current cosmological models.  Upcoming observations by the \\textit{James Webb Space Telescope} and \\textit{Thirty Meter Telescope} might discriminate between these two scenarios."
    },
    "0801/0801.4840_arXiv.txt": {
        "abstract": "This paper investigates the possibility that UCD galaxies in the Fornax cluster are formed by the threshing of nucleated, early-type dwarf galaxies (hereafter dwarf galaxies).  Similar to the results of \\citet{CPF06} for the Virgo cluster, we show that the Fornax Cluster observations are consistent with a single population in which all dwarfs are nucleated, with a ratio of nuclear to total magnitude that varies slowly with magnitude.  Importantly, the magnitude distribution of the UCD population is similar to that of the dwarf nuclei in the Fornax cluster.  The joint population of UCDs and the dwarfs from which they may originate is modelled and shown to be consistent with an NFW profile with a characteristic radius of 5\\,kpc.  Furthermore, a steady-state dynamical model reproduces the known mass profile of Fornax.  However, there are a number of peculiarities in the velocity dispersion data that remain unexplained.  The simplest possible threshing model is tested, in which dwarf galaxies move on orbits in a static cluster potential and are threshed if they pass within a radius at which the tidal force from the cluster exceeds the internal gravity at the core of their dark matter halo.  This fails to reproduce the observed fraction of UCDs at radii greater than 30\\,kpc from the core of Fornax. ",
        "introduction": "In recent years considerable evidence has accumulated that disruptive processes play an important role in galaxy evolution as well as the more dominant hierarchical merging. Observational evidence for these disruptive processes is particularly evident in the dense environment of galaxy clusters. The evidence includes populations of individual intra-cluster objects such as planetary nebulae \\citep{Arn04,FCJ04} and red giant stars \\citep{Dur02}, as well as the general diffuse light now thought to make up a significant fraction of the total stellar mass in clusters \\citep*{FCJ04,FMM04,GZZ05,Zib05}.  In this paper we focus on a relatively new component of intra-cluster space, ultra-compact dwarf (UCD) galaxies. These are compact systems of old stars akin to globular clusters but they are 10--100 times more luminous than Galactic globular clusters and they are located in intra-cluster space between galaxies. The first UCDs were discovered in the Fornax Cluster independently in studies of globular clusters \\citep*{Min98,HIV99} and in studies of compact dwarf galaxies \\citep{DJG00,PDG01}. The UCDs are unlike any known galaxies in terms of luminosity, morphology and size \\citep{DGH03}. Several hypotheses have been suggested to explain the origin of UCDs ranging from them being the high-luminosity end of a putative intra-cluster globular cluster distribution, to being the evolved super star clusters formed in galaxy merger events. In this paper we focus on the model that UCDs are formed by the global tidal field of a cluster which can strip, or ``thresh'', the outer stellar envelopes of nucleated dwarf galaxies (dE,Ns and dS0,Ns) as they pass repeatedly through the inner regions of a cluster leaving just the bare nucleus to survive as a UCD \\citep*{BCD01,BCD03,GMK07}  The motivation for the current work is the subsequent discovery of a larger population of fainter UCDs in the central region of the Fornax Cluster \\citep{DGC04,Gre08}. This sample of 60 UCDs is large enough to permit us to test several aspects of the threshing hypothesis using a statistically significant sample. Our focus will be to test simple aspects of the distributions of the UCD and galaxy populations. An alternative approach based on the internal properties of the UCDs is also in progress \\citep[e.g.~][]{EGD07}.  Our basic premise for this paper is that if UCDs are descendants of disrupted galaxies, then the UCD parent population can be modelled by the combined {\\em current} population of Fornax cluster UCDs and dwarf galaxies.  We test whether the observed spatial and velocity distributions of the two populations are consistent with this hypothesis and conclude that they are.  We then model the orbits of UCDs/galaxies drawn from this joint population to determine what fraction of them pass close enough to the cluster centre to lead to threshing.  The relative fraction UCDs to % dwarfs seen at large radii in Fornax is inconsistent with this static threshing model.  In Section~\\ref{sec:data} we define the UCD and galaxy samples for our analysis. In Section~\\ref{sec:luminosity} we test if the luminosity function of the UCDs is consistent with them having been drawn as random sample from the nuclei of dwarf galaxies in the cluster. Section~\\ref{sec:dynamical} develops a dynamical model for the joint population, and Section~\\ref{sec:thresh} calculates the fraction of threshed orbits at each radius.  Finally, in Section~\\ref{sec:conc}, we summarise our results and draw conclusions about the plausibility  of the threshing hypothesis.  We adopt a distance of 20\\,Mpc to the Fornax Cluster \\citep*{DGC01} corresponding to a distance modulus of 31.51 magnitudes.  In this paper we are not concerned with late-type galaxies.  To avoid endless repetition, we use the terms galaxy and dwarf to refer to early-type objects only, as defined in Section~\\ref{sec:early}. ",
        "conclusions": "\\label{sec:conc}  In this paper we have investigate the possibility that UCD galaxies are formed by the threshing of nucleated, early-type dwarf galaxies.  We first contrast the distribution of nucleated and non-nucleated dwarfs, which are indistinguishable apart from a small excess of bright, nucleated dwarfs at small clustocentric radii.  We concur with the conclusion of \\citet{CPF06} that the observations are consistent with a single population in which all dwarfs are nucleated, with a ratio of nuclear to total magnitude that varies slowly with magnitude.  However, we need to flatten their relation in order to obtain a good fit when extrapolating to fainter magnitudes.  Given this hypothesis, we can reproduce the magnitude distribution of the UCD population, except at bright magnitudes where the model predicts more UCDs than are observed.  Under the threshing model, the UCDs are likely to have originated from dwarfs with magnitudes brighter than about $M_B=-15$.  We use the joint UCD plus bright dwarf population in the modelling that follows.  The threshing model predicts that over half of all dwarf galaxies must be disrupted: 38 surviving dwarfs have nuclei of similar magnitude to the 49 observed UCDs.  This may seem excessive but corresponds to an intracluster light fraction of just 8 per cent, well within the observed range for clusters of this mass \\citep{FCJ04,FMM04,GZZ05,Zib05}.  The distribution of dwarf galaxies in Fornax follows that of the total mass distribution and shows no evidence for disruption of dwarfs near the cluster core.  Nevertheless, the UCD population is more centrally-concentrated than the dwarfs, as would be expected in the threshing model.  If we assume that the joint population is in a steady-state dynamically, then it should also satisfy the Jeans Equations.  We show that the joint population is well-described by a density distribution of the form \\begin{equation} \\rho=\\frac{\\rho_0}{x\\,(1+x)^{s-1}}, \\end{equation} where $x=r/a$, and $a$ and $s$ are constants, with $s$ lying between about 3 and 4.5.  The velocity dispersion of UCD galaxies shows a sharp decline with radius that is hard to explain.  It may in part be due to a radial bias in the orbits, but this is not enough in itself to explain the effect.  The velocity dispersion of bright dwarfs is greater than that of the UCDs.  When the two are combined, then the joint population with density slope $s=3$ provides a marginally-acceptable fit to the mass profile of Fornax.  We have tested the simplest possible threshing model, in which dwarf galaxies move on orbits in a static cluster potential and are threshed if they pass within a radius at which the tidal force from the cluster exceeds the internal gravity at the core their dark matter halo.  This fails to reproduce the observed fraction of UCDs at radii greater than 50\\,kpc from the core of Fornax.  There are many deficiencies in the model but these are unlikely to raise the threshing radii by a factor of 5, as is required, and so we conclude that this static mode is unviable.  Our results have several points of agreement with the earlier work by \\citet{BCD03} despite a very different approach: we have used analytic descriptions of the cluster dynamics compared to their numerical computations.  In our work we have based our prediction on a parent sample of dwarf galaxies generated directly from the known Fornax galaxies, whereas Bekki et al.\\ generated their galaxy sample from more general empirical relations for the luminosity functions and radial profiles of galaxies within clusters.  In particular, they used a King profile with a core radius of 50\\,kpc for the density distribution, very different from our model.  They demonstrate that dwarf galaxies are disrupted if they pass inside their critical threshing radius when orbiting the cluster centre. They then use this radius to estimate the population of threshed galaxies (UCDs) in the Fornax Cluster.  They find this to be consistent with the known distribution of the 7 very luminous UCDs known in the cluster at that time.  Our conclusion (refuting the simple threshing model) differs from that of Bekki et al.\\ for a number of reasons.  Firstly, we use the measured positions of galaxies in Fornax, rather than a generic King model.  We also have many fewer dwarf galaxies than predicted by their Schechter model for the cluster luminosity function.  In addition, we have extended the analysis to much lower luminosities of both the UCDs (as new data have become available) and the parent galaxies (due to the greater difference in nuclear to total luminosity now used).  This new analysis has clearly revealed a disagreement between the number of UCDs at large clustocentric distances and the threshing predictions.  A recent paper by GMK07 undertook an extensive series of simulations to investigate the disruption of UCD host galaxies within a cluster potential similar to that of the Virgo cluster.  They considered two different models for the host galaxy with very different degrees of central concentration and followed their threshing in a static potential over 5 Gyr.  They then looked at the orbits of particles in a cosmological simulation of cluster formation to assess which of those orbits would lead to threshing.  This latter step follows the dynamical evolution of the halo and is much more realistic than a static potential.  GMK07 state that their model ``leads to the observed spatial distribution of UCDs'', in apparent disagreement with our results above.  In fact our theoretical UCD fractions as a function of radius agree with theirs and are bracketed by their upper and lower predictions.  The difference in the conclusion arises from the very different observed threshing fractions that we adopt.  GMK07 use only 15 UCDs in both Virgo and Fornax combined, whereas we use a new sample of 49 UCDs from Fornax alone.  Also, GMK07 do not say how they define the nucleated dwarfs corresponding to the parent sample, whereas we are careful to select only those dwarfs that would have nuclei that match those of the observed UCDs.  In conclusion, the origin of UCDs as dwarf galaxy nuclei remains unproven.  Our modelling has revealed a number of attractive features: \\begin{description} \\item The distribution of nuclear magnitudes for dwarf galaxies roughly matches that of known UCD galaxies. \\item UCDs are more centrally-concentrated within Fornax than are dwarf galaxies.  (However, this would also be true if the UCDs constituted an extended globular cluster population around NGC\\,1399.) \\item The joint UCD plus bright dwarf population has a smooth density profile with a recognisable (NFW) form and appears to sit in dynamical equilibrium within the Fornax cluster. \\end{description} At the same time, there are several major deficiencies in the model: \\begin{description} \\item The model requires that more dwarf galaxies must have been disrupted in Fornax than currently remain.  However, the spatial distribution of dwarfs matches that of the total mass profile of the cluster and shows no sign of galaxy disruption near the cluster core. \\item The very low velocity dispersion of UCDs as compared to bright dwarfs is unexplained, as is the sharp decline in velocity dispersion of the UCDs with radius.  (However, this would prove true for any dynamical model of the UCD population, regardless of its origin.) \\item A static threshing model for UCD formation, based upon orbits within the current cluster potential, is a hopeless failure.  It predicts far too few UCDs at radii greater than about 30\\,kpc. \\item The simulations of GMK07 within an evolving cluster potential also give too few UCDs at large radii. \\end{description} The balance of evidence would seem to be against the threshing model. Before dismissing the model altogether, however, we note that the threshing may have occurred within smaller sub-clusters that later fell into Fornax and have not yet reached dynamical equilibrium. This mechanism is suggested by the spatial distribution of UCDs in the Fornax Cluster: they show some association with normal galaxies and, in particular, lie in a band across the cluster \\citep{Gre08}.  In considering the threshing hypothesis for UCD formation we should not discuss the dynamical properties of the objects in isolation from their internal properties.  \\citet{EGD07} studied the stellar populations of Virgo Cluster UCDs and concluded that the Virgo UCDs have stellar populations the globular clusters of the central galaxies M87 and M49 (old ages, a range of metallicity, and supersolar alpha-abundances).  On this basis, the Virgo UCD stellar populations are not consistent with simple threshing model.  On the other hand, \\citet{MHI06} found metallicities and (a range of) ages in Fornax Cluster UCDs, which are more in agreement with the hypothesis that the Fornax UCDs are threshed nuclei.  A detailed analysis of the structure and colours of both Virgo and Fornax UCDs \\citep{Evs08} concluded that their structural properties could be consistent with either globular clusters or dwarf galaxy nuclei, with the interesting observation that UCDs are about twice as extended (in effective radius) as the nuclei of dwarf galaxies at the same luminosity.  Most of these observational results, as well as our own analysis in this paper, argue against the simple threshing hypothesis for UCD formation."
    },
    "0801/0801.0591_arXiv.txt": {
        "abstract": "Continuum observations at 350\\,$\\mu$m are presented of seven nearby elliptical galaxies, for which CO-gas disks have recently been resolved with interferometry mapping. These SHARC\\,II mapping results provide the first {\\em clearly resolved} far-infrared(FIR) to submillimeter(submm) continuum emission from cold dust (with temperatures 32\\,K $> T >$ 22\\,K) of any elliptical galaxy at a distance $>40$\\,Mpc. The measured FIR excess shows that the most likely and dominant heating source of this dust is not dilute stellar radiation or cooling flows, but rather star-formation, that could have been triggered by an accretion or merger event and fueled by dust-rich material that has settled in a dense region co-spatial with the central CO-gas disks. The dust is detected even in two cluster ellipticals that are deficient in H\\,I, showing that, unlike the H\\,I, cold dust and CO in ellipticals can survive in the presence of hot X-ray gas, even in galaxy clusters. No dust cooler than 20\\,K, either distributed outside the CO disks, or co-spatial with and heated by the entire dilute stellar optical galaxy (or very extended H\\,I), is currently evident. ",
        "introduction": "\\label{sec:intro_es}  Cold gas and dust in nearby elliptical galaxies were detected only fairly recently \\citep[][]{sad85, kna89, lee91, wik95}. Compared to that in spiral galaxies, the cold Interstellar Medium (ISM) in ellipticals is present in relatively small amounts and is seen in only 50\\% to 80\\% of nearby ellipticals.  The source and content of the cold ISM in these galaxies is still uncertain, with optical and far-infrared (FIR) dust-mass estimates differing by $\\sim$ 10 to 100 \\citep[e.g.,][]{gou95_b}.  There is some evidence for an internal origin of the dust arising in circumstellar envelopes of red giant stars. This internal origin is supported by the detection of $\\sim 10\\,\\mu$m emission in excess of the stellar emission, with the same de Vaucouleurs profile as the optical stellar images and near- to mid-IR photospheric emission \\citep[e.g.,][]{ath02}. Some have suggested that this dust would come in contact with the hot ISM ($\\sim 10^7$\\,K) and be eroded in thermal collisions with ions and destroyed in only $\\sim 10^7$ to $10^8$\\, years \\citep[e.g.,][]{sch98_w}. However, \\citet[][]{mat03} have argued that the dust may induce efficient cooling of the hot gas on a timescale shorter than the destruction lifetimes of the dust in the hot gas, allowing some dust to survive and settle primarily in the centers of the galaxies while the hot gas cools.  They propose that if this dust-induced cooling is the dominant evolutionary path of the dust from stars, the thermal energy released by the hot ISM should be re-radiated by dust in the FIR and could be responsible for the central dusty disks and clouds seen in ellipticals.  It has also been suggested that mergers or accretion events that include a gas-rich galaxy may be the dominant source of dust in ellipticals.  The total far-IR luminosities \\citep[e.g.,][]{tem04} and cold gas content \\citep[][]{kna85, lee91} of ellipticals are observed to be uncorrelated with their optical luminosities.  This lack of correlation, together with HST detections of irregular dust lanes and filaments with random orientations with respect to the optical major axes \\citep[e.g.,][]{tra01} is usually interpreted as evidence that the dust has been externally acquired. The external origin is likely to be accretion via the interaction of an elliptical galaxy with a gas-rich spiral, or in mergers of gas-rich spirals that evolve through an FIR-bright (starburst) phase as they relax to form ellipticals \\citep[e.g.,][]{too72, sch98_w}. In fact, if elliptical galaxies are formed through spiral mergers, then starbursts are necessary to produce the high phase-space densities seen in the centers of elliptical galaxies \\citep[e.g.,][and references therein]{rot04}. Numerical simulations by \\citet[][]{bar02} have shown that 20\\% to 60\\% of the gas from two spirals merging can form relaxed disks of sizes up to 20\\,kpc, with some of the gas maybe losing angular momentum, falling to the nucleus of the galaxy, and perhaps leading to star-formation.  Observations of dust emission from elliptical galaxies have been made using $IRAS$ or $ISO$ primarily \\citep[e.g.,][]{kna89, tem04}. These measurements favored the warm dust components of the ISM and produced unresolved images of the dust emission.  The $IRAS$ and $ISO$ derived SEDs of these galaxies clearly show that there are multi-temperature components to the dust. Indeed, the presence of an additional, low-level, very cold ($<20$\\,K) dust component in ellipticals has been suggested by limited detections of unresolved mm and submm emission in some nearby galaxies \\citep[][]{kna91, wik95b, lee04, tem04}, as well as clearly extended FIR to submm detections from the dust lane in the nearest giant elliptical galaxy NGC\\,5128 that hosts the active galactic nucleus (AGN) Centaurus\\,A \\citep[e.g.,][]{lee02, qui06}. Observations of CO emission from a number of ellipticals show that many of them contain significant amounts of molecular gas \\citep[][]{lee91, wik95}, but many of these observations do not spatially resolve the CO emission.  Recently, interferometric measurements by \\citet[][]{you02, you05} and \\citet[][]{wik97} have spatially resolved the CO (1-0) emission in seven elliptical galaxies with spatial resolutions of 7 and 3 arcsec, respectively. These data imply that the molecular gas is contained mostly in rotating disks lying within the central 5\\,kpc of the galaxies. This prompted the study by our group to spatially resolve the dust emission in these seven elliptical galaxies. Our goals were (a) to determine the physical properties of the dust coincident with the molecular gas; (b) to see if there are dust associations with other spatial distributions (e.g., stellar distribution); and (c) to determine the most likely heating source (or sources) for the dust. We wanted also to address the question of whether these observed gas disks are the result of external merger/accretion events, as suggested by the optical morphological signatures of the galaxies, or internal mass-losing stars that constitute their elliptical optical components.  The seven galaxies in our sample (NGC\\,83, NGC\\,759, NGC\\,807, UGC\\,1503, NGC\\,3656, NGC\\,4476, and NGC\\,5666) were all classified by \\citet[][]{wik95} as ellipticals, based on the determination that their luminosity profiles follow the $R^{-1/4}$ de Vaucouleurs law. However, morphologically the sample includes galaxies from different extragalactic environments and represents a spread of possible merger/accretion tracers or ages, from galaxies that have been classified as on-going or early-age major mergers \\citep[e.g., NGC\\,3656,][]{bal01} to very-late accretion or quiescent systems \\citep[e.g., NGC\\,807,][]{mur00}. Recent CCD imaging of some of the galaxies, e.g., NGC\\,5666 by \\citet[][]{don03}, have disputed the classification as ellipticals in favor of SO galaxies, while new optical observations of others, e.g., NGC\\,83 by \\citet[][]{you05}, have confirmed the elliptical classification.   The sample's updated environments, general properties, and molecular gas content are summarized in Table 1. The observational merger or accretion signatures of the sample are worthy of note as they fortuitously allow an evolutionary analysis in the context of merger/accretion stages and cold dust and gas properties. These merger/accretion signatures are described further with other observational characteristics of the galaxies in Section~\\ref{sec:submm_morph} below. ",
        "conclusions": "\\label{sec:sum_es}  SHARC\\,II continuum observations at 350\\,$\\mu$m have been presented of seven elliptical galaxies with known CO-gas disks \\citep[][]{wik97, you02, you05}. The SHARC\\,II mapping provides the first {\\em clearly resolved} FIR-submm continuum emission from cold dust (with temperatures $\\sim 31$\\,K $> T >$ 23\\,K) of any elliptical galaxy $>40$\\,Mpc. Calculations of the measured FIR excess show that the dust's most likely and dominant heating source is not dilute stellar radiation or cooling flows but rather star-formation, that could have been triggered by an accretion or merger event and fueled by dust-rich material that has settled in a dense region co-spatial with the central CO-gas disks. The dust is detected even in two cluster ellipticals that are deficient in H\\,I, showing that, unlike the H\\,I, cold dust and CO in ellipticals can survive among hot X-ray gas, even in galaxy clusters.  Above the 50\\% contour levels, where the submm detections are of high significance, the submm dust continuum and optical stellar intensities of similar spatial resolution have the same general distribution, suggesting that the dust, CO, and stellar components in the central regions are co-spatial. The early-merger elliptical NGC\\,3656, with its dense optical dust lane that is thought to be in a highly inclined warped disk, is an exception to this coupling of submm and optical central flux distributions.  Below the 50\\% contour levels, and especially below the 30\\% levels, the current maps are not detected at high enough significance that proper dust distributions can be discerned.  No dust cooler than 20\\,K, distributed outside the CO disks, or co-spatial with and heated by the entire, dilute stellar optical galaxy (or very extended H\\,I) is currently evident. These observations clearly show that the FIR-submm emission in our sample ellipticals is primarily from cold dust that is associated with star-formation located in the central parts of these galaxies. The results support the hypothesis that if ellipticals are formed through mergers of gas-rich spirals, their starbursts (1) are necessary to produce the high phase-space density seen in the center of ellipticals and (2) would power the FIR-bright stage as the merger relaxes. Combining the current data set together with anticipated $Spitzer$ and other upcoming sensitive FIR-to-submm mapping observations may detect or rule-out the existence of larger-scale, diffusely distributed cold dust with an average temperature that is perhaps $<20$\\,K (and maybe associated with H\\,I).  Assuming the grain mass absorption coefficient $k_{\\rm d}^{350\\,\\mu {\\rm m}} = 1.92 \\times 10^{-1}\\,{\\rm m}^2 {\\rm kg}^{-1}$, as for Galactic-like dust \\citep[][]{dra03}, the measured submm fluxes and dust temperatures yield dust masses for the sample that range from $\\sim 9 \\times 10^5 {\\rm M}_{\\odot}$ to $ \\sim 2 \\times 10^7 {\\rm M}_{\\odot}$. The sample has diverse H$_2$-mass to dust-mass ratios that cover $60 \\la M$(H$_2$)/$M_{\\rm dust} \\la 310$, suggesting that in the central galactic regions, where the H$_2$ and dust masses are respectively measured from the resolved CO and our CSO/SHARC\\,II continuum measurements, some of the sample galaxies are relatively under-massive in H$_2$ and over-massive in dust content, while others are the opposite. The possible implications of these mass ratios on star-formation or, at least, gas-phases in the central regions of these galaxies will be investigated further in our upcoming paper that uses $Spitzer$ IRS data to study the detailed MIR spectral and/or star-formation properties of these galaxies.  The total cold-gas-mass to dust-mass ratios for the sample are in a narrower range of $230 \\la M$(H$_2$ + H\\,I)/$M_{\\rm dust} \\la 400$, with a total cold-gas-mass to dust-mass ratio for the sample that, at face value, may appear to be about 3 times the value for the Milky Way Galaxy (i.e., $ \\sim 100$). However, in sample objects where atomic gas has been resolved \\citep[e.g.,][]{lak87, bal01}, the H\\,I has been observed to be much more extended than the inner galactic regions where our observations have mapped dust emission that is associated with molecular gas. A cold and extended dust component may yet be detected that is associated with the atomic hydrogen, and perhaps massive enough to bring the total cold gas-to-dust mass ratio for the sample in line with the Milky Way value. As noted above, this may be oberved with future more sensitive FIR/submm detectors such as those that will be on SOFIA, $Herschel$, JCMT, or ALMA."
    },
    "0801/0801.2594_arXiv.txt": {
        "abstract": "We use the previously-identified 15 infrared star-cluster counterparts to X-ray point sources in the interacting galaxies NGC 4038/4039 (the Antennae) to study the relationship between total cluster mass and X-ray binary number.  This significant population of X-Ray/IR associations allows us to perform, for the first time, a statistical study of X-ray point sources and their environments.  We define a quantity, $\\eta$, relating the fraction of X-ray sources per unit mass as a function of cluster mass in the Antennae.  We compute cluster mass by fitting spectral evolutionary models to $K_s$ luminosity. Considering that this method depends on cluster age, we use four different age distributions to explore the effects of cluster age on the value of $\\eta$ and find it varies by less than a factor of four. We find a mean value of $\\eta$ for these different distributions of $\\eta$ = 1.7$\\times$10$^{-8}$ $M_{\\sun}^{-1}$ with $\\sigma_{\\eta}$ = 1.2$\\times$10$^{-8}$ $M_{\\sun}^{-1}$.  Performing a $\\chi^2$ test, we demonstrate $\\eta$ could exhibit a positive slope, but that it depends on the assumed distribution in cluster ages.  While the estimated uncertainties in $\\eta$ are factors of a few, we believe this is the first estimate made of this quantity to ``order of magnitude'' accuracy.  We also compare our findings to theoretical models of open and globular cluster evolution, incorporating the X-ray binary fraction per cluster. ",
        "introduction": "The Antennae are a pair of colliding galaxies with an unusually large number of X-ray point sources.  High resolution X-ray observations taken with {\\it Chandra} revealed 49 new individual sources \\citep{fab01}, where previous observations only indicated extended filamentary structure \\citep{fab97}.  These X-ray sources range in luminosity from $10^{38}-10^{40}$ ergs s$^{-1}$.  Most of these are thought to be X-ray binaries (XRBs) with a black hole compact companion \\citep{fab01}.  In addition to numerous X-ray sources, the Antennae contain many bright, massive young clusters that are evident in both optical, {\\it HST} \\citep{whi99} and infrared (IR) \\citep[henceforth Paper I]{bra05} images.  This makes this pair of interacting galaxies an ideal target for studying the environments of XRBs.  In \\citet[henceforth Paper II]{cla07} we performed an extensive study of the XRB environments using {\\it Chandra} X-ray images and $J$ and $K_s$ IR images (Paper I).  Our present paper will expand on our previous study by exploring the relationship between XRBs and cluster mass in the Antennae.  Recent theoretical models of young, massive cluster evolution provide a framework for comparison to our observational study.  Two in particular, \\citep{osk05,sep05}, incorporate the fraction of XRBs per cluster in their models.  \\citet{osk05} use a population synthesis code to study the evolution of X-ray emission in young, massive clusters.  \\citet{sep05} investigate the role of supernova kicks in XRB expulsion from the parent cluster using the population synthesis code, StarTrack.  They also incorporate the number of XRBs for a range in cluster mass.  We will compare our measurements of the fraction of XRBs per cluster in the Antennae to those predicted by these models.  We organize our paper as follows: In \\S2 we give a brief summary of our previous work on the Antennae and then define a quantity, $\\eta$, relating the XRB fraction to cluster mass in the Antennae.  In our analysis, we estimate cluster mass using $K_s$ luminosity, which depends non-trivially on the assumed cluster age.  We explore cluster age/luminosity relations and their impact on our mass estimates in \\S3.  We compare $\\eta$ to the measured value predicted by theoretical cluster evolutionary models and present conclusions in \\S4. ",
        "conclusions": "Through the quantity $\\eta$, our investigation revealed conflicting results with respect to the relationship between observed number of XRBs and cluster mass.  We performed a $\\chi^2$ test on the slope in $\\eta$ for a variety of assumed star cluster age distributions and found some cases where the slope was insignificant, while other cases showed a distinctly positive slope.  No slope indicates the observed number of XRBs per unit mass is independent of cluster mass, while a positive slope in $\\eta$ suggests more XRBs per unit mass are produced in more massive clusters.  In the following discussion, we will consider both a constant value in $\\eta$ as well as a slope in $\\eta$, and the implications of each result.  Initially, we estimated cluster mass by fitting BC spectrophotometric models to cluster $M_{K_s}$, assuming all clusters are $\\sim$20 Myr. Recognizing that this method depends on cluster age, we explored how different assumptions of a cluster age distribution for the Antennae affect $\\eta$ and showed that $\\eta$ varies by a factor of roughly four; although including individual ages for the instant bursts case will increase the variations in $\\eta$ to a factor of 10.  We now proceed by comparing the mean value of $\\eta$ for the four different assumed age distributions, $\\eta$ = 1.7$\\times$10$^{-8}$ $M_{\\sun}^{-1}$, to that predicted by models of young, massive clusters.  We will compare $\\eta$ to theoretical models discussed in \\citet{osk05} and \\citet{sep05}.  In a recent study presented in \\citet{osk05}, the author modeled X-ray emission from young, massive star clusters, assuming a closed system with constant mass, no dynamics and all stars are coeval, with cluster metallicities of either z = 0.02 or z = 0.008.  These models predict $\\sim$2-5\\% of all OB stars in a cluster should produce high mass X-ray binaries (HMXBs).  In the models in \\citet{osk05}, all clusters are assumed to have masses of M$_{cl}$ = 10$^6$ $M_{\\sun}$ with stellar masses ranging from 1 -- 100 $M_{\\sun}$.  Considering the Salpeter initial mass function (IMF) of the form $\\xi(M)$ = $M_0M^{-2.35}$ and defining stars with masses $>8$ $M_{\\sun}$ as ``OB stars'', for our purposes here, we estimated 6\\% of all stars in the model clusters are OB stars.  Therefore, 1 -- 3$\\times$10$^{-3}$ of all stars in a cluster with an initial mass of 3$\\times$10$^6$ $M_{\\sun}$ (set by the Salpeter IMF) should produce an XRB.  Since the Salpeter IMF implies there are 7$\\times$10$^5$ stars in a cluster, then these stars should produce 7 -- 22$\\times$10$^2$ XRBs -- orders of magnitude greater than the $\\sim$49 observed in the Antennae. Expressing $\\eta$ as a fraction of XRBs-to-cluster mass, the models in \\citet{osk05} suggest $\\eta$ ranges from 3--7$\\times$10$^{-4}$ $M_{\\sun}^{-1}$.  These values are greater by at least a factor of 1000 from our estimates for $\\eta$.  Clearly, this predicts a much larger number of compact object binaries than what we observed in the Antennae.  \\citet{osk05} note that they were unable to detect HMXBs in three massive ($\\sim10^4$ $M_{\\sun}$) clusters which they predict should contain between 1-3 HMXBs.  If our measured value for $\\eta$ accurately describes the number of HMXBs formed, then it is not surprising that \\citet{osk05} fail to find any.  As pointed out by these authors, future modeling of HMXB formation is needed to understand the discrepancy between the predictions and observations of XRB populations in starburst galaxies.  In another study, \\citet{sep05} use the binary evolution and population synthesis code, StarTrack \\citep{bel02}, to investigate the rate of XRB formation and ejection from young, massive clusters.  This program tracks stellar parameters such as radius, luminosity, mass and core mass.  The simulations are stopped at the formation of a compact object.  The models include mass transfer in binaries and include transient XRBs.  \\citet{sep05} consider cluster masses ranging from $5\\times10^4$ $M_{\\sun}$ to $5\\times10^6$ $M_{\\sun}$ and cluster ages from 1 to $\\sim$20 Myr and compute the average number of XRBs within 1--1000 pc of the cluster center.  Considering the typical cluster age in the Antennae is 20 Myr, \\citet{sep05} predict a $5\\times10^4$ $M_{\\sun}$ cluster should contain 0.13 XRBs, while 15 XRBs should reside in a $5\\times10^6$ $M_{\\sun}$ cluster.  Here we assume that an XRB is associated with a cluster if it is within 100 pc.  This separation is equivalent to $1\\farcs0$ at the distance of the Antennae (for $H_{0}$=75 km s$^{-1}$ Mpc$^{-1}$), which is our criteria for an XRB-cluster association (Paper II).  These model predictions for XRB detections assume a limiting X-ray luminosity of $L_X$ = 5$\\times$10$^{35}$ ergs s$^{-1}$, but the observed limiting luminosity in the Antennae is 2$\\times$10$^{37}$ ergs s$^{-1}$.  Using the X-ray luminosity function (XLF) for the Antennae defined in \\citet{zez02c}, we scaled the XRB results of \\citet{sep05} to estimate what these models would predict for the observed number of XRBs in the Antennae clusters.  Using a XLF power law slope of $\\alpha$ = -0.45 \\citep{zez02c}, the models predict 0.02 XRBs are observed in a $5\\times10^4$ $M_{\\sun}$ cluster, while 2.7 XRBs should be seen in a $5\\times10^6$ $M_{\\sun}$ cluster, at the luminosity limits of the X-ray observations.  Expressing these model results as a fraction of XRBs-to-cluster mass, we can directly compare them to our measured value for $\\eta$ in the Antennae.  Doing so, \\citet{sep05} predict $\\eta$ ranges from $4-5\\times10^{-7}$ $M_{\\sun}^{-1}$, with in a factor of five from our predictions for $\\eta$.  As mentioned by \\citet{sep05}, several caveats exist for their models including: 1) assumed binary fraction of unity which could lead to over estimates of the mean number of XRBs per cluster, 2) the stellar, power-law IMF can affect the XRB fraction per cluster, and 3) changes in the half-mass radius can strongly influence the median XRB distance from the cluster.  These factors could potentially explain the discrepancies between their models and our observations.  Since some forms of $\\eta$ exhibit a distinctly positive slope, these models may not always apply to $\\eta$.  More importantly, a positive slope has implications for star formation scenarios in clusters.  As mentioned above, a positive slope implies more XRBs per unit mass are produced in more massive clusters.  If an abnormally large number of XRBs exist in a star cluster, then the progenitors of their compact objects should also posses an over abundance.  Since the progenitors are massive stars, this implies star formation in massive clusters favors stars at the heavier end in mass.  This is not unusual.  Work by \\citet{sto05} suggest some of the largest clusters in the Milky Way could have a top heavy mass function.  In this paper we introduced the quantity, $\\eta$, relating the fraction of X-ray sources per unit mass as a function of cluster mass. Applying this function to the Antennae, we revealed several important environmental implications for the X-ray sources in the Antennae. Specifically, $\\eta$ predicts a far different relationship between XRB formation and cluster mass than that predicted by \\citet{osk05} and is broadly consistent with that predicted by \\citet{sep05}.  Clearly, future cluster modeling with particular emphasis on the relationship between the number of XRBs in a galaxy and the galactic cluster environment is essential to explain our current observations. Furthermore, a $\\chi^2$ test demonstrated the functional form of $\\eta$ did not always remain consistent with a single value, but for some assumptions for a cluster age distributions, the slope had a significantly positive value.  While this could imply a top heavy mass function in massive clusters, our statistics are small.  We plan to enlarge our statistical base by extending our observational study to additional starburst galaxies.  We can then address whether the properties of $\\eta$ depend on an individual galaxy or are consistent across all galactic environments."
    },
    "0801/0801.0244_arXiv.txt": {
        "abstract": "Over the last decade, considerable effort has been made to measure the proper motions of the pulsars B1757$-$24 and B1951$+$32 in order to establish or refute associations with nearby supernova remnants and to understand better the complicated geometries of their surrounding nebulae.  We present proper motion measurements of both pulsars with the Very Large Array, increasing the time baselines of the measurements from 3.9~yr to 6.5~yr and from 12.0~yr to 14.5~yr, respectively, compared to previous observations.  We confirm the nondetection of proper motion of PSR~B1757$-$24, and our measurement of ($\\mu_{\\alpha}, \\mu_{\\delta}) = (-11 \\pm 9, -1 \\pm 15)$~mas~yr$^{-1}$ confirms that the association of PSR~B1757$-24$ with SNR~G5.4$-$1.2 is unlikely for the pulsar characteristic age of 15.5~kyr, although an association cannot be excluded for a significantly larger age.  For PSR~B1951$+$32, we measure a proper motion of ($\\mu_{\\alpha}, \\mu_{\\delta}) = (-28.8 \\pm 0.9, -14.7 \\pm 0.9)$~mas~yr$^{-1}$, reducing the uncertainty in the proper motion by a factor of 2 compared to previous results.  After correcting to the local standard of rest, the proper motion indicates a kinetic age of $\\sim$51~kyr for the pulsar, assuming it was born near the geometric center of the supernova remnant.  The radio-bright arc of emission along the pulsar proper motion vector shows time-variable structure, but moves with the pulsar at an approximately constant separation $\\sim$2.5\\arcsec, lending weight to its interpretation as a shock structure driven by the pulsar. ",
        "introduction": "Neutron stars (NSs) are born from the core collapse of massive stars, whose deaths are marked by supernova remnants (SNRs).  Measuring the proper motions of young NSs allows their trajectories to be traced back, and can thus provide strong tests of proposed NS-SNR associations.  In cases where such associations can be confirmed, an independent age estimate can be derived for both the NS and its associated SNR.  However, such an exercise is subject to several pitfalls.  On one hand, a typical radio pulsar might be detectable for $\\gtrsim 10^7$~yr, while the associated SNR fades into the interstellar medium (ISM) in $\\lesssim 10^5$~yr, leaving behind young pulsars with no detectable SNRs.  On the other hand, there are also several SNRs with no associated pulsars.  Not all NSs are radio pulsars, but when young, they are all likely to be bright in thermal X-ray emission, as illustrated, for example, by the detection of point sources in the SNRs Kes 73 \\citep{vg97} and Kes 79 \\citep{ssss03}. The absence of such X-ray detections in a set of nearby, young SNRs \\citep{kfg+04,kgks06} aggravates the problem.  While unusual cooling scenarios may be required for NSs in these young SNRs, a partial explanation for the ``missing'' NSs might lie in their birth velocities.  Radio pulsars have characteristic birth velocities $\\sim$ 400-500~\\kms\\ \\citep{acc02,hllk05}, and they are likely to escape their natal remnants once the expansion of the SNR is decelerated by the ISM, making identification of pulsar-SNR pairs problematic.  Thus, cases where a young radio pulsar can be associated with a SNR are particularly important.  The pulsars B1757$-$24 and B1951$+$32 present two such situations in which associations might be possible. Both are young pulsars with high spin-down energy-loss rates, and in both cases, the ram pressure balance between the NS relativistic wind and the ISM produces a pulsar wind nebula (PWN) with a bow shock structure.  PSR~B1757$-$24 appears to be exiting the approximately circular SNR~G5.4$-$1.2, whose asymmetric brightness, in conjunction with the PWN produced by the pulsar, produces the structure known as ``the Duck'' (see Fig.~\\ref{fig:duck}). Meanwhile, PSR~B1951+32 drives a complex interaction within the SNR~CTB~80 (G69.0$+$2.7; see Fig.~\\ref{fig:ctb80}), including a radio-bright structure resembling a bow shock in the direction of its motion.  We present Very Large Array (VLA) observations that improve the time baseline for the proper motion measurements of these two pulsars.  In \\S\\ref{sec:duck} we detail the observations, analysis, and results for PSR~B1757$-$24, and likewise for PSR~B1951$+$32 in \\S\\ref{sec:ctb80}. We discuss the implications of our results in \\S\\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  \\subsection{The Association between SNR~G5.4$-$1.2  and PSR~B1757$-$24} \\label{sec:discussduck}  The identity of the remnant associated by birth with the pulsar B1757$-$24 is still a mystery.  Based on morphology and location, SNR~G5.4$-$1.2 is the only known remnant likely to be related to the pulsar, but that association has serious problems (see Blazek et al. [2006] for a recent discussion). To remain tenable, such an association requires a pulsar age much larger than the characteristic age $\\tau_c = 15.5$~kyr, an explosion site much closer to the current location of the pulsar, or some combination of the two.  For example, G5.4$-$1.2 could be expanding into a region with a density gradient \\citep[e.g.,][]{gvar04}.  The brightening of the remnant on the western limb, the side nearer to the Galactic plane, may be evidence for such an asymmetric expansion.  However, even in an optimistic scenario, the birth site could be no closer than the westernmost limb of G5.4$-$1.2.  In such a case, for a true age $\\le \\tau_c$, a westward proper motion $\\gtrsim 18$~\\mpy\\ is required for an association, but all three measurements reported in Table~\\ref{tab:summary}, where we summarize our results, imply westward motions no greater than 15~mas~yr$^{-1}$.  It thus appears that either the true age is larger than $\\tau_c$, probably significantly so, or the association of the pulsar with G5.4$-$1.2 must be excluded.  We note that the characteristic ages of some young pulsars with well-determined kinetic ages are overestimates, for example, J1811$-$1925 \\citep{krv+01}, J0538$+$2817 \\citep{nrb+07}, and B1951$+$32 \\citep[our \\S\\ref{sec:discussctb}]{mig02}.  Thus, the true age of B1757$-$24 might be no greater than 15.5~kyr, rendering the association unlikely.  But Vela, for example, may have a true age greater than $\\tau_c$ \\citep{lpg+96}, and \\citet{bgc+06} argue that a growing surface magnetic field, which decouples the true age from the characteristic age, could account for such an age discrepancy in the Duck while preserving the association.  Given a long enough time baseline, the pulsar proper motion (or the motion of the nebula) will surely be detectable, allowing a firm conclusion about this vexing association.  \\begin{deluxetable*}{lccccc}\\vspace{0mm} \\tablecolumns{6} \\tablewidth{0pc} \\tablecaption{Summary of Proper Motion Measurements\\label{tab:summary}} \\tablehead{ & \\multicolumn{3}{c}{B1757$-$24}& \\multicolumn{2}{c}{B1951$+$32}\\\\ & \\colhead{This work} & \\colhead{Ref.\\ 1} & \\colhead{Ref.\\ 2}& \\colhead{This work}& \\colhead{Ref.\\ 3} } \\startdata Target                       & B1757$-$24     & B1757$-$24   & G5.27$-$0.9  & B1951$+32$   & B1951$+$32 \\\\ $\\mu_\\alpha$ (mas yr$^{-1}$) & $-11\\pm9$ & $-2.1\\pm7.0$&  $>-$7.9      & $- 28.8\\pm0.9$& $-29\\pm2$  \\\\ $\\mu_\\delta$ (mas yr$^{-1}$) & $-1\\pm15$ & $-14\\pm13$   &  \\ldots      & $- 14.7\\pm0.9$& $-8.7\\pm1.3$\\\\ Time Baseline (yr)           & 6.5            & 3.9          & 12.0         & 14.5         & 11.9       \\\\ Frequency (GHz)              & 1.4            & 1.4          & 8.5          & 1.4          & 1.4        \\\\ \\\\ RA of Pulsar & \\multicolumn{3}{l}{\\ \\ 18 01 00.016$\\pm$0.008} & \\multicolumn{2}{l}{19 52 58.206$\\pm$0.001}\\\\ DEC of Pulsar& \\multicolumn{3}{l}{$-$24 51 27.5$\\pm$0.2} & \\multicolumn{2}{l}{32 52 40.51$\\pm$0.01}\\\\ Epoch of RA, DEC&  \\multicolumn{3}{l}{\\ \\ 2004 Dec 09} &\\multicolumn{2}{l}{2003 Jun 03} \\enddata \\tablecomments{All provided uncertainties and limits are are 68\\% confidence intervals. \\citet{bgc+06} do not measure $\\mu_\\delta$, and their 68\\% $\\mu_\\alpha$ has been inferred from their 5 $\\sigma$ limit; results of \\citet{mig02} have been adjusted to remove the differential Galactic rotation applied in the published result. Proper motions and positions are reported in J2000.0 coordinates and do not include LSR corrections. Units of right ascension are hours, minutes, and seconds, and units of declination are degrees, arcminutes, and arcseconds. } \\tablerefs{(1) \\citet{tbg02}; (2) \\citet{bgc+06}; (3) \\citet{mig02}} \\end{deluxetable*} \\subsection{The Age and Bow Shock of B1951$+$32} \\label{sec:discussctb}  For B1951$+$32, we measure a proper motion of $(\\mu_{\\alpha}, \\mu_{\\delta})=(-28.8\\pm0.9, -14.7\\pm0.9)$~mas~yr$^{-1}$, confirming and improving on the results of \\citet{mig02}.  \\citet{mig02} filter the {\\it u-v} plane to remove all spatial scales greater than 4\\arcsec\\ and to leave only compact sources. We use a different approach to remove the complexity of the PWN, subtracting the linear slope of the PWN from the region around the pulsar in the image domain.  The inclusion of a longer time baseline and a larger field of reference sources increases both the precision and the accuracy of this measurement over that of \\citet{mig02}.  The new measurement corresponds to an LSR-corrected transverse velocity of $(274\\pm12)\\; d_2$~km~s$^{-1}$.  At a position angle of 249\\arcdeg, the proper motion vector is well matched to the long symmetry axis of the X-ray PWN \\citep{moon04}.  The measurement of the proper motion allows a revised estimate for the age of the pulsar.  The LSR-corrected proper motion of the pulsar is shown in Figure~\\ref{fig:ctb80}.  In projection, its closest approach to the expansion center estimated by \\citet{Koo} is $\\sim$330\\arcsec, a 5.5 $\\sigma$ deviation assuming a 1\\arcmin\\ uncertainty in the reported expansion center.  A geometric center at coordinates 19h54m50s, 33\\arcdeg00\\arcmin30\\arcsec\\ (J2000.0) was estimated by superposing a circle on the remnant shell seen in the left panel of Figure~\\ref{fig:ctb80}.  The pulsar's closest approach to this point implies an age $t_p \\sim 51$~kyr with a minimum separation of 74\\arcsec, consistent with the 56\\arcsec\\ positional uncertainty based solely on the proper motion uncertainty.  Our measurement of the pulsar motion is therefore not consistent with the $\\sim$77~kyr age and the expansion center estimated by \\citet{Koo} for the CTB~80 SNR, but it is consistent with the pulsar age determined by \\citet{mig02}. We note that the true explosion center may be masked by asymmetric expansion due, for example, to density gradients and structure in the ISM, and neither estimate for the age may be correct. However, the true age is almost certainly less than $\\tau_c \\sim 107$~kyr, since the pulsar reaches the far edge of the SNR when its measured proper motion is projected backwards over that time interval, as shown in Figure~\\ref{fig:ctb80}.  As B1951+32 moves through its ambient medium, its relativistic wind outflow is confined by ram pressure, and as seen in some other PWNe, synchrotron emission from the confined pulsar wind produces a cometary X-ray nebula. On the other side of the contact discontinuity, the shocked ISM produces H$\\alpha$ emission \\citep[see, e.g., the illustration by][]{gvc+04}. In such a picture, the radio-bright arc would be produced by the confined pulsar wind at the location of the contact discontinuity.  As we have shown here, the structure of the bright arc is dynamic, changing from epoch to epoch, but overall, it moves with the pulsar at an approximately steady separation $\\sim$2.5\\arcsec, corresponding to a standoff distance of $\\sim$0.024$\\;d_2$~pc.  \\subsection{Two Similar Pulsars in Very Different Environments}  Currently, over 1600 radio pulsars have measured values for $P$ and $\\dot{P}$, as listed in the ATNF pulsar catalog\\footnote{See http://www.atnf.csiro.au/research/pulsar/psrcat} \\citep{mhth05}.  In this large ensemble, B1757$-$24 and B1951$+$32 are among the fifty most energetic pulsars, with $\\dot{E} = 10^{36.4}$ and $10^{36.6}$~erg~s$^{-1}$, respectively. They are also among the hundred youngest known pulsars, with characteristic ages $\\tau_c = 15.5$ and $107$~kyr, respectively.  On one hand, the characteristic age of B1951$+$32 is likely to be an overestimate, since its current spin period is not much greater than the typical birth spin periods of pulsars \\citep{fk06}, and its kinematic age $t_p \\sim 51$~kyr, assuming birth near the center of CTB~80.  On the other hand, the upper limits on its westward proper motion suggest that B1757$-$24 may be substantially older than its characteristic age.  Although the true ages of the pulsars differ in relation to their respective characteristic ages, these two pulsars are quite similar in energetics and age.  Both pulsars also exhibit cometary PWNe.  B1757$-$24 is associated with a $\\sim10$\\arcsec\\ elongated nebula visible at both radio (as in Fig.~\\ref{fig:duck}) and X-ray wavelengths \\citep{kggl01}.  B1951+32 shows a cometary PWN in X-rays \\citep{moon04}, but its radio structure resembles a bow shock and an extended, limb-brightened bubble, and it is enclosed by a complex nebula visible in H$\\alpha$ \\citep{hester00}, with extended lobes and a bow shock structure. \\citet{hk88} suggest that the pulsar is interacting with the wall of its evolved SNR, and we find that the radio shock structure changes with time while moving with the pulsar.  On the other hand, B1757$-$24 appears to be outside any parent remnant. If indeed it is associated with G5.4$-$1.2, it would have interacted with the wall of the remnant in the past (as B1951+32 is at present), possibly re-energising the shell and leaving behind the structure we see as the Duck.  Alternatively, the parent remnant of B1757$-$24 is no longer visible, and we are witnesses to a cosmic coincidence."
    },
    "0801/0801.0134_arXiv.txt": {
        "abstract": "This paper reviews the rich corpus of observational evidence for tidal effects, mostly based on photometric and radial-velocity measurements. This is done in a period when the study of binaries is being revolutionized by large-scaled photometric surveys that are detecting many thousands of new binaries and tens of extrasolar planets.  We begin by examining the short-term effects, such as ellipsoidal variability and apsidal motion.  We next turn to the long-term effects, of which circularization was studied the most: a transition period between circular and eccentric orbits has been derived for eight coeval samples of binaries.  The study of synchronization and spin-orbit alignment is less advanced. As binaries are supposed to reach synchronization before circularization, one can expect finding eccentric binaries in pseudo-synchronization state, the evidence for which is reviewed. We also discuss synchronization in PMS and young stars, and compare the emerging timescale with the circularization timescale.  We next examine the tidal interaction in close binaries that are orbited by a third distant companion, and review the effect of pumping the binary eccentricity by the third star. We elaborate on the impact of the pumped eccentricity on the tidal evolution of close binaries residing in triple systems, which may shrink the binary separation.  Finally we consider the extrasolar planets and the observational evidence for tidal interaction with their parent stars. This includes a mechanism that can induce radial drift of short-period planets, either inward or outward, depending on the planetary radial position relative to the corotation radius.  Another effect is the circularization of planetary orbits, the evidence for which can be found in eccentricity-versus-period plot of the planets already known.  Whenever possible, the paper attempts to address the possible confrontation between theory and observations, and to point out noteworthy cases and observations that can be performed in the future and may shed some light on the key questions that remain open. ",
        "introduction": "Stars in close binary systems are subject to mutual tidal forces that distort their stellar shape, breaking their spherical and axial symmetry, which leads to different observational effects.  This paper reviews the rich corpus of observational evidence for the tidal effects, which has been accumulated mostly by photometric and radial-velocity studies of short-period binaries.  One can divide the observational effects of the stellar distortion caused by tidal interaction into three classes:  \\begin{itemize}  \\item Direct observations of the distorted shape of the two components of the binary system. One example is the ellipsoidal periodic modulation of the photometric intensity of the binary with the orbital period, caused by the rotation of the stars with the ellipsoidal shape.  \\item Deviation of the orbital motion of the binary from pure Keplerian orbit. This is caused by departures of the gravitational attraction from the inverse-square law, due to the distortion of the two stars from spherical symmetry. One example is the famous apsidal motion observed in many eccentric eclipsing binaries.  \\item Long timescale evolution of the binary orbital elements and the stellar rotation, induced by the tidal interaction.  The long-term interaction will take place as long as the tides raised by the mutual interaction vary in size, shape and location on the surface of the two components. The tides do not vary only when  \\begin{itemize}  \\item the binary motion is circularized,  \\item the stellar rotation is synchronized with the binary orbital motion, and  \\item the stellar rotation axis is aligned with the normal to the orbital plane of motion.  \\end{itemize} Observational evidence for this long-term evolution can be found, for example, through the circular orbits of short-period binaries.  \\end{itemize}  The timescales of the three classes are substantially different. The ellipsoidal effect varies with the orbital period, which is of the order of a few days. The timescale of the deviations of the binary motion from the Keplerian orbit is in the range between tens and thousands of years, whereas the last class timescale is of the order of millions or billions of years. Therefore, we can witness close binaries for which the effects of the first and the second class are in action, while observation of the third class of effects can be made only retroactively, in recognition of the results of a long timescale processes.  Observing one or more of these effects allows us to learn about the response of each star to the tidal forces exerted by its companion. This response depends on the stellar internal structure, and therefore observations of these effects can provide us with a unique opportunity to study that structure in a manner which is not possible with single stars.  In the last few decades numerous observations of the tidal interaction in close binaries have been accumulated.  In many cases, observations of binaries performed for other reasons could not have been accounted for without the tidal effects. One example is the lightcurves of eclipsing binaries, which are usually observed in order to derive the geometrical elements of the systems. However, the lightcurves of these close systems can not be fully understood without taking into account the distortion of the two stars. In other cases, special observational effort has been made in order to detect some evidence for tidal effects. One example is the search for the circularization cutoff period in samples of coeval binaries. Both types of cases give us an opportunity to confront the tidal theory, based on stellar structure and evolution, with the observations. This paper will review both types of observations.  Sections 2 \\& 3 review the observational evidence for the ellipsoidal modulation and the apsidal motion of eccentric binaries. Section 4, 5, and 6 discuss the evidence for long-term circularization, synchronization and alignment of short-period binaries. Section 7 covers the tidal interaction in triple systems, when the third distant companion can have an impact on the tidal evolution of the close binary.  The study of tidal interaction between the recently discovered extrasolar planets and their parent stars has gained rapid momentum since 1995, following the discovery of the first confirmed extrasolar planet. Section 8 reviews some aspects of this fascinating new field.  This paper has been written at a juncture when the study of close binaries has been undergoing a transitional phase, the result of large-scaled photometric surveys which are starting to produce large samples of eclipsing binaries. In acknowledgment of the new era, Section 9 reviews some of the recent results from two of these surveys. Finally, Section 10 summarizes the previous sections and discusses some aspects of future studies of tidal interaction in close binaries.  This paper is too short and my knowledge too incomplete to cover all the studies in the different aspects of the subject. I made a great effort, though, to cover in each of the sections at least the open questions in that subfield. I apologize to those whose work was not included in this review.  I have made an effort to keep the notation of each section consistent with the other sections of this review. This forced me to adopt some notations that are not so common in the different fields covered by this paper. Each section was written as a separate piece of study, so many definitions were repeated throughout this paper.  Finally, I hope that this review will be of some use to young researchers in the field. It would fulfill my highest expectations if some research is initiated as a result of reading this review and considering some of questions that it raises. ",
        "conclusions": "This paper has attempted to review a large corpus of observational evidence for tidal interaction in close binary systems, including short-period extrasolar planets. We first discussed two tidal effects that can be observed directly --- the ellipsoidal variation and the apsidal motion. Unfortunately, these two effects have been detected in the past in only a small number of stars, because the detection required special photometric observations which involved a considerable effort. The small number of known binaries with detected ellipsoidal modulation or apsidal motion allowed for only a limited confrontation between the theory and the observations. To perform a comprehensive confrontation we need many more binaries with accurate measurements of these two effects.  Ellipsoidal modulation was mostly used in the past to study known spectroscopic binaries (e.g., HD 149420 --- Fekel \\& Henry 2006)\\nocite{fekel2006}, or binaries with compact objects.  This situation is being changed dramatically, as ground-based large-scaled photometric surveys have been operated in the last decade. We already discussed two such 1-m class exhaustive projects, MACHO (Alcock \\etal\\ 1993\\nocite{alcock1993}) and OGLE (e.g., Udalski \\etal\\ 1997\\nocite{udalski1997}).  Furthermore, a new wave of observational effort to search for transiting planets by small, 10-cm class, telescopes is now at its full thrust (e.g., HATnet --- Bakos \\etal\\ \\nocite{bakos2004}2004; XO --- McCullough \\etal\\ \\nocite{mccullough2005}2005; WASP --- Pollacco \\etal\\ \\nocite{pollacco2006}2006; TrES --- Alonso \\etal\\ \\nocite{alonso2007}2007). These projects are collecting an unprecedented number of photometric datasets that harbor thousands of ellipsoidal binaries.  The amplitude of the ellipsoidal modulation can be of the order of a few percent, and therefore should be easily detected in these large photometric datasets.  Detecting ellipsoidal variables can become a tool for discovering new binaries. In such cases, one must distinguish between the ellipsoidal modulation and other possible variations, particularly stellar spots and pulsation. For that purpose, one must rely on the theoretical shape of the modulation and on observations in different passbands. A first step towards this goal has been performed already by Derekas \\etal\\ (2006\\nocite{derekas2006}), who identified ellipsoidal variables in the sample of variable pulsating red giants within the MACHO data.  The huge datasets are having another effect on the study of close binaries: they are already revealing a large number of new eclipsing binaries (e.g., Devor 2005). Therefore the number of known eclipsing binaries will increase in the next few years by one or two orders of magnitude. The tens of thousands of new eclipsing binaries will revolutionize our understanding of close binaries in general and their statistical features in particular. These datasets will reveal many other effects of eclipsing binaries, such as apsidal motion and third distant companions through the time-travel effect. We will then be well equipped to confront the theory of apsidal motion with observations, and learn more about the frequency and distribution of triple systems. We will be able to verify conjectures regarding the correlation between short-period binaries and distant companions.  The influx of photometric data is about to increase, as more automated photometric surveys with 1-m class telescopes will set into action. This includes SkyMaper (Bayliss \\& Sackett 2007\\nocite{bayliss2007}), Las Cumbres Observatory Global Telescope (Brown \\etal\\ \\nocite{brown2007}2007), Pan-STARRS (e.g., Holman \\etal\\ \\nocite{holman2007}2007), and finally LSST (e.g., Strauss \\etal\\ 2006\\nocite{strauss2006}). One such binary was discovered already by Blake \\etal\\ (\\nocite{blake2007}2007), using the Sloan Digital Sky Survey II.  This influx of photometric data will be complemented by extremely accurate datasets from space missions, which will have an immediate impact on binary studies. This trend has already been started by HIPPARCOS (e.g., Jorissen \\etal\\ \\nocite{jorrisen2004}2004), and will be joined by CoRoT (Maceroni and Ribas \\nocite{maceroni2007}2007), Kepler (Koch \\etal\\ \\nocite{koch2007}2007) and GAIA (e.g., Niarchos 2006\\nocite{niarchos2006}) missions. For the future of binary studies in the 'brave new era,' see a review by Guinan \\& Engle (\\nocite{guinan2006}2006).  We have discussed the long-term tidal processes of circularization, synchronization and spin alignment. As was reviewed, binary circularization has attracted many studies, particularly of the transition period between circular and eccentric binaries. Nevertheless, no consensus has been reached with regard to the theoretical timescale of the tidal processes, and even the dominant dissipation process behind the circularization is still being debated.  Synchronization and alignment were studied much less. For example, the number of eclipsing binaries with observed RM effect by which we study the spin-orbit alignment is only six, an extremely small number. Modern spectrographs with efficient digital detectors can observe many more eclipsing binaries, and the development of tools to analyze the observations can enable us to accurately derive the spin-orbit inclination. The results may well be crucial for our understanding of binary formation, as this will help us estimate the primordial alignment of binary population. In case binaries are formed with some spread of spin-orbit inclinations, we should also expect coeval samples of binaries to exhibit a transition period between aligned and non-aligned binaries.  Such an observed transition can serve as another confirmation of the theory of tidal interaction.  It would be of extreme importance to compare the aligned/non-aligned transition period with the pseudo-synchronized/non-pseudo-synchronized and circularized/eccentric transition periods.  An additional feature of the new large accurate photometric datasets is their long time span, of the order of a few years, that will enable us to detect minute variations which disclose long-term variations. The search of Michalska \\& Pigulski (2005) and Michalska (2007) for evidence of apsidal motion discussed above is one example.  The works of Kubiak \\etal\\ (\\nocite{kubiak2006}2006) and Pilecki \\etal\\ (\\nocite{pilecki2007}2007) are two other examples of studies that utilized the long time span of the new data sets. Some of the yet-to-come studies will be able to find eventually direct evidence for long-term tidal processes in close binaries.  Finally, we have discussed in brief some ideas about tidal evolution of short-period extrasolar planets. Many of the short-period planets are being discovered by photometry which detects how these planets eclipse their parent stars. The transiting planets hold a higher potential for the study of tidal effects, as we know so much more about them, including their masses and radii. They can be used as clean and unique laboratories to study tidal effects, as the complexity of their systems is reduced because of the small masses and radii of the planets. As of October 2007, 27 transiting planets are known, which is about half of the short-period extrasolar planets, the others have been discovered by radial-velocity measurements. We anticipate that the balance between the radial-velocity and transiting planets will change in the near future in favour of the transiting planets.  The rapidly increasing number of transiting planets will enable us to study the possible correlation between presumed tidal effects and the stopping mechanism, the inflated radii of some planets, synchronization of the stars and planets and their spin-orbit alignment. A comprehension of these features will help us assess the role of tidal interaction in the orbital evolution of short-period extrasolar planets."
    },
    "0801/0801.3889_arXiv.txt": {
        "abstract": "We derive an analytic model for the redshift evolution of linear-bias, allowing for interactions and merging of the mass-tracers, by solving a second order differential equation based on linear perturbation theory and the Friedmann-Lemaitre solutions of the cosmological field equations. We then study the halo-mass dependence of the bias evolution, using the dark matter halo distribution in a $\\Lambda$CDM simulation in order to calibrate the free parameters of the model. Finally, we compare our theoretical predictions with available observational data and find a good agreement. In particular, we find that the bias of optical QSO's evolve differently than those selected in X-rays and that their corresponding typical dark matter halo mass is $\\sim 10^{13} \\; h^{-1} \\; M_{\\odot}$ and $\\magcir 5 \\times 10^{13} \\; h^{-1} \\; M_{\\odot}$, respectively. ",
        "introduction": "The distribution of matter on large scales, based on different extragalactic objects, can provide important constraints on models of cosmic structure formation. However, a serious problem that hampers such a straight forward approach is our limited knowledge of how luminous matter traces the underlying mass distribution. It is well known that the large scale clustering pattern of different extragalactic objects (galaxies, AGN, clusters, etc) trace the underlying dark matter distribution in a biased manner (Kaiser 1984; Bardeen et al. 1986). Such a biasing is assumed to be statistical in nature; with galaxies and clusters being identified as high peaks of an underlying, initially Gaussian, random density field. Furthermore, the biasing of galaxies with respect to the dark matter distribution was also found to be a necessary ingredient of Cold Dark Matter (CDM) models of galaxy formation in order to reproduce the observed galaxy distribution (eg. Davis et al. 1985; Cole \\& Kaiser 1989; Benson et al. 2000).  Furthermore, the bias redshift evolution is very important in order to relate observations with models of structure formation and it has been shown that the bias factor, $b(z)$, is a monotonically increasing function of redshift. Indeed, Mo \\& White (1996) and Matarrese et al. (1997) have developed a model for the evolution of the correlation bias, defined as the ratio of the halo to the mass correlation function, the so called galaxy merging bias model. This model takes into account, via the Press-Schechter formalism, the collapse of different mass halos at the different epochs. In this framework and at high $z$'s one finds, in the EdS universe, that $b(z)-1 \\propto (1+z)$  (see also Sheth et al. 2001). According to Matarrese et al (1997), if the only halos that exist at any epoch are those that have just formed via a merging process of smaller halos then the bias evolution, in the high density peak limit, becomes $b(z) \\propto (1+z)^{2}$ (see also Bagla 1998). There are many studies of the merging bias model in the context of a $\\Lambda$ cosmology (e.g. Jing 1998; Sheth \\& Tormen 1999; Sheth, Mo \\& Tormen 2001), while the bias has been found to be a key ingredient of the halo model (e.g. Peacock \\& Smith 2000; Seljak 2000; Ma \\& Fry 2000).  The so-called galaxy conserving bias model (Nusser \\& Davis 1994; Fry 1996; Tegmark \\& Peebles 1998; Hui \\& Parfey 2007) assumes only that the different mass-tracer (acting as ``test particles'') fluctuation field is proportional to that of the underlying mass and it predicts a bias evolution according to: $b(z)=1+(b_{0}-1)(1+z)$ for $\\Omega_{\\rm m}=1$, where $b_{0}$ is the bias factor at the present time. Coles, Melott \\& Munshi (1999) have developed a bias model within the hierarchical clustering scenario which gives interesting scaling relations for the galaxy bias. Another approach was proposed by Basilakos \\& Plionis (2001; 2003), in which the linear bias evolution was described via the solution  of a second order differential equation, derived using linear perturbation theory and the basic cosmological equations. Also, in this model the mass tracer population is assumed to be conserved in time.   In this paper we extend the original Basilakos \\& Plionis linear and scale-independent bias evolution model to include the effects of an evolving mass-tracer population, possibly due to interactions and merging. We then derive the dependence of the linear-bias evolution on halo mass, with the help of N-body simulations, in the framework of the concordance $\\Lambda$CDM cosmological model. Additionally we compare our model with observations. ",
        "conclusions": "In this work we extend our original Basilakos \\& Plionis (2001) bias evolution model, based on  linear perturbation theory and the Friedmann-Lemaitre solutions of the cosmological field equations, to take into account the redshift evolution of an interaction term that could affect the mass-tracer number density. Parameterizing our analytical model with N-body simulations of the concordance $\\Lambda$CDM cosmological model ($\\Omega_{\\rm m}=1-\\Omega_{\\Lambda}=0.3$), we investigate the halo-mass dependence of our linear bias evolution model, with or without the interaction term.  Assuming that the extragalactic mass tracers are hosted by a dark matter halo of given mass, we can predict analytically their present bias parameter and its $z$-evolution. Inversely, fitting our bias evolution model to the observed bias of different extragalactic mass-tracers, we can identify the typical mass of the DM halo within which they live.  Comparing our theoretical predictions with the observed bias of different mass-tracers (galaxies, AGNs, DRGs, EROs, etc) we find a very good agreement with our bias evolution functional form and therefore we also identify the typical mass of the DM halo which they inhabit. This has allowed us to infer that X-ray and optically selected AGNs inhabit DM halos of different mass (with X-ray AGNs being associated with higher DM halo masses)."
    },
    "0801/0801.4116_arXiv.txt": {
        "abstract": "We present new high resolution ($R \\simeq 18,000$) near-infrared spectroscopic observations of a sample of classical FU~Orionis stars (FUors) and other young stars with FUor characteristics that are sources of Herbig-Haro flows.  Spectra are presented for the region $\\lambda =$ 2.203 -- 2.236~$\\mu$m which is rich in absorption lines sensitive to both effective temperatures and surface gravities of stars.  Both FUors and FUor-like stars show numerous broad and weak unidentified spectral features in this region.  Spectra of the 2.280 -- 2.300~$\\mu$m region are also presented, with the 2.2935~$\\mu$m $v$=2--0~CO absorption bandhead being clearly the strongest feature seen in the spectra all FUors and Fuor-like stars.  A cross-correlation analysis shows that FUor and FUor-like spectra in the 2.203 -- 2.236~$\\mu$m region are not consistent with late-type dwarfs, giants, nor embedded protostars.  The cross-correlations also show that the observed FUor-like Herbig-Haro energy sources have spectra that are substantively similar to those of FUors.  Both object groups also have similar near-infrared colors.  The large line widths and double-peaked nature of the spectra of the FUor-like stars are consistent with the established accretion disk model for FUors, also consistent with their near-infrared colors.  It appears that young stars with FUor-like characteristics may be more common than projected from the relatively few known classical FUors. ",
        "introduction": "FU~Orionis stars (FUors) are rare.  Classical FUors have been identified by large ($\\sim 5$ mag) increases in brightness and luminosity followed by fading over decades, and only a few young stars have been confirmed as such, notably FU Ori, V1057 Cyg, V1515 Cyg, V1735 Cyg \\citep{H77}, V346 Nor \\citep{GF85}, and V733 Cep \\citep{RABBCH07}.  Additionally, while in outburst, they exhibit optical and near-IR spectra similar to FU~Orionis itself \\citep{H77, HK96, HPD03}.  The absorption line widths (interpreted as rotational velocities) and derived spectral types of FUors change with wavelength, varying from F or G in the visible to M type in the near-IR. Their spectra also indicate supergiant or giant surface gravities.  These characteristics have been modeled as arising from young stars with massive accretion disks, the latter dominating the spectra and luminosities of these objects \\citep[see][]{HK96}.  The accretion disk model has been quite successful in explaining the basic visible-to-IR spectral features and energy distributions of FUors \\citep{HK85, KHH88, HHC04, GHCWIFSF06}, although in detail there are still discrepancies \\citep{HPD03}.  The only published (and modeled) high resolution near-IR spectra of FUors have been in the vicinity of the first vibrational overtone $v = 0 - 2$ of CO ($\\lambda \\gtrsim 2.294 \\mu$m), providing temperature and rotation information.  New observations at other near-IR wavelengths may aid in determining the range of effective temperatures and surface gravities over which FUor spectral features arise in the near-IR, providing more constraints on applicable physical models.  The general symmetry and relatively even spacing often seen in Herbig-Haro object knots strongly suggest that they are somehow linked to episodic outbursts from their parent star plus disk systems \\citep{D78,R89}.  This suggests that the sources of HH objects and FUor outbursts may possibly be related.  \\citet{RA97} found that five young stars associated with HH flows (termed Herbig-Haro energy sources or HHENs) had low resolution K-band spectra that are very similar to those of FUors.  If episodic outbursts do generally indicate FUor activity, then there may be many more young stars with FUor characteristics than the number predicted by extrapolating numbers from the few known classical FUors.  In the following discussions, we use the term FUor-like star to indicate objects that have spectral similarities to the classical FUors, but for which no eruption has been witnessed \\citep{RHKSB02}.  How similar are classical FUors and FUor-like stars?  Both are young stars that are likely undergoing episodes of high accretion.  However, their luminosities may be powered by different physical processes or they may be at different phases of their outburst and decay cycles. It is also important to understand how similar these objects are to embedded (Class~I) protostars accreting at fairly high rates (e.g., $\\dot{M} \\sim 10^{-6} M_{\\odot}$ yr$^{-1}$) and to investigate whether there are any spectral similarities between protostars, FUors, and FUor-like objects.  Also, do any of these objects show spectral characteristics similar to other young eruptive variables such as EX~Lupi-type (EXors) young stars?  EXors undergo multiple short duration optical outbursts of 1 -- 5~mag but have K or M dwarf optical spectra dominated by emission lines, much different from FUor spectra \\citep[e.g., see][]{HAGI01}.  We believe that such fundamental questions can be addressed with new high resolution spectra.  Such spectra must cover near-IR wavelengths since many FUor-like stars are highly extinguished and cannot be observed in visible light.  The low resolution near-IR spectra of \\citet{RA97} did show strong similarities between FUors and FUor-like objects, but those spectra were also consistent with giant stars.  One object with such a near-IR spectrum in close proximity to a young star was recently found to be a background giant star \\citep{AG07}. Linking other young stars to FUors will improve our knowledge of the number and types of young stars with high accretion rates, and detailed modeling may reveal more information on physical structures and accretion mechanisms or processes.  We have conducted a new near-IR spectroscopic study of FUors and FUor-like stars covering a relatively broad range of wavelengths near 2~$\\mu$m, including spectral features sensitive to both effective temperature and surface gravity.  We present these new data in $\\S 2.$ The properties of the individual FUor-like stars are discussed in $\\S 3$.  In $\\S 4$ we analyze the similarity of the spectra of FUor-like stars to those of {\\it i)} classical FUors, {\\it ii)} late-type spectral standards, and {\\it iii)} Class~I protostars.  We discuss the likely physical similarities of these stars and the possible origins of their attributes in $\\S 5$. ",
        "conclusions": "The similarity of the FUor and FUor-like spectra and their significant cross-correlations indicate that the spectral lines of these objects must form in physically similar environments.  \\citet{RA97} had also postulated this based on similarities of low resolution near-IR spectra, and the high resolution data presented in this paper confirm that the FUor-like HHENs are more similar to FUors than dwarfs, giants, supergiants, or embedded protostars.  Given the cross-correlation results, the near-IR spectra of FUor-like stars do not appear to be produced in normal stellar photospheres. The broad cross-correlation peaks indicate rotational velocities $v$~sin~$i \\sim 100$~km~s$^{-1}$, much larger than measured for T~Tauri stars or even the most rapidly spinning embedded protostars \\citep{DGCL05, CGDL06}.  The cross-correlations with late-type dwarfs and the embedded protostar IRS~63 are also generally poor (low significance), but they are somewhat better with late-type giants and supergiants (Table~\\ref{tbl-2}).  This suggests that the spectral features of FUor-like stars originate in cool regions with low surface gravities that differ from normal stellar photospheres, and the generally good cross-correlations with FUor spectra reinforce this. The strong CO absorptions of the FUors and FUor-like stars (Figures~\\ref{fig1} and \\ref{fig2}) indicate relatively low effective temperatures, T$_{eff} \\sim 3500$~K \\citep[e.g.,][]{KH86}.  A giant or supergiant star of that effective temperature would have a radius greater than about 40 $R_{\\odot}$ and a mass of approximately 6 $M_{\\odot}$ or more.  Such stars have rotational breakup velocities on the order of 100 km~s$^{-1}$.  Therefore, the CO strengths and projected rotation velocities of FUor and FUor-like spectra (from cross-correlation and CO bandhead widths; Fig.~\\ref{fig2} and \\ref{fig3}) are consistent with giants rotating at breakup.  However, the marginal cross-correlation significance with giants or supergiants (Table~\\ref{tbl-2}) indicates that the FUors and FUor-like stars are not likely to be ordinary giants or supergiants.   The combined visible and IR spectra of FUors are reasonably well fit by models of line formation in active circumstellar accretion disks \\citep[e.g.,][]{KHH88, HHC04}; the decrease in their line widths (and implied rotational velocities) at longer wavelengths is recognized as perhaps the best evidence of line formation in disks.  The double-peaked nature of the cross-correlation of FUors and FUor-like near-IR spectra with those of late type dwarfs (Figure~\\ref{fig3}) is also consistent with a disk rotational velocity profile, and the significant near-IR excesses exhibited by most of these objects (e.g., $\\S 4.3$ and Fig.~\\ref{fig4}) are also consistent with hot circumstellar disk material.  If the spectra of FUors and FUor-like HHENs are indeed produced in accretion disks, then why are they not more similar to the spectra of embedded protostars with active disks?  Even Class~I protostars with large veiling and relatively high accretion rates $\\dot{M} \\simeq 10^{-6} M_{\\odot}$~yr$^{-1}$ do not have FUor-like spectra \\citep[see also][]{GL02, DGCL05}.  In a circumstellar disk model atmosphere, high accretion rates are needed to induce a temperature gradient with a hot midplane and cooler exterior to produce a spectrum similar to a FUor \\citep[e.g.,][]{KHH88}.  \\citet{CHS91} found that an irradiated accretion disk model with mass accretion rate of $1.6 \\times 10^{-4} M_{\\odot}$ yr$^{-1}$ reproduced the near-IR spectrum of FU~Ori well, including its CO and H$_{2}$O absorptions.  This accretion rate is approximately 2 orders of magnitude higher than that of embedded low mass protostars such as IRS 63.  \\citet{GL02} estimated that the mass accretion rate of the embedded protostars YLW 15A was $2.3 \\times 10^{-6} M_{\\odot}$ yr$^{-1}$, and the same modeling would produce a very similar (within $\\sim50$\\%) value for IRS 63 since the two protostars have virtually identical effective temperatures, near-IR veilings, rotational velocities \\citep{DGCL05} and IRAS fluxes.  It appears that higher accretion rates are needed to produce luminous disks with absorption features similar to those found in FUors. Future discoveries and observations of weaker FUors with lower accretion rates and more active protostars with higher accretion rates would be valuable in constraining further the accretion rate at which a circumstellar disk develops a strong enough temperature gradient to produce absorption lines that dominate those of the stellar photosphere.   The very high degree of similarity of the FUor-like and FUor spectra (and their good cross-correlations) suggest that FUor activity may be more common than indicated from the small number of known classical FUors alone; more young stars have spectra likely to originate in luminous accretion disks.  Taking only the list of 20 classical and FUor-like objects by \\citet{AKCKMP04} and adding the two other HH sources found to have FUor-like spectra by \\citet{RA97} (HH~381~IRS and HH~354~IRS; also confirmed in this work) yields a sample of 22 objects.  Of these, five are widely acknowledged to be bona fide classical FUors due to their observed outbursts and spectral properties (FU~Ori, V1057 Cyg, V1515 Cyg, V1735 Cyg, and V346 Nor). Therefore FUor-like stars may outnumber FUors.  However, FUors are certainly not ubiquitous.  Very few FUors or stars with FUor-like spectra are found in regions of clustered star formation such as Tau-Aur, $\\rho$~Oph, NGC~1333, IC~346, or the Orion Nebula Cluster.  Instead, they are mostly found in regions of low star formation activity.  This implies that either FUors are much less frequent or have many fewer outburst cycles in high density regions.  One possible explanation is that interactions between stars in clusters could disrupt disks, perhaps eliminating much FUor activity during the Class 0 or Class I protostellar evolutionary phases.  Alternatively, FUor events may be triggered by the presence of a companion star \\citep{BB92}, in which case a dense cluster environment might accelerate the orbital evolution, causing FUor eruptions in clusters to occur mostly during the Class 0 or Class I stages \\citep{RA04}.  However, these are certainly not definitive explanations, and the causes of FUor events must be understood better before the inhomogeneity in the distribution of FUors and FUor-like stars can be understood."
    },
    "0801/0801.1505_arXiv.txt": {
        "abstract": "A  long  standing problem  for  hierarchical  disk galaxy  formation models has been  the simultaneous matching of the  zero point of the Tully-Fisher relation  and the galaxy luminosity  function (LF).  We illustrate this problem for a  typical disk galaxy and discuss three solutions: low  stellar mass-to-light ratios, low  initial dark halo concentrations,  and no  halo contraction.   We speculate  that halo contraction may  be reversed through a combination  of mass ejection through  feedback and  angular  momentum exchange  brought about  by dynamical friction  between baryons and dark matter  during the disk formation process. ",
        "introduction": "The  relation between  the rotation  velocity and  luminosity  of disk galaxies, commonly  known as the Tully-Fisher (TF)  relation (Tully \\& Fisher  1977), is  one  of  the most  fundamental  properties of  disk galaxies.  This  is because it is  a link between  luminous mass (i.e. baryons) and dynamical  mass (baryons and dark matter)  and because it has    very   little    intrinsic   scatter    $\\sigma_{\\ln   V}\\simeq 0.1$. Furthermore,  unlike the  Faber-Jackson relation for  early type galaxies,  the scatter  in the  TF  relation is  {\\it independent}  of surface brightness, size, or offsets from the size-luminosity relation (Courteau \\& Rix 1999; Courteau \\etal 2007).  CDM  based disk  galaxy formation  models  are able  to reproduce  the slope, and  in some cases the  amount of scatter, in  the TF relation. However, a  long standing  problem has been  the matching of  the zero point  of  the TF  relation,  with  the  generic problem  of  galaxies rotating  too fast  at  a given  luminosity.   This has  been seen  in analytic  models (van den  Bosch 2000;  Mo \\&  Mao 2000;  Dutton \\etal 2007),   semi-analytic  models   (e.g.   Benson   \\etal  2003   )  and cosmological  N-body  simulations (e.g.   Eke,  Navarro, \\&  Steinmetz 2001) ",
        "conclusions": ""
    },
    "0801/0801.4927_arXiv.txt": {
        "abstract": "The concordance Cold Dark Matter model for the formation of structure in the Universe, while remarkably successful at describing observations on large scales, has a number of problems on galaxy scales.  The Milky Way and its satellite system provide a key laboratory for exploring dark matter (DM) in this regime, but some of the most definitive tests of local DM await microarcsecond astrometry, such as will be delivered by the Space Interferometry Mission (SIM Planetquest). I discuss several tests of Galactic DM enabled by future microarcsecond astrometry. ",
        "introduction": "Since the seminal study of Searle \\& Zinn (1978) the notion of accretion of ``subgalactic units\", including ``late infall\", has been a central concept of stellar populations studies. N-body simulations of the formation of structure in the Universe in the presence of dark matter (and dark energy) also show galaxies (and all large structures) building up hierarchically. However, while the active merging history on all size scales demonstrated by high resolution, Cold Dark Matter (CDM) numerical simulations have had remarkable success in matching the observed properties of the largest structures in the Universe (like galaxy clusters), they are a challenge to reconcile with the observed properties of structures on galactic scales.  The Milky Way (MW) and its satellite system are a particularly important laboratory for testing specific predictions of the CDM models.  The era of microarcsecond astrometry such as will be delivered by SIM Planetquest, will enable a number of definitive dynamical tests of local CDM by way of determining the distribution of Galactic dark matter (DM). We discuss several of these tests here. ",
        "conclusions": ""
    },
    "0801/0801.4861_arXiv.txt": {
        "abstract": "We present numerical simulations of cold, axisymmetric, magnetically driven relativistic outflows. The outflows are initially sub-Alfv\\'enic and Poynting flux-dominated, with total--to--rest-mass energy flux ratio up to $\\mu \\sim 620$. To study the magnetic acceleration of jets we simulate flows confined within a funnel with rigid wall of prescribed shape, which we take to be $z\\propto r^a$ (in cylindrical coordinates, with $a$ ranging from 1 to 2). This allows us to eliminate the numerical dissipative effects induced by a free boundary with an ambient medium. We find that in all cases they converge to a steady state characterized by a spatially extended acceleration region. For the jet solutions the acceleration process is very efficient - on the outermost scale of the simulation more than half of the Poynting flux has been converted into kinetic energy flux, and the terminal Lorentz factor approached its maximum possible value ($\\Gamma_\\infty \\simeq \\mu$). The acceleration is accompanied by the collimation of magnetic field lines in excess of that dictated by the funnel shape. The numerical solutions are generally consistent with the semi-analytic self-similar jets solutions and the spatially extended acceleration observed in some astrophysical relativistic jets. In agreement with previous studies we  also find that the acceleration is significantly less effective for wind solutions suggesting that pulsar winds may remain Poynting dominated when they reach the termination shock. ",
        "introduction": "\\label{introduction}  There is strong evidence for relativistic motions in jets that emanate from active galactic nuclei (AGNs). The Lorentz factors of blazar jets lie in the range $\\sim 5-40$ \\cite{J01,J05,C07}.   In the case of AGNs there have indeed been indications from a growing body of data that the associated relativistic jets undergo the bulk of their acceleration on scales that are of the order of those probed by very-long-baseline radio interferometry. In particular, the absence of bulk-Comptonization spectral signatures in blazars has been used to infer that jet Lorentz factors $\\geqslant 10$ are only attained on scales $\\geqslant 10^{17}\\ {\\rm cm}$ \\cite{S05}.  The first theoretical clues to the necessity of relativistic motion in GRB's arose from the compactness problem \\cite{rud75}. The requirements that the source be optically thin can be used to obtain direct limits on the minimal Lorentz factor, $\\Gamma \\gtrsim 100$ \\cite{kp91,pir99}. Recently, superluminal expansion was observed in the afterglow emission of GRB 030329 \\cite{tay04}.  The main source of power of AGN and GRB jets is the rotational energy of the central black hole \\cite{L76,BZ77} and/or its accretion disk. The naturally occurring low mass density and hence high magnetization of black-hole magnetospheres suggests that the relativistic jets originate directly from the black-hole ergosphere as Poynting-dominated outflows, whereas the disk surface launches a slower, possibly non-relativistic wind that surrounds and confines the highly relativistic flow. Although, dissipative processes may directly transfer the electromagnetic energy to emitting particles, the commonly held view is that it is first converted into the bulk kinetic energy and only subsequently channeled into radiation through shocks and other dissipative waves \\cite{BR74,BBR84,VK03a,VK03b,bn06, kbvk2007}. In the limit of ideal MHD such conversion can be achieved only via magnetic forces and the efficiency of this mechanism is the main subject of our investigation.  Since most of the acceleration takes place far away from the source the space-time is basically flat. We also use an isentropic equation of state $  p=Q\\rho^s$, where $Q=$const and $s=4/3$. Since we are interested in the magnetic acceleration of cold flows, we make $Q$ very small, so the gas pressure is never a dynamical factor. This relation enables us to exclude the energy equation from the integrated system and overcome the stiffness of relativistic MHD in magnetically dominated regime.  Given the condition of axisymmetry the poloidal magnetic field is fully described by the so-called magnetic flux function $\\Psi(r,z)$, the total magnetic flux enclosed by the axisymmetric loop circle $r,z=$const. Moreover, for stationary flows there are 5 quantities that propagate unchanged along the magnetic field lines and thus are functions of $\\Psi$ alone. These are $\\massint$, the rest-mass energy flux per unit magnetic flux; $\\Omega$, the angular velocity of magnetic field lines; $l$, the total angular momentum flux per unit rest-mass energy flux; $\\enint$, the total energy flux per unit rest-mass energy flux; and $Q$, the entropy per particle. For cold flows ($Q=0$, $w=\\rho c^2$) we have $  \\massint = {\\rho \\Gamma v_p}/{B_p} $, $  \\Omega r=v^{\\hat{\\phi}}-v_pB^{\\hat{\\phi}}/B_p $, $    l = -I/{2\\pi\\massint c}+r u^{\\hat\\phi}$, and $    \\enint = \\Gamma \\left(1+\\sigma\\right),$ where $\\Gamma$ is the Lorentz factor,  $v_p$ is the magnitude of the poloidal velocity, $B_p$ and $B^{\\hat{\\phi}}$ are the magnitudes of poloidal and azimuthal magnetic field, $ I = c r B^{\\hat\\phi}/2 $ is the total electric current flowing through a loop of cylindrical radius $r$, $\\sigma$ is the ratio of the Poynting flux to the matter (kinetic plus rest-mass) energy flux, and $ \\Gamma \\sigma=-\\Omega I/{2\\pi\\massint c^3}$ is the Poynting flux per unit rest-mass energy flux. It is easy to see that the Lorentz factor $\\Gamma$ cannot exceed $\\enint$.  \\begin{figure}[H] \\begin{center} {\\epsfig{figure=figures/barkov-fig1.eps,width=7.5cm,angle=-90}} \\end{center} \\caption{Evolution of $\\sigma$ along the magnetic field line $\\Psi=0.8\\Psi_{max}$ in models with $\\sigma = $ 620 (solid line), 310 (dashed line), 155 (dash-dotted line), 78 (dotted line) and 39 (dash-triple-dotted line). } \\label{sigma-ev-a1} \\end{figure} ",
        "conclusions": "\\label{conclusion}  In validating the basic features of the simplified semi-analytic solutions, our numerical results go a long way toward establishing an ``MHD acceleration and collimation paradigm'' for relativistic astrophysical jets. In particular, they demonstrate that even jets with extremely high initial magnetization can be effectively accelerated via the ideal magnetic mechanism with more than a half of the Poynting flux converted into the bulk motion energy of the flows. The highest Lorentz factor reached in the simulations is $\\Gamma=300$ which is well within the range deduced for GRB jets. The slow character of magnetic acceleration allows dissipationless energy transport over large distances, the property which is deduced from observations and which is rather difficult to explain in other models of jet generation.   Interested reader can find more details about the method of our simulations and results for AGN jets in \\cite{kbvk2007}."
    },
    "0801/0801.4050_arXiv.txt": {
        "abstract": "Normal field stars located behind dense clouds are a valuable resource in interstellar astrophysics, as they provide continua in which to study phenomena such as gas-phase and solid-state absorption features, interstellar extinction and polarization. This paper reports the results of a search for highly reddened stars behind the Taurus Dark Cloud complex. We use the Two Micron All Sky Survey (2MASS) Point Source Catalog to survey a $\\sim$\\,50~deg$^2$ area of the cloud to a limiting magnitude of $K_{\\rm s} = 10.0$. Photometry in the 1.2--2.2\\mic\\ passbands from 2MASS is combined with photometry at longer infrared wavelengths (3.6--12\\mic) from the {\\it Spitzer Space Telescope\\/} and the {\\it Infrared Astronomical Satellite\\/} to provide effective discrimination between reddened field stars and young stellar objects (YSOs) embedded in the cloud. Our final catalog contains 248~confirmed or probable background field stars, together with estimates of their total visual extinctions, which span the range $2<\\av<29$~mag. We also identify the 2MASS source J04292083+2742074 (IRAS~04262+2735) as a previously unrecognized candidate YSO, based on the presence of infrared emission greatly in excess of that predicted for a normal reddened photosphere at wavelengths $>5$\\mic. ",
        "introduction": "\\label{intro}  Observations of many interstellar phenomena rely on the presence of background field stars, i.e.\\ stars lying beyond the region of interest that provide sources of continuum radiation. Examples include absorption-line spectroscopy of ionic, atomic and molecular gas, spectroscopy of solid-state absorption features in dust grains, and the continuum effects of interstellar extinction and polarization introduced by dust, at wavelengths extending from the far ultraviolet to the mid-infrared. Observations at ultraviolet and visible wavelengths are generally limited to lines of sight with low to moderate extinction and therefore sample mostly diffuse phases of the interstellar medium (ISM). Observations in the near to mid-infrared allow studies of these absorptive phenomena to be extended to dense molecular clouds, where they provide information complementary to that obtained at much longer wavelengths were cool dust and gas can be observed in emission in their own right.  Observing absorption features in discrete lines of sight toward background stars places an obvious limitation on spatial resolution. Consider, for example, the distribution of molecular gas in cold, dense clouds, and the depletion of molecules onto the surfaces of dust grains to form icy mantles. Radio astronomers can readily map the distribution of molecular gas by observing its intensity in a chosen spectral emission line (most commonly of CO) at a spatial resolution limited only by the beam of the telescope (see Mizuno \\etal\\ 1995 for an example relevant to the current work). However, to map the corresponding distribution of molecular solids one must study an infrared absorption feature, such as the 3.0\\mic\\ feature of solid H$_2$O, against a continuum provided by a background star. The first attempt to map the ice distribution by this method, that of Murakawa \\etal\\ (2000), was limited to detections of the 3.0\\mic\\ feature in some~25 lines of sight in the Heiles Cloud~2 region of Taurus. At mid-infrared wavelengths, it is possible to study interstellar absorptions against extended background emission, as described by Sonnentrucker \\etal\\ (2008) for the 15\\mic\\ feature of CO$_2$; but this opportunity arises only in a limited number of cases, where cold material is located fortuitously in front of warmer material heated to appropriate temperatures by nearby luminous stars. In regions of low-mass star formation a point-like stellar continuum source is generally the only option. Interstellar absorption features are observed routinely in the spectra of young stellar objects (YSOs) embedded in the clouds (e.g.\\ Boogert \\etal\\ 2004), but in these lines of sight ambiguity often exists between interstellar matter and material local to the source that may be modified by its presence. Only background stars provide the means of studying absorptions arising in quiescent regions of the dense ISM, remote from sources of radiation.  A primary motivation for infrared surveys of dark clouds is to conduct a census of star formation, in which YSOs are carefully distinguished from field stars. The Taurus dark-cloud complex is a key region in studies of low-mass star formation (e.g.\\ Luhman \\etal\\ 2006; G\\\"udel \\etal\\ 2007; Padgett \\etal\\ 2007), and also a valuable laboratory for studying interstellar molecules and dust (e.g.\\ Pratap \\etal\\ 1997; Dickens \\etal\\ 2001; Whittet \\etal\\ 2001). It is nearby ($\\sim 140$\\,pc; Kenyon \\etal\\ 1994; Loinard \\etal\\ 2005) and situated at moderate Galactic latitude ($b \\sim 15^\\circ$), circumstances that aid both types of study: there is relatively little risk of confusion between stellar associations at different distances along the line of sight, compared with a region in the Galactic disk, and almost all of the interstellar extinction and absorption in the direction of the cloud arises in the cloud itself. The first infrared survey of the Taurus region to include a significant number of reddened background stars was that of Elias (1978), who identified about a dozen such objects that have since been observed intensively for interstellar absorption features (e.g.\\ Whittet \\etal\\ 2007 and references therein).  In recent years, the task of identifying reddened field stars has been greatly facilitated by the availability of data from the {\\it Two-Micron All-Sky Survey\\/} (2MASS; Skrutskie \\etal\\ 2006) and from surveys at longer infrared wavelengths made with the {\\it Spitzer Space Telescope\\/} (e.g.\\ Luhman \\etal\\ 2006). The goal of this paper is to use these resources to construct a much larger catalog of field stars reddened by dust in the Taurus dark-cloud complex than has previously been available. ",
        "conclusions": "\\label{conc} The main product of this work is the catalog presented in Table~1, which includes 248~confirmed or probable background field stars (and one probable new YSO). This catalog should prove to be a valuable resource for future observing programs. For example, virtually all the included sources are expected to show 3\\mic\\ \\water-ice absorption, based on the previously-known $\\av>3$ threshold in the correlation of ice optical depth with extinction (Whittet \\etal\\ 1988, 2001). The distribution of ice may thus be mapped in more detail than was possible in the previous study of Murakawa \\etal\\ (2000). The 3\\mic\\ feature also has potential for mapping magnetic fields within the clouds, by virtue of excess polarization observed at this wavelength when ice-mantled grains are aligned (Hough \\etal\\ 1988). This method is potentially more reliable than studies utilizing continuum polarization, as the ice feature is an unambiguous tracer of dense material in the line of sight (Whittet \\etal\\ 2008)."
    },
    "0801/0801.1675_arXiv.txt": {
        "abstract": "We compare recently measured rest-frame $V$-band luminosity functions (LFs) of galaxies at redshifts $2.0 \\leq z \\leq 3.3$ to predictions of semianalytic models by De Lucia \\& Blaizot and Bower et al. and hydrodynamic simulations by Dav\\'e et al. The models succeed for some luminosity and redshift ranges and fail for others. A notable success is that the Bower et al.\\ model provides a good match to the observed LF at $z\\sim 3$. However, all models predict an increase with time of the rest-frame $V$-band luminosity density, whereas the observations show a decrease. The models also have difficulty matching the observed rest-frame colors of galaxies. In all models the luminosity density of red galaxies increases sharply from $z\\sim3$ to $z\\sim2.2$, whereas it is approximately constant in the observations. Conversely, in the models the luminosity density of blue galaxies is approximately constant, whereas it decreases in the observations. These discrepancies cannot be entirely remedied by changing the treatment of dust and suggest that current models do not yet provide an adequate description of galaxy formation and evolution since $z\\sim 3$. ",
        "introduction": "In the current paradigm of structure formation, dark matter (DM) halos build up in a hierarchical fashion through the dissipationless mechanism of gravitational instability. The assembly of the stellar content of galaxies is instead governed by much more complicated physical processes, often dissipative and non-linear, which are generally poorly understood. To counter this lack of understanding, prescriptions are employed in the galaxy formation models. One of the fundamental tools for constraining the physical processes encoded in these models is the luminosity function (LF), since its shape retains the imprint of galaxy formation and evolution processes.  The faint end of the LF can be matched with a combination of supernova feedback and the suppression of gas cooling in low-mass halos due to a background of photoionizing radiation (e.g. \\cite{benson02}). Matching the bright end of the LF has proven more challenging. Very recent implementation of active galactic nucleus (AGN) feedback in semianalytic models (SAMs) has yielded exceptionally faithful reproductions of the observed local rest-frame $B$- and $K$-band global LFs \\cite{bower06,croton06}, including good matches to the local rest-frame $B$-band LFs of red and blue galaxies (although with some discrepancies for faint red galaxies; \\cite{croton06}).  The excellent agreement between observations and models at $z\\sim0$ is impressive but is partly due to the fact that the model parameters were adjusted to obtain the best match to the local universe. A key question is therefore how well these models predict the LF at earlier times. The SAMs of \\cite{bower06}, \\cite{croton06}, and \\cite{delucia06} have been compared to observations at $0<z<2$ \\cite{bower06,kitzbichler06}. Although the agreement is generally good, \\cite{kitzbichler06} infer that the abundance of galaxies near the knee of the LF at high redshift is larger in the SAMs than in the observations (except possibly for the brightest objects), in an apparent reversal of previous studies (e.g. \\cite{cimatti02}).  We have compared the observed rest-frame optical LFs of galaxies at $2.0 \\leq z \\leq 3.3$ \\cite{marchesini07} to those predicted by theoretical models in this redshift range, in order to test the predictive power of the latest generation of galaxy formation models. We have also compared the observed LF to predictions from smoothed particle hydrodynamics (SPH) simulations, which have so far only been compared to data at $z\\sim6$ \\cite{dave06}. The details of this comparison can be found in \\cite{marchesini07b}; here we summarize the main results from our study. All magnitudes are in the AB system, while colors are on the Vega system. ",
        "conclusions": "\\label{disc}  The main results of our comparison between the observed and the model-predicted rest-frame $V$-band LFs of galaxies at $z \\geq 2$ are (1) the SAM of Bower reproduces well the observed LF at $z\\sim 3$; (2) the models predict an increase with time of the rest-frame $V$-band luminosity density, whereas the observations show a decrease; (3) the models predict strong evolution in the red galaxy population, whereas in the observations most of the evolution is in the blue population; (4) the models greatly underpredict the abundance of galaxies with $B-V \\geq0.5$ at $z\\sim 3$.  The different results obtained for $U-V$ and $B-V$ colors are interesting, as they may hint at possible ways to improve the models. We further investigate the disagreement between observed and predicted colors in the SAM of Bower by comparing the observed and predicted colors of galaxies in the $B-V$ versus $U-B$ diagram. While the SAM seems to broadly reproduce the observed $U-B$ distribution, it predicts galaxies that are systematically bluer in $B-V$ than the observed galaxies. As shown in \\cite{marchesini07b}, the differences between observed and predicted colors could be due to larger amount of dust and/or to more complex star-formation histories in the observed galaxies. The ad hoc treatment of dust absorption is a significant and well-known source of uncertainty in the models.  Modifications to the specific dust model could partly resolve the differences between observations and SAM predictions. By simply multiplying the $A_{\\rm V}$ in the SAM by a fixed factor, we were able to better reproduce the observed LFs at $z\\sim2.2$ (although making the faint-end slope of the red galaxy LFs quantitatively too steep) and to have a better agreement between observed and predicted colors. However, at $z\\sim3$ this simple remedy is not able to solve the disagreement between the predicted and the observed number of $B-V \\geq 0.5$ galaxies. We conclude that, while ad hoc modifications of the dust treatment might help to alleviate some of the found disagreements, it does not seem to be sufficient to resolve the problem with global colors at $z\\sim3$."
    },
    "0801/0801.2783_arXiv.txt": {
        "abstract": "AzTEC is a mm-wavelength bolometric camera utilizing 144 silicon nitride micromesh detectors.  Herein we describe the AzTEC instrument architecture and its use as an astronomical instrument.  We report on several performance metrics measured during a three month observing campaign at the James Clerk Maxwell Telescope, and conclude with our plans for AzTEC as a facility instrument on the Large Millimeter Telescope. ",
        "introduction": "\\label{sec:int} In the era of ALMA, the scientific niche for large mm-wave telescopes comes in the ability to use array receivers to make large surveys of the sky.  The last decade has seen a series of major advances in mm-wavelength bolometric instruments and it is now commonplace for ground-based telescopes to have mm/submm wavelength cameras with tens to hundreds of detectors: for example, SCUBA on the JCMT~\\citep{Holland99}, MAMBO on the IRAM 30~m telescope~\\citep{Kreysa98}, and Bolocam on the Caltech Submillimeter Observatory~\\citep{Haig2004,Glenn2003}.  AzTEC is a millimetre-wavelength bolometer camera designed for the Large Millimeter Telescope (LMT). Its 144 silicon nitride micromesh bolometers operate with a single bandpass centered at either 1.1, 1.4, or 2.1~mm, with one bandpass available per observing run. AzTEC was commissioned at 1.1~mm during an engineering run at the James Clerk Maxwell Telescope (JCMT) in June 2005 and completed a successful observing run at the JCMT from November 2005-February 2006 during the JCMT05B semester.  In 2007 the instrument was installed on the 10~m Atacama Submillimeter Telescope Experiment (ASTE) where it will reside as a facility instrument through 2008.  AzTEC will be installed on the LMT in early 2009.  In its 1.1~mm wavelength configuration, AzTEC is sensitive to the Rayleigh-Jeans tail of the thermal continuum emission from cold dust grains. Consequently, AzTEC is well-suited for extragalactic surveys of dusty, optically obscured starburst (or AGN-host) galaxies detected by their sub-mm/mm emission (so-called Sub-Millimetre Galaxies, or SMGs).  As the far-infrared (FIR) peak of the cold dust emission is increasingly redshifted to mm-wavelengths with increasing distance, AzTEC is equally sensitive to dusty galaxies with redshift $z\\ge 1$. The AzTEC/JCMT05B surveys of the Lockman Hole and the Subaru Deep Field~(Austermann {\\it et al.} in preparation), COSMOS~\\citep{scott2008}, and GOODS-North~(Perera {\\it et al.} in preparation) represent the largest SMG surveys with uniform high sensitivity ($1\\sigma \\sim 1$ mJy) ever conducted, successfully demonstrating the superb mapping speed and stability of the instrument and providing a new insight on the clustering and cosmic variance of the SMG population.  In a similar vein, AzTEC's high mapping speed and high sensitivity make it an interesting instrument for studies of cold dust in the Milky Way and in nearby galaxies.  Configured for 1.4~mm and 2.1~mm observations, AzTEC will be tuned to make high resolution images of the Sunyaev-Zel'dovich effect in clusters of galaxies.  AzTEC on the LMT will have a per-pixel resolution of 10$^{\\prime\\prime}$ at 2.1~mm (5$^{\\prime\\prime}$ at 1.1~mm) and will therefore be an unprecedented instrument for the study of the energetics of the free electron gas in clusters.  In this paper we report on the design and performance of the AzTEC instrument. In Section~\\ref{sec:system}, we describe the components of the instrument and its configuration. The observing software, mapping strategies, and practical observing overheads are addressed in Section~\\ref{sec:obs}. We describe the calibration of AzTEC data in Section~\\ref{sec:cal} and in Section~\\ref{sec:per} we list the sensitivity and mapping speeds of AzTEC as measured at the during the JCMT05B run. We conclude with a brief discussion of the future of AzTEC in Section~\\ref{sec:fut}. ",
        "conclusions": ""
    },
    "0801/0801.0265_arXiv.txt": {
        "abstract": "We consider grains with superparamagnetic inclusions and report two new condensed matter effects that can enhance the internal relaxation of the energy of a wobbling grain, namely, superparamagnetic Barnett relaxation, as well as, an increase of frequencies for which nuclear relaxation becomes important. This findings extends the range of grain sizes for which grains are thermally trapped, i.e. rotate thermally, in spite of the presence of uncompensated pinwheel torques. In addition, we show that the alignment of dust grains by radiative torques gets modified for superparamagnetic grains, with grains obtaining perfect alignment with respect to magnetic fields as soon as the grain gaseous randomization time gets larger than that of paramagnetic relaxation. The same conclusion is valid for the mechanical alignment of helical grains. If observations confirm that the degrees of alignment are higher than radiative torques can produce alone, this may be a proof of the presence of superparamagentic inclusions. ",
        "introduction": "Alignment of dust in magnetic fields is one of the astrophysical problems of the longest standing. Discovered by Hiltner (1949) and Hall (1949) via observation of the polarization of starlight attenuated by interstellar dust, aligned grains are shown to trace magnetic fields in various environments, from interplanetary space (see Rosenbush et al. 2007) to circumstellar regions (see Aitken et al. 2002; Cho \\& Lazarian 2005).  The presently-accepted dominant alignment process is based on the action of radiative torques (RATs) on irregular grains (see Lazarian 2007 for a review). This process was proposed by Dolginov \\& Mytrophanov (1976) and extended, corrected and elaborated in the subsequent publications (Lazarian 1995; Draine 1996; Draine \\& Weingartner 1996, 1997, henceforth, DW96 and DW97; Weingartner \\& Draine 2003, Weingartner 2006; Lazarain \\& Hoang 2007a, henceforth LH07; Hoang \\& Lazarian 2008). This has induced theory-based modeling of polarized radiation from molecular clouds (Cho \\& Lazarian 2005; Pelkonen et al. 2007, Bethell et al. 2007) and accretion disks (Cho \\& Lazarian 2007).  While DW96 professed the omnipotence of RATs for aligning interstellar grains, the physics involved happens to be more complex than it was initially anticipated. In particular, an analytical model (AMO) for RATs proposed in LH07 shows that, for a variety of the directions between the magnetic field and the radiation flux, the alignment is not perfect (see also a particular example of this in Weingartner \\& Draine 2003 and more details in Hoang \\& Lazarian 2008). This causes the justified worries whether all the essential physical processes has been accounted for to explain the high degrees of interstellar polarization (see Whittet et al. 2008).  In LH07a we showed that ordinary paramagnetic relaxation does not affect the RAT alignment in most astrophysical environments, e.g., in interstellar gas. This does not preclude inclusions with strong magnetic response, e.g. superparamagnetic inclusions, to affect the RAT alignment. Such grains were discussed in Jones \\& Spitzer (1967) and also Martin (1995). Given the abundance of iron in dust, this would not be a far fetched suggestion. Moreover studies of interstellar grains caught in the Earth atmosphere (see Bradley 1994; Goodman \\& Whittet 1995), as well as, laboratory studies of the silicate films (see Djouadi et al. 2007) are suggestive of the abundance of such inclusions. Below in \\S 2 we consider both the effect of superparamagnetic inclusions on grain internal relaxation, i.e. on the relaxation that  may induce grain to rotate about its axes of maximal inertia (Purcell 1979; Lazarian 1994; Lazarian \\& Roberge 1997; Lazarian \\& Efroimsky 1998; Lazarian \\& Draine 1999ab, henceforth LD99ab)  and on the RAT alignment (in \\S 3). ",
        "conclusions": "In terms of the RAT alignment it is possible to show that even inclusions with $N\\sim 10^3$ can result in the appearance of a new high-$J$ attractor point. As a result, the RAT alignment can be enhanced substantially, becoming nearly perfect. This alignment is different from the alignment described in Draine \\& Weingartner (1996), however. In that paper, it was assumed that RATs provide the spin-up similar to pinwheel torques in Purcell (1979), while ordinary paramagnetic dissipation does the job of alignment. The latter interpretation of the RAT action is erroneous, as RATs provide the alignment of their own, making the weak paramagnetic torques irrelevant. It is only after being enhanced through the presence of superparamagnetic inclusions, that the paramagnetic dissipation can affect the alignment. If observations prove that perfect alignment of grains is required, {\\it SRAT alignment} can achieve this. Interestingly enough, the synthesis of pinwheel torques and superparamagnetic inclusions proposed in Spitzer \\& McGlynn (1979), does not work for interstellar medium, as superparamagnetic grains are strongly thermally trapped and cannot rotate suprathermally, unless RATs are present. Note, that  superparamagnetic inclusions will also enhance the mechanical alignment of helical grains. Such an alignment was proposed in Lazarian \\& Hoang (2007b) where it was shown that  torques arising from the gas-grain drift may act on irregular grains the same way as RATs do.    Our results above can be briefly summarized as follows:  $\\bullet$ Superparamagnetic inclusions affect both Barnett and nuclear relaxation, but in a different way: Barnett relaxation is being enhanced by the increase of magnetic susceptibility, while the nuclear relaxation does not changes its amplitude, getting, nevertheless extended for higher frequencies.  $\\bullet$ Supraparamagnetic inclusions modify the RAT alignment and RAT-superparamagnetic mechanism, which we termed {\\it SRAT alignment}, may provide higher degree of alignment compared to the RAT mechanism alone.  $\\bullet$ The results above are also applicable to mechanical alignment of irregular grains.   {\\bf Acknowledgments} The work was supported by the NSF Center for Magnetic Self-Organization in Laboratory and Astrophysical Plasmas and NSF grant AST 0507164."
    },
    "0801/0801.2867_arXiv.txt": {
        "abstract": "{ We present chemical evolution models for the Galactic disk. We also present a new determination of $X$, $Y$, and $Z$ for M17 a Galactic metal-rich H~{\\sc ii} region. We compare our models for the Galactic disk with the Galactic H~{\\sc ii} regions abundances. The $\\Delta Y/\\Delta O$ ratio predicted from the galactic chemical evolution model is in very good agreement with the $\\Delta Y/\\Delta O$ value derived from M17 and the primordial helium abundance, $Y_p$, taking into account the presence of temperature variations in this H~{\\sc ii} region. {From} the M17 observations we obtain that $\\Delta Y/\\Delta Z = 1.97 \\pm 0.41$, in excellent agreement with two $\\Delta Y/\\Delta Z$ determinations derived from K dwarf stars of the solar vicinity that amount to $2.1 \\pm 0.4$ and $2.1 \\pm 0.9$ respectively. We also compare our models with the solar abundances. The solar and Orion nebula O/H values are in good agreement with our chemical evolution model. }  \\addkeyword{galaxies: abundances} \\addkeyword{galaxies: evolution} \\addkeyword{H~II regions} \\addkeyword{ISM: abundances} \\addkeyword{ISM: individual: M17, Orion nebula} \\addkeyword{Sun: abundances}  \\begin{document} ",
        "introduction": "\\label{sec:intro}   The main purpose of this work is to study the evolution of the helium abundance with respect to the heavy elements as a function of time and position in the Galactic disk. For this purpose we will use the Galactic chemical evolution model by \\citet{car05} that has been successful in explaining: the observed O/H and C/H abundance gradients in the interstellar medium, ISM, the present gaseous distribution in the Galactic disk, the current star formation rate, the stellar mass as a function of the Galactic radius, and most of chemical properties of the solar vicinity.  To have useful observational constraints we need accurate $X$, $Y$, and $Z$ determinations or at least $\\Delta Y$/$\\Delta Z$ determinations to compare the models with observations. In this paper we recompute the $X$, $Y$, and $Z$ values for the H~{\\sc ii} region M17, the best Galactic H~{\\sc ii} region for which it is possible to compute an accurate enough $Y$ value, for this purpose we make use of the best observations available and the new He I atomic data needed for the helium abundance determination. We also compare the chemical evolution model with the $\\Delta Y$/$\\Delta Z$ determination derived from K dwarf stars of the solar vicinity with metallicities similar or higher than solar by \\citet{jim03} and \\citet{cas07}.  We also compare our model with the initial solar O/H value and with the Orion O/H value. The initial solar O/H value is representative of the ISM, 4.5 Gyr ago when the Sun was formed, and the Orion nebula O/H value is representative of the present day ISM.   The model together with the primordial helium determination is also used to provide an equation  between the $Y$ enrichment and the $O$ enrichment of the ISM. This equation can be used to provide the initial $Y$ values for those stars for which we can derive their initial oxygen abundances. These initial $Y$ values provide meaningful initial abundances for a set of stellar evolutionary models with different heavy elements content.  We adopt the usual notation $X$, $Y$, and $Z$ to represent the hydrogen, helium and heavy elements abundances by mass, respectively. Based on our models we study the increase of helium,  $\\Delta Y$, as a function of the increase of $C$, $O$, $Fe$, and $Z$ by mass.  In Section~\\ref{sec:models} we discuss the general properties of the chemical evolution models, we discuss inflow models for the Galaxy with two sets of stellar yields and present the chemical abundances for the disk at 5 galactocentric distances. For the two models discussed  we present the increase of helium by mass $\\Delta Y$, relative to the increase of carbon, oxygen, iron and heavy elements by mass, $\\Delta C$, $\\Delta O$, $\\Delta Fe$, and $\\Delta Z$. We also discuss the evolution of the $Y$ and $O$ abundances for our models,  and present an equation that predicts for the Galaxy the $Y$ enrichment as a function of the $O$ enrichment of the ISM.  In Section~\\ref{sec:M17} we present a new determination of the helium abundance for the metal rich H~{\\sc ii} region M17; this abundance is compared with our Galactic chemical evolution models.  In Section~\\ref{sec:Orion} we compare abundances for the Orion nebula with those of B stars of the Orion association. In Section~\\ref{sec:solar O} we discuss the absolute calibration of the O/H ratio in the local ISM, this ratio is one of the most important observational constraints for Galactic chemical evolution models, the absolute calibration of the O/H ratio is obtained based on the Orion nebula and the solar abundances. In Section~\\ref{sec:He} we compare the solar initial helium abundance, inferred from the standard solar models by \\citet{bah06}, with our Galactic chemical evolution models considering the time since the Sun was formed and the presence of gravitational settling and diffusion. The conclusions are presented in Section~\\ref{sec:conclusions}.  Throughout this paper we will use the primordial helium abundance by mass, $Y_p$, derived by \\citet{pei07} based on direct helium abundance determinations of metal poor extragalactic H~{\\sc ii} regions that amounts to $0.2477 \\pm 0.0029$. This result is in excellent agreement with the $Y_p$ determination by \\citet{dun08} that amounts to $0.2484 \\pm 0.0003$. This determination is based on the $\\Omega_bh^2$ value derived from WMAP observations, the assumption of standard big-bang nucleosynthesis, and the neutron lifetime, $\\tau_n$, of $885.7 \\pm 0.8$ sec obtained by \\citet{arz00}. Following \\citet{mat05} the value of $Y_p$ derived from WMAP is revised downwards to $0.2468 \\pm 0.0003$ by adopting the $\\tau_n = 878.5 \\pm 0.8$ sec derived by \\citet{ser05}, and to $0.2475 \\pm 0.0006$ by adopting for $\\tau_n$ the new world average that amounts to $881.9 \\pm 1.6$ sec, average that includes the results by \\citet{arz00} and \\citet{ser05}. ",
        "conclusions": "\\label{sec:conclusions}  Based on the HWY model we find the following equation to estimate the initial helium abundance with which stars form in the Galactic disk  $$Y = Y_p + (3.3 \\pm 0.7) O,$$ for $O < 4.3\\times10^{-3}$, and  $$Y = Y_p + (3.3 \\pm 0.7) O + (0.016 \\pm 0.003)(O/4.3 \\times10^{-3} - 1)^2,$$ for $ 4.3\\times 10^{-3} < O < 9\\times10^{-3}$.  The increase of {$\\Delta Fe/\\Delta Z$} has to be taken into account in order to determine the $\\Delta Y/\\Delta Z $ value based on the [Fe/H] abundances of stars in the solar vicinity.  High mass loss rates due to stellar winds should be adopted in the evolutionary stellar models for massive stars of high metallicity, because only the Galactic chemical evolution models with HWY can reproduce simultaneously the O/H  and C/O Galactic gradients, the $C/O$ versus $O$ relation in the solar vicinity, and the $\\Delta Y/\\Delta O$ value in the inner Galactic disk.  Based on the O/H value of the Orion nebula and the solar photospheric value together with a chemical evolution model of the Galaxy we recommend for the present day local ISM a value of 12 + log O/H = $8.79 \\pm 0.08$, where both the gaseous and the dust components of O are taken into account.  By comparing the O/H value of the Orion nebula with the solar value we find that the nebular ratio derived using O recombination lines, that is equivalent to the use of the forbidden O lines under the adoption of a $t^2 \\neq 0.00$, is in considerably better agreement with the initial solar value than the Orion nebula value derived adopting $t^2 = 0.00$.  The stellar O/H and Ar/H abundance ratios derived by \\citet{cun06} and \\citet{lan08} for B stars of the Orion association are in excellent agreement with the nebular abundance ratios derived from the $t^2 \\neq 0.00$ values for the Orion nebula.  The $\\Delta Y/\\Delta Z  = 1.97 \\pm 0.41$ value derived from observations of M17 for $t^2 = 0.036$ is in very good agreement with the $2.1 \\pm 0.4$ and the $2.1 \\pm 0.9$ values derived by \\citet{jim03} and \\citet{cas07} from K dwarf stars of the solar vicinity. On the other hand the value $\\Delta Y/\\Delta Z  = 4.00 \\pm 0.75$ derived from observations of M17 for $t^2 = 0.000$ is not.  Both Galactic chemical evolution models with the HWY set and the LWY set are in agreement with the observed $\\Delta Y/\\Delta Z$ for $t^2 = 0.036$ but not with the $\\Delta Y/\\Delta Z$ for $t^2 = 0.000$. Higher accuracy determinations of $\\Delta Y/\\Delta Z$ for high metallicity objects are needed to discriminate between the HWY model and the LWY model predictions.       \\vskip 1cm  We are grateful to Brad Gibson and Antonio Peimbert for several fruitful discussions. We are also grateful to the anonymous referee for some excellent suggestions. This work was partly supported by the CONACyT grants 46904 and 60354."
    },
    "0801/0801.0323_arXiv.txt": {
        "abstract": "The influence of a possible non zero chemical potential $\\mu$ on the nature of dark energy is investigated by assuming that the dark energy is a relativistic perfect simple fluid obeying the equation of state (EoS), $p=\\omega \\rho$ ($\\omega <0, constant$). The entropy condition, $S \\geq 0$, implies that the possible values of $\\omega$ are heavily dependent on the magnitude, as well as on the sign of  the chemical potential. For $\\mu >0$, the $\\omega$-parameter must be greater than -1 (vacuum is forbidden) while for $\\mu < 0$ not only the vacuum but even a phantomlike behavior ($\\omega <-1$) is allowed. In any case, the ratio between the chemical potential and temperature remains constant, that is, $\\mu/T=\\mu_0/T_0$. Assuming that the dark energy constituents have either a bosonic or fermionic nature, the general form of the spectrum is also proposed. For bosons $\\mu$ is always negative and the extended Wien's law allows only a dark component with $\\omega < -1/2$ which includes vacuum and the phantomlike cases. The same happens in the fermionic branch for $\\mu <0$. However, fermionic particles with $\\mu >0$ are permmited only if $-1 < \\omega < -1/2$. The thermodynamics and statistical arguments constrain the EoS parameter to be $\\omega < -1/2$, a result surprisingly close to the maximal value required to accelerate a FRW type universe dominated by matter and dark energy ($\\omega \\lesssim -10/21$). ",
        "introduction": "The  current idea of an accelerating Universe is based on a large convergence of independent observations, and its explanation constitutes one of the greatest challenges for our current understanding of fundamental physics \\cite{SN,CMB}. In the context of a Friedmann-Robertson-Walker (FRW) cosmology, dominated by pressureless matter with density $\\rho_m$ plus an extra component of density $\\rho$ and pressure $p$,  the scale factor evolution  is governed by the equation 3${\\ddot a}/a= -4\\pi G(\\rho_m + \\rho + 3p)$. This means that a hypothetical component with large negative pressure may drive the evolution of an expanding accelerating Universe. This exotic component is usually termed dark energy or quintessence, and it represents about 70\\% of the total content in the Universe. The origin and the nature of dark energy is still a mystery, however, there is no doubt that its existence is beyond the domain of the standard model of particle physics \\cite{review}.  Among a number of possibilities to describe the dark energy component, the simplest and most theoretically appealing way is by means of a cosmological constant $\\Lambda$, which acts on the FRW equations as an isotropic and homogeneous source with a constant equation of state parameter $p/\\rho=-1$. On the other hand, although cosmological scenarios with a $\\Lambda$ term might explain most of the current astronomical observations, from the theoretical viewpoint they are plagued with some fundamental problems thereby stimulating the search for alternative dark energy models driven by different candidates \\cite{model,list1}.  In the XCDM scenario, the dark energy component is phenomenologically described by an equation of state $p = \\omega\\rho$. The case $\\omega = -1$ reduces to the cosmological constant. The imposition $\\omega \\geq -1$ is physically motivated by the classical fluid description \\cite{HElis82}. This hyphothesis introduces a strong bias in the $\\omega$-parameter determination from observational data. In order to circunvent such a difficulty, superquintessence or phantom dark energy cosmologies have been recently considered where such a condition is relaxed \\cite{faraoni02}. In contrast to the usual quintessence model, a decoupled phantom component presents an anomalous evolutionary behavior. For instance, the existence of future curvature singularities, a growth of the energy density with the expansion, or even the possibility of a rip-off of the structure of matter at all scales are theoretically expected. Although possessing such strange features, the phantom behavior is theoretically allowed by some kinetically scalar field driven cosmology \\cite{chi00}, as well as, by brane world models \\cite{Shani03}, and, perhaps, more important to the present work, a PhantomCDM cosmology is not ruled out by the present type Ia Supernovae and other observations \\cite{ChouPadm,Alc04}.  In a series of papers \\cite{limamaia,LA04}, we have studied some thermodynamics and statistical properties of dark energy with no chemical potential ($\\mu =0$). By using standard thermodynamics for a relativistic simple fluid, we concluded that the case of phantom energy is ruled out because the total entropy of a dark component with $\\omega < -1$ is negative. In addition, by combining thermodynamics and statistical arguments  the EoS was restricted to the interval  $-1 \\leq \\omega < -1/2$ and a fermionic nature to the dark energy particles was favored.  Later on, thermodynamics arguments in favor of the phantom hypothesis were put forward by Gonz\\'alez-D\\'{i}az and Sig\\\"uenza \\cite{gonzalez}. They claimed that the temperature of a phantomlike fluid is always negative in order to keep its  entropy positive definite (as statistically required). This new viewpoint was justified by arguing that the scalar field representation of a phantom field has a negative kinetic term ${\\dot \\phi}^{2}$ which quantifies the translational kinetic energy of the associated fluid system,  and, as such, its temperature (a measure of the average kinetic energy) should be negative. Although temptative to some degree, both approaches have been considered in the literature (see, for instance, \\cite{pavon,jorge} and Refs. therein). More important to the present work, they share a common property, namely, the chemical potential of the dark energy fluid was set to be zero from the very beginning.  In this article we reanalyze the thermodynamic and statistical properties of the dark energy scenario by considering the existence of a non-zero chemical potential. In order to clarify some subtleties present in the earlier results, we rederive the physical quantities in the presence of $\\mu$ by adopting the full thermostatistic approach proposed in Refs. \\cite{limamaia,LA04}. This means that all thermodynamic and statistical properties follow directly from the EoS plus the hypothesis that $\\mu \\neq 0$. In particular, the temperatures  of the dark energy fluids must be always positive definite.  This is an interesting aspect of the present work since there are many scalar field representations for dark energy fluids, however, the thermodynamic laws are independent to what happens at a microscopic level as long as the equation of state has been  defined.  As we shall see, one of the main consequences of a negative chemical potential is that the phantom scenario is recovered without the need to appeal to negative temperatures. In addition, a bosonic or fermionic nature of the dark energy component now becomes possible.  The paper is planned as follows. In section 2 we discuss the thermodynamic constraints when a chemical potential is introduced. In  section 3 we consider a statistical analysis  by assuming that the dark energy particles are massless and can have either a bosonic or a fermionic nature. In the conclusion section, a joint analysis involving the thermodynamic and statistical constraints on the EoS $\\omega$-parameter is presented ",
        "conclusions": "In this paper we have investigated the thermodynamic and statistical properties of a dark energy fluid with equation of state, $p=\\omega \\rho$, by assuming that its chemical potential is different from zero.  In Figure 4, we summarize the main results of the present analysis by combining both approaches.  As discussed in section 2, the regions with $S < 0$ are  thermodynamically forbidden. Note also that  many dark energy fluids satisfy the combined constraints regardless of the $\\mu$ sign, that is, a large interval of  negative $\\omega$ values is allowed by thermodynamic and statistical considerations.  However, a phantom like behavior ($\\omega < -1$) is permitted only for $\\mu < 0$, and the corresponding massless particles can have either a bosonic or fermionic nature.  It was also proved (see also Fig. 4) that the EoS parameter of a dark energy fluid obeying a generalized Wien's law always satisfies the constraint $\\omega<-0.5$ (a thermostatistics limit).  This upper limit is surprisingly close to the maximal value of the EoS $\\omega$-parameter  necessary to accelerate the present universe. Actually, in order to accelerate  a FRW universe dominated by matter and dark energy, the EoS parameter must satisfies the inequality, $\\omega <  -(1+ \\Omega_m/\\Omega_x)/3$. Therefore, for $\\Omega_m\\cong0.3$ and $\\Omega_x\\cong0.7$, as indicated by the present observations \\cite{SN,CMB}, one finds the dynamic constraint $\\omega \\lesssim - 10/21$.  Finally, it should be stressed that for $\\mu=0$ one finds  $\\omega_{min}=-1$ (see Eq. (\\ref{accord})) in accordance to the results previously derived by Lima and Alcaniz \\cite{LA04}. The present analysis with $\\mu \\neq 0$  also opens  the possibility for an EoS parameter $\\omega < -1$, thereby recovering the idea of a phantom dark energy without to appeal to negative temperatures. Perhaps more interesting, unlike the results  for $\\mu=0$  which favored only a fermionic nature to the dark energy fluid (phantom and nonphantom), it was demonstrated that a bosonic kind of dark energy becomes possible from a thermostatistics viewpoint."
    },
    "0801/0801.0115_arXiv.txt": {
        "abstract": "s{ In the limit of infinite external magnetic field $\\bf B$ the static field of an electric charge is squeezed into a string parallel to $\\bf B$. Near the charge the  potential grows like $|x_3|(\\ln |x_3| + const)$ with the coordinate $x_3$ along the string. The energy of the string breaking is finite and very close to the effective photon mass.}      It is known that the attraction force between two colored charges, when calculated using the Wilson loop method on a lattice, is concentrated in a string with its width being of the order of the lattice spacing $l$. The string potential consists of three additive components: $(i)$ the Coulomb term $1/R$ that dominates at small distance $R$ from the charge, $(ii$) a constant term, corresponding to the infinite mass renormalization (recall that $l$ is the ultraviolet cutoff parameter in the lattice theory), turns to infinity as $1/l$, when  $l$ is taken close to zero,  $(iii$) the term linearly growing with $R$ corresponding to a constant string tension and providing the confinement according to the Wilson area criterion (see, $e.g.$, \\cite{kondo}).  We show that in quantum electrodynamics with external magnetic field $\\bf B$ the Coulomb potential of a point-like static charge, when corrected by the vacuum polarization, acquires a similar string-like form in the infinite-magnetic-field limit, the electron Larmour radius $L_{\\rm B}=(\\sqrt{eB})^{-1}=(1/m\\sqrt{b})\\rightarrow 0$ playing, in a way, the same r\\^{o}le as the spacing $l$ does in the lattice theory. Here $b$ stands for the magnetic field, $b=B/B_0,$ measured in the units of $B_0=m^2/e=4.4\\times 10^{13}$G, $e$ and $m$ are electron charge and mass, resp.  Electric potential $A_0$ of a point charge $q$ when calculated as a sum of chains of electron-positron loops in a magnetic field $B$, so strong that the Larmour radius is much smaller than the electron Compton length, $\\Lar \\ll m^{-1}$ (this implies $B\\gg B_0$), is represented as a sum \\bee\\label{twoparts}A_0({\\bf x})= A_{\\rm s.r.}({\\bf x})+A_{\\rm l.r.}({\\bf x}).\\eend The analytic expressions for these functions are given in \\cite{SU1, SU2}. The argument  $\\bf x$ here is the radius-vector with its origin in the charge, its components across and along the magnetic field being ${\\bf x}= ({\\bf x}_\\perp, x_3)$.   The function $A_{\\rm s.r.}({\\bf x})$ has the exact scaling property\\bee\\label{scaling}\\hspace{-1.5cm} A_{\\rm s.r}({\\bf x})= \\frac{\\widetilde{A}(\\widetilde{\\bf x})}{L_{\\rm B}},\\eend where the dimensionless function $\\widetilde{A}$ contains the magnetic field through its argument $\\widetilde{\\bf x}={\\bf x}\\Lar^{-1}$ only . The function $A_{\\rm s.r.}({\\bf x})$ is short-range: for distances from the charge, large in the Larmour scale, $|x_3|\\gg \\Lar$, or $|x_\\perp|\\gg \\Lar,$ it  reduces to the Yukawa law\\bee\\label{yuk} A_{\\rm s.r.}({\\bf x})\\simeq \\frac {q}{4\\pi L_{\\rm B}}\\frac {\\exp\\{-\\left(\\frac{2\\alpha}\\pi\\right)^{\\frac 1{2}} {\\sqrt{\\widetilde{x}_\\perp^2+\\widetilde{x}_3^2}} \\}}{\\sqrt{\\widetilde{x}_\\perp^2+\\widetilde{x}_3^2}}=\\frac q{4\\pi}\\frac{\\exp\\{-\\left(2\\alpha b/\\pi\\right)^{\\frac 1{2}}m|{\\bf x}|\\}} {|{\\bf x}|},\\eend  where $\\alpha=e^2/4\\pi=1/137$. This equation reflects the Debye screening of the charge by the polarized vacuum. The effective photon mass (inverse Debye radius) in (\\ref{yuk}) $M= (2\\alpha/\\pi)^{1/2}\\;L_{\\rm B}^{-1}$ tends to infinity together with the magnetic field. The \"photon mass $\\widetilde{M}\\;$\" in the dimensionless function $\\widetilde{A}(\\widetilde{\\bf x})$ is finite, ${\\widetilde{M}}= (2\\alpha/\\pi)^{1/2}$, and corresponds to the topological photon mass in the two-dimensional massless electrodynamics of Schwinger \\cite{schwinger}. Anisotropic corrections to (\\ref{yuk}) were pointed in \\cite{tehran}.  The second term in (\\ref{twoparts}) is long-range: it slowly decreases at the distances of the order of the Compton length $m^{-1}$ and larger, following the anisotropic Coulomb law \\bee\\label{llargex}A_{\\rm l.r.}(x_3,x_\\perp)\\simeq \\frac 1{4\\pi}\\frac {q}{\\sqrt{(x^\\prime_\\perp)^2+x_3^2}},\\qquad x^\\prime_\\perp=x_\\perp\\beta,~~~~\\beta=\\left(1+\\frac{\\alpha b}{3\\pi}\\right)^{\\frac 1{2}}.\\eend It represents the whole potential $A_0({\\bf x})$ (\\ref{twoparts}) there, since $A_{\\rm s.r.}({\\bf x})$ is already negligible at such distances. This potential decreases with the growth of perpendicular distance $x_\\perp$ from the charge faster than along the field. The equipotential surface is an ellipsoid $(x^\\prime_\\perp)^2+x_3^2=\\; $const, which contracts to a string in the limit $b=\\infty$. The family of electric lines of force, parameterized by the angle $\\phi,$  is for each value of the magnetic field given as $x_3=(x_\\perp)^{\\beta^{-2}}\\tan\\phi$ (see Fig. 1). All of them gather together inside a string passing through the charge and directed along the external magnetic field. \\begin{figure}[t] \\centering \\includegraphics[scale=1.0]{ShabFig1.eps} \\caption{\\hspace{0.5cm}Lines of force of electric field coming from a point charge.\\qquad \\qquad \\qquad \\qquad \\qquad (a) No magnetic field, $B$=0, (b) $B=10^4B_0$} \\end{figure}  The short-range part of the potential has the following asymptotic expansion near the point $x_3=x_\\perp=0,$ where the charge is located, \\bee\\label{expand}A_{\\rm s.r.}({\\bf x})\\simeq \\frac{q}{4\\pi}\\left(\\frac 1 {|\\bf x|}-2mC_{\\rm s.r.}+o(x_3^2)+o(x_\\perp^2)\\right).\\eend  For large magnetic fields $b\\gg 2\\pi/\\alpha \\sim 10^3$ we have $2mC_{\\rm s.r.}\\simeq 0.9595\\;\\Lar^{-1}\\sqrt{2\\alpha/\\pi}$. It is infinite for $\\Lar=0$ or $b=\\infty$, as stated at the beginning. The corresponding constant in the expansion of $\\widetilde{A}(\\widetilde{\\bf x})$ is finite: $2\\widetilde{C}=0.9595\\;\\sqrt{2\\alpha/\\pi}$. The constant 0.9595 is calculated using the experimental value of $\\alpha$. For $\\alpha=0$ it turns into unity. The growth of $C_{\\rm s.r.}$ with the magnetic field following the square root law provides the narrowing of the potential and, in the end, the finiteness of the ground state energy of a hydrogen-like atom in infinite magnetic field (see \\cite{SU1, SU2} and also the discussion \\cite{wang}). In the limit $b=\\infty$ the short-range part of the potential becomes the Dirac $\\delta$-function: \\bee\\label{deltaxperp} \\left.A_{\\rm s.r.}(0,x)\\right|_{b=\\infty}=2.178~\\frac q{2\\pi}\\delta(x).\\eend The behavior of the long-range part $A_{\\rm l.r.}(x_3,x_\\perp=0)$ at the string $x_\\perp=0$ is shown in Fig.2. (See \\cite{SU2} for analytical equations). The limiting form of this function at $\\Lar =0$ is finite. It decreases in agreement with Eq. (\\ref{llargex}) at large distances, and has the following behavior near the charge, $|x_3|\\ll m^{-1}$, \\bee\\label{confinement}\\left.A_{\\rm l.r.}(x_3,0 )\\right|_{b=\\infty}\\simeq\\left.A_{\\rm l.r.}(0,0 )\\right|_{b=\\infty}\\nonumber\\\\+\\frac{qm}{4\\pi}\\left(1-{\\alpha\\over\\pi} f(\\alpha)\\right)2m|x_3|\\left[\\ln (2m|x_3|)-{1\\over 2}\\ln 2+\\gamma-1\\right],\\quad \\eend where $\\gamma=0.577$ is the Euler constant and the coefficient $f$ depends on the fine structure constant,  $f\\left(\\alpha={1/137.036}\\right)=4.533$. It is nonanalytic in $\\alpha$ :\\\\$\\left.f(\\alpha)\\right|_{\\alpha\\rightarrow 0}\\simeq -\\ln\\alpha$. We see that the growth is not linear. It provides \"confinement\" within the Compton distances, where the approximation (\\ref{confinement}) is valid. The first constant term in (\\ref{confinement}) is defined by an integral depending on the fine structure constant. If calculated with its experimental value, $\\alpha=1/137.036,$ it makes $\\left.A_{\\rm l.r.}(0,0 )\\right|_{b=\\infty}=1.4152(qm/2\\pi)$. The numerical coefficient here is rather close to $\\sqrt{2}=1.4142$. Bearing in mind that $\\left.A_{\\rm l.r.}(\\infty,0 )\\right|_{b=\\infty}=0$ we note that $\\left.A_{\\rm l.r.}(0,0 )\\right|_{b=\\infty}$ is the increment of the long-range part of the potential along the string between the point where the charge is located and the infinitely remote point, i.e. is equal to the work needed for removing a  unit test charge to infinity. It may be referred to as the energy of the string breaking. It is remarkable that this quantity is finite. If we set $q=e$, we find that the  energy density of breaking of a string associated with the electron is with surprising accuracy equal to the dimensionless photon mass : $\\left.A_{\\rm l.r.}(0,0 )\\right|_{b=\\infty}/m=1.0007~\\widetilde{M}$. The coincidence would be exact for the value of the fine-structure constant equal to $1/121$. By postulating this coincidence as a a physical principle we may obtain an approach for calculating $\\alpha$, to which end an approximation, better than one-loop approximation referred to here, would be needed. \\begin{figure}[t] \\centering \\includegraphics[scale=0.7]{ShabFig2.eps} \\caption{Energy $-eA_{\\rm l.r.}(x_3,0)$ of the electron in the long-range part  of the potential of the point charge $q=eZ$ for $b=10^6,\\; 10^5$, $3\\times 10^4$, $10^4$ (dashed lines from bottom to top). The dashed thick line corresponds to the limit $b=\\infty$ and represents the string potential.} \\end{figure}   Formation of a string in QED discussed in the present paper is not completely unexpected, since it was noted \\cite{kondo} that the Abelian theory does have a topologically nontrivial two-dimensional sector (described by a nonlinear $\\sigma$-model). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.3927_arXiv.txt": {
        "abstract": "{} {In this paper we present the first results of a detailed spectroscopic and photometric analysis of the $V = 11.7^{\\rm m}$ eclipsing binary \\object{ASAS~J052821+0338.5}. } {With the FIES spectrograph at the Nordic Optical Telescope we have obtained a series of high-resolution spectra \\mbox{($R \\approx 47000$)} covering the entire orbit of the system. In addition we obtained simultaneous broadband photometry from three small aperture telescopes. From these spectroscopic and photometric data we have derived the system's orbital parameters and determined the fundamental stellar parameters of the two components. } {Our results indicate that ASAS~J052821+0338.5 is a K1/K3 pre-main-sequence eclipsing binary, with component masses of $1.38\\,{\\rm M}_{\\sun}$ and $1.33\\,{\\rm M}_{\\sun}$ and a period of 3.87 days, located at a distance of $280 \\pm 30$ pc. The kinematics, physical location and the evolutionary status of the two stars suggest that ASAS~J052821+0338.5 is a member of the $\\sim 11$ Myr old Orion OB1a subassociation. The systems also exhibits smooth $\\sim 0.15^{\\rm m}$ out-of-eclipse variations that are similar to those found in RS~CVn binaries. Furthermore the parameters we derived are consistent with the $10$--$13$ Myr isochrones of the popular Baraffe stellar evolutionary models. } {} ",
        "introduction": "The advantageous geometry of detached double-line eclipsing binary systems makes it possible to measure stellar radii and masses to very high accuracy (1-2\\%, Andersen \\cite{andersen91}). Such accurate parameters are of crucial importance in providing observational constraints on stellar evolutionary models. However, very few of the known double-lined eclipsing binaries are stars at early stages of evolution, leaving the models in this domain with only few empirical points of reference.  The list of known double-lined eclipsing binaries with components that are in the pre-main-sequence (PMS) phase of evolution is not long. Up to now, only five low-mass double-lined eclipsing binaries in which both components are PMS objects have been reported, RXJ~0529.4+0041A (Covino et al. \\cite{covino00}, Covino et al. \\cite{covino04}), V1174~Ori (Stassun et al. \\cite{stassun04}), the brown dwarf eclipsing binary 2MASS~J05352184-0546085 (Stassun et al. \\cite{stassun06, stassun07}) and, very recently, JW 380 (Irwin et al. \\cite{irwin07b}) and Par 1802 (Cargile et al. \\cite{cargile07}). In addition, EK~Cep and TY~CrA (Popper \\cite{popper80}; Popper \\cite{popper87}; Andersen \\cite{andersen91}; Casey \\cite{casey98}) are known to harbour one PMS component and one main-sequence component.  In this paper, we announce the discovery that \\mbox{ASAS J052821+0338.5} (also known as \\mbox{GSC2.2 N300311274}, \\mbox{2MASS 05282082+0338327} and \\mbox{USNO-B1.0~0936-0073790}) is a PMS eclipsing binary system. In addition, the system could be associated with the ROSAT X-ray source \\mbox{1RXS J052820.4+033823}, located only $11^{\\prime\\prime}$ away, within ROSAT's positional error margin of $14^{\\prime\\prime}$. Using new high-cadence photometric and high-resolution spectroscopic observations we have determined the system's fundamental properties and confirm the PMS nature of both components. ",
        "conclusions": "\\label{sec:discussion}  \\begin{figure} \\centering \\includegraphics[width=0.95\\columnwidth]{stempelsfig8.pdf} \\caption{The modified Hertzprung-Russel diagram showing radius versus effective temperature for the components of ASAS~J052821+0338.5 (black solid circles) and other PMS EBs in Orion (grey solid circles). Evolutionary tracks (1.4, 1.3, 1.2, 1.1, 0.7 M$_{\\sun}$) and isochrones (3, 6, 10, 13, 16, 300 Myr) from Baraffe et al. \\cite{baraffe98} are overplotted as grey lines. The black solid line shows the evolutionary track interpolated at the mass of the primary star (1.375~M$_{\\sun}$). The black dashed line shows the evolutionary track interpolated at the mass of the secondary star (1.329~M$_{\\sun}$). } \\label{fig:hrplot} \\end{figure}  The photometric and spectroscopic observations we obtained of ASAS~J052821+0338.5 identify this system as a new pre-main-sequence eclipsing binary with a period of $3.87$ days. The masses and radii of the components place the system in an as yet unsampled region of the mass-radius and HR diagrams, and thereby provide a new pair of reference points for PMS stellar evolutionary models. Both components of the system are more massive than all other members of known PMS eclipsing systems, except for TY~CrA~B, the 1.64 ${\\rm M}_{\\sun}$ and $\\sim 3$ Myr old companion to a main-sequence B-type primary (Casey et al. \\cite{casey98}).  The physical location and kinematic properties of ASAS~J052821+0338.5 suggest this system to be a member of the $\\sim 11.4 \\pm 1.9$~Myr Orion OB1a subassociation. This age is also compatible with the position of the two components in the mass-radius diagram shown in Fig.~\\ref{fig:mrplot}, where both stars fall close to the 13~Myr isochrone of the Baraffe et al. (\\cite{baraffe98}) evolutionary models. At this age, the two 1.35 ${\\rm M}_{\\sun}$ stars have just left the fully-convective Hayashi stage. The temperature of the two components will then increase at almost constant luminosity, first slightly growing in radius before their final contraction towards the main sequence, settling at a final temperature of $\\sim 6250$\\,K (for a more detailed description of pre-main-sequence evolution, see for example Palla \\& Stahler \\cite{palla93}). We also find a that the locations of the stars in the HR diagram agree quite well with the predicted 13 Myr evolutionary tracks of Baraffe et al. (\\cite{baraffe98}, see Fig.~\\ref{fig:hrplot}). Although there appear to be small differences between the observed radii and the predicted values, this does not indicate any inconsistency since the only well-constrained parameter is the sum of the radii (see Sect. \\ref{sec:lcmodel}), and full consistency of both stars with the predicted tracks requires only small antagonal adjustements (of about 3\\%) to the stellar radii. Therefore, a critical next step to further improve these parameters will be a multi-band lightcurve analysis with a more detailed spot treatment.  Most of the known PMS eclipsing binaries are located in Orion. ASAS~J052821+0338.5 and RXJ~0529.4+0041A are thought to be located in the Orion OB1a subassociation (Covino et al. \\cite{covino00}, Covino et al. \\cite{covino04}). V1174~Ori is considered to be a member of Orion OB1c (Stassun et al. \\cite{stassun04}) and 2MASS~J05352184-20130546085, JW 380 and Par 1802 are found in the Orion Nebula Cluster (Stassun et al. \\cite{stassun06}; Irwin et al. \\cite{irwin07b}; Cargile et al. \\cite{cargile07}). These OB subassociations in Orion form an evolutionary sequence, with ages of $\\sim$\\,1--12 Myr. This age range covers a key period in star formation when stars experience rapid evolution in both radius and temperature. The components of the PMS eclipsing binaries in Orion have masses between 0.03 and 1.35~${\\rm M}_{\\sun}$ and have well-determined radii and temperatures. As such they form an easily accessible and homogenous sample of late-type PMS stars that can be used to perform relative calibrations of PMS stellar evolutionary models.  As described in Sect. \\ref{sec:phot}, we detected considerable ($\\sim~0.15^{\\rm m}$) out-of-eclipse variations, occuring around the same orbital phase throughout our four-month campaign of photometric observations, not unlike the `photometric wave' that is characteristic of the chromospherically active RS~CVn binaries (see the review by Rodon{\\'o} \\cite{rodono92}). Like in an RS~CVn binary, the components of ASAS~J052821+0338.5 have relatively short rotational periods and extensive convective envelopes, which help to drive a stellar dynamo and enhance magnetic activity. The photometric variations may then be explained by a large cool spot (or region of spots) located on the back of the primary, rotating in and out of view in phase with the binary period. However, this configuration requires the primary to rotate synchronously with the binary orbit. According to the theory of Zahn \\& Bouchet (\\cite{zahn89}), tidal breaking during the Hayashi phase, when stars are fully convective, is a very efficient mechanism for inducing orbital synchronization and circularization in short-period late-type binaries. Once the convective envelope starts decreasing, tidal breaking loses its efficiency, and further stellar contraction may lead to (temporary) departures from synchronous rotation. ASAS~J052821+0338.5 seems to follow this pattern. The orbit is almost perfectly circularized, and the observed rotational velocities of the primary and secondary are $24.5 \\pm 0.8$ km/sec and $24.5 \\pm 0.7$ km/s, respectively. Since the expected synchronous rotational velocities are $24.0$ km/sec and $23.4$ km/sec, the primary is most likely in synchronous rotation with the binary orbit, while the smaller secondary may rotate slightly super-synchronously.  While an extended cool spot on the back of the primary can fully explain the observed out-of-eclipse photometric variations, we did not detect any signatures of a cool spot in our current high-resolution spectra, such as changes in the observed line-ratios or the appearance of molecular features. However, the brightness contrast between the spotted and the unspotted areas is large, and since the signal-to-noise of our spectra is moderate it is unclear whether such an effect would be detectable. An alternative explanation to the out-of-eclipse variability might be periodic eclipses of the system by a warped circumbinary disk, or by circumstellar material in co-rotation with the binary. Such concentrations of circumstellar matter localized around the stable co-rotating Lagrangian points (L2 and L3) have been detected in the close PMS binary V4046 Sgr (Stempels \\& Gahm \\cite{stempels04}). We are currently pursuing multi-band photometry of the system in order to determine the nature of the out-of-eclipse variations, as well as to further improve the orbital and stellar parameters of ASAS~J052821+0338.5, in particular the temperatures and radii of the components."
    },
    "0801/0801.4446_arXiv.txt": {
        "abstract": "A clear detection of excess of power, providing a substantial evidence for solar-like oscillations in the G5 subgiant \\muher{}, is presented. This star was observed over seven nights with the SARG echelle spectrograph operating with the 3.6-m Italian TNG Telescope, using an iodine absorption cell as a velocity reference.  A clear excess of power centered at 1.2\\,mHz, with peak amplitudes of about 0.9\\,\\ms in the amplitude spectrum is present.  Fitting the asymptotic relation to the power spectrum, a mode identification for the $\\ell=0,1,2,3$ modes in the frequency range $900-1600\\; \\muHz$ is derived. The most likely value for the large separation turns out to be  56.5 \\muHz, consistent with theoretical expectations. The mean amplitude per mode ($l=0,1$) at peak power results to be $0.63 \\;\\rm m\\,s^{-1}$, almost three times larger than the solar one. ",
        "introduction": "The search for solar-like oscillations in stars has experienced a tremendous growth in recent years   (See \\citealt{B+K2006b},  for a summary).  Most of the results came from high-precision Doppler measurements using spectrograph such as CORALIE, HARPS, UCLES, UVES. In particular, recent measurements obtained with the SARG spectrograph have led to the determination of frequencies, mode amplitudes, lifetime and granulation noise on Procyon A \\citep{cla05,sil07}, a star for which the nature of oscillation spectrum is still debated \\citep{natureHD,nature2004}.  In this paper we report the detection of excess of power, providing evidence for oscillations, in the G5 subgiant \\muher{} (HR~6623, $V=3.417\\pm0.014$, G5~IV), a slightly evolved  solar-type star with mass $1.1\\,M_\\sun$, effective temperature $T_{\\rm e}=5596 \\pm 80 \\; \\rm K$,  and $\\rm log \\; g = 3.93\\pm 0.10$ \\citep{fu98}. ",
        "conclusions": "Our observations of \\muher{} show an evident excess of power in the PS, clearly disentangled from the low frequency increase, and with a position and amplitude that are in agreement with expectations.  Although hampered by the single-site window, the  comb analysis and the echelle diagram show clear evidence for regularity in the peaks at the spacing expected from asymptotic theory. Moreover, our results provide valuable confirmation that oscillations in solar-like stars really do have the amplitudes that scales as $L/M$  by extrapolating from the Sun. We hope that in the near future a multi-site observing campaign will allow us to explore further the oscillation spectrum of \\muher."
    },
    "0801/0801.3791_arXiv.txt": {
        "abstract": "The role of hypernuclear physics for the physics of neutron stars is delineated. Hypernuclear potentials in dense matter control the hyperon composition of dense neutron star matter. The three-body interactions of nucleons and hyperons determine the stiffness of the neutron star equation of state and thereby the maximum neutron star mass. Two-body hyperon-nucleon and hyperon-hyperon interactions give rise to hyperon pairing which exponentially suppresses cooling of neutron stars via the direct hyperon URCA processes. Non-mesonic weak reactions with hyperons in dense neutron star matter govern the gravitational wave emissions due to the r-mode instability of rotating neutron stars. ",
        "introduction": "Neutron stars are born in spectacular core collapse supernova explosions. These compact, massive objects have typical radii of about 10 km and masses of $(1-2)M_\\odot$. Matter in the core of neutron stars is compressed to extreme densities, several times normal nuclear matter density, i.e.\\ $\\rho\\gg \\rho_0 = 3\\cdot 10^{14}$~g/cm$^3$. The relation of supernova explosions being the birthplace of neutron stars is exemplified by the historic supernova remnant of AD 1054, the crab nebula, and the crab pulsar, a rotation-powered neutron star, sitting in its center.  More than 1700 pulsars are recorded in the publicly available pulsar data base at the Australian National Telescope Facility \\cite{Manchester:2004bp}. The number of discovered pulsars is continuously growing with the ongoing pulsar surveys at radio telescopes worldwide (Arecibo, Green Bank Telescope, Parkes Multibeam). The best determined mass is still the one of the Hulse-Taylor pulsar with $M=(1.4411\\pm0.00035)M_\\odot$ \\cite{Weisberg:2004hi}, the fastest rotating one is the pulsar PSR J1748-2446ad with 716 revolutions per second \\cite{Hessels:2006ze}.  Recently, indications of an extremely massive neutron star have been published \\cite{Freire:2007jd,Freire:2007ux}. If the measured periastron advance is due to effects from general relativity only, the pulsar PSR J1748-2021B has a mass of $M\\geq (2.74 \\pm 0.21) M_\\odot$ and is more massive than $M\\geq 2.0 M_\\odot$ with a confidence level of 99\\%. Note, that it can not be excluded at present that the system measured actually contains two neutron stars, not a single one. Redshifted spectral lines have been claimed to be extracted from the analysis of x-ray bursts from EXO 0748--676 \\cite{Cottam:2002cu}, which give a constraint on the mass-radius ratio of the compact star. A recent analysis comes to the conclusion that the compact star mass is $M\\geq (2.10\\pm 0.28) M_\\odot$ with a radius of $R\\geq (13.8 \\pm 1.8)$~km \\cite{Ozel:2006km} claiming that 'unconfined quarks do not exist at the center of neutron stars'! However, this conclusion was put into perspective in a follow-up reply \\cite{Alford:2006vz} which demonstrated that those limits rule out soft equations of state, but not quark stars or hybrid stars. The interactions between quarks can be quite strong so that the presence of quark matter in the core stabilises the compact star. On the other hand, the mass limit provides indeed a strong constraint for hyperons in dense neutron star matter. Hyperons are likely to appear at moderate densities, which will substantially decrease the maximum mass. This conclusions is guided by hypernuclear data and present model calculations. If such massive neutron stars are confirmed in the future, say with masses above $2M_\\odot$, then it seems that our present understanding of hypernuclear physics of compact stars will be in conflict with the pulsar data, as we will outline in more detail below.  Constraints on the mass and radius of neutron stars can be derived by observations in the optical as well as in the x-ray band, a booming field of exploration since the launch of the x-ray satellites Chandra and XMM-Newton in 1999. The best studied isolated neutron star is RXJ 1856.35-3754, the closest one known. A two-component blackbody fit to the combined optical and x-ray spectra results in a low soft temperature, so as not to be in contradiction with the observed x-ray flux. This low temperature implies a rather large radiation radius, the radius observed at infinity, $R_\\infty = R/\\sqrt{1-2GM/R}$, so that the optical flux comes out right. A conservative lower limit was given in \\cite{Trumper:2003we} and confirmed in a detailed modelling of the neutron star atmosphere \\cite{Ho:2006uk,Ho:2007aw} as being $R_\\infty \\approx 17$~km for an assumed distance of $d=140$~pc. With the derived gravitational redshift of $z_g\\approx 0.22$, the true radius of the neutron star would be about $R\\approx 14$~km with a corresponding mass of $M\\approx 1.55M_\\odot$.  The large radiation radius implies generally a large neutron star radius which could only be explained with a stiff nuclear equation of state.  The biggest uncertainty, besides the systematical one, is the distance to the neutron star.  The spectra of the neutron star X7 in the Globular Cluster 47 was fitted in Ref.~\\cite{Rybicki:2005id} with an improved hydrogen atmosphere model.  The radius of the neutron star was estimated to be $R_{\\rm ns}=(14.5^{+1.8}_{-1.6})$~km, which is a little bit larger than for non-relativistic nuclear models but right in the range of standard relativistic mean-field models, see e.g.~\\cite{Lattimer:2006xb}.  The authors of Ref.~\\cite{Rybicki:2005id} state, that for a radius of 10 km the mass should be in the range $M_{\\rm ns}=(2.20^{+0.03}_{-0.16}) M_\\odot$.  For a radius of 14~km, however, any mass between 0.5 and $2.3M_\\odot$ is allowed by the fit.  On the other hand, atmosphere fits to the spectra of M13 lead to rather small radii, a radius of only $R=9.77^{+0.09}_{-0.29}$~km was derived in Ref.~\\cite{Webb:2007tc}.  The allowed ranges in the mass-radius diagram for the fit to the spectra of M13 and X7 are nearly mutually exclusive on the 99\\% confidence level. However, one should keep in mind, that the whole mass-radius curve for neutron stars just has to reach somewhere those two regions. In fact, many of the mass-radius curves shown in \\cite{Webb:2007tc} pass the two constraints from the fits to the spectra of M13 and X7, except for the curves of the most stiffest models, in particular for the relativistic field theoretical models without hyperons.  Another way of probing neutron star matter properties is by cooling observations of supernova remnants, see e.g.\\ \\cite{Kaplan:2004tm,Kaplan:2006mb}. The observational limits points towards fast cooling processes in the interior of neutron stars, i.e.\\ direct URCA reactions. Standard conventional cooling curves are too high, so that either a large nuclear asymmetry energy or strange exotic particles are needed to generate efficient and fast cooling (see \\cite{Yakovlev:2004iq} for a theoretical review).  The basic structure of the low-density region of neutron stars is fairly well-known. The outer crust consists of a lattice of nuclei with free electrons and is a few 100 meters thick. The sequence of nuclei is controlled by their binding energies and follows mainly along the neutron magic numbers 50 and 82 (for a most recent investigation of the outer crust see \\cite{Ruster:2005fm}). Similar features will be discussed in the context of hypernuclei below. The inner crust starts at the neutron drip density at $n\\approx 4\\cdot 10^{11}$~g/cm$^3$ and consists of a lattice of nuclei with free neutrons and electrons. The core starts at the end of the inner crust which occurs around half times normal nuclear matter density. In this core region, hyperons can populate the interior of neutron stars. The implications of the presence of hyperons for the properties of neutron star will be outlined in this review, which is an update and an extension of a preliminary version of Ref.~\\cite{SchaffnerBielich:2007tj}. ",
        "conclusions": "Hyperons have a substantial impact on neutron star properties. There is a sizable decrease in the maximum mass of neutron stars due to the presence of hyperons in the core. The $\\Lambda$ hyperons appear at $n\\approx 2n_0$ in neutron star matter.  The population of $\\Sigma$ hyperons hinges crucially on their in-medium potential. They are likely to be absent for a repulsive potential, but the negatively charged $\\Sigma^-$ could be the first exotic component in neutron star matter for an attractive potential.  A tiny amount of hyperons can suffice to cool neutron stars rapidly by the hyperon direct URCA process, which is controlled by hyperon pairing gaps. A strongly attractive YY interaction, between $\\Xi$ hyperons, results in a first order phase transition from neutron-rich to hyperon-rich matter. This transition allows for a new, stable solution for compact stars, hyperon stars, with similar masses but smaller radii. The presence of the nonmesonic weak decay reactions with hyperons in neutron stars determines the bulk viscosity of neutron star matter and leads to an enhanced stability window with regard to r-modes of pulsars.  It is obvious, that hypernuclear physics provides essential input for compact star physics. The YN interactions, in particular the potential depth in bulk nuclear matter, controls the population of hyperons for massive neutron stars, the first exotic component likely to appear for supranuclear densities present in the core. The emergence of hyperons softens the nuclear equation of state and the maximum neutron star mass possible considerably which depends on the YN coupling strength and sensitively on the hyperon three-body forces.  Two-body YY interactions regulate the cooling behaviour of massive neutron stars, as the hyperon direct URCA reaction is suppressed by hyperon gaps. Nonmesonic weak decay of hyperons in the dense medium as well as hyperon superfluidity controls the bulk viscosity of neutron star matter which regulates the stability of r-modes of pulsars and the emission of gravitational waves. In addition, hyperons can generate a new class of compact stars, hyperon stars, for a suitably attractive YY potential. The ongoing and future experimental hypernuclear programs at DA$\\Phi$NE, Jefferson Lab, KEK, J-PARC, MAMI-C, and at GSI, Darmstadt, in particular the HypHI program and HYPER-GAMMA with PANDA at FAIR, will provide here the decisive inputs for addressing the macrophysics and microphysics of neutron stars as hyperons play such an important role for many compact star observables."
    },
    "0801/0801.3758_arXiv.txt": {
        "abstract": "We discuss the classic theorem according to which a gravitational lens always produces at least one image with a magnification greater than unity. This theorem seems to contradict the conservation of total flux from a lensed source.  The standard solution to this paradox is based on the exact definition of the reference `unlensed' situation, in which the lens mass can either be removed or smoothly redistributed.  We calculate magnifications and amplifications (in photon number and energy flux density) for general lensing scenarios not limited to regions close to the optical axis. In this way the formalism is naturally extended from tangential planes for the source and lensed images to complete spheres. We derive the lensing potential theory on the sphere and find that the Poisson equation is modified by an additional source term that is related to the mean density and to the Newtonian potential at the positions of observer and source.  This new term generally reduces the magnification, to below unity far from the optical axis, and ensures conservation of the total photon number received on a sphere around the source.  This discussion does not affect the validity of the \\emph{focusing theorem}, in which the unlensed situation is defined to have an unchanged affine distance between source and observer. The focusing theorem does not contradict flux conservation, because the mean total magnification (or amplification) directly corresponds to different areas of the source (or observer) sphere in the lensed and unlensed situation. We argue that a constant affine distance does not define an \\emph{astronomically} meaningful reference.  By exchanging source and observer, we confirm that magnification and amplification differ according to Etherington's reciprocity law, so that surface brightness is no longer strictly conserved. At this level we also have to distinguish between different surface brightness definitions that are based on photon number, photon flux, and energy flux. ",
        "introduction": "Already \\citet{einstein36} calculated the magnification of a background star lensed by a foreground compact mass and found the classic result for the total magnification of both images, which decreases for large impact angles (far away from the optical axis) but always stays above unity.  At that time, gravitational lensing was regarded as a theoretical curiosity with very little chance of leading to \\emph{observable} effects. Much later, after the first actual lens systems had been found, \\citet{schneider84} presented and proved a theorem stating that a gravitational lens with arbitrary mass distribution will always produce at least one image with a magnification greater than 1.  Since the theorem is thought to be valid for all configurations of source, lens and observer, one might naively think that this necessarily leads to an increase of the total number of photons from the source. This would be in contradiction with the fact that lensing can neither create nor destroy photons.  It was already pointed out by \\citet{schneider84} that such an interpretation is not correct. In order to compare the lensed with the unlensed situation, it is essential to define the latter in a meaningful way, which is far from trivial. In general relativity, the lens will slightly distort the geometry of the whole Universe, so that it is plainly impossible to add a lens without changing the configuration of source and observer. The `reference situation' in the standard explanation is the one in which the mass of the lens is redistributed evenly over the Universe.  It can be shown that the comparison finds no net increase of photon number then \\citep[e.g.][]{weinberg76}. The reason for this is that two differently lensed scenarios are compared.  Some level of uneasiness remained in the minds of several authors, because one may think of situations in which the general validity of the theorem would still lead to a violation of photon number conservation.  \\citet{avni88} calculated the amplification caused by a point-mass lens by dropping the small-angle approximation of light rays with respect to the optical axis (defined by observer and lens), which is generally used in the standard formalism.  One might think that far from the optical axis the deflection becomes so weak that the full calculation would add only higher-order terms which are not relevant in the end. Still, slightly surprising, these authors find that the amplification actually sinks below 1, which corrects for the amplification excess close to the optical axis.  \\citet{avni88} worked with a point-mass metric in Schwarzschild coordinates, and fixed the radial coordinate relative to the lens for the comparison of lensed and unlensed situation.  \\citet{sef} question the value of such a comparison. They argue that there is generally no unique way to compare mean fluxes in different space-times.  Nevertheless, approximating the calculations of \\citet{avni88} for regions close to the optical axis would directly lead to the tangential plane approach of \\citet{schneider84}. It is not entirely clear why (if at all) the theorem loses its validity far from the optical axis.  \\citet{jaroszynski96} make similar calculations, but keep the \\emph{metric distance} between source and observer constant for the comparison. They find that the effect of the lens changes the total area of the sphere around the source and can thus (without contradicting photon conservation) lead to changes in the averaged amplification. They also make a distinction between \\emph{magnification} and \\emph{amplification}, which becomes relevant at the accuracy levels required far from the optical axis.  None of these publications actually presents the magnification to second order (in terms of the deflection potential) for arbitrary angles to the optical axis\\footnote{Close to the optical axis, the second-order terms become dominant for compact masses.}. They either approximate for small angles or include only first-order terms.  One aim of our work is to extend the calculations to include second-order terms also for large impact angles.  Even more important is the discussion of appropriate unlensed situations to act as a reference in order to define magnifications and amplifications.  We will also develop the formalism to treat off-axis lenses similarly to the standard small-angle approximation and learn which part of the proof of the magnification theorem cannot be generalised for this situation. This will be done for arbitrary mass distributions and not only for point-mass lenses.   The outline of this paper is as follows.  We start in Sec.~\\ref{sec:planar} with a brief review of the standard planar lensing theory, which is described in a tangential plane at the position of the lens. This includes the magnification theorem and a simple version of its proof. In Sec.~\\ref{sec:solid} we explain our approach of discussing the apparent magnification in terms of solid angles. In this way we can avoid the explicit definition of an unlensed reference situation. Solid angle must always integrate to $4\\upi$, so that the magnification theorem (in these terms) cannot hold everywhere. This discussion is equivalent to a reference situation in which the area of the source sphere stays constant.  Section~\\ref{sec:point-mass} comprises calculations for a point-mass lens at infinity. We explicitly find that the magnification sinks below 1 far from the optical axis.  General mass distributions (still with source at infinity) are discussed in Sec.~\\ref{sec:general mass sphere}. There we develop the potential theory of lenses on the sphere as an extension of the standard potential theory in the tangential plane. We find that the Poisson equation is modified by an additional source term, which ensures the validity of Gauss' theorem on the closed sphere.  The exact form of magnification matrices on the sphere for finite (and possibly large) deflection angles is investigated in Appendix~\\ref{sec:mag sph}. For our current discussion we only need the result that the magnification matrix in terms of second-order derivatives of the potential differs from the usual form only to second order in the deflection angle, so that this difference can be neglected.  In Sec.~\\ref{sec:finite ds} we extend our calculations to finite values of the source distance and explicitly calculate deflection angles and magnifications for point-mass lenses. By adapting the potential theory to this situation, we find that the correction terms in Poisson equation and magnification are directly related to the Newtonian potential at the source and observer.  The subject of reciprocity is covered in Sec.~\\ref{sec:recip}, where we exchange the role of source and observer to relate \\emph{magnification} with \\emph{amplification}. These two are \\emph{not} the same anymore, so that the surface brightness is no longer conserved if we define it in a standard physical sense. This is a direct consequence of the reciprocity theorem of \\citet{etherington33} and of the space-time modification at the position of source and observer, caused by the gravitational field of the lens. In anticipation of this result we try to make a clear distinction between amplification and magnification from the beginning.  After a brief discussion of light travel times in Sec.~\\ref{sec:light time} we summarise and discuss our results in Secs.~\\ref{sec:summary} and \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  The total magnification excess obtained from a naive interpretation of the magnification theorem is of the same order of magnitude as the perturbation of the metric at the position of observer or source caused by the gravitation of the lens. For a discussion of the apparent paradox, terms of this order cannot be neglected. This means that the lensing process can no longer be described in a tangential plane, but full spheres have to be used for the lensed images and the unlensed source. Extending the planar formalism naturally leads to additional terms, which compensate the magnification excess and lead to a mean magnification of exactly 1. This is the case for an unlensed reference situation in which the source sphere has the same area as the lensed source sphere. The fact that total flux (or solid angle) is conserved in a formalism which is based on this very assumption is clearly tautological. However, without the additional spherical correction terms, this consistency could not have been achieved.  It is a very satisfying result that generalising the formalism in order to assure total flux conservation automatically avoids a violation of Gauss' theorem for the deflection angle on the sphere, by modifying Poisson's equation for the deflection potential. It is found that the source function for the deflection is not the surface mass density itself, but the density \\emph{contrast} relative to the mean density. This result is entirely plausible, because a constant surface mass density can, for symmetry reasons, not lead to a light deflection.  In our formalism, magnifications are defined as directly being caused by light deflection and not by the indirect effect of metric perturbations at the positions of source or observer. It is thus clear that a constant surface mass density cannot lead to magnifications, and that the magnification theorem does not hold on the complete sphere.  In contrast to this, a different comparison situation is explicitly defined for the \\emph{focusing theorem} \\citep[e.g.][]{sef,perlick04}, in which the affine distance is kept the same in the lensed and unlensed situation. With this definition, the theorem is \\emph{always} valid, even under much more general circumstances. However, we argue that keeping the affine distance constant is not the appropriate way to compare lensed and unlensed sources in an astronomical context.  In addition we learn that at the level of accuracy needed for a discussion of the magnification theorem(s), we also have to distinguish between magnification and amplification, and that for the latter we have to decide if we define it in terms of total photon number, photon flux, total energy or energy flux. This distinction is well known in the context of cosmological redshifts. As a result of these effects, surface brightness is generally \\emph{not} conserved."
    },
    "0801/0801.4387_arXiv.txt": {
        "abstract": "We investigate the influence of magnetic fields upon the dynamics of, and resulting gravitational waves from, a binary neutron star merger in full general relativity coupled to ideal magnetohydrodynamics~(MHD). We consider two merger scenarios, one where the stars have aligned poloidal magnetic fields magnetic fields and one without.  Both mergers result in a strongly differentially rotating object. In comparison to the non-magnetized scenario, the aligned magnetic fields delay the full merger of the stars.   During and after merger we observe phenomena driven by the magnetic field, including Kelvin-Helmholtz instabilities in shear layers, winding of the field lines, and transition from poloidal to toroidal magnetic fields. These effects not only mediate the production of electromagnetic radiation, but also can have a strong influence on the gravitational waves. Thus, there are promising prospects for studying such systems with both types of waves. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.4452_arXiv.txt": {
        "abstract": "{The planets Uranus and Neptune with small apparent diameters are primary calibration standards.} {We investigate their variability at $\\sim90\\,$GHz using archived data taken at the IRAM 30m telescope during the 20\\,years period 1985 to 2005.} {We calibrate the planetary observations against non-variable secondary standards (NGC\\,7027, NGC\\,7538, W3OH, K3-50A) observed almost simultaneously.  } {Between 1985 and 2005, the viewing angle of Uranus changed from south-pole to equatorial. We find that the disk brightness temperature declines by almost 10\\% ($\\sim2\\sigma$) over this time span indicating that the south-pole region is significantly brighter than average. Our finding is consistent with recent long-term radio observations at 8.6\\,GHz by Klein \\& Hofstadter (2006).  Both data sets do moreover show a rapid decrease of the Uranus brightness temperature during the year 1993, indicating a temporal, planetary scale change.  We do not find indications for a variation of Neptune's brightness temperature at the 8\\% level.} {If Uranus is to be used as calibration source, and if accuracies better than 10\\% are required, the Uranus sub-earth point latitude needs to be taken into account. } ",
        "introduction": "At mm-wavelengths, the planets Uranus and Neptune with small apparent diameter are frequently used for calibration of astronomical sources and telescope parameters. In this context it is often tacitly assumed that they are constant radiators, at least at short term intervals.  Observations over a decade or a longer period, preferably with the same telescope and receivers and traceable modifications, may reveal a long--term variability of Uranus and Neptune.  For a long time, Uranus and Neptune have been separated by a distance of $\\sim$\\,22$^{\\rm o}$ or less in the sky and thus allow precise comparative measurements.  \\citet{griffin1993} observed Uranus and Neptune in 1990 and 1992 to derive their brightness temperature in the millimeter and submillimeter regime.  However, recent microwave as well as observations at visible and near-infrared wavelengths indicate that Uranus is variable on time scales of several years \\citep[see e.g.][KH06 in the following]{kh06}.  Due to Uranus' large obliquity of $82^\\circ$, measurements from the earth alternate between observations of the poles and the equator, i.e. the sub-earth point (SEP) latitude varies.  KH06 argue that the variations they observed at 3.5\\,cm (8.6\\,GHz) are partly caused by the geometrical effect, but partly also by temporal variations deep inside the Uranus troposphere. However, this variability has not yet been studied at millimeter wavelengths.   We investigate whether these long-term variations described above are noticeable in 20 years pointing observations \\footnote{This collection of pointing measurements provided also the basis for several compilations of quasar flux densities published by \\citet{steppe1988,steppe1992,steppe1993} and \\citet{reuter1997}. Unfortunately, most of the data of 1988 and 1989 are lost, for unknown reason.}  made with the IRAM 30m telescope at 90\\,GHz. By relating the observations of the planets to nearly simultaneous observations of the constant secondary calibrators NGC\\,7027, NGC\\,7538, W3OH and K3--50A, the measurements of the planets are free from changes of the telescope performance (reflector adjustments, receiver upgrades, etc.) but contain, on the other hand, the errors of the secondary measurements.  The secondaries are small relative to the half power beam width (HPBW) of $27''$ at 90\\,GHz. ",
        "conclusions": "By careful selection of data from the archived pointing measurements made at the IRAM 30m telescope during the period 1985 to 2005, we have been able to study the long-term behaviour of Uranus and Neptune at 90\\,GHz. The large number of selected observations allowed a statistically meaningful analysis.  \\begin{enumerate}  \\item We obtained 286 observations of Uranus. The scatter of the derived normalized brightness temperatures is 6\\%.  We find a systematic variation of the Uranus brightness temperature of $\\sim10$\\% with sub-earth point (SEP) latitude when the orientation of Uranus changes from south-pole view to equator view. This effect is detected at the $\\sim2\\sigma$ level, the mean scatter of the data averaged in $15^\\circ$ bins is 5\\%. A similar change is indicated by the 8.6\\,GHz observations of KH06.  The year 1993 shows a rapid decrease of brightness temperatures at 90\\,GHz of $\\sim7\\%$, which corresponds to a similar decrease seen by KH06, indicating a temporal change of a large fraction of the Uranus atmosphere.  More observations at other frequencies and improved models of the Uranus atmosphere, taking into account its latitude structure, are needed to better determine the variations with SEP.  If Uranus is to be used as calibration source at 90\\,GHz, the latitude dependence of its brightness temperature needs to be taken into account.  \\item For Neptune, we only obtained 71 observations tied to secondary calibrators and 81 simultaneous observations with Uranus. Both show that its brightness temperature stays constant at a level of better than 8\\%.  \\end{enumerate}"
    },
    "0801/0801.4208_arXiv.txt": {
        "abstract": "We discuss a Lagrangian reconstruction method of the velocity field from galaxy redshift catalog that takes its root in the Euler equation. This results in a ``functional'' of the velocity field which must be minimized. This is helped by an algorithm solving the minimization of cost-flow problems. The results obtained by applying this method to cosmological problems are shown and boundary effects happening in real observational cases are then discussed. Finally, a statistical model of the errors made by the reconstruction method is proposed. ",
        "introduction": "Introduction}  Cosmologists are highly interested in studying galaxy peculiar velocities. Indeed, their study is a direct way to measure the dynamical state of a system and would thus permit to better understand dark matter distribution in our local Universe. The main difficulty is that measured velocities are only available sparsely and hence does not provide a good probe of the matter distribution. One must then devise an algorithm that is able to predict, under fair hypotheses, galaxy peculiar velocities from their present positions, which are their sky coordinates and their redshift, and compare the result to the measurement. Jim Peebles \\cite{Peebles89} first tried to do full orbit reconstruction by evolving the present system back in time. This method proved to be quite accurate for very small volume and number of objects. However, whenever one tries to reconstruct orbits of a large number of galaxies, the method fails because the number of plausible solution is blowing up. A simplification of this problem is presented: 3D galaxy positions are assumed to be known and a simpler gravitational dynamic model is going to be assumed. We will also assume that the dynamics of galaxies is mostly driven by collisionless dark matter particles.  This proceeding is organized as follows. In \\S~\\ref{sec:recons}, we recall the principal result of the reconstruction method developed by \\cite{Brenier2002} (see also the companion paper Mohayaee \\& Sobolevskii, hereafter MS). The method requests the using of a special fast algorithm to solve the problem. This algorithm is presented in \\S~\\ref{sec:auction}. The method is then applied to a dark matter distribution obtained from a cosmological simulation and the reconstructed velocities are checked against the simulated ones (\\S~\\ref{sec:apply_cosmo}). Finally, a discussion on problems with bad boundary conditions, as usually met in observational cosmology, is quickly discussed in \\S~\\ref{sec:boundary}. ",
        "conclusions": ""
    },
    "0801/0801.3334_arXiv.txt": {
        "abstract": "\\vspace{1cm} We study how the properties of the four neutralino states, $\\chi_i$ (i = 1, 2, 3, 4), can be investigated at the Large Hadron Collider (LHC), in the case when the lightest one, $\\chi_1$, has a mass $m_{\\chi} \\lsim 50$ GeV and is stable. This situation arises naturally in supersymmetric models where gaugino masses are not unified at a Grand Unified (GUT) scale and R-parity is conserved.  The main features of these neutralino states are established by analytical and numerical analyses, and two scenarios are singled out on the basis of the cosmological properties required for the relic neutralinos.  Signals expected at LHC are discussed through the main chain processes started by a squark, produced in the initial proton-proton scattering.  We motivate the selection of some convenient benchmarks, in the light of the spectroscopical properties (mass spectrum and transitions) of the four neutralino states. Branching ratios and the expected total number of events are derived in the various benchmarks, and their relevance for experimental determination of neutralino properties is finally discussed. ",
        "introduction": "\\label{sec:intro}  Light neutralinos, {\\it i.e.} neutralinos with a mass $m_{\\chi}  \\lsim 50$ GeV, arise naturally in supersymmetric models where gaugino masses are not unified at a Grand Unified (GUT) scale. When R-parity conservation is assumed, these neutralinos offer a very rich phenomenology under various cosmological and astrophysical aspects \\cite{lowneu,lowdir}.  In the present paper we study how these light neutralinos  can be investigated at the Large Hadron Collider (LHC), which will soon start being operated at CERN \\cite{atltdr,cmstdr2}.  To this purpose we first delineate what are the main properties of the four neutralino states, $\\chi_i$ (i = 1, 2, 3, 4), in the case when the lightest one, $\\chi_1$ (or $\\chi$ in short), has a mass $m_{\\chi}  \\lsim 50$ GeV. This neutralino is stable, in case it occurs to be  the Lightest Supersymmetric Particle and R-parity is conserved. In particular, special asymptotic schemes for the four neutralino states are analysed: two hierarchical schemes (with so-called normal and inverted hierarchy, respectively) and an almost degenerate one.  Once the properties of  all the neutralino states are established by analytical and numerical analyses, two main scenarios are  singled out  on the basis of the cosmological properties required for the relic neutralinos.  Then, we analyse the signals expected at LHC, in terms of the main chain processes which are initiated by a squark and lead finally to $\\chi_1$ through an intermediate production of $\\chi_i$ (i = 2, 3, 4). We evaluate the branching ratios for the various processes and discuss their features in terms of the spectroscopic properties (mass spectrum and transitions) of the four neutralino states. Various benchmarks of special physical interest are considered.  The total number of events expected for these processes at LHC is estimated for a typical representative value of the integrated luminosity and their relevance for experimental determination of neutralino properties is finally discussed.   The scheme of our paper is the following. The features of the employed supersymmetric model are presented in Sect. \\ref{sec:susy}, where also the main structural properties of the four neutralino states are described. In Sect. \\ref{sec:scenarios} we delineate the scenarios imposed by cosmology on light relic neutralinos. In Sect. \\ref{sec:signals} we derive branching ratios and the total number of events expected at LHC. Finally, conclusions are drawn in Sect. \\ref{sec:conclusions}. ",
        "conclusions": "  i) The scenario $\\mathcal{A}$ should be easily explorable both through sequential decay chains and through branched chains. For the sequential case good perspectives are offered by the $\\bar{e} e$ (or $\\bar{\\mu} \\mu$) signal and by the $\\bar{\\tau} \\tau$ one. For the branched case the scenario  $\\mathcal{A}$ gives good measurement perspectives for production of a $\\bar{b} b$ pair through Z and h, A Higgs bosons, combined with light lepton pairs through Z.  ii) In scenario $\\mathcal{B}$ large signals are expected in terms of $\\bar{e} e$  (or $\\bar{\\mu} \\mu$) and  $\\bar{\\tau} \\tau$ pairs in sequential processes. High statistics will be required to explore the scenario  $\\mathcal{B}$ by branched decay; the most favorable processes being represented by those mediated by Z or h, A Higgs bosons when $M_2 > |\\mu|$.  These results show that LHC has a strong potential in the investigation  of the supersymmetric parameter region compatible with a light neutralino. Due to the characteristic features of this region, the measurements of LHC should easily prove or disprove our model.  Finally, we wish to recall that our conclusions  are based on the evaluated total number of events, only. To ascertain the actual potentiality of LHC in the study of our model, the investigation has to be pursued to include  analysis of the signal/background ratios and of specific kinematical distributions. This further investigation will be presented in a subsequent publication \\cite{future}."
    },
    "0801/0801.3599.txt": {
        "abstract": "We present results from a statistical analysis of 173 bright radio-quiet AGNs selected from the {\\sl Chandra} Deep Field-North and {\\sl Chandra} Deep Field-South surveys (hereafter, CDFs) in the redshift range of 0.1~$\\lesssim~z~\\lesssim~$4. We find that the X-ray power-law photon index ($\\Gamma$) of radio-quiet AGNs is correlated with their 2--10~keV rest-frame X-ray luminosity ($L_X$) at the $>$~99.5\\% confidence level in two redshift bins, $0.3~\\lesssim~z~\\lesssim~0.96$, and $1.5~\\lesssim~z~\\lesssim~3.3$ and is slightly less significant in the redshift bin $0.96~\\lesssim~z~\\lesssim~1.5$. The X-ray spectral slope steepens as the X-ray luminosity increases for AGNs in the luminosity range $10^{42}$ to $10^{45}$ \\lumin. Combining our results from the CDFs with those from previous studies in the redshift range $1.5~\\lesssim~z~\\lesssim~3.3$, we find that the $\\Gamma-L_{\\rm X}$ correlation has a null-hypothesis probability of 1.6 $\\times10^{-9}$. We investigate the redshift evolution of the correlation between the power-law photon index  and the hard X-ray luminosity  and find that the slope and offset of a linear fit to the correlation change significantly (at the $>$~99.9\\% confidence level) between redshift bins of $0.3~\\lesssim~z~\\lesssim~0.96$ and $1.5~\\lesssim~z~\\lesssim~3.3$. We explore physical scenarios explaining the origin of this correlation and its possible evolution with redshift in the context of steady corona models focusing on its dependency on variations of the properties of the hot corona with redshift. ",
        "introduction": "It is important to extend the study of quasars to high redshifts in order to understand their evolution and environments. A relevant conclusion from modern studies is that the quasar luminosity function evolves positively with redshift, having a comoving space density strongly peaked at $z \\approx 2$ \\citep[e.g.,][]{Sch68,Boy87,War94}. More recent findings suggest that the evolution of the space density of AGNs is strongly dependent on X-ray luminosity ($L_{\\rm X}$), with the peak space density of AGNs moving to higher redshifts for more luminous AGNs \\citep[e.g.,][]{Ued03,Has05}.  The X-ray band probes the innermost region of the central engines of AGNs. The study of AGNs in the X-ray band provides important insights about their central engines and the evolution of the AGN luminosity function. In most AGNs, the observed X-ray continuum can be modeled using a power-law of the form $N(E) = N_{0}(E/E_{0})^{-\\Gamma}$, where $\\Gamma$ is the photon index. This power-law is attenuated by material in our Galaxy as well as material intrinsic to the host galaxy. Several recent studies have centered on estimating the distribution of intrinsic column densities ($N_{\\rm H}$) and the fraction of AGNs having $N_{\\rm H} \\gtrsim 10^{22}$ \\cmsq. Recent theoretical studies of AGNs \\citep[e.g.,][]{Hop05} suggest that the distribution of $N_{\\rm H}$ is luminosity dependent; this is supported observationally with the detection of an anti-correlation between the obscuration fraction and luminosity \\citep[e.g.,][]{Stef03,Ued03,LaF05,Aky06}. We note, however, recent work by \\cite{Dwe06} reporting that the obscuration fraction may be independent of luminosity. The dependence of the obscuration fraction on redshift is a controversial issue. Some authors detect an increase of the obscuration fraction with redshift \\citep[e.g.,][]{LaF05,Tre06,Toz06}, while others do not find any evidence for evolution \\citep[e.g.,][]{Ued03,Aky06,Dwe06}.  A recent mini-survey of relatively high-redshift $(1.5 < z < 4)$ gravitationally lensed radio-quiet quasars (RQQs) observed with {\\it Chandra} and {\\it XMM-Newton} \\citep{Dai04} indicated a possible correlation between the X-ray power-law photon index and X-ray luminosity. This correlation, characterized by an increase of $\\Gamma$ with $L_{\\rm X}$, was found for RQQs with 2--10~keV luminosities in the range $10^{43}$ to $10^{45}$ \\lumin. Such a correlation is not found in nearby $z \\lesssim 0.1$ quasars \\citep[e.g.,][]{Geo00}. Several studies to date of high-redshift quasars do not have large enough sample sizes in the 2--10~keV luminosity range $10^{43}$ to $10^{45}$ \\lumin~ to place any significant constraints on a possible $\\Gamma-L_{\\rm X}$ correlation \\cite[e.g.,][]{Ree00,Pag05}.  One of the concerns with the \\cite{Dai04} analysis was that the limited number of quasars in the sample, combined with the poor signal-to-noise ratio (S/N) available for several of the observations and the relatively large fraction of BAL quasars, may have led to problematic systematic effects. The number of available lensed radio-quiet quasars used by \\cite{Dai04} was limited to a total of 25 sources, of which the brightest 11 had X-ray observations. In order to increase the size of the high-redshift radio-quiet quasar sample, we have compiled a sample of 173 high-redshift AGNs with moderate-to-high S/N spectra available from the \\chandra\\ Deep-Field-North and \\chandra\\ Deep-Field-South surveys \\citep[\\cdfn$ $ and \\cdfs, respectively; jointly CDFs;][]{Gia02,Ale03}.   The main scientific goal of this work is to constrain better the $\\Gamma-L_{\\rm X}$ correlation found by Dai et al. (2004). The significant increase in sample size allows us to place tighter constraints on the significance of the correlation. We also test the correlation in narrower redshift bands which will allow us to determine the epoch by which possible changes in the average emission properties of AGNs occurred. Currently, the two deepest X-ray surveys are the \\cdfn$ $ and \\cdfs$ $ with $\\approx$2~Ms and $\\approx$1~Ms exposures, respectively. Both surveys cover $\\approx$300~arcmin$^2$ areas and target different regions of the sky characterized by low Galactic column densities and an absence of bright stars \\citep{Gia02,Ale03}. The CDFs pointings have sufficient sensitivity to detect the X-ray emission from AGNs with moderate luminosities ($L_{\\rm X} \\approx10^{43}-10^{44}$ \\lumin) out to $z\\approx2-6$.   Radio-quiet AGNs (RQ AGNs) correspond to the majority of active galaxies $(\\sim90\\%)$ that contain a central active nucleus and show several differences in their spectral properties compared to radio-loud AGNs. Radio-loud AGNs have powerful sub-parsec jet-linked X-ray synchrotron self-Compton (SSC) emission, which introduces an additional component to their spectra. As a consequence RQ AGNs are observed to have, on average, steeper X-ray power-laws than radio-loud AGNs \\citep[e.g.,][]{Ree97}. We therefore have chosen to exclude radio-loud AGNs from this study. Throughout this paper we adopt a flat $\\Lambda$-dominated universe with $H_0$=70~km~s$^{-1}$Mpc$^{-1}$, $\\Omega_\\Lambda=0.7$, and $\\Omega_M=0.3$. The \\chandra\\ data were reduced using the CIAO version 3.3 software tools provided by the {\\sl Chandra} X-ray Center (CXC), and the spectral analysis was performed using XSPEC version 12.   \\begin{figure} \\epsscale{1} \\plotone{f1.eps} \\centering \\caption{(upper panel) Number of sources with more than $S$ photons (\\oOF) versus $S$.  The thick line corresponds to all CDFs sources with measured redshifts of $z\\gtrsim0.1$. The dotted line shows sources of the CDF-N survey with $z\\gtrsim 0.1$, and the dashed line shows sources of the CDF-S survey with $z\\gtrsim 0.1$. The vertical dotted line corresponds to $S$=170. Note that the sample used to generate this figure contains AGNs (both RQ and radio-loud AGNs), normal galaxies and starburst galaxies. (lower panel) Number of radio-quiet AGNs with $z\\gtrsim 0.1$ vs. the number of photons ($S$; 0.5--8~keV) in their spectra. The solid line represents sources with fits performed in the \\oOF $ $ band and the dotted line represents sources with fits performed in the \\rRF $ $ band.} \\label{cnts} \\end{figure} ",
        "conclusions": ""
    },
    "0801/0801.1107_arXiv.txt": {
        "abstract": "A white dwarf (WD) gains substantial angular momentum during the accretion process that grows it toward a Chandrasekhar mass. It is therefore expected to be quickly rotating when it ignites as a Type Ia supernova. The thermal and shearing profile are important for subsequent flame propagation. We highlight processes that could affect the WD shear, during accretion as well as during the $\\sim1000$ years of pre-explosive simmering. Baroclinic instabilities and/or the shear growth of small magnetic fields provide sufficient torque to bring the WD very close to solid body rotation during accretion. The lack of significant shear makes it difficult to grow a WD substantially past the typical Chandrasekhar mass. Once carbon ignites, a convective region spreads from the WD's center. This phase occurs regardless of progenitor scenario, and therefore it is of great interest for understanding how the WD interior is prepared before the explosive burning begins. We summarize some of the key properties of the convective region, which includes demonstrating that the mass enclosed by convection at any given time depends most sensitively on a single parameter that can be expressed as either the ratio of temperatures or densities at the top and bottom of the convection zone. At low Rossby numbers the redistribution of angular momentum by convection may result in significant shearing at the convective/non-convective boundary. ",
        "introduction": "The fundamental role Type Ia supernovae (SNe Ia) play in determining the expansion history of the universe \\citep[][and references therein]{rie04} has added new urgency to understanding the progenitors of these events. Theoretical work on modeling the flame dynamics of exploding white dwarfs (WDs) has demonstrated that the composition and energy of ejecta depend sensitively on the competition between flame propagation, instabilities driven by turbulence, and expansion of the WD \\citep[e.g.,][]{hn00}. This has sparked new interest in understanding the evolution that occurs up until the onset of explosive burning. A critical aspect is the shear profile present within the WD.  SNe Ia are expected to occur exclusively in accreting binary systems. Therefore the role of angular momentum may be unique or at least especially important for these SNe. If accretion drives shear in the WD, it may be an important source of viscous heating and material mixing. \\citet{yl04} considered the accretion and evolution of a WD up until the point of ignition, including the effects of rotation. In their study angular momentum was transported primarily via the Kelvin-Helmholtz instability, which lead to significant shear throughout the WD. \\citet{sn04} focused on whether rotational effects can prevent accretion induced collapse (AIC) at high accretion rates ($\\dot{M}=3\\times10^{-6}-10^{-5}\\ M_\\odot\\ {\\rm yr^{-1}}$). They found that an AIC was not averted, but perhaps more interestingly, they also found the WD should be rotating nearly uniformly when the baroclinic instability is included, a process that was neglected by \\citet{yl04}.  There is further opportunity for shear to develop during the carbon simmering phase. At carbon ignition, a convective region grows at the WD center. This envelops $\\sim1\\ M_\\odot$ over a timescale of  $\\sim1000\\ {\\rm yrs}$, until a burning wave commences. This phase has received increasing attention in recent years due to the realization that it sets the initial temperature and density for the explosive burning \\citep{les06}, as well as the distribution of ignition points \\citep{woo04,ww04,kuh06}. Nuclear reactions during this time may change the neutron excess of the WD core, which is crucial for determining the ratio of radioactive to non-radioactive nickel produced in the Type Ia explosion \\citep{pb07b,cha07}.  In the present study we consider the opportunity for developing shear during both of these stages. In \\S \\ref{sec:accretion} we revisit estimates for the shear of accreting WDs and confirm that the baroclinic instability limits the shear before the Kelvin-Helmholtz instability initiates. In addition we explore whether magnetohydrodynamic effects could reduce the shear even further. We conclude that the WD is nearly uniformly rotating at the onset of carbon ignition, and that viscous heating is negligible. In \\S \\ref{sec:convection} we explore the properties of convection present prior to explosive burning. The growth of this convective region allows further opportunity for shearing. The three-dimensional nature of this problem makes it hard to definitively determine what occurs during this stage. We illustrate some general features expected for the interaction of spin and convection by appealing to observations and numerical experiments. In \\S \\ref{sec:theend} we conclude with a summary of our results and a discussion of future research. ",
        "conclusions": "\\label{sec:theend}  In this study we have investigated the internal rotation profile for SN Ia progenitors, focusing on the stages of accretion and simmering. Since these progenitors are formed exclusively in binary systems, they contain a enormous amount of angular momentum. This is important for the WD structure and the subsequent flame propagation during the explosive burning of a SN Ia.  During the accretion phase we compared three different mechanisms for angular momentum transport. Either the baroclinic instability or the Tayler-Spruit dynamo limit the shear to values so low that the Kelvin-Helmholtz instability cannot occur. We argue that the WD will have nearly uniform rotation, consistent with the results of \\citet{sn04}.  This conclusion has important implications for extremely energetic SNe Ia such as SN2003fg \\citep[SNLS-03D3bb,][]{how06}, which has been argued to be from a super-Chandrasekhar WD \\citep{jef06}. One well-known way to get a mass this high is via strong differential rotation \\citep[][and references therein]{st83}. If uniform rotation is required then the WD mass can only be increased by a few percent over the normal Chandrasekhar limit. We should also mention that just because the WD mass is higher, the chemical yields of the deflagration or detonation of a strongly differentially rotating WD will not necessarily match the high nickel mass needed for SN2003fg \\citep{ste92,pfa06}. For this reason, other interpretations have been suggested for this SN Ia \\citep{hil07}.  The simmering phase is another opportunity for differential rotation to develop within the WD. This occurs due to heating and expansion as well as angular momentum mixing by convection. The full impact of the convection is hard to determine with certainty. Given the characteristic convective velocities, we expect $Ro\\lesssim1$, thus the influence of rotation cannot be discounted. We consider two rotation laws for the convective zone, a uniform rotation case and another case where the spin increases outward. In either case the most dramatic shearing appears at the convective/non-convective boundary. This shearing may have an impact on the propagation of flames or bubbles, and in some extreme circumstances may hold enough free energy to alter the thermal profile of the non-convective envelope.  Given the impact rotation could potentially have for SNe Ia, the time appears ripe for including such effects in simmering and flame propagation calculations. Simulations such as those performed by \\citet{bro04} are not yet able to reach the characteristic fluid parameters expected during the simmering stage \\citep[such as the Reynolds and Rayleigh numbers; see the summary in][]{woo04,ww04}. Nevertheless, $Ro\\approx0.1$ is well within reach, as has been shown by the simulations of \\citet{kuh06}. The biggest difficulty for the future will be following the secular evolution of the WD as the convective zone grows. Since the carbon burning is so temperature sensitive, the majority of the evolution takes place at late times (see Fig. \\ref{fig:shear_solid} and \\ref{fig:shear_incr}). This will help limit the duration that must be simulated."
    },
    "0801/0801.2104_arXiv.txt": {
        "abstract": "We present the first X-ray monitoring observations of the X-ray bright FR~I radio galaxy NGC~6251 observed with \\rxte\\ for 1 year. The primary goal of this study is to shed light on the origin of the X-rays, by investigating the spectral variability with model-independent methods coupled with time-resolved and flux-selected spectroscopy. The main results can be summarized as follows: 1) Throughout the monitoring campaign, NGC~6251 was in relatively high-flux state with an average 2--10 keV absorbed flux of the order of 4.5$\\times 10^{-12} ~{\\rm erg~cm^{-2}~s^{-1}}$ and a corresponding intrinsic luminosity of 6$\\times 10^{42}~{\\rm erg~s^{-1}}$. 2) The flux persistently changed with fluctuations of the order of $\\sim$2 on time scales of 20-30 days. 3) When the hardness ratio is plotted against the average count rate, there is evidence for a spectral hardening as the source brightens; this finding is confirmed  by a flux-selected spectral analysis. 4) The fractional variability appears to be more pronounced in the hard energy band (5--12 keV) than in the soft one (2.5--5 keV). 5) 2-month averaged and flux-limited energy spectra are adequately fitted by a power law. A Fe K$\\alpha$ line is never statistically required, although the presence of a strong iron line cannot be ruled out, due to the high upper limits on the line equivalent width. The inconsistency of the spectral variability behavior of \\ngc\\ with the typical trend observed in Seyfert galaxies and the similarity with blazars lead support to a jet-dominated scenario during the RXTE monitoring campaign. However, a possible contribution from a disk-corona system cannot be ruled out. ",
        "introduction": "Non-blazar radio-loud active galactic nuclei (AGNs) --objects with jets forming large viewing angles to the line of sight-- are traditionally divided into two classes: Fanaroff-Riley II (FR~II), with 178 MHz powers $> 2 \\times 10^{25}$ W/Hz and edge-darkened radio morphologies, and FR~I galaxies with lower powers and more compact morphologies (Fanaroff \\& Riley 1974). For the same host galaxy optical magnitude, FRIs produce about one order of magnitude less optical line emission than FRIIs (Baum et al. 1995) and have fainter or negligible UV continuum fluxes (Zirbel \\& Baum 1995, Ho 1999).  While earlier models for the origin of the FRI/II dichotomy focused mainly on accounting for their large-scale radio morphologies, more recently new ideas have emerged concerning the nature of the central engine in the two types of radio galaxies in an attempt to explain the nuclear properties. One school of thought is that the nuclear X-ray properties of FR~I and FR~II are related to a different accretion rate onto the central supermassive black hole, with FRII being dominated by relatively large values of $L/L_{Edd}$, while FRI would be accreting at sub-Eddington rates, $L/L_{Edd} \\ll 10^{-3}$ (e.g., Reynolds et al. 1996; Ghisellini \\& Celotti 2001).  A fundamental step to gain insight into the nature of the central engine in radio-loud AGNs is to understand whether the X-ray radiation is produced by disk/corona systems as in Seyfert galaxies or by jets as in blazars. While the time-averaged spectroscopy (due to spectral degeneracy) and the pure temporal analysis (due to the fact that both radio-quiet and radio-loud show strong variability) cannot firmly discriminate between the two competing scenarios, their combination, i.e., time-resolved spectral analysis and energy-selected temporal analysis, offers in principle a better way to distinguish between accretion-dominated and jet-dominated systems. This conclusion is supported by the strikingly different spectral variability behavior shown by Seyfert-like objects (e.g., Papadakis et al. 2002; Markowitz \\& Edelson 2001) and by blazars (e.g., Zhang et al. 1999; Fossati et al. 2000; Gliozzi et al. 2006).  Here, we concentrate on the X-ray nuclear properties of the nearby radio galaxy NGC~6251 ($z$=0.024), which is a giant elliptical galaxy hosting a supermassive black hole with mass $M_{\\rm BH}\\sim 4-8 \\times10^8 M_{\\odot}$ (Ferrarese \\& Ford 1999). Based on its radio power at 178 MHz, NGC~6251 is classified as an FR~I (e.g., Owen \\& Laing 1989), whereas in the optical, it is classified as a type-2 AGNs (e.g., Shuder \\& Osterbrock 1981).  Despite the intensive study of this source at all wavelengths, the nature of the accretion process in NGC~6251 is still a matter of debate. Based on the radio-to-X-ray spectral energy distribution, Ho (1999) suggested that an Advection-Dominated Accretion Flow (ADAF) is present in the nucleus of NGC~6251. On the other hand, Ferrarese \\& Ford (1999) and Melia et al. (2002) favored a standard accretion disk. Finally, Mukherjee et al. (2002) and  Chiaberge et al. (2003) advocate a jet origin for the broad-band emission, based on the possible association of \\ngc\\ with the {\\it EGRET} source 3EG~J1621+8203 and on its spectral energy distribution, respectively.  In the X-ray band, NGC~6251 has been previously observed with various satellites. For example, \\rosat\\ showed the presence of an unresolved nuclear source embedded in a diffuse thermal emission with temperature $kT \\sim 0.9$ keV, associated to the galaxy's halo (Birkinshaw \\& Worrall 1993). In addition, the existence of a correlation between soft X-rays and radio fluxes prompted Hardcastle \\& Worrall (2000) to hypothesize a jet origin for the soft X-rays. On the other hand, the prominent \\feka\\ line (EW$\\simeq$ 400--500 eV) detected by \\asca\\ in 1994 (Turner et al. 1997; Sambruna et al. 1999) suggested a standard Seyfert-like scenario with accretion-dominated X-rays. \\sax\\ observed \\ngc\\ in July 2001 during a  high-flux state (the 2--10 keV flux, $F_{\\rm X}=4.7\\times 10^{-12}$ \\flux, was  $\\sim$3 times larger than the value measured by \\asca\\ 7 years earlier). Based on the absence of a prominent \\feka\\ line (EW $<$ 100 eV) or other reprocessing features, Guainazzi et al. (2003) proposed a scenario with two main spectral components: a blazar-like spectrum dominating the high-flux state and a Seyfert-like spectrum emerging during the low-flux state. More recently, in March 2002, \\ngc\\ was observed with \\xmm\\ when the 2--10 keV flux was $\\sim$15\\% lower than the \\sax\\ value. Despite the relatively high flux, the \\xmm\\ spectral results seem to support the picture emerged from the \\asca\\ observation with the presence of a prominent (EW$\\sim$220 eV) and possibly broad ($\\sigma>0.3$ keV) \\feka\\ line (Gliozzi et al. 2004). However, the existence of a broad \\feka\\ line in the \\xmm\\ spectrum is still a matter of debate (see Evans et al. 2005 and Gonz\\'alez-Mart\\'in et al. 2006 for a discording and a supporting view, respectively).  The controversial results derived from previous X-ray studies highlight the impossibility of firmly determining the origin of the X-rays in NGC~6251 based solely on time-averaged spectral results. Here, we present the results from a systematic study of the long-term X-ray flux and spectral variability of \\ngc\\ using one year long proprietary Rossi X-ray Timing Explorer (\\rxte) observations in the 2-12 keV range. We use model-independent methods and time-resolved spectroscopy to study the X-ray temporal and spectral properties of this source. The main purpose of this analysis is to shed light on the origin of the X-rays and in particular on the role played by a jet in the X-rays. Once (if) the jet contribution is properly assessed, the physical parameters characterizing the accretion process onto the supermassive black hole can be better constrained, and hence it is possible to discriminate between competing theoretical models for the accretion process.  The outline of the paper is as follows. In $\\S~2$ we describe the observations and data reduction. The main characteristics of the X-ray light curve are described in $\\S~3$. In $\\S~4$ we study the X-ray spectral variability of \\ngc\\ with two model-independent methods. In $\\S~5$ we describe the results of a time-resolved and flux-selected spectral analyses. In $\\S~6$ we discuss the results and their implications. In $\\S~7$ we summarize the main results and conclusions. ",
        "conclusions": "We have used data from a year-long \\rxte\\ monitoring campaign to study the spectral variability of \\ngc\\, following model-independent and spectral model fitting methods. The main results can be summarized as follows:  \\begin{itemize}  \\item Throughout the monitoring campaign and especially during the last 4 months, \\ngc\\ was in a relatively high-flux state, with values of the 2--10 keV absorbed flux comparable to that observed by \\sax\\ in 2001.  \\item The light curves show persistent variations by a factor of 2 on timescales weeks/months in the total (2.5--12 keV), soft (2.5--5 keV), and hard (5--12 keV) energy bands.  \\item The fractional variability, computed over the soft and hard energy bands, reveals an enhanced variability at harder energies ($F_{\\rm var} \\sim$20\\%) compared to the soft band ($F_{\\rm var}\\sim$10\\%).  \\item There is evidence of a positive trend in the $HR-ct$ plot (or, analogously, of a negative trend in the $\\Gamma$-flux plot); in other words the spectrum hardens as the flux increases.  \\item The 2-month averaged spectra are well fitted by a power-law model, with $\\langle\\Gamma\\rangle\\simeq2.5$ and no indication for an \\feka\\ line. Combining all the 2-month averaged spectra yielded EW$<$ 104 eV. Only for the lowest flux-selected spectrum, there is marginal evidence for a \\feka\\ line. However, the low S/N does not allows one to put reliable constraints on the putative Seyfert-like component. \\end{itemize}  The inconsistency of the spectral variability behavior of \\ngc\\ with the typical trend observed in Seyfert galaxies and the similarity with blazars lead support to a jet-dominated scenario. However, based on the \\rxte\\ observations, a substantial contribution from a disk-corona system cannot be ruled out."
    },
    "0801/0801.4758_arXiv.txt": {
        "abstract": "{\\it{HST}} ACS/HRC images in UV (F250W), V (F555W), and I (F814W) resolve three isolated OB associations that lie up to 30 kpc from the stellar disk of the S0 galaxy NGC 1533.  Previous narrow-band H$\\alpha$ imaging and optical spectroscopy showed these objects as unresolved intergalactic H~II regions having H$\\alpha$ luminosities consistent with single early-type O stars. These young stars lie in stripped H~I gas with column densities ranging from 1.5 - 2.5 $\\times10^{20}$ cm$^{-2}$ and velocity dispersions near 30 km s$^{-1}$. Using the HST broadband colors and magnitudes along with previously-determined H$\\alpha$ luminosities,  we place limits on the masses and ages of each association, considering the importance of stochastic effects for faint (M$_{V}$ $>-8$) stellar populations.  The upper limits to their stellar masses range from 600 M$_{\\odot}$ to 7000 M$_{\\odot}$, and ages range from 2 - 6 Myrs.   This analysis includes an updated calculation of the conversion factor between the ionizing luminosity and the total number of main sequence O stars contained within an H~II region. The photometric properties and sizes of the isolated associations and other objects in the HRC fields are consistent with those of Galactic stellar associations, open clusters and/or single O and B stars. We interpret the age-size sequence of associations and clustered field objects as an indication that these isolated associations are most likely rapidly dispersing. Furthermore, we consider the possibility that these isolated associations represent the first generation of stars in the H~I ring surrounding NGC 1533. This work suggests star formation in the unique environment of a galaxy's outermost gaseous regions proceeds similarly to that within the Galactic disk and that star formation in tidal debris may be responsible for building up a younger halo component. ",
        "introduction": "\\label{intro}  H~II regions in gaseous tidal arms, far beyond the rotating disks of galaxies, present an opportunity to study the star formation process in conditions vastly different from those in typical galactic spiral arms. They may even contribute to our understanding of satellite galaxy formation and evolution, halo stellar populations, and the enrichment of the intergalactic medium.  Such intergalactic H~II regions have been found in galaxy clusters, compact groups, and in the halos of galaxies (e.g. \\nocite{boquien07, walter,emma,oosterloo,cortese,sakai,gerhard,arnaboldi, mendes04, gallagher01} Boquien et al. 2007; Walter et al. 2006; Ryan-Weber et al. 2004; Oosterloo et al. 2004; Mendes de Oliveira et al. 2004; Cortese et al.  2003; Sakai et al. 2002; Gerhard et al. H~II 2002;  Arnaboldi et al. 2002; Gallagher et al. 2001). In each case, they represent young stellar associations forming ostensibly where no stars existed previously. Intergalactic H~II regions  are distinct from the H~II regions in the outer arms of spiral galaxies (e.g., \\nocite{ferguson98} Ferguson et al. 1998) in that they are further from the host galaxy, they do not show significant continuum emission, and thus have much higher H$\\alpha$ equivalent widths \\citep{emma,gerhard,mendes04}.  These regions may represent H$\\alpha$-emitting counterparts to the extended UV-disk population of star clusters revealed by GALEX, as they lie at similar distances from the host galaxies, and exhibit similar levels of star-formation \\citep{Thilker07}. Much speculation exists as to whether intergalactic H~II regions represent isolated stars, what role tidal interactions play in their formation, whether their stellar populations are typical of young star clusters in quiescent systems, and whether they contribute significantly to the enrichment of the IGM.  \\begin{figure*} \\begin{center} \\plotone {f1_astroph.eps} \\figcaption {The H$\\alpha$ SINGG image of NGC 1533 overlaid with Australia Telescope Compact Array (ATCA) {\\ion{H}{1}} contours and marked with the locations of the individual spectroscopically confirmed H~II regions studied here. The detection images  are shown to the right of the SINGG image. The ATCA {\\ion{H}{1}} contours are 1.0 (in bold; 3$\\sigma$), 1.5, 2.0, 2.5, and 3.0 $\\times 10^{20}$ cm$^{-2}$ and have a resolution of $\\sim 1$ arcminute. There is no detectable H~I emission on the optical core of NGC 1533.} \\end{center} \\end{figure*}   Five intergalactic H~II regions, three of which have been spectroscopically confirmed, lie in the vicinity of the galaxy NGC 1533. \\cite{emma} discovered these H~II regions in the continuum-subtracted narrow-band image of NGC 1533 taken by the Survey for Ionization in Neutral Gas Galaxies (SINGG, \\nocite{SINGG} Meurer et al. 2006). The peculiar ring-like distribution of {\\ion{H}{1}} around NGC 1533 as seen in Figure 1 (and figure 2 of Ryan-Weber et al. 2004 \\nocite{emma} with higher H~I cut-off levels) may be the result of an interaction with the nearby dwarf galaxy seen in the NW corner, IC2038 (with the SINGG identification J0409-56:S2). No obvious optical counterpart to the disturbed {\\ion{H}{1}} appears in either the DSS image or the SINGG image. The velocities of the three spectroscopically confirmed regions coincide well with the {\\ion{H}{1}} gas, which is bound to NGC 1533 \\citep{emma}. The projected radius from the center of NGC 1533 to the isolated H~II regions ranges from 19 kpc (region 5) to 31 kpc (regions 1 and 2), up to four times the R-band 25$^{th}$ magnitude isophotal radius of NGC 1533. Table 1 summarizes the properties of the three H~II regions presented in \\cite{emma} and imaged here. It gives the regions'  positions, projected radii from NGC 1533, velocities determined from their H$\\alpha$ emission lines, H$\\alpha$ luminosities, and  H~I columns at each region's position. Note that the values in our Table 1 differ from those presented in \\cite{emma} because of the recent revision in the measurement of the distance to NGC 1533 using images taken in parallel to those discussed here (see Barber DeGraaff et al. 2007). In this paper, we present {\\it{Hubble Space Telescope}} Advanced Camera for Surveys/ High Resolution Channel (ACS/HRC) images and photometry of the three spectroscopically-confirmed isolated H~II regions associated with NGC 1533. These observations enable us to resolve the star-forming regions, identify their ionizing sources, and carry out an in-depth study of their stellar populations.  \\begin{deluxetable*}{ccccccc} \\tabletypesize{\\small} \\tablewidth{0pc} \\tablecaption{Properties of the Confirmed Isolated H~II Regions in NGC 1533 \\tablenotemark{a}} \\tablehead{\\colhead{H~II}& \\colhead{RA}& \\colhead{Dec}& \\colhead{Separation}&\\colhead{Velocity}& \\colhead{Log L$_{H\\alpha}$ \\tablenotemark{b}}&\\colhead{N$_{H~I}$}\\\\ \\colhead{Region}&\\colhead{}&\\colhead{}& \\colhead{kpc}&\\colhead{km s$^{-1}$}& \\colhead{ergs $s^{-1}$}&\\colhead{$10^{20}$cm$^{-2}$} } \\startdata  1&          04 10 13.7&     -56 11 37.5 &    31&   $846\\pm50$  &$37.67\\pm0.26$& 2.4\\\\ 2&          04 10 14.5&     -56 11 35.9 &    31&   $831\\pm50$  &$37.50\\pm0.26$&2.4\\\\ 5&           04 10 15.7&     -56 06 16.2 &    19&   $901\\pm50$  &$37.35\\pm0.26$&1.5\\\\    \\enddata \\tablenotetext{a} {\\scriptsize Summary of properties originally presented in Ryan-Weber et al. (2004): Positions are given in J2000 coordinates, with RA in hours, minutes, and seconds, and Declination in degrees, arcminutes, and arcseconds. Separation is the projected separation in kiloparsecs to the optical center of NGC 1533. The H~I column densities are measured from ATCA maps, at the position of each H~II region. NGC 1533 has a heliocentric velocity measured by HIPASS to be 785 km s$^{-1}$.  The numbers for projected separation and H$\\alpha$ luminosity differ from those presented in \\cite{emma} because we use the most recent measurement of the distance to NGC 1533, 19.4 Mpc (Blakeslee, private communication).} \\tablenotetext{b}{\\scriptsize As measured in the SINGG H$\\alpha$ images} \\end{deluxetable*}    In combination with the HST data, we make use of an ATCA H~I synthesis map of NGC 1533 \\citep{emma03} to help constrain some of the details of the formation and evolution of the isolated H~II regions.  Intergalactic H~II regions are often found forming in relatively low column density H~I gas. There is no H~I detected to a limit of 2  $\\times$10$^{19}$ cm$^{-2}$ at the location of the H~II region near NGC 4388 (\\nocite{hioosterloo}Oosterloo \\& van Gorkom, 2005), and the H~I surface density in the vicinity of the three confirmed systems, regions numbered 1, 2, and 5 in Figure 1 is $ 1.5-2.5 \\times 10^{20}$ cm$^{-2}$ \\citep{emma}. Furthermore, the H~I peaks in the neutral gas ring surrounding NGC 1533 and in the stripped material around NGC 4388 do not exhibit any detectable level of star formation.   These observations offer further support for the lack of correlation between star formation rate densities and H~I surface densities on local scales found by \\cite{kennicutt07}, unlike what is seen on global galactic scales.  The overall kinematics of the gas in the region of the intergalactic H~II regions may provide clues as to the trigger of the star formation and the potential growth of the stellar associations into a tidal dwarf galaxy.   Independent of the gaseous properties of cluster environments, star-cluster ``infant mortality\" suggests that the isolated, young star clusters presented here may be short-lived \\citep{rz05,fcw,cfw}. Will supernova and mass loss in these clusters remove enough interstellar matter to leave the stars freely expanding, resulting in the clusters' rapid dissolution? Studies of the Antennae galaxies \\citep{fcw,cfw} and the Small Magellanic Cloud \\citep{rz05} purport that over half of clusters formed dissolve less than 10 Myr after their formation independent of their initial masses. However, a recent study by \\cite{gieles07} concludes that there is little or no ``infant mortality\" for ages beyond 10-20 Myrs in the Small Magellanic Cloud, and \\cite{degrijs07} find less than 30\\% of young ($<20$ Myrs) SMC clusters undergo ``infant mortality.\" Generally, studies of this kind require a statistical sample of clusters. Instead, resolving a few star clusters in an isolated environment in a period during which they may be actively dissolving offers a different approach to examining their fates.   The photometric properties of the underlying stars in these intergalactic H~II regions can help to answer some of the questions surrounding their formation and evolution. Previous studies have used stellar population synthesis models and evolutionary tracks to infer cluster ages and masses. For example, \\cite{whitmore99} identify 3 distinct populations of star clusters in and around NGC 4038/4039 (Antennae Galaxy) using their integrated photometric properties. \\cite{gerhard} determine the mass and age of an isolated H~II region in the Virgo Cluster from its measured H$\\alpha$ and $V$ band luminosities.  Thus far,  studies of these sorts of objects have focused on what the star clusters' integrated light reveals about their formation histories, and have been unable to resolve the clusters into individual stars or stellar components. With the discovery of more and more young, faint clusters in tidal tails of merging galaxies, intracluster space and the intergalactic medium, there arises a need for information about the nature of the forming stars. Yet, there is an unavoidable, inherent uncertainty in the properties of lower mass clusters derived from models that assume a fully populated IMF \\citep{cervino}. Resolving the star clusters presents an opportunity to examine their populations in greater detail.  We organize this paper as follows: Section 2 contains a description of our observations and measurements; Section 3 presents the results and model comparisons; Section 4.1 explores comparisons between these intergalactic H~II regions and other extragalactic and Galactic young stellar populations; in Sections 4.2 and 4.3, we discuss the gas properties in the vicinity of the intergalactic H~II regions and examine the likely fates of the stellar associations; and Section 5 briefly summarizes our results.  The Appendix contains the new calculation of the conversion factor between the ionizing luminosity and the total number of main sequence O stars contained within an H~II region following the methodology of Vacca (1994). ",
        "conclusions": "HST ACS/HRC images in UV, V and I bands resolve the three very young stellar associations powering the intergalactic H~II regions that lie up to 30 kpc from the stellar disk of the S0 galaxy NGC 1533. We have catalogued and performed circular aperture photometry on 48 individual sources in the HRC images, 27 of which appear to be associated with the ongoing star formation. From background galaxy statistics and broadband colors, we have determined that the other 21 sources are likely distant background galaxies. Furthermore, we have defined a range of sources in the images: the term ``associations\" refers to the total stellar populations powering the H~II regions; the term ``clumps\" refers to the individual resolved sources composing the associations; and the term ``field objects\" refers to the individual sources detected separately  from the associations. Using two different methods,  (Starburst99; \\nocite{starburst99}Leitherer et al. 1999, and an independent calculation derived from \\nocite{vacca94, martins} Vacca 1994 and Martins 2005, see appendix) in combination with our broadband colors and magnitudes, and H$\\alpha$ luminosities from SINGG \\citep{SINGG}, we have placed upper limits on the the total numbers of stars, masses and ages of each association. The masses of the isolated associations range from 600 M$_{\\odot}$ to 7000 M$_{\\odot}$, and the ages range from 2 - 6 Myrs. Given the faint luminosities of the stellar associations, we have also considered the importance of stochastic effects upon our calculations. Colors and absolute magnitudes of many of the clumps within the associations and field objects are consistent with single O and B stars dominating the output radiation.  There are three results from our analysis that deserve emphasis: \\begin{itemize} \\item{Despite their location in an environment that is vastly different from that of a Galactic spiral arm, the photometric properties and sizes of the associations, clumps, and blue field objects are consistent with those of Galactic stellar associations and clusters and/or stars. This result suggests something uniform about the star-formation process on scales of tens of parsecs, independent of environment.} \\item{The level of star formation in the SE H~I cloud of NGC 1533 is lower than that observed in other systems (e.g. NGC 5291; \\nocite{boquien07} Boquien et al. 2007).  Furthermore, low H~I column densities (10$^{20}$cm$^{-2}$) and high H~I velocity dispersions ($\\sim30$km s$^{-1}$) characterize its environment. These data hint that there may be some relation between SFR, H~I column density and velocity dispersion that holds outside of massive gas disks, on local scales. These data are also consistent with the recent finding that H~I surface density and star formation rate densities do not correlate on sub-kiloparsec scales \\citep{kennicutt07}.  The kinematics of the H~I in the vicinity of the H~II regions do not suggest the stellar associations will evolve into a larger tidal dwarf galaxy.} \\item{The stellar associations detected here are either in the process of rapid dissolution or represent the first generation of stars in the gaseous ring around NGC 1533.  If the star formation in the SE part of the H~I ring has been continuous over the last 20 Myrs to 1 Gyr, we are directly witnessing the rapid dissolution of young clusters and associations. Not only do the associations display a range of ages and sizes consistent with cluster dissolution, but the age constraints and locations of several clustered blue field objects indicate they are the last visible remnants of an already dispersed association. If instead, these isolated associations and blue field objects represent a first generation of stars in the stripped H~I of NGC 1533, we cannot determine the ultimate fate of the associations. This possibility of a first generation of stars is then uniquely suited for exploring star formation triggers and initial conditions without the influence of stellar feedback.   Given that their Galactic OB association analogs are generally thought to be gravitationally unbound and freely expanding \\citep{blaauw,hills80}, along with evidence for star cluster ``infant mortality'' in various clustered stellar systems (e.g. \\nocite{ladalada, fcw} Lada \\& Lada 2003; Fall et al. 2005), we find it most likely that the stellar populations of the intergalactic H~II regions will soon ($< 20$ Myrs or so) no longer be detectable nor recognizable as associations of stars; rather, they will be building a background population of young objects in the halo of NGC 1533. }  \\end{itemize}"
    },
    "0801/0801.4933_arXiv.txt": {
        "abstract": "We implement an efficient method to quantify time-dependent orbital complexity in gravitational $N$-body simulations. The technique, which we name DWaTIM, is based on a discrete wavelet transform of velocity orbital time series. The wavelet power-spectrum is used to measure trends in complexity continuously in time. We apply the method to the test cases $N=3$ Pythagorean- and a perturbed $N=5$ Caledonian configurations.  The method recovers the well-known time-dependent complexity of the dynamics in these small-$N$ problems. We then apply the technique to an equal-mass collisional $N=256$ body simulation ran through core-collapse.  We find that a majority of stars evolve on relatively complex orbits up to the time when the first hard binary forms, whereas after core-collapse, less complex orbits are found on the whole as a result of expanding mass shells. ",
        "introduction": "\\label{Sect:Introduction}   Computer simulations of the gravitational $N$-body problem aim to solve the set of $3N$ second order ordinary differential equations \\begin{equation} \\label{Eqn:grav} \\bmath{F}_{i}=m_{i}\\bmath{\\ddot{r}}_{i}=-Gm_{i}\\sum_{l=1,l\\neq i}^{l=N}\\frac{m_{l}(\\bmath{r}_{i}-\\bmath{r}_{l})}{|\\bmath{r}_{i}-\\bmath{r}_{l}|^{3}},\\, (i=1,...,N). \\end{equation} Here $G$ is the gravitational constant, $\\bmath{F}_{i}$ and $\\bmath{\\ddot{r}}_{i}$ denote the force and the acceleration exerted by the $N-1$ particles $l$ on particle $i$ of mass $m_{i}$ at position $\\bmath{r}_{i}$. A system with $N>2$ is in general chaotic, i.e., the solution to Eq.\\,(\\ref{Eqn:grav}) is known to be sensitive to the initial conditions. Miller (1964) was the first to show that individual orbits calculated from neighbouring initial configurations diverge on a short time-scale proportional to $4\\xi/N$, where $\\xi$ is the mean time between two subsequent close encounters of a particle. Even when computed with double precision arithmetic, the differences between the two integrations become comparable to the characteristic length- and velocity-scales of the cluster within only a few crossing times. Thus one may argue that the positions and velocities obtained from an $N$-body simulation are a fair rendition of the problem at hand in a statistical sense only (\\citealt{Aarseth75}, \\citealt{Quinlan92}).  \\citet{Goodman93} identify three mechanisms responsible for the exponential growth of the distance between nearby trajectories. Aside from numerical errors\\footnote{A round-off error at one time step can be considered as a change in the initial conditions for the next time step.}, they point out that the exponential instability of the solutions also arises through a fluctuating mean gravitational field, or through inherently chaotic orbits in a static field. However, for systems in dynamical equilibrium and with near-spherical symmetry (such as e.g. globular clusters), they conclude that the principal mechanism responsible for chaos is the cumulative effect of near neighbour interactions, i.e. two-body encounters. \\\\ \\indent The exponential divergence of the $N$-body problem of Eq.\\,(\\ref{Eqn:grav}) has been studied by several authors for systems with up to $N=10^{5}$ particles (\\citealt{Kandrup91}; \\citealt{Kandrup01}; \\citealt{Hemsendorf02}). The debate whether Miller's instability is formally caused by chaos is still on-going (see e.g., \\citealt{Kandrup03}; \\citealt{Helmi07}). The assumed chaoticity of such systems may be evaluated by computing the variational equations associated to Eq.\\,(\\ref{Eqn:grav}) \\citep{Miller71} and by retrieving indicators of chaos such as, for instance, Lyapunov characteristic numbers \\citep{Benettin76}. Other indicators of chaos commonly found in the literature are e.g., the relative Lyapunov indicator \\citep{Sandor00}, the mean exponential divergence of nearby orbits \\citep{Cincotta00} or the small-alignment index \\citep{Skokos01}. All these indicators accurately quantify the exponential divergence of nearby trajectories.\\\\ \\indent In this work, we propose to investigate time-dependent orbital complexity in an $N$-body simulation. We define orbital complexity to be a measure of the richness and non-triviality of the frequency spectrum of an orbit at a given time $t$. The possible connection between orbital chaos and orbital complexity has been pointed out by \\citet{Kandrup97}. In that work, the authors find a strong correlation between their measure of orbital complexity and short time Lyapunov exponents. The notion of complexity provides information about the orbital content of a gravitational system governed by Eq.\\,(\\ref{Eqn:grav}) that is complementary to the classical indicators of chaos mentioned above. For instance, the concept of complexity is exploited by \\citet{Sideris02} to study the continuum limit in the case of large-$N$ simulations. It is used to compare the orbital behavior between particles evolving in smooth potentials and bodies orbiting in the corresponding frozen $N$-body configurations. The notion of complexity here allows to contrast the discreteness effects of the $N$-body configuration to the orbital evolution obtained in the smooth case. However, whereas a global measure of orbital complexity has been implemented in several works (see e.g., Kandrup et al. 1997), a method to measure the impact of instantaneous changes in orbital complexity has not. A trajectory computed from Eq.\\,(\\ref{Eqn:grav}) can show multiple, qualitatively distinct regimes in time. For instance, an orbit may display arcs of relatively smooth motion such as e.g., an unperturbed parabolic or hyperbolic orbit. At other times, the body may be gravitationally bound in a binary system with a particle of approximately the same mass. Likewise, the body may be temporarily trapped in a complicated higher-order resonance, orbiting about a massive central body. All these states of motion can be identified in time by a suitable measure of complexity. The goal of this paper is precisely to discuss such time-resolved complexity by introducing a dedicated tool for complexity evaluation. The issue of formally relating complexity to chaos will not be addressed here.\\\\ \\indent Classical spectral methods and Fourier transform based techniques (see \\citealt{Laskar93}; \\citealt{Carpintero98}; \\citealt{Valluri98}) are of short execution time and able to provide an accurate frequency domain representation of a given orbit. A global measure of complexity can be retrieved by such an approach. However, the Fourier transform suffers from the problem of losing any time-dependent information on the motion of the particle. In this paper, we present an accurate, easy-to-implement and time-resolved complexity-detection tool for individual orbits in $N$-body simulations. A state of dynamical equilibrium is assumed throughout. We implement a specially-adapted method of time series analysis that is based on a discrete wavelet transform. The method post-processes the orbital data of a simulation and provides a time-resolved measure of complexity. The present work mainly focuses on a detailed description of the method and is structured as follows. In Section 2 we introduce the identification technique for time-dependent complexity. Section 3 presents applications to $N=2$, $3$ and $5$-body problems. The analysis of a low-resolution $N=256$ equal-mass Plummer model is presented in Section $4$ and a brief discussion of the results is given in Section $5$. ",
        "conclusions": "We presented a method to compute the time-dependent orbital complexity in $N$-body simulations. The gravitational $N$-body problem is described by the $3N$ second order ordinary differential equations of Eq.\\,(\\ref{Eqn:grav}). We extract a discrete wavelet transform information measure (DWaTIM) from the velocity time series of the individual particles of the $N$-body simulation. We apply the technique to several few-body problems and to a larger $N=256$ particle simulation. The method captures the time-dependent changes in the dynamics of three- and five-body systems and furthermore quantifies orbital complexity continuously in time. For example, we recovered and quantified the dynamically more complex phases of the well-studied Pythagorean problem (see \\S\\ref{Subsect:pyth}) as well as the complex dynamics of a perturbed Caledonian configuration (see \\S \\ref{Subsect:caledonian}).  We also applied the method to a set of $N=256$ equal-mass Plummer spheres (see \\S \\ref{Sect:plummer}). We found that, on a global scale, orbital complexity decreases during the evolution. The occurrence of core-collapse and the subsequent core expansion causes a considerable drop in overall DWaTIM complexity $\\Upsilon_{\\rm tot}$. Furthermore, we observed that the complexity of individual orbits with a DWaTIM $\\Upsilon_{i}\\geq 0.5\\, (i=1,...,N)$ at the instant of core-collapse tends to increase until the occurrence of core expansion.\\\\ \\indent We opted for a DWaTIM implementation that allowed 1) to identify qualitative changes in orbital dynamics in a quasi-instantaneous manner and 2) to provide a continuous quantification of time-dependent complexity on a well-defined gauge between $0$ and $1$. One should however bear in mind that an \\textit{absolute} measure of complexity of a dynamical system cannot be obtained by the method. Complexity ultimately depends on the range of scales over which the system is studied. The DWaTIM is a band-limited measure and therefore strongly depends on the resolution of the time series that is subjected to analysis. Experience drawn from several test cases shows that a repeat calculation with truncated frequency range is desirable to confirm the convergence of the diagnostics.\\\\ \\indent Border effects are an unwanted artifact of the wavelet transform. In order to reduce the bias so introduced we opted to implement a \\textit{discrete} wavelet transform and to select cubic spline functions as mother wavelets.  This DWaT implementation avoids redundant wavelet coefficients. This helps in obtaining a proper measure of complexity for the two limiting cases of a unique base frequency sinusoid (DWaTIM = 0) and of a white noise signal (DWaTIM =~1), and thus to recover a well-defined gauge of complexity. We obtain reliable diagnostics of complexity whenever the orbit is integrated over a minimum of $8$ complete oscillations (see \\S\\ref{Subsect:pyth}).  The method is computationally inexpensive. For example, the analysis of a time series of length $Q=2^{13}$ data points takes approximately $1$ second on a Pentium IV 2.4 GHz workstation with 1.2 GB RAM memory. It is then possible to obtain a complexity diagnostic of an $N=10^5$ body system in about $80$ processor-hours. A parallel implementation of the scheme is currently under development."
    },
    "0801/0801.1511_arXiv.txt": {
        "abstract": "This paper is devoted to exploring how we can discover and study {\\it nearby} ($< 1-2$ kpc) planetary and binary systems by observing their action as gravitational lenses. Lensing can extend the realm of nearby binaries and planets that can be systematically studied to include dark and dim binaries, and face-on systems. As with more traditional studies, which use light from the system, orbital parameters (including the total mass, mass ratio, and orbital separation) can be extracted from lensing data. Also in common with these traditional methods, individual systems can be targeted for study. We discuss the specific observing strategies needed in order to optimize the discovery and study of nearby planetary and binary systems by observing their actions as lenses. ",
        "introduction": "Microlensing can be a powerful tool to study binaries and planets \\citep{MaoPaczynski, GouldLoeb, DiStefanoMao, DiStefanoScalzo1, DiStefanoScalzo2}. Until recently it has been assumed that most lenses, whether single or multiple stars, are likely to be at least several kpc away, either in the Galactic Halo or in the distant system of monitored sources. It therefore seemed likely that follow-up observations would be of limited value for multiple systems discovered through their action as lenses. In this paper we turn things around and consider lensing by binaries and planetary systems that are relatively nearby, within $1-2$ kpc. We refer to these nearby lenses as {\\it mesolenses} because (1) a combination of spatial and time lensing signatures may be exhibited by the events, (2) they have a high probability, relative to more distant lenses, of producing detectable lensing events, and (3) targeted observations designed to study the lensing action of some selected individual masses are possible. (See Di\\thinspace Stefano 2005.)  The lensing science is well-understood. We therefore concentrate on contrasts between nearby and more distant multiple lenses, framing the issues in terms of possible observing programs. Section 2 focuses on photometric monitoring, including ongoing and planned programs such as Pan-STARRS \\citep{PanSTARRS} and LSST \\citep{LSST}. \\S 3 discusses ground-based astrometric monitoring, which will be carried out by both Pan-STARRS and LSST. In \\S 4 we consider the possibility of selecting individual stars to be targets of sensitive lensing studies. In \\S 5 we discuss what lensing can reveal about companion stars and planets, listing our expectations for the next decade in \\S 5.1, and, in \\S 5.2,  emphasizing the contrasts and interplay with other types of investigations, such as Doppler and transit studies. ",
        "conclusions": ""
    },
    "0801/0801.2809_arXiv.txt": {
        "abstract": "% Recently, a dark energy model characterized by the age of the universe, dubbed ``agegraphic dark energy'', was proposed by Cai. In this paper, a connection between the quintessence scalar-field and the agegraphic dark energy is established, and accordingly, the potential of the agegraphic quintessence field is constructed. ",
        "introduction": "\\label{sec1} Many cosmological experiments, such as observations of large scale structure (LSS)~\\cite{LSS}, searches for type Ia supernovae (SNIa)~\\cite{SN} and measurements of the cosmic microwave background (CMB) anisotropy~\\cite{CMB}, all suggest that our universe is experiencing an accelerating expansion at the present time. Following the standard Friedmann-Robertson-Walker (FRW) cosmology such an expansion implies the existence of a mysterious dominant component, dark energy, with large enough negative pressure. Although we can affirm that the ultimate fate of the universe is determined by the feature of dark energy, the nature of dark energy as well as its cosmological origin remain enigmatic (for reviews, see e.g. \\cite{Weinberg88cp,Carroll00fy,Peebles02gy}). It is fair to say that disclosing the nature of dark energy is one of the central problems in the research of both cosmology and theoretical physics at present. Researchers proposed many candidates to interpret or describe the properties of dark energy. The simplest candidate for dark energy is the famous cosmological constant which has the equation of state $w=-1$. However, the cosmological constant scenario, is plagued with the so-called ``fine-tuning problem'' and ``coincidence problem''~\\cite{lamada}. Theorists have made lots of efforts to try to resolve the cosmological constant problem, but all these efforts seem to be unsuccessful. Of course the theoretical consideration is still in process and has made some progresses.  An alternative proposal for dark energy is the dynamical dark energy scenario which is often realized by scalar field mechanism. It suggests that the energy form with negative pressure is provided by a scalar field evolving down a proper potential.  A famous example of scalar-field dark energy is the so-called ``quintessence'' \\cite{quintessence}. Provided that the evolution of the field is slow enough, the kinetic energy density is less than the potential energy density, giving rise to the negative pressure responsible to the cosmic acceleration. So far, besides quintessence, a host of scalar-field dark energy models have been studied such as $K$-essence \\cite{kessence}, hessence \\cite{hessence}, tachyon \\cite{tachyon}, phantom \\cite{phantom}, ghost condensate \\cite{ghost} and quintom \\cite{quintom}, and so forth. But we should note that the mainstream viewpoint regards the scalar field dark energy models as an effective description of an underlying theory of dark energy. In addition, other proposals on dark energy include interacting dark energy models \\cite{intde}, braneworld models \\cite{brane}, and Chaplygin gas models \\cite{cg}, etc..  However, we still can make some attempts to probe the nature of dark energy according to some principles of quantum gravity although a complete theory of quantum gravity is not available today. The holographic dark energy model \\cite{holo1,holo2,holo3,holo4} is just an appropriate and interesting example, which is constructed in the light of the holographic principle of quantum gravity theory. That is to say, the holographic dark energy model possesses some significant features of an underlying theory of dark energy. More recently, a new dark energy model, dubbed agegraphic dark energy, has been proposed~\\cite{rgcai}, which takes into account the uncertainty relation of quantum mechanics together with the gravitational effect in general relativity.  In the general relativity, one can measure the spacetime without any limit of accuracy. However, in the quantum mechanics, the well-known Heisenberg uncertainty relation puts a limit of accuracy in these measurements. Following the line of quantum fluctuations of spacetime, K\\'{a}rolyh\\'{a}zy and his collaborators~\\cite{r1} (see also~\\cite{r2}) made an interesting observation concerning the distance measurement for Minkowski spacetime through a light-clock {\\it Gedanken experiment}, namely, the distance $t$ in Minkowski spacetime cannot be known to a better accuracy than \\begin{equation} \\delta t=\\lambda t_p^{2/3}t^{1/3}~,\\label{eq1} \\end{equation} where $\\lambda$ is a dimensionless constant of order unity. We use the units $\\hbar=c=k_B=1$ throughout this paper. Thus, one can use the terms like length and time interchangeably, whereas $l_p=t_p=1/m_p$ with $l_p$, $t_p$ and $m_p$ being the reduced Planck length, time and mass respectively.  The K\\'{a}rolyh\\'{a}zy relation~(\\ref{eq1}) together with the time-energy uncertainty relation enables one to estimate a quantum energy density of the metric fluctuations of Minkowski spacetime~\\cite{r3,r2}. Following~\\cite{r3,r2}, with respect to the Eq.~(\\ref{eq1}) a length scale $t$ can be known with a maximum precision $\\delta t$ determining thereby a minimal detectable cell $\\delta t^3\\sim t_p^2 t$ over a spatial region $t^3$. Such a cell represents a minimal detectable unit of spacetime over a given length scale $t$. If the age of the Minkowski spacetime is $t$, then over a spatial region with linear size $t$ (determining the maximal observable patch) there exists a minimal cell $\\delta t^3$ the energy of which due to time-energy uncertainty relation can not be smaller than~\\cite{r3,r2} \\begin{equation} E_{\\delta t^3}\\sim t^{-1}~.\\label{eq2} \\end{equation} Therefore, the energy density of metric fluctuations of Minkowski spacetime is given by~\\cite{r3,r2} \\begin{equation} \\rho_q\\sim\\frac{E_{\\delta t^3}}{\\delta t^3}\\sim \\frac{1}{t_p^2 t^2}\\sim\\frac{m_p^2}{t^2}~.\\label{eq3} \\end{equation}  In~\\cite{r3} (see also~\\cite{rgcai}), it is noticed that the K\\'{a}rolyh\\'{a}zy relation~(\\ref{eq1}) naturally obeys the holographic black hole entropy bound. In fact, the holographic dark energy~\\cite{holo1} also stems from the the idea of holographic black hole entropy bound. It is worth noting that the form of energy density Eq.~(\\ref{eq3}) is similar to the one of holographic dark energy~\\cite{holo1}, i.e., $\\rho_\\Lambda\\sim l_p^{-2}l^{-2}$. The similarity between $\\rho_q$ and $\\rho_\\Lambda$ might reveal some universal features of quantum gravity, although they arise from different ways. But, the agegraphic dark energy model avoids the causality problem which exists in the holographic dark energy models. For extensive studies on agegraphic dark energy, see~\\cite{Wei:2007ut,Wei:2007zs,Wu:2007wu,Wei:2007ty,Zhang:2007ps,Wei:2007xu}.  By far, hundreds of dark energy models have been constructed. However, it is hard to say that we are close to the Grail of revealing the nature of dark energy (that almost means we can understand the quantum gravity well). Although going along the holographic principle of quantum gravity may provide a hopeful way towards the aim, it is hard to believe that the physical foundation of agegraphic dark energy is convincing enough. Actually, it is fair to say that almost all dynamical dark energy models are settled at the phenomenological level, neither holographic dark energy model nor agegraphic dark energy model is exception.  Though, under such circumstances, the models of holographic and agegraphic dark energy, to some extent, still have some advantage comparing to other dynamical dark energy models because at least they are built according to some fundamental principle---holographic principle---in quantum gravity. We thus may as well view that this class of models possesses some features of an underlying theory of dark energy. Now, we are interested in that if we assume the agegraphic dark energy scenario as the underlying theory of dark energy, how the low-energy effective scalar-field model can be used to describe it. In this direction, we can establish a correspondence between the agegraphic dark energy and quintessence scalar field, and describe agegraphic dark energy in this case effectively by making use of quintessence. We refer to this case as ``agegraphic quintessence''. The quintessence potential $V(\\phi)$ can be reconstructed from supernova observational data \\cite{Saini:1999ba,simplescalar}. In addition, from some specific parametrization forms of the equation of state $w(z)$, one can also reconstruct the quintessence potential $V(\\phi)$ \\cite{Guo:2005at}. The reconstruction method can also be generalized to scalar-tensor theories \\cite{scalartensor}, $f(R)$ gravity \\cite{frgrav}, a dark energy fluid with viscosity terms \\cite{viscosity}, and also the generalized ghost condensate model \\cite{Zhang:2006em}. For a reconstruction program for a very general scalar-field Lagrangian density see \\cite{general}. For the scalar-field effective description of the holographic dark energy see \\cite{holoscalar}. In this paper, we shall reconstruct the quintessence potential and the dynamics of the scalar field in the light of the agegraphic dark energy.   In the next section, we briefly review the original agegraphic dark energy model proposed in~\\cite{rgcai} and the new agegraphic dark energy model proposed in~\\cite{Wei:2007ty}. In Sec.~\\ref{sec3}, we establish the correspondence between the agegraphic dark energy and the quintessence. The quintessence potential and the dynamics of the scalar field are also constructed in the light of the agegraphic dark energy. Conclusion is given in Sec.~\\ref{sec5}. ",
        "conclusions": "\\label{sec5} In conclusion, we suggest in this paper a correspondence between the agegraphic dark energy scenario and the quintessence scalar-field model. We adopt the viewpoint that the scalar field models of dark energy are effective theories of an underlying theory of dark energy. A new dark energy model, named as ``agegraphic dark energy'', has been proposed by Cai \\cite{rgcai}, based on the K\\'{a}rolyh\\'{a}zy uncertainty relation, which arises from the quantum mechanics together with general relativity. Perhaps, one may argue that the agegraphic dark energy lacks convincing physical foundation and the two models with energy densities in Eqs. (5) and (12) considered in this paper are just two types of the phenomenologically parametrized equation of state of dark energy. However, anyway, one should admit that it is an interesting attempt to consider the nature of dark energy based on some combination of quantum mechanics and general relativity. %   If we regard the scalar-field model (such as quintessence) as an effective description of such a theory, we should be capable of using the scalar-field model to mimic the evolving behavior of the agegraphic dark energy and reconstructing this scalar-field model according to the evolutionary behavior of agegraphic dark energy. We show that the agegraphic dark energy with $n\\geq 1$ can be described totally by the quintessence in a certain way. A correspondence between the agegraphic dark energy and quintessence has been established, and the potential of the agegraphic quintessence and the dynamics of the field have been reconstructed."
    },
    "0801/0801.0451_arXiv.txt": {
        "abstract": "We describe a web-based cgi calculator for constructing synthetic color-magnitude diagrams for a simple stellar population (SSP) using the Yonsei-Yale (YY) isochrone data base. This calculator is designed to be used interactively.  It creates quick look CMD displays in (B-V) and (V-I) colors.  Stochastic effects on the CMDs are included.  Output in tabular form is also provided for special purpose displays, or for combining the CMDs of different stellar populations. This research tool has applications in studies of the stellar content of our Galaxy and external systems.  It  provides an easy way to interpret the CMDs in resolved stellar populations.  It offers the means to explore the dependence of the integrated properties of unresolved stellar systems on stellar parameters (ages, chemical composition, binarity) and on the characteristics of their parent population (IMF slope and mass range). ",
        "introduction": "This program creates a synthetic color-magnitude diagram (CMD) for a simple stellar population (SSP), i.e. for a stellar population of a given age and chemical composition, that obeys a specified initial mass function (IMF) based on the prescription by Park \\& Lee (1997). Stochastic effects on the CMD are included. As an additional feature, the presence of a population of unresolved binaries can be taken into account (Woo et al. 2003).  The code is designed for use with the Yonsei-Yale (YY) isochrones, and covers the pre-main sequence and hydrogen burning phases of evolution.  The YY isochrones and luminosity functions are described in detail in papers by Yi et al. (2001, 2003; Paper~1 and Paper~3) and Kim et al. (2002; Paper~2).  Evolutionary tracks are available in Demarque et al. (2004; Paper~4).  The YY isochrone tables cover metallicities $Z$ (heavy element content by mass) in the range 0.0 to 0.08, and values of $\\alpha$ enhancement corresponding to  $[\\alpha/Fe]$ = 0.0, 0.3 and 0.6. Helium diffusion and convective core overshoot have been taken into account in calculating the evolutionary tracks.  This first set of isochrones (Paper ~1) was for the scaled solar mixture. The completing sets (for twice and four-times $\\alpha$-enhanced) were released in Paper~2. Two significant features of these isochrones are that (1) the stellar models start their evolution from the pre-main sequence birthline instead of from the zero-age main sequence, and (2) the color transformation has been performed using both the tables of Lejeune et al. (1998), and the older, but now modified, Green et al. tables (1987).  \\begin{deluxetable}{lccc} \\tablewidth{0pt} \\tablecaption{The assumptions for $\\alpha$-enhancement \\label{tbl1}} \\tablehead{\\colhead{Element} &\\colhead{[$\\alpha$/Fe]=0.0}\\tablenotemark{a} & \\colhead{[$\\alpha$/Fe]=+0.3} &\\colhead{[$\\alpha$/Fe]=+0.6} } \\startdata C & 8.55 &  & \\\\ N & 7.97 &  & \\\\ O & 8.87 & 9.17 &9.47 \\\\ Ne & 8.08 & 8.38 &8.68 \\\\ Na & 6.33 & 6.63 &6.93 \\\\ Mg & 7.58 & 7.88 &8.18 \\\\ Al & 6.47 & 6.17 &5.87 \\\\ Si & 7.55 & 7.85 &8.15 \\\\ P & 5.45 & 5.75 &6.05 \\\\ S & 7.21 & 7.51 &7.81 \\\\ Cl & 5.50 & 5.80 &6.10 \\\\ Ar & 6.52 & 6.82 &7.12 \\\\ K & 5.12 &  & \\\\ Ca & 6.36 & 6.66 &6.96 \\\\ Ti & 5.02 & 5.32 &5.62 \\\\ Cr & 5.67 &  & \\\\ Mn & 5.39 & 5.24 &5.09 \\\\ Fe & 7.50 &  & \\\\ Ni & 6.25 &  & \\\\ \\enddata \\tablenotetext{a}{ The scaled-solar abundance ratios of metals are taken from Grevesse \\& Noels 1993.} \\tablecomments{The abundance of the elements in logarithmic scale, $\\log N_{el}/N_H +12$, where $N_{el}$ is the abundance by number. } \\end{deluxetable}   The web interface to the CMD calculator can be found at the two YY web sites: \\begin{center} \\underline{\\textrm{\\textcolor{blue}{http://www.astro.yale.edu/demarque/yyiso.html}}}\\\\or\\\\ \\underline{\\textrm{\\textcolor{blue}{http://csaweb.yonsei.ac.kr/~kim/yyiso.html}}}\\\\ \\end{center} In addition, these two web sites provide links to the four YY  papers listed above, and codes to interpolate between isochrones in age and chemical composition within the available parameter range. ",
        "conclusions": "This note describes a web-based CMD calculator designed for the study of stellar populations in the Galaxy and distant stellar systems.  This web based calculator, which is based on the YY isochrones, is a useful tool for studies of resolved star clusters, and for explorations of the sensitivity of the integrated light of stellar systems to stellar and population parameters.  Future versions of this calculator will be based on a more comprehensive YY database covering an extended range of helium abundances, the inclusion of helium burning phases of evolution, often very significant in the CMDs of stellar systems, and the availability of a simple spectral energy distribution (SED) library.  This research was supported in part by NASA grant HST-GO-10505.03-A."
    },
    "0801/0801.2348_arXiv.txt": {
        "abstract": "We construct a scalar field based cosmological model, possessing a cosmological singularity characterized by a finite value of the cosmological radius and an infinite scalar curvature. Using the methods of the qualitative theory of differential equations, we give a complete description of the cosmological evolutions in the model under consideration. There are four classes of evolutions, two of which have finite lifetimes, while the other two undergo an infinite expansion. ",
        "introduction": "The discovery of the phenomenon of cosmic acceleration \\cite{cosmic} has stimulated a study of a huge variety of cosmological models, based on perfect fluids, scalar fields, tachyons etc \\cite{dark}. The cosmological models based on scalar  fields  have a long history, being used for exploration of possible inflationary scenarios \\cite{inflation} and for description of dark energy \\cite{quintessence}. This has attracted a special attention to the technique of reconstruction of scalar potentials reproducing a given cosmological evolution \\cite{reconstruct}.  On the other hand this development of model-designing art has revealed  cosmological evolutions possessing various types of singularities, sometimes very different from the traditional Big Bang and Big Crunch. The most popular between them is, perhaps, the Big Rip cosmological singularity \\cite{Star-rip,Rip} arising in superaccelerating models driven by some kind of phantom matter \\cite{phantom}. Other types of singularities are sudden singularities \\cite{sudden}, Big Brake \\cite{we-tach}, and so on \\cite{other}. Here we would like to study the singularities which are close to the known Big Bang, Big Crunch and Big Rip singularities, but arising at finite values of the cosmological factor (different from zero and infinity as well). Similar singularities were recently considered in \\cite{freeze}.  In the present paper we construct potentials which can drive the cosmological evolution towards (or from) such singularities. Combining qualitative and numerical methods we study the set of possible cosmological histories in the suggested models to show that the presence of such singularities in a cosmological model under consideration depends essentially on initial conditions and that the same model can accomodate qualitatively different cosmological scenarios.  The structure of the paper is the following: in Sec. 2 we construct some potentials corresponding to evolutions with ``soft'' singularities. In the third section we analyze their dynamics. The conclusion is devoted to an interpretation of the obtained results. ",
        "conclusions": "Let us sum up our results. Wishing to describe a cosmological evolution beginning from (or ending in) a singularity characterized by a nonvanishing initial (or final) cosmological radius and an infinite value of the scalar curvature due to the infinite value of the Hubble parameter, we have constructed a scalar field potential, providing such an evolution. Then, using the methods of qualitative analysis of the differential equations, we have shown that the proposed model accomodates four different classes of cosmological evolutions, depending on initial conditions. Numerical simulations have confirmed our predictions.  The main results of this work are: (i) the realization of a concrete cosmological model with a scalar field, where finite cosmological radius singularities are present and (ii) the complete description of all the possible evolutions of this model depending on the initial conditions. It is important to remark that, given the fixed scalar field potential, one has different types of evolutions encountering different kinds of singularities. The trajectories, belonging to the regions $I$ and $II$ have an infinite time of expansion, while the trajectories belonging to the regions $III$ and $IV$ begin and end their histories in the singularities and have a finite lifetime.  Let us note, that there are some studies \\cite{sudden,other,freeze} devoted to the general analysis of various kinds of new cosmological singularities. In this work we were not looking for an exhaustive classification of different possible cosmological models possessing   singularities, but rather we wanted to study in a complete way a particular cosmological model, having some interesting properties."
    },
    "0801/0801.2381_arXiv.txt": {
        "abstract": "We present a complete catalogue of 27 \\OVI\\ absorbers at low redshift ($0.12 < z < 0.5$) from a blind survey of 16 QSO echelle spectra in the HST/STIS data archive.  These absorbers are identified based only upon matching line profiles and the expected doublet ratio between the $\\lambda\\lambda\\,1031, 1037$ transitions. Subsequent searches are carried out to identify their associated transitions.  Here we present all relevant absorption properties. By considering absorption components of different species which are well-aligned in velocity-space, we derive gas temperatures and non-thermal broadening values, $\\bnt$. We show that in all 16 cases considered the observed line width is dominated by non-thermal motion and that gas temperatures are well below those expected for \\Oion\\ in collisional ionization equilibrium.  This result reaffirms previous findings from studies of individual lines of sight, but are at odds with expectations for a WHIM origin. At least half of the absorbers can be explained by a simple photoionization model. In addition, in some absorbers we find evidence for large variation in gas density/metallicity across components in individual absorbers. Comparisons of multiple associated metal species further show that under the assumption of the gas being photoionized by the metagalacitic background radiation field, the absorbing clouds have gas densities $<\\hden> < -2.9$ and sizes $L > 1\\kpc$. Finally, we compare our absorber selection with the results of other independent studies. ",
        "introduction": "\\label{sec: introduction} One of the current key questions in observational cosmology is the location of the missing baryons. The {\\it total} baryon content of the universe is well constrained, with various measurements in relatively good agreement \\citep[e.g.][]{burles-etal-01-BBN-omega, spergel-etal-03-WMAP-yr1, omeara-etal-06-deuterium}.  In the high-redshift universe, the \\Lya\\ forest dominates the baryon census \\citep{rauch-etal-97-Lya-Omegab}, but in the present day universe, only 1/3 of the baryons have been identified in known components \\citep{fukugita-peebles-04-mean-densities}.  Cosmological simulations indicate that up to 50\\% of the baryons exist in a warm-hot intergalactic medium (WHIM), in which gas is shock-heated to $\\sim 10^5 - 10^7\\K$ by the accretion onto large-scale structures, and remains hot owing to the low gas density and inefficient cooling \\citep{cen-ostriker-99-missing-baryons, dave-etal-01-WHIM-baryons, cen-ostriker-06-II-winds}. This gas is a result of accretion onto large-scale filamentary structures, where cooling is inefficient due to the low densities, and (in some models) the operation of large scale winds that shock-heat outflowing gas to $\\sim10^6\\K$ \\citep[e.g.][]{cen-ostriker-06-II-winds}.  It is therefore important to identify this gas, and constrain its contribution to the baryon fraction.  At the temperature range in question ($\\logT = 5.0 - 7.0$), the best observational window for this hot gas is X-ray absorption lines \\citep[e.g.\\ \\OVII\\,K$\\alpha$, \\OVIII\\,K$\\alpha$, \\NeIX\\,K$\\alpha$; ][]{gnat-sternberg-07-non-equilibrium-cooling}, but the resolutions of current X-ray spectrographs (e.g.\\ Chandra and XMM/Newton) are an order of magnitude too coarse to be useful, and the sensitivities are more than an order of magnitude too low \\citep[e.g.][]{fang-etal-06-galactic-OVII}. The best tool currently available, therefore, is UV absorption spectroscopy. The \\OVI\\ doublet ($\\lambda\\lambda\\,1031.9261, 1037.617\\A$) offers the best hot gas tracer for a number of reasons: the transition has a large oscillator strength; Oxygen is a relatively common metal; the abundance of the \\Oion\\ ion peaks in collisional ionization equilibrium (CIE) at $\\logT \\sim 5.3$\\footnote{It also has an ionization fraction greater than $10^{-2}$ up to $\\logT \\sim 5.7$; see e.g.\\ \\citet{gnat-sternberg-07-non-equilibrium-cooling}.}; the transitions occurs longward of the Lyman limit; and absorption can be detected down to limiting column densities $\\logN \\approx 13.5 (\\Wr = 30\\mA)$ with current instruments such as STIS.  If the absorption lines of an element are resolved and unsaturated, we can measure directly the column density ($N$) and Doppler parameter ($b$) of the absorbing gas by voigt profile fitting\\footnote{In cases where multiple transitions of the same species are available, a curve-of-growth analysis can also be used.}.  The Doppler parameter is an oft-used measure of the temperature of the gas, through the well-known relation $b^2 = \\bnt^2 + 2kT/m$, where \\bnt\\ accounts for non-thermal broadening of the line due to e.g.\\ turbulence. The column-density, meanwhile, may be used in conjunction with the path-length to derive the contribution to the cosmological mass density of the O$^{5+}$ ions, \\OmegaOVI\\ \\citep[see][hereafter \\pI]{thom-chen-08-I-OVI-statistics}.  In the optical band, ground-based high resolution (echelle) spectra have made \\OVI\\ detection possible at $2.3 \\simlt z \\simlt 3$ \\citep[e.g.][]{norris-etal-83-OVI, carswell-etal-02-OVI, simcoe-etal-02-OVI-search}. The lower limit is set by the plummeting transmissivity of the atmosphere to UV photons below $\\sim 3500\\A$, while the upper limit is set by confusion with the \\Lya\\ forest. In the high-redshift regime, \\Oion\\ is predominantly photoionized by the UV background radiation field, much the same as the \\Lya\\ clouds \\citep{carswell-etal-02-OVI, simcoe-etal-04}. To push to lower redshifts, space-based UV spectrographs are required. Using near-UV spectra from the Faint Object Spectrograph on-board HST, \\citet{burles-tytler-96-OVI-Omega_b} identified 12 \\OVI\\ doublets with $W_{r}(1031) > 0.21\\AA$ in the range $0.5 \\simlt z \\simlt 2$ and derived a cosmological mass density $\\OmegaOVI\\,h \\geq 7\\times10^{-8}$, providing the first constraint on the mass density of this highly ionized gas.  In the low-redshift universe ($z \\simlt 0.5$), the \\OVI\\ doublet transition remains in the far-UV ($\\lambda_{obs} \\simlt 1600\\A$).  The terminated Far Ultraviolet Spectroscopic Explorer \\citep[{\\it FUSE}; ][]{moos-etal-00-FUSE} could probe \\OVI\\ absorption out to $z \\sim 0.15$ with good signal-to-noise, but at a resolution of only $\\sim 20\\kms$ \\citep[e.g.][]{danforth-etal-05-OVI_baryon_census}. This compares poorly with the thermal line width of \\OVI\\ ($\\sim 3\\kms$ at $\\logT = 4.0$ and $\\sim 10\\kms$ at $\\logT = 5.0$).  The (currently suspended) Space Telescope Imaging Spectrograph \\citep[{\\it STIS}; ][]{woodgate-etal-98-STIS} on-board the Hubble Space Telescope ({\\it HST}) offers a resolution of $\\sim 6-7\\kms$ and its data are useful for studies of \\OVI\\ absorption at slightly higher redshifts than FUSE ($0.12 < z < 0.5$).  Early work on \\OVI\\ absorption in the low-$z$ universe typically concentrated on single lines of sight, due to the paucity of data. Single-sightline or single-absorber analyses have been conducted on the various lines of sight, beginning with the QSO H\\,1821$+$643 \\citep{tripp-etal-98-Lya-galaxies, tripp-etal-00-H1821+643, tripp-etal-01-H1821+643_z0.1212, oegerle-etal-00-H1821+643-FUSE}. As more data become available, further sightline were analysed: PG0953$+$415 \\citep{savage-etal-02-PG0953+415}, PG1259$+$593 \\citep{richter-etal-04-PG1259+593}, PG1116$+$215 821\\citep{sembach-etal-04-PG1116+215}, PKS0405$-$123 \\citep{prochaska-etal-04-PKS0405-12}, HE0226$-$4110 \\citep{lehner-etal-06-HE0226-4110} and PKS\\,1302$-$102 \\citep{cooksey-etal-08-PKS1302-102}. Now, with a database of UV data available\\footnote{More data are likely soon to be available with the upcoming HST servicing mission.}, statistical approaches have become possible \\citep{danforth-etal-05-OVI_baryon_census, danforth-etal-06-OVI_survey, thom-chen-08-I-OVI-statistics, tripp-etal-08-OVI, danforth-shull-08-OVI}.   This is the second in a series of papers reporting the results of our search for \\OVI\\ absorption systems in the STIS E140M archive. Unlike other searches, we employ a {\\it blind} search for \\OVI\\ doublets, which is independent of {\\it a priori} knowledge of the presence of other transitions such as \\Lya. In \\pI\\ we reported results on the statistics of the \\OVI\\ absorbers. The major results were: a) a measurements of the number of absorbers per unit redshift, $d\\,{\\cal N}(W\\ge 30\\,\\mA)/dz = 10.4 \\pm 2.2$; b) a measurement of the cosmological mass density of the \\Oion\\ gas, $\\OmegaOVI\\,h = (1.7 \\pm 0.3) \\times10^{-7}$; c) $<5$\\% of \\OVI\\ absorbers originate in underdense regions that do not show a significant trace of \\HI; d) \\HI\\ column densities of \\OVI\\ absorbers span more than 5 orders of magnitude, and a moderate correlation exists between \\NHI\\ and \\NOVI; and e) the number density of \\OVI\\ absorbers along a given line of sight appears to be inversely correlated with the number density of \\HI\\ absorbers. In this paper we present our catalogue of \\OVI\\ absorbers upon which the results of \\pI are based. We also address the physical conditions of the \\Oion\\ bearing gas. The nature of the IGM, the \\OVI\\ absorbers in particular, and their relation to the WHIM is an area that has seen much recent progress. At least two other groups have contemporaneously reported results of similar analyses. We refer the interested reader to the works of \\citet{tripp-etal-08-OVI} and \\citet{danforth-shull-08-OVI}. Specifically, see \\citet{tripp-etal-08-OVI} Sec~\\ref{sec: not_recovered} for comments on the differences between both works.  We recall the description of our search and selection technique from \\pI\\ in Sec~\\ref{sec: data}, presenting the full table of absorbers. We discuss individual lines-of-sight, and present the measured quantities for each system, in Sec~\\ref{sec: los}. In Sec~\\ref{sec: not_recovered} we discuss those systems reported in \\citet{tripp-etal-08-OVI} that are not accepted by our selection criteria. The physical properties of the ionized gas selected via \\OVI\\ absorption are discussed in Sec~\\ref{sec: properties}. Sec~\\ref{sec: discussion} contains a summary and concluding remarks. ",
        "conclusions": "\\label{sec: discussion}  We have presented a catalogue of \\OVI\\ absorbers selected from high-resolution HST/STIS echelle data. Our selection technique followed a blind-search for the \\OVI\\,1031, 1037 doublet feature. Relying solely on the presence of the \\OVI\\ doublet, and the ratio of doublet line strength, we detect 27 \\OVI\\ absorption systems, independent of other transitions.  The statistics of these absorbers were presented in \\pI.  We note that 16 systems reported in a contemporaneous work by \\citet{tripp-etal-08-OVI} do not satisfy our selection criteria and therefore are not included in our sample of 27 absorbers.  In cases where only one transition is found or the doublet ratio appears to be inconsistent with model expectations, these authors include the presence of other transitions, such as \\Lya\\ absorption, for justifying the identifications of \\OVI.  In our absorbers, it is common to find multiple absorption components in any given system, corresponding to physically distinct gas structures. By matching these absorption components from different species, we can identify different transitions likely due to the same gas. Under this assumption, we can then derive both the temperature of the gas, and the non-thermal contribution to the Doppler parameter, \\bnt. This analysis demonstrates that, for well-matched \\OVI/\\HI\\ components, gas temperatures are in the range $\\logT = 3.7 - 5.0$, well below the temperature range at which \\OVI\\ is expected to be found in collisional ionization equilibrium ($\\logT = 5.5$). We thus advise caution in identifying the \\OVI\\ absorbers with the hot WHIM gas that simulations predict will contain a significant fraction of the baryons. Our finding based on all available absorbers in the current HST/STIS data archive reaffirms previous findings from studies of individual lines of sight. Future generations of X-ray spectrographs are likely to be necessary to solve this question conclusively.  If galactic winds are the dominant source of \\OVI\\ formation, shock heating the gas to the WHIM regime, and pushing it out to $\\sim 1\\Mpc$ from galaxies, we may expect to see this in a variety of ways. A class of \\HI\\ free \\OVI\\ absorbers may arise, as both the hydrogen and oxygen are ionized in strong winds. We see no evidence of a large number of \\HI-free \\OVI\\ absorbers, with only two systems present in our survey. Reasoning that the \\HI\\ absorption will best trace the local over-density, we compared the velocity difference between our \\OVI\\ components and the nearest \\HI\\ component. A wind origin may leave an imprint on the kinematics of the absorbers, manifest as a systematic offset between the \\HI\\ and \\OVI\\ positions. We see no such effect (see also comparisons in T08), but this test is likely a weak one. Clearly systematic galaxy redshift surveys around the QSO lines of sight are required to properly address this.  The location of the missing baryons in the low-redshift is clearly an important issue in modern cosmology and astronomy. The WHIM is a leading candidate reservoir for this mass with simulations and suggesting that it contains as much as 50\\% of all $z = 0$ baryons. While some have claimed that the observations support this view, we argue that caution is required. High-resolution spectra that cover a broad wavelength range are crucial for constraining the ionization, metallicity, and the temperature of the gas involved, and care must be taken to deduce non-thermal broadening mechanisms. If the baryons are to be unambiguously located in the WHIM, further observations are crucial."
    },
    "0801/0801.0384_arXiv.txt": {
        "abstract": "We present the results of the first complete survey of the Large and Small Magellanic Clouds for 6668-MHz methanol and 6035-MHz excited-state hydroxyl masers. In addition to the survey, higher-sensitivity targeted searches towards known star-formation regions were conducted. The observations yielded the discovery of a fourth 6668-MHz methanol maser in the Large Magellanic Cloud (LMC), found towards the star-forming region N160a, and a second 6035-MHz excited-state hydroxyl maser, found towards N157a. We have also re-observed the three previously known 6668-MHz methanol masers and the single 6035-MHz hydroxyl maser. We failed to detect emission from either transition in the Small Magellanic Cloud. All observations were initially made using the Methanol Multibeam (MMB) survey receiver on the 64-m Parkes telescope as part of the MMB project and accurate positions have been measured with the Australia Telescope Compact Array (ATCA). We compare the maser populations in the Magellanic Clouds with those of our Galaxy and discuss their implications for the relative rates of massive star-formation, heavy metal abundance, and the abundance of complex molecules. The LMC maser populations are demonstrated to be smaller than their Milky Way counterparts. Methanol masers are under-abundant by a factor of $\\sim$45, whilst hydroxyl and water masers are a factor of $\\sim$10 less abundant than our Galaxy. ",
        "introduction": "Maser emission from star-forming regions in the Large Magellanic Cloud (LMC) has been detected from a number of molecular transitions. Initially, ground-state hydroxyl (OH) and water (H$_{2}$O) maser transitions were detected (Caswell \\& Haynes 1981; Haynes \\& Caswell 1981; Scalise \\& Braz 1982; Whiteoak et al. 1983; Whiteoak \\& Gardner 1986) and more recently methanol (CH$_{3}$OH) maser emission at 6668 MHz (Sinclair et al. 1992; Ellingsen et al. 1994; Beasley et al. 1996). In all cases the detection rates and strength of maser emission are lower than that seen in our own Galaxy.  Before the present survey, only three 6668-MHz methanol masers were known to exist outside the Milky Way, and all lay within the LMC. The first to be detected was towards the nebula N105a/MC23 (Sinclair et al. 1992).  A second was detected towards N11/MC18 as part of a targeted search of 48 H{\\sc ii} regions in the Magellanic Clouds (Ellingsen et al. 1994). A later search targeted 55 IRAS colour-selected sources and 12 known regions of bright H$\\alpha$ emission and found a third 6668-MHz methanol maser toward IRAS 05011-6815 (Beasley et al. 1996). A recent, very sensitive, search of M33 failed to detect any further extragalactic methanol masers (Goldsmith, Pandian \\& Deshpande 2007). The excited-state OH maser transition at 6035 MHz is often seen to be coincident with 6668-MHz methanol in our own Galaxy (e.g. Caswell 1997), but only one extragalactic example has been reported; it is in the LMC, detected towards the nebula N160a/MC76 (Caswell 1995) and is not associated with any of the known methanol masers.  Searches of the SMC for 6668-MHz methanol masers by Ellingsen et al. (1994) towards 13 IRAS positions,  and by Beasley et al.  (1996) towards 12 H{\\sc ii} regions, failed to detect any emission. The only known masers in the SMC are 22-GHz H$_{2}$O masers at S7 and S9 (Scalise \\& Braz 1982).  The SMC and LMC are both significantly less massive than the Milky Way.  Stanimirovi\\'c, Staveley-Smith \\& Jones (2004) found a dynamical H{\\sc i} mass for the SMC of  2.4 $\\times$ 10$^{9}$ M$_{\\odot}$, while the mass of the LMC is between 6 $\\times$ 10$^{9}$ and 1.5 $\\times$ 10$^{10}$ M$_{\\odot}$ (Meatheringham et al. 1988; Schommer et al. 1992), compared to the Milky Way which has a mass of the order of 5 $\\times$ 10$^{11}$ M$_{\\odot}$.  In addition the LMC has a star formation rate one tenth that of the Milky Way (Israel 1980) and a lower metallicity, at 0.3 to 0.5 times the solar metallicity (Westerlund 1997). This will directly affect the ISM chemistry, and produce a lower dust-to-gas ratio, resulting in a higher ambient UV field (Gordon et al. 2003).  Beasley et al. (1996) conclude that lower oxygen and carbon abundances in the LMC would lead to an under abundance of methanol by a factor of 6$-$12. Additionally, based on the work of Cragg et al. (1992), they believed, due to the significant role played by dust mid-IR emission in pumping this maser, that the lower dust abundance would contribute to fewer occurrences of methanol masers. The pumping model is now supported by even more accurate modelling (e.g. Cragg, Sobolev \\& Godfrey 2002; Cragg, Sobolev \\& Godfrey 2005) and is widely accepted.  To provide a more complete and uniform view of the maser population in the LMC and SMC, we report on the first complete systematic survey of the LMC and SMC made using the Parkes 64-m radio telescope. We present the results of the survey, supplemented with some deeper targeted Parkes observations, and some follow-up measurements with the Australia Telescope Compact Array (ATCA). ",
        "conclusions": "A complete survey of the Magellanic Clouds has been conducted as part of the Methanol Multibeam project and supplemented by a targeted search of known star-formation regions in the Large Magellanic Cloud. This has detected two new extragalactic masers: a fourth 6668-MHz methanol maser towards the star-forming region N160a in the LMC; and a second extragalactic 6035-MHz OH maser towards N157a. We have also re-observed the three previously known 6668-MHz methanol masers and the single 6035-MHz OH maser in the LMC. All these masers have had their positions determined with the ATCA and exhibit stability in spectral structure and intensity, even over the $\\sim$10 year separation between the current and previous observations. The current lower limits on the LMC maser populations are demonstrated to be up to a factor of $\\sim$45 smaller than the Milky Way, with the methanol showing the largest discrepancy. Even after correcting for lower star formation rates (section 4.1), methanol masers in the LMC are still under-abundant by a factor of $\\sim$4-5. This remaining disparity, in agreement with the speculation of Beasley et al. (1996), may be due to the lower oxygen and carbon abundances in the LMC.  These observations, only made feasible by the speed provided by the 7-beam MMB receiver, have produced the most accurate determination yet of the luminosity distribution of extragalactic masers. When combined with the complete MMB Galactic catalogue, they will allow for an improved comparison of Galactic star-formation regions with those of lower metallicity found in the LMC. This will also allow us to better determine whether the LMC 6668-MHz methanol and 6035-MHz OH maser populations simply scale directly with the star-formation rate or if other factors are involved."
    },
    "0801/0801.4751.txt": {
        "abstract": "The intriguing observations of {Swift}/BAT X-ray flash XRF 060218 and the BATSE-BeppoSAX gamma-ray burst GRB 980425, both with much lower luminosity and redshift compared to other observed bursts, naturally lead to the question of how these low-luminosity (LL) bursts are related to high-luminosity (HL) bursts. Incorporating the constraints from both the flux-limited samples observed with CGRO/BATSE and Swift/BAT and the redshift-known GRB sample, we investigate the luminosity function for both LL- and HL-GRBs through simulations.  Our multiple criteria, including  the $\\log N - \\log P$ distributions from the flux-limited GRB sample, the redshift and luminosity distributions  of the redshift-known sample, and the detection ratio of HL- and LL- GRBs with Swift/BAT, provide a set of stringent constraints to the luminosity function. Assuming that the GRB rate follows the star formation rate, our simulations show that a simple power law or a broken power law model of luminosity function fail to reproduce the observations, and a new component is required. This component can be modeled with a broken power, which is characterized by a sharp increase of the burst number at around $L < 10^{47}$ erg s$^{-1}$. The lack of detection of moderate-luminosity GRBs at redshift $\\sim 0.3$ indicates that this feature is not due to observational biases. The inferred local rate, $\\rho_0$, of LL-GRBs from our model is $\\sim 200$ Gpc$^{-3}$ yr$^{-1}$ at $\\sim 10^{47}$ erg s$^{-1}$, much larger than that of HL-GRBs. These results imply that LL-GRBs could be a separate GRB population from HL-GRBs. The recent discovery of a local X-ray transient 080109/SN 2008D would strengthen our conclusion, if the observed non-thermal emission has a similar origin as the prompt emission of most GRBs and XRFs. ",
        "introduction": "Long duration Gamma-ray bursts (GRBs) are believed to be tied to the death of massive stars (Colgate 1974; Woosley 1993; Stanek et al. 2003; Hjorth et al. 2003; Campana et al. 2006).  Of the roughly 6000 bursts observed since the late 1960s (see e.g. http://heasarc.gsfc.nasa.gov/grbcat/), almost 100 have redshift measurements. Observations show that long GRBs are scattered over a large redshift and luminosity\\footnote{Throughout the text the burst luminosity refers to the isotropic equivalent value, which does not include the possible correction of the unknown beaming factor of the GRBs.} range, from $z=0.0085\\sim$ 6.7\\footnote{This record redshift measurement is from GRB 080913.  It remains unclear what category this particular burst belongs to, Type I or Type II, and is not included in the statistical analysis of this work.} and $L=10^{46}\\sim 10^{54}~{\\rm erg~s}^{-1}$. Most of these bursts have high luminosity (HL, $L>$ several $10^{48}$ erg s$^{-1}$) with the exception of two peculiar bursts, GRBs 980425 and 060218, which have extremely low redshift and luminosity measurements, $(z,L) = (0.0085, 4.7\\times 10^{46}~{\\rm erg~s}^{-1}$) and (0.033, 6.03$\\times 10^{46}~{\\rm erg~s}^{-1}$) respectively (Tinney et al. 1998; Mirabal et al 2006). It remains unclear whether the LL-GRBs are due to unusual progenitor properties or a unique population with an intrinsic difference in the central engine, i.e., black hole versus magnetar (Mazzali et al. 2006; Soderberg et al. 2006; Toma et al. 2007).  One empirical way to look into this problem is to see whether observational data collected so far are still consistent with LL-GRBs as a natural extention of HL-GRBs to low luminosities in a continuous luminosity function (LF), or if LL-GRBs form a distinct new LF component. We have suggested that the latter possibility (two-component LF model) is necessary based on the redshift-known sample of GRBs after the discovery of GRB 060218 (Liang et al. 2007). Such a possibility has also been considered as a hypothesis in the BeppoSAX era (Coward 2005). On the other hand, although Guetta \\& Della Valle (2007) agreed that the two-component model is possible, they also argued that the $z$-known GRB sample may be also consistent with a single component model with a steeper slope in the luminosity function so that more LL-GRBs are accounted for.  Comparing observational data with simulations is a useful way to address the relationship between LL-GRBs and HL-GRBs. It is an important task to constrain the luminosity function, $\\Phi(L)$, and local rate of GRBs, $\\rho_0$, in a manner that can self-consistently reproduce various observations. In particular, the current LL-GRB population studies have been focused on the $z$-known sample only. They were not confronted with the existing $\\log N - \\log P$ distribution of the BATSE GRB sample. The $\\log N-\\log P$ distribution (or $V/V_{\\rm max}$ distribution) addresses the statistical properties of GRBs regardless of their redshift, and carries essential information of the GRB LF. For HL-GRBs, this criterion has been utilized extensively (e.g. Schmidt 2001; Stern et al. 2002; Lloyd-Ronning et al. 2002; Norris 2002; Guetta et al. 2005) and was confronted with simulations (Lloyd-Ronning et al. 2004; Dai \\& Zhang 2005; Daigne et al. 2006). The conclusion has been that $\\Phi(L)$ of HL-GRBs is generally characterized by a one-component broken power law model with $\\rho_{0, HL}\\sim 1 ~{\\rm Gpc^{-3}~yr^{-1}}$ (e.g. Schmidt 2001; Guetta et al. 2004, 2005). On the other hand, the $\\rho_0$ of LL-GRBs inferred from the two detections (GRBs 980425 and 060218) in a decade suggests a much higher local rate than that of HL-GRBs, i.e., $\\rho_{0,LL}=100 \\sim 1000$ Gpc$^{-3}$ yr$^{-1}$ (Coward 2005; Cobb et al. 2006; Pian et al. 2006; Soderberg et al. 2006; Liang et al. 2007; Chapman et al. 2007). Guetta et al. (2004) propose an extension to lower luminosities which would increase ${\\rho_{\\rm 0, HL}}$ from roughly $1.1~{\\rm Gpc^{-3}~yr^{-1}}$ to $10~{\\rm Gpc^{-3}~yr^{-1}}$, at roughly $10^{48}~{\\rm erg~s^{-1}}$. This local rate, however, even extrapolated down to $10^{45}~{\\rm erg~s^{-1}}$ is not sufficiently large to produce the observed LL events. This was the main motivation of the two-component model in our previous analysis (Liang et al. 2007, see also Coward 2005; Le \\& Dermer 2007). The analysis with the $z$-known sample makes an arguable case (Liang et al. 2007), but is by no means conclusive (cf. Guetta \\& Della Valle 2007).  In this paper we extend our previous analysis to include a more complete set of observational constraints. In particular, we introduce the important BATSE and Swift $\\log N-\\log P$ distribution criteria along with the previously considered multiple criteria involving the $z$-known sample (1-D $z$-distribution, 1-D $L$-distribution, 2-D $z-L$ distribution, and the observed number ratio of HL- vs. LL-GRBs). Although the number of LL-GRBs remains the same, the $z$-known sample has grown since our last analysis in Liang et al. (2007),  a firmer conclusion drawn in this paper is largely due to the additional observational criterion included in this analysis. In order to confront multiple criteria with different LF models and a wide range of LF parameters, we utilize a series of Monte Carlo simulations (MCSs). Instrument observational selection effects are difficult to model, and we introduce some empirical formulae to roughly reflect gamma-ray detector trigger sensitivity and the selection effect of redshift measurement. Various models are presented in \\S 2. Our siumlation results are shown in \\S 3, and the conclusions and discussion are presented in \\S 4. The concordance cosmology with parameters $H_{\\rm 0} = 71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm{m}}=0.3$ and $\\Omega_\\Lambda = 0.7$ is assumed throughout.  Throughout the paper we do not touch on another distinctly different group of bursts, namely short-hard (Kouveliotou et al. 1993), or more general Type I (see e.g. Zhang et al. 2007 for a discussion of the multiple criteria needed to classify GRBs) bursts, which are found to be consistent with the compact-star-merger origin (Gehrels et al. 2005; Fox et al. 2005; Barthelmy et al. 2005; Berger et al. 2005c, see M\\'esz\\'aros 2006, Nakar 2007, Zhang 2007 for reviews). The analysis of these GRBs applying the same technique will be presented elsewhere. ",
        "conclusions": "By utilizing Monte Carlo simulations and various test criteria we are able to further constrain the form and parameters of the LF of long GRBs. After confronting model results with various observational criteria, including 1-D and 2-D $L-z$ distributions of the redshift-known GRBs and the $\\log N - \\log P$ distributions of both CGRO/BATSE and Swift/BAT GRBs, we conclude that various one-component LF models discussed by previous authors (e.g. Guetta et al. 2004, 2005; Guetta \\& Della Valle 2007) are insufficient to account for all the data.  As the luminosity function parameters are modified to accommodate for the observed LL-bursts, these models cause an overproduction of bursts at low and intermediate redshifts that is irreconcilable with the observed distributions.  Although a PL or BPL would seem a simple and straight-forward solution to the LF problem, our reasoning and analysis imply that a two-component LF model (Liang et al. 2007, Coward 2005) is necessary. The latter model implies an event rate of local LL-GRBs of $\\sim 100-400 \\rm~Gpc^{-3}~yr^{-1}$ at an $L_{B,LL}$ of $L\\sim 10^{47}~{\\rm erg~s^{-1}}$, which is much larger than that of HL-GRBs ($\\sim 1 \\rm~Gpc^{-3}~yr^{-1}$).  In addition, as mentioned in Liang et al. (2007), the functional form of the LL-bursts LF is quite uncertain, especially below the break luminosity.  Most constraints are drawn from the more numerous HL observations, although important information about local rate and break luminosities for both distributions can be drawn simply from the number of LL-events detected within the time of Swift's operation.  Other effects that add difficulty to the analysis, but are addressed as fully as possible, include the inhomogeneity of the burst sample, redshift detection selection effects, and the relatively small sample size.   A recent development that may affect the local rate determination is the serendipitous discovery of  a very low luminosity (peak luminosity $\\sim 6.1 \\times 10^{43}~{\\rm erg ~s^{-1}}$) X-ray transient, XRF 080109, by Swift XRT (Soderberg et al. 2008). This event is associated with SN 2008D. Although it has been suggested that the X-ray emission may be related to shock breakout (Soderberg et al. 2008), the possibility that the X-ray emission is the jet emission from a very low luminosity X-ray flash has been suggested and cannot be ruled out from the data (Xu et al. 2008; Li 2008). In particular, the non-thermal spectrum of XRF 080109 makes it different from the thermal X-ray emission component discovered in XRF 060218, another shock-breakout emission candidate claimed in the literature (Campana et al. 2006). We therefore regard the origin of XRF 080109 inconclusive. If it is indeed a very low luminosity LL-GRB, its very high event rate (Soderberg et al. 2008; Xu et al. 2008) is consistent with the conclusion of this paper that LL-GRBs form a distinct new component in LF, and the event rate increases to even higher values towards low luminosities. This model also predicts the existence of X-ray flashes in the luminosity range of $10^{44}-10^{46}~{\\rm erg~s^{-1}}$ that bridge XRF 080109 and XRF 060218.  The current LL-GRB sample is too small to address whether they follow the same empirical correlations as HL-GRBs, such as the lag-luminosity relation (Norris et al. 2005), variability - luminosity relation (Reichart et al. 2001),  spectral peak energy - isotropic energy relation (Amati et al. 2002), etc. Apparent discrepancy exists for some (e.g. GRB 980425 as an outlier of the Amati-relation), but consistency exists for others (e.g. GRB 060218 satisfies the Amati and lag - luminosity relations, Liang et al. 2006; Amati 2006). On the other hand, these correlations are closely related to the radiation physics, which is not directly related to the central engine and progenitor of the bursts. The GRB fireball picture is very generic. Bursts with different types of progenitors may share the same emission physics and hence, similar empirical correlations. More data are needed to draw firmer conclusions in this direction.  In the analysis above we show that this 2-component LF model can interpret observations from both {Swift} and CGRO/BATSE in various criteria. Although some accommodations have been made to find common ground within all tests employed, these criteria shed light on the luminosity function problem and imply that a two-component LF is necessary. Various effects that are not considered in this work may further affect the luminosity and redshift distribution of the observed bursts. Changes to the form of the star formation rate appear often in the literature and are essential to the basic assumptions of long GRBs as being associated with the death of massive stars.  Monte Carlo Simulations provide a useful tool for probing this effect, either as different functional forms of the SFR (Porciani and Madau 2001; Rowan-Robinson 1999; Hopkins and Beacom 2006) or deviations and evolutions with redshift (Kistler et al. 2007).   Other effects that might affect the distributions include an evolution of the luminosity function with redshift (Lloyd-Ronning et al. 2002) or a dependence on cosmic metallicity (Li, 2007). These processes might provide solutions to the deficit of simulated bursts at high-luminosity and high-redshift, since most of these effects produce a larger rate of bursts at high redshift. Consequently, this redshift dependence does not affect the nearby LL-population.  Understanding how and to what extent each of these processes affects the luminosity and redshift distributions is a necessary next step in the constraints of the luminosity function of GRBs, and we plan to explore them in full in a future work."
    },
    "0801/0801.0541_arXiv.txt": {
        "abstract": "We present Doppler imaging and Balmer line analysis of the weak-line T Tauri star TWA 6. Using this data we have made one of the first attempts to measure differential rotation in a T Tauri star, and the first detection of a slingshot prominence in such a star. We also show the most direct evidence to date of the existence of solar-type plages in a star other than the Sun.  Observations were made over six nights: 11--13th February 2006 and 18--20th February 2006, when spectra were taken with the UCL Echelle Spectrograph on the 3.9-m Anglo--Australian Telescope. Using least-squares deconvolution to improve the effective signal--to--noise ratio we produced two Doppler maps. These show similar features to maps of other rapidly rotating T Tauri stars, i.e. a polar spot with more spots extending out of it down to equator. Comparison of the two maps was carried out to measure the differential rotation. Cross-correlation and parameter fitting indicates that TWA 6 does not have detectable differential rotation.  The Balmer emission of the star was studied. The mean H$\\rm \\alpha$ profile has a narrow component consistent with rotational broadening and a broad component extending out to $\\pm$250\\kms. The variability in H$\\rm \\alpha$ suggests that the chromosphere has active regions that are cospatial with the spots in the photosphere, similar to the `plages' observed on the Sun. In addition the star has at least one slingshot prominence $3R_{*}$ above the surface -- the first such detection in a T Tauri star. ",
        "introduction": "\\subsection{T Tauri Stars}  T Tauri stars are pre-main sequence equivalents of Sun-like stars. They are young, have low masses (0.3 -- 1.0~\\mdot) and spectral types F--M~\\citep{hayashi66}. The T Tauri phase is considered to begin when the star becomes optically visible and continues until hydrogen burning begins and the star has reached the post-T Tauri phase at an age of the order of $10^{7}$ years (depending on spectral type).  Very young stars fall on Hayashi tracks on the Hertzsprung-Russell diagram, i.e. vertical tracks of constant temperature. As the star approaches the zero-age main sequence (ZAMS) it can either continue to follow the Hayashi track all the way onto the ZAMS (and remain fully convective), or develop a radiative core and turn onto the Henyey track and evolve onto the ZAMS with constant luminosity and increasing temperature. Which of these evolutionary paths the T Tauri star follows is thought to depend on the mass of the star but observations of TTS with ages close to this turn-off will allow us to characterise this stage of stellar evolution more fully.  T Tauri stars fall into two subcategories -- classical and weak-line (\\citealp{bertout89} and references therein). Observations of classical TTS (CTTS) show excess ultraviolet and infrared emission and strong emission line activity, characteristics which are due to the presence of a circumstellar disc. The star--disc interactions lead to strong, irregular variability that is predominately due to accretion of material from the disc onto the surface. The disc has a braking effect on the star, causing CTTS to rotate more slowly than their discless equivalents \\citep{bouvier93,cc93}.  Weak-line T Tauri stars (WTTS) have weaker emission-line activity and show less evidence of accretion. WTTS appear to be older than CTTS on average \\citep{armitage03},  suggesting that when TTS form they have circumstellar discs which dissipate as the star evolves, but disc lifetimes are variable, ranging from 0.1 -- 10~Myr \\citep{bouvier97}. WTTS are variable, but in contrast to CTTS the variation is usually periodic as it is due to cool spots that go in and out of view as the star rotates. These cool spots form when magnetic flux tubes emerge through the photosphere.  Doppler imaging \\citep{vogt83} maps spots on stellar photospheres. Doppler imaging of rapidly rotating, magnetically active stars, such as WTTS, often finds spots at high latitudes and even at the poles, see e.g. \\cite{v41096}, \\cite{strass02}, in contrast to the Sun where they are normally within $\\pm30^\\circ$ of the equator. One of the favoured explanations is that Coriolis forces in rapidly rotating stars cause the flux tubes to emerge at higher latitudes \\citep{schu92,caligari94}.  In the past Doppler imaging of cool stars has been limited by its challenging signal-to-noise requirements. The technique of least-squares deconvolution \\citep[LSD,][]{jf97} has made it possible to produce Doppler images of stars for which the signal-to-noise ratio (SNR) would be too low if individual spectral lines were used. In LSD the underlying rotationally broadened line profile for each spectrum is found by deconvolving a line list from the spectrum. This makes it possible for late--K T Tauri stars to be imaged, where previously the SNR requirements would have been prohibitive.  Sunspots are accompanied by prominences, i.e. loops of material that follow the magnetic field lines emerging through the surface of the star. High--lying prominences, known as `slingshot prominences', typically appear as fast--moving absorption transients and have been observed in a number of young stars such as Speedy Mic \\citep{duns06} and AB Dor \\citep{cc89,abdor99}, but not in any TTS to date. These prominences are viewed as a tracer for large-scale magnetic fields. If a T Tauri star was shown to be able to support such high-lying prominences for several rotational periods it would suggest the presence of a large-scale field that does not decay away quickly outside the stellar radius.  \\subsection{Differential Rotation in Stars with Deep Convection Zones}\\label{intdiff}  Fully convective stars are believed to have a distributed magnetic dynamo ($\\alpha^{2}$), arising from turbulent motions or Coriolis forces, rather than a solar-type dynamo ($\\alpha\\Omega$) that is due to both convection and differential rotation, as seen in stars with radiative cores.  Differential rotation, when rotational velocity is a function of latitude, is expected to be less apparent with increasing convection zone depth \\citep{kitch99}, so stars with larger convection zones are expected to rotate more like a solid body. Differential rotation in the Sun gives rise to the $\\Omega$ effect that produces a large-scale toroidal field from the small-scale poloidal field. Most models for an $\\alpha^{2}$ dynamo rule out the possibility of there being differential rotation in fully convective stars \\citep[e.g.][]{chab06}. However \\cite{dobler06} suggest fully convective stars have significant differential rotation, although this would not be Sun-like, and takes the form of a meridional circulation.  Observations to date have only attempted to detect the latitudinal differential rotation seen in the Sun. A recent measurement for the fully-convective M4 dwarf V374 Peg \\citep{morin07} found a weak differential rotation (around 0.1 of the solar value). Differential rotation has been observed to fall with decreasing photospheric temperature. For example, figure 2 in \\cite{barnes05a} plots differential rotation against temperature. The trend seems to imply that differential rotation approaches zero below 4000K but it is uncertain where it becomes undetectable as there has been no measurement to date of stars with temperatures between 3800--4500K. In addition there have not been any reliable measurements of differential rotation in T Tauri stars. Such measurements are needed to confirm that TTS follow these trends too. \\subsection{TW Hydrae Association}  Nearby young co-moving groups are a rich source of T Tauri stars for observation. Our chosen targets lie in the TW Hydrae association (TWA), a group of pre-main sequence stars lying at a distance of about $51\\pm 5$pc and with an age of 8 Myr \\citep{zuck04}. Of the 24 likely members of the TWA, 16 have had their periods measured in \\cite{lawson05}. The age of this association is an important one as it is corresponds to the mean age of circumstellar discs and hence the time when T Tauri stars can begin to spin up, so the members of this association are interesting objects for study. The star that gives the association its name (TW Hya) is a classical T Tauri star, spectral type K7. Its disc has been directly imaged in \\cite{twdisk} and its magnetic field was detected in \\cite{yang07}. Although it is rotating fairly rapidly (P = 2.8d) it is almost pole--on and is therefore not a suitable candidate for Doppler imaging.  This paper concerns another member of the TW Hydrae association -- TWA 6. It has a spectral type of K7 \\citep{webb99}, corresponding to a temperature of about 4000K for a T Tauri star \\citep{cohen79}. Models by \\cite{siess00} indicate that it is fully convective or just beginning to develop a radiative core. Photometry has revealed a large lightcurve amplitude, with a variation of 0.49 magnitudes in the V band \\citep{lawson05}. Its rapid rotation, \\vsini~ = 55~\\kms~ \\citep{webb99} and period of 0.54 days \\citep{lawson05} make it an excellent target for Doppler imaging. In this paper we present Doppler images of TWA 6, and a discussion of whether it has differential rotation. We also analyse the Balmer line emission in order to study the behaviour of the chromosphere. ",
        "conclusions": "\\label{con}  Spectra of TWA 6 were taken at AAT/UCLES during two epochs separated by approximately nine rotational periods. Using these spectra and Doppler imaging two maps of its spot distribution were produced. The spot distribution of TWA 6 is qualitatively similar to other T Tauri stars, e.g LO Peg, \\citep{barnes05b} and V410 Tau, \\citep{v410a}, with the predominate spot emergence at high latitudes and spots all the way to the equator.  In the Sun photospheric spots are often coincident with `plages', active regions in the chromosphere, usually observed using a H$\\rm \\alpha$ filter. The H$\\rm \\alpha$ emission of TWA 6 indicates that the active chromospheric regions are cospatial with the photospheric spots  - similar to the behaviour of the Sun. This supports previous findings that H$\\rm \\alpha$ emission in young stars increases when spot coverage is larger (\\cite{jf97b}, \\cite{fern04}) but here we have, for the first time, a strong indication that the plages are overlying the spots.  TWA 6 appears to support slingshot prominences, the first such observation in a T Tauri star. At least one prominence is observed, lying at a velocity of 285\\kms~which corresponds to a radial distance of $4R_{*}$ from the axis. There is evidence that the prominence survives for three days after the initial detection - suggesting that the structure is stable over at least 5.5 rotational periods. The presence of a stable slingshot prominence at a height of $3R_{*}$ from the surface suggests that the star has a large-scale field. Further observations of TWA 6 are justified in order to fully characterise its prominence system, as well as of other stars of similar age and spectral type to determine whether the behaviour set out in this paper is typical.  The star has no differential rotation (upper limit $0.003~\\rm{rad~day^{-1}}$), in line with the previously observed trend that differential rotation decreases with decreasing effective temperature and increasing rotational velocity. Earlier differential rotation measurements have been for older objects, here we have extended the observations to younger stars. The null--measurement is consistent with expectations for the spectral type and angular speed of TWA 6 and indicates that T Tauri stars exhibit solid--body rotation."
    },
    "0801/0801.0893_arXiv.txt": {
        "abstract": "We present a deep [O~III] $\\lambda\\lambda$4959,5007 image of the northern filamentary jet in the Crab Nebula taken with the 8.2m Subaru telescope.  Using this image and an image taken with the KPNO 4m in 1988 \\citep{fes93}, we have computed proper motions for 35 locations in the jet. \\textbf{The results suggest that when compared to the main body of the remnant, the jet experienced less outward acceleration from the central pulsar's rapidly expanding synchrotron nebula.} The jet's apparent expansion rate yields an undecelerated explosion date for the Crab Nebula of 1055 $\\pm$24 C.E., a date much closer to the appearance of the historic 1054 C.E.\\ guest star than the 1120 -- 1140 C.E.\\ dates estimated in previous studies using filaments located within the remnant's main nebula. \\textbf{Our proper motion measurements suggest the jet likely formed during the 1054 supernova explosion and represents the remnant's highest velocity knots possibly associated with a suspected N-S bipolar outflow from the supernova explosion.} ",
        "introduction": "The first object in Messier's catalog of nebulous sources, the Crab Nebula (M1), is a young galactic supernova remnant (SNR) and one of the most well known astronomical objects. It possesses a variety of remarkable properties and structures including: a highly complex structure of optically bright filaments as seen in recent high-resolution Hubble Space Telescope images \\citep{hes96,loll04}, an unusually luminous central pulsar (see Harden \\& Seward 1984 and Rots et al. 2004 and references therein) with accompanying X-ray jets \\citep{bri85}, and a radio and optically bright synchrotron nebula which exhibits fine-scale structural changes on timescales of just days or weeks \\citep{sca69,sca70,hes95}.  The Crab Nebula is generally considered the archetype of a pulsar driven synchrotron wind nebula (PWN) and of the so-called  ``plerionic'' remnants in general (see review by \\citealt{dav85}). The spinning magnetic field of the pulsar and strong relativistic wind of particles off the neutron star have measurably accelerated the remnant's optical filaments such that their velocities, when projected back, give a date of $\\sim$1130 C.E. \\citep{tri70, nug98}.  This date is $\\sim$80 years more recent than the 1054 C.E. date, the year a ``guest star\" associated with the Crab supernova is reported to have been observed by the Chinese and other cultures in several ancient astronomical texts \\citep{cla77,ste02}. This apparent filament acceleration, confirmed by several independent proper motion studies (see Table 1), due to an outward pressure contained within the remnant's rapidly expanding PWN is described in detail by \\citet{TR70}.  \\setcounter{table}{0} \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{Results from Previous Proper Motion Studies} \\begin{tabular}{@{}lcccc@{}} \\hline Reference & Number of Knots & Image Epochs & Convergence Date & $\\sigma$ (yr)\\\\ \\hline Duncan 1939             & 20    & 1909, 1938                          & 1172        & ...    \\\\ Trimble 1968            & 132   & 1939, 1950, 1953, 1964, 1966, 1966 & 1140        & 15         \\\\ Nugent 1998             & 50    & 1939, 1960, 1976, 1992              & 1130        & 16         \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*}  \\setcounter{figure}{0} \\begin{figure*} \\centering \\includegraphics[width=.9\\textwidth]{f1_gimp.ps} \\caption{Subaru [O~III] image of the Crab Nebula using a log intensity stretch to show the position of the northern filamentary jet in relation to the remnant's brighter inner filament complex.} \\label{fig1} \\end{figure*}  A particularly intriguing feature of the Crab is a faint, northern filamentary `chimney' or `jet' of emission filaments, which we will henceforth refer to simply as the ``jet'', but which should not to be confused with the physically smaller X-ray synchrotron jets off the central pulsar.  First detected in the optical by \\citet{van70}, this optical filamentary jet consists of line-emitting filaments, especially prominent in [O~III] $\\lambda\\lambda$4959,5007 \\citep{che75}. Although primarily a line-emission feature, very weak coincident non-thermal continuum emission was subsequently discovered, first in the radio \\citep{vel84} and then in the optical \\citep{wol87}, suggesting the jet itself is also filled with relativistic particles off the pulsar.  One especially puzzling aspect of the jet is its morphology. The first deep [O~III] images of the jet showed it to be a complex network of filaments that appear to be cylindrically symmetric, but with a central axis unaligned with either the Crab's center of expansion or pulsar \\citep{gul82}. Follow-up imaging and kinematic studies \\citep{shu84,ver85,mar90} lead to several proposed explanations including a mass loss trail from a red giant progenitor \\citep{bla83}, interaction of ejecta with an interstellar cloud \\citep{mor85}, and transverse synchrotron driven shocks \\citep{san97}.  However, preliminary proper motion studies of the jet indicated ordinary radial motions away from the center of expansion \\citep{fes86, fes93}.  Due to the jet's greater distance from the pulsar and the remnant's synchrotron nebula, the jet may have experienced much less acceleration than most other regions of the remnant. In that case, the jet might lead to a dynamically greater remnant age leading to a date closer to the 1054 C.E. historic supernova sighting.  Here we present a deep and high resolution [O~III] image of the Crab jet taken with the 8.2 m Subaru telescope. This image shows a network of extremely fine and detailed filaments. We have compared this image to a 1988 epoch image obtained with the KPNO 4m telescope \\citep{fes93} to obtain improved proper motion measurements of its optical filaments and an estimate of the jet's expansion age. ",
        "conclusions": "\\subsection{Dynamical Age of the Crab Nebula}  Our computed convergence date for the jet is 1055 $\\pm 24$ C.E., in excellent agreement with historical records which indicate the Crab SN occurred in the year 1054.  We attribute the reason for the difference between our result and earlier age estimates to be chiefly the jet's northernmost knots experiencing little appreciable acceleration by the pulsar's rapidly expanding synchrotron nebula. All previous age estimates based upon proper motion studies used knots and filaments located within the main body of the Crab remnant. Consequently, interior knot motions reflected both an expansion due to the SN event itself and a post-SN acceleration resulting from the interaction of the pulsar wind nebula.  Although the 1054 explosion date for the Crab Nebula's SN has not been in serious doubt, it is nonetheless valuable to have a direct confirmation based on expansion dynamics.  Although the jet's knots and filaments appear not to have been significantly accelerated radially away from the remnant's central region, our measured jet knot proper motions may indicate a small, non-radial component.  As shown in Figure 5 (right hand panel), the projected motions of the knots intersect north of the position of the center of expansion. This effect may not represent the actual nature of the jet as there are several sources of error which could contribute to this motion. First, it is possible that there may be a slight rotational shift between the two World Coordinate System fits for the images.  Because the original motions of the jet ejecta were so close to directly northward, a slight shift east or west would cause a much more noticeable change in the E-W component of their proper motion than would a similarly sized shift in the N-S direction. Thus, the directions of the knots would be greatly affected by such a shift, while the velocities would show considerably less error.  Secondly, an error in centroiding the knots can account for a slight shift as well.  Our computed errors show a change in direction of approximately 3$^{\\circ}$ is possible which could account for the non-radial motions of the knots.  On the other hand, it is possible that the measured motions do reflect the jet's true expansion properties. The jet is essentially a hollow structure and as shown by \\citet{vel84} and \\citet{wol87}, the jet does contain some faint synchrotron emission, most visible at its southern base.  If this plasma induced an acceleration that was largely against the jet's filamentary walls, this might act to expand the jet in the east-west plane and explain why the northwestern jet knots end up east of the center of expansion while most knots located at the northeastern portion of the jet project back west of the center of expansion.  Since our derived expansion date for the jet is most sensitive to north-south motions, such an acceleration would not significantly effect our estimated dynamical age.  \\subsection{The Nature of the Jet}  Our proper motion measurements provide a valuable constraint on the origin of the Crab's northern jet of ejecta. A successful jet formation model should explain the jet's non-radial appearance, its peak expansion velocity, specific location along the remnant's outer periphery, and the lack of any obvious counter-jet to the south. Below, we briefly review some of the proposed theories for the jet's origin and  discuss them in relation to our dynamical measurement results.  \\subsubsection{Plasma Instabilities Models }  Several authors have suggested the jet is a post-SN feature resulting from an interaction between the ejecta and the pulsar wind nebula. Many variations of plasma instability models have been suggested including: 1) an aneurysm of ejecta pushed aside by the pulsar's plasma flow \\citep{byc75,ver85,mar90}, 2) Rayleigh-Taylor instabilities \\citep{che75,kun83} or differential acceleration \\citep{che75}, 3) pressure from swept-up magnetic field lines \\citep{shu84}, 4) ejecta entrainment in magnetic fields formed by highly-ordered outflows from the polar caps of the pulsar itself \\citep{ben84}, and 5) synchrotron driven shocks \\citep{san97}.  A particularly attractive feature of such models is their ability to explain both the non-radial appearance of the jet and the lack of a counter-jet to the south. However, our proper motion measurements indicate an essentially radial expanding jet. Moreover, some of these plasma instability models call for greater acceleration of the jet's filaments than those found in the main body of the nebula \\citep{che75,shu84}, or a younger age compared to the rest of the nebula \\citep{kun83}. \\textbf{ While a greater degree of acceleration or a younger age for the jet is not supported by our data which suggests a current jet age around 950 yr similar to the Crab's historic age, a more modest expansion breakout acceleration is not completely ruled out by our proper motion measurements. }  \\textbf{ For most plasma instability (`blowout') models, the exact location of a high-velocity ``jet'' along the remnant's outer boundary is undefined but with a greater likelihood along the pulsar's NW--SE wind injection axis. } However, the remnant's filamentary jet is some 45 degrees away from the pulsar's synchrotron jets, which lie along a NW--SE axis (PA = 300 deg; \\citealt{ng04}). Along this NW--SE axis, the PWN appears to have pushed aside the remnant's outer ejecta filaments creating a virtual hole in the filamentary ejecta along the NW limb \\citep{law95,cad04}. There the synchrotron emission can be seen to leak out of the remnant's thick thermal plasma shell in optical, X-rays, and radio images of the remnant. The lack of extened filamentary line emission ejecta in this region similar to the northern `jet' is suggestive that the northern jet was not the result of a breakout of the pulsar's synchrotron wind nebula.  \\subsubsection{An Asymmetrical Explosion?}  \\textbf{ Despite a lack of an obvious southern counter-jet, the remnant's northern jet might represent high-velocity ejecta resulting from a north-south bipolar expansion asymmetry.} In 1989, MacAlpine et al. showed that N-S oriented spectral drift scans of the Crab revealed a nearly N-S bipolar structure. The data showed that the pinched waist of the  hourglass-like expansion  coincided with a band of helium rich filaments (a ``high-helium band'') in which the computed helium mass fraction was at least 75\\% \\citep{uom87}. \\citet{mac89} suggested a helium rich circumstellar torus or disk could have resulted in constrained velocities in the horizontal region around the progenitor. Subsequent drift scans of the remnant by \\citet{fes97} and \\citet{Smith03} confirmed the presence of a strong bipolar expansion.  Additional support for a circumstellar torus was presented by \\citet{sch79} and \\citet{fes92} who examined the east and west indentations seen in the synchrotron emission (the synchrotron bays) and suggested these may have be caused by an E-W magnetic torus associated with the E-W band of He-rich filaments. Such a magnetic torus might constrain the east-west expansion of the PWN, leaving the bay indentations seen today.  In an asymmetrical explosion model for the jet, its location along the northern edge of the remnant is not accidental \\citep{fes93}.  The alignment of the jet with the northern extent of the north expansion bubble seen in spectral drift scans of the remnant plus its major axis aligned nearly perpendicular to the synchrotron bays suggest a causal link between the jet and a possible constrained E-W expansion of the SN ejecta. The lack of an equally obvious southern `counter-jet' might be explained by either a weaker ejection of material to the south and/or increased ejecta confinement, and indeed drift scans of the remnant do show higher velocities to the north compared to the south, relative to the pinched central zone \\citep{mac89,fes97,Smith03}.  The jet's filaments, being farther from the pulsar's energetic wind, would experience less acceleration compared to knots and filaments located in the main body of the nebula. Our measurements of other locations at the base of the jet and to the east and west appear consistent with this picture, showing \\textbf{larger} dynamical ages correlating with their increased distance from the pulsar (See Table 3).  Finally, we note that our computed expansion age of the jet, although yielding an expansion age consistent with the historic SN sighting in the year 1054, could also be viewed consistent with \\textbf{an initial} acceleration by the pulsar and \\textbf{subsequent} deceleration through an interaction with a suspected outer halo of unshocked ejecta \\citep{che77,mur81,lun86,hes96}.  The presence of appreciable outer ejecta, however, remains highly controversial with several radio, optical, and X-ray observations failing to find any supportive evidence for its existence \\citep{fra95,fes97,wal99,sew06}. \\textbf{In any case, we view it as unlikely that the degree of any PWN acceleration and possible outer ejecta induced deceleration would be so nearly offsetting as to lead to the currently derived jet expansion date of around C.E. 1054}."
    },
    "0801/0801.2258_arXiv.txt": {
        "abstract": "Gravitational lenses with anomalous flux ratios are often cited as possible evidence for dark matter satellites predicted by simulations of hierarchical merging in cold dark matter cosmogonies. We show that the fraction of quads with anomalous flux ratios depends primarily on the total mass and spatial extent of the satellites, and the characteristic lengthscale $\\Rfifty$ of their distribution.  If $\\Rfifty \\sim100\\kpc $, then for a moderately elliptical galaxy with a line-of-sight velocity dispersion of $\\sim250 \\kmpersec$, a mass of $\\sim3\\E{9} \\Msun$ in highly-concentrated (Plummer model) satellites is needed for 20\\% of quadruplets to show anomalous flux ratios, rising to $\\sim1.25 \\E{10} \\Msun$ for 50\\%. Several times these masses are required if the satellites have more extended Hernquist profiles. Compared to a typical elliptical, the flux ratios of quads formed by typical edge-on disc galaxies with maximum discs are significantly less susceptible to changes through substructure -- three times the mass in satellite galaxies is needed to affect 50\\% of the systems.  In many of the lens systems with anomalous flux ratios, there is evidence for visible satellites (e.g., B2045+265 or MG0414+0534). We show that if the anomaly is produced by substructure with properties similar to the simulations, then optically identified substructure should not be preponderant among lens systems with anomalies. There seem to be two possible resolutions of this difficulty. First, in some cases, visible substructure may be projected within or close to the Einstein radius and wrongly ascribed as the culprit, whereas dark matter substructure is causing the flux anomaly. Second, bright satellites, in which baryon cooling and condensation has taken place, may have higher central densities than dark satellites, rendering them more efficient at causing flux anomalies. ",
        "introduction": "The abundance of substructure in galaxy halos is emerging as a key test in theories of galaxy assembly.  In Cold Dark Matter cosmogonies, dark matter overdensities collapse to form cusped halos, with the smallest and least massive halos being the densest. The simulations of \\citet{Kl99} and \\citet{Mo99} predicted hundreds of small Galactic satellite halos, in contrast to the nine then known satellites around the Milky Way. \\citet{Ef92} had already suggested that photoionisation may lengthen the cooling times of gas in haloes with low circular speeds. This effect suppresses the formation of satellite galaxies, but produces a large population of entirely dark satellites~\\citep[see e.g.][]{Kr04,Mo06}.  In strong lensing, it has been known for some years that simple, smooth models of galaxy lenses usually fitted the image positions well, but the flux ratios of the images poorly. In a bold paper, \\citet{Da02} argued that the flux anomalies in gravitational lens systems could be interpreted as evidence for entirely dark substructures. They used 7 of the then available four-image lens systems to claim detection of substructure amounting to a total mass of $0.6 \\% - 6 \\%$ of the lens galaxy mass. However, this interpretation was challenged as flux anomalies could arise from alternative sources, such as absorption, scattering, or scintillation by the interstellar medium of the lens galaxy, or by higher order harmonics in the ellipsoidal models used to fit lens systems, or stellar microlensing~\\citep[see e.g.][]{Ev03,Ko04,Ma04}. In particular, the surface mass density in substructure as judged from simulations seems to be lower than that required by gravitational lensing, at least within the Einstein radius which probes primarily the inner parts of halos~\\citep[e.g.,][]{Ma04}. It is also hard to reproduce the observed statistics on cusp violations with substructure ~\\citep{Ma06}. This argues against substructure as a primary cause of anomalous flux ratios.  Nonetheless, there is good evidence in favour of substructure, and, in some cases, visible substructure can be identified.  One of the gravitational lens systems with a flux ratio anomaly is the radio-loud quadruple CLASS B2045+265 discovered by \\citet{Fa99}. Recent deep {\\it Hubble Space Telescope} and {\\it Keck} imaging of this system by \\citet{Mc07} have revealed the presence of a tidally disrupted dwarf galaxy G2.  This may be the cause of the flux ratio anomaly, although caution is needed as modelling suggests that G2 must be very highly flattened ($q = 0.13$).  There is also evidence for visible structure -- possibly a small galaxy -- in the radio-loud quadruple MG0414+0534. \\citet{Sc93} already argued that the perturbation caused by this object may account for the relatively poor agreement between the observed data on this lens and the theoretical models. The quadruple lens systems CLASS B1608+656 \\citep[see e.g.,][]{Fa96} has a loose group of galaxies at the same redshift as the main lensing galaxy, a phenomenon that also occurs in the six-image radio-loud system B1359+154~\\citep{Ru01}. It seems that visible substructure may well be responsible for some of the flux ratio anomalies.  In addition, there have been striking observational developments nearer to home. The last two years have seen the discovery of 10 faint, new Milky Way satellites in data from the Sloan Digital Sky Survey (SDSS, see Willman et al. 2005, Zucker et al. 2006a,b; Belokurov et al. 2006, 2007).  It seems likely that a population of ultra-faint, dwarf galaxies does surround the Milky Way.  These may be representatives of the ``missing satellites'', as predicted by the simulations of \\citet{Kl99} and \\citet{Mo99}, or they may be a population of tidal dwarf galaxies or even star clusters~\\citep[see e.g.,][]{Be07}. Such ultra-faint objects are only detectable nearby, and so would be -- to all intents and purposes -- dark at the typical redshifts of strong lenses.  All this suggests that it is worth re-examining the effects of substructure on strong lenses.  In this paper, we answer the following questions.  Given what we know about the satellite populations, how frequently might we expect anomalous flux ratios for elliptical and spiral galaxies?  If luminous satellite galaxies represent the bright and massive end of a predominantly faint population of objects, how frequently might we expect to attribute flux ratio anomalies to visible objects? Three (B2045+265, MG0414+0534, B1608+656) out of the sample of six quadruplets originally identified by \\citet{Da02} as anomalous have optically identified companions that are possible causes. At the outset, this seems surprisingly high, if the substructure is predominantly dark.  The paper is arranged as follows. In \\S 2, models of elliptical and spiral galaxies are briefly introduced, together with their satellite galaxy populations. In \\S 3, the frequency with which anomalous flux ratios occur is shown to depend primarily on the scalelength of the satellite distribution, the mass model used for the satellites and the total mass in substructure, while depending only weakly on how the mass is distributed between satellites. The simulations reported in \\S 4 give the expected fraction of anomalous flux ratios for both ellipticals and spiral lenses, together with the typical numbers caused by high mass and luminous satellite galaxies. ",
        "conclusions": "It remains unclear whether dark matter satellites and substructure are responsible for anomalous flux ratios in strong lensing. \\citet{Da02} originally studied a sample of 7 radio-loud four-image lens systems and claimed evidence of anomalies in 6 of them. This seemingly suggests that anomalous flux ratios are very common. Here, we have carried out a theoretical study of the frequency of flux ratio anomalies as a function of lensing galaxy and dark matter substructure parameters.  The likelihood that satellites affect flux ratios in strong lenses depends on their mass profile. Here, we considered compact Plummer spheres and diffuse, tidally stripped Hernquist profiles for our satellites, with the satellite size scaling as a power of mass.  As the Hernquist satellites are more extended than the Plummer models, they therefore have a smaller effect on image fluxes -- typically about 3 times the mass is needed to generate the same numbers of anomalous flux ratios.  The probability that strong lensing flux ratios are affected by satellites is crucially dependent on their spatial distribution. The characteristic lengthscale $\\Rfifty$ has a large effect on the fraction of lenses with affected flux ratios.  Our spatial distributions of satellite galaxies are inspired by the observational data on the Milky Way, for which the satellite number density falls off as $\\rhat^{-3.5}$ in three-dimensions with a lengthscale of $\\Rfifty \\sim 100 $ kpc. For such distributions, most satellites are too far out to affect the fluxes of images. For example, even with Plummer satellites, a mass of $\\sim 3\\E{9} \\Msun$ is needed for 20 \\% of quadruplets to show anomalous flux ratios for a typical elliptical galaxy, rising to $\\sim 1.25 \\E{10} \\Msun$ for 50 \\%.  Lenses that are edge-on spiral galaxies with maximum discs (like the Milky Way) are more resistant to flux changes by satellites, so the mass in satellites and substructure has to be roughly a factor of 3 times as great for the same proportion of quads to be affected. To obtain anything like the apparent abundance of anomalous flux ratios, then the scalelength of the substructure has to be different to what is known for the Milky Way satellites.  Whether the flux ratios in a lens system are affected by satellites is sensitive to the total mass in satellites, but more weakly dependent of how this mass is apportioned between them. For Plummer model satellites, the probability that a given satellite changes a flux ratio increases with its mass only slightly faster than linearly, at least when its mass is between $\\sim5\\times10^{6}\\Msun$ and $\\sim10^{9} \\Msun$. For example, satellites of mass $\\sim10^{9}\\Msun$ are only responsible for the causing $\\sim50\\%$ more flux ratio anomalies than those of mass $\\sim10^{7}\\Msun$. For Hernquist model satellites, more massive ones seem no more efficient (per unit mass) at changing fluxes; indeed, more massive dwarfs, being more prone to tidal disruption, might even be \\textit{less} efficient than lighter, more compact ones. One interesting consequence is that, if matter in dark satellites is not predominantly in the most massive ones, then the contribution of the most massive satellite galaxies to flux ratio anomalies should not be predominant.  In the light of this, the fact that so many anomalous flux ratios systems have optically identified substructure seems at outset surprising. There seem to be two possible explanations. First, the visible substructure may have been wrongly identified as the cause, whereas dark substructure may be the true culprit. A large dwarf galaxy may by chance be projected close to the Einstein radius, whereas unrelated dark substructure may be the major cause of the anomaly.  Second, visible satellites may be more concentrated than their dark cousins, a physical effect that may naturally arise from baryon condensation. Compression factors causing an enhancement of the central density by a factor of 10 in bright satellites seem to be ample to give a satisfactory explanation of the observed statistics. Nonetheless, such high compression factors are probably implausible except for the largest satellite galaxies. This seems to be in accord with the results of \\citet{Ma07}, who found that including baryons in numerical simulations did not help in reconciling simulation results with the statistics of anomalous flux ratios.   Finally, we remark that the likelihood that flux ratios are affected depends on the ellipticity of the main lens galaxy, but this dependence is much stronger in quads than doublets. This is a consequence of high magnification images being more easily affected by a dwarf galaxy than low magnification ones. Quads are more highly magnified than doublets (the disc triplets of spirals are in between), and changing the ellipticity of the main galaxy changes the typical magnification of quads more than it changes that of doublets. Generally, the greater the ellipticity the less the effect of satellite galaxies. For example, a given mass $\\Mdw$ of dwarf satellites around a typical edge-on maximum-disc spiral galaxy is significantly less likely to change image flux ratios than $\\Mdw$ around a typical elliptical galaxy, even though the spiral is less massive than the elliptical."
    },
    "0801/0801.1401_arXiv.txt": {
        "abstract": "We present semi-analytic models of galactic outflows that are constrained by available observations on high redshift star formation and reionization. Galactic outflows are modeled in a manner akin to models of stellar wind blown bubbles. Large scale outflows can generically escape from low mass halos ($M\\lesssim 10^9~M_\\odot$) for a wide range of model parameters while this is not the case in high mass halos ($M\\gtrsim 10^{11}~M_\\odot$). The flow generically accelerates within the halo virial radius, then starts to decelerate, and traverses well into the intergalactic medium (IGM), before freezing to the Hubble flow. The acceleration phase can result in shell fragmentation due to the Rayleigh-Taylor instability, although the final outflow radius is not significantly altered. The gas phase metallicity of the outflow and within the galaxy are computed assuming uniform instantaneous mixing. Ionization states of different metal species are calculated and used to examine the detectability of metal lines from the outflows. The global influence of galactic outflows is also investigated using porosity weighted averages and probability density functions of various physical quantities. Models with only atomic cooled halos significantly fill the IGM at $z\\sim3$ with metals (with $-2.5\\gtrsim [Z/Z_\\odot]\\gtrsim -3.7 $), the actual extent depending on the efficiency of winds, the initial mass function (IMF) and the fractional mass that goes through star formation. The reionization history has a significant effect on the volume filling factor, due to radiative feedback. In these models, a large fraction of outflows at $z\\sim3$ are supersonic, hot ($T\\ge 10^5$K) and have low density, making metal lines difficult to detect. They may also result in significant perturbations in the IGM gas on scales probed by the Lyman-$\\alpha$ forest. On the contrary, models including molecular cooled halos with a normal mode of star formation can potentially volume fill the universe at $z\\ge 8$ without drastic dynamic effects on the IGM, thereby setting up a possible metallicity floor ($-4.0\\le [Z/Z_\\odot]\\le-3.6$). The order unity fluctuations at $z\\sim8$ that becomes the mildly non-linear fluctuations traced by Lyman-$\\alpha$ forest at $z<4$ will then have this metallicity. Interestingly, molecular cooled halos with a ``top-heavy'' mode of star formation are not very successful in  establishing the metallicity floor because of the additional radiative feedback, that they induce. ",
        "introduction": "The rapid growth of observations of the high redshift universe has raised several intriguing questions regarding physics of galaxy formation and the physical state of the intergalactic medium (IGM). Some of the important issues are:  how and when the dark ages ended with the reionization of the IGM, the origin of the metals, and temperature of the low density  IGM traced by Lyman-$\\alpha$ forest at $z\\simeq2.5$. In this work we concentrate on the issue of IGM metal enrichment.   Presence of metals in the Lyman-$\\alpha$ forest (with $\\tau\\ge1$) is now well established though C~{\\sc iv}  (Tytler et al. 1995; Songaila \\& Cowie, 1996) and O~{\\sc vi} (Carswell et al. 2002; Simcoe et al. 2002;  Bergeron et al. 2002) absorption lines  detected in the echelle spectra of high redshift QSOs. The measured N(C~{\\sc iv})/N(H~{\\sc i}) at $z\\sim3$ are consistent with [C/H]$\\sim-2.5$ with large errors (see Rauch et al. 1997). Songaila (2001) has shown that the C~{\\sc iv} column density distribution function is consistent with being invariant  between $z=1.5$ and $z=5$ and a minimum metallicity  of [C/H]$\\sim-3.5$ is already in place at $z \\sim 5.5$ (see also Ryan-Weber et al. 2006). As under-dense space occupy most of the volume, measuring metallicity in the low density region is very important. Given the  expected low metallicities and large ionization corrections, direct detection of metals from such low neutral hydrogen optical depth (i.e $13\\le$~log~[N(H~{\\sc i})~{cm}$^{-2}$]~$\\le 14$) is currently impossible and usually the metallicity is estimated using pixel statistics \\cite{ellison00,schaye03,aracil04}.  Schaye et al. (2003) obtained [C/H]$\\sim-3.5$ for $\\log \\delta >-0.5$ at $z = 1.8-4.1$. Here, $\\delta$ being the over density defined as $\\delta =\\rho/\\bar{\\rho}\\ge 10 (N(H~{\\sc i})/10^{15} {\\rm cm}^{-2})^{2/3} [(1+z)/4]^{-3}$ (Schaye 2001) and $\\bar{\\rho}$ is the mean IGM density. However, Aracil et al. (2004) have not detected C~{\\sc iv} from $\\tau < 1$ Lyman-$\\alpha$ systems in their UVES data. Thus,  what fraction of the  IGM is filled with metals is a subject of ongoing debates and investigations (Schaye et al. 2003; Aracil et al. 2004; Scannapieco et al. 2006). In any case, in the standard framework of LCDM models the Lyman-$\\alpha$ forest absorption lines originate from density fluctuations that are either in the linear or in the quasi-linear regime (Bi \\& Davidsen, 1997). Thus, metals that may be present in these regions need to be transported from the neighboring star-forming regions. The amount and  distribution of metals in the Lyman-$\\alpha$ forest provides information on different feedbacks from star forming galaxies.  Observations of high-$z$ Lyman break galaxies (LBGs) frequently show galactic scale superwinds (Pettini et al. 2001), high value for the escape fraction of the UV photons (Steidel et al. 2001) and a strong correlation between the C~{\\sc iv} absorption systems and the LBGs at $z\\sim3$ within an impact parameter of $\\sim 50$ kpc (Adelberger et al. 2005).  Profiles of C~{\\sc iv} and O~{\\sc vi} absorption lines seen in the high redshift damped Lyman-$\\alpha$ systems (DLAs) are also consistent with them originating from outflows in DLA protogalaxies (Fox et al. 2007, 2007a). Adelberger et al. (2005) have argued that considerable fraction of high  column density C~{\\sc iv} systems may originate from the large scale galactic outflows. Clustering properties of strong C~{\\sc iv} absorption lines (Scannapieco et al. 2006) and the velocity profiles C~{\\sc iv} absorption lines (Songaila et al. 2006) are consistent with a good fraction of the high column density C~{\\sc iv} systems originating from the region near massive galaxies. However, the important question is whether these metals are freshly emitted by the bright galaxies or due to biased clustering of gas ejected by low mass galaxies from early (say $z >6$) epochs (Porciani \\& Madau, 2005). Aguirre et al. (2005) by comparing spectra predicted by various simulations that includes winds found that the predicted metal distribution in the models are highly inhomogeneous and can not reproduce the observed probability distribution of C~{\\sc iv} absorption. They suggested that strong winds from galaxies at $z \\le6$ cannot fully explain the observed enrichment and additional pre-enrichment from higher redshift galaxies are needed.   Note that star formation is the key element in controlling the outflows from galaxies. Star formation also has an important effect on the reionization of the universe which in turn affects the star formation in low mass halos through radiative feedback. There exists a growing body of data constraining the star formation rate (SFR) density in the high redshift universe \\cite{bouwens05,richard06,hopkins06}. Also observational constraints on the reionization are provided by the spectra of highest redshift QSOs (Fan et al. 2006) and ongoing WMAP satellite observations of Cosmic Microwave Background (CMB) polarization (Spergel et al. 2007). Therefore one must deal with reionization and galactic outflows in a manner consistent with the observed star formation history of the universe. In our previous paper Samui, Srianand \\& Subramanian (2007; here after Paper I) we have described a set of models which correctly produce the observed luminosity functions of LBGs and hence the SFR density at different redshifts. Here we take those constrained models of star formation and reionization from Paper I and used them to predict the various properties galactic outflows and their impact on the IGM such as volume filling factor of wind/metals, mean metallicity, temperature of these polluted region etc.  There are several attempts to model the galactic outflow using semi-analytic calculations (for example, Madau, Ferrara \\& Rees 2001; Scannapieco, Ferrara \\& Madau 2002; Furlanetto \\& Loeb 2003; Scannapieco 2005). The model presented here will be broadly in line with these above studies with significant improvements. In particular, the star formation is continuous and duration of the star formation and the amount of gas going into stars are constrained by the observed high-$z$ luminosity function (see Paper I and section 2 below). This is very different from the previous attempts where the star formation is usually in the form of bursts. Our model takes into account the reionization and the radiative feedback to the star formation in a self-consistent way. This is very important, in particular, to estimate the metal pollution due to low mass halos and constrain the initial mass function (IMF) at high redshifts. We discuss various other improvements and their effects in detail.  The paper is organized in following manner. In the next section we first briefly outline the model of star formation. In section 3 we introduce our model of galactic outflows in detail. The structural properties of the outflow and  sensitivity of our results to the adopted initial conditions are discussed in detail in section 4.  In section 5 we discuss the dependence of outflow properties on the model parameters. In section 6 we investigate the growth of Rayleigh-Tayler instabilities in the accelerating phase of the outflow and study its consequences. The ionization correction and the detectability of the winds in different stages of its evolution are discussed in section 7. The global effects of the outflow are studied in section 8 and the summary and conclusions are provided in section 9. In most of this work we use the cosmological parameters consistent with the recent WMAP data ($\\Omega=1$, $\\Omega_m = 0.26$, $\\Omega_\\Lambda = 0.74$, $\\Omega_b=0.044$, $h = 0.71$, $\\sigma_8=0.75$ and $n_s =0.95$). ",
        "conclusions": "We have examined in some detail here semi-analytic models of galactic outflows and their consequences for the intergalactic medium. Our models are constrained by available observations of both star formation and reionization. The nature of star formation is one of the key elements which controls the energetics of galactic outflows and also their metal content. At the same time, it also decides the reionization history which is an important input for the radiative feedback that suppresses star formation in low mass halos. We improve on earlier semi-analytical modeling of galactic outflows in several important ways. Galaxies form stars continuously with the duration and fraction of baryons going into stars being constrained by the observed high-$z$ UV luminosity functions (see Paper I for details). We take into account of existing constraints on reionization and most importantly, the implied radiative feed back on the star formation in a self-consistent manner. We adopt the WMAP 3rd year cosmological parameters. We model galactic outflows in a manner similar to stellar wind blown bubbles (cf. Weaver et al. 1977), following the dynamics of both the outer shock $R_s$ and also a possible inner reverse shock at $R_1$. This can naturally incorporate possible smooth transitions from pressure to momentum driven outflows. We use the modified PS formalism of \\cite{sasaki94} formalism to calculate the formation rate of dark matter halos and the global consequences of outflows, instead of taking just the time derivative of the PS function, which does not account for the destruction rate of halos.  Outflows generically accelerate initially due to the increasing energy input from the galaxy and decreasing halo density profile. The outer shock velocity reaches a peak value which increases with halo mass, typically ranging from $\\sim 100-400$ km s$^{-1}$ for halo masses $\\sim 10^7-10^{12} M_\\odot$ respectively. This phase lasts for a dynamical time-scale after which the outflow decelerates, till it becomes subsonic and freezes to the Hubble flow. The hot bubble of shocked wind material, has initial temperatures $\\gtrsim 10^6$ K but subsequently cools due to adiabatic expansion. If there is significant mass loading from the galaxy, it can also cool radiatively to transit to a momentum driven flow.  However, for most model parameters here, this does not occur. The swept up shell gas typically cools efficiently while the outflow traverses the halo. During the acceleration phase the thin shell is also unstable to R-T instability leading to shell fragmentation (see also Ferrara \\& Ricotti (2006)) without significantly altering the final outflow radius. Such shell fragmentation could enhance the detectability of metal lines from outflows by providing evaporating interfaces where the hot metal-enriched bubble gas mixes with cool, dense shell material.  During its evolution, individual outflows can travel well beyond the virial radius of the host halo to proper distances $\\sim 200-1000$ kpc for the above halo mass range. The inner shock at $R_1$ keeps up with the outer shock while the galaxy is actively forming stars, with typical values of $R_1/R_s \\sim 0.4-0.6$. This is very similar to the wind structure seen in the simulations of Fujita et al. (2004). By the time the outflow becomes subsonic and freezes onto the Hubble flow, both the hot bubble and shell temperatures are $\\sim 10^4$~K (determined now by photoheating), and the shell is likely to fragment and mix with the bubble and IGM gas.   We explored in some detail the dependence of outflow properties on the assumed initial conditions and various model parameters. We show that the initial conditions play very little role in deciding the nature of the outflows.  We also check this by comparing our model predictions with scale-free solutions.  We find that outflows can generically escape from the low mass halos ($ M \\lesssim 10^9 M_\\odot$), that dominantly contribute to the volume filling of the IGM. For galactic scale halos, as expected, having higher halo density $f_h$, higher wind mass loading $\\eta$ or lower energy input efficiency ($\\epsilon_w\\nu f_*$) makes it more difficult for outflows to escape. A burst mode of star formation generically leads to a smaller outflow radius, even for the same values of other parameters. As the outflow properties of low mass halos are less sensitive to above mentioned parameters the nature and efficiency of star formation in these objects decide the feedback due to galactic winds.  The detection of metals from the outflows and in the IGM is one of the crucial issues in our paper. The gas phase metallicity in the ISM and the wind are self-consistently computed for a given star formation rate and IMF assuming instantaneous uniform mixing in the ISM. The detection of the expelled metals either in different stages of the outflow or in the IGM depends crucially on the ionization state of the gas.  Using photoionization calculations performed using Cloudy we have shown that the metals in free wind and low density bubble will be very difficult to detect through standard UV  absorption lines of C~{\\sc iv} and Si~{\\sc iv} and O~{\\sc vi}. We need clumped high density gas either coming from the ISM in the form of free wind or from the R-T instabilities to detect these absorption lines. In particular some of these high ionization species can be detected in the conductive interfaces between the cold clumps and hot bubble material. We show that the metals in the underdense regions ($\\rho\\sim 0.1\\bar{\\rho}$) will be very difficult to detect if the temperatures are higher than 10$^5$ K. Thus fresh outflow from galaxies that enter the IGM will be very difficult to detect.  C~{\\sc iv} and O~{\\sc vi} absorption lines are easily detectable if they originate from overdense regions ($\\rho>10\\bar{\\rho}$) that are typically probed by high column density Lyman-$\\alpha$ absorption line with $T\\sim10^4$~K. This is possible if IGM is already filled with pre-enriched gas by the time these over densities were order unity fluctuations (say $z\\ge 8$).   One of the important issues is to understand is how outflows impact on the physical properties of the intergalactic medium. We therefore computed the volume of the IGM affected by outflows, the porosity weighted averages and PDFs of several important outflow characteristics, for a number of atomic and molecular cooling models.  For our fiducial atomic cooling model A, more than $30\\%$ of the universe is affected by the outflows even at $z\\sim 6$ and this increases to $\\sim 60\\%$ by $z \\sim 2$. Galaxies with mass range $10^{7}-10^{9}~M_\\odot$ dominantly contribute to the volume filling factor; higher masses only dominate at low redshifts. This is consistent with the suggestion of Madau, Ferrara and Rees (2001), of the dominant influence of halos of $\\sim 10^8 M_\\odot$ in filling the universe with outflows; our work however includes halos of all mass ranges. Further, these galaxies which dominate in filling the IGM are not yet detected directly in the high-$z$ UV luminosity functions (see Fig.~4. in Paper I), The porosity averaged outflow comoving radius, peculiar velocity and bubble temperature, evolve from $\\sim 100 $~kpc, $\\sim 60 $~km~s$^{-1}$ and $10^6$~K, respectively, at $z \\sim 10$ to $\\sim 500$~kpc $20$~km~s$^{-1}$ and $\\sim 10^5$~K at $z \\sim 3 $. The median value of the outflow radius is $60 $~kpc (proper) at $z =5$ which increases to $100$~kpc at $z=3$. Whereas at $z=5$ , more than 60\\% of the volume filled by bubbles are at a temperature higher than $10^5$~K this fraction decreases to less than 15\\% by $z=3$. This is mainly due to the adiabatic expansion of the bubbles. The void regions these bubbles fill will nevertheless be at a higher temperature than the photoionized IGM. Most of the bubbles have metallicities between $0.01-0.1~Z_\\odot$, with the differential PDF peaked around the lower value. The average global metallicity evolution is also of interest. For our fiducial model, this gradually builds up from about $\\bar{Z} \\sim 10^{-5}~Z_\\odot$ at $z\\sim 10$ to $\\bar{Z} \\sim 2 \\times 10^{-3}~Z_\\odot$ at $z \\sim 2$, by which time the porosity has just exceeded unity.  We have examined several other atomic cooling models which are all consistent with the constraints on star formation obtained in paper I. These models also lead to significant filling of the IGM at $z\\sim3$ with metals (with $-2.5\\gtrsim [Z/Z_\\odot]\\gtrsim -3.7$), the actual extent depending on the efficiency of winds, the initial mass function (IMF) and the fractional mass that goes through star formation and cosmological parameters. The reionization history has a significant effect on the volume filling factor, due to radiative feedback (see also Pieri et al. 2007). Further, a large fraction of outflows at $z\\sim3$ are supersonic, hot ($T\\ge 10^5$~K) and have low density, making metal lines difficult to detect. These models may also result in significant perturbations in the IGM gas on scales probed by the Lyman-$\\alpha$ forest.  On the other hand, we find that models including star formation in molecular cooled halos with a normal mode of star formation (or a lower UV escape fraction) can potentially volume fill the universe at $z\\ge 8$ without drastic dynamic effects on the IGM, thereby setting up a possible metallicity floor ($-4.0\\le [Z/Z_\\odot]\\le-3.6$). In fact, the hubble frozen outflows dominate below this redshift and the $Q$ contributed by such outflows itself becomes of order unity below $z\\sim 6$. On the average, very low mass halos, with mass $\\sim 10^6-10^7~M_\\odot$ at $z=8$ to $\\sim 10^8-10^9~M_\\odot$ at $z=3$ dominantly contribute to the volume filling factor. The porosity averaged comoving radius of the outflows is less than 100 kpc at redshifts where the $Q\\sim 1$. The bubble temperature and peculiar velocities are also smaller than for the fiducial atomic cooling model. The above features make this model ideal to spread metals into the regions which will subsequently collapse to form the Lyman alpha forest regions, without unduly disturbing these regions dynamically. To some extent this scenario is the extrapolation of the scenario of Madau, Ferrara and Rees (2001) which was applied to outflows from dwarf galaxies, to even lower mass halos. Interestingly, molecular cooled halos with a ``top-heavy'' mode of star formation are not very successful in  establishing the metallicity floor because of the additional radiative feedback, that they induce.  As we discussed above we use a functional form for SFR (constrained by the observations of high-$z$ UV luminosity functions) without doing self-consistent calculations. However, to get reliable results from the self-consistent calculations one needs to be very clear about various physical processes that are involved. This is reflected in the fact that two recent simulations addressing this issue conclude differently. Scannapieco et al. (2006) find that supernova feedback and the resulting outflows decrease the fraction of baryons which are turned into stars, by factor ranging from 2-4 as one changes the mass of the object, and also a related more rapid fall in the star formation rate with time. However, Koboyashi et al. (2007) find somewhat different results. They note that the two effects of supernovae, the increased metal line cooling due to the chemical feedback and the increased heating due to the energy feedback both have opposite effects on the star formation rate in the galaxy. And indeed these two effects seem to cancel to produce no net effect due to supernovae feedback on the star formation rate (see their Fig.~2 and the discussion). Therefore it is not entirely clear from these works the extent to which one needs to change $f_*$ and $\\kappa$. If SNe produce strong feedback effects as found by Scannapieco et al. (2006) the star formation rate will fall sharply as a function of time. This corresponds to our models with low $\\kappa$ that requires low value of $f_*$ in order to reproduce the UV luminosity functions.   Future improvement of our work would involve replacing Eq.~(\\ref{eqnsf}) with a model of star formation in a multiphase ISM including various heating and cooling processes, possible effects of dark halo clustering on the outflow properties  and importantly setting up our semi-analytical model in the framework of a large-scale structure simulation."
    },
    "0801/0801.4362_arXiv.txt": {
        "abstract": "We discuss how rotation and binary interactions may be related to the diversity of type Ibc supernovae and long gamma-ray bursts. After presenting recent evolutionary models of massive single and binary stars including rotation, the Tayler-Spruit dynamo and binary interactions, we argue that the nature of SNe Ibc progenitors from binary systems may not significantly differ from that of single star progenitors in terms of rotation, and that most long GRB progenitors may be produced via the quasi-chemically homogeneous evolution at sub-solar metallicity. We also briefly discuss the possible role of magnetic fields generated in the convective core of a massive star for the transport of angular momentum, which is potentially important for future stellar evolution models of supernova and GRB progenitors. ",
        "introduction": "Rotation influences not only the evolution of massive stars, but also their supernova (SN) explosions \\citep{Maeder00,Heger00}. In particular, recently many asymmetric supernovae with unusually large energy (broad-lined SNe or hypernovae) have been discovered (e.g. \\citealt{Mazzali07}), which shows evidence for rapidly rotating progenitors. The most spectacular example may be long gamma-ray bursts (GRBs), which are generally believed to be produced by deaths of some massive stars retaining extremely large angular momenta in their cores ($j \\ge \\sim 10^{16}~\\mathrm{cm^2~s^{-1}}$; \\citealt{Woosley93, Macfadyen99}; see, however, \\citealt{Dessart08}). Interestingly, such energetic core-collapse events seem to only occur in Wolf-Rayet (WR) stars: all of the broad-lined SNe/Hypernovae and the supernovae associated with long GRBs have been observationally identified as Type Ic (see \\citealt{Woosley06b} for a review). This raises the question which WR stars can produce broad-lined SNe Ic or long GRBs while most WR stars die as normal SNe Ibc. Here we present recent evolutionary models of massive stars that include binary interactions and the transport of chemical species and angular momentum via rotationally induced hydrodynamic instabilities and magnetic torques, and discuss how rotation and binary interactions may be related to the diversity of SNe Ibc and long GRBs. ",
        "conclusions": "\\begin{figure} \\begin{center} \\resizebox{0.85\\hsize}{!}{\\includegraphics{visit.eps}} \\caption{Mean radial fields $B_\\mathrm{r} (r, \\theta)$ on the meridional plane in a $12~\\mathrm{M_\\odot}$ rotating star on the main sequence in a MHD simulation with a 3-D anelastic code \\citep{Glatzmaier84}. The adopted angular velocity is $10^{-5}~\\mathrm{Rad~s^{-1}}$. The inner region of $r \\le 7 \\times 10^{10}~\\mathrm{cm}$ ($r \\le 0.5$ in the code units) is the convective core. }\\label{fig:bfields} \\end{center} \\end{figure}  Our stellar evolution models including the Tayler-Spruit dynamo indicate that most SNe Ibc progenitors should explode with a similar amount of angular momentum in their cores, regardless of their single or binary star origin. This is due to the self-regulationary nature of the Tayler-Spruit dynamo. Loss of the hydrogen envelope due to stellar winds or mass transfer results in both removal of angular momentum from the core and weakening of the core braking by the extended hydrogen enveloped due to magnetic torques, and vice versa. Therefore, most different pre-supernova evolutionary paths may not contribute much to the diversity of Type Ibc supernovae in terms of core angular momentum, although different iron core masses, and thus different spin rates of young neutron stars may result \\citep{Heger05}. GRB progenitors, which require unusually large angular momenta, may undergo the chemically homogeneous evolution, which may not be unusual at low metallicity. On the other hand, recent observations indicate that not all broad-lined SNe Ic are associated with long GRBs (e.g. \\citealt{Modjaz08}), which needs a theoretical explanation in future work.  Although the predicted spin rates of the stellar remnants from the magnetic models are consistent with observations, the validity of the Tayler-Spruit dynamo has been recently questioned by several authors \\citep{Denissenkov07, Zahn08}. Furthermore, we might have ignored potentially  important physical ingredients in simulating the evolution of rotating massive stars. These include gravity waves (Townsend in this volume), and magnetic fields generated by the convective core. For instance, our recent 3-D simulations with an anelastic magnetohydrodynamics code \\citep{Glatzmaier84} show that the strength of poloidal fields generated in the convective core in a young massive star may amount to several thousand Gauss on average (Fig.~\\ref{fig:bfields}), and its influence on the transport processes might be comparable to what the Tayler-Spruit dynamo predicts. This issue will be addressed in Yoon, Woosely \\& Glatzmaier (2008, in prep.).  This work was, in part, supported by the DOE Program for Scientific Discovery through Advanced Computing and NASA."
    },
    "0801/0801.3297_arXiv.txt": {
        "abstract": "We present a measurement of the rate of type Ia supernovae (SNe Ia) from the first of three seasons of data from the \\sns. For this measurement, we include 17 SNe Ia at redshift $z\\le0.12$. Assuming a flat cosmology with $\\Omega_m = 0.3=1-\\Omega_\\Lambda$, we find a volumetric SN Ia rate of $[2.93^{+0.17}_{-0.04}({\\rm systematic})^{+0.90}_{-0.71}({\\rm statistical})] \\times 10^{-5}~{\\rm SNe}~{\\rm Mpc}^{-3}~h_{70}^3~{\\rm year}^{-1}$, at a volume-weighted mean redshift of 0.09. This result is consistent with previous measurements of the SN Ia rate in a similar redshift range. The systematic errors are well controlled, resulting in the most precise measurement of the SN Ia rate in this redshift range. We use a maximum likelihood method to fit SN rate models to the \\sns~data in combination with other rate measurements, thereby constraining models for the redshift-evolution of the SN Ia rate. Fitting the combined data to a simple power-law evolution of the volumetric SN Ia rate, $r_V \\propto (1+z)^{\\beta}$, we obtain a value of $\\beta = 1.5 \\pm 0.6$, i.e. the SN Ia rate is determined to be an increasing function of redshift at the $\\sim 2.5 \\sigma$ level. Fitting the results to a model in which the volumetric SN rate, $r_V=A\\rho(t)+B\\dot \\rho(t)$, where $\\rho(t)$ is the stellar mass density and $\\dot \\rho(t)$ is the star formation rate, we find $A = (2.8 \\pm 1.2) \\times 10^{-14} ~\\mathrm{SNe} ~\\mathrm{M}_{\\sun}^{-1} ~\\mathrm{year}^{-1}$, $B = (9.3^{+3.4}_{-3.1})\\times 10^{-4} ~\\mathrm{SNe} ~\\mathrm{M}_{\\sun}^{-1}$. ",
        "introduction": "\\label{sec:intro} Type Ia supernovae (SNe Ia) have gained increasing attention from astronomers, primarily due to their remarkable utility as cosmological distance indicators. There is now broad consensus that a type Ia supernova is the thermonuclear explosion of a Carbon-Oxygen white dwarf star that accretes mass from a binary companion until it reaches the Chandrasekhar mass limit (e.g.~\\citet{Branch_95}). However, much remains to be learned about the physics of SNe Ia, and there is active debate about both the nature of the progenitor systems and the details of the explosion mechanism. For example, the binary companion may be a main-sequence star, a giant or sub-giant, or a second white dwarf. The type of the companion star determines in part the predicted time delay between the formation of the binary system and the SN event~\\citep{Greggio_05}. The time delay can be constrained observationally by comparing the SN Ia rate as a function of redshift to the star formation history (SFH)~\\citep{Strolger_04,Cappellaro_07}.  The insight into the nature of the progenitor systems that SN Ia rate measurements provide can also potentially strengthen the utility of SNe Ia as cosmological distance indicators. Although the strong correlation between SN Ia peak luminosity and light curve decline rate was found purely empirically \\citep{Psk_77,Phillips_93}, the physics underlying this relation has been extensively studied \\citep{Hoeflich_95,Hoeflich_96,Kasan_07}. There is hope that improved physical understanding and modeling of SN Ia explosions, coupled with larger high-quality observational data sets, will lead to improved distance estimates from SNe Ia. As part of this program, deeper understanding of the nature of the progenitor systems can help narrow the range of initial conditions that need to be explored in carrying out the costly simulations of SN Ia explosions that in principle predict their photometric and spectroscopic properties.  Measurement of the SN Ia rate may also have a more direct impact on the determination of systematic errors in SN Ia distance estimates. \\citet{Mannucci_05,ScanBildsten,Neill_06} and \\citet{Sullivan_06} have argued that a two-component model of the SN Ia rate, in which a prompt SN component follows the star formation rate and a second component follows the total stellar mass, is strongly favored over a single SN Ia channel. In this picture, since the cosmological star formation rate increases sharply with lookback time, the prompt component is expected to dominate the total SN Ia rate at high redshift. \\citet{Mannucci_05} and \\citet{Howell_07} pointed out that this evolution with redshift can be a potential source of systematic error in SN Ia distance estimates, if the two populations have different properties.  In order to test such a model for the evolution of the SN Ia rate, improved measurements of the rate as a function of redshift and of host galaxy properties are needed. The Supernova Legacy Survey (SNLS) has recently presented the most precise measurement of the SN Ia rate at high redshift ($z \\sim 0.5$) based on 73 SNe~Ia \\citep{Neill_06,Sullivan_06}. At low redshifts ($z\\sim 0.1$), SN Ia rate measurements~\\citep{Cappellaro_99,Hardin_00,Madgwick_03,Blanc_04} have suffered from small sample sizes and also from systematic errors associated with heterogeneous samples \\citep{Cappellaro_99} and with selection biases due to the targeting of known, relatively luminous galaxies \\citep{Hardin_00,Blanc_04}. The low-redshift measurement of \\citet{Madgwick_03}, based on SNe discovered fortuitously in SDSS galaxy spectra, is affected by different systematic uncertainties than traditional photometric searches, e.g., due to the finite aperture of the SDSS spectroscopic fibers.  In this paper, we present a new measurement of the SN Ia rate at low redshift, based upon the first season of data from the \\sns. The \\sns~\\citep{Frieman_08} offers several advantages for this measurement. It covers a larger spatial volume than previous SN surveys, a result of the combination of intermediate-scale (2.5-m) telescope aperture, wide field of view (3 square degrees), modest effective sidereal exposure time (54 sec), and use of drift-scanning to efficiently cover a large sky area ($\\sim 300$ square degrees). The \\sns~is a rolling search, with new SNe discovered simultaneously with the follow-up of previously discovered SNe. Unlike SN searches that target known galaxies, the \\sns~is not biased against finding SNe in low-luminosity host galaxies. Well-calibrated photometry in the SDSS {\\it ugriz} passbands \\citep{Fukugita_96}, with a typical interval between observations of four days, yields well-sampled, multi-band light curves that enable photometric typing of SNe with high confidence. Moreover, rapid on-mountain photometric reduction and image processing coupled with an extensive spectroscopic follow-up program enable spectroscopic confirmation of a very high fraction of the low-redshift SN Ia candidates.  The \\sns~was carried out over three three-month seasons, during Sept.-Nov. 2005-7. The results presented here are based on the Fall 2005 season. The \\sns~measures light curves for SNe Ia to redshift $z \\simeq 0.4$, with a median redshift of $\\langle z \\rangle = 0.22$ for spectroscopically confirmed SNe Ia. However, in this first paper we limit the analysis to low redshift, $z \\leq 0.12$, since our spectroscopic follow-up is essentially complete over this redshift range, and the uncertainty due to spectroscopically unobserved (and untyped) SNe is therefore negligible. In presenting a SN Ia rate measurement, one must decide whether to include peculiar SNe Ia, i.e., events that are photometrically and/or spectroscopically unusual, since it is not clear that they are members of the same population as the ``normal'' SNe Ia. Formerly, the peculiar designation included events such as 1991T and 1991bg, which are highly overluminous and underluminous events, respectively. However, since these SNe appear to follow the standard peak-luminosity/decline-rate relation, they are now generally considered extreme members of the normal SN Ia population \\citep{Nugent_95}. Other events, such as 2002ic \\citep{Hamuy_03} and 2002cx \\citep{Li_03}, exhibit more pronounced peculiarities and do not fit the luminosity-decline relation. The first season  of the \\sns~included two such truly peculiar events at low redshift, 2005hk~\\citep{Phillips_07} and 2005gj~\\citep{Prieto_07, Aldering_06}. Although these peculiar events may arise from the same evolutionary path as normal SNe Ia, which would argue for including them in a SN Ia rate measurement, we have chosen to include only SNe with light curves that obey the standard brightness-decline relation. More specifically, we include in our rate measurement sample only SNe with light curves that are well described by the MLCS2k2 SN Ia light curve model~\\citep{Riess_96,Jha_07}, see \\S \\ref{sec:ssample}. Regardless of the physical arguments surrounding peculiar events, we exclude them from this analysis primarily because we do not yet have a robust determination of our efficiency for detecting them.  The rest of the paper is organized as follows. In \\S\\ref{sec:survey} we provide a brief outline of the survey observing strategy and operations as they relate to the rate determination. In \\S\\ref{sec:sample} we define selection criteria and present the sample of SNe Ia used in this measurement, based on spectroscopic and photometric measurements. In \\S\\ref{sec:effs} we present estimates of the detection efficiency for low-redshift SNe Ia, based on artificial SNe inserted into the survey images and on Monte Carlo simulations. We present our measurement of the SN Ia rate and discuss the SN Ia rate as a function of host galaxy type in \\S\\ref{sec:results}. In \\S\\ref{sec:rate_models} we compare our result to other SN Ia rate measurements and combine rate measurements to fit semi-empirical models of rate evolution. ",
        "conclusions": ""
    },
    "0801/0801.1547_arXiv.txt": {
        "abstract": "{This paper reviews the theory, phenomenology, and observational constraints on the coupling parameters of Einstein-\\ae ther gravity, i.e.\\ General Relativity coupled to a dynamical unit timelike vector field. Several new remarks and a discussion of open issues are included.}  \\FullConference{From Quantum to Emergent Gravity: Theory and Phenomenology\\\\ June 11-15 2007\\\\ Trieste, Italy}  \\begin{document}   \\def\\H{{\\cal H}} \\def\\ttheta{\\tilde{\\theta}}   \\def\\beq{\\begin{equation}} \\def\\eeq{\\end{equation}} \\def\\bea{\\begin{eqnarray}} \\def\\eea{\\end{eqnarray}} \\def\\ben{\\begin{enumerate}} \\def\\een{\\end{enumerate}} \\def\\la{\\langle} \\def\\ra{\\rangle} \\def\\a{\\alpha} \\def\\b{\\beta} \\def\\g{\\gamma}\\def\\G{\\Gamma} \\def\\d{\\delta} \\def\\e{\\epsilon} \\def\\phi{\\varphi} \\def\\k{\\kappa} \\def\\l{\\lambda} \\def\\m{\\mu} \\def\\n{\\nu} \\def\\o{\\omega} \\def\\p{\\pi} \\def\\r{\\rho} \\def\\s{\\sigma} \\def\\z{\\zeta} \\def\\t{\\tau} \\def\\L{{\\cal L}} \\def\\S{\\Sigma } \\def\\gsim{\\; \\raisebox{-.8ex}{$\\stackrel{\\textstyle >}{\\sim}$}\\;} \\def\\lsim{\\; \\raisebox{-.8ex}{$\\stackrel{\\textstyle <}{\\sim}$}\\;} \\def\\gtrsim{\\gsim} \\def\\lessim{\\lsim} \\def\\loc{{\\rm local}} \\def\\vm{v_{\\rm max}} \\def\\bh{\\bar{h}} \\def\\del{\\partial} \\def\\nab{\\nabla} \\def\\half{{\\textstyle{\\frac{1}{2}}}} \\def\\fourth{{\\textstyle{\\frac{1}{4}}}} ",
        "introduction": "There are reasons to suspect that the vacuum in quantum gravity may determine a preferred rest frame at the microscopic level. However, if such a frame exists, it must be very effectively concealed from view. Numerous observations severely limit the possibility of Lorentz violating physics among the standard model fields~\\cite{Mattingly:2005re}. The constraints on Lorentz violation in the gravitational sector are generally far weaker.  To allow for gravitational Lorentz violation without abandoning the framework of general relativity (GR), the background tensor field(s) breaking the symmetry must be dynamical. Einstein-\\ae ther theory is of this type. In addition to the spacetime metric tensor field $g_{ab}$ it involves a dynamical, unit timelike vector field $u^a$. Like the metric, and unlike other classical fields, the unit vector cannot vanish anywhere, so it breaks local Lorentz symmetry down to a rotation subgroup. It defines a congruence of timelike curves filling all of spacetime, like an omnipresent fluid, and so has been dubbed the ``\\ae ther\". This paper aims to provide a review of what has been learned about this theory, the status of observational constraints, and some currently open issues.  A primary motivation for studying Einstien-\\ae ther theory---``\\ae -theory \" for short---is the quantum gravity suspicion mentioned above. A secondary aim is to develop a viable and reasonably natural foil against which to compare gravitational observations, in an era when numerous alternate gravity theories have already been either ruled out or severely constrained. A third source of interest is the theoretical laboratory it offers for studying diffeomorphism invariant physics with preferred frame effects.  The action involving metric and \\ae ther is highly constrained by diffeomorphism invariance, locality, and the unit constraint on $u^a$. The only term with no derivatives is the cosmological constant, there are no terms with one derivative (other than a total divergence), and there are five terms with two derivatives, \\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 Here $R$ is the Ricci scalar, and the tensor $K^{ab}_{mn}$ is defined by \\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 where the $c_i$ are dimensionless coupling constants.  Like general relativity, pure classical \\ae -theory defined by this action is scale-free. (The metric signature is $({+}{-}{-}{-})$, the speed of light defined by the metric $g_{ab}$ is unity, and the \\ae ther is taken to be dimensionless.) The term $R_{ab}u^au^b$ can be expressed as the difference of the $c_3$ and $c_2$ terms, up to a total derivative, so is not independent. In computations the unit timelike constraint on the \\ae ther is usually imposed by adding a Lagrange multiplier term $\\l(g_{ab}u^au^b-1)$ to the action. The covariant derivative operator $\\nabla_a$ involves derivatives of the metric through the connection components, and the unit vector is nowhere vanishing, hence the terms quadratic in $\\nabla_a u^m$ are quadratic in derivatives of both \\ae ther and metric perturbations, so the metric and \\ae ther modes are coupled. The field equations are written out in detail in many of the references, beginning with Ref.~\\cite{Eling:2003rd}. It is noteworthy that the \\ae ther stress tensor includes second derivative terms arising from the variation of the metric in the connections. In all phenomenology work to date, it has been assumed that the \\ae ther is aligned at large scales with the rest frame of the microwave background radiation.   The restriction to no more than two derivatives is motivated by the standard precepts of effective field theory~\\cite{Burgess:2003jk}: higher derivatives would be suppressed by one power of a small length scale for each extra derivative. The natural size of the coupling constants $c_i$ depends on unknown physics at high energies. If the Planck scale is the only relevant scale then the $c_i$ are naturally all of order unity. If on the other hand there is an additional scale or scales characterizing the Lorentz violating physics, then the $c_i$ might be naturally smaller and could differ from each other in order of magnitude. Note that although the $c_4$ term in the action is quartic in the \\ae ther field, it makes a quadratic contribution to the kinetic terms for metric and \\ae ther perturbations when expanded around a flat background with a constant \\ae ther of unit norm.  Einstein-\\ae ther theory is similar to the vector-tensor gravity theories studied long ago 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 at being 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}. Related vector-tensor gravity theories are that of Kostelecky and Samuel~\\cite{Kostelecky:1989jw} which corresponds to the Maxwell-like special case of \\ae -theory ($c_3=-c_1$, $c_2=c_4=0$), both with a fixed norm and with a symmetry breaking potential for the vector, of Gripaios~\\cite{Gripaios:2004ms} which has all two-derivative terms and a symmetry breaking potential for the vector, and the generalized Einstein-\\ae ther theories of Zlosnik, Ferreira and Starkman~\\cite{Zlosnik:2006zu} and of Zhao~\\cite{Zhao:2007ce}, in which the \\ae ther kinetic terms of (\\ref{action}) are replaced by functions thereof. Kanno and Soda~\\cite{Kanno:2006ty} introduced dependence of the coupling parameters on a scalar field that has its own dynamics, and studied inflationary cosmology for one such model. Other theories involving scalar fields in addition to a vector field with a Maxwell-like action for the vector are Bekenstein's TeVeS~\\cite{Bekenstein:2004ne} and Moffat's STVG~\\cite{Moffat:2005si}.   \\subsection{Matter couplings} Lorentz-violating (LV) effects in the matter sector are produced by couplings of the matter to the \\ae ther. (More complicated Lorentz symmetry breaking patterns require LV extensions of general relativity with other symmetry breaking fields~\\cite{Bluhm:2004ep}.) Such effects have been highly constrained by observations~\\cite{Mattingly:2005re}, so if they exist they are very weak. In this review I will therefore assume that matter couples universally to the metric $g_{ab}$. This assumption is motivated by phenomenology, but goes against the precepts of effective field theory since the Lorentz violation in the gravitational sector would presumably induce Lorentz violating terms in the matter action via loop effects. Supersymmetry could conceivably provide a natural suppression of these terms, as discussed briefly in section~\\ref{Frontiers}.   \\subsection{Metric redefinitions}  When investigating aspects of the theory that do not involve matter, it can be helpful to exploit the metric redefinition $g'_{ab} = g_{ab}+(\\z-1) u_a u_b$, which ``stretches\" the metric tensor in the \\ae ther direction by a positive factor $\\z$. (A negative factor would return a Euclidean signature metric.) The action (\\ref{action}) for ($g'_{ab}$, $u'^a = u^a/\\sqrt{\\z}$) takes the same form as that for ($g_{ab}$, $u^{a}$), with new coefficients $c'_i$. The relation between the $c'_i$ and $c_i$ was worked out in~\\cite{Foster:2005ec}, and is conveniently given in terms of certain combinations with simple scaling behavior: \\bea c'_{14} &=& c_{14}\\nonumber\\\\ c'_{123} &=& \\z c_{123}\\nonumber\\\\ c'_{13}-1 &=& \\z (c_{13}-1)\\nonumber\\\\ c'_1-c'_3-1 &=& \\z^{-1}(c_1-c_3-1). \\label{redef} \\eea Note that in the absence of matter couplings a one-parameter set of Einstein-\\ae ther theories is actually pure vacuum GR in disguise~\\cite{BarberoG.:2003qm,Foster:2005ec}: the parameters $c_i'$ all vanish if $c_{14}=c_{123}=2c_1-c_1^2+c_3^2=0$, and $\\z>0$ provided $c_{13}<1$. ",
        "conclusions": "Einstein-\\ae ther theory provides a theoretical laboratory in which gravitational effects of possible Lorentz violation can be meaningfully studied without abandoning the generally covariant framework of general relativity. It is striking that \\ae -theory can match observations to the degree that it can. While of course starting with four free parameters $c_{1,2,3,4}$ makes this easier, the freedom to choose these parameters might have already been used up in setting the PPN parameters to agree with those of GR. Instead, only two parameters need be fixed at that stage, leaving $c_1$ and $c_3$ free. Then, despite possessing three distinct types of linearized wave modes, all positive energy, stability and vacuum \\v{C}erenkov requirements on the modes are met within a large common region of the $(c_1,c_3)$ space. Within this same parameter space, the dynamics of the cosmological scale factor and perturbations differ little from GR, and non-rotating neutron star and black hole solutions are quite close to those of GR (but may be distinguishable with future observations).  Additional constraints arise from the observed radiation damping rate in binaries. For systems with weak self-fields, a constraint of order $10^{-3}$ would be imposed on one combination of the two parameters $c_{1}$ and $c_3$. Current observations involve pulsars which possess strong self-fields. Presuming that the sensitivity parameters for neutron stars turn out to have the expected magnitude, such binaries will exhibit effects constraining both the parameters to be less than around $10^{-2}$, in addition to the stronger radiation damping constraint on one combination. It may turn out that observations of multiple systems with different combinations of sensitivities will constrain both parameters separately at the $10^{-3}$ level."
    },
    "0801/0801.3542_arXiv.txt": {
        "abstract": "{Using radio polarimetry we study the connection between the transport of cosmic rays (CR's), the three-dimensional magnetic field structure, and features of other ISM phases in the halo of NGC\\,253. We present a new sensitive radio continuum map of NGC\\,253 obtained from combined VLA and Effelsberg observations at $\\lambda6.2\\,{\\rm cm}$. We find a prominent radio halo with a scaleheight of the thick radio disk of 1.7\\,kpc. The linear dependence between the \\emph{local} scaleheight of the vertical continuum emission and the cosmic ray electron (CRE) lifetime requires a vertical CR bulk speed of $270\\,{\\rm km\\,s^{-1}}$.\\\\ The magnetic field structure of NGC\\,253 resembles an ``X''-shaped configuration where the orientation of the large-scale magnetic field is plane-parallel only in the inner regions of the disk and at small distances from the galactic midplane. At larger galactocentric radii and further away from the midplane the vertical component becomes important. This is most clearly visible at the location of the ``radio spur'' southeast of the nucleus, where the magnetic field orientation is almost vertical. We made a simple model for the dominant toroidal $(r,\\phi)$ magnetic field component using a spiral magnetic field with prescribed inclination and pitch angle. The residual poloidal $(r,\\phi,z)$ magnetic field component which was revealed by subtracting the model from the observations shows a distinct ``X''-shaped magnetic field orientation centered on the nucleus. The orientation angle of the poloidal magnetic field is consistent with a magnetic field transport described by the superposition of the vertical CR bulk speed and the rotation velocity.\\\\ Hence, we propose a \\emph{disk wind} which transports cosmic rays, magnetic field, and (partially) ionized gas from the disk into the halo.}  \\FullConference{From planets to dark energy: the modern radio universe\\\\ October 1-5 2007\\\\ University of Manchester, Manchester, UK}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.2152_arXiv.txt": {
        "abstract": "Polarimetric observations are affected by leakage of unpolarized light into the polarization channels, in a way that varies with the angular position of the source relative to the optical axis.  The off-axis part of the leakage is often corrected by subtracting from each polarization image the product of the unpolarized map and a leakage map, but it is seldom realized that heterogeneities in the array shift the loci of the leaked radiation in a baseline-dependent fashion.  We present here a method to measure and remove the wide-field polarization leakage of a heterogeneous array.  The process also maps the complex voltage patterns of each antenna, which can be used to correct all Stokes parameters for imaging errors due to the primary beams. ",
        "introduction": "\\label{sec:intro}  The hardware typically used in radio telescopes has the great benefit of observing the Stokes $Q$, $U$, and $V$ parameters simultaneously with Stokes $I$, but always allows some mixing between the polarization channels, as in Figs.~\\ref{fig:intdiag}~and~\\ref{fig:nolc}.  This ``leakage'' is particularly troublesome when it goes from $I$ into $Q$, $U$, or $V$, since the polarized signals are usually a small fraction of the total intensity, and therefore easily swamped by similarly strong leakages from $I$.  \\begin{figure} {\\centering \\plotone{intdiag} \\caption{Conceptual diagram of polarization leakage in an interferometer. Each antenna measures two nominally orthogonal polarizations $p$ and $q$, but they are partially mixed before entering the correlator $C$. } \\label{fig:intdiag} \\begin{tabular}{r@{: }p{0.85\\linewidth}} $\\boldsymbol{b}$ & Baseline (separation) between antennas.\\\\ $\\boldsymbol{O\\!A}$ & The optical axis (i.e.\\ pointing direction).\\\\ $(l, m)$ & Longitudinal and latitudinal offsets perpendicular to $\\boldsymbol{OA}$.\\\\ $\\boldsymbol{n}$ & An arbitrary offset in $(l, m)$.\\\\ \\jones{\\Gamma}{A} & Voltage pattern of antenna $A$, factored to exclude polarization leakage.\\\\ HPBW & Half Power Beam Width.\\\\ \\jones{D}{B} & $\\boldsymbol{n}$ independent factor of the polarization leakage of antenna $B$.\\\\ \\jones{\\Psi}{} & $\\boldsymbol{n}$ dependent factor of the polarization leakage.\\\\ $V_{\\mathrm{obs},AB}$ & Observed vector of visibilities in each polarization. \\end{tabular}} Although the diagram places \\jones{\\Psi}{} and \\jones{\\Gamma}{} above the receiver and \\jones{D}{} below, each includes effects from the feed, reflector surface, and receiver support struts. \\end{figure}  A radio interferometer uses two or more antennas to measure the amplitudes and phases of the electric field impinging on their receivers.  The measurements are stored as ``visibilities'', which are the correlations of the receiver voltages.  Given some conditions which this article will assume to have been met, the visibilities sample the Fourier transform of the sky multiplied by the primary beam (directional sensitivity function) of the antennas \\citep{bib:siss_clark, bib:siss_thompson}.  Mathematically, the effect of a pair of antennas $A$ and $B$ on the visibilities they observe, $V_{\\mathrm{obs},AB}$, is conveniently expressed using the Hamaker-Bregman-Sault \\citep{bib:hbs1} formalism, where the four polarizations are combined into a column vector.  The true visibilities are multiplied on the left by a set of Jones matrices, each one the outer product of Jones matrices for antennas $A$ and $B$, i.e. \\begin{equation*}  % \\jones{D}{AB} = \\jones{D}{A} \\otimes \\jones{D}{B}^* \\end{equation*} represents the on-axis mixing between the nominally orthogonal polarization channels, often called the ``$D$ terms''.  The outer product of two matrices $\\mathbf{\\mathsf{M}}$ and $\\mathbf{\\mathsf{N}}$, $\\mathbf{\\mathsf{M}} \\otimes \\mathbf{\\mathsf{N}}$, is formed by multiplying each entry of $\\mathbf{\\mathsf{M}}$ with all of $\\mathbf{\\mathsf{N}}$, and is used, along with a complex conjugation of the second factor, to bring together the elements from each antenna in a correlation.  \\citet{bib:hbs1} explain the algebraic properties, including coordinate transformations, of Jones matrices and the outer product in more detail.  As in \\citet{bib:bhatnagar_evla100}, direction dependent effects can also be included, but they must go \\emph{inside} the Fourier integral: \\ifPP \\begin{gather} V_{AB}^\\textrm{obs} = \\jones{D}{AB}\\!\\int\\! \\jones{\\Psi}{\\! AB}(\\boldsymbol{n}) \\jones{\\Gamma}{\\! AB}(\\boldsymbol{n}) \\mathbf{\\mathsf{S}}I_S(\\boldsymbol{n} + \\boldsymbol{n}_c)\\, \\mathe^{\\mi \\boldsymbol{n} \\cdot \\boldsymbol{b}_{AB}}\\dif n \\label{eq:js} \\end{gather} \\else \\begin{align} V_{AB}^\\textrm{obs} = \\jones{D}{AB}\\!\\int\\! & \\jones{\\Psi}{\\! AB}(\\boldsymbol{n}) \\jones{\\Gamma}{\\! AB}(\\boldsymbol{n})\\cdot\\nonumber\\\\ & \\mathbf{\\mathsf{S}}I_S(\\boldsymbol{n} + \\boldsymbol{n}_c)\\, \\mathe^{\\mi \\boldsymbol{n} \\cdot \\boldsymbol{b}_{AB}}\\dif n \\label{eq:js} \\end{align} \\fi where $\\boldsymbol{n}$ is a direction on the sky relative to the ``phase tracking center'', $\\boldsymbol{n}_c$.  $\\boldsymbol{n}_c$ is set electronically, but usually it is chosen to coincide with the pointing direction of the antennas.  $\\mathbf{\\mathsf{S}}$ is the Stokes matrix, which transforms the sky's Stokes parameters, $I_S = (I, Q, U, V)$, into the observational polarization basis, typically correlations of either circular or linear polarizations.  \\jones{\\Psi}{AB} and \\jones{\\Gamma}{AB} are respectively the wide-field leakage pattern and primary beam for the correlation of antennas $A$ and $B$.  They are sometimes multiplied together to form a single Jones matrix which is a generalization of the primary beam, but the magnitudes of the effects are more easily assessed if they are kept separate.  With the separation, \\jones{\\Gamma}{AB} is diagonal since it does not mix polarizations in the observational basis, and the diagonal elements of \\jones{\\Psi}{AB} are all one.  The on-axis portion of the leakage, \\jones{D}{AB}, is dealt with by standard polarimetric calibration techniques, but the leakage varies with direction, growing worse toward the edges of the primary beam, as in Fig.~\\ref{fig:nolc}.  This paper is concerned with the wide-field polarization leakage, \\jones{\\Psi}{AB}, and will assume that $V_{\\mathrm{obs},AB}$ has already been corrected by multiplication with $\\jones{D}{AB}^{-1}$.  \\begin{figure*} \\begin{minipage}[t]{1.0\\columnwidth} \\plotone{urawpm5.png} \\caption{An example of leakage from Stokes $I$ into Stokes $U$ in a CGPS field containing the supernova remnant IC443.  The thin circle is the 75\\arcmin\\ radius (24\\% power) cutoff of the usable part of the beam in polarization.  The image has not been corrected for the sensitivity dropoff of the primary beam, and only includes ST data.  Note that it has been CLEANed, so the arcs are mostly leakage.  The grayscale goes from -5 (black) to 5 (white) mJy/beam. \\label{fig:nolc}} \\end{minipage}\\hfill\\begin{minipage}[t]{1.0\\columnwidth} \\plotone{uoipcpm5.png} \\addtocounter{figure}{1}\\caption{Fig.~\\ref{fig:nolc} (same grayscale) after image-based leakage correction.  The ``on-source'' correction of unresolved sources is accurate to 1\\% of $I$, close to the theoretical precision of the measured leakage map, but the arcs around strong leakage remain unaffected.  Leakage amplitude differences between antennas produce rings, and phase differences produce asymmetric arcs. \\label{fig:oldlc}} \\end{minipage}  \\begin{minipage}[t]{1.0\\columnwidth} \\plotone{im12p350.png} \\addtocounter{figure}{-2}\\caption{The sources of the $I$ radiation that leaked into Fig.~\\ref{fig:nolc}.  The grayscale goes from -12 (white) to 350 (black) mJy/beam.} \\end{minipage}\\hfill\\begin{minipage}[t]{1.0\\columnwidth} \\plotone{unipcpm5.png} \\addtocounter{figure}{1}\\caption{Fig.~\\ref{fig:nolc} (same grayscale) after correcting leakage using measured patterns for each antenna.  Leakage measurements were made only inside the circle, but they have been extrapolated to the edge of the image, which works well for clearing up the arcs of sources slightly outside the limit.  The remaining arcs are primarily due to differences between the primary voltage patterns of the antennas.  \\label{fig:newlc}} \\end{minipage} \\end{figure*}  If all of the antennas in an array are identical, the effects of the primary beam and leakage patterns can be removed in the image plane.  At the Dominion Radio Astrophysical Observatory (DRAO) we previously corrected the wide-field contamination or polarization by multiplying the Stokes $I$ image with ``leakage maps'', and subtracting the results from the measured $Q$ and $U$ images, as in Fig.~\\ref{fig:oldlc}.  The leakage maps were measured by observing the apparent $Q/I$ and $U/I$ of an intrinsically unpolarized source in a grid of offsets from the primary beam center \\citep{bib:peracaula1999}. This correction is performed completely in the image plane, so we call it the ``image-based'' leakage removal method.  It was immediately applicable for DRAO's Synthesis Telescope (ST, \\cite{bib:ST2000}) since its antennas are equatorially mounted and thus its leakage patterns never rotate relative to the sky.  The leakage patterns of an altitude-azimuth mounted telescope such as the Very Large Array (VLA) rotate relative to the sky over the course of an observation, but the image-based method can still be applied if the data are first broken up into a series of snapshots \\citep{bib:cotton_aips86}.  Unfortunately, there are differences, known or unknown, between the antennas of any real interferometer.  The array may be a combination of antennas from originally separate telescopes, such as the Combined Array for Research in Millimeter-Wave Astronomy (CARMA, \\citet{bib:carma2006}), and most very long baseline interferometers.  It could also be in a transition period where only some antennas have been modified, like the partially Enhanced Very Large Array, and/or have serious surface errors as at (sub)mm wavelengths.  The ST is an example of an array where the antennas are similar to each other, but with known differences between them.  The two outermost antennas have 9.14\\,m diameters with four metal struts supporting their receivers, while the other five are 8.53\\,m in diameter with three struts, made of either metal or fiberglass.  The differences in antenna diameter obviously create differences in the half-power beamwidths (HPBWs), which at 1420\\,MHz are 101.8\\arcmin\\ for the two outer antennas and 108.8\\arcmin\\ for the rest.  The variation in the number and composition of the struts affects the scattering of incoming light, which is an important component of polarization leakage (most of the rest comes from the feeds).  In polarization images the differences between antennas are seen as mismatches between the standard point spread function (PSF, or ``dirty beam'') and the PSF of the leakage.  When there are phase differences between the leakages of the antennas, the effective PSF of the leakage is asymmetric \\citep{bib:siss_ekers} and offset from the peak of the unpolarized emission.  The effective PSF of the leakage also varies across the field, meaning that subtracting a multiplication of the Stokes $I$ map with a leakage map cannot fully correct the polarization leakage of a heterogeneous array.  Additionally, the response of each antenna in an array, both in leaked and true radiation, depends on the scale of the source(s).  Resolved features have less power at high spatial frequencies, so antennas that only participate in long baselines will contribute little leakage to them.  Unresolved features have no such attenuation with baseline length, and elicit an equally weighted mix of leakage from all antennas.  Usually leakage maps are measured using a bright unresolved object, so in the case of a heterogeneous array their corrections are only accurate for unresolved sources.\\footnote{Even then, only if the images are made with the same baseline weighting as used for the leakage map.}  Fig.~\\ref{fig:oldlc} exhibits both of these problems, as can be seen in comparison with Fig.~\\ref{fig:newlc}, the same image corrected with the method described in Section~\\ref{sec:corr}.  The main change for the unresolved sources is the presence or lack of surrounding arcs, but the supernova remnant also shows a strong difference in the on-source residual leakage.  The work described here aims to improve polarization imaging from the DRAO ST. The telescope is engaged in an extensive survey of the major constituents of the Interstellar Medium, the Canadian Galactic Plane Survey (CGPS, \\cite{bib:cgps2003}) which includes imaging in Stokes parameters $Q$ and $U$ at 1420 MHz along the plane of the Milky Way.  We recently measured the real and imaginary parts of the leakage patterns for each antenna of the ST (similar to a hologram measurement, but with finer spacing over a smaller area) and have started using them to correct its $Q$ and $U$ observations, as seen in Fig.~\\ref{fig:newlc}. ",
        "conclusions": "\\label{sec:dis}  The measured leakage patterns, Figs.~\\ref{fig:qlvpre} to \\ref{fig:ulvpim}, show that although there is some overall consistency in the patterns, their details are unpredictable, both from antenna to antenna and from band to band in frequency.  Most noticeably, the antennas with quadrupod receiver supports, 1 and 7, are structurally nearly identical, but their leakage patterns do not show any more similarity to each other than they do to those of the tripod antennas.  Likely this is because most of the leakage comes not from the struts, but from the feeds.  The feeds are nominally identical, and their individual flaws are neither easily apparent to visual inspection nor tied to the type of antenna they are mounted on.  This suggests that wide-field polarimetry with even nominally homogeneous arrays requires measuring the leakage patterns of each antenna, if the needed fidelity warrants it.  Variation of the leakage patterns from band to band is prominent in the real parts of the leakage patterns. This rapid change with frequency seems surprising at first glance: one might expect properties of a waveguide feed to vary quite slowly with frequency, and hardly at all across a band that is only 2\\% of the center frequency. The cause appears to be the probes used to feed the reflector at 408 MHz; they are housed within the 1420 MHz feed \\citep{bib:Veidt85}. Computed simulations (B.G. Veidt, private communication) indicate that these probes cause some fine structure in the performance at 1420 MHz.  The primary voltage patterns, Figs.~\\ref{fig:pvpre} and \\ref{fig:pvpim}, reassuringly exhibit only the expected dependence on wavelength; namely their angular scales are proportional to the observing wavelength.  Their apparent tight link to antenna structure suggests that primary voltage pattern errors are more amenable to correction by adjusting the antennas, as is often done using holograms.  Once the primary voltage patterns are known, their effect can also be reduced post-observation, even for an inhomogeneous array \\citep{bib:bhatnagar_evla100}.  Currently such errors are attacked with direction dependent self-calibration (modcal, \\citep{bib:modcal}, also called peeling), which is vulnerable to confusing true features on the sky with unwanted artifacts.  Measuring the antenna patterns with a bright unresolved calibration source instead of through self-calibration with a potentially complicated fainter science target removes that vulnerability.  \\begin{figure*} \\begin{minipage}[t]{1.0\\columnwidth} \\ifOnline \\includegraphics[width=\\lmswidth,clip]{pvpre_c} \\else \\includegraphics[width=\\lmswidth,clip]{pvpre_g} \\fi \\caption{Real parts of the primary voltage patterns of antennas 1 (top) to 7 (bottom) for bands A (left) to D (right).  The grayscale goes from 0.4 (white) to 1.0 (black), and the contours go from 0.4 to 0.9 in steps of 0.1.} \\label{fig:pvpre} \\end{minipage}\\hfill\\begin{minipage}[t]{1.0\\columnwidth} \\ifOnline \\includegraphics[width=\\lmswidth,clip]{pvpim_c} \\else \\includegraphics[width=\\lmswidth,clip]{pvpim_g} \\fi \\caption{Imaginary parts of the primary voltage patterns of antennas 1 (top) to 7 (bottom) for bands A (left) to D (right).  The \\rampname\\ goes from -0.1 (\\negcolor) to 0.1 (\\poscolor).  The symmetry of antenna 6's patterns suggests that it is out of focus.} \\label{fig:pvpim} \\end{minipage} \\end{figure*}  Although we have only tested heterogeneous array leakage correction with equatorially mounted antennas, in principle it would be even easier to adapt it to antennas on altitude-azimuth mounts than the image-based leakage map method. Since the leakage voltage pattern method already deals with visibilities on an individual basis, the only modification needed would be make $x_j$ and $y_j$ in Equation~\\ref{eq:visleakage} functions of time to account for the rotation of the antennas about the optical axis relative to the sky as the Earth turns.  An implicit, but difficult to avoid, assumption in correcting for the effect of beam patterns is that the patterns do not change with time or observing elevation.  The prospect of spending observing time on frequent antenna pattern remeasurements, possibly for a set of elevations and frequencies, is unappealing, so there is considerable pressure to engineer antennas that are stable enough for occasional measurements to capture most of the effects.  The ST antennas were not expected to change significantly with time or observing direction, but we confirmed their behavior by comparing recent leakage measurements to the measurements made by Peracaula of the ST's overall leakage amplitude maps at 21 cm wavelength.  There was little change over the intervening 10 years, despite some surface modifications to a few of the antennas.  Stability is expected to be a more serious problem for larger (as measured in wavelengths) dishes, especially if standing waves create a noticeable resonance effect in the leakage at certain observing frequencies. Interpolation, or theoretical modeling, may be useful for extending the applicability of measured maps to additional elevations and/or frequencies. Alternatively, if an extremely accurate correction is only needed for one bright source within the field of an observation, the antenna patterns could be measured at that spot immediately before and after the science observation, as opposed to mapping the entire main lobe of the antenna patterns."
    },
    "0801/0801.2963_arXiv.txt": {
        "abstract": "The 11th magnitude star {\\LSIV} has been classified previously as an OB star in the Luminous Stars survey, or alternatively as a hot subdwarf.  It is actually a binary star.  We present spectroscopy, spectroscopic orbital elements, and time series photometry, from observations made at the Kitt Peak National Observatory 2.1m, Steward Observatory 2.3m, MDM Observatory 1.3m and 2.4m, Hobby-Eberly 9.2m, and Michigan State University 0.6m telescopes.  The star exhibits emission of varying strength in the cores of H and {\\HeI} absorption lines.  Emission is also present at 4686\\,\\AA\\ ({\\HeII}) and near 4640/4650\\,\\AA\\ ({\\NIII}/{\\CIII}).  Time-series spectroscopy collected from 2005 July to 2007 June shows coherent, periodic radial velocity variations of the {\\Ha} line, which we interpret as orbital motion with a period of 0.1952894(10)\\,days.  High-resolution spectra show that there are two emission components, one broad and one narrow, moving in antiphase, as might arise from an accretion disk and the irradiated face of the mass donor star. Less coherent, low-amplitude photometric variability is also present on a timescale similar to the orbital period.  Diffuse interstellar bands indicate considerable reddening, which however is consistent with a distance of $\\sim$100--200 pc.  The star is the likely counterpart of a weak {\\em ROSAT} X-ray source, whose properties are consistent with accretion in a cataclysmic variable (CV) binary system.  We classify {\\LSIV} as a new member of the UX UMa subclass of CV stars. ",
        "introduction": "\\label{sec:Intro}  The star {\\LSIV} has $V = 11.475$, $(\\bv) = +0.190$ {\\citep{RN00}}, and was first classified ``OB'' in the {\\em Luminous Stars in the Northern Milky Way} catalog {\\citep{LSIV}}.  Later it was reclassified ``sdB'' by {\\citet{KB92}} based on Reticon spectrograms and Str\\\"{o}mgren photometry.  Additional photometry of this object is available in the literature {\\citep{RCL,Reed,RN00}}, but no additional spectroscopic observations have been reported. Astrometric and photometric data for this star, collected from the literature, are reported in Table~\\ref{table:data}.  We obtained spectra of {\\LSIV} as part of our program on composite spectrum hot subdwarf stars {\\citep{SW,SW04,Thesis,SW06}}, owing to its unusually red $V-K_s$ and $J-K_s$ colors as observed in the {\\em Two Micron All-Sky Survey} {\\citep[2MASS;][]{Skrutskie06}}.  We found that {\\LSIV} exhibits broad absorption lines with core emission, which is variable in both strength and profile shape.  We obtained time-series photometry and additional spectroscopy in order to correctly classify this star. ",
        "conclusions": "\\label{sec:Discussion}  \\subsection{Novalike CV Interpretation}\\label{sec:novalike}  We have already suggested (\\S\\ref{sec:Steward}) that {\\LSIV} could have been misclassified as a hot subdwarf (sdB) star, given only photometry and spectroscopy confined to wavelengths blueward of {\\Hb}. Indeed, the dwarf nova FO Per (aliases: RL 92 and RWT 92) was classified by \\citet{Chromey79} as an sdB star, when, unbeknownst to him, it must have been near maximum light in its outburst cycle.  Two Steward Observatory 2.3m spectra of FO Per, also near maximum light, obtained in 2004 Dec and 2006 Dec with the identical setup as described in \\S\\ref{sec:Steward}, show a spectrum that is almost indistinguishable from a somewhat reddened sdB star, except for emission cores in H$\\alpha$, H$\\beta$, and, very faintly, in the next few higher order Balmer lines, i.e., at almost exactly the same level as seen in our 2007 May spectrum of {\\LSIV}.  The SED of {\\LSIV} is considerably redder than those of known single sdB stars. Dereddening by the equivalent of $E(\\bv) \\sim 0.6$ or more would be required to bring the Steward SED at optical wavelengths into agreement with sdBs that have $T_{\\rm eff}$ in the range 24,000 --- 35,000 K.  Likewise, the $E(\\bv)$ needed to bring the 2MASS infrared colors into the range of single sdB stars is 0.6 mag or more.  On the other hand, {\\LSIV} shares many properties with novalike variables of the UX UMa subclass.  These systems show persistent broad Balmer absorption-line spectra. Ironically, these lines were once thought to indicate pressure-broadening in the atmosphere of a hot subdwarf or white dwarf --- e.g., \\citet{WH54} --- and only later were reinterpreted as arising from an optically thick accretion disk, where they are (at least in part) kinematically broadened.  UX UMa stars also exhibit a wide range in strengths of the emission line components.  Emission from combinations of {\\HI}, {\\HeI}, {\\HeII}, {\\NIII}, and {\\CIII} have all been observed at a variety of strengths in various UX UMa stars, or variable over time in a single star of the class.  Emission is seen in all of these lines in {\\LSIV}, and the features are consistent with those observed in the UX UMa stars RW Sex and IX Vel {\\citep[][respectively]{BSS92,BeuermannThomas90}}.  For UX UMa stars, the Balmer decrement in the emission components is steeper than in the absorption lines.  This can result in a situation where {\\Ha} is purely in emission, while higher members show progressively stronger (less filled-in) absorption troughs.  This effect is seen in {\\LSIV} (Fig.~\\ref{fig:betsyflux2002}).  Narrow Balmer emission has been found in many UX UMa systems, superimposed on the broader emission from the disk. This narrow component is attributed to irradiation of the atmosphere of the secondary star, with reprocessing of the light to Balmer emission.  As discussed in \\S\\ref{sec:HRS}, the asymmetry and variations in the emission line profile shape (Figs.~{\\ref{fig:HRS-keyphases}} {\\&} {\\ref{fig:trailedspec}}) suggest the presence of such a component, moving in anti-phase to the broad component, as required for this interpretation.  The orbital period of $\\sim0.195$\\,d ($\\sim4.7$ hours) determined for {\\LSIV} is comparable to other UX UMa systems.  \\subsection{Distance and Reddening}\\label{sec:reddening}  {\\LSIV} lies at galactic coordinates $(l,b) = (11\\fdg0, +20\\fdg8)$. The \\citet{Schlegel} reddening maps show that the total line-of-sight extinction varies on small angular scales in this direction.  We attempted to infer the $E(\\bv)$ color excess directly from the strength of the DIBs, using the linear relations, $W = a_0 E(\\bv)$, established by {\\citet{HerbigDIB}}.  The $\\lambda$5780 DIB has an EW of $0.32 \\pm 0.03$\\,\\AA, suggesting $E(\\bv) \\approx 0.50$ mag with an uncertainty of at least 10\\%, if we are measuring the narrow component of this feature only, or $E(\\bv) \\approx 0.24$ mag if we are measuring the sum of the narrow and broad DIBs (the latter is more likely). The $\\lambda$4430 DIB has an EW of $0.26 \\pm 0.04$\\,\\AA, corresponding to $E(\\bv) \\approx 0.12$ with at least 15\\% uncertainty.  Some of the uncertainty arises, depending on whether ``all-sky'' or ``Sco-Oph'' values of $a_0$ are adopted.  Both DIBs imply reddening less than the full \\citet{Schlegel} column: $E(\\bv) = +0.47$.  The heliocentric RV of the interstellar \\ion{Na}{1} D lines is $-14.5~{\\rm km~s^{-1}}$ (average of D$_1$ and D$_2$).  Most of this can be explained by the reflex of the Sun's space motion with respect to the Local Standard of Rest, so kinematic information about the distance of the absorbing material between the Sun and {\\LSIV} is unavailable.  If {\\LSIV} is assumed to be a UX UMa variable star, how far away is it, and is this consistent with the measured reddening?  We adopt an absolute magnitude $M_V \\approx +5$ and a representative apparent magnitude $V = 11.50$ for {\\LSIV}. We use the standard relation $A_V = 3.1 E(\\bv)$ to compute distance $d$ as a function of assumed $E(\\bv)$. Values between 100 and 200 pc result for $E(\\bv)$ in the range 0.0 to 0.5 mag, smaller distances corresponding to larger reddenings. \\citet{Ak07} have put forward a calibration of $M_J$ for CVs, using as inputs the orbital period and the extinction-corrected 2MASS $J$ magnitude and $J-H$, $H-K_{s}$ colors.  Again varying the assumed $E(\\bv)$, we find $d$ in the range 140--180 pc by this method (minimum error at fixed reddening is 40\\%).  Comparing the 2MASS colors of {\\LSIV} and IX Vel directly, we find a reasonable match if $E(\\bv) \\approx 0.10$ mag.  We conclude that, if $E(\\bv)$ in the range 0.10--0.20 mag can be accommodated within distances of 100--200 pc, the colors and magnitudes of {\\LSIV} are consistent with those of a UX UMa variable.  Such reddening values within such distances are plausible along the line of sight to {\\LSIV}, according to the maps of, e.g., {\\citet{PerryJohnston82}}.  \\subsection{{\\em ROSAT} X-ray Observations}  A weak X-ray source detected in the {\\em ROSAT} All-Sky Survey (RASS) appears to be coincident with {\\LSIV} {\\citep[$1$RXS~J$165630.2$$-083442$;][]{Voges}}.  The count rate in the position sensitive proportional counter (PSPC) was 0.11$\\pm$0.02~s$^{-1}$, yielding an estimated 0.1--2.4\\,keV flux of $1.5 \\times 10^{-12}~{\\rm erg~cm^{-2}~s^{-1}}$.  Taking {\\LSIV} to be the optical counterpart, the X-ray/optical flux ratio is given by $\\log f_X/f_{\\rm opt} \\approx -2.0$.  This source was also observed serendipitously by {\\em ROSAT} {\\citep[$1$WGA~J$1656.4$$-0834$;][]{White}} during a pointed observation of another target.  In this $\\sim$8000\\, s exposure, 150 net counts were collected, and the count rate was 0.02$\\pm$0.002~s$^{-1}$ (5 times lower than in the RASS), corresponding to $f_X \\approx 3.0\\times 10^{-13}~{\\rm erg~cm^{-2}~s^{-1}}$ and $\\log f_X/f_{\\rm opt} \\approx -2.6$.  A best-fitting blackbody spectrum has $kT \\approx 0.2$~keV with an intervening hydrogen column density of $1.3\\times 10^{21}$~cm$^{-2}$, corresponding to $E(\\bv) \\approx 0.2$.  The unabsorbed flux is $F_{\\mathrm{0.5-2~keV}} \\approx 2.7\\times10^{-13}$~erg\\,s$^{-1}$\\,cm$^{-2}$.  Fits using a bremsstrahlung or thermal plasma emission model would result in a similar $kT$ value, given the highly absorbed spectrum and the PSPC's limited energy range and resolution. The inferred total X-ray flux is quite uncertain.  X-ray emission is not associated with sdB stars.  On the other hand, CVs of all subclasses show X-ray emission.  {\\citet{VBRP97}} discuss the CVs detected in the RASS.  A typical luminosity in the 0.5--2.5~keV band for a novalike variable is $L_X \\sim 10^{31}~{\\rm erg~s^{-1}}$. The required distance for {\\LSIV} to have this $L_X$, given its modeled flux, is $d \\sim 500$ pc.  Using a 2 keV bremsstrahlung spectrum to model the ROSAT counts, {\\citet{VBRP97}} find $\\log F_X/F_{\\rm UV+opt} < -2.5$ to be typical for novalikes. Given the uncertainties in the input spectrum and absorption column, we conclude that {\\LSIV} has X-ray properties consistent with membership in the UX UMa subclass of CVs.    Given the many similarities between our observations and other UX UMa novalike variables, we propose that {\\LSIV} should be classified with the UX UMa stars.  Evidence supporting this classification includes the presence of emission reversals in the Balmer absorption series; other emission lines including \\ion{He}{1}, \\ion{He}{2}, and \\ion{N}{3}/\\ion{C}{3}; the variability of the emission line strengths; the intrinsically blue continuum; and the mild photometric variability.  The key evidence in favor of a CV interpretation lies in the periodically modulated Doppler shifts and profiles of the emission lines, with a period of 4.7 hr, consistent with a low-mass quasi-main sequence star that fills its Roche lobe and transfers mass to a luminous accretion disk around a white dwarf.  The inner hemisphere of the donor star may be irradiated by the accretion disk.  Reddening and distance estimates are consistent with the CV interpretation, from the standpoint of the expected luminosity of the system and the distribution of interstellar dust in the direction of {\\LSIV}. Finally, weak X-rays are apparently associated with the system, and the X-ray luminosity and X-ray/optical flux ratio are consistent with observations of other novalike CVs.  We refrain from deriving the dynamical mass ratio of the binary or individual masses of the stars in {\\LSIV} from any of the various RV semiamplitudes that we have presented above, mainly because the HRS spectra on which such estimates would rely were obtained on widely separated dates, and the EW of the emission lines varied significantly among the various epochs.  It is thus possible that one or both sites of the emission varied as well. Work aimed at determining the mass ratio is best undertaken with intensive time-series spectroscopy that can provide full orbital phase coverage in one or two nights, to minimize such variations.  Likewise, we cannot offer a definitive interpretation of the photometric variability, given our present limited data.  {\\LSIV} thus remains as an attractive and interesting target for further study.  Despite the apparent line-of-sight extinction to the source, {\\LSIV} is one of the brighter UX UMa stars so far discovered. {\\citet{Downes}} list only three ``UX Uma'' systems (of two dozen catalogued) that are brighter: IX Vel, V3885 Sgr, and RW Sex.  Only seven CVs of any ``novalike'' sub-type (89 systems total) are brighter than {\\LSIV} at their maximum light.  Bright, relatively nearby systems are not only appropriate for follow-up studies, but contribute disproportionately to studies of the space density and kinematics of the various subclasses of CVs. Even though hundreds of thousands of stars that are much fainter have been observed spectroscopically by, e.g., the Sloan Digital Sky Survey, the sky has not been thoroughly explored at $V\\sim 11$."
    },
    "0801/0801.2687_arXiv.txt": {
        "abstract": "We present new results with a prototype detector that is being developed by the DMTPC collaboration for the measurement of the direction tag (\\headtail) of dark matter wind. We use neutrons from a \\Cf source to create low-momentum nuclear recoils in elastic scattering with the residual gas nuclei. The recoil track is imaged in low-pressure time-projection chamber with optical readout. We measure the ionization rate along the recoil trajectory, which allows us to determine the direction tag of the incoming neutrons. ",
        "introduction": "\\label{sec::introduction} The non-baryonic dark matter in the form of weakly interacting massive particles (WIMPs) still eludes detection despite recent achievements in the detection technology~\\cite{DMreview}. Aside from scaling up the size of existing detectors, the improvement in the detection sensitivity is possible by  detecting  the direction of the incoming dark matter particles. As the Earth moves in the galactic halo, the dark matter particles appear to come from Cygnus constellation. The direction tag of the of the incoming particle, often referred to as the \\headtail effect, increases the sensitivity of a directional detector by one order of magnitude~\\cite{agreen}.  In this paper we present improved results for tagging the direction of low-energy nuclear recoils created by neutrons from a $^{252}$Cf source by using a time-projection chamber with optical readout. The neutrons are used in lieu of the dark matter particles because they create similar distributions of recoil energies and angles. The measurement of directionality  tag relies on the fact that the ionization rate of  recoiling nuclei depends on their residual energy, and therefore the direction of the recoil can be tagged from the light distribution along the track. ",
        "conclusions": "\\label{sec::summary}   Since the measured light yield is proportional to the energy of the recoil segment and the length is proportional to the track range projected to the wire, these two quantities should be correlated. Figure~\\ref{fg::recoil_energy_vs_length} shows  clear correlation between the light yield  versus length of the recoil segments.    \\begin{figure}[hb] \\includegraphics[width=7.5cm]{run00251-00257b_range_vs_energy.pdf} % \\hspace{2pc}% \\begin{minipage}[b]{7.5cm}\\caption{ Correlation between the energy and the projected range of the recoil segments in data. \\label{fg::recoil_energy_vs_length} } \\end{minipage} \\end{figure}    We collect 1~day of live-time of data without sources and find two events that pass our standard selection cuts. We verify good rejection of gammas by collecting 1/3~day of live-time of data with $^{137}$Cs source (8~$\\mu$Ci) placed near the sensitive area of our detector and find zero events passing the cuts.    We assign a  conservative error of 10\\% to the density of the \\cf4 gas. The statistical uncertainty on the energy measurements is about 10\\%. The systematic error on the energy comes from non-uniformity in wire gain, stability of the gain over time, the pressure measurement and the calibration method that assumes the energy-independent proportionality of the stopping power with the ionization rate. The error on the recoil range comes from the analysis technique that overestimates the range for low-energy recoils with the range close to the diffusion width.         We have presented improved results for tagging the direction of low-momentum nuclear recoils generated in the elastic scattering of low-energy neutrons with \\cf4 gas. We have shown that in our current experimental setup the tag of incoming particle can be determined for recoil energies above 200~keV. This threshold can be further reduced with expected improvements in the detector preformance. This study has profound implications for the development of  dark matter detectors, as the directionalty will be essential to produce convincing evidence for dark matter particles in the presence of backgrounds."
    },
    "0801/0801.1017_arXiv.txt": {
        "abstract": "An implicit algorithm  for solving the equations of general relativistic hydrodynamics in conservative form in three-dimensional axi-symmetry is presented. This algorithm is a direct extension of the pseudo-Newtonian implicit radiative magnetohydrodynamical solver -IRMHD- into the general relativistic regime.  We adopt  the Boyer-Lindquist coordinates and formulate the hydrodynamical equations in the fixed background of a Kerr black hole. The set of equations are solved implicitly using the hierarchical solution scenario (HSS).  The HSS is efficient, robust and enables the use of a variety of solution procedures that range from a purely explicit up to fully implicit schemes. The discretization of the HD-equations is based on the finite volume formulation and the defect-correction iteration strategy for recovering higher order spatial and temporal accuracies. % Depending on the astrophysical problem, a variety of relaxation methods can be applied. In particular the vectorized black-white Line-Gauss-Seidel relaxation method is most suitable for modeling accretion flows with shocks, whereas the Approximate Factorization Method is for weakly compressible flows.  The results of several test calculations that verify  the accuracy and robustness of the algorithm are shown. % Extending the algorithm to enable solving the non-ideal MHD equations in the general relativistic regime is the subject of our ongoing research. ",
        "introduction": "The field of astrophysical fluid dynamics (henceforth AFD) deals with the macroscopic evolution of gaseous-matter and plasmas in a wide variety of circumstances in astrophysics. The scope of AFD is broad, encompassing topics such as star formation, accretion phenomena onto compact and young stellar objects, dynamics of the interstellar medium, jets, winds and outflows emerging from around young stellar objects, from quasars and microquasars, supernovae explosion, $\\gamma-$ray bursts and structure formation in the universe.   One of the ultimate aims of numerical astrophysics is to develop a black box algorithm which contains numerical solvers that are unconditionally stable, robust, efficient, Newtonian, relativistic and capable of treating flows that are strongly compressible, weakly incompressible, self-gravitating, radiating, magnetized multi\\--com\\-po\\-nent-plasmas with high spatial and temporal accuracies on unstructured meshes and to provide the required solutions within the scale of hours. While this goal is unlikely to be achieved within the next few years, the increasing number of sophisticated numerical algorithms developed during the last two decades is remarkably encouraging. % In particular, significant improvements have been achieved in increasing the spatial dimensions and enhancing  the efficiency and accuracy of numerical algorithms \\citep[see e.g.][]{Nagel_etal2006}. \\\\ On the other hand, the problem of robustness of the solvers in AFD has been barely considered  nor even seriously discussed. % In this paper we discuss the robustness problem in AFD, present enhancement strategies and address the necessity of constructing robust general relativistic implicit and radiative MHD solvers. For completeness we review the concepts of efficiency and robustness of numerical solvers in computational fluid dynamics.  A numerical solver is said to be relatively efficient if the corresponding number of algebraic manipulations per time step per computer-processor can be made respectively small. As a consequence, using  high performance computers with large number of computing processors, a significant progress has been achieved in improving runtime efficiencies,  provided the computing load is distributed appropriately. Thus, by means of modern hardware technology, the efficiency of computer-codes could be enhanced even without modifying the employed mathematical approach. This  has led to uncoordinated developments of relatively large number of computer-codes, that may differ in efficiency and integrated physical processes but are almost essentially identical with respect to the mathematical solution procedure. This phenomenon can be attributed to the lack of robustness. By robustness, we mean the capability of a solver to be applied to a large class of problems without modifying the algorithmic core significantly.   In an attempt to enhance both efficiency and robustness, the hierarchical solution strategy (HSS) has been suggested (Fig. \\ref{ClusterMatrix} and \\ref{HSS}, also see \\citet{Hujeirat2005} and the references therein). The HSS relies on the fact that any set of hydrodynamical, magnetohydrodynamical or radiative equations are linearize-able and therefore can be re-written in a simple matrix equation \\( A q = b,\\) where $A,q,b$ are the coefficient matrix, vector of unknown variables and the vector of known quantities, respectively. Applying the defect-correction strategy \\citep{Stetter1978, BoehmerStetter1984}, we may then re-write this equation as: \\(A\\delta q = d,\\) where $\\delta q$ and $d$ denote respectively the vectors of small-time corrections and the defect, provided  $d$ is consistent with the real mathematical equations by construction. The matrix A in the latter equation can be replaced by a variety of matrices that correspond to a sequence of numerical approaches that ranges from purely explicit to strongly implicit \\citep{Hujeirat2005}. In this formulation, explicit methods arise as a special case, in which A is being replaced by the most easiest invertible-matrix: the identity matrix I. Based on this formulation, the Courant condition follows  from the requirement that the matrix A should be stable-invertible.  Therefore,  strongly implicit and explicit methods are different variants of the same algebraic problem. While the former retain almost all the entries of the matrix A in the inversion procedure, the latter rely on neglecting all off-diagonal entries as well as crudely simplifying the diagonal elements. These methods are well-unified within the HSS, and that depending on the physical properties of the flow, a directive will carefully select the entries of A that are relevant for the problem.   In Table \\ref{Table1} we have summarized properties of  several numerical methods. Thus, as long as efficiency is concerned, explicit methods are unrivalled candidates, provided the flows are strongly time-dependent and compressible. {However, due to the relatively large sound speed, these methods may stagnate both in modeling incompressible or even weakly compressible flows. To clarify this point, we mention that the time step-size in explicit methods must satisfy the Courant-Friedrichs-Lewy (CFL) condition}: \\[\\delta t_{exp} < \\min\\{\\DD{\\Delta x}{U + V_s}\\} = \\min\\{\\DD{\\Delta x}{U(1 + 1/\\cal{M})}\\}, \\] where the minimum function runs over all points of the domain of calculation. Here $\\delta t_{exp}, \\Delta x, U, V_s, \\cal{M}$ correspond to the explicit time step-size,  space increment, velocity, sound speed and the Mach number, respectively. Therefore, as \\({\\cal{M}}\\rightarrow 0 \\) the flow becomes incompressible and the time step size approaches zero; hence a stagnation of the time-advancement procedure. We note that although in this case using consistent implicit solution procedures is necessary, by no means it is sufficient. Here it has been verified that standard implicit solvers experience serious difficulties in simulating low Mach number flows, typically found in the interior layers of stars, planets as well as  around moving vehicles in the Earth atmosphere.  The above discussion addresses the following questions: 1) Relativistic fluid motions typically occur on the dynamical time scale. The advantages of still using implicit solvers should be clarified. 2) Multigrid methods have been shown to display convergence which depends  weakly on the number of unknowns in the finite space. In combination with nested iteration, the multigrid method can solve these problems to truncation error accuracy in a number of operations that is proportional to the number of unknowns. Therefore the reason for still favouring the prolongation strategy over multigrid or adaptive mesh refinement needs to be explained. 3) The storage capacities of modern computers to date are capable of handling the entries of large matrices that correspond to the 3D MHD equations. Thus, the reliability and credibility  of 3D axi-symmetric algorithms should be justified.  In fact, there are several reasons that justify using implicit numerical procedures for modeling relativistic flows. In particular: \\bit \\item The set of relativistic MHD equations is generically a highly coupled non-linear system, which gives rise to fast growing perturbations due to non-linearities, thereby imposing a further restriction on the size of the time step. \\item In most of the cases the horizon of a black hole represents a geometrical singularity.  The deformation of the geometry grows non-linearly when approaching the black hole. Thus, in order to capture flow-configurations  in  the vicinity of a black hole accurately, a non-linear distribution of the grid points is necessary, which, again, may destabilize explicit schemes. \\item  Depending on the evolutionary conditions, non-relativistic flows may become ultra-relativistic or vice versa. However, almost all non-relativistic astrophysical flows known to date are considered to be dissipative and diffusive. Therefore, in order to track their time-evolution reliably, the  employed numerical solver should be capable of treating  the corresponding second order viscous terms properly. \\item  The timescales of most astrophysical flows are considered to have a great disparity. Stability requirement of conditionally stable methods however requires that the time step size should be a small fraction of the shortest possible timescale. This implies that, in order to cover a timescale of astrophysical relevance,  an extremely large number of time steps would be required, which would give rise to prohibitive computational costs. Furthermore, the accumulated round off errors resulting from performing a large number of time-extrapolations for time-advancing a numerical hydrodynamical solution may easily cause divergence. \\item The initial conditions of most astrophysical phenomena are not known. Therefore, in carrying out global hydrodynamical simulations, the end solution should weakly depend on the initial flow configuration. Conditionally stable numerical methods, however, rely on time-advancing of the initial conditions. \\eit The latter reason may explain also why using the prolongation strategy is advantageous over classical multigrid. Worth noting  is that the main building blocks of multigrid methods are: \\ben \\item Restriction, i.e., down sampling of the residual errors into coarser meshes. \\item Residual smoothing: reducing the high frequency errors by performing several iterations, using a computationally efficient iterative procedure such as Jacobi or Gauss-Seidel. \\item Prolongation, which relies on the interpolation from the coarse onto finer meshes. \\een The high-frequency errors here are reduced by cheap smoothing on the  fine meshes, whereas the low-frequency errors are reduced by defect correction on the coarser meshes. As the bulk of the algebraic operations are made on the coarser meshes, the combined solution procedure is considerably efficient. However, multigrid methods display satisfactory convergence, if the underlying flows are predominantly diffusion-dominated. In the case of advection-dominated flows, errors,  that are responsible for the slow convergence on the fine meshes, can be easily advected  by the flow on the coarser meshes, thereby reducing the coarse grid correction. In the case of astrophysical flows, the corresponding equations may change their character from Newtonian into ultra-relativistic or vice versa. Unlike Newtonian flows that are generically diffusion-dominated, relativistic flows may become predominantly advection-dominated, depending on how large the corresponding Lorentz factors are. Hence, multigrid methods may fail to provide the expected rate of convergence.  Finally we note that in order to model the formation and acceleration of relativistic flows in the vicinity of ultra-compact objects accurately, it is necessary to cover the domain of calculation by a strongly stretched mesh. Furthermore,  Lorentz factors enhance the  inner-coupling of the relativistic equations and give rise to a larger spectrum of non-linearities. These numerical difficulties combined with the need to include sophisticated magnetic and radiative processes make the construction of  a fully 3D algorithm, at the moment,  a computationally unrealizable numerical task.  Therefore, in this paper, we do not intend to perform 3D global simulations, but rather focus on the algorithmic structure of unconditionally stable and robust 3D axi-symmetric solvers.  These algorithms should enable us to search for stationary or quasi-stationary solutions for the fully-coupled radiative MHD equations in which detailed physical processes are taken into account. \\\\ The paper runs as follows: in Sec. 2 we describe the relativistic hydrodynamical equations, in Sec. 3   the transformation between the primitive and conservative variables is described. The numerical solution and the discretization methods are presented in Sections 4 and 5. In Sec. 6 we present the results of several test calculations and end up with a summary in Sec. 7.     \\begin{table} \\begin{minipage}{\\linewidth} \\begin{tabular}{l|c|c|c} & Explicit & Implicit & HSS  \\\\\\hline \\parbox[c][6ex][c]{0.15 \\hsize}{solution\\\\ method} & $q^{n+1} = q^n + \\delta t\\,d^n$ &  $q^{n+1} = q^n + \\delta t \\tilde{A}^{-1}d^*$ &  {\\scriptsize $q^{n+1}  = \\alpha q^n + (1-\\alpha) \\delta t \\tilde{A}_d^{-1} d^*$} \\\\ \\hline Type of flows & \\parbox{0.23 \\hsize}{Strongly time-\\\\dependent,\\\\ compressible,\\\\ weakly dissipative\\\\ HD and MHD\\\\ in 1, 2 and 3 dimensions} & \\parbox{0.23 \\hsize}{Stationary, \\\\quasi-stationary,\\\\ highly dissipative,\\\\ radiative and\\\\ axi-symmetric MHD-flows in 1, 2 and 3 dimensions} & \\parbox[c][23ex][c]{0.25 \\hsize}{Stationary,\\\\ quasi-stationary,\\\\ weakly compressible,\\\\ highly dissipative,\\\\ radiative and\\\\ axi-symmetric MHD-flows in 1, 2 and 3 dimensions} \\\\ \\hline Stability & conditioned & unconditioned & unconditioned \\\\ \\hline Efficiency & $1$ (normalized/2D) & $\\sim m^2$ & $\\sim m_d^2$ \\\\ \\hline \\parbox[c][9ex][c]{0.15 \\hsize}{Efficiency:\\\\ Enhancement\\\\ strategies} & Parallelization & \\parbox{0.23 \\hsize}{Parallelization,\\\\ preconditioning,\\\\ multigrid} & \\parbox{0.23 \\hsize}{HSS, parallelization, preconditioning, prolongation} \\\\ \\hline \\parbox{0.15 \\hsize}{Robustness:\\\\ Enhancement\\\\ strategies} & \\parbox[c][12ex][c]{0.23 \\hsize}{i. subtime-stepping \\\\ ii. stiff terms\\\\ are solved\\\\ semi-implicitly} & \\parbox{0.23 \\hsize}{i. multiple iteration\\\\ ii. reducing the time step size} & \\parbox{0.25 \\hsize}{i. multiple iteration\\\\ ii. reducing the time step size, HSS} \\\\ \\hline \\parbox{0.15 \\hsize}{Numerical Codes\\\\ Newtonian} & \\parbox[c][12ex][c]{0.23 \\hsize}{Solvers1${}^{a}$\\\\ {\\scriptsize ZEUS\\&ATHENA${}^{b}$,\\\\ FLASH${}^{c}$, NIRVANA${}^{d}$,\\\\ PLUTO${}^{e}$, VAC${}^{f}$ } } &  Solver2${}^{g}$ &  IRMHD${}^{h}$ \\\\ \\hline \\parbox{0.15 \\hsize}{Numerical Codes\\\\ Relativistic} & \\parbox[c][15ex][c]{0.23 \\hsize}{Solvers3${}^{i}$\\\\ {\\scriptsize GRMHD${}^{j}$, ENZO${}^{k}$,\\\\ PLUTO${}^{l}$, HARM${}^{m}$,\\\\ RAISHIN${}^{n}$, RAM${}^{o}$,\\\\ GENESIS${}^{p}$, WHISKY${}^{q}$ } } & Solver4${}^{r}$ & GR-I-RMHD${}^{s}$ \\\\ \\hline \\end{tabular} \\\\ \\caption{\\label{Table1} A list of only a part of the grid-oriented codes in AFD and their range of applications. The matrix-equations in the first row are illustrated in Sec. 4. In these equations,  $q^{n,n+1}$, ${\\delta t}$, ${\\tilde{A}}$, $\\alpha$ and $d^*$ denote the vector of variables from the old and new time levels, time step size, a preconditioning matrix, a switch on/off parameter and a time-modified defect vector, respectively. ``m\" in row 4 denotes the bandwidth of the corresponding matrix.} \\smallskip \\hrule \\parbox{\\hsize}{\\footnotesize ${}^{a}$\\citet{Bodenheimer_etal1978, Clarke1996}, ${}^{b}$\\citet{StoneNorman1992, GardinerStone2006}, ${}^{c}$\\citet{Fryxell_etal2000}, ${}^{d}$\\citet{Ziegler1998}, ${}^{e}$\\citet{MignoneBodo2003, Mignone_etal2007}, ${}^{f}$\\citet{Toth_etal1998}, ${}^{g}$\\citet{Wuchterl1990, Swesty1995}, ${}^{h}$\\citet{Hujeirat1995, Hujeirat2005, Falle2003}, ${}^{i}$\\citet{Koide_etal1999, Komissarov2004, Anninos2005, Meliani2007, DelZanna2007}, ${}^{j}$\\citet{DeVilliersHawley2003}, ${}^{k}$\\citet{Wang2007}, ${}^{l}$\\citet{Mignone_etal2007}, ${}^{m}$\\citet{Gammie_etal2003}, ${}^{n}$\\citet{Mizuno_etal2006}, ${}^{o}$\\citet{ZhangMacFadyen2006}, ${}^{p}$\\citet{Aloy1999}, ${}^{q}$\\citet{Baiotti_etal2003}, ${}^{r}$\\citet{Liebensdoerfer_etal2002}, ${}^{s}$the present algorithm. } \\end{minipage} \\end{table} ",
        "conclusions": "In this paper we have extended the previous Newtonian implicit algorithm to enable solving the hydrodynamical equations in general relativity. The 3D axi-symmetric hydrodynamical equations have been presented in the background of a Kerr metric of a black hole using the Boyer-Lindquist coordinates. The equations have been formulated in conservative form and subsequently solved numerically, using the finite volume formulation. The new extension can be well accommodated within the hierarchical solution scenario, in which the degree of implicitness can be made dynamical, depending on the hydrodynamical problem in hand. In particular, for modeling strongly time-dependent astrophysical flows, such as moving shocks, the pre-conditioners used are tri-diagonal matrices that are solved successively. Although the computational costs per time step may be one order of magnitude larger than their explicit counterparts, this can be compensated through a reduction of the overall number of time steps required to recover a physically reliable time scale. %  On the other hand, the efficiency and robustness of the HSS are superior, if the solutions sought are stationary or quasi-stationary, irrespective of whether the flow is dissipative or not.  Finally, a unification scheme for various numerical methods has been presented. In particular, the HSS algorithm enables the construction of a large variety of solvers, in which the degree of implicitness may range from purely explicit up to strongly implicit, depending on the physical properties of the underlying flow problem. Thus, the HSS is actually a unified algorithm for treating weakly compressible, incompressible, time-dependent, time-independent, radiative, magnetized non-dissi\\-pative or strongly dissipative flows. As a consequence, using the HSS algorithm, we are able to save a large number of working hours which otherwise would go in designing different solvers for different physical problems.\\\\ In a subsequent paper, we intend to discuss and describe the inclusion of the magnetohydrodynamical equations in general relativity into the present solver. %  {\\bf Acknowledgment} A.H. thanks Prof. J. Du\\v{s}ek for reading the manuscript and for the hospitality during his visit to the Institut de mechanique des fluides et des solides, CNRS and the Louis Pasteur University in Strasbourg. This work has been partially supported by the Klaus-Tschira Stiftung under the project number 00.099.2006."
    },
    "0801/0801.1684_arXiv.txt": {
        "abstract": "Stellar dynamos are driven by complex couplings between rotation and turbulent convection, which drive global-scale flows and build and rebuild stellar magnetic fields.  When stars like our sun are young, they rotate much more rapidly than the current solar rate. Observations generally indicate that more rapid rotation is correlated with stronger magnetic activity and perhaps more effective dynamo action.  Here we examine the effects of more rapid rotation on dynamo action in a star like our sun.  We find that vigorous dynamo action is realized, with magnetic field generated throughout the bulk of the convection zone.  These simulations do not possess a penetrative tachocline of shear where global-scale fields are thought to be organized in our sun, but despite this we find strikingly ordered fields, much like sea-snakes of toroidal field,  which are organized on global scales. We believe this to be a novel finding. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.2861.txt": {
        "abstract": "Gamma-ray bursts (GRBs) are one of the candidates of ultra-high-energy ($\\gtrsim {10}^{18.5}$ eV) cosmic-ray (UHECR) sources. We investigate high-energy cosmic-ray acceleration including heavy nuclei in GRBs by using Geant 4, and discuss its various implications, taking both of high-luminosity (HL) and low-luminosity (LL) GRBs into account. This is because LL GRBs may also make a significant contribution to the observed UHECR flux if they form a distinct population. We show that not only protons but also heavier nuclei can be accelerated up to ultra-high energies in the internal, (external) reverse and forward shock models. We also show that the condition for ultra-high-energy heavy nuclei such as iron to survive is almost the same as that for $\\sim$ TeV gamma-rays to escape from the source and for high-energy neutrinos not to be much produced. The multi-messenger astronomy by neutrino and GeV-TeV gamma-ray telescopes such as IceCube and KM3Net, GLAST and MAGIC will be important to see whether GRBs can be accelerators of ultra-high-energy heavy nuclei. We also demonstrate expected spectra of high-energy neutrinos and gamma rays, and discuss their detectabilities. In addition, we discuss implictaions of the GRB-UHECR hypothesis. We point out, since the number densities of HL-GRBs and LL-GRBs are quite different, its detemination by UHECR observations is also important. ",
        "introduction": "Introduction} Gamma-ray bursts (GRBs) and supernovae (SNe) are the most powerful phenomena in the universe. The latter is believed to be the origin of high-energy cosmic rays below the knee $\\approx {10}^{15.5}$ eV. The former could also accelerate baryons up to high energies if the dissipation process is due to shock dissipation. Waxman, Milgrom \\& Usov and Vietri suggested that ultra-high-energy cosmic rays (UHECRs) are produced in GRBs \\cite{Wax1,Mil1,Vie1}. Their suggestions were based on the two arguments. First is the requirement that relativistic outflows that make GRBs satisfy various conditions for baryons to be accelerated up to greater than ${10}^{20}$ eV. Second is that the energy generation rate required to account for the observed UHECR flux is comparable to the energy generation rate of observed gamma rays. The latter argument depends on the local GRB rate which is not well known observationally. %Recent \\textit{Swift} observations have suggested that the %cosmological high-luminosity GRB (HL GRB) rate increases rapidly %with the redshift $z$, and the local GRB rate %may be somewhat smaller compared to some estimations in the %pre-\\textit{Swift} era (see, e.g., \\cite{Le1,Gue1,Kis1}). If the local cosmological high-luminosity GRB (HL GRB) rate is not high enough, which may be suggested by recent observations \\cite{Le1,Gue1,Kis1}, the required baryon loading becomes larger \\cite{Wic1,Mur1,Ber1}. In addition, if the UHECR spectrum at the source is steeper than that with the spectral index $p \\sim 2.0-2.3$ (as expected in the dip model, $p \\sim 2.4-2.7$ \\cite{Ber1,Alo1}), the total nonthermal cosmic-ray energy of GRBs, necessary for explaining the observed UHECR flux, would be larger than the radiation energy. Despite such caveats, the GRB-UHECR hypothesis is still one of the most interesting possibilities for explanation of observed UHECRs, and it should be examined by future observations.  Recent observations have tentatively suggested that some GRBs may form a different class (see, e.g., \\cite{Lia1,Gue2}). %Some long-duration GRBs like %GRB 980425 and GRB 060218 are associated with core-collapse supernovae %(SNe of Type\\Roman{ichi}c to be more specific). As Guetta et al. %\\cite{Gue3} argued, no bright bursts at $z<0.17$ should have been %observed by a HETE-like instruments within the 20 years, %if the logN-logP relationship of HL GRBs continues to the %low-luminosity region. Therefore, the unexpected discovery of %low-luminosity GRB 060218 lead to the %idea that they form a different new class of GRBs %from the conventional HL GRBs, and GRBs in this class are called low-luminosity GRBs (LL GRBs) or sub-energetic GRBs, which may be more common than HL GRBs. If this speculation is true, they can provide the significant energy as high-energy cosmic rays, neutrinos and gamma rays. Murase et al. \\cite{Mur2} suggested that cosmic rays can be accelerated up to ultra-high energies ($\\gtrsim {10}^{18.5}$ eV) in LL GRBs, and LL GRBs may make an important contribution to the observed high-energy cosmic-ray flux and the diffuse neutrino background \\cite{Mur2,Gup1}. Such LL GRBs were associated with energetic supernovae called hypernovae. As in cases of supernovae, cosmic-ray acceleration is expected at shocks formed by the high-velocity ejecta. In fact, if the magnetic field is sufficiently amplified, cosmic-rays can be accelerated up to very high energies, even to ultra-high energies. Wang et al. \\cite{Wan1,Wan2} suggested that they could explain observed high-energy cosmic rays above the second knee, under the assumption that LL GRBs have the enough local rate higher than that of HL GRBs. %Therefore, if high-energy cosmic rays are produced by %GRB outflows and/or hypernovae, we can expect that not only usual HL %GRBs but also LL GRBs can make an important contribution.  %The origin of UHECRs has been one of the biggest mysteries in the %universe. In order to reveal the origin of UHECRs, UHECRs observatories %with larger exposures have been constructed. Now, the total exposure of Recently UHECR observations have been greatly advanced thanks to (Sourth) Pierre Auger Observatory (PAO). %already exceeds past detectors %such as Akeno Giant Air Shower Array (AGASA). PAO collects many %events, and t The UHECR spectrum deduced by PAO suggests the existence of the Greisen-Zatsepin-Kuz'min (GZK) cutoff energy which arises from the photomeson production process between UHECRs and cosmic microwave background (CMB) photons \\cite{Gre1,Zat1}. In addition, PAO recently revealed the anisotropy and rejected the isotropy of UHECRs above 60 EeV at 3 $\\sigma$ level \\cite{PAO1,PAO2}. %They also suggested %that the positional correlation between nearby sources such as %Active galactic nuclei (AGN). Although It suggests that the distribution of UHECRs traces the matter distribution in the universe, although the origin of UHECRs is still unknown, that is, AGNs, GRBs and other possibilities remain (see the latest review for AGN and GRBs, e.g., \\cite{Der1}). Interestingly, the preliminary elongation rate data by PAO showed the break around $\\sim {10}^{18.35}$ eV \\cite{PAO3}. Whether heavier nuclei becomes more important above this break as the energy increases (which seems conflict with HiRes data) is still under debate. It is because there are large uncertainties in hadronic interactions. On the other hand, recent results on UHECRs above 60 EeV by PAO may infer that protons are dominant sources \\cite{PAO1,PAO2}, although some authors dispute \\cite{Gor1,Wib1}. Although the issue of the composition of UHECRs requires more and more studies, present data seem to allow the presence of significant fractions of ultra-high-energy (UHE) nuclei around the GZK cutoff energy, even if protons are dominant. Therefore, it is one of the important issues whether UHE nuclei can survive (i.e., UHE nuclei are not photodisintegrated and not depleted due to photomeson production during the dynamical time) in UHECR production sites.  In this paper, we study acceleration of high-energy cosmic-ray nuclei in very detail by exploiting numerical calculations, and discuss various implications for multi-messenger astronomy. In Sec. \\Roman{ni}, we show that cosmic rays (protons and heavier nuclei) can be accelerated up to ultra-high energies in both of HL GRBs and LL GRBs. We consider the internal, (external) reverse and forward shock models.  If cosmic rays can be accelerated up to sufficiently high energies, high-energy neutrino production via the photomeson production process and/or $pp$ reaction is unavoidable. Now, large neutrino detectors such as IceCube \\cite{Ahr1}, KM3Net \\cite{Kat1} are being constructed. ANITA \\cite{Bar1} and PAO \\cite{Van1} can also detect very high-energy neutrinos. We have chances to see such high-energy neutrino signals, which are important as a probe of cosmic-ray acceleration, in the future. High-energy neutrino signals from HL GRBs were predicted in the context of the standard scenario of GRBs assuming that observed UHECRs come from GRBs \\cite{Wax2,Wax3,Dai1,Der2}. We also investigated the neutrino emission from GRBs in very detail \\cite{Mur1,Mur2,Mur3,Mur4}, and various predictions in the \\textit{Swift} era are summarized in Ref. \\cite{Mur4}. In Sec. \\Roman{san}, we demonstrate predicted neutrino spectra in some of our models and discuss implications for neutrino astronomy. We can see that neutrino signals highly suppressed when very heavy nuclei like iron can survive.  In Sec. \\Roman{yon}, we also demonstrate gamma-ray spectra for some models when heavy nuclei can be accelerated up to ultra high energies and survive. In this section, we also show that very high-energy gamma rays with GeV-TeV energies can escape from the source under conditions where UHE heavy nuclei such as iron can survive. We can expect to detect high-energy gamma-ray signals by upcoming Gamma-Ray Large Area Telescope (GLAST) and Cherenkov telescopes such as MAGIC as long as the source is nearby. Otherwise, high-energy gamma rays above $\\sim$TeV are attenuated by the cosmic microwave/infrared background (CMB/CIB) photons (\\cite{Mur6} and references there in).  In Appendix F, We also discuss various implications of the GRB-UHECR hypothesis (including the HLGRB-UHECR, LLGRB-UHECR and Hypernova-UHECR hypotheses) for UHECR astronomy. We see that the determination of the source number density of UHECR sources and extragalactic magnetic field (EGMF) strength is essential to test the hypothesis. In this paper, we shall hereafter use notations such as $Q_x=Q/10^x$ in cgs units. ",
        "conclusions": "In this paper, we have shown that not only protons but also heavy nuclei can be accelerated up to ultra-high energies in both of HL GRBs and LL GRBs. We exploit the internal shock model, (external) reverse shock model and forward shock model. We have also discussed various implications for neutrino, gamma-ray and UHECR astronomy. Especially, we have studied cosmic-ray acceleration in LL GRBs which may play an important role as high-energy cosmic-ray sources.  Let us summarize this paper below. First, we have shown that UHECR nuclei can be produced in both of HL GRBs and LL GRBs.  (A1) In the internal shock model, UHE protons and heavier nuclei can be produced in both of HL GRBs and LL GRBs. However, the allowed parameter range is limited. At smaller radii, most of the UHE nuclei are depleted, and a significant fraction of the nonthermal cosmic-ray energy is transferred into neutrinos. On the other hand, survival of UHE nuclei is possible only at large radii and/or for large Lorentz factors. Typically, relatively large radii $r \\gtrsim {10}^{15}$ cm will be required. For HL GRBs with $L_{b} \\sim {10}^{51-52} \\, \\rm{erg} \\rm{s}^{-1}$, both of protons and heavier nuclei can be accelerated up to ultra-high energies above ${10}^{20}$ eV. For LL GRBs with $L_{b} \\sim {10}^{46-47} \\, \\rm{erg} \\rm{s}^{-1}$, heavy nuclei can be accelerated above ${10}^{20}$ eV, while proton acceleration above such high energies may be difficult. In addition, we have shown that the effect of synchrotron self-absorption suggested by Wang et al. \\cite{Wan2} is not so important due to the cross section in the non-resonant region.  (A2) In the (external) reverse shock model, UHE nuclei can be produced in both of HL GRBs and LL GRBs. Not only heavy nuclei but also protons achieve ultra-high energies above $\\sim {10}^{20}$ eV. Whether heavy nuclei can survive or not depends on various parameters, and the sufficiently low photon density is required for survival. That is, small $E_{\\rm{ej}}$, $\\epsilon_{e}$, $\\epsilon_{B}$ and $n$ are needed. Conversely, if values of these parameters are large enough, in which optical/infrared are often expected, survival of UHE nuclei is impossible. Note that UHE nuclei produced at internal shocks could be photodisintegrated by photons generated at the reverse shock, when the ejecta is in the thick ejecta regime. Hence, we could not expect survival of UHE nuclei, when we see strong optical/infrared flashes.  (A3) In the (external) forward shock model, not only heavy nuclei but also protons can achieve ultra high energies above $\\sim {10}^{20}$ eV in both of HL GRBs and LL GRBs. When the photon density is sufficiently low, survival of UHE nuclei is possible. The relevant parameters $f_{N \\gamma}$ and maximum energy $E_{N}^{\\rm{max}}$ depend on time. In the ISM case, $f_{N \\gamma}$ monotonically increases with time before the jet break occurs. On the other hand, in the wind-medium case, $f_{N \\gamma}$ monotonically decreases with time. We have shown that the jet-break effect can be important for survival of UHE nuclei, which was not pointed out previously. After the jet break, the photon density decreases, so that survival of UHE nuclei becomes easier.  In the internal shock model and reverse shock model, we have to assume that heavy nuclei are contained in relativistic outflows of GRBs. It is natural to consider that only light elements can be synthesized in such a highly relativistic outflows due to the high photon to baryon ratio~\\cite{pruet02,lemoine02}. However, there may be possibility that such an outflow is contaminated with heavy nuclei that are produced by explosive nucleosynthesis~\\cite{nagataki03,maeda03,nagataki06,tominaga07}. Also, such heavy nuclei can be entrained from the stellar surroundings due to Kelvin-Helmholtz instabilities and/or oblique shocks \\cite{Zha2}. Whether entrained heavy nuclei can survive or not is the issue that should be carefully examined. On the other hand, in the forward shock model and hypernova model, heavy nuclei will be supplied by the swept material. They could come from the stellar wind of the progenitor star, e.g., the Wolf-Rayet star.  (B) Neutrino astronomy is useful as a direct probe of cosmic-ray acceleration, although the detection is not so easy. High-energy neutrinos from GRBs have been predicted in various contexts. If detected, we can obtain very useful information on cosmic-ray acceleration in GRBs. In this paper, we have calculated neutrino spectra from HL GRBs and LL GRBs in the internal shock model. We have shown that the neutrino detection by near-future telescopes such as IceCube and KM3Net would be difficult, when acceleration and survival of UHE heavy nuclei (such as iron) are possible. In the internal shock model, for example, we can expect $N_{\\mu} \\sim 1$ event from one GRB at $z=0.1$ and $N_{\\mu} \\sim 50$ events/yr as the cumulative background for one of the typical parameter sets (where UHE heavy nuclei cannot survive), with large baryon loading factors required in the GRB-UHECR hypothesis. On the other hand, survival of UHE heavy nuclei leads to much smaller neutrino fluxes, as we have demonstrated. Hence, second generation neutrino telescopes would be needed for neutrino detection in the latter case, although observations by IceCube and KM3Net would be useful to constrain the GRB-UHECR hypothesis.  (C) We can also expect high-energy gamma-ray emissions (including leptonic and hadronic gamma rays) from GRBs as well as high-energy neutrinos. We have shown that high-energy gamma rays above TeV can escape from the source, when UHE heavy nuclei (such as iron) can survive. As an example, we have calculated and evaluated gamma rays coming from accelerated cosmic rays in the internal shock model. If the source is nearby, we have seen that high-energy gamma rays from cosmic-ray synchrotron radiation would be detected by GLAST, MAGIC and VERITAS, with large baryon loading factors required in the GRB-UHECR hypothesis. (However, we should keep in mind that the synchrotron self-inverse Compton radiation by accelerated electrons could be more important.) Hence, the high-energy gamma-ray detection will also be one of the important clues to testing the GRB-UHECR hypothesis. When the intrinsic pair-creation optical depth becomes high enough (which infers UHE nuclei cannot survive), escape of TeV gamma rays becomes impossible and the intrinsic pair-creation cutoff should exist. But we may still find signatures of UHECR acceleration \\cite{Asa2}, although we cannot expect survival of UHE heavy nuclei in GRBs.  In this paper, we have demonstrated spectra of high-energy neutrinos and gamma rays, which are expected if GRBs are UHECR accelerators, by using detailed numerical calculations. We have also discussed detectabilities of such high-energy emission in the near future. Not only neutrino and gamma-ray astronomy but also UHECR astronomy are becoming important to test the GRB-UHECR hypothesis. In Appendix F, we have discussed implications of the GRB-UHECR hypothesis for UHECR astronomy by exploiting qualitative arguments. Future observations of UHECRs may allow us to distinguish between bursting sources and steady sources. If UHECR sources are bursting sources such as GRBs, we expect to obtain information on the effective EGMF and source number density by using the observed anisotropy of UHECRs. If we can, we could constrain the local bursting rate, for example, which may allow us to distinguish between HL GRBs and LL GRBs.  When we are preparing this work, we found that Wang et al. \\cite{Wan2} studied acceleration of UHE nuclei in GRBs independently. In this paper, we have studied in more detail by making use of the detailed cross section and spectra of afterglows, and calculate spectra of high-energy neutrinos and gamma rays. Furthermore, we have also studied cosmic-ray acceleration in LL GRBs.  %% If you wish to include an acknowledgments section in your paper, %% separate it off from the body of the text using the"
    },
    "0801/0801.4542_arXiv.txt": {
        "abstract": "{This paper is the written version of the rapporteur talk on Section HE-2, {\\em muons and neutrinos}, presented at the 30th International Cosmic Ray Conference, Merid\\'{a}, Yucatan, July 11, 2007.  Topics include atmospheric muons and neutrinos, solar neutrinos and astrophysical neutrinos as well as calculations and instrumentation related to these topics. } \\begin{document} ",
        "introduction": "There were 5 sections of contributed papers on muons and neutrinos with a total of 107 papers distributed as shown in Table~\\ref{table1}.  The most active category is HE 2.3, {\\em astrophysical neutrinos}.  Particularly in this area, there were also many papers presented in OG 2.5 sessions, on {\\em high-energy astrophysical neutrinos}. I include discussion of these topics to the extent necessary to present a coherent overview of the field as of mid-year 2007.  \\begin{table}[htb] \\label{table1} \\caption{Papers on muons and neutrinos} \\begin{tabular}{l|l|l}\\\\ \\hline Session & Topic & \\# \\\\ \\hline HE 2.1 & Muon experiments & 17 \\\\ HE 2.2 & Observations of solar & \\\\ & \\& atmospheric $\\nu$ & 16 \\\\ HE 2.3 & Observations of & \\\\ & astrophysical $\\nu$ & 37 \\\\ HE 2.4 & Theory and simulations & 19 \\\\ HE 2.5 & New experiments & \\\\ &\\& instrumentation & 18 \\\\ \\hline \\end{tabular} \\end{table}   \\begin{figure} \\vspace{-.5cm} \\begin{center} \\noindent \\includegraphics[width=0.5\\textwidth]{fig16june14.eps} \\end{center} \\caption{Muon charge-ratio from Ref.~\\protect\\cite{MINOS}.  The line ``$\\pi K$\" model corresponds to a fit to the ratio of $\\mu^+/\\mu^-$ from Eq.~\\ref{muons}. } \\label{fig1} \\end{figure} ",
        "conclusions": ""
    },
    "0801/0801.1960_arXiv.txt": {
        "abstract": "{Observations of protostellar disks indicate the presence of the magnetic field of thermal (or superthermal) strength. In such a strong magnetic field, many MHD instabilities responsible for turbulent transport of the angular momentum are suppressed.} {We consider the shear-driven instability that can occur in protostellar disks even if the field is superthermal.} {This instability is caused by the combined influence of shear and compressibility in a magnetized gas and can be an efficient mechanism to generate turbulence in disks.} {The typical growth time is of the order of several rotation periods.} {} ",
        "introduction": "Protostellar disks require sufficiently strong turbulence to enhance the efficiency of angular momentum transport. The origin of turbulence is often attributed  to hydrodynamic and hydromagnetic instabilities that can arise in differentially-rotating, stratified gaseous disks. One of the candidates is  magnetorotational instability (MRI), which  can operate in a conductive flow if the angular velocity decreases with the cylindrical radius and the magnetic field is not strong (Velikhov 1959). The MRI has been studied in depth in the case of  stellar and accretion disk conditions (see, e.g., Fricke 1969; Safronov 1969; Acheson 1978, 1979; Balbus \\& Hawley 1991; Kaisig et al. 1992; Kumar et al. 1994; Zhang et al. 1994). Simulations of the MRI in disks (Hawley et al. 1995, Matsumoto \\& Tajima 1995, Brandenburg et al. 1995, Torkelsson et al. 1996, Arlt \\& R\\\"udiger 2001) show  the turbulence generated can significantly enhance  the angular momentum transport.  Most likely, however, the number of instabilities that can arise in astrophysical disks is quite large. An analysis of MHD modes in stratified accretion disks demonstrates a wide variety of instabilities even in the case of simple magnetic geometry (Keppens, Casse \\& Goedbloed 2002). Therefore, the current point of view on the origin of turbulence in disks is likely highly simplified. Even pure, hydrodynamic origin of turbulence cannot be excluded (see Urpin 2003, Arlt \\& Urpin 2004, Dubrulle et al. 2005) despite the most efficient local linear hydrodynamic instabilities advocated to date are not sufficiently efficient. Global instabilities, such as the baroclinic-like instability of Klahr \\& Bodenheimer (2003), are sensitive to boundary conditions (Johnson and Gammie 2006) and, therefore, are unlikely to drive turbulence in disks. { Note, however, that such factors as fast radiative cooling, high thermal diffusion, and large radial temperature gradients can strengthen the baroclinic feedback and make the baroclinic instability more efficient (Petersen et al. 2007).} Astrophysical disks are stratified and stratification can change stability properties of shear flows, providing either a stabilizing or destabilizing effect, depending on details of the disk structure. Recently, Dubrulle et al. (2005) have considered one of the possible ``strato-rotational'' instabilities arising in the presence of both differential rotation and stable stratification. However, Brandenburg \\& R\\\"udiger (2005) demonstrated that the growth rate of this instability decreases with an increasing Reynolds number, rendering the instability less relevant for astrophysical applications. Note also that convection can drive turbulence in disks with unstable stratification, but it induces inward turbulent transport of the angular momentum instead of the required outward one (see, e.g., Stone \\& Balbus 1996).  As it has been argued by Lesur \\& Longaretti (2005), nonlinear hydrodynamic instability that often occurs in linearly-stable flows at sufficiently large Reynolds numbers is likely inefficient in disks even for the most optimistic extrapolations of numerical data. At least, the subcritical transition to hydrodynamic turbulence cannot occur in quasi-keplerian flows at Reynolds numbers up to $\\sim 10^6$ (Ji et al. 2006). Of course, Reynolds numbers are much higher in real disks; nevertheless it seems that there is little room for nonlinear instability.  The stability properties of protostellar disks are much different from those of accretion disks. The magnetic Reynolds number is likely not very large in cold- and dense-protostellar disks because of a low electrical conductivity, and so the field cannot be treated as \"frozen\" into the gas (Gammie 1996). The effect of Ohmic dissipation on the MRI has been considered in the linear (Jin 1996) and nonlinear regimes (Sano, Inutsuka \\& Miyama 1998; Drecker, R\\\"udiger \\& Hollerbach 2000). Fleming \\& Stone (2003) and Turner et al. (2006) treated turbulent mixing caused by the MRI in protostellar disks,  taking into account magnetic diffusivity. They found that the midplane is shielded from cosmic rays, and that MRI does not occur even under the most favorable conditions. Nevertheless, turbulence can mix a fraction of the weakly-conducting surface material into the interior, providing some coupling of the midplane gas to the magnetic field. As a result, weak stresses can appear in the disk midplane.  As first pointed out by Wardle (1999), poorly conducting protostellar disks can be strongly magnetized if electrons are the main charge carriers. Therefore, transport must be  anisotropic with substantially different properties along and across the magnetic field. In strongly magnetized gas the Hall component provides the main contribution to the resistivity tensor, which produces  the electric field perpendicular to both the magnetic and electric current. The stability analysis by Wardle  (1999) shows that the Hall effect can provide either stabilizing or destabilizing influences depending on the direction of the field. Balbus \\& Terquem (2001)  have conducted a more general study of the role of the Hall term on MRI.  They found that the Hall effect qualitatively changes the stability properties of protostellar disks and can lead to instability even if the angular velocity increases outward.  These authors, however, did not take into account the effect of gravity that is crucial for disks. Urpin \\& R\\\"udiger (2005) considered MRI under the combined influence of the Hall effect and gravity, and also derived also the criteria of several other instabilities that can occur in protostellar disks. Sano and Stone (2002) investigated the effect of the Hall term on the evolution of the MRI in weakly-ionized disks using local axisymmetric simulations. These authors concluded  the properties of the  MRI depend essentially on the direction of the magnetic field as it is anticipated from the dispersion equation in a linear stability analysis. Salmeron \\& Wardle (2005) have also considered the properties of the MRI modified by the Hall effect. These authors argued that the MRI is active in protoplanetary disks over a wide range of field strengths and fluid conditions. The Hall conductivity results in a faster growth of perturbations and extends the region of instability. Recently, Livertz et al. (2007) and Shtemler et al. (2007) have considered the Hall MRI in a non-axisymmetric case. This type of instability is proposed as a viable mechanism for the azimuthal fragmentation  of the protoplanetary disks and planet formation. The non-axisymmetric instability is caused  by the combined effect of the radial stratification and Hall electric field. Note that the MRI and its modifications can be completely suppressed if the magnetic field is sufficiently strong (see Urpin 1996, Kitchatinov \\& R\\\"{u}diger 1997).  If rotation is cylindrical and $\\Omega = \\Omega(s)$, where $\\Omega$ is the angular velocity and $s$ is the cylindrical radius, the critical magnetic field that suppresses the MRI is given approximately by the condition $c_{A} / H \\sim s \\Omega' \\sim \\Omega$ where $c_{A} = B/ \\sqrt{ 4 \\pi \\rho}$ is the Alfv\\'en velocity. Since $\\Omega \\sim c_s /H$ in the standard disk, where $c_s$ is the sound speed, we find that the MRI is suppressed if the magnetic field is superthermal, $c_{A} > c_s$. Recent measurements (Hutawarakorn \\& Cohen 1999, 2005, Donati et al. 2005) indicate that the magnetic field can be strong in protostellar disks. This particularly concerns the innermost regions where the field strength reaches $\\sim 1$ kG at the radius $\\sim 0.05$ AU (Donati et al. 2005). The field configuration includes the azimuthal component whose direction agrees with the radial field  sheared by the disk differential rotation. The derived ratio of the azimuthal and radial fields in the surface layers is $\\sim 0.5$. Note that the azimuthal component can be substantially stronger than the radial one in the disk interiors because dynamo theories (such as the $\\alpha - \\Omega$ dynamo) usually predict that the generated field should have a strong toroidal component. According to estimates by Donati et al. (2005), the equipartition field strength with roughly equal thermal and magnetic pressures should be $\\sim 10^3$ Gauss at the radius $\\sim 0.05$ AU in protostellar disks. This indicates the detected field is approximately of the thermal or even superthermal strength and can substantially modify or even suppress the MRI.  Recently, Bonanno \\& Urpin (2006, 2007) have argued that a compressible, differentially-rotating flow is unstable if the magnetic field has a non-vanishing, radial component. A remarkable feature of this shear-driven instability is that it can arise even if the magnetic field is superthermal. Bonanno \\& Urpin (2006, 2007) have considered the instability in the case of a highly conductive plasma when the magnetic diffusivity plays no role. However, the instability can also arise in weakly-ionized protostellar disks where the magnetic field has a radial component detected in observations (Donati et al. (2005). The effect of a finite diffusivity can be important for the onset of instability and for the properties of generated turbulence, particularly near the the mid-plane where conductivity is extremely low. In this paper, we consider the shear-driven instability in the conditions of protostellar disks and show that the conditions of instability in weakly and highly ionized plasmas differ substantially. Nevertheless, the shear-driven instability can manifest itself even in weakly-ionized protostellar disks.  The paper is organized as follows. In Section 2, we consider the basic equations governing instability in compressible dissipative fluids and we then derive the dispersion relation. In Section 3, we derive the criteria of instability and discuss  conditions under which instability can occur in protostellar disks. The growth rate of instability is calculated in Section 4. Finally, in Section 5, we discuss the results obtained. ",
        "conclusions": "We have considered the new instability that can arise in protostellar disks under the combined influence of the magnetic field, differential rotation, and compressibility. To illustrate the main qualitative features of the instability, we analyzed a particular case of perturbations with the wavevector $\\vec{k}$ perpendicular to the magnetic field $\\vec{B}$. In this case, the standard MRI does not occur because its growth rate is $\\propto (\\vec{k} \\cdot \\vec{B})$. The considered instability is related to shear and compressibility of a magnetized gas. In the incompressible limit ($c_{s} \\rightarrow \\infty$) we have, from Eq.~(27), \\begin{equation} (\\sigma^2 + \\kappa^2) (\\sigma + \\omega_{\\eta})^{2} =0, \\end{equation} and the instability does not occur. This is a major difference from other instabilities caused by differential rotation, such as the Rayleigh or magnetorotational instabilities that can occur in incompressible limit.  The magnetic field is likely thermal or superthermal in protostellar disks (see Donati et al. 2005). The considered instability can arise even in a very strong magnetic field when the MRI is suppressed completely, It is known that the MRI does not occur if the magnetic field satisfies the condition $B > B_{cr}= (s \\Omega' /k) \\sqrt{4 \\pi \\rho}$. The shear-driven instability can operate even if the field is stronger than $B_{cr}$. The growth time of instability is still rather short (several rotation periods) in the gas where the magnetic pressure is greater than the thermal pressure.  The criteria for the considered instability can be satisfied in protostellar disks. Likely, all disks have both radial and azimuthal magnetic fields. These magnetic fields are found in observations, as well as in numerical simulations. Superthermal magnetic fields do not suppress the instability  crucial in cold protostellar disks.  The instability can arise even in a low conductive gas where the magnetic Reynolds number associated with rotation is relatively small, Re$_m \\sim 100$.  Our paper considers only a local axisymmetric instability. Most likely, the same mechanism of instability will efficiently destabilize perturbations with the wavelength comparable to the disk scale height. Dissipative effects are less important for these perturbations, and the unstable modes can grow even faster. A global instability caused by a combined influence of differential rotation and compressibility will be considered elsewhere. The turbulence that could be generated by the considered instability can be strongly anisotropic in the $(s,z)$-plane because both the criteria and growth rate are sensitive to the direction of the wave vector. The generated turbulence, however,  might be efficient in the radial transport of angular momentum in strongly magnetized protostellar disks."
    },
    "0801/0801.1151_arXiv.txt": {
        "abstract": "We continue our work on developing techniques for studying turbulence with spectroscopic data. We show that Doppler-broadened absorption spectral lines, in particularly, saturated absorption lines,  can be used within the framework of the earlier-introduced technique termed the Velocity Coordinate spectrum (VCS). The VCS relates the statistics of fluctuations along the velocity coordinate to the statistics of turbulence, thus it does not require spatial coverage by sampling directions in the plane of the sky. We consider lines with different degree of absorption and show that for lines of optical depth less than one, our earlier treatment of the VCS developed for spectral emission lines is applicable, if the optical depth is used instead of intensity. This amounts to correlating the logarithms of absorbed intensities. For larger optical depths and saturated absorption lines, we show, that the amount of information that one can use is, inevitably, limited by noise. In practical terms, this means that only wings of the line are available for the analysis. In terms of the VCS formalism, this results in introducing an additional  window, which size decreases with the increase of the optical depth. As a result, strongly saturated absorption lines carry the information only about the small scale turbulence. Nevertheless, the contrast of the fluctuations corresponding to the small scale turbulence increases with the increase of the optical depth, which provides advantages for studying turbulence combining lines with different optical depths. We show that, eventually, at very large optical depths the Lorentzian profile of the line gets important and extracting information on velocity turbulence, gets impossible. Combining different absorption lines one can tomography turbulence in the interstellar gas in all its complexity. ",
        "introduction": "Turbulence is a key element of the dynamics of astrophysical fluids, including those of interstellar medium (ISM), clusters of galaxies and circumstellar regions. The realization of the importance of turbulence induces sweeping changes, for instance, in the paradigm of ISM. It became clear, for instance, that turbulence affects substantially star formation, mixing of gas, transfer of heat. Observationally it is known that the ISM is turbulent on scales ranging from AUs to kpc (see Armstrong et al 1995, Elmegreen \\& Scalo 2004), with an embedded magnetic field that influences almost all of its properties.  The issue of quantitative descriptors that can characterize turbulence is not a trivial one (see discussion in Lazarian 1999 and ref. therein). One of the most widely used measures is the turbulence spectrum, which describes the distribution of turbulent fluctuations over scales. For instance, the famous Kolmogorov model of incompressible turbulence predicts that the difference in velocities at different points in turbulent fluid increases on average with the separation between points as a cubic root of the separation, i.e. $|\\delta v| \\sim l^{1/3}$. In terms of direction-averaged energy spectrum this gives the famous Kolmogorov scaling $E(k)\\sim 4\\pi k^2 P({\\bf k})\\sim k^{5/3}$, where $P({\\bf k})$ is a {\\it 3D} energy spectrum defined as the Fourier transform of the correlation function of velocity fluctuations $\\xi ({\\bf r})=\\langle \\delta v({\\bf x})\\delta v({\\bf x}+{\\bf r})\\rangle$. Note that in this paper we use $\\langle  ...\\rangle$ to denote averaging procedure.  Quantitative measures of turbulence, in particular, turbulence spectrum, became important recently also due to advances in the theory of MHD turbulence. As we know, astrophysical fluids are magnetized, which makes one believe that the correspondence should exist between astrophysical turbulence and MHD models of the phenomenon (see Vazquez-Semadeni et al. 2000, Mac Low \\& Klessen 2004, Bellesteros-Paredes et al. 2007, McKee \\& Ostriker 2007 and ref. therein).   In fact, without observational testing, the application of theory of MHD turbulence to astrophysics could always be suspect. Indeed, from the point of view of fluid mechanics astrophysical turbulence is characterized by huge Reynolds numbers, $Re$, which is the inverse ratio of the eddy turnover time of a parcel of gas to the time required for viscous forces to slow it appreciably. For $Re\\gg 100$ we expect gas to be turbulent and this is exactly what we observe in HI (for HI $Re\\sim 10^8$). In fact, very high astrophysical $Re$ and its magnetic counterpart magnetic Reynolds number $Rm$ (that can be as high as $Rm\\sim 10^{16}$) present a big problem for numerical simulations that cannot possibly get even close to the astrophysically-motivated numbers. The currently available 3D simulations can have $Re$ and $Rm$ up to $\\sim 10^4$.  Both scale as the size of the box to the first power, while the computational effort increases as the fourth power (3 coordinates + time), so the brute force approach cannot begin to resolve the controversies related, for example, to ISM turbulence.  We expect that observational studies of turbulence velocity spectra will provide important insights into ISM physics. Even in the case of much more simple oceanic (essentially incompressible) turbulence, studies of spectra allowed to identify meaningful energy injection scales. In interstellar, intra-cluster medium, in addition to that, we expect to see variations of the spectral index arising from the variations of the degree of compressibility, magnetization, interaction of different interstellar phases etc.  How to get the turbulence spectra from observations is a problem of a long standing. While density fluctuations are readily available through both interstellar scincillations and studies of column density maps, the more coveted velocity spectra have been difficult to obtain reliably until very recently.  Turbulence is associated with fluctuating velocities that cause fluctuations in the Doppler shifts of emission and absorption lines. Observations provide integrals of either emissivities or opacities, both proportional to the local densities, at each velocity along the line of sight. It is far from trivial to determine the properties of the underlying turbulence from the observed spectral line.   Centroids of velocity (Munch 1958) have been an accepted way of studying turbulence,  although it was not clear to when and to what extend the measure really represents the velocity. Recent studies (Lazarian \\& Esquivel 2003, henceforth LE03, Esquivel \\& Lazarian 2005, Ossenkopf et al 2006, Esquivel et al. 2007) have showed that the centroids are not a good measure for supersonic turbulence, which means that while the results obtained for HII regions (O'Dell \\& Castaneda 1987) are probably OK, those for molecular clouds are unreliable.  An important progress in analytical description of the relation between the spectra of turbulent velocities and the observable spectra of fluctuations of spectral intensity was obtained in Lazarian \\& Pogosyan (2000, henceforth LP00). This description paved way to two new techniques, which were later termed Velocity Channel Analysis (VCA) and Velocity Coordinate Spectrum (VCS).  The techniques provide different ways of treating observational data in Position-Position-Velocity (PPV) data cubes. While VCA is based on the analysis of channel maps, which are the velocity slices of PPV cubes, the VCS analyses fluctuations along the velocity direction. If the slices have been used earlier for turbulence studies, although the relation between the spectrum of intensity fluctuations in the channel maps and the underlying turbulence spectrum was unknown, the analysis of the fluctuations along the velocity coordinate was initiated by the advent of the VCS theory.  With the VCA  and the VCS one can relate both observations and simulations to {\\it turbulence theory}. For instance, the aforementioned turbulence indexes are very informative, e.g. velocity indexes steeper than the Kolmogorov value of $-5/3$ are likely to reflect formation of shocks, while shallower indexes may reflect scale-dependent suppression of cascading (see Beresnyak \\& Lazarian 2006 and ref. therein).  By associating the variations of the index with different regions of ISM, e.g. with high or low star formation, one can get an important insight in the fundamental properties of ISM turbulence, its origin, evolution and dissipation.  The absorption of the emitted radiation was a concern of the observational studies of turbulence from the very start of the work in the field (see discussion in Munch 1999). A quantitative study of the effects of the absorption was performed for the VCA in Lazarian \\& Pogosyan (2004, henceforth LP04) and for the VCS in Lazarian \\& Pogosyan (2006, henceforth LP06). In LP06 it was stressed that absorption lines themselves can be used to study turbulence. Indeed, the VCS is a unique technique that does not require a spatial coverage to study fluctuations. Therefore individual point sources sampling turbulent absorbing medium can be used to get the underlying turbulent spectra.  However, LP06 discusses only the linear regime of absorption, i.e. when the absorption lines are not saturated. This substantially limits the applicability of the technique. For instance, for many optical and UV absorption lines, e.g. Mg II, SII, SiII the measured spectra show saturation. This means that a part of the wealth of the unique data obtained e.g. by HST and other instruments cannot be handled with the LP06 technique.  The goal of this paper is to improve this situation. In particular, in what follows, we develop a theoretical description that allows to relate the fluctuations of the absorption line profiles and the underlying velocity spectra in the saturated regime.  Below, in \\S 2 we describe the setting of the problem we address, while our main derivations are in \\S 3. The discussion of the new technique of turbulence study is provided in \\S 4, while the summary is in \\S 5. ",
        "conclusions": "\\subsection{Simplifying assumptions}  In the paper above we have discussed the application of VCS to strong absorption lines. The following assumption were used. First of all, considering the radiative transfer we neglected the effects of stimulated emission. This assumption is well satisfied for optical or UV absorption lines (see Spitzer 1979). Then, we assumed that the radiation is coming from a point source, which is an excellent approximation for the absorption of light of a star or a quasar. Moreover, we disregarded the variations of temperature in the medium.  Within our approach the last assumption may be most questionable. Indeed, it is known that the variations of temperature do affect absorption lines. Nevertheless, our present study, as well as our earlier studies, prove that the effects of the variations of density are limited. It is easy to see that the temperature variations can be combined together with the density ones to get effective renormalized ``density'' which effects we have already quantified.  Our formalism can also be generalized to include a more sophisticated radiative transfer and the spatial  extend of the radiation source. In the latter case we shall have to consider both the case of a narrow and a broad telescope beam, the way it has been done in LP06. Naturally, the expressions in LP06 for a broad beam observations can be straightforwardly applied to the absorption lines, substituting the optical depth variations instead of intensities. The advantage of the extended source is that not only VCS, but also VCA can be used (see Deshpande et al. 2000). As a disadvantage of an extended source is the steepening of the observed v-coordinate spectrum for studies of unresolved turbulence. This, for instance, may require employing even higher order structure functions, if one has to deal with windows arising from saturation of the absorption line.   In LP06 we have studied the VCS technique in the presence of absorption and formulated the criterion for the fluctuations of intensity to reliably reflect the fluctuations in turbulent velocities. In this paper, however, we used the logarithms of intensities and showed that this allows turbulence studies beyond the regime, at which fluctuations of intensity would be useful.  A similar approach, potentially, is applicable to the VCS for studies of emitted radiation in the presence of significant absorption. Indeed, the equation of the radiative transfer provides us in this case with \\begin{equation} I_{\\bf X}(v)= \\frac{\\epsilon}{\\alpha}\\left[1-{\\mathrm e}^{-\\alpha \\rho_s({\\bf X},v)}\\right]~~~, \\label{simplified} \\end{equation} where the first term is a constant, that can be subtracted. If this is done, the logarithm can be taken of the intensity. Potentially, this should allow VCS studies of emission lines where, otherwise, absorption distorts the statistics.  The difficulty of such an approach is the uncertainty of the base level of the signal. Taking logarithm is a non-linear operation that may distort the result, if the base level of signal is not accounted for properly. However, the advantage of the approach that potentially it allows studies of velocity turbulence, when the traditional VCA and VCS fail. Further research should clarify the utility of this approach.     \\subsection{Prospects of Studying ISM Turbulence with Absorption Lines}  The study of turbulence using the modified VCS technique above should be reliable for optical depth $\\tau$ up to $10^3$. For this range of optical depth, the line width is determined by Doppler shifts rather than the atomic constants. While formally the entire line profile provides information about the turbulence, in reality, the flat saturated part of the profile will contain only noise and will not be useful for any statistical study. Thus, the wings of the lines will contain signal.  As several absorption lines can be available along the same line of sight, this allows to extend the reliability of measurements combining them together. We believe that piecewise analyses of the wings belonging to different absorption lines is advantageous.  The actual data analysis may employ fitting the data with models, that, apart from the spectral index, specify the turbulence injection scale and velocity dispersion, as this is done in Chepurnov et al. (2006).  Note, that measurements of turbulence in the same volume using different absorption lines can provide complementary information. Formally, if lines with weak absorption, i.e. $\\tau_0<1$ are available, there is no need for other measurements. However, in the presence of inevitable noise, the situation may be far from trivial. Naturally, noise of a constant level, e.g. instrumental noise, will affect more weak absorption lines. The strong absorption lines, in terms of VCS sample turbulence only for sufficiently large $k_v$. This limits the range of turbulent scales that can be sampled with the technique. However, the contrast that is obtained with the strong absorption lines is higher, which provides an opportunity of increasing signal to noise ratio for the range of $k_v$ that is sampled by the absorption lines. If, however, a single strong  absorption line is used, an analogy with a two dish radio interferometer is appropriate. Every dish of the radio interferometer  samples spatial frequencies in the range approximately $[1/\\lambda, 1/d]$, where $\\lambda$ is the operational wavelength, $d$ is the diameter of the dish. In addition, the radio interferometer samples the spatial frequency $1/D$, where $D$ is the distance between the dishes. Similarly, a strong absorption line provides with the information on turbulent velocity at the largest spatial scale of the emitting objects, as well as the fluctuation corresponding to the scales $k_{window}, k_{abs}$.  In LP06 we concentrated on obtaining asymptotic regimes for studying turbulence. At the same time in Chepurnov et al. (2006) fitting models of turbulence to the data was attempted. In the latter approach non-power law observed spectra can be used, which is advantageous for actual data, for which the range of scales in $k_v$ is rather limited. Indeed, for HI with the injection velocities of 10 km/s and the thermal velocities of 1 km/s provides an order  of magnitude of effective \"inertial range\". Correcting for thermal velocities one can increase this range by a factor, which depends on the signal  to noise ratio of the data. Using heavier species rather than hydrogen one can increase the range by a factor $(m_{heavy}/m_{H})^{1/2}$. This may or may not be enough for observing good asymptotics.  We have seen in \\ref{} that for absorption lines the introduction of windows determined by the width of the line wings introduces additional distortions of the power spectrum. However, this is not a problem if, instead of asymptotics, fitting of the model is used. Compared to the models used in Chepurnov et al. (2006) the models for absorption lines should also have to model the window induced by the absorption. The advantage is, however, that absorption lines provide a pure pencil beam observations.   \\subsection{Comparison and synergy with other techniques}  Formally, there exists an extensive list of different tools to study turbulence that predated our studies (see Lazarian 1999 and ref. therein). However, a closer examination shows that this list is not as impressive as it looks. Moreover, our research showed that some techniques may provide confusing, if not erroneous, output, unless theoretical understanding of what they measure is achieved. For instance, we mentioned in the introduction an example of the erroneous application of velocity centroids to supersonic molecular cloud data. Note, that clumps and shell finding algorithms would find a hierarchy of clumps/shells for synthetic observations obtained with {\\it incompressible} simulations. This calls for a more cautious approach to the interpretation of the results of some of the accepted techniques.   For instance, the use of different wavelets for the analysis of data is frequently treated in the literature as different statistical techniques of turbulence studies (Gill \\& Henriksen 1990, Stutzki et al. 1998, Cambresy 1999, Khalil et al. 2006), which creates an illusion of an excessive wealth of tools and approaches. In reality, while Fourier transforms use harmonics of $e^{i{\\bf kr}}$, wavelets use more sophisticated basis functions, which may be more appropriate for problems at hand.  In our studies we also use wavelets both to analyze the results of computations (see Kowal \\& Lazarian 2006a) and synthetic maps (Ossenkopf et al. 2006, Esquivel et al. 2007), along with or instead of Fourier transforms or correlation functions. Wavelets may reduce the noise arising from inhomogeneity of data, but we found in the situations when correlation functions of centroids that we studied were failing as the Mach number was increasing, a popular wavelet ($\\Delta$-variance) was also failing (cp. Esquivel \\& Lazarian 2005, Ossenkopf et al. 2006, Esquivel et al. 2007).  While in wavelets the basis functions are fixed, a more sophisticated technique, Principal Component Analysis (PCA), chooses basis functions that are, in some sense, the most descriptive. Nevertheless, the empirical relations obtained with PCA for extracting velocity statistics provide, according to Padoan et al. (2006), an uncertainty of the velocity spectral index of the order $0.5$ (see also Brunt et al. 2003), which is too large for testing most of the turbulence theories. In addition, while our research in LP00 shows that for density spectra $E_{\\rho}\\sim k^{-\\alpha}$, for $\\alpha<1$ both velocity and density fluctuations influence the statistics of PPV cubes,  no dependencies of PPV statistics on density have been reported so far in PCA studies. This also may reflect the problem of finding the underlying relations empirically with data cubes of limited resolution. The latter provides a special kind of shot noise, which is discussed in a number of papers (Lazarian et al. 2001, Esquivel et al. 2003, Chepurnov \\& Lazarian 2006a).  {\\it Spectral Correlation Function (SCF)} (see Rosolowsky et al. 1999 for its original form) is another way to study turbulence. Further development of the SCF technique in Padoan et al. (2001) removed the adjustable parameters from the original expression for the SCF and made the technique rather similar to VCA in terms of the observational data analysis. Indeed, both SCF and VCA measure correlations of intensity in PPV ``slices'' (channel maps with a given velocity window $\\Delta v$), but if SCF treats the outcome empirically, the analytical relations in Lazarian \\& Pogosyan (2000) relate the VCA measures to the underlying velocity and density statistics. Mathematically, SCF contains additional square roots and normalizations compared to the VCA expressions. Those make the analytical treatment, which is possible for simpler VCA expressions, prohibitive. One might speculate that, similar to the case of conventional centroids and not normalized centroids introduced in Lazarian \\& Esquivel (2003), the actual difference between the statistics measured by the VCA and SCF is not significant.  In fact, we predicted several physically-motivated regimes for VCA studies. For instance, slices are ``thick'' for eddies with velocity ranges less than $\\Delta v$ and ``thin'' otherwise.  VCA relates the spectral index of intensity fluctuations within channel maps to the thickness of the velocity channel and to the underlying velocity and density in the emitting turbulent volume\\footnote{We showed that much of the earlier confusion stemmed from different observational groups having used velocity channels of different thicknesses (compare, e.g. , Green 1993 and Stanimirovic et al. 1999).}. In the VCA these variations of indexes with the thickness of PPV ``slice'' are used to disentangle velocity and density contributions. We suspect that similar ``thick\" and ``thin\" slice regimes should be present in the SCF analysis of data, but they have not been reported yet.  While the VCA can be used for all the purposes the SCF is used (e.g. for an empirical comparisons of simulations and observations), the opposite is not true. In fact, Padoan et al. (2004) stressed that VCA eliminates errors inevitable for empirical attempts to calibrate PPV fluctuations in terms of the underlying 3D velocity spectrum.  {\\it VCS} is a statistical tool that uses the information of fluctuations along the velocity axis of the PPV. Among all the tools that use spectral data, including the VCA, it is unique, as it {\\it does not} require spatial resolution. This is why, dealing with the absorption lines, where good spatial coverage is problematic, we employed the VCS. Potentially, having many sources sampling the object one can create PPV cubes and also apply the VCA technique. However, this requires very extended data sets, while for the VCS sampling with 5 or 10 sources can be sufficient for obtaining good statistics (Chepurnov \\& Lazarian 2006a).  We feel that dealing with the ISM turbulence, it is synergetic to combine different approaches. For the wavelets used their relation with the underlying Fourier spectrum is usually well defined. Therefore the formulation of the  theory (presented in this work, as well as, in our earlier papers in terms of the Fourier transforms) in terms of wavelets is straightforward. At the same time, the analysis of data with the wavelets may be advantageous, especially, in the situations when one has to deal with window functions."
    },
    "0801/0801.1198_arXiv.txt": {
        "abstract": "If the gas in filaments and halos shares the same velocity field than the luminous matter, it will generate measurable temperature anisotropies due to the Kinematic Sunyaev-Zeldovich effect. We compute the distribution function of the KSZ signal produced by a typical filament and show it is highly non-gaussian.  The combined contribution of the Thermal and Kinematic SZ effects of a filament of size $L\\simeq 5$Mpc and electron density $n_e\\simeq 10^3m^{-3}$ could explain the cold spots of $\\delta\\sim -200\\mu$K on scales of $30'$ found in the Corona Borealis Supercluster by the VSA experiment. PLANCK, with its large resolution and frequency coverage, could provide the first evidence of the existence of filaments in this region. The KSZ contribution of the network of filaments and halo structures to the radiation power spectrum peaks around $l\\sim 400$, a scale very different from that of clusters of galaxies, with a maximum amplitude $l(l+1)C_l/2\\pi\\sim 10-25 (\\mu K)^2$, depending on model parameters, i.e., $\\sigma_8$ and the Jeans length. About 80\\% of the signal comes from filaments with redshift $z\\le 0.1$. Adding this component to the intrinsic Cosmic Microwave Background temperature anisotropies of the concordance model improves the fit to WMAP 3yr data by $\\Delta\\chi^2\\simeq 1$. The improvement is not statistically significant but a more systematic study could demonstrate that gas could significantly contribute to the anisotropies measured by WMAP. ",
        "introduction": "At low redshift, about 50 \\% of all baryons have not been identified (Fukugita \\& Peebles, 2004). The highly ionized intergalactic gas, that has evolved from the initial density perturbations into a complex network of mildly non-linear structures in the redshift interval $0<z<6$ (Rauch 1998; Stocke, Shull \\& Penton 2004),  could contain most of the baryons in the Universe (Rauch et al. 1997, Schaye 2001, Richter et al 2005). With cosmic evolution, the fraction of baryons in these structures decreases as more matter is concentrated within compact virialized objects. The Ly$\\alpha$ forest absorbers at low redshifts are filaments (with low HI column densities) containing about 30\\%  of all baryons (Stocke et al. 2004). Hydrodynamical simulations predict that another large fraction of all baryons resides within mildly-nonlinear structures with temperatures $0.01-1$KeV, called  Warm-Hot Intergalactic Medium (WHIM). The amount of baryons within this WHIM is estimated to reach 20 - 40\\% (Cen \\& Ostriker 1999; Dave et al. 1999, 2001).  Being highly ionized the WHIM could generate temperature anisotropies by means of the thermal or kinematic Sunyaev-Zeldovich effect (TSZ, KSZ; Sunyaev \\& Zeldovich 1972, Sunyaev \\& Zeldovich 1980), adding to the contribution coming from gas in halos. Hallman et al (2007) carried out a numerical simulation and found that about a third of the SZ flux would be generated by unbound gas. In Atrio-Barandela \\& M\\\"ucket (2006) -hereafter ABM- we estimated the CMB anisotropies generated by the WHIM due to the TSZ effect. In this letter we compute the imprint on CMB temperature anisotropies due to the peculiar motions of gas in filaments and halos. The relative KSZ contribution with respect to the TSZ increases with decreasing temperature and the anisotropies generated could be used to trace the gas distribution on large scales.  Briefly, in Sec. 2 we show the effect of individual filaments on the CMB temperature anisotropies. We extend an early calculation of the KSZ probability distribution function due to halos by Yoshida, Sheth \\&  Diaferio (2001) by including gas not bound in halos and generalize the work of Juszkiewicz, Fisher \\& Szapudi (1998). In Sec. 3 we compute the radiation power spectrum, showing that it reaches its maximum at $l\\sim 400$, an angular scale much larger than that of clusters of galaxies; we also show that a statistical analysis including the anisotropies generated by the peculiar motions in the WHIM improves the fit to WMAP 3yr data.  Finally, we present our main conclusions. ",
        "conclusions": "If the 'missing baryons' are in the form of a highly ionized WHIM, and this medium shares the same velocity field than galaxies and dark matter, then it will produce a contribution to the CMB temperature anisotropies that could be detectable.  Gas in halos and filaments with bulk flows aligned with the $los$ could explain a significant fraction of the signal measured in the decrement H of the Corona Borealis supercluster observed by the VSA interferometer. Also, we computed the effect of the overall population of filaments on CMB temperature anisotropies. Their contribution has a caracteristic angular scale at $l\\sim 400$, very different from that of clusters of galaxies. Uncertainties in our model parameters translate into an uncertainty of a factor $5$ in the amplitude of the power spectrum.  Intrinsic and KSZ temperature anisotropies have the same frequency dependence, they can not be separated directly from the data.  Using a MCMC analysis we have shown that this extra component improves the fit of the concordance model to WMAP 3yr data, but the improvement can not be considered significant. WMAP 5yr data or the future PLANCK observations could have enough statistical power to indicate the existence of this component. Further, PLANCK has enough angular resolution and frequency coverage as to measure the KSZ contribution of structures such as decrement H in the VSA data of the Corona Borealis Supercluster."
    },
    "0801/0801.3230_arXiv.txt": {
        "abstract": "Most of the presently identified exoplanets have masses similar to that of Jupiter and therefore are assumed to be gaseous objects.  With the ever-increasing interest in discovering lower-mass planets, several of the so-called super-Earths (1 M$_\\earth$$<M<$10~M$_\\earth$), which are predicted to be rocky, have already been found. Here we report the possible discovery of a planet around the M-type star GJ~436 with a minimum mass of 4.7$\\pm$0.6~M$_\\earth$ and a true mass of $\\sim$5~M$_\\earth$, which would make it the least massive planet around a main-sequence star found to date. The planet is identified from its perturbations on an inner Neptune-mass transiting planet (GJ~436b), by pumping eccentricity and producing variations in the orbital inclination. Analysis of published radial velocity measurements indeed reveals a significant signal corresponding to an orbital period that is very close to the 2:1 mean motion resonance with the inner planet. The near-grazing nature of the transit makes it extremely sensitive to small changes in the inclination. ",
        "introduction": "Hundreds of exoplanets have been discovered over the past decade using a variety of techniques, notably precision radial velocities, most of them being Jupiter-like. While such planets are extremely useful to understand the morphology and evolution of planetary systems and star-planet connections, there is an obvious interest in eventually identifying an Earth-like object. Model calculations indicate that planets with masses in the interval 1~M$_{\\earth}$$<M<$10~M$_{\\earth}$ are of terrestrial type \\citep[e.g.,][]{IL04,Vea07}.  Such planets have often been dubbed ``super-Earths''. In this quest for smaller planets there have already been some successful detections of planets that fall in the super-Earth domain \\citep{Rea05,Bea06,Uea07}, and even their habitability may be evaluated \\citep{Sea07}.  Besides the direct detection from radial velocities or transits, methods for the discovery of low-mass planets as perturbers to other planets have been suggested (although without any successful detections yet). These are mostly based on monitoring variations in the timing of the transit over a relatively long time scale so that perturbations by planets as small as a few Earth masses could be detected \\citep{M02,S04,HM05,Aea05}.  In this {\\em Letter} we make use of yet another method that has allowed the possible detection of a super-Earth perturbing the transiting planet GJ~436b. In this case, it is not perturbations to the transit mid-time but to the overall orbital elements of the inner transiting object that cause variations on the transit shape and depth. We show that, for near-grazing transits, their duration is strongly sensitive to small changes in the orbital inclination. ",
        "conclusions": "The proposed existence of GJ~436c is based on three independent pieces of evidence: {\\em a)} An orbital inclination change supported by the lack of transit detection in 2004 and the grazing transit observed in 2007; {\\em b)} an inner planet with significant orbital eccentricity and a tidal dissipation timescale much smaller than the age of the system; and {\\em c)} a low-amplitude radial velocity signal that is fully consistent with a planet in the 2:1 mean motion resonance with the inner planet. A key element of our study is a dynamical investigation that explains the observed effects {\\em a)} and {\\em b)}, and predicts the existence of a perturbing planet with constraints on its mass and semi-major axis. Reanalysis of the radial velocity data indeed reveals such planet closely matching the predicted properties. In other words, the eccentricity of the orbit clearly reveals the existence of a perturbing planet, and a grazing transit will be most sensitive to even mild non-coplanarity between the two objects. We find such possible inclination change and we identify a radial velocity signal matching the predictions of our dynamical integrations.  The resulting planet system around GJ 436 strengthens the trend of a relatively large number of hot-Neptunes and super-Earths around stars of low mass \\citep{Bea07}. The presence of a long-term radial velocity trend of $\\sim$1~m~s$^{-1}$ per year could still be indicative of further planets in the system with wider orbits. Indeed, the system around GJ~436 shows striking resemblances to that around the M-type star Gl~581 \\citep{Uea07}, and thus its planets may experience changes in the orbital elements, perhaps eventually undergoing transits in spite of a previously null result \\citep{LMea06}. Our study provides yet another illustration of the variety of exoplanet systems and highlights the potential for complex dynamical histories that imply sizeable variations of the planets' orbital elements over timescales of decades.  The method of using near-grazing transits should be of special interest to discover small planets. Even for mild non-coplanarity, objects as small as a few M$_{\\earth}$ can lead to moderately high values of the induced eccentricity and to orbital inclination changes in the order of tenths of degrees per year. This will be especially effective for transit search missions from space, which could overcome the lower detection probability of near-grazing transits, and push the detection limits to even lower-mass objects than those responsible for the transit events."
    },
    "0801/0801.2236_arXiv.txt": {
        "abstract": " ",
        "introduction": "The disk wind plays a clue role in the process of removing the angular momentum excess from accretion disks (Blandford and Payne 1982 (BP82)), and works towards the accretion of the disk matter towards the star. A physical connection between accretion and outflow processes is confirmed by a decrease of the outflow activity with the age of stars, accompanying with a decrease of accretion rates (Calvet et al. 2000) and a frequency of accretion disks (see Andr\\'{e} et al. (2000); Mundy et al. (2000).)  Most of the modern models (see., e.g., review by Pudritz et al. (2007)) suggest that the disk wind does not contain the dust. Nevertheless, already in 1993 Safier (1993a) in his generalized version of the BP82 self-similar wind model argued that the weakly ionized wind arising from the accretion disk surface lifts the dust grains due to their collisions with the neutral atoms, i.e. the disk wind is dusty. According to Safier, the maximum size of the grain which are capable to lift with the wind is about of 1 mm.  Apparently, the presence of the dust in the disk wind has to affect the circumstellar (CS) extinction, the spectral energy distribution and polarization properties of the young stellar objects. Based on the results of Safier (1993a), in our previous papers (Grinin and Tambovtseva 2002; Grinin et al. 2004; Tambovtseva et al. (2006)) we consider photometric effects produced by the dusty disk winds in the young binaries. In the present paper we consider in more details than in the cited papers an interaction of the dust component of the wind with the radiation of the star as well with the hot gas for the single TTSs and HAEs. ",
        "conclusions": "Thus, we showed that in the case of T Tauri stars the disk wind can absorb and scatter the radiation of the star within a rather large segment of the solid angle 4$\\pi$ in the wide range of the accretion rates $\\dot{M_a}= 10^{-8} - 10^{-6}M_\\odot$ yr$^{-1}$. This means that the disk wind, in fact, is capable to play the same role as the puffed-up inner rim in the dust sublimation zone of the accretion disk (Natta et al. 2001). This inner rim screens the adjoining regions of the accretion disk from the direct stellar radiation (Dullemond et al. 2001), and under the certain inclination angles of the disk to the line of sight may be a source of the variable CS extinction (Dullemond et al. 2003). The dust component of the disk wind is able to produce the same effect.  In the case of Herbig Ae stars a screening effect produced by the dust component of the disk wind at the same range of accretion rates is significantly less than that for TTSs. Therefore, in HAEs the contribution of the disk wind into the thermal radiation of the CS dust can be comparable with that from the inner rim only at the high values of the accretion rate $\\geq 10^{-6}M_\\odot$ yr$^{-1}$.  It should be noted that we considered optical properties of the dust component of the wind using the model of Garcia et al (2001) with $\\xi = 0.01$. In the disk wind theory this important parameter governs an efficiency of the magneto centrifugal mechanism of the gas acceleration. In particular, a growth in $\\xi$ leads to an increase of the mass loss rate in the disk wind as well to a decrease of the terminal velocity of the wind (Garcia et al. 2001). Both effects works in the same direction: they increase the density of the matter in the wind. Therefore in the models with the large $\\xi$ the disk wind has to be more opaque on dust in comparison with the model considered above.  \\subsection{Structural disk wind and variable circumstellar extinction} Basing on the existing disk wind models we suggested that the wind possesses an axial symmetry and azimuthal homogeneity. In fact, this is a model simplification, and in reality, it seems hardly feasible. In conditions of supersonic turbulent motions the disk wind cannot be a continuous outflow in the hydrodynamical sense. It has to consist of an aggregate of gas and dust streams of the different power arising from the disk surface. In such a case the filling factor $q$ would be one of the wind parameters. It implies a fraction of the wind volume filled in with streams of the matter. Now we can only note that $q$ has to be less than unity.  Thus, under the real conditions, the column density of the dust along the line of sight passed through the disk wind, may be a complex function of time. It may fluctuate due to the motion of the gas and dust streams. Besides its changes can reveal quasi-periods caused by the repeated intersection of the line of sight by the same dominant wind stream. Note, that the quasi-periods in the brightness changes have been really observed in some UX Ori type stars (see e.g., Shevchenko et al. 1993). The rotation of the nonhomogeneous disk wind could be the reason of the spectral variability of some young stars (e.g. Kozlova et al. 2003).  Changes in the CS extinction may vary not only the radiation flux coming to the observer directly from the star, but the radiation flux from that region of the disk which is illuminated by the star through the disk wind. Shadows from the wind in this part of the CS disk have to move along the disk following the motion of the gas and dust streams. Since these streams looks like spinning-up spirals, their shadows projected on the disk have to be also spiral-like. Detection and investigation of such moving shadows on the images of the CS disks would be important for the theory of the disk winds.  {\\bf HH 30}. It is likely that namely such a mechanism of the variability is realized in the case of HH 30. Comparison of the images of this object obtained in the different time with the Hubble Space Telescope showed that they are variable. Both the type of the object's asymmetry  and the integral flux from it were variable (Burrows et al. 1996; Stapelfeldt et al. 1999, Wood et al. 2000). Wood and Whitney (1998) supposed that changes in conditions of illumination of the CS disk by the spotted rotating star could be the reason of the HH 30 variability. However, new data on the variability of the object (Watson and Stapelfeldt 2004, 2007) did not confirm the presence of the period connected with the rotation of the spotted star. According to these authors, variability of HH 30 has a more complex character and caused by changes in the CS extinction in the inner regions of the disk. A structural disk wind consisted of the separate gas and dust streams starting from the surface of the CS disk corresponds well to this role.  {\\bf RW Aur}. Another example of the young star whose variability is difficult to explain without appealing to a hypothesis about a dusty disk wind is the classical TTS RW Aur. This star relates to the most studied young stars. It is characterized with a large amplitude photometric activity (Herbst et al. 1994) and a complex type of variability of the emission line profiles and intensities (Petrov et al. 2001; Alencar et al. 2005). Recently Petrov and Kozak (2007) analyzed in detail a long-term series of the spectral and photometric observations of RW Aur and showed that there is a correlation in variations in the emission lines with the different excitation potential, which can be explained only if to assume that spectral variability is due to screening the emission region by the CS dust clouds. It is known from observations that the symmetry axis of the RW Aur's CS disk is inclined to the line of sight under 46 $\\pm 3^\\circ$ (this angle was derived very accurately with the help of the radial and space (a projection on the sky plane) velocities of the moving details in the optical jet (Lopez-Martin et al. 2003)). Under such an inclination disk cannot screen the star even if to take into account the rim in the sublimation zone. Therefore, an appearance of the dust on the line of sight (and hence, on the high latitudes in the star's coordinate system) Petrov and Kozak connected with the dust fragments of the disk wind.  An applicability of the theory of the dusty disk winds is not limited by examples given above. The calculations show (Grinin and Tambovtseva 2002; Tambovtseva et al. 2006) that the photometric effects caused by the dust component of the disk wind can be observed in the young binaries. In particular, obscuration by the extended disk win could cause abnormally long lasting eclipses observed in some binaries."
    },
    "0801/0801.4695_arXiv.txt": {
        "abstract": "The problem of the resolution of turbulent flows in adaptive mesh refinement (AMR) simulations is investigated by means of 3D hydrodynamical simulations in an idealised setup, representing a moving subcluster during a merger event. AMR simulations performed with the usual refinement criteria based on local gradients of selected variables do not properly resolve the production of turbulence downstream of the cluster. Therefore we apply novel AMR criteria which are optimised to follow the evolution of a turbulent flow. We demonstrate that these criteria provide a better resolution of the flow past the subcluster, allowing us to follow the onset of the shear instability, the evolution of the turbulent wake and the subsequent back-reaction on the subcluster core morphology. We discuss some implications for the modelling of cluster cold fronts. ",
        "introduction": "\\label{intro}  The role of turbulence in astrophysics and, in particular, in cosmic structure formation has been attracting increasing attention, along with observational and theoretical indications of its relevance. There are different observational hints converging towards the turbulent character of the intra-cluster medium (ICM) (\\citealt{sfm04,cfj04,ev06,rcb05,rcb06}, but see also \\citealt{fsc03}). From an observational point of view, the next generation of X-ray observatories will be able to observe the Doppler shifts of emission lines of heavy ions due to bulk motions in the ICM, thus providing information about the turbulent state of the flow. Numerical studies predict that gas bulk motions and turbulence in the ICM can be a consequence of cluster merging \\citep{r98,nb99,t00,rs01,dvb05}. In their model, \\citet{ssh06} identify three physical regimes for turbulence production and decay in clusters, the latest one being dominated by turbulent production in the wakes of minor mergers. This phase would play a key role also for the magnetic field amplification in the ICM.  Merging is a crucial process in cosmic structure formation. It is predicted in the framework of hierarchical clustering and can be studied theoretically by means of numerical simulations (e.g.~\\citealt{rbl96,rsm98,mbb07}). Observationally, modern X-ray observatories are able to detect substructures in the ICM that are indicative of merging \\citep{snb03}. Among the first results of galaxy cluster observations with {\\it Chandra} was the discovery of ``cold fronts'' in the merging clusters A2142 and A3667, and later in many other clusters (see \\citealt{mv07} for a review). The origin of cold fronts in merger clusters has been ascribed to infalling substructures, although the cold fronts in the cores of cooling core clusters have a different interpretation \\citep{th05,am06}.  Cosmological simulations of galaxy cluster formation have confirmed the merging scenario for the transient formation of cold fronts \\citep{bem02,nk03}. The problem has also been addressed with simplified setups, which allow a better control over the involved physical parameters, in 2D \\citep{asp03,hcf03,afm04,xcd07} and 3D simulations \\citep{t05a,t05b,afm05,afm07,dp07}.  The theoretical study of turbulence in the framework of full cosmological simulations is a challenging task and presents the typical problems of numerical simulations of strongly clumped media. Adaptive mesh refinement (AMR) is a viable tool for saving computational resources and modelling the large dynamic range in a proper way \\citep{n04}. The choice of the most suitable mesh refinement criterion is thus a delicate compromise between following the flow structure in the most accurate way and exploiting the advantage of saving computational time and memory.  This study (Paper I) is the first part of a broader project on the physics of galaxy clusters, focused on investigating the generation of turbulence in the ICM by means of AMR simulations. In Paper I, we present new AMR criteria that are more suitable for resolving the production of turbulence than those commonly used in cosmological simulations. They were tested in a simple setup, reminiscent of a wind tube experiment, but designed in order to be representative of an idealised subcluster merger. This test problem is a useful and controlled testbed for the new tools, far from the complexity of a more realistic cosmological simulation. Nevertheless, it is similar to previous studies of subcluster mergers and can provide useful insights about the physics of the minor mergers, relevant for the generation of turbulence in the ICM. In contrast with earlier work, we focus our analysis on the turbulent wake of the subcluster rather than on the cold front, showing the difference of the evolution of the Kelvin-Helmholtz instability (KHI) resulting from the use of different AMR criteria.  The second step, described in \\citet{in08} (hereafter Paper II), is to apply the same tools to simulations of structure formation and to study the turbulent flow in the ICM. The results of the present work will thus be compared with the outcomes of more realistic simulations of minor mergers.  The paper is structured as follows: the new refinement criteria are defined in in Sec.~\\ref{simulations}. The numerical setup of the performed simulations and the features of the KHI are introduced in Sects.~\\ref{setup} and \\ref{theory}, respectively. The results are presented in Sec.~\\ref{results} and discussed in Sec.~\\ref{discussion}. ",
        "conclusions": "\\label{discussion}  We presented an interesting case of new refinement techniques, developed for resolving turbulent flows in AMR hydrodynamical simulations, which may have useful applications in astrophysical problems. The AMR is not only a tool used for this study, but has itself been investigated and tested in its capabilities. New refinement criteria based on the regional variability of control variables of the flow were introduced and shown to be superior for simulations of turbulent media.%  As a first application, we presented a set of simulations of an idealised substructure in a wind in order to study a subcluster merger in a simplified fashion. In this setup, the novel AMR criteria are suitable for the study of the turbulent wake, which is better resolved than using refinement criteria based on local gradients of selected hydrodynamical variables. The optimal control of the refining procedure ensures high resolution where desired, as pointed out by the comparison with an equivalent static grid simulation.  The overall evolution of the subcluster % is similar to the hydrodynamical simulations of \\citet{asp03}, \\citet{t05a} and \\citet{t05b}, where the development of shear instability leads to the formation of a turbulent wake. The typical rms velocity in the wake is of the order of $500\\ \\rmn{km\\ s^{-1}}$, similar to the above cited simulations and to the theoretical predictions \\citep{ssh06}. This value is obtained only in runs using the new AMR criteria and tends to the result from an equivalent static grid simulation, whereas in the reference run $A$ the inferred turbulent velocity is smaller by a factor of 2.  We investigated the proper refining of the turbulent wake with several numerical experiments. The choice of the AMR criteria has an impact both on the turbulent velocity and on the morphology of the subcluster core as a result of the back-flow. We claim that this issue should be carefully considered in numerical simulations of subcluster mergers, as for instance in studies of cold fronts. If the problem is addressed with an inadequate numerical setup (as in run $A$), the turbulent back-flow will be underestimated significantly. In extreme cases (run $B$, or Sect.~\\ref{zeus}), the simulation might fail to reproduce the flow downstream of the subcluster. We further showed that refining turbulence in the wake does not depend only on AMR but, as expected, also on the hydrodynamic solver. A detailed code comparison in this setup, like the one performed for the blob test by \\citet{ams06}, would certainly be useful.  As already stated, we do not intend to discuss the features of the cold front which forms ahead of the subcluster because our simulations do not include the magnetic field evolution and/or heat conduction \\citep{afm04,afm05,afm07,xcd07}. The observed cold front temperature profiles (the best studied example is A3667) call for the suppression of the thermal conduction at the interface between the cold subcluster gas and the hot ICM background \\citep{mv07}. This suppression, in turn, would be naturally explained by a magnetic flux tangential to the cold front surface. \\citet{dp07} show, with MHD simulations using AMR, that the flow past a moving bullet generates vorticity, even if a magnetic field stabilises its surface against the KHI. The presence of MHD turbulence and further amplification of the magnetic field is also suggested. Despite the simplified approach of both setups, we infer that the generation of a turbulent wake and a proper modelling of this phase are therefore crucial {\\em also} for the subcluster MHD simulations, and particular care in modelling the turbulent flow should be taken also in this class of simulations (with AMR or not). Only a clear assessment of the capability of a code in dealing with turbulent flows can allow to distinguish, for example, between the instability suppression caused by a magnetic field, and an unreliable modelling of the KHI onset.  Our simplified approach to the minor merger case shares the limitations of previous, similar simulations cited above. However, it has the advantage of presenting a well controlled setup, which would be difficult to study in detail in the framework of a full cosmological simulation. In view of the results it is necessary to note that, in the performed simulations, the merging subcluster was embedded in a background with a very simple velocity field. In a less idealised flow (cf.~\\citealt{nb99,dvb05}), it is admittedly unlikely that the turbulent tail would extend for several subcluster radii without being mixed. Nonetheless, it would not prevent from testing the turbulent character of the flow immediately downstream of the subcluster \\citep{t05b}, both observationally and in cosmological simulations.  From an observational point of view, the turbulence level of the subcluster wake can be an interesting test for supporting the studied scenario. The rms velocity behind the subcluster can be used as an observable for turbulence for future X-ray spectrometers. This test would pertain to the wider, long-standing problem of detecting turbulence in galaxy clusters \\citep[][and reference therein, for an overview]{sc06}.  Although significant efforts are undertaken for modelling turbulent flows in SPH simulations \\citep{dvb05,vtc06}, grid-based AMR simulations are a very powerful approach to simulations in this field. The AMR criteria used in this work will be further applied to astrophysical problems, the next step being the simulation of galaxy cluster formation and evolution (cf.~Paper II). According to the analysis of \\citet{ssh06}, besides the minor merger phase there are several other physical regimes when turbulence evolves and can play an important role in the energy budget of a galaxy cluster. The present study corroborates the importance of the minor mergers in the production of turbulence in galaxy clusters. Moreover, it provides new tools for the theoretical study of this problem, which will be continued in Paper II. The final goal of our investigation is a consistent AMR modelling of turbulent flows based on a subgrid scale model approach (cf.~\\citealt{snh06,snhr06}). This will be the topic of forthcoming work."
    },
    "0801/0801.2285_arXiv.txt": {
        "abstract": "{Robotic telescopes and grid technology have made significant progress in recent years. Both innovations offer important advantages over conventional technologies, particularly in combination with one another. Here, we introduce robotic telescopes used by the Astrophysical Institute Potsdam as ideal instruments for building a robotic telescope network. We also discuss the grid architecture and protocols facilitating the network integration that is being developed by the German AstroGrid-D project. Finally, we present three user interfaces employed for this purpose.} ",
        "introduction": "In recent years the first robotic telescopes came to operation and many of them already routinely collect data without human interaction. A natural next step in automated astronomy is the integration of these telescopes into a network. A network of distributed telescopes can perform new types of observations which cannot be accomplished with individual instruments. Important examples include continuous long-term monitoring of objects independent of day and night cycles and weather, rapid response to transient objects and multiwavelength campaigns.  The Astrophysical Institute Potsdam (AIP) currently operates five robotic telescopes and is pursuing their integration into a network. This effort is supported by AstroGrid-D, which is the German astronomy community grid. Here {\\it grid} refers to a distributed network of loosely coupled resources with a common, user-friendly infrastructure. The AstroGrid-D collaboration consists of 20 people from fourteen German institutes under the leadership of the AIP. The project has been funded by the German Ministry of Education and Research (BMBF) for a period of three years. The network will be built on grid technology, which provides an ideal framework for a robotic telescope network. For example, it provides solutions for the management of Virtual Organization\\footnote{In grid computing, a Virtual Organization defines a collaboration of users with same access rights to grid resources.}, grid resources, computational jobs and observation, data and metadata. In addition, it allows the immediate access to computational and storage resources for data analysis.  \\phantom{ }  \\phantom{ } ",
        "conclusions": "Five robotic telescopes are operated by the AIP and are prepared for network integration. An architecture for the integration of these telescopes to a network using grid technology has been developed within AstroGrid-D and its main components are available. RTML metadata of telescopes and observations can be provided to, and received from, the information service. User interfaces are available for easy access to this information."
    },
    "0801/0801.1678_arXiv.txt": {
        "abstract": "We have used extensive libraries of model and empirical galaxy spectra (assembled respectively from the population synthesis code of Bruzual and Charlot and the fourth data release of the Sloan Digital Sky Survey) to interpret some puzzling features seen in the spectra of high redshift star-forming galaxies. We show that a stellar He\\,{\\sc ii}\\,$\\lambda 1640$ emission line, produced in the expanding atmospheres of Of and Wolf-Rayet stars, should be detectable with an equivalent width of $0.5-1.5$\\,\\AA\\ in the integrated spectra of star-forming galaxies, provided the metallicity is greater than about half solar. Our models reproduce the strength of the He\\,{\\sc ii}\\,$\\lambda 1640$ line measured in the spectra of Lyman break galaxies for established values of their metallicities. With better empirical calibrations in local galaxies, this spectral feature has the potential of becoming a useful diagnostic of massive star winds at high, as well as low, redshifts.   We also uncover a relationship in SDSS galaxies between their location in the [O\\,{\\sc iii}]/H$\\beta$ vs. [N\\,{\\sc ii}]/H$\\alpha$ diagnostic diagram (the BPT diagram) and their excess specific star formation rate relative to galaxies of similar mass. We infer that an elevated ionisation parameter $U$ is at the root of this effect, and propose that this is also the cause of the offset of high redshift star-forming galaxies in the BPT diagram compared to local ones. We further speculate that higher electron densities and escape fractions of hydrogen ionising photons may be the factors responsible for the systematically higher values of $U$ in the H\\,{\\sc ii} regions of high redshift galaxies. The impact of such differences on abundance determinations from strong nebular lines are considered and found to be relatively minor. ",
        "introduction": "\\label{sec:introduction}  \\begin{table*} \\begin{minipage}[c]{0.6\\textwidth} \\centering \\caption{Adopted emission-line luminosities of individual Wolf-Rayet stars in units of $10^{35}$ erg s$^{-1}$} \\begin{tabular}{@{}lccccc} \\hline \\hline Line name & \\multicolumn{2}{c}{WNE stars} & \\multicolumn{2}{c}{WNL stars} & \\multicolumn{1}{c}{OIf stars}\\\\ & $Z<0.2 Z_\\odot$ & $Z \\ge 0.2 Z_\\odot$& $Z<0.2 Z_\\odot$ & $Z \\ge 0.2 Z_\\odot$ \\\\ \\hline He\\,{\\sc ii}\\,$\\lambda 4686$ &  1.7$^{a}$  & 8.4 &   4.3    &  24.7 & 1.3 \\\\ He\\,{\\sc ii}\\,$\\lambda 1640$ & 17    & 84  &   43     & 247 & 20 \\\\ N\\,{\\sc iii}\\,$\\lambda 4640$ &  0    & 0      & 4.0      & 6.3  & 0 \\\\ N\\,{\\sc v}\\,$\\lambda 4603-4620$   & 0.24$^{a}$ & 1.60     & 0 & 0 & 0 \\\\ \\hline \\end{tabular}  \\medskip $^a$ Includes binaries, see Crowther \\& Hadfield (2006). \\label{tab:WR_line_lum} \\end{minipage} \\end{table*}  The spectra of galaxies, particularly at optical and ultraviolet (UV) wavelengths, contain a wealth of information on their stellar populations and interstellar media. The veritable explosion in the size of galaxy surveys which we have witnessed in recent years has been accompanied by the development of increasingly sophisticated models for the analysis and interpretation of their spectra (e.g. Kewley et al. 2001a; Bruzual \\& Charlot 2003; Heavens et al. 2004; Ocvirk et al. 2006). The same analysis tools are also increasingly being applied to the spectra of galaxies at redshifts $z > 1$, even though their detailed characteristics may well differ from those of their counterparts at lower redshifts. Considering, for example, star-forming galaxies, it now seems well established that at high redshifts a larger fraction of the star formation activity took place in conditions similar to those encountered today in the relatively rare `luminous infrared galaxies' (e.g. Reddy et al. 2006 and references therein), and that the transition from starburst dominated to a more quiescent mode of star formation occurred at redshifts between $z \\simeq 2$ and 1 (e.g. Papovich et al. 2005).  It is thus reasonable to question to what extent diagnostic tools developed to interpret local galaxies may need to be modified in order to decipher correctly the clues encoded in the spectra of high redshift galaxies. In this paper we use state-of-the-art spectral synthesis models and observations of local galaxies to investigate two apparent `anomalies' which have been noted in the spectra of high redshift star-forming galaxies. The first is the common detection of the He\\,{\\sc ii}\\,$\\lambda 1640$ emission line with a width which is suggestive of an origin in the expanding atmospheres of luminous stars rather than in H\\,{\\sc ii} regions. This spectral feature is clearly seen in the composite rest-frame UV spectrum of more than 800 Lyman break galaxies at $z \\simeq 3$ constructed by Shapley et al. (2003). As those authors point out, its strength is underestimated by continuous star formation models generated by the Starburst99 code (Leitherer et al. 1999), particularly at sub-solar metallicities. Locally, the He\\,{\\sc ii}\\,$\\lambda 1640$ emission line is prominent in the UV spectra of starburst galaxies, but only during a brief period when the number of Wolf-Rayet (W-R) stars is at a maximum (e.g. Chandar, Leitherer \\& Tremonti 2004), so that its strength in the composite spectrum of hundreds of galaxies which span a wide range of ages, from $10^7$ to $10^9$ years (Shapley et al. 2001; Erb et al. 2006b), is at first sight puzzling. Here we reassess this `problem' using the most recent models and observations of Wolf-Rayet stars in galaxies of differing metallicities, and taking into account a range of star formation histories.  The second issue concerns the indication that star-forming galaxies at $z \\simgt 1$ may exhibit systematically higher ratios of collisionally excited to recombination lines than  present-day  H\\,{\\sc ii} regions (Shapley et al. 2005; Erb et al 2006a; Kriek et al. 2007). Although the data samples are still small, most of the galaxies at $z \\simgt 1$ where the relevant emission lines have been measured, are offset---relative to lower $z$ galaxies---in the [O\\,{\\sc iii}]$\\lambda 5007$/H$\\beta$ vs. [N\\,{\\sc ii}]/H$\\alpha$ diagram first proposed by Baldwin, Phillips \\& Terlevich (1981, BPT) as a way of classifying emission line regions and their sources of ionisation/excitation. A variety of different physical processes may be at the root of such differences, as discussed by Shapley et al. (2005, see also Liu et al 2008); we investigate this question further here with detailed modelling and comparison with the large body of data on local galaxies provided by the Sloan Digital Sky Survey (SDSS).  Where relevant, we have adopted today's consensus cosmology with density parameters $\\Omega_{\\rm M}=0.3$, $\\Omega_\\Lambda=0.7$ and a Hubble parameter $H_0=70$\\,km~s$^{-1}$~Mpc$^{-1}$. We have also adopted throughout the stellar initial mass function (IMF) of Chabrier (2003), unless otherwise specified. ",
        "conclusions": ""
    },
    "0801/0801.1352_arXiv.txt": {
        "abstract": "The central star of this nebula has an observed intense magnetic field and the fast wind is no longer present, indicating that a back flow process has probably developed. Long-slit, spatially resolved echelle spectra have been obtained across the main body of NGC 1360 and over its system of bipolar jets. Deep images of the knotty structures of the jets have also been obtained. The data allow a detailed study of the structure and kinematics of this object and the results are modeled considering the effects of a magnetic collimation process in the development of the nebula and then switching off the fast stellar wind to follow its evolution to its current state. The model is able to successfully reproduce many key features of NGC 1360 under these premises. ",
        "introduction": "NGC 1360 is an interesting planetary nebula for several reasons. Firstly, this is one of the few planetary nebulae where an intense stellar magnetic field, of the order of kilogauss, has been detected in its core \\citep{jor05}. Secondly, the nebula shows an elongated and nearly featureless morphology with no apparent bright rim, sharp inner boundary or central cavity, indicating the current absence of a significant stellar wind, this is further confirmed by archive IUE spectra of the core of NGC 1360 that reveal the absence of P Cygni-type line profiles. Furthermore, the nebula has a system of fast-expanding bipolar jets outside the main nebular body.  Although NGC 1360 had been classified by \\citet{men77} as containing a close binary nucleus, \\citet{weh79} were unable to confirm such claim after a detailed radial velocity study of the central star. More recently, \\citet{gol04} have described NGC 1360 as a thick shell without a sharp inner edge that indicates a lack of ongoing compression by a fast stellar wind. They also have brought attention to the presence of the bipolar jets and studied the kinematics of the shell and the northern jet. The central region of NGC 1360 shows no [N~II] $\\lambda$~6548, 6584 emission lines, these lines are only detected from the strings of knots that form the system of bipolar collimated outflows and from an isolated knot located on the northern edge of the nebula in our data. Since collimated, bipolar outflows in planetary nebulae are thought to originate either from the action of binary nuclei  or  magnetic fields and given the lack of evidence of the former and the high significance of the observations of the latter in this case, NGC 1360 represents a unique case to test the possible effects of a magnetic collimation process in the development of this object. Furthermore, recently \\citet{gar06} have started exploring the expected structure of planetary nebulae once the fast wind has faded, and NGC 1360 represents also an opportunity to examine the characteristics of this stage.  Relevant data for NGC 1360 are its high electron temperature T$_e$ [O~III] = 16500 K   and low electron density N$_e$ $\\lesssim$ 100 cm$^{-3}$ \\citep{kal90}; high effective temperature T$_{eff}$ = 97 000 K, log L = 3.64 L$_\\odot$, M = 0.65 M$_\\odot$, main sequence mass M$_{ms}$= 2.7 M$_\\odot$ \\citep{tra05, jor05}.  A distance estimate determined from  {\\it Hipparcos} data \\citep{ack98} d =${350} {\\pm} {1000 \\atop 180}$ pc is adopted here. ",
        "conclusions": "NGC 1360 is a large planetary nebula with a mild expansion velocity, broad line components and a system of bipolar, fast, collimated outflows that increase in velocity with distance. The fast stellar wind from the central star has died away at least a few thousand years ago and a back filling process has modified its structure producing a smooth, nearly featureless and elongated high excitation nebula. Past attempts to identify a binary core in this object have yielded negative results. The central star conserves a strong magnetic field  indicating that a magnetized stellar wind could have been responsible for axial collimation in the development of this planetary nebula. We have therefore studied the history of this nebula considering a process of magnetic collimation and then switching off the stellar wind and letting it evolve to its present time. The resulting model is able to successfully reproduce many of the observed key features in NGC 1360."
    },
    "0801/0801.3482_arXiv.txt": {
        "abstract": "The next generation of X-ray telescopes have the potential to detect faint quasars at very high redshift and probe the early growth of massive black holes (BHs). We present modelling of the evolution of the optical and X-ray AGN luminosity function at $2<z<6$ based on a CDM merger-driven model for the triggering of nuclear activity combined with a variety of fading laws. We extrapolate the merger-driven models to $z>6$ for a range of BH growth scenarios. We predict significant numbers of sources at $z \\sim 6$ with fluxes just an order of magnitude below the current detection limits and thus detectable with XEUS and Constellation-X, relatively independently of the fading law chosen.  The predicted number of sources at even higher redshift depends sensitively on the early growth history of BHs.  For passive evolution models in which BHs grow constantly at their Eddington limit, detectable BHs may be rare beyond $z \\sim 10$ even with Generation-X. However, in the more probable scenario that BH growth at $z>6$ can be described by passive evolution with a small duty cycle, or by our merger driven accretion model, then we predict that XEUS and Generation-X will detect significant numbers of black holes out to $z \\sim 10$ and perhaps beyond. ",
        "introduction": "The next generation of X-ray satellites promise to be powerful probes of quasar activity out to very high redshifts. The European lead project XEUS\\footnote{http://www.rssd.esa.int/index.php?project=XEUS} is designed to reach a sensitivity $\\sim 100$ better than the deepest observations to date. Constellation-X\\footnote {http://constellation.gsfc.nasa.gov/} would offer at least an order of magnitude improvement in sensitivity, hopefully paving the way for a mission like Generation-X\\footnote {htt://www.cfa.harvard.edu/hea/genx.html} which could probe 50 times fainter again.  In this paper we aim to assess the prospects for probing the early growth of supermassive black holes (SMBHs) at $z \\gtrsim 6$ with these telescopes or similar future X-ray missions [see, e.g., \\citet{haiman1999} and \\citet{wyithe2003} for earlier work on this]. The observed correlations between SMBH mass and galaxy properties, which are now well established \\citep[e.g.][]{gebhardt2000, ferrarese2000}, strongly suggest a tight link between the build-up of the stellar mass in galactic bulges and the mass of the central BH. Most models of the growth history of SMBHs assume that the frequent merging of galaxies predicted by CDM-like hierarchical models of galaxy formation plays an important role in this process \\citep[e.g.][]{kh2000, wyithe2003, volonteri2003, croton2006, malbon2006}. Such models, although subject to considerable degeneracies, have been reasonably successful in reproducing the observed luminosity function of AGN as well as the inferred BH mass function and their respective evolution with redshift. These models generally struggle, however, to combine the very efficient high redshift growth of BHs, suggested by the SMBHs detected at $z \\sim 6$, with the strong feedback required to reproduce the rapid decline in the fuelling rate of the most massive BHs at low redshift \\citep{bromley2004,malbon2006}.  Furthermore, at low redshift there is considerable observational evidence that some, and maybe even most, of the fuelling of SMBHs is not connected to major mergers between galaxies \\citep[e.g.][]{kauffmann2003,li2007}. This may be related to the fact that merger-driven models of quasar activity have difficulty reproducing the luminosity dependent density evolution (LDDE) observed for faint AGN at redshifts below $z \\sim 3$ \\citep[e.g.][]{ueda2003,hasinger2005,silverman2007,bongiorno2007} [see, e.g., \\citet{marulli2007}], although more sophisticated models which follow the simultaneous growth of galaxies and AGN can incorporate feedback in such a way to orchestrate this effect \\citep{fontanot2006}.   We explore predictions for a variety of models for the early growth of SMBHs which range from merger driven accretion to continuous Eddington-limited accretion.  The BH growth scenario is crucial in determining the number of active BHs that we expect to detect as faint X-ray sources. However, as X-ray observations push to fainter flux limits they are expected to discover a mixture of low-luminosity objects at low to intermediate redshift, as well as distant objects previously at undetectable redshifts. Predictions for the redshift distribution of faint sources, which encrypt the growth history of SMBHs, therefore require a simultaneous study of the faint end of the luminosity function at all redshifts.  We have adopted a hybrid approach to modelling the redshift evolution of the X-ray emission associated with the fuelling of SMBHs. At low redshift ($z \\lesssim 2$) we use the observed X-ray luminosity functions, which are now well established down to very faint luminosities. At intermediate redshifts the faint end of the X-ray luminosity function is still subject to some uncertainty due to the difficultly of assigning redshifts to faint sources \\citep{aird2008}. At $2<z<6$ we have therefore adopted a CDM merger-driven model for the evolution of the emission from SMBHs, constrained by the available optical and X-ray data.  The efficiency of BH formation, the quasar lifetime and accretion rate, and the quasar spectral energy distribution (SED) are all free parameters in such a model.  The relevant assumptions should obviously be guided by our empirical and physical understanding of quasars, but there is still much flexibility, particularly near the observational limits in redshift and luminosity.   Observationally, faint AGN are more likely to display signatures of obscuration by gas and/or dust than brighter sources \\citep{ueda2003,simpson2005,treister2006,maiolino2007}. Most semi-analytic models of quasar activity do not model obscuration effects in sufficient detail to be able to explain these trends. \\citet{hopkins2005b,hopkins2005d,hopkins2006b,hopkins2006c} have, however, recently  presented a model for the quasar luminosity function (QLF) at $z<5$, which includes prescriptions for quasar luminosity dependent fading and absorption, to successfully reconcile the optical and X-ray QLFs.  We follow a very similar approach and adopt many of the same assumptions. However, whereas \\citet{hopkins2006b} extract the quasar formation rate from the constraints provided by the observed QLF, we use a cosmological galaxy merger rate from the extended Press-Schechter formalism \\citep[e.g.][]{lacey1993}. This will allow us to extend our model to the redshifts ($z>6$) we are predominantly concerned with in a well-motivated way. Semi-analytic models are an efficient way of exploring the wide range of physically plausible evolutionary scenarios for the QLF. As we will discuss in some detail, the faint end of the luminosity functions, and therefore the redshift distribution of faint X-ray sources, is very sensitive to the assumed time dependence of the fading of the emission in different wave-bands during galaxy mergers.  We utilise the flexibility of our model to study the constraints on this fading rate provided by the optical, soft and hard X-ray QLFs, and the unresolved Cosmic X-ray background (CXRB).  We then go on to explore several models for BH growth at $z>6$ which are consistent with the observed data at $0<z<6$ and span a wide range of possibilities for the rate of BH growth at high redshift.   Our paper is structured as follows. In Section 2 we briefly review the constraints on the optical and X-ray QLFs. We describe our merger driven model for $0<z<6$ in Section~3, and the high redshift extensions in Section~4.3. In the main body of Section~4 we calibrate our model against the measured optical and X-ray QLFs. In Section~5 we examine the consistency of our models with the CXRB and observed ${\\rm logN}-{\\rm logS}$ X-ray source distribution, and present our prediction for the source density detectable with future missions, before concluding in Section~6.  Throughout this paper we adopt a cosmological matter density $\\Omega_m = 0.27$, baryonic matter density $\\Omega_b = 0.044$, cosmological constant $\\Omega_{\\Lambda} = 0.73$, present day Hubble constant $H_o \\equiv  100h$~km~s$^{-1}$~Mpc$^{-1} = 71$km~s$^{-1}$~Mpc$^{-1}$, a mass variance on scales of $8 h^{-1}$~Mpc $\\sigma_8 = 0.84$ and a scale invariant primordial power spectrum (slope $n=1$). ",
        "conclusions": "We have presented a hybrid model for the redshift evolution of the X-ray emission connected with the fuelling of supermassive black holes with the aim of assessing the prospects of detecting quasars via their X-ray emission at $z>6$. At $z \\lesssim 2$ we have used the observed X-ray luminosity functions, at $2<z<6$ we have adopted a CDM merger-driven model for the evolution of the emission from supermassive black holes combined with a slow and a rapid fading law and we have explored a range of assumptions for the growth of supermassive black holes at $z>6$.  For appropriate choices for the efficiency of black hole formation in a dark matter halo and the characteristic quasar lifetime, our model is in good agreement with the observed optical and soft and hard X-ray quasar luminosity functions at $0<z<6$ for both the slow and the rapid fading law. The only disagreement that occurs is for the faint end of the X-ray luminosity functions where our merger model combined with the slow fading model suggested by \\citet{hopkins2005b} based on detailed numerical simulations predicts too many faint objects. It consequently also overpredicts the soft X-ray background suggesting that the slow fading model of \\citet{hopkins2005b} cannot extend to black hole masses much below $10^{7}$~M$_{\\sun}$.  With a more rapid fading law our models are well within the limits of the integrated CXRB, leaving room for the population of star-forming galaxies expected at low flux densities, and the population of Compton-thick sources needed to explain the overall spectral shape of the CXRB.  The main aim of our study is to assess the prospects of planned and proposed X-ray missions like Constellation-X, XEUS and Generation-X to study the build-up of black holes at $z>6$. With a point source sensitivity of $3\\times 10^{-17}$~ergs~s$^{-1}$~cm$^{-2}$ for a 1~Ms exposure, Constellation-X will detect significant numbers of black holes at $z\\sim 5-6$ but will probably not yet reach the necessary sensitivity to detect significant numbers of sources at $z>6$. This conclusion holds independently of our assumptions for the growth of BHs at $z>6$.  For XEUS with its anticipated ten times superior point source sensitivity, the prospects for detecting sources at $z>6$ is much more favourable. Up to 17 per cent of the approximately 100 sources in the 49 square arcminute FOV expected for the rapid fading model with a minimum active black hole mass of $\\sim 3 \\times 10^4$~M$_{\\odot}$ are expected to be at $z>6$. If our merger driven model can be extrapolated to $z>6$ or if black holes grew with a duty cycle of about 10~per~cent, XEUS would detect significant numbers of black holes with a rather flat redshift distribution out to $z\\sim 10$.  In either case, observable X-ray emission would have to accompany the growth of supermassive black holes starting from seed black holes of $\\ga 10^{4}$~M$_{\\sun}$ or smaller.  For the more remote prospect of Generation-X, with an anticipated point source sensitivity yet a factor 100 better again, the flat redshift distribution could extent to $z\\sim 15$ and beyond.  If black holes grew much faster at $z>6$ as in our Eddington limited growth model with a duty fraction of unity, which would be required for Eddington limited growth of the most massive black holes from stellar mass seed black holes, neither XEUS nor Generation-X would detect many black holes at $z \\gtrsim 10$. However, while such rapid growth may occur and may indeed be necessary for the most massive black holes at $z \\sim 6$, it appears unlikely that this will be the norm for the majority of black holes.  The prospects that XEUS (and Generation-X) will unravel the growth history of black holes at $z>6$ and beyond are thus excellent.   \\begin{figure} \\vspace*{110mm} \\special{psfile=plot11.ps hscale=90 vscale=90 hoffset=-30 voffset=-380 angle=0} \\caption{Observed $0.5-2$~keV flux from a BH with the \\citet{hopkins2006c} SED shining at its Eddington limit as a function of source redshift. BH masses as labelled. } \\end{figure}"
    },
    "0801/0801.1955_arXiv.txt": {
        "abstract": "We introduce the contour process to describe the geometrical properties of merger trees. The contour process produces a one-dimensional object, the contour walk, which is a translation of the merger tree. We portray the contour walk through its length and action. The length is proportional to to the number of progenitors in the tree, and the action can be interpreted as a proxy of the mean length of a branch in a merger tree.  We obtain the contour walk for merger trees extracted from the public database of the Millennium Run and also for merger trees constructed with a public Monte-Carlo code which implements a Markovian algorithm. The trees correspond to halos of final masses between $10^{11} h^{-1}$ M$_\\odot$ and $10^{14} h^{-1}$ M$_\\odot$. We study how the length and action of the walks evolve with the mass of the final halo. In all the cases, except for the action measured from Markovian trees, we find a transitional scale around $3 \\times 10^{12} h^{-1}$ M$_\\odot$. As a general trend the length and action measured from the Markovian trees show a large scatter in comparison with the case of the Millennium Run trees. ",
        "introduction": "The current paradigm of large scale structure formation in the Universe is hierarchical, meaning that structures are formed by merger aggregation.  The dominant driver for this merger dynamic is thought to be the non-baryonic and non-collisional cold dark matter. Therefore, research in large scale structure formation is principally done through the study of dark matter structures, which on large scales are thought to serve as scaffolding for baryonic structures \\citep{LSS_SFW}.  The dark matter merger process has been usually represented by a tree, keeping the analogy with the genealogical tree of an individual. Great effort has been put into the methods of construction of merger trees, as they are a way to understand dark matter aggregation and are a necessary input for codes of semi-analytic galaxy formation. The methods range from the analytical approach of Monte-Carlo methods (\\cite{2007IJMPD..16..763Z} for a review) passing through hybrid approaches that mix numerical realizations of a density field with analytical approximations for its evolution \\citep{MORGANA} to the fully numerical approach that identifies the dark matter halos from different snapshots in a $N$-body simulation to construct the merger history\\citep{galicsI}.  Usually, the validity of the trees constructed in the analytical and hybrid way is stated from comparisons with trees constructed from numerical simulations \\citep{1999MNRAS.305..946S,2007arXiv0708.1382P,2007arXiv0708.1599N}. Unfortunately, the quantities used to compare trees from two approaches usually sacrifice the complexity inherent to the tree structure in the sake of simple tests. The most common simplification is to select one branch of the tree (the most massive) to make the analysis. Another approach, measures the abundance of structures of a given mass among the halos in all the merger tree. In all these cases the geometrical information is suppressed, mostly because of the lack of simple structures to describe that kind of information.  In this paper we present a new way, in the astrophysical context, to translate the geometrical information from a merger tree into a 1-dimensional structure. The translation is based on the encoding of the tree information into its \\emph{contour walk}.  We apply this description to the merger trees from a large dark matter numerical simulation, the Millennium Run \\citep{Millennium}. We use its public database to select halos in different bin masses at redshift $z=0$ to extract its merger histories and build the contour walk. We analyze each tree in terms of two simple statistics extracted from these walks.  We have also performed this kind of analysis on merger trees obtained with the algorithm described by \\cite{2007arXiv0708.1599N}, using the source code they kindly made public. With this code we have constructed trees with two resolutions for the minimum mass halo, one mimicking the Millennium Run, and the other with nearly $8$ times lower resolution.   This paper is divided as follows. In Section 2 we explain the construction of a contour walk from a merger tree. In Section 3 we present the simulation that produced the public data we used in this paper and the Monte Carlo code for the construction of Markovian merger trees. In Section 4 we perform an immediate implementation of these concepts to the available merger trees. We calculate global statistics from the walks, and discuss its possible physical meaning. In the last section we discuss our results and suggest how the tool we have proposed can be used to tackle more complex questions about merger trees. ",
        "conclusions": "We introduced from the mathematical literature the contour process of a tree, and we applied this concept to the description of merger trees. We used a large dark matter simulation (the Millennium Run) and a Monte-Carlo code (implementing a Markovian approach) to obtain merger trees  in these two approximations. Furthermore, the Markovian trees were obtained  with two different values for the minimal mass of a parent halo. One  resolution mimicked exactly the MR, and the other resolution had a $8$  times more massive minimal halo mass. We refer to these Markovian runs as the high resolution and low resolution runs respectively.   We extracted simple statistics from these walk: the length $N$ (proportional to the total number of halos in the tree) and the action $S$ (which can be loosely interpreted as the mean longitude of a branch). We report our results of walk length and action emphasizing its evolution as a function of the physical halo mass,  and not its absolute values.  From the length and the action, we found in the Millennium Run a transitional mass scale at $\\sim 3.0\\times 10^{12} h^{-1}$ M$_{\\odot}$. In the case of the walk length, $N$, this transitional mass-scale marks the change between a dependence $N\\propto M^{0.5}$ for low halo masses and $N\\propto M^{0.8}$ for high mass halos, where $M$ is the mass of the final dark matter halo in the tree. With the action, the scale marks the highest value for the action as function of the halo mass.    For the Markovian trees the dependence of the walk length on the halo mass is almost the same everywhere, although much steeper than the dependence found in the Millennium Run. Nonetheless, we found the same transitional scale from the length statistics for the two resolutions. It does not prove that the transitional scale is independent on the resolution mass (in fact, it should be dependent) but suggests that the scale does not have a strong dependence with the minimal mass resolution used to describe the merger tree. The evolution of the action as a function of halo mass is completely different from the MR case. The Markovian run with mimicking MR resolution does not show a sharp transitional scale, instead it shows a large plateau ranging for two orders of magnitude in mass.  Perhaps the biggest difference between the two ways of constructing the trees is that the Markovian approach always show a bigger dispersion on its contour walk statistics, which seem best defined in the Millennium Run, judging from its low dispersion. This could be interpreted in fact as a higher geometrical variability in the Markovian trees.    We have shown how simple statistics from the contour walk can give a new handle on the description of merger tree, exploring the geometrical information of merger trees. Even if the length of the walk could have been obtained without an intermediating contour process, the approach of extracting information from higher order statistics, as was the case for the normalized action $S$, can be extended in complexity. For instance, is relevant to point out that the information of the mass in each node is also encoded in the contour walk. The mass information is used as an ordering criteria to walk the tree. One could define sections defined by $\\{\\tau_{i}|x(\\tau_{i}=t_{1})\\}$, i.e. the points where the walk touches the next to last snapshot, and the \\emph{length ratio} of the sections delimited by these points would include information of the \\emph{mass ratio} of the mergers at time $t_{1}$. In general the study of crossing paths, defined as the the section of the contour walks within some boundary $a<\\{x_{i}\\}<b$ might produce useful statistics to further classify the complex merger trees of massive halos."
    },
    "0801/0801.1487_arXiv.txt": {
        "abstract": "LS 5039 is one of a handful of X-ray binaries that have been recently detected at high-energy $\\gamma$-rays, in this case, by the High-Energy Stereoscopy Array (H.E.S.S.). The nature of this system is unknown: both a black hole and a pulsar have been invoked as possible compact object companions. Here we work with a model of the high energy phenomenology of the system  in which it is assumed that the companion object is a pulsar rotating around an O6.5V star in the $\\sim 3.9$ days orbit. The model assumes two different sets of power-law spectral parameters of the interacting primary leptons corresponding to the two orbital phase intervals defined by H.E.S.S. as having different gamma-ray spectra and very-high-energy (VHE) cutoffs.  We show the H.E.S.S. phenomenology is completely explained by this model. We present predictions for photons with lower energies (for $E>1 $ GeV), subject to test in the forthcoming months with the GLAST satellite. We find that GLAST is able to judge on this model within one year. ",
        "introduction": "The Gamma-ray Large Area Space Telescope (GLAST) is the next generation  $\\gamma$-ray observatory due for launch in May 2008. Its primary instrument is the Large Area Telescope (LAT), which will measure $\\gamma$-ray flux and spectra from 100 MeV to $\\sim$100 GeV. LAT is the successor of the EGRET experiment that flew onboard the Compton satellite. LAT will have better angular resolution, greater effective area, wider field of view, and broader energy coverage than EGRET. Developing predictions for the GLAST energy domain upon which to test models of LS 5039 is then timely and worth exploring.  In the previous work (Sierpowska-Bartosik \\& Torres 2007, where we refer the reader for  references), under the assumption that LS 5039 is composed by a pulsar rotating around an O6.5V star in the $\\sim 3.9$ days orbit, we presented the results of a theoretical modeling of the high energy phenomenology observed by  H.E.S.S.\\footnote{H.E.S.S. found a periodicity in the $\\gamma$-ray flux --consistent with the orbital timescale as determined by Casares et al. (2005)-- and fast variability displayed on top of this periodic behavior, both in flux and spectrum (Aharonian et al. 2006).} The main difference between the previous model and the current one is that the former assumed a constant spectrum describing the interacting particles along the orbit, whereas now it is let to vary. The previous model (including detailed account of the system geometry, Klein-Nishina inverse Compton (IC), $\\gamma \\gamma$ absorption, and cascading)  was able to describe reasonably well the rich details found in the system, both flux and spectrum-wise. However, almost the complete spectral energy distribution at superior conjunction, as well as other features of  the observed orbital variation of the H.E.S.S. energy spectra,  remained to be consistently explained. With models that do not entirely match the observed data at the highest energies, the study of lower energy predictions lacks testing power (strictly speaking, these models are already ruled out by H.E.S.S. observations). Here, we analyze and motivate the description of the VHE data allowing for a variable spectrum of interacting primary leptons along the orbit and indicate that the new model can nicely explain the H.E.S.S. observations.  We then present predictions for the flux and spectrum as function of phase for photons with lower energies, where H.E.S.S. is not sensitive but GLAST is. LS 5039 is near the galactic center and has two other $\\gamma$-ray sources nearby, as well as a significant galactic diffuse background: it is not an easy target for GLAST (see Dubois 2006). In order to study orbital variability, the need to integrate long time intervals and make harder energy cuts to bring the signal above background was reported. ",
        "conclusions": "A detailed modeling of LS 5039, under the assumption that it contains a pulsar, is able to fully describe the challenging H.E.S.S.-observed phenomenology found in the system. Observed lightcurve, spectra (Fig.  1), and short-timescale spectral variability (Fig.  2) are matched with a variable spectrum of primary particles along the orbit, where two phase intervals are considered. It is interesting to see that the influence of the orbital inclination angle is minor along all energy intervals within our theoretical model. Predictions of the model at lower $\\gamma$-ray energy domains are ready for testing with GLAST. We find that GLAST is able to judge on this model within one year. Features both at the lightcurve and spectral level can be recognized. Additional tests will come at higher energy yet beyond the current reach, e.g., behavior of the spectra at shorter phase intervals to be compared with data from the next generation of ground-based instruments: H.E.S.S. II, and the planned Cerenkov Telescope Array (CTA)."
    },
    "0801/0801.1164_arXiv.txt": {
        "abstract": "We present K-band polarimetric images of several massive young stellar objects at resolutions $\\sim$ 0.1-0.5 arcsec. The polarization vectors around these sources are nearly centro-symmetric, indicating they are dominating the illumination of each field. Three out of the four sources show elongated low-polarization structures passing through the centers, suggesting the presence of polarization disks. These structures and their surrounding reflection nebulae make up bipolar outflow/disk systems, supporting the collapse/accretion scenario as their low-mass siblings. In particular, S140 IRS1 show well defined outflow cavity walls and a polarization disk which matches the direction of previously observed equatorial disk wind, thus confirming the polarization disk is actually the circumstellar disk. To date, a dozen massive protostellar objects show evidence for the existence of disks; our work add additional samples around MYSOs equivalent to early B-type stars. ",
        "introduction": "With the wide availability of sub-arcsec observations, the presence of disks around massive young stellar objects (MYSOs) has become a hot topic recently. The reason for this, is the dispute about the way how massive stars are formed. That massive stars are formed in a scaled-up version of how low-mass stars form \\citep{shu87} would be a natural thought. However, it has been proposed that, if the mass of the central star exceeds 8 \\msun, the tremendous radiation pressure from the central star will halt the mass accretion and keep it from growing \\citep[e.g.,][]{kahn74,pal93}. Such a consideration leads to an entirely different mechanism, i.e., massive stars may be formed through mergers of lower-mass protostars \\citep{bon98}. On the other hand, a number of solutions have been proposed to account for the radiation pressure problem, such as non-isotropic accretion\\citep{york02}, different dust opacity \\citep{kahn74}, and redirection of radiation by disk/outflow system\\citep{kru05}. The presence of disk/outflow system around MYSOs will provide key evidence to evaluate these two kind of scenarios.  Disks around MYSOs have been detected over a wide range of wavelengths with different tracers \\citep{pat05,jiang05,ces06,bel06}. These results suggest that up to early B type stars \\textbf{can} be formed through a similar route as solar type stars. In this work we report sub-arcsec near-infrared (NIR) polarimetric imaging observations which show three polarization disks that can be best interpreted as disks or toriods around the MYSOs. ",
        "conclusions": "There are several possible causes that make a low-polarization lane passing through the bright source. In addition to a flattened structure around the illuminating source, a magnetic field may also produce a similar morphology by dichroic extinction. We checked the polarization of point sources in our observed fields, and find no large-scale alignment of polarization vectors in the direction of PDs. Another possibility is that, when a nearly spherical envelop illuminated by a source inside and a much brighter source outside, the envelop will show a bipolar polarization structure. However, this happens when the source outside is hundreds to thousands of times brighter than the source inside, so that the photons from the two sources shining on the envelop are roughly equal. It is not possible in our case since our sources are brighter than 10$^3$ \\lsun. It is also possible that a foreground filament which is just in front of the sources dilutes the polarization, but it is of very low probability and is in fact impossible to happen for all three targets. Therefore, taking the morphologies of the observed reflection nebulae into consideration, the most likely interpretation of the PDs here would be real disks or toriods. We note however, since the polarimetric observations do not give velocity information, it is not possible to distinguish between a Keplerian disk and a toriod structure. A confirmation of the conclusion have to await sensitive molecular line observations.   Some parameters of the sources are listed in Table 1.  Of the four sources, three (S140 IRS1, S255 IRS1 and I23033) show evidence for the presence of a disk or toroid.  All of them are deeply embedded (A$_V$ $>$ 15 for IRAS23033+5951, $>$20 for S255 IRS1 and $>$30 for S140 and NGC7538, Fig. 2a). They also show very large near infrared color excesses. Their near infrared colors suggest the evolutionary status could be comparable to low-mass Class I objects. Except for S255 IRS1, whose spectral type is still controversial \\citep{hey89,how97,ito01}, these sources are MYSOs with masses ranging between 10 to 30 solar masses according to their spectral types. A J vs J-H plot would also suggest a similar conclusion about spectral types of these objects (Fig. 2b).   Up to date, a dozen of MYSOs show evidence of the existence of disks \\citep{ces06,zin07}. These candidate disks are found in various evolutionary status of their host objects, which are detected from (sub)millimeter wavelengths the NIR. However, the candidate disks happen around MYSOs with their luminosities less that 10$^5$ \\lsun, typical of early B-type main-sequence stars \\citep{ces06}, suggesting that disks might be common occurrence for this kind of MYSOs; we still lack of evidence whether still higher-mass YSOs host circumstellar disks. Our observations add samples to this kind. And it is interesting to note that our most massive sample, NGC7538 IRS1 ($\\sim$ 30 \\msun), does not show strong evidence of PD.  We have also shown that high-resolution polarimetric imaging can provide a sensitive tool to detect circumstellar disks around MYSOs, even if the disks are too faint to be detectable directly. Our future work will concentrate on searching for polarization disks around MYSOs of higher masses and different evolutionary status. We aim to set up a sample for future observations with higher resolution and sensitivity. This kind of work will give crucial hints to answer the question at what mass and evolutionary phase the disks around MYSOs are truncated."
    },
    "0801/0801.2484_arXiv.txt": {
        "abstract": "Observations of Type Ia supernovae used to map the expansion history of the Universe suffer from systematic uncertainties that need to be propagated into the estimates of cosmological parameters. We propose an iterative Monte-Carlo simulation and cosmology fitting technique (SMOCK) to investigate the impact of sources of error upon fits of the dark energy equation of state. This approach is especially useful to track the impact of non-Gaussian, correlated effects, e.g.~reddening correction errors, brightness evolution of the supernovae, K-corrections, gravitational lensing, etc. While the tool is primarily aimed for studies and optimization of future instruments, we use the ``Gold'' data-set in  Riess et al.~(2007) to show examples of potential systematic uncertainties that could exceed the quoted statistical uncertainties. ",
        "introduction": "The direct evidence for dark energy, which were obtained through observations of Type Ia supernovae (SN~Ia) in the 1990s, marked the beginning of a new era in observational cosmology~\\cite{goo95,garnavich98,riess98,schmidt98,perl99}. Significantly improved SN~Ia data sets have since then been reported in e.g.~\\cite{knop03,tonry03,barris04,riess04}. These data sets have been followed by the impressive results from the ongoing dedicated SN~Ia surveys: SNLS~\\cite{astier06} and ESSENCE~\\cite{wood}. An important recent compilation of published SNe~Ia, including e.g.~the supernovae from the first year of SNLS, but also adding several high redshift ($z \\gtrsim 1$) SNe~Ia from the GOODS survey, is reported in \\cite{riess06}.  As the SN~Ia sample sizes keep growing at an increased rate, the relative importance of statistical uncertainties is steadily decreasing, which brings systematic uncertainties to focus.  It is thus of key importance to estimate the impact of systematic uncertainties as we enter the era of ``precision cosmology''.  In this work we propose a Monte Carlo simulation approach to quantify the propagated effect of known (or suspected) systematic effects in supernova cosmology. The emphasis here is to quantify the degradation of the precision to estimate cosmological parameters due to systematic errors, possibly correlated in arbitrary groupings of supernovae: in redshift bins, by locations in the sky, observational instruments, time of observations, etc.   The current study includes investigations of sources of error from: uncertainties in the extinction by dust, K-corrections, redshift uncertainties, gravitational lensing, Malmquist bias, possible brightness evolution of the ``calibrated'' standard candle (e.g. metallicity effects), spectral template differences, and miss-calibrations between the low and the high redshift supernova sample. Systematic effects caused by bulk motion have been specifically targeted in a number of recent studies, e.g.~\\cite{neill07}, and were not specifically included in this study. However, some general systematic effects, e.g. drift of estimated peak brightness at low-z, would mimic the main effect expected from peculiar velocities.  The proposed method, implemented in the SMOCK package and presented in~\\S\\ref{sec:method}, allows us to quantify the sensitivity of cosmological results, obtained from real or simulated data, to specific (multiple) systematic effects. For example, we can quantify at which level a systematic effect gives rise to a bias exceeding the statistical uncertainty. The tool could hence be used to define requirements of any future mission aimed at improving our knowledge of the dark sector using SNe~Ia. As an example of how this technique can be used we have performed a dedicated study of the ``Gold'' sample presented in Reference~\\cite{riess06}. This study demonstrates how cosmological parameter fits are affected by a number of different systematic errors.  In~\\S\\ref{sec:snmag} the model of the Universe which we have used and the cosmological parameters are introduced. Our methodology for studying the effects of systematic uncertainties is described in~\\S\\ref{sec:method}. In~\\S\\ref{sec:sysquant} we show how the changes in the cosmological results due to systematic uncertainties can be quantified and the trends parametrized. A number of different systematic effects, their implementation, and their effects on the Gold set are the topics of~\\S\\ref{sec:syseff}. Some observations regarding systematic effects and the Gold set are presented in~\\S\\ref{sec:disc}. The paper is concluded by some implications for current and future SN~Ia surveys in~\\S\\ref{sec:concl}.  Throughout the paper units are used where the speed of light is unity, i.e.~$c=1$. ",
        "conclusions": "\\label{sec:concl} SMOCK, a statistical Monte Carlo simulation approach towards quantification of systematic effects is presented. This method can give concrete answers to how sensitive SN~Ia samples are to a wide variety of different systematic effects. We can estimate \\emph{both} systematic and statistical uncertainties in a time efficient way.   Current SN~Ia data-sets are likely to be influenced by unknown systematics, which quite possibly already are larger than the statistical errors.  In order for future SN~Ia surveys to be successful, the propagated effects into the cosmolgical fits must be understood. With the methods presented here this will be possible through a process where we first define acceptable systematic uncertainties to achieve some target confidence level in e.g. the impact on the derived dark energy properties. The SMOCK technique can be used to further define the instrumental requirements necessary to detect and possibly correct for any putative source of systematic uncertainty.  Studies of a range of different systematics were performed on the Gold data-set, demonstrating the possibility of quantifying sensitivity to systematic effects. The importance of excluding the effects of a possible $R_V$ or color bias and obtaining a good spectral template were highlighted. For increased precision future surveys will need to demonstrate that maximal systematic color effects are at the $0.01$ magnitude level and that host galaxy dust properties do not have any systematic shifts over the $\\Delta R_V = 1$ level (or that possible such shifts are under control)."
    },
    "0801/0801.0953_arXiv.txt": {
        "abstract": "During the last few years a class of enigmatic sub-luminous accreting neutron stars has been found in our Galaxy. They have peak X-ray luminosities (2--10 keV) of a few times $10^{34}$ erg s$^{-1}$ to a few times $10^{35}$ erg s$^{-1}$, and both persistent and transient sources have been found. I present a short overview of our knowledge of these systems and what we can learn from them. ",
        "introduction": "Instrumental limitations (e.g., sensitivity, spatial resolution) of past X-ray instruments restricted the study of accreting neutron stars in low-mass X-ray binaries to the relatively bright systems, with 2--10 keV X-ray luminosities of $10^{36-38}$ erg s$^{-1}$. The existence of fainter sources was well known, but until recently they could not be studied in detail. Luckily, this is not true for the current X-ray satellites and a sizable number ($\\sim$15) of sub-luminous (X-ray luminosities of $10^{34-35}$ erg s$^{-1}$) accreting neutron stars have been found in the last few years (called sub-luminous NS LMXBs; I will only consider the systems in which the companion is suspected to have a low mass [$<$1 $M_\\odot$] but there exist also a population of systems which have a suspected high-mass companion). Their neutron star nature has been confirmed by the detection of type-I X-ray bursts.  Two of them are persistent sources, while the others are transients with a large variety in outburst durations and recurrence times.\\   \\vspace{0.2cm}  \\noindent {\\bf The persistent sub-luminous NS LMXBs:} Currently, two such sources are known: 1RXS J171824.2--402934 \\citep{1999A&A...349..389V,2000A&A...358L..71K} was shown to be persistent with a 0.5--10 keV X-ray luminosity of only $\\sim 5\\times 10^{34}$ erg s$^{-1}$ by \\cite{2005A&A...440..287I}. After AX J1754.2--2754 \\citep{2002ApJS..138...19S, 2007ATel.1094....1C} exhibited a type-I X-ray burst, \\cite{2007ATel.1143....1D} detected it with Swift at a 2--10 keV luminosity of a few times $10^{34}$ erg s$^{-1}$ establishing it too as a persistent sub-luminous NS LMXB.\\  \\vspace{0.2cm}  \\noindent {\\bf The quasi-persistent sources:} Recently, two quasi-persistent (outburst durations of years) sub-luminous NS LMXBs have been identified. The best example is XMMU J174716.1--281048 \\citep{2003ATel..147....1S}. \\cite{2006ATel..970....1B}, using INTEGRAL, reported on a new type-I X-ray burster near the Galactic center which was very soon identified with XMMU J174716.1--281048 \\citep{2006ATel..972....1W}. After analyzing the burst properties and several XMM-Newton observations, \\cite{2007A&A...468L..17D} identified the source as a quasi-persistent sub-luminous NS LMXB with a luminosity of $\\sim 5 \\times 10^{34}$ erg s$^{-1}$; a classification which was later confirmed by a series of Swift and Chandra observation of the source in 2007 \\citep{2007ATel.1078....1D,2007ATel.1136....1D,2007ATel.1174....1S}.   \\cite{1996PASJ...48..417M} first discovered the second system, AX J1745.6--2901, as a bursting transient near Sgr A*, with a luminosity of a few times $10^{35}$ erg s$^{-1}$ and with deep eclipses every $\\sim$8.4 hours (the orbital period). They suggested that AX J1745.6--2901 could be the old, bright transient A 1742--289 (albeit this time in a faint outburst), but \\cite{1996PASJ...48L.117K} could not find eclipses in the Ariel V data of A 1742--289 indicating that AX J1745.6--2901 is a new transient sub-luminous NS LMXB. Quasi-daily observations of the Galactic center with Swift in February 2006 yielded a new transient, Swift J174535.5--290135.6 \\citep{2006ATel..753....1K}, which XMM-Newton conclusively identified with AX J1745.6--2901 \\citep[showing eclipses with the right period;][]{2007ATel.1058....1P}.  The source turned of after staying on for several months \\citep{deegwijnands}, but it was detected in outburst again in February 2007 \\citep{2007ATel.1005....1K,2007ATel.1006....1W}.  The Swift quasi-daily monitoring campaign continued also in 2007 showing that the source remained active throughout 2007 \\citep{deegwijnands}.\\   \\vspace{0.2cm} \\noindent {\\bf The transient sources:} In addition, about a dozen, short-duration, transient sub-luminous NS LMXBs are known \\citep[see, e.g.,][]{2006A&A...449.1117W}. In general, they have peak outburst luminosities (2--10 keV) of $10^{34-35}$ erg s$^{-1}$ and their outbursts last from a few days to at most a few weeks. The combination of their very-faint peak luminosity and their transient nature makes them difficult to discover and study, but dedicated programs \\citep[e.g.,][]{2006A&A...449.1117W} are increasing the number of these systems known and our understanding of them. ",
        "conclusions": ""
    },
    "0801/0801.4950_arXiv.txt": {
        "abstract": "% I discuss different theories of massive star formation: formation from massive cores, competitive Bondi-Hoyle accretion, and protostellar collisions. I summarize basic features of the Turbulent Core Model (TCM). I then introduce the Orion Kleinmann-Low (KL) region, embedded in the Orion Nebula Cluster (ONC) and one of the nearest regions of massive star formation. The KL region contains three principal radio sources, known as ``{\\it I}'', ``{\\it n}'' and ``BN''. BN is known to be a runaway star, almost certainly set in motion by dynamical ejection within the ONC from a multiple system of massive stars, that would leave behind a recoiling, hard, massive, probably eccentric binary. I review the debate about whether this binary is $\\Theta^1C$, the most massive star in the ONC, or source {\\it I}, and argue that it is most likely to be $\\Theta^1C$, since this is now known be a recoiling, hard, massive, eccentric binary, with properties that satisfy the energy and momentum constraints implied by BN's motion. Source {\\it n} is a relatively low-mass protostar with extended radio emission suggestive of a bipolar outflow. Source {\\it I}, located near the center of the main gas concentration in the region, the Orion Hot Core, is the likely location of a massive protostar that is powering the KL region, and I discuss how its basic properties are consistent with predictions from the TCM. In this scenario, the radio emission from source {\\it I} is the base of a bipolar outflow that is ionized by the massive protostar and should be elongated along the axis of the outflow. ",
        "introduction": "Understanding the formation of massive stars is an important problem for many areas of astrophysics including high redshift Population III star formation, galaxy formation and evolution, galactic center environments and supermassive black hole formation, star and star cluster formation, and planet formation around stars in clusters, which may be relevant to our own solar system. The problem is challenging because of the wide range of spatial and temporal scales, the complicated interplay of gravity, thermal pressure, magnetic fields, radiation and ``turbulent'' (i.e. nonthermal) motions, including bipolar outflows from surrounding stars, and the uncertain initial conditions, caused by high obscuration to the typically distant and crowded regions where massive stars form. Thus progress requires close testing of theoretical models against detailed observational data, such as is available for the nearest massive star-forming regions.  I adopt the following definitions (see also McKee \\& Tan 2003; Beuther et al. 2007): a {\\it protostar} is a star (i.e. a near-equilibrium gaseous configuration in which gravity is balanced by thermal and radiation pressure and possibly rotation) accreting matter so that its mass, $m_*$, has not yet reached its maximum value, $m_{*f}$, that it is born with; a {\\it massive protostar} (MP) or star has a mass $\\geq 8 M_\\odot$; a {\\it pre-(massive-)stellar core} is the self-gravitating (negative total energy), topologically-connected, matter surrounding the location where the (eventually massive) protostar first forms (at time $t=0$); at later times when the protostar and rotationally-supported disk have formed, this system is a {\\it (massive-)star-forming core}, which includes the mass of the protostar and disk (mass may join or leave the core during the growth of the protostar); in principle the evolution of the pre-stellar core at earlier times ($t<0$) could be followed, defining its center to be at the minimum of the gravitational potential; a {\\it star-forming clump} or {\\it protocluster} is the system made up of a forming cluster of stars and the self-gravitating gas surrounding them. The minimum number of stars, which may be protostars, to define a cluster, ${\\cal N}_{\\rm *,min}$, can be debated: I suggest ${\\cal N}_{\\rm *,min} = 10 \\gg 1$. For clumps, self-gravity should first be assessed for all the mass in the minimum convex volume that includes the ${\\cal N}_*$ stars, then surrounding stars and gas (including effects of external pressure) can be assessed to see if they are bound to, and therefore a part of, this structure.  Note, since the mass of a protostar when it first forms is expected to be very small ($\\ll M_\\odot$), massive protostars must have gone through a stage when they had low and intermediate masses. Beuther et al. (2007) refer to these as ``Low to Intermediate Mass Protostars destined to become Massive Protostars'' and here we suggest the abbreviation LIMP-MP, so the expected evolutionary sequence if massive star formation is a scaled-up version of low-mass star formation is Pre-Massive-Stellar Core $\\rightarrow$ Massive-Star-Forming Core containing LIMP-MP $\\rightarrow$ Massive-Star-Forming Core containing MP $\\rightarrow$ Massive Star. Much of the theoretical debate in the field centers on whether the pre-massive-stellar core is itself {\\it massive} (i.e. with mass comparable to or larger than $m_{*f}$), whether the massive-star-forming core containing the LIMP-MP is {\\it massive}, and whether there is a stage where the mass of the massive-star-forming core containing the MP is gas dominated. The alternative is that the core has a small mass when the protostar first forms and then accumulates the bulk of its mass during the star-forming stage while maintaining a relatively small gas mass fraction.  Star-forming clumps, which are typically supersonically turbulent, form new cores and protostars by ``gravitational fragmentation'' (but note that the material in the new cores will still be part of the clump, so the clump itself has not strictly speaking fragmented). New cores and protostars may also form inside existing star-forming cores. I distinguish ``disk fragmentation'' from ``turbulent fragmentation''. In disk fragmentation, which only occurs inside star-forming cores, the new core (and protostar) is in a low eccentricity orbit about the original protostar, i.e. is gravitationally bound to it with infall mostly resisted by rotational support, and is likely to have its boundaries set by tidal forces from the original protostar. The original star-forming core maintains its identity, but now with a binary and circumbinary disk at its center. Turbulent fragmentation (e.g. Padoan \\& Nordlund 2002) can occur both inside and outside of star-forming cores. If it takes place inside a star-forming core, the new core will be bound to the previous core (it is part of it), but will not have its infall supported significantly by rotation. One expects that the enclosed mass for the position of the new core inside the old core will be dominated by the protostar in the case of disk fragmentation and (non-stellar) gas in the case of turbulent fragmentation, but these are not discrete categories.  If a star-forming core undergoes turbulent fragmentation, it is possible to view it as forming a sub-cluster within the main protocluster. This could be a major difference from the isolated mode of star formation that forms single stars (or binaries or higher multiples via disk fragmentation). Nevertheless the system can approximate that of isolated star formation from cores if a significant fraction, say $\\geq 1/2$, of the {\\it stellar} mass produced in the core is in a single star (or binary or higher multiple formed by disk fragmentation). In this situation there is still, approximately, a one-to-one correspondence of the core with the final principal star. Thus, another important discriminator between star formation models is the relative importance of turbulent fragmentation that occurs in gas that is already part of a star-forming core compared to gas that is part of the clump, but not yet associated with any particular protostar.  It should be noted that since massive stars are rare (they contain a small fraction of the total young stellar mass), pre-massive-stellar cores are also rare: they require relatively special conditions to form.  Pre-stellar cores may form in a clump that is already rich in (proto)stars so that, especially for large, massive cores, some stars are embedded in the core volume. The orbits of these stars in the clump may be such that they are not bound to the new core, so that their mass is not formally part of the core by the above definitions, even though their mass played some typically minor role in defining its potential. As these stars orbit through and out of the core they may accrete some mass by Bondi-Hoyle accretion, with gas streamlines typically being intercepted by a pre-existing remnant disk from an earlier accretion phase. This Bondi-Hoyle accretion of the core gas is a continuation of Bondi-Hoyle accretion of clump gas that these protostars would have been experiencing prior to the formation of the core.  Note that Bondi-Hoyle accretion involves accretion of gas that is initially not bound to the protostar (strictly speaking not bound to the star-forming core, which is composed mostly of the protostar, together with a remnant accretion disk). The mass affecting the gravitational cross-section here is dominated by the protostar.  This is to be contrasted with accretion of gas to cores in which the gas mass fraction of the core is still significant and the gas distribution still extended. This latter case is less likely to be affected by protostellar feedback.  What if a pre-existing star is bound to the new core? Most stars in these environments will still have at least some remnant accretion disk and so will be protostars and thus would have had remnant cores of their own. This situation can be viewed in a number of ways. From the point of view of a new pre-massive-stellar core, it is just about to form with its new central, defining protostar (that later will become a massive star), but finds itself infested with older, pre-existing stars that may steal some of the massive core gas. From the point of view of the pre-existing protostars, they will likely see their boundaries grow, and maybe merge, as the massive core forms or grows around them, but their boundaries would not include the new LIMP-MP\\footnote{A pre-massive-stellar core could in principle form inside a star-forming core by fragmentation, but then would be expected to have a low mass and to not typically contain pre-existing stars.}. They would likely experience enhanced accretion, possibly contributing significantly to their final mass.  Given these possibilities, one basic question to answer is ``how do massive protostars typically build up their mass?'' The turbulent core model (McKee \\& Tan 2002, 2003) posits that massive stars form from massive, gas-dominated, relatively-near-equilibrium cores, that form by turbulent fragmentation from the magnetized clump medium or by pre-stellar-core agglomeration processes, have some significant fraction of their support from turbulent motions and that during their collapse channel a large fraction of the star-forming gas via a central disk into just one star (or a few stars formed by disk fragmentation).  An alternative possibility is competitive Bondi-Hoyle accretion without involving massive gas dominated cores around the massive protostar (Bonnell et al. 2001; Schmeja \\& Klessen 2004; Bonnell, Vine, \\& Bate 2004; Bonnell \\& Bate 2006). The gas in the core is a relatively small fraction of the massive protostellar mass, but is continuously being replenished by efficient accretion of gas that was previously not bound to the protostar. In these numerical models, this gas is typically being funneled to the center of a star-forming clump that is undergoing global collapse. For some of the massive stars formed in these models -- those whose protostellar seeds are amongst the first to form in the cluster -- their formation can be described as involving a massive gas core, but one which subsequently undergoes efficient turbulent fragmentation to put most of its mass into a cluster of low-mass stars. In neither of these descriptions is there a close correspondence between massive core mass and resulting massive star mass. Since these simulations only include thermal pressure and no magnetic pressure, it is not surprising that they do not see massive near-equilibrium cores, with masses much greater than the thermal Jeans mass, which is $<M_\\odot$ in typical massive star-forming regions.  A third possibility, requiring extremely high stellar densities, is the growth of massive protostars by stellar mergers (Bonnell et al. 1998), which is effectively a merger of star-forming cores, followed by the merger of the protostars. Bally \\& Zinnecker (2005) invoked this mechanism to explain the ``explosive'' nature of the outflow from the Orion KL region (Allen \\& Burton 1993), although we shall see that another more likely possibility is the tidally-enhanced accretion and accretion-powered outflow from close passage of a fast-moving star (BN) with a massive protostar and accretion disk (source {\\it I}) (Tan 2004). ",
        "conclusions": "It is remarkable that, given its astrophysical importance, massive star formation remains so poorly understood. Theories that involve basic differences in the accretion mechanism are actively debated. After defining the terminology and physical properties expected of cores forming together in a star-forming clump, we see that the theoretical differences boil down to whether massive star formation proceeds from massive gravitationally bound gas cores or from competitive accretion to protostellar seeds that are already well-formed before much of their gas is accumulated - i.e. in this latter case the local core potential is mostly determined by the stellar mass rather than the gas mass. We favor the theory of massive star formation from massive gas cores, and suspect numerical models will support this view once they include the physics of magnetic fields (that can help support massive cores) and feedback from protostellar outflows and radiation pressure (that inhibit Bondi-Hoyle accretion and small scale fragmentation near massive protostars - Krumholz 2006).  However, even in the context of models of formation from massive gas cores, understanding what is going on in even the nearest massive protostar and core remains challenging. The Orion KL region is crowded with young stars, and close dynamical interactions definitely occur between them, such as must have accelerated the BN object. I have argued $\\Theta^1C$ is responsible for this event, because it satisfies all the required properties expected of the recoiling, eccentric, hard, massive binary system that must be left behind. As proper motion measurements improve, this issue should be resolved definitively. Source {\\it I} is likely to be an actively accreting massive protostar, that was perturbed by BN's close passage. However, even the orientation of the disk/outflow axis of this system is debated from two orthogonal possibilities. Outflow-confined HII regions are a prediction of massive star formation models that are scaled-up versions of low-mass star formation models. The radio spectrum and morphology of source {\\it I} can be explained by this type of model, but the kinematics of the excited SiO maser spots on 20-100~AU scales remain mysterious."
    },
    "0801/0801.1869_arXiv.txt": {
        "abstract": "We present the spectral analysis of duration-integrated broadband spectra (in $\\sim30\\;$keV$-200\\;$MeV) of 15 bright BATSE gamma-ray bursts (GRBs). Some GRB spectra are very hard, with their spectral peak energies being above the BATSE LAD passband limit of $\\sim$2~MeV. In such cases, their high-energy spectral parameters (peak energy and high-energy power-law indices) cannot be adequately constrained by BATSE LAD data alone. A few dozen bright BATSE GRBs were also observed with EGRET's calorimeter, TASC, in multi-MeV energy band, with a large effective area and fine energy resolution. Combining the BATSE and TASC data, therefore, affords spectra that span four decades of energy ($30\\;$keV$-200\\;$MeV), allowing for a broadband spectral analysis with good statistics. Studying such broadband high-energy spectra of GRB prompt emission is crucial, as they provide key clues to understanding its gamma-ray emission mechanism. Among the 15 GRB spectra, we found two cases with a significant high-energy excess, and another case with a extremely high peak energy (\\epeak~$\\gtrsim$ 170~MeV). There have been very limited number of GRBs observed at MeV energies and above, and only a few instruments have been capable of observing GRBs in this energy band with such high sensitivity. Thus, our analysis results presented here should also help predict GRB observations with current and future high-energy instruments such as AGILE and GLAST, as well as with ground-based very-high-energy telescopes. ",
        "introduction": "High-energy observations of gamma-ray burst (GRB) prompt emission with various detectors have indicated that GRB continuum spectra extend up to MeV--GeV energies \\citep{han94,hur94,cat98,din98,har98,kip98}. GRB spectra in the MeV to GeV range are usually well-described by a single power-law with an index in the approximate range of $-1$ to $-4$ \\citep{sch92, kwo93, hur94, han94, sch95, cat98, kip98, bri+99, wre02}. This range is in agreement with the distributions of the high-energy power-law indices observed with BATSE Large Area Detectors (LADs) in $\\sim$30~keV$-$2~MeV \\citep[\\kan~hereafter]{kan06}. However, due to the power-law nature of GRB spectra, photon counts above $\\sim$1~MeV are usually very low, and this, combined with the fact that the field of view of a high-energy detector is generally limited, results in much fewer GRBs observed in multi-MeV band than keV-band observations.  The Burst and Transient Source Experiment (BATSE) on board the {\\it Compton Gamma-Ray Observatory (CGRO)} has provided the largest GRB spectral database to date in the passband of sub-MeV range. Previously, \\kan~showed that there are some BATSE GRB spectra whose peak energy (\\epeak) lies very close to or above the upper energy bound ($\\sim$2~MeV) of BATSE LADs. For such cases, the LAD data alone cannot adequately determine either \\epeak~or the high-energy power law index. Moreover, the ability to identify the high-energy power law component using LAD data alone is limited, since the LAD sensitivity decreases significantly toward the upper energy bound \\citep{fis89}. Therefore, observations of spectra extending to much higher energies with reasonable sensitivity are needed for these spectral parameters to be well determined. Combining BATSE LAD data with multi-MeV observations by another high-energy detector enables such a broadband study.  Broadband spectral analyses of a few BATSE GRBs have been presented in the literature, using the data obtained with the \\textit{CGRO} instruments \\citep{sch98, bri+99}. However, those broadband spectra were superpositions of the deconvolved photon spectra that were obtained by analyzing each dataset separately with various photon models. Deconvolved photon counts are model dependent, and a spectrum constructed by combining individually deconvolved spectra can be quite different from that obtained properly by simultaneously fitting a common model to all datasets.   The Total Absorption Shower Counter (TASC), the calorimeter of another experiment also aboard the {\\it CGRO} -- the Energetic Gamma-Ray Experiment Telescope (EGRET) spark chamber -- was one of the few instruments capable of observing GRBs in the multi-MeV energy band with excellent energy resolution \\citep{tho93}. Due to its sensitivity to all directions, some bright, BATSE-triggered GRBs were also observed with the TASC apart from the EGRET spark chamber events. Since the TASC provided spectra in the range $\\sim$1$-$200 MeV, broadband GRB spectra spanning four decades of energy can be obtained by combining LAD and TASC data. Such spectra, in at least one case, resulted in the discovery of a distinct multi-MeV spectral component apart from the extrapolated sub-MeV BATSE component \\citep{gon03}. In this paper, we present duration-integrated broadband spectral analysis of 15 bright BATSE GRBs, using LAD and TASC data. We first describe the instuments and available data types of TASC in \\S\\ref{sec:instruments}, and the selection and analysis methodology of our study in \\S\\ref{sec:tasc_selection} \\& \\S\\ref{sec:tasc_analysis}. Then in \\S\\ref{sec:tasc_result}, we present the results, and finally discuss the results in \\S\\ref{sec:sum}. ",
        "conclusions": "\\label{sec:sum} In this study, we extended the BATSE GRB spectral analysis (\\kan) to high-energy broadband spectra obtained by combining BATSE LAD data and EGRET TASC data. Time-integrated joint spectra of 15 hard BATSE GRBs were analyzed in order to probe high-energy spectral properties of prompt emission. The TASC data in multi-MeV energy band with fine energy resolution, clearly confirmed that the GRB spectra do extend up to $\\sim$200~MeV and probably beyond. The joint broadband spectra in the energy range of $\\sim$30~keV$-$200~MeV indeed constrain the high-energy spectral indices and break energies of strong GRBs that have significant MeV emission much better than single detector analysis. In most cases, \\peakenergy~(and \\eb) values derived with LAD data alone and values found by joint analysis were consistent within 1$\\sigma$ uncertainty. However, in a few cases, the jointly-fit indices differed significantly from those determined with the LAD data alone. This indicates the possibility that some high-energy indices obtained with BATSE LADs alone may not reflect the intrinsic values, and thus highlights the importance of broadband spectral analysis.  We identified one case (GRB~930506; trigger 2329) in which the \\peakenergy~could be extremely high, with a lower limit of $\\sim$167 MeV. We note here, however, that this possibly very high \\peakenergy~should be interpreted with caution for two reasons: \\begin{enumerate} \\item{While the high-energy power law index being $> -2$ is statistically very significant (by 7.5$\\sigma$) as determined by the {\\it statistical} error bar, there is the possibility of a systematic error existing in a joint fit between two instruments, which is not taken into account here. Including the systematic uncertainties in the analysis could bring down the index value to $< -2$, in which case the \\peakenergy~would probably become a few MeV.} \\item{Another possibility is that there are two spectral components in this spectrum; the ``regular\" one with \\peakenergy~of a few MeV and an additional high-energy multi-MeV component. In addition, if the spectra evolved similar to the case of GRB~941017, the duration-integrated spectrum could be a superposition of two evolving components. In such cases, the single model used here could hinder correct identification of the \\peakenergy~as well as the high-energy power law index. We found, however, no statistical evidence for such high-energy component in the spectrum of GRB~930506. } \\end{enumerate} In any case, the \\peakenergy~of GRB~930506 is clearly above the BATSE LAD energy range. This, combined with the fact that in some cases the \\peakenergy~values were found to be slightly higher than those determined by LAD data only, may indicate that there exists a tail population of \\peakenergy~values extending to a few hundred MeV, especially given the limited sensitivity of TASC.  In addition, while the broadband spectra were mostly consistent with broken power laws (BAND or SBPL), indications of high-energy excess were also found in two events (GRBs~941017 \\& 980923; triggers 3245 \\& 7113). For one of them (GRB~941017; trigger 3245), the distinct MeV component has been clearly identified with time-resolved spectral analysis \\citep{gon03}, with much stronger significance than with the time-integrated spectrum. For this particular GRB, the existence of the extra component was further reinforced by the COMPTEL spectra, which covered an energy range of 30\\,keV$-$10\\,MeV \\citep{kan04}.  Possible explanations that have been proposed for such a high-energy component include the Compton upscattering of synchrotron photons in the reverse shock by the synchrotron-emitting relativistic electrons \\citep{gra03}, and the electromagnetic cascade emission of ultra-relativistic baryons through photo-pion interactions and subsequent pion decay \\citep{der04}. In the synchrotron shock model \\citep{kat94}, the synchrotron self-Compton emission associated with the synchrotron component that peaks in the keV band is expected to be observed in the MeV to TeV energy band, depending on the Lorentz factor of the electrons \\citep{gue03}. However, the brightness and delayed nature of the high-energy component are inconsistent with the synchrotron-self Compton. It is particularly interesting if the component is indeed due to the relativistic baryons, since the observed component may be direct evidence of baryonic acceleration, namely, cosmic rays.  The high-energy power-law component observed in GRB~941017 by \\citet{gon03} was described with an additional power law of index $\\sim -1$. This requires that there exists another break energy (and a true \\peakenergy) above 200~MeV, in order to avoid the energy divergence. Although the redshift for this event is unknown, a rough upper limit energy for such a break could be placed at $\\sim$1~TeV, solely by the total (isotropic equivalent) energy constraint of $\\sim 10^{53}$ ergs, and by assuming the burst originated nearby ($z \\ll 1$).  As was the case for GRB~941017, time-resolved spectral analyses of high-energy GRBs are needed in order to explicitly identify similar distinct spectral component.  We have indeed performed such time-resolved analyses for all GRBs presented here, the result of which is the subject of another paper that will follow this work (M.M.~Gonz\\'alez et al. in preparation). It is possible that the time-resolved spectral analysis of broadband spectra would reveal such high-energy spectral components in many more GRBs. Also in the MeV energy band and above, there should be a spectral cutoff due to pair-production attenuation \\citep{bar00}. Such a cutoff energy would determine the bulk Lorentz factor $\\Gamma$, which is one of the main factors, along with a redshift, preventing the determination of the intrinsic \\peakenergy~values in the rest frame of expanding matter \\citep{zham04}.  Unfortunately, the TASC data are only available for the brightest of BATSE GRBs due to its sensitivity limitation. Although GRB spectra with a distinct MeV component seem to be rare, and we observed no pair-production cutoff in our analysis, these high-energy spectral features could be sought after using currently existing very high-energy telescopes, such as AGILE\\footnote{http://agile.iasf-roma.inaf.it/}, MAGIC\\footnote{http://wwwmagic.mppmu.mpg.de/}, VERITAS\\footnote{http://veritas.sao.arizona.edu/}, and HESS\\footnote{http://www.mpi-hd.mpg.de/hfm/HESS/HESS.html}. Upcoming observations by {\\it GLAST}\\footnote{http://glast.gsfc.nasa.gov/} with much higher sensitivity in an unprecedented broad energy band of 10~keV$-\\sim$30~GeV, are anticipated to reveal broadband spectral characteristics of GRBs, which holds significant clues to understanding GRB emission mechanism."
    },
    "0801/0801.4027_arXiv.txt": {
        "abstract": "We present deep, wide-field $g$ and $r$ photometry of the transition type dwarf galaxy Leo T, obtained with the blue arm of the Large Binocular Telescope. The data confirm the presence of both very young ($<$1 Gyr) as well as much older ($>$5 Gyr) stars. We study the structural properties of the old and young stellar populations by preferentially selecting either population based on their color and magnitude. The young population is significantly more concentrated than the old population, with half-light radii of 104$\\pm$8 and 148$\\pm$16 pc respectively, and their centers are slightly offset. Approximately 10\\% of the total stellar mass is estimated to be represented by the young stellar population. Comparison of the color-magnitude diagram (CMD) with theoretical isochrones as well as numerical CMD-fitting suggest that star formation began over 10 Gyr ago and continued in recent times until at least a few hundred Myr ago. The CMD-fitting results are indicative of two distinct star formation bursts, with a quiescent period around 3 Gyr ago, albeit at low significance. The results are consistent with no metallicity evolution and [Fe/H]$\\sim -$1.5 over the entire age of the system. Finally, the data show little if any sign of tidal distortion of Leo T. ",
        "introduction": "\\label{sec:intro}  Leo T clearly stands out from the large number of new dwarf galaxies recently discovered \\citep{UMaI,Wil1,CVnI,UMaII,Boo,5pack,LeoT,BooII} around the Milky Way using Sloan Digital Sky Survey \\citep[SDSS,][]{york00,dr5} data. It is by far the most distant, located approximately 420 kpc from the Galaxy and probably not bound to it. Furthermore, it is the only one that contains very young ($<$1 Gyr) stars, apart from an old or intermediate age population, and has detected HI associated with it \\citep{LeoT}. In terms of its properties it seems to be a transitional system, as it combines the round, regular structure of dwarf spheroidal galaxies with the presence of gas and recent star formation common in dwarf irregulars.  The other known so-called transition type dwarf galaxies, DDO 210, Phoenix, Pisces, Antlia, Pegasus and Leo A \\citep[e.g.][]{mateo98,grebel01}, are also isolated Local Group members, but much brighter than Leo T. Because of its large distance from the Milky Way, it is unlikely that Leo T has been significantly affected by the tidal forces of the Milky Way, suggesting that its current low luminosity is intrinsic. From the velocity dispersion of 19 red giants \\cite{simongeha} infer a dark halo mass of $\\sim 10^7 M_\\sun$ and a mass-to-light ratio of $\\sim 140~ M_\\sun/L_{V,\\sun}$ for Leo T. From HI observations, \\cite{ryanweber08} infer a lower limit for the total dynamical mass of $\\sim 3\\times10^6 M_\\sun$ and a mass-to-light ratio of $\\gtrsim 56~ M_\\sun/L_{V,\\sun}$. They also detect two distinct components in the neutral hydrogen, one cold ($\\sim$500 K) and one warm ($\\sim$6000 K). It is interesting to see that such a low luminosity system still contains gas and very young stars at the current epoch.  In this work we present deep photometry of Leo T, obtained with the Large Binocular Telescope. We use these data, the deepest data available on Leo T, to study its stellar populations and structural properties. The outline of this paper is as follows. In \\S \\ref{sec:data} we will first describe the data used in this work. Based on color-magnitude diagram (CMD) morphology we describe the stellar populations present in Leo T in \\S \\ref{sec:pops}. The structural properties of Leo T are analyzed in \\S \\ref{sec:struct} and in \\S \\ref{sec:sfh} we use CMD-fitting techniques to constrain its star formation history and metallicity evolution. Finally, our conclusions are presented in \\S \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  Based on deep $g_0$ and $r_0$ photometry obtained with the LBT, we have studied the stellar populations and structural properties of the Leo T dwarf galaxy. A comparison of the background-subtracted Hess diagram with theoretical isochrones confirms the presence of very young stars with ages between $\\sim$200 Myr and 1 Gyr, and an older stellar population ($>5$ Gyr), the latter with with a likely metallicity of [Fe/H]$\\simeq -$1.7.  The stellar mass in the young population is estimated to be $\\sim$10\\% of the total stellar mass.  Based on the apparent magnitude of the RC stars we determine the distance modulus of Leo T to be 23.1$\\pm$0.2, confirming the value from \\cite{LeoT}. The total luminosity of Leo T is estimated to be $M_{V,total}\\simeq -$8.0 mag, almost one magnitude higher than the estimate from the discovery paper. Using the mass estimate for Leo T from \\cite{simongeha} this implies a mass-to-light ratio of $\\sim 60~ M_\\sun/L_{V,\\sun}$.  A more sophisticated analysis of the photometry using CMD-fitting techniques yields similar results. The distance determination gives a distance modulus of $m-M=$23.0$\\pm$0.2. CMD fits with two different software packages both indicate two distinct episodes of star formation, one at $\\sim$6 -- 12 Gyr and one at $\\sim$0.4 -- 1 Gyr, although the drop in star formation rate in between has a low significance. From the CMD fits no strong metallicity evolution is apparent, with an approximate metallicity of [Fe/H]$\\sim -$1.5 at all times. The metallicity seems to have a small dip around $\\log(t)=9.2$, but since the metallicity constraints at that age come from regions of the CMD where the stellar models are least reliable, such as the RC and asymptotic giant branch, this is not significant.  Old and young stars were preferentially selected using CMD-selection boxes in order to study their individual spatial distributions. The young stars are more strongly concentrated near the center of Leo T, but both components display the same spatial extent.  There seems to be a slight offset between the young and old stars, but from these data this is hardly significant. Since stars are formed in clusters, young stars are usually spatially more strongly confined than older stars. This has also been found in other dwarf galaxies \\citep[e.g.][]{hurley99,martinez99,harbeck01,battaglia06,mcconnachie06,martin08}. The half-light radius for all stars is found to be 120$\\pm$7 pc, slightly smaller than the value from \\cite{LeoT}.  Comparing our results with the properties of the other transition type dwarf galaxies in the Local group, DDO 210, Phoenix, Pisces, Antlia, Pegasus and Leo A \\citep[e.g.][]{mateo98,grebel01}, shows a lot of similarities. The presence of both old or intermediate ($>$5 Gyr) age stars and very young ($<$1 Gyr) stars is common, as is the metallicity of [Fe/H]$\\sim-$1.5. A similar apparent lack of metallicity evolution was found in Leo A by \\cite{cole07}. While in terms of total mass and HI mass Leo T is similar to most other transition galaxies, it is much fainter. Our luminosity estimate of $M_{V,total}\\simeq -$8.0 is two magnitudes fainter than any of the other transition objects \\citep[e.g.][]{mateo98}. Therefore Leo T remains the least luminous galaxy with recent star formation known.  We find no evidence for tidal distortion of Leo T. This is not unexpected, since it is at a large distance from the Milky Way and M31, although based on the limited information about its orbit it cannot be excluded that Leo T has come close to the Milky Way some time in the past. However, the lack of distortion implies that its low mass and luminosity are intrinsic, rather than the result of interaction or disruption. Thus, the case of Leo T suggests that isolated halos with masses as low as $10^7 M_\\sun$ are able to accrete and/or retain gas and form stars for at least a Hubble time."
    },
    "0801/0801.3411.txt": {
        "abstract": "The quite different behaviors exhibited by microscopic and macroscopic systems with respect to quantum interferences suggest that there may exist a naturally frontier between quantum and classical worlds. The value of the Planck mass (22$\\mu$g) may lead to the idea of a connection between this borderline and intrinsic fluctuations of spacetime. We show that it is possible to obtain quantitative answers to these questions by studying the diffusion and decoherence mechanisms induced on quantum systems by gravitational waves generated at the galactic or cosmic scales. We prove that this universal fluctuating environment strongly affects quantum interferences on macroscopic systems, while leaving essentially untouched those on microscopic systems. We obtain the relevant parameters which, besides the ratio of the system's mass to Planck mass, characterize the diffusion constant and decoherence time. We discuss the feasibility of experiments aiming at observing these effects in the context of ongoing progress towards more and more sensitive matter-wave interferometry. ",
        "introduction": "When he introduced the relativistic conception of space-time~\\cite{Einstein05}, Einstein emphasised that remote clocks have to be compared through the transfer of time references encoded on electromagnetic pulses. This clock synchronisation procedure between distant observers is nowadays a routine for metrological applications~\\cite{TFmetrology} as well as the basic building block of Global Navigation Satellite Systems~\\cite{GNSS} or the definition of reference systems~\\cite{refsystems}. The tracking observables used in most gravity tests performed in the solar system are also built up on similar protocols. This statement is valid for example for radio tracking of planetary probes~\\cite{radioscience}, in particular Pioneer~\\cite{pioneer} and Cassini probes~\\cite{cassini}, as well as lunar laser ranging~\\cite{LLR}. In the present section, we present the basic tracking observables which are built up on electromagnetic time transfer signals exchanged between remote observers and compared to locally available atomic clocks. We will see in the next section how these observables are affected by gravitational waves.  To make the discussion more concrete, let us consider a synchronisation between an atomic clock on board a probe in space and another one colocated with a station on ground. Both clocks (observers) are supposed to be equipped with transfer capabilities corresponding to the emission and reception of a time reference encoded as the ``center of energy'' of an electromagnetic pulse or a marked phase front. The time reference transferred from one observer to the other is by necessity a quantity encoded on the field and preserved by the law of propagation. In special relativity, one may define this time reference as the light cone variable which labels the various phase fronts along the line of sight. To be more precise, let us consider an emission event with positions $t_\\e$ in time and $x_\\e$ on the space axis along which the transfer takes place. Then the transfer connects the reception event with positions $t_\\r$ in time and $x_\\r$ on the space axis with the emission event with positions $t_\\e$ and $x_\\e$ if they share the same value of the light-cone variable \\beq \\label{Minkowski_synchro} t_\\e-\\frac {x_\\e}c = u = t_\\r-\\frac {x_\\r}c \\eeq This equation also means that the spatial distance $\\left|x_\\r-x_\\e\\right|$ between the two events is directly determined by the elapsed time $c\\left(t_\\r-t_\\e\\right)$. As a consequence, one of the observers may measure its distance to the other one as follows~: he emits a pulse to be reflected by the remote observer and measures the time elapsed till he gets the pulse back; he then deduces the spatial distance as half this lapse time; this well known principle of radar ranging produces what is called a ``range''.  In presence of gravity fields, we have to replace the simple relation (\\ref{Minkowski_synchro}) by a more general theoretical description of the light cones which connect emission and reception events. To this aim, we define the infinitesimal metric element $\\d s$ from the metric tensor $g_{\\mu\\nu}$ characterizing space-time and displacements $\\d x^\\mu$ in space-time \\beq \\label{metric_element} \\d s^2\\equiv g_{\\mu\\nu}\\d x^\\mu \\d x^\\nu \\eeq The light cones are the solutions of the eikonal equation written for a massless probe $\\d s^2=0$. In the weak gravitational field of the solar system, there is a unique solution for an electromagnetic ray going from one point to the other. For any couple of massive observers, labelled 1 and 2, the light-cone solution can be written as a relation between the times of the emission and reception events. This relation can be obtained in a variety of representations, for example through an explicit solution of the light-cone equations~\\cite{explicit}, or using the so-called world function~\\cite{worldfunction} or time delay function~\\cite{JRcqg06}. In the latter point of view, the light-cone equation is written as \\beq t_2-t_1=\\mathcal{T} \\left(t_1,\\mathbf{x}_1(t_1);t_2,\\mathbf{x}_2(t_2)\\right) \\eeq The timedelay function $\\mathcal{T}$ measures the time taken to propagate the pulse between the emitter and receiver, the motions of which are described by $\\mathbf{x}_1(t_1)$ and $\\mathbf{x}_2(t_2)$.  In order to illustrate the method, we write the timedelay function $\\mathcal{T}$ in the solar system with a few simplifying assumptions~: we ignore other gravity sources than the Sun, taken as pointlike and motionless; we use the Eddington gauge convention where spatial coordinates are isotropic (see~\\cite{Varenna1} for a more detailed discussion) and treat the effect of gravity up to first order in $GM$ where $G$ is the gravitational constant and $M$ the mass of the Sun. With these assumptions (and taking $\\gamma\\equiv1$ for general relativity; see discussions in \\cite{Varenna1}), the timedelay is the following function of the spatial positions of the two endpoints ($R_{12} \\equiv|\\mathbf{x}_2-\\mathbf{x}_1|,\\, r_1\\equiv|\\mathbf{x}_1|,\\, r_2\\equiv|\\mathbf{x}_2|$) \\beq \\label{Shapiro} &&c\\mathcal{T} \\left(\\mathbf{x}_1;\\mathbf{x}_2\\right) = R_{12} + \\frac{2GM}{c^2} \\ln\\frac{r_1+r_2+R_{12}}{r_1+r_2-R_{12}} \\eeq In this function, the first term $R_{12}$ represents the Minkowskian approximation (zeroth-order in $GM$), while the second term is the (first-order) effect of gravity on electromagnetic propagation, known as the Shapiro time delay~\\cite{Shapiro}. The order of magnitude of this term is fixed by $2GM/c^2\\simeq3$km.  At this point, it is worth emphasizing that the light-cone relation between emission and reception endpoints is independent of the coordinate system (\\textit{i.e.} gauge invariant) although the explicit expression of the timedelay function depends on the specific coordinate system (specific gauge convention). There are some conveniency reasons for choosing the Eddington gauge convention for which the spatial part of the metric is isotropic, so that the propagation of electromagnetic field also respects isotropy (with the simplifications evoked in the preceding paragraph). But other gauge choices would be perfectly as respectable as this one, and they would lead to the same final expressions for the physical observables. This is true in particular for the range defined in the foregoing paragraphs as a gauge invariant spatial distance between remote observers.  To fix ideas, let us suppose we study time transfer between two atomic clocks, one brought on board a space probe, the other one colocated with a radio or laser station on Earth~\\cite{SAGAS}. Such transfers can be performed on up- as well as down-links. The uplink signal emitted connects two points at positions $(t_1^\\u,\\bx_1(t_1^\\u))$ on ground and $(t_2^\\u,\\bx_2(t_2^\\u))$ in space while the downlink signal connects $(t_1^\\d,\\bx_1(t_1^\\d))$ on ground and $(t_2^\\d,\\bx_2(t_2^\\d))$ in space. When the two links are crossing at the space endpoint ($t_2^\\u\\equiv t_2^\\d$), the observer on ground may define a spatial distance to the remote probe from the times of emission and reception at his station. This distance is a gauge invariant observable (independent of the coordinate system) as soon as it is written in terms of clock indications \\beq \\label{range} \\ell_{12} \\equiv {\\Delta s \\over2} \\quad,\\quad \\Delta s \\equiv s_1^\\mathrm{d} - s_1^\\mathrm{u} = \\int_{t_1^\\u} ^{t_1^\\d} \\d s_1 \\quad&\\leftrightarrow&\\quad \\left[ \\mathrm{ranging\\;with\\;} s_2^\\d = s_2^\\u\\right] \\eeq $s_1$ is the proper time (multiplied by $c$) as measured by the clock on ground, and $\\ell_{12}$ is half the proper time elapsed during the roundtrip of the signal to and from the probe.  Note that it can also be written as the half sum of quantities defined on the up- and down-links (with $s_2^\\d = s_2^\\u \\equiv s_2$) \\beq \\label{range_updown} \\ell_{12} \\equiv {\\ell_{12}^\\u + \\ell_{12}^\\d \\over2} \\quad&,&\\quad \\ell_{12}^\\u \\equiv s_2 - s_1^\\u \\quad,\\quad \\ell_{12}^\\d \\equiv s_1^\\d - s_2 \\eeq The latter are proper relativistic observables, which do not depend on the coordinate system, but however depend on the choice of the origins of proper times for the two clocks. This dependence can be fixed by a convention stating that the two up- and down- one-way ranges are equal, at some point chosen for any conveniency reason.  The Doppler tracking observables which have been over years the main source of information on the navigation of remote probes \\cite{radioscience} are directly related to the time derivatives of these ranges. In order to make this point explicit, let us differentiate (\\ref{range_updown}) with respect to a commonly defined time, for example $\\d s_2$. We thus get the ratios of frequencies measured on the up- and down-link at the ground station as \\beq \\label{Doppler} &&{\\omega_1^\\d\\over\\omega_1^\\u} = {\\d s_1^\\u\\over\\d s_1^\\d} = { 1 - \\dot\\ell_{12}^\\u \\over 1 + \\dot\\ell_{12}^\\d  } \\quad,\\quad \\dot\\ell_{12}^\\u \\equiv {\\d\\ell_{12}^\\u\\over \\d s_2} \\quad,\\quad \\dot\\ell_{12}^\\d \\equiv {\\d\\ell_{12}^\\d\\over \\d s_2} \\eeq This Doppler observable is directly given by time derivatives of the ranges. In the non relativistic limit, where velocities and gravity fields are small, it can be written in terms of the range (\\ref{range}) only \\beq \\label{Doppler_approx} &&{\\omega_1^\\d\\over\\omega_1^\\u} = {\\d s_1^\\u\\over\\d s_1^\\d} \\simeq 1 - \\dot\\ell_{12}^\\u - \\dot\\ell_{12}^\\d  = 1 - 2 \\dot\\ell_{12} \\quad,\\quad \\dot\\ell_{12} \\equiv {\\d\\ell_{12} \\over \\d s_2} \\eeq It follows that the observable $\\dot\\ell_{12} $ can be regarded as a properly defined relative velocity between the two remote observers. Note that the Doppler observable (\\ref{Doppler}) contains not only what is usually called the Doppler effect (relative motion of one observer with respect to the other) but also the Einstein effect (effect of gravity on the clock rates) as well as the Shapiro effect (effect of gravity on the propagation of light). In the non relativistic limit, it can be written as a sum of terms each corresponding to one of these interpretations. It must however be kept in mind that the individual terms are no longer gauge invariant whereas the sum is gauge invariant from its very definition.  Before ending this introduction on the status of tracking observables, let us recall that the Doppler observable is the only source of information on the navigation of Pioneer probes~\\cite{pioneer} (see the discussions and references in \\cite{Varenna1}) as well as the primary signal in the Cassini relativity experiment~\\cite{cassini}. ",
        "conclusions": ""
    },
    "0801/0801.2571_arXiv.txt": {
        "abstract": "We study the constraints which the next generation of radio telescopes could place on the nature of dark energy, dark matter and inflation by studying the gravitational lensing of high redshift 21~cm emission, and we compare with the constraints obtainable from wide-angle surveys of galaxy lensing.  If the reionization epoch is effectively at $z\\sim 8$ or later, very large amounts of cosmological information will be accessible to telescopes like SKA and LOFAR. We use simple characterizations of reionization history and of proposed telescope designs to investigate how well the two-dimensional convergence power spectrum, the three-dimensional matter power spectrum, the evolution of the linear growth function, and the standard cosmological parameters can be measured from radio data. The power spectra can be measured accurately over a wide range of wavenumbers at $z\\sim 2$, and the evolution in the cosmic energy density can be probed from $z\\sim 0.5$ to $z\\sim 7$. This results in a characterization of the shape of the power spectra (i.e. of the nature of dark matter and of inflationary structure generation) which is potentially more precise than that obtained from galaxy lensing surveys. On the other hand, the dark energy parameters in their conventional parametrization ($\\Omega_\\Lambda, w_o, w_a$) are somewhat less well constrained by feasible 21~cm lensing surveys than by an all-sky galaxy lensing survey. This is because dark energy is felt primarily at relatively low redshifts in this model; 21~cm surveys would be more powerful than galaxy surveys for constraining models with ``early'' dark energy.  Overall, the best constraints come from combining surveys of the two types. This results in extremely tight constraints on dark matter and inflation, and improves constraints on dark energy, as judged by the standard figure of  merit, by more than an order of magnitude over either survey alone. ",
        "introduction": "\\label{sec:introduction}  The now widely accepted cold dark matter (CDM) model combines with inflation-inspired scale-invariant primordial density fluctuations to provide a consistent explanation for cosmic microwave background (CMB) fluctuations, for type Ia supernova luminosity distances, for the clustering of galaxies in redshift surveys, for the galaxy cluster abundance and its evolution, and for the statistics both of weak gravitational lensing and of Ly$\\alpha$ forest absorption in quasar spectra. This success comes at the cost of introducing several mysterious and apparently {\\it ad hoc} constituents.  Most of the matter is supposed to be dark, a weakly interacting neutral particle so far detected only through its gravitational effects.  All structure is supposed to have originated during a very early period of accelerated expansion driven by an inflaton field which has been posited purely for this purpose.  Finally the energy density of the Universe is apparently dominated today by a different and even more unexpected field, dark energy, which is similarly accelerating the present cosmic expansion. The nature of the dark energy and its relation to the rest of physics are unknown. Detailed measurements of the recent expansion history and of the corresponding gravitationally driven growth of structure provide the only known route to narrow down the possibilities.  The techniques mentioned above can measure the expansion precisely out to $z\\sim 4$ and the growth of structure out to $z\\sim 1.5$.  In this paper we investigate a new technique that complements, refines and greatly increases the precision of these methods by directly constraining both cosmic expansion and structural growth out to $z\\sim 6$.   The spin temperature of neutral hydrogen during and before the epoch of reionization ($8 \\simlt z \\simlt 300$) fell out of thermal equilibrium with the CMB radiation, resulting in the absorption and emission of 21~cm radiation.  There has been a great deal of interest in the prospect of detecting and mapping this radiation using low-frequency radio telescopes, Several suitable instruments are now under construction or in planning stages \\citep[see][for an extensive review]{astro-ph/0608032}.  This radiation provides an excellent source for gravitational lensing studies.  Structure is expected in the 21~cm emission down to arcsecond scales, and at each point on the sky there will be $\\sim 1000$ statistically independent regions, separated in redshift (and thus frequency) that can in principle be observed.  Gravitational lensing coherently distorts the 21~cm brightness temperature maps at different frequencies.  This coherent distortion can be distinguished from intrinsic structure in the HI gas if enough independent redshifts are observed.  In this way a map of the foreground density can be constructed \\citep{ZandZ2006,metcalf&white2006,HMW07,2007arXiv0710.1108L}.  Observing the 21~cm radiation at high redshift will be challenging. It will require large arrays of radio telescopes working at low frequencies ( $\\sim 100$ to 200 MHz). At these frequencies foregrounds -- terrestrial, galactic and extragalactic -- are very large and will need to be subtracted by complex and as yet untested procedures \\citep{astro-ph/0608032}. This challenge is currently being addressed by several observational teams (see section \\ref{sec:model-telescopes}).  In addition to the noise associated with mapping the brightness temperature, the lensing signal has an additional {\\it intrinsic noise} component which comes from the unknown intrinsic structure of the 21~cm brightness temperature distribution.  This noise cannot be reduced by increasing the collecting area of the telescope, by increasing the integration time or by improving the removal of foregrounds.  \\cite{metcalf&white2006} showed that if the signal-to-noise in the brightness temperature map at each frequency is greater than one, then the noise in the mass map will be close to the intrinsic value.  Increasing the frequency resolution of the radio observations increases the number of effectively independent regions along the line-of-sight until the bandwidth becomes smaller than the correlation length in the redshift direction of the brightness temperature distribution.  If the bandwidth is matched to the correlation length, the intrinsic noise is minimized.  The correlation length in turn depends on beam size, and is smaller for smaller beams.  Thus unlike galaxy lensing surveys, the intrinsic noise for 21~cm lensing {\\it decreases} with increasing telescope resolution.  In practice, there is a trade-off because smaller bandwidth means less flux, but this can be compensated by increasing collecting area and/or integration time.  The noise is also affected by the range of frequency (i.e. redshift) over which the 21~cm radiation can be detected.  The low-frequency limit is set by the telescope and/or the ability to subtract foregrounds.  The high-frequency limit is typically set by the reionization of the universe, after which the amount of intergalactic HI is small.  In \\cite{HMW07} we simulated how well an idealized optimal telescope would be able to map the 2 dimensional distribution of matter.  In this paper, we study how well radio telescopes with specifications similar to those of instruments currently being built or planned will be able to constrain cosmological parameters, in particular those related to the nature of dark matter, dark energy and inflation.  This paper is organized as follows.  The next two sections introduce the formalism we use to study lensing in general and lensing of the 21~cm radiation in particular.  In section~\\ref{sec:param-estim-dens} we develop a formalism for extracting cosmological information and density maps from such lensing data.  In section~\\ref{sec:model-reionization} we discuss the relevant aspects of reionization and the simple reionization model we use to make quantitative predictions.  The telescope parameters used in our predictions are given in section~\\ref{sec:model-telescopes}. Section~\\ref{sec:2-dimens-inform} contains predictions for the noise levels in various quantities.  A summary of our results and of future prospects is given in the last section. ",
        "conclusions": "\\label{sec:discussion}  We have shown that a very large amount of cosmological information could be extracted from the gravitational lensing of pregalactic 21~cm radiation.  The planned low-frequency radio telescopes, an upgraded LOFAR or SKA, will already have sufficient sensitivity to do this successfully as long as reionization occurs sufficiently late, foreground contamination can be accurately removed, and enough sky can be observed. Under these conditions, the matter power spectrum should be well measured in two and three dimensions, and the growth of structure will be probed from $z=0.5$ to $z=7$.  Such observations would constrain many cosmological quantities, in particular the form of the primordial power spectrum and presence of early dark energy, to unprecedented accuracy.  Combining 21~cm lensing surveys with surveys of lensing of foreground galaxies would dramatically improve constraints on structure growth at $z\\simlt 1$ where the strongest effects occur in the most popular parameterizations of dark energy.  Increasing the collecting area of the telescopes by a factor of $\\sim 2$ would also dramatically increase their sensitivity to lensing, as would extending their frequency range to lower values, if foregrounds permit.  \\vspace{0.3cm} \\leftline{\\bf Acknowledgments} We would like to thank Adam Amara for allowing use of his own software to compare with results from our own parameter constraint codes.  We would like than J. Weller for help in calculating the expected fisher matrix for the Planck satellite mission."
    },
    "0801/0801.2086_arXiv.txt": {
        "abstract": "We have processed the data accumulated with \\textit{INTEGRAL} SPI during 4 years ($\\sim$ 51 Ms) to  study the  Galactic ``diffuse'' emission morphology in the 20 keV to 8 MeV energy range. To achieve this objective, we have derived simultaneously an all-sky census of  emitting sources and images of the Galactic Ridge (GR) emission. In the central radian, the  resolved point source emission amounts to 88\\%, 91\\% and 68\\%  of the  total emission in the 25-50, 50-100 and 100-300 keV domains respectively. We have compared the  GR emission spatial distribution to those obtained from CO and NIR maps, and quantified our results through latitude and longitude profiles. Below 50 keV, the SPI data are better traced by the latter, supporting a stellar origin for this emission. Furthermore, we found that the GR emission spectrum follows a power law with  a photon index $\\sim$ 1.55 above 50 keV while an additional component is required below this energy. This component shows a cutoff around 30 keV, reinforcing a stellar origin, as proposed by Krivonos et al. (2007).\\\\ The  annihilation diffuse emission component   is extracted simultaneously, leading to the determination of the related parameters (positronium flux and fraction). A specific discussion is devoted to the annihilation line distribution since a significant emission is detected over a region as large as $\\sim 80^\\circ$ by  $\\sim 10^\\circ$  potentially associated with the disk or halo surrounding the central regions of our Galaxy. ",
        "introduction": "The soft $\\gamma$-ray GR Emission ($>$ 20 keV) has  been previously studied essentially with the CGRO and GRANAT missions (Purcell et al., 1996, Skibo et al.,  1997, Kinzer, Purcell \\& Kurfess, 1999). The main conclusion was that point sources explain at least 50 \\% of  the total emission up to $\\sim$  ~200 keV, allowing the possibility to anticipate that unresolved sources could explain a major part of what was seen as ``diffuse continuum emission''. Soon after its launch, first INTEGRAL results indicated that the GR diffuse emission is less than 15 \\% of the total in the 20-200 keV domain (Lebrun  et al., 2004, Terrier et al., 2004, Strong et al., 2005, Bouchet et al., 2005). In  more recent works, Revnivtsev et al. (2006), with RXTE~PCA data, and Krivonos et al. (2007), using the data collected with the \\textit{INTEGRAL} IBIS  telescope, suggest that the GRXE (Galactic Ridge X-ray Emission) between 3 and  $\\sim$ 60 keV is explained in terms of sources belonging to the population of accreting white dwarfs binaries. \\\\ In a previous paper using one year of SPI data (Bouchet et al., 2005, hereafter paper I), we had used Galactic tracer morphologies to extract the spectrum of the diffuse continuum emission as well as that of the positronium/annihilation in the central radian of our Galaxy. We thus derived the relative contributions of these emissions along with that of the point source emission to the total Galactic emission. In this paper, we  present a global view of the soft $\\gamma$-ray sky emission based on a larger amount of data covering the whole sky and perform  an imaging analysis of the diffuse component, with an estimate of its spectrum between 20 keV and 8 MeV. ",
        "conclusions": "The \"diffuse emission\" generic term gathers a set of various processus depending on the energy domain, and making the determination of its origin rather difficult. In the soft $\\gamma$-ray domain, the diffuse emission represents a complex problem, as its presence is swamped with the individual sources emission and difficult to disentangle.\\\\ To investigate it, we have  developed an imaging algorithm dedicated to extended structures. This code first determines and takes into account the point sources emitting in the same energy domain, allowing us to then better estimate the  geometry of the regions producing diffuse emissions.\\\\ The SPI all-sky survey analysis reveals that 173 sources can be identified in the 25-50 keV with 30 of them emitting a significant flux ($>$ $\\sim$ 2.5 $\\sigma$) above 100 keV. It is clear now that the sources dominate the Galactic emission up to $\\sim$ 300 keV (Fig.~\\ref{fig:spectrum}). Finally, upper limits on the diffuse emission are estimated to be one tenth of the total emission below 100 keV and one third in the 100-300 keV band. \\\\  The $e{^-}/e{^+}$ annihilation produces an emission which has already been the subject of many studies. The main point  emerging from our analysis concerns the spatial distribution of this emission since a large structure, potentially associated with the Galactic disk/halo  has been found to contain a significant fraction of the 511 keV line total flux. While a deeper analysis is required to refine this result, it is clear that the  $\\sim 8^\\circ$  Gaussian distribution previoulsy  reported does not represent the complexity of the  511 keV line emission morphology.\\\\ Once the source and annihilation process contributions have been taken into account, we can access to the \"diffuse\" Galactic emission, whose origin remains problematic.\\\\ Its  morphology is of prime importance as it is thought to trace the CR electron population, observed through bremsstrahlung or Compton emission. The \\textit{INTEGRAL} SPI data allow us to investigate this particular topic, and we present for the first time images of the diffuse Galactic emission up to a few MeV.\\\\ This diffuse emission is clearly detected between 50 keV and 2 MeV with a power law relatively hard (photon index around 1.55), and an additional component, much steeper, required below 50 keV.\\\\ Toward high energy, the SPI spectrum joins with the OSSE and EGRET measurements. All together, they can be compared to various models to understand the nature of this emission. The inverse Compton interaction is a good candidate and will be investigated in a forthcoming paper. Other interstellar processes have been invoked to explain the emission below 1 MeV (Dogiel et al., 2002; Masai et al., 2002; Schr\\\"oder et al., 2005). In another side, unresolved point sources are the most probable origin below 50 keV (Krivonos et al., 2007) and are also a possibility at higher energies as AXPs or other pulsars (see for exemple Kuiper et al. 2004; Strong, 2007). Observations of such sources up to a few MeV would allow to quantify their potential contribution and bring a new piece to the Galactic diffuse emission riddle."
    },
    "0801/0801.2615_arXiv.txt": {
        "abstract": "We investigate the tidal interactions of a red giant with a main sequence in the dense stellar core of globular clusters by Smoothed Particle Hydrodynamics method. Two models of $0.8 \\msun$ red giant with the surface radii $20$ and $85 R_\\sun$ are used with 0.6 or $0.8M_\\sun$ main sequence star treated as a point mass. We demonstrate that even for the wide encounters that two stars fly apart, the angular momentum of orbital motion can be deposited into the red giant envelope to such an extent as to trigger rotational mixing and to explain the fast rotation observed for the horizontal branch stars, and also that sufficient mass can be accreted on the main sequence stars to disguise their surface convective zone with the matter from the red giant envelope. On the basis of the present results, we discuss the parameter dependence of these transfer characteristics with non-linear effects taken into account, and derive fitting formulae to give the amounts of energy and angular momentum deposited into the red giant and of mass accreted onto the perturber as functions of stellar parameters and the impact parameter of encounter. These formulae are applicable to the encounters not only of the red giants but also of the main sequence stars, and useful in the study of the evolution of stellar systems with the star-star interactions taken into account. ",
        "introduction": "In the core of globular clusters, it is thought that star-star interactions play an important role because of very high stellar density and of relatively low velocity \\citep{hills76b}. There is growing evidence for the modifications of stellar properties and populations under the influence of close encounters and collisions. For example, the overabundance of low-mass X-ray binaries and millisecond pulsars is regarded as consequent upon tidal captures of an environment star by neutron stars, and upon the exchange encounters involving a neutron star \\citep{fabian75,hills76a}: the smaller relative frequency of red giants in the core is attributed to the deprival of their envelope during close encounters with environment stars \\citep{djorgoski91, beer2004}, and blue-stragglers, which are main sequence stars more massive than the turn-off stars, may result from direct collisional coalescence and/or binary merge of two or more stars \\citep[][see Mapelli et al. 2006, and Leigh et al. 2007 and references therein for recent works]{leonard89}; see also reviews by \\citet{bailyn95}, \\citet{hut03} and \\citet{ferraro06}. In particular, the inflation of the number of known blue stragglers boosted by the observations using Hubble Space Telescope \\citep{ferraro97, ferraro99, paltrinieri01, ferraro03, ferraro04, sabbi04, beccari06, warren06}, suggests that a significant fraction of stellar populations undergo such encounters with neighboring stars. Recent observations with the {\\sl Chandra} X-ray Observatory indicate a link between the numbers of X-rays binaries and the stellar encounter rates in globular clusters \\citep{pooley03, pooley06}.  Furthermore, there is a longstanding problem of large star-to-star variations in the surface abundances of light elements such as C, N, O, Ne, Mg and Al. Some giants in globular clusters exhibit the anomalous surface abundances that cannot be explain in terms of the nucleosynthesis and material mixing in the stars within the current standard framework of stellar evolution \\citep[e.g., see reviews by][]{kraft94, dacosta97}. Since these anomalies are observed only in globular clusters but not from field giants in the Galactic halo, it is natural to search for their origin(s) in the differences between the environment in the globular clusters and in the Galactic halo, and hence, to consider them as a evidence of the star-star interactions. In fact, \\citet{fujimoto99} have proposed a scenario for the formation mechanism of these abundance anomalies involving hydrogen shell flashes in red giants, as a result of internal mixing, triggered by the deposition of angular momentum into their envelopes during a close encounter with other stars. It is demonstrated that this extra-mixing model can reproduce the observed relationship such as correlation and scatter in the anomalous abundances of Na and Al and the Mg-Al anti-correlation \\citep{aikawa01,aikawa04}.  Recently, similar abundance variations are found among unevolved stars of turn-off and sub-giant \\citep{gratton01}. It is true that the variations and anti-correlations between CN and CH bands have been reported not only for giants but for stars of upper main sequence, which may be taken to suggest the presence of abundance anomalies in unevolved stars \\citep{suntzeff91,briley92, cannon,cohen}. It has been argued that these facts refute the evolutionary scenario that the abundance anomalies are produced during the evolution along the giant branch and favor the primordial scenario that the stars were born of gas already subject to the anomalous abundances \\cite[e.g.,][]{sneden04}. As a possible compromise, the recycling scenarios have been proposed that the inhomogeneity is due to the surface pollution by accreting the ejecta of anomalous abundances from the erstwhile AGB stars of intermediate masses \\citep{thoul02}, or that the second-generation stars were born from the gas polluted with the ejecta of AGB stars \\citep{dantona04,dantona04b}. \\citet{ventura01} argue that the burning at the bottom of convective zone in low metallicity AGB stars can produce the observed O-Na and Mg-Al anti-correlations. On the other hand, \\citet{Fenner04} cast doubt upon the relevance of AGB ejecta to the observed anomalies. In addition, the scenarios involve serious difficulties both in the mass supply and in the overabundances of CN and s-process elements, attendant with the third dredge-up during the evolution; the necessary amount of mass ejecta only to cover and disguise the surface convection of giants may well exceed the total envelope mass that can be ejected from the erstwhile AGB stars, and the enrichment of s-process elements is never observed \\citep{James04}.  For the evolutionary scenario, it can also be argued that the abundance anomalies are printed onto the surface of unevolved stars through the mass transfer during the close encounters with such giants that have already developed the abundance anomalies; the anomalous abundances are themselves attributed to the deposition of angular momentum into the convective envelope of giants during prior close encounters with environment stars \\citep{shimada03}. The surface convection of population II main sequence stars contain mass of $3 \\times 10^{-3} M_\\odot$ near the turn-off \\citep[e.g., see][]{fujimoto95}, and hence, the accretion of mass of this order may suffice to disguise their surface abundances with those of accreted matter. The evolutionary scenario with star-star interactions during the red giant branch taken into account are free from the above difficulties and have a fair prospect of giving a satisfactory explanation to these inhomogeneous anomalies. Recently, the helium production by this extra-mixing mechanism is discussed \\citep{suda07} in relevance to the splitting of main sequence branch, observed from $\\omega$ Cen \\citep{Bedin04} and from NGC 2808 \\citep{piotto07}.  One of the aims of this paper is to investigate whether the orbital angular momentum can be transferred into the envelope of giant from the orbital motion and whether the main sequence star can accrete the envelope mass from giants enough to disguise their surface layer with the accreted matter through star-star interactions. It is argued that the rotation-induced mixing requires (differential) rotation of $\\sim 0.01$ times the local critical rate from the energetic viewpoint \\citep[e.g., see][]{fujimoto99}, although the proper theory is yet to be established. From the observations, the horizontal branch stars are known to display a bimodal distribution of rotation velocity with the fastest rotators at velocity $v_{\\rm rot} \\sin i \\gtrsim 30 \\hbox{ km s}^{-1}$ (where $i$ is the inclination angle of spin axis) on the cooler side of horizontal branch where $T_{\\rm eff} < 15000^\\circ$ K \\citep{peterson83, peterson95, cohen97, behr00a, behr00b, recio02}. If we neglect the angular momentum loss during the transition to the horizontal branch, such rapid rotations require the angular momentum corresponding to the rotation rate of an order of $\\Omega \\simeq 0.01 \\Omega_{\\rm RG}$ at the tip of red giants ($\\Omega_{\\rm RG}$ being the critical rotation rate at the surface). This poses a problem of the origin of angular momentum since in the low-mass stars, the angular momentum is effectively extracted by magnetic braking during the main sequence phase and by mass loss during the red giant phase \\citep[e.g., see][]{recio02,suda06}.  The star-star interactions have been proposed as the mechanism(s) to form the unusual stellar objects discussed above and studied by many authors. Among the analytical approaches, \\citet{fabian75} first presented an idea and evaluated the possibility that the low-mass X-ray binaries are produced through tidal dissipation during the two body encounters involving a neutron star or low-mass black hole. \\citet{press77} developed the linear perturbation theory of the two-body tidal capture mechanism to derive a general formula for the amount of orbital energy, deposited into the oscillatory modes of stellar envelope during a periastron passage. \\citet{lee86} and \\citet{mcmillan87} worked out the cross sections for the binary formation via tidal capture of a main sequence star and  \\citet{mcmillan90} that of a red giant, respectively. These studies are, however, limited to the linear regime and can not deal with the non-linear effects such as the mass transfer between the stars and the mass loss from the stars owing to large deformations of the stars by tidal force.  In order to estimate the non-linear effects during the close encounters, numerical simulations are necessary. Among the numerical approaches, most studies have been devoted to understanding the resultant offspring of the stellar interactions. Simulations of tidal encounters have been performed for the various combinations of stars, e.g., a main sequence and a red giant star \\citep{benz91}:  a neutron star and a main sequence star or a red giant star in an attempt to explain the formation of the low-mass X-ray binaries and the millisecond pulsars \\citep{davies92, rasio91, davies95, lee96}: main sequence stars in encounter and collision, aiming at the formation of blue stragglers \\citep{lai93, lombardi02}: and a red giant star and a neutron star in relation to the formation of pulsars or ultra-compact X-ray binaries \\citep{rasio91, lombardi06}. These studies have been performed exclusively by using the smoothed particle hydrodynamics (SPH) method, except for the encounters involving a massive black hole and a star, which were calculated by using a three-dimensional Euler hydrodynamic code \\citep{khokhlov93a, khokhlov93b}. Recently SPH simulations are applied to the evolution of a giant planet through the tidal interactions with a sun-like star \\citep{faber04, ivanov04}.  The former hydrodynamic simulations, especially those of the tidal encounters between a red giants and a main sequence star by \\citet{davies91}, have been carried out with a relatively small number of SPH particles and for limited range of parameters. Their results are thought to be subject to limitations arising from low-mass resolutions since mass involved in the interactions decreases as the periastron distance increases, and the SPH method may not give a valid description of such situations where mass scales are as small as that allotted to each particles. In our problems of surface pollution, we deal with the accretion of mass $ \\sim 10^{-3} M_\\odot$. Simulations with finer mass resolutions, and hence with larger particle numbers, are necessary to investigate such encounters involving the transfer of mass of this order. It is also desirable to perform simulations for a wide range of parameters, such as the periastron distance, the red giant models in different evolutionary stages, and the mass of main sequence stars in order to obtain a realistic and general information about the characteristics of the tidal interactions.  In this paper, we first carry out simulations of tidal interactions between a red giant and a main sequence star by using SPH method. We make a detailed analysis of the amounts of energy and angular momentum, transferred from the orbital motion to the oscillation and spin of red giant, and the amount of mass, lost from the red giant and accreted onto the main sequence stars. Based on the numerical experiments, we then attempt to clarify the parameter dependence of these characteristics and to formulate the quantitative outcome as simple functions of the stellar parameters and impact parameter of encounters. The present results are applied to investigate the relevance of the scenario that the star-star interactions give rise to the abundance anomalies observed among not only giants but also main sequence stars in globular clusters. The derived formulae will be useful to perform simulations of dynamical evolution of stellar systems with the effects of stellar interactions taken into account.  The organization of the paper is as follows. In next section, we describe our numerical methods, including the set-up of the initial conditions, the models of the red giant, the treatment of accretion and the determination of the viscosity of red giant models. In \\S 3, we present the results from our simulations with the discussion of the non-linear effects of tidal interactions. In \\S 4, we derive fitting formulae for the energy and angular momentum deposited into the envelope of the red giants and the mass accreted onto the main sequence stars during the tidal encounters. The conclusions follow in \\S 5, with the discussion about the application to the globular clusters. ",
        "conclusions": "We have performed the SPH simulations of tidal encounter of red giants with environment stars and investigate the characteristics of stellar interactions for a variety of sets of parameters, the evolutionary stages of red giant, the mass of perturber stars and the assumed strength of viscosity as well as the orbital parameters of encounter. Based on our results and the other extent models, we discuss the dependences of interactions in the non-linear regime on the stellar and encounter parameters, and proposed formulae to describe the energy and angular momentum deposition to red giants, and the mass accretion onto the perturber stars in simple and convenient forms as a function of these parameters. Our main quantitative results are as follows;   \\begin{enumerate} \\item We obtain the both amounts of energy and angular momentum, transferred from the orbital motion into the oscillation and rotation of red giants, during tidal encounters by numerical simulations. The angular momentum deposited in the red giants can be large enough to rotate the envelope at rate $\\Omega \\gtrsim 0.01 \\Omega_{\\rm RG}$ for the encounter of periastron distance $r_p / R_{\\rm RG} \\lesssim 2.5$ (or the impact parameter $b / R_{\\rm RG} \\lesssim 15.7$) and hence for such encounters that ended in the two stars flying apart. For still closer encounters, it increases to give the rotation rates significantly exceeding $\\Omega \\simeq 0.1 \\Omega_{\\rm RG}$. Accordingly, the tidal encounter works as the source of angular momentum necessary to trigger rotational mixing in the red giants and also to explain the origin of fast rotators observed among the horizontal branch stars. For larger radius of red giant model, and hence, later stage of evolution, the transferred angular momentum increases while the energy deposition decreases since the transferred quantities are scaled with the stellar parameters. The fitting formulae are derived to describe these quantities as a function of the mass and radius of red giants, subject to the perturbation, the mass of main sequence stars as the perturber, and the impact parameter. Further we show that these transferred quantities, when normalized with respect to the momentum of inertia of models, are given solely as functions of the encounter closeness parameter $\\eta$, Our fitting formulae agree well with the results of encounter simulations by other authors, and can even reproduce the results for the encounters of main sequence stars approximated to a polytrope of index $N = 1.5$.  \\item With aid of fine mass resolution, we demonstrate that the main sequence stars can capture gas from the red giant envelope sufficiently to disguise their surface with accreted matter even for the encounters that ended in fly-by with an evolved red giant. The accreted mass onto the main sequence as a perturber during the tidal encounters is shown to be in a direct relationship with the energy deposition into the red giants. We also derive the formula, which predicts the accreted mass as a function of impact parameters for given stellar parameters and are applicable to the encounters involving not only the red giants but also the main sequence stars. \\end{enumerate}  The derived formulae are useful in determining the periastron distance of the tidal capture limit for the encounter of various model parameters. They also be useful in inquiring whether some stellar objects in the globular clusters, for example the red giants and the main sequence stars with abundance anomalies and the fast rotating horizontal brunch stars, which cannot be explained through the framework of the normal stellar evolution, can be produced through the stellar interactions. In our computations, the mass accretion rate may exceed the Eddington limit (${\\dot M}_{\\rm Edd} = 4 \\pi c R_{\\rm MS} / \\kappa_e$) on the surface of main sequence star for very close encounters, but since it remains below the Eddington limit at the accretion radius, we may well assume that the accreted mass mostly settles on the surface of main sequence stars with loosing thermal energy. In the following, we discuss the application of these formulae and the possibility that such stellar objects have their origins in the stellar encounters.   \\subsection{Tidal Capture Limit and Comparisons with Other Works}  In our simulations, the tidal capture limits, $r_{p, \\rm cap}$ in periastron distance and $\\eta_{\\rm cap}$, in $\\eta$, are estimated from the condition that $\\Delta E (t_E) = (1/2) \\mu v_\\infty (10 \\hbox{ km s})^2$ with use of eq.~(\\ref{eq:Ecrit}) or eq.~(\\ref{eq:delEfit-unique}). Our estimates are given in Table~\\ref{tb:capturelimit} for the two red giant models with the perturber masses of $0.6 M_{\\sun}$, $0.8 M_{\\sun}$ and $1.4 M_{\\sun}$. For the red giant of larger radius, the tidal capture limits decrease slightly ($\\sim 10 \\% $) when normalized with respect to the radius of red giant, but increase in the physical dimensions nearly in proportion to the surface radius; they slightly increase with the mass of perturber. For the encounter of $M_{\\rm RG} = 0.8M_{\\sun}$ red giant with a $M_{\\rm MS}=0.6M_{\\sun}$ main sequence star, \\citet{davies91,davies92} give the periastron distances for the tidal capture limit in the range of $2.00 < r_{p, \\rm cap} / R_{\\rm RG} < 2.25$, and our estimate, $r_{p, \\rm cap} / R_{\\rm RG}= 2.1$, as seen from Table~\\ref{tb:capturelimit}, resides in their range. On the other hand, our estimates turn out to be larger by $\\sim 20\\%$ than those obtained from the linear analysis by \\citet{mcmillan90}, as listed in the table.  On the other hand, \\citet{khokhlov93a} and \\citet{khokhlov93b} compute the encounter of a polytrope star of mass $M_{*}=0.8M_{\\sun}$ with a black hole of mass $M_B \\gg M_*$ for various values of polytropic index (the relative velocity is 100 km/s at infinity). The tidal capture limit decreases from $\\eta_{\\rm cap}= 2.75$ for $N = 1.5$ to 2.20 and 1.55 for $N = 2.0$ and 3.0, respectively. This demonstrates that for larger polytrope index, the energy deposition rate of the star becomes smaller, because of the increase in the mass concentration toward the center and of decrease in the moment of inertia for a given mass and radius.  \\subsection{Relevance to the Origin of Stars with Anomalous Abundances in Globular Clusters}  We may apply our fitting formulae to examine the possibility that the abundance anomalies observed both for red giants and main sequence stars in some globular clusters can be explained in terms of the stellar interactions during the close encounters. As the origin of these objects, we propose the following scenario; 1) The anomalies of red giants are generated through the flash-assisted deep mixing mechanism that is triggered by the injection of angular momentum into their envelope during the encounter with environment stars \\citep{fujimoto99}: 2) The main sequence stars gain the abundance anomalies as a result of the surface pollution by accreting matter from the red giants which have already developed these anomalies. We evaluate whether these work under the conditions prevailing in globular clusters.  First as for the point 1), \\citet{fujimoto99} argue that the angular momentum of $\\Delta L_{\\rm RG} / L_{\\rm crit} \\gtrsim 1/100$ is necessary to induce the flash-assisted deep mixing, and the transfer of angular momentum of this order occurs during the encounter of $\\eta \\lesssim 3$, and hence, $r_p \\sim 2.6 R_{\\rm RG}$ from our formulae. The transfer of angular momentum may have relevance to the a bimodal distribution of rotation velocity that horizontal branch stars display with the fastest rotators distributed on the cooler (redder) side of the branch whereas with the slower rotators spread over wider range on the branch. \\citet{suda06} suggest that the different modes of helium mixing mechanism may result according to when the stars undergo the close encounter and the deposition of angular momentum on the RGB and influence the horizontal branch morphology; the injection of angular momentum may invoke hydrodynamical instabilities due to differential rotation and invokes turbulent mixing to trigger the hydrogen-flash driven deep mixing, and the resultant helium enrichment in the envelope accelerates the evolution of RGB. The stars that experience the helium enrichment at an earlier epoch on the RGB have a smaller mass of helium core and hence, located on redder side of HB and those with later mixing epoch shifts to blue-ward: on the other hand, if the stars experience the close encounter near the tip of RGB, the helium-flash driven mixing, rather than the hydrogen-flash driven deep mixing, takes places and causes the largest decrease of helium core so that the stars are situated at the reddest end on the HB. If the stars experience close encounter at an earlier stage of RGB, then, they become slow rotators, due to angular momentum loss through mass loss on the RGB, and hence, settle on redder side of horizontal branch, and the stars, if experiencing it at later stage of RGB, become faster rotators and settle on bluer side. Finally, the stars, which undergo close encounters very close to the tip of RGB, become the fastest HB rotators, located on redder-most side of the branch. The close encounters at $\\eta \\gtrsim 3$ can explain the fastest rotation rates observed from HB stars of $\\Omega \\sim 0.1 \\Omega_{\\rm K}$ if the angular momentum is conserved during the contraction from the RGB to the HB. While the pristine angular momentum is lost effectively for such low mass stars, such fast rotation as observed for HB stars may be expected also from the synchronization of red giants in the binary systems of separations (more than a several au), and yet, it is difficult for such binaries to survive without suffering encounters in the dense stellar environment of globular clusters (see below eq.~[\\ref{eq:tau-enc}]). In other word, we may take the existence of HB stars of these fastest rotations as an evidence that such close encounters as $\\eta \\lesssim 3$ take place in these clusters.  Next as for the point 2), since the mass in the surface convective zone is of $\\sim 3 \\times 10^{-3} M_\\odot$, the accreted mass of an order of $\\sim 10^{-3} \\msun$ suffices to disguise the surface abundances with those transferred from the red giant envelope with anomalous abundances. From the present results, it is possible to estimate the range of periastron distance, $\\eta$, that can allow the accreted mass of this order at $\\eta =1.8 \\sim 2.2$ and $1.5 \\sim 2.0$, which correspond to $r_p = 1.9 \\sim 2.1 R_{\\rm RG}$ and $1.7 \\sim 2.0 R_{\\rm RG}$, for $R_{\\rm RG} = 20 R_\\odot$ and $85 R_\\odot$, respectively, for masses range of main sequence stars of $0.6 M_\\odot$ and $0.8 M_\\odot$.  The timescale of tidal interactions in the environment where the stellar density is $n_f \\hbox{ pc}^{-3}$ and the velocity dispersion $v_{\\ \\infty}$, is estimated (with the gravitational focusing taken into account, because of low velocity dispersion of environment stars in the core) at: \\begin{eqnarray} \\tau_{\\rm enc} \\sim 7.4 \\times 10^{9} (\\frac{10^{4} \\hbox{ pc}^{-3}}{n_f}) (\\frac{100R_{\\sun}} {r_p}) (\\frac{v_{ \\infty}}{10 \\hbox{km s}^{-1}})(\\frac{M_1+M_2}{2M_{\\sun}} )^{-1} {\\rm yr} \\label{eq:tau-enc} \\end{eqnarray} On this basis, the timescales of tidal encounters, which can bring about the mass accretion of a order of $10^{-3}M_{\\sun}$ and the angular momentum transfer of $\\Delta L_{\\rm RG} / L_{\\rm crit} \\gtrsim 1/100$, are estimated at $\\sim 1.5 \\cdot 10^{10}$ yr for $20 R_{\\sun}$ and $\\sim 3 \\cdot 10^{9}$ yr for $85R_{\\sun}$ in the environment of $n_f = 10^{4} \\hbox{ pc}^{-3}$ and $v_{\\infty} = 10 \\hbox{ km s}^{-1}$. These time-scales seem to be too large to explain the observed these objects by tidal interactions as compared with the lifetimes on the corresponding stages of red giants, $\\sim 10^8$ yr for $20R_\\sun$  and $\\sim 10^7$ yr for $85R_\\sun$, and the accumulated number of close encounters, attendant with the abundance anomalies, seems to be no more than a few, too small to explain the observations even in rough estimates. In order for the encounters to be viable, these time-scales have to be shorter more than an order of magnitude, and accordingly, there need some mechanism(s) to enhance the frequency of tidal encounter in globular clusters.  \\citet{sugimoto96} points out the importance of mass segregation in modeling dynamical evolution of star clusters to explain the observable number of mill-second pulsars in 47 Tuc; without the mass segregation, theoretical estimation indicates that the formation probability of binary with a neutron star is higher in $\\omega$ Cen having no collapsed core than 47 Tuc having collapsed core, although in actuality, the former has not been reported to contain any pulsars. Moreover, mass segregation promotes the core collapse and make it more rapidly than in the case of single-mass component. \\citet{portegies01} confirm that mass segregation enriches the core of star cluster in giants and white dwarfs by N-body simulations for an open star cluster. In order to explain the observable number of giants with abundance anomalies, about a half of total number of giants have to experience tidal encounters with field stars. It is necessary to see whether the mass segregation can gather most of giants, though not all, in the core and can increase the two-body encounter rate by more than an order of magnitude. In addition to the mass segregation, the gravo-thermal oscillations of star clusters, which has been proposed by \\citet{bettwieser84} and confirmed by N-body simulation \\citet{makino96}, may influence the rate; since the interactions are expected to occur in the evolutional stage near the high density peak, it is necessary to investigate in a correct manner how deeply and how long such a high density state reaches and lasts to affect the encounter rates.  For the encounter with the red giants of later stages, the pollution of main sequence stars are possible even when the two stars fly apart after the encounters. For the encounter that ends in the formation of binary, it depends on the fate of two stars whether the abundance anomalies imprinted onto the companions are observable or not, and it is necessary to pursue the details of evolution of tidally captured binaries. There are many effects to influence the binary evolution, such as the spin-up and mode damping rates of red giant after the encounter,  the mass transfer to the companion at subsequent periastron passages \\citep[ex.][]{sep07}, the mass loss from the bound system, and also the expansion of red giant as it ascends the red giant branch. In addition, we should also take into account the tidal interactions of formed binaries with the environment stars. It is conceivable as the destination of tidally captured binary that either one component is liberated through the exchange encounter with a third body, or the two stars eventually coalesce as a result of Roche lobe overflow. The possibility of the exchange event depends on the evolution timescale of giants and the encounter timescale; since the exchange encounter rate is proportional nearly to a semi-major axis of a binary under the assumption of point mass limit \\citep{ heggie96}, then the former is expected to occur more frequently. Accordingly, if tidal capture of two body encounters can contribute the modification of stellar populations, then, the exchange encounters follow at larger rates and make greater contributions. This increase in the encounter rates with red giant in the binary may affect the statistics of the main sequence stars with the abundance anomalies. Further, through the exchange encounters, the red giants and horizontal branch stars, now losing the mass due to the mass loss and becoming lighter than the heaviest main sequence stars, can be ejected from the binary systems to fly apart as single stars.  In summary, the enhancement of tidal interactions, necessary to explain the observed abundance anomalies, is expected to be provided by the formation of high density core due to the gravo-thermal oscillations and by the mass segregations which enlarge the fractions of stars of the upper mass end in the core. The proper understandings of these effects wait for N-body simulations of globular clusters with the stars of multi-mass spectra taken into account since star-star interactions as studied in the present work will play a critical role. In these studies of dynamical evolution of stellar systems, the formulae for the transfer of energy, angular momentum, and accreted mass derived in the present work serve the purpose for incorporation these effects into the simulations. It is also necessary to pursue the binary evolution and accurately explore the fate of red giants and subsequent horizontal branch stars with the interactions with the environment stars taken into account."
    },
    "0801/0801.2938_arXiv.txt": {
        "abstract": "% Edge-on spiral galaxies offer a unique perspective on disks. One can accurately determine the height distribution of stars and ISM and the line-of-sight integration allows for the study of faint structures. The Spitzer IRAC camera is an ideal instrument to study both the ISM and stellar structure in nearby galaxies; two of its channels trace the old stellar disk with little extinction and the 8 micron channel is dominated by the smallest dust grains (Polycyclic Aromatic Hydrocarbons, PAHs). \\cite{Dalcanton04} probed the link between the appearance of dust lanes and the disk stability. In a sample of bulge-less disks they show how in massive disks the ISM collapses into the characteristic thin dust lane. Less massive disks are gravitationally stable and their dust morphology is fractured. The transition occurs at 120 km/s for bulgeless disks. Here we report on our results of our Spitzer/IRAC survey of nearby edge-on spirals and its first results on the NIR Tully-Fischer relation, and ISM disk stability. ",
        "introduction": "For 32 edge-on spiral galaxies, spanning Hubble type and mass, we fit the edge-on infinite disk model by \\cite{vdKruit81a} on the IRAC mocaics: $\\mu(R, z) = \\mu (0,0) ~ \\left({R/h}\\right)K_1\\left({R/h}\\right) ~ sech^2\\left({z/z_0}\\right)$. We applied stellar dominated 3.6 and 4.5 $\\mu$m  channels and the PAH emission at 8.0 $\\mu$m, with the stellar contribution subtracted \\citep[using the prescription from ][]{Pahre04a}. The data is from our dedicated GO program (GO 20268) and the {\\em Spitzer} archive. Our program aims to populate the lower range of disk masses. From the fit we obtain the scale-length ($h$), the scale-height ($z_0$), and the face-on central surface brightness ($\\mu_0$). The edge-on disk model fits the stellar channels very well but the PAH channel less so because of star-formation structures, such as HII regions, in the PAH maps. The disk's total luminosity is inferred from the fitted model: $L_{disk} = 2 \\pi h^2 \\mu_0$, with $h$ the scale-length and $\\mu_0$ the face-on central surface brightness \\citep[See][]{Kregel02}. Combined with the rotational velocities from HyperLeda, we can construct the Tully-Fisher relation for our galaxy disks and a plot of disk oblateness as a function of rotational velocity (and hence dynamical mass). ",
        "conclusions": "Hence, from our disk models of IRAC data, we conclude:  \\begin{itemize} \\item[1] The Tully-Fisher relation has a shallower slope ($\\alpha=3.5$) than naively expected from trends with filter (Figure \\ref{f:tf}), which may be a metallicity effect of the stellar population's M/L. \\item[2] The PAH disk is flatter than the stellar one, especially in more massive spirals (Figure \\ref{f:hz}). The effect is not as pronounced as found by \\cite{Dalcanton04} and \\cite{Obric06} based on dust lanes. \\end{itemize}"
    },
    "0801/0801.3425_arXiv.txt": {
        "abstract": "Einstein's equations of gravitation are not invariant under geodesic mappings, i. e.  under a certain class of mappings of the Christoffel symbols and the metric tensor  which leave the geodesic equations  in a given coordinate system invariant. A theory in which geodesic mappings play the role of gauge transformations is considered. ",
        "introduction": "If we start with Einstein's beautiful hypothesis that test particles in gravitational field move along geodesic lines of some Riemannian spaces $V_{4}$, it is natural to expect that the differential equations for finding the metric tensor $g_{\\alpha\\beta}(x)$ for a given distribution of matter also should be invariant under any transformations at which the geodesic equations remain invariant. However, the geodesic equations are invariant not only under arbitrary transformation of coordinates (it is rather obvious) but also under geodesic mappings $\\Gamma_{\\beta\\gamma}^{\\alpha}(x)\\rightarrow\\overline{\\Gamma}_{\\beta\\gamma}^{\\alpha}(x)$ of the Christoffel symbols in any fixed coordinate system  \\cite{Weyl}-\\cite{Eisenhart}.  If $\\Gamma_{\\beta\\gamma}^{\\alpha}$ are Christoffel symbols in some coordinate system $(x^{0}, x^{1}, x^{2}, x^{3}) $, and we use coordinate time $t=x^{0}/c$ ($c$ is speed of light) as a  parameter along geodesic lines, then the differential equations of  a geodesic line are of the form \\begin{equation} \\ddot{x}^{\\alpha}+(\\Gamma_{\\beta\\gamma}^{\\alpha}-c^{-1}\\Gamma_{\\beta\\gamma }^{0}\\dot{x}^{\\alpha})\\dot{x}^{\\beta}\\dot{x}^{\\gamma}=0, \\label{EqMotionOfTestPart}% \\end{equation} where $\\dot{x}^{\\alpha}=dx/dt,$ $\\ddot{x}^{\\alpha}=d\\dot{x}^{\\alpha}/dt$. It easily to verify that these equations are invariant under the mapping \\begin{equation} \\label{GammaGeodesTransformations}\\overline{\\Gamma}_{\\beta\\gamma}^{\\alpha }(x)=\\Gamma_{\\beta\\gamma}^{\\alpha}(x)+\\delta_{\\beta}^{\\alpha}\\ \\phi_{\\gamma }(x)+\\delta_{\\gamma}^{\\alpha} \\phi_{\\beta}(x), \\end{equation} where $\\phi_{\\alpha}(x)$ is  a continuously differentiable vector field. In Riemannian space-time $\\phi_{\\alpha}(x)$ can be expressed in terms of determinants $g$ and $\\overline{g}$ of the metric tensors before and after the geodesic mapping of space-time $V$ into $\\overline{V}$ as follows:%  \\begin{equation} \\phi_{\\alpha}=\\frac{1}{n+1}\\left(  \\Gamma_{\\alpha\\gamma}^{\\gamma}-\\overline {\\Gamma}_{\\alpha\\gamma}^{\\gamma}\\right)  =\\frac{1}{2(n+1)}\\frac{\\partial }{\\partial x^{\\alpha}}\\ln\\left\\vert \\frac{\\overline{g}}{g}\\right\\vert. \\end{equation}   Such mappings of the Christoffel symbols in a given coordinate system induce  some transformations not only  of the curvature and Ricci tensors  but also  transformations $g_{\\alpha\\beta}\\rightarrow\\overline{g}_{\\alpha\\beta}$ of the metric tensor  which are obtained by solving of the partial differential equation \\begin{equation} \\overline{g}_{\\beta\\gamma;\\alpha}(x)=2\\phi_{\\alpha}(x)\\overline{g}% _{\\beta\\gamma}(x)+\\phi_{\\beta}(x)\\overline{g}_{\\gamma\\alpha}(x)+\\phi_{\\gamma }(x)\\overline{g}_{\\alpha\\beta}(x), \\label{gGeodesTransformation}% \\end{equation} where the semicolon denotes a covariant derivative with respect to the metric in $V_{4}$.  It is clearly that all solutions of these equation are, in principle, equivalent physically. However  Einstein's equations even in vacuum are not invariant with respect to geodesic mappings \\cite{Petrov}.  It follows from (\\ref{GammaGeodesTransformations}) that the components $\\overline{\\Gamma}^{i}_{00}=\\Gamma^{i}_{00}$. Therefore, in Newtonian limit geodesic-invariance is not an essential fact. However this cannot be said about the relativistic case. For this reason is of interest to explore a theory in which the gravitation equations, as well as the equations of motion of test particles, are geodesic- invariant i.e. in which geodesic invariance plays the role of gauge invariance.  In Sect. 2 a geodesic-invariant generalization of Einstein's equations is considered. It turns out that the simplest equations of such a kind are some bimetric equations.  In Sects\u02d9 3 to 7, proceeding   from new sight at  Poincar\\'{e}  old fundamental ideas, we argue that the both spaces, Riemannian (with the curvature other than zero) and \ufb02at,  can have physical sense depending on which kind of  reference frame is used. In Sects. 8,  9 we find the spherically-symmetric solution of the above equations. Finally, in Sect. 10 it is shown that features of the gravitational force resulting from this solution, well explain properties of the Hubble diagram obtained from the latest observations. ",
        "conclusions": "The necessity of a discussion of  the physical description of gravity through geodesic-invariant equations is quite obvious. However, the satisfactory implementation of such a programme is a difficult task. The equations of gravitation considered in the given paper do not contradict observation data, and their some predictions seem rather attractive.     However it is still unclear  whether they are  only possible equations. Besides, it appears that the correct equations of this type should be the bimetric ones. However physical interpretation of such bimetricity involves    deep problems of the space-time theory which presently have no  a satisfactory understanding and a complete solution."
    },
    "0801/0801.4519_arXiv.txt": {
        "abstract": "{We compare magnetic field extrapolations from a photospheric magnetogram with the observationally inferred structure of magnetic loops in a newly developed active region. This is the first time that the reconstructed 3D-topology of the magnetic field is available to test the extrapolations.  We compare the observations with potential fields, linear force-free fields and non-linear force-free fields.  This comparison reveals that a potential field extrapolation is not suitable for a reconstruction of the magnetic field in this young, developing active region.  The inclusion of field-line-parallel electric currents, the so called force-free approach, gives much better results. Furthermore, a non-linear force-free computation reproduces the observations better than the linear force-free approximation, although no free parameters are available in the former case. ",
        "introduction": "Due to the low plasma $\\beta$ the magnetic field is the dominating quantity in the low solar corona. Thus, the 3-D magnetic field structure is of basic importance for physical processes in the solar atmosphere, such as flares, coronal mass ejections and X-ray jets.  Direct observations of chromospheric and coronal magnetic fields are difficult, but significant progress has been made within the last few years, e.g. \\cite{lee99,lin00,kundu01,white02,raouafi02,solanki03,lagg04,lin04}. Here we compare magnetic loops reconstructed from magnetic field measurements by \\cite{solanki03} with magnetic fields extrapolated from a photospheric magnetogram. The measured fields allow a much more sensitive test of extrapolations than other observations. ",
        "conclusions": "\\label{results} We have compared the observationally inferred structure of magnetic loops in the upper chromosphere with magnetic fields extrapolated from photospheric measurements.  We find that the simplest model, potential fields, is not sufficient to reproduce the observations.  The inclusion of field line parallel currents, the so called force-free approach, gives much better results. Among force-free models, a linear one gives less accurate results than the non-linear force-free extrapolation.  We find that the observed and extrapolated loops agree quite well for almost $2/3$ of the loops, while the remaining $1/3$ might suffer from the limited field of view of the available vector magnetogram.  The investigated active region is quite young. With a vertical upflow speed of $v=1.5{^+_-}0.5 {\\rm km/s}$ at the loop apex and a maximum loop height of $10 {\\rm Mm}$ we can estimate the time elapsed since the loop tops first emerged as $2 {\\rm h}{^+_-}40 {\\rm min}$ assuming a constant rise speed. A horizontal shear flow of $1 {\\rm km/s}$ on the photosphere would give a shear of $7.2 {\\rm Mm}$. This value is comparable to the difference of the footpoint locations between potential field loops and observed loops. It is therefore not clear, whether the electric current has been caused by shear flow motion or if the magnetic loops already contain the current during their emergence. The rise of the loops may also explain some of the discrepancy between observed and extrapolated loops, since the loops may change somewhat during the time needed for the instrument to scan the region."
    },
    "0801/0801.3080_arXiv.txt": {
        "abstract": "We discuss the stability of the higher-dimensional de Sitter (dS) brane solutions with two-dimensional internal space in the Einstein-Maxwel theory. We show that an instability appears in the scalar-type perturbations with respect to the dS spacetime. We derive a differential relation which has the very similar structure to the ordinary laws of thermodynamics as an extension of the work for the six-dimensional model \\cite{ksm}. In this relation, the area of dS horizon (integrated over the two internal dimensions) exactly behaves as the thermodynamical entropy. The dynamically unstable solutions are in the thermodynamically unstable branch. An unstable dS compactification either evolves toward a stable configuration or two-dimensional internal space is decompactified. These dS brane solutions are equivalent to the accelerating cosmological solutions in the six-dimensional Einstein-Maxwell-dilaton theory via dimensional reduction. Thus, if the seed higher-dimensional solution is unstable, the corresponding six-dimensional solution is also unstable. From the effective four-dimensional point of view, a cosmological evolution from an unstable cosmological solution in higher dimensions may be seen as a process of the transition from the initial cosmological inflation to the current dark energy dominated Universe. ",
        "introduction": "Six-dimensional braneworld models have attracted particular interests in recent years\\cite{p}. From the cosmological aspects, these models may be useful as a way of resolution of the cosmological constant problem, since a codimension two brane helps the sideslip of the brane vacuum energy into the bulk \\cite{cc+}. It has been pointed out that the original proposal of the resolution of this problem actually does not work well \\cite{cc-}. Nevertheless, six dimensional models has been recognized as an important playground to unsderstand cosmology and gravity in higher-dimensional theory with non-trivial fluxes. The flux stabilization of extra dimensions would be a powerful tool to obtain realistic phenomenology and cosmology from string theory. In the simplest realization of the flux compactifications in six dimensions, the internal space has the shape of a rugby ball \\cite{cc+,cc-}, where codimension two branes are located at the positions of the poles. The warped generalizations of the rugby ball solutions have also been reported in the context of the six-dimensional Nishino-Sezgin (Salam-Sezgin), gauged supergravity \\cite{ggp} (see \\cite{ns} for the original supergravity) and pure Einstein-Maxwell theory \\cite{msyk}. \\footnote{In Ref. \\cite{cdgv}, a study on warped codimension-two braneworld solutions was presented on the analogy of the classical mechanics.}   It also has been recognized that a 3-brane in six or higher dimensions generically have the problems on localication of matter on the brane due to its stronger self-gravity. One well motivated way to circumvent this problem is to regularize the brane, by taking the microscopic structure of the brane into account. Several ways of regularization of codimension two branes have been proposed in \\cite{thick,pst,vc}. Based on these regularizations, low energy cosmology \\cite{vc,cosmo,cck} and effective gravity on the brane \\cite{grad} have been studied. On the other hand, the exact solutions \\cite{6dt, km, shock,btz} will help to obtain unique observational/experimental predictions from six dimensions.    Stability of six-dimensional flux compactifications is an important issue and several analyses of linear perturbations have been reported. It has been reported that the Minkowski brane solutions in the supergravity \\cite{ns} are marginally stable \\cite{st1} and those in the Einstein-Maxwell theory are stable \\cite{st2}. On the other hand, in the de Sitter (dS) brane solutions in the Einstein-Maxwell theory an instability appears in the scalar sector of perturbations with respect to the symmetry of dS spacetime for a relatively higher brane expansion rate \\cite{ksm}. This type of instability is commonly known in the dS spacetime~\\cite{kp} with an internal space compactified by a flux~\\cite {fr} \\footnote{Other instabilities in the quadrupole or higher multipole modes were reported, which gives rise to deformation of the internal space geometry, see e.g., Ref. \\cite{martin}. However, these instabities appear in dS compactifications with more than four-dimensional internal space and are not relevant for the models discussed in this paper.} We will see that such an instability also appears in the dS brane solutions in higher-dimensional Einstein-Maxwell theory. The important fact is that a class of the six-dimensional Einstein-Maxwell-dilaton theories has the equivalent structure to that of the higher-dimensional Einstein-Maxwell theory via dimenional reduction \\cite{km}. Thus instability of a solution in the higher dimensional theory suggests that of the corresponding solution in six-dimensional theory.     Also in Ref. \\cite{ksm}, an important relation which has very similar structure to the ordinary laws of thermodynamics was found in the dS brane solutions in the six-dimensional Einstein-Maxwell theory. In this relation, the area of dS horizon (integrated over the internal space) exactly behaves as the usual thermodynamical entropy. It was shown that dynamically unstable solutions are also {\\it thermodynamically} unstable. This may be seen as an example to support the conjecture that claims the equivalence of these two instabilities \\cite{gm}, which originally has been discussed for black brane solutions. We will see that such a thermodynamical relation can be easily extended to the case the higher dimensional dS brane solututions.   The dynamical evolution from unstable dS flux compactifications \\cite{fr} have been investigated in e.g., Ref. \\cite{krish}. Based on these arguments, we will discuss the fate of unstable dS brane solutions and corresponding cosmological solutions in six dimensions. We will see that an initially unstable dS brane solutions evolves to another stable dS or anti-de Sitter (AdS) brane solution unless the internal space is decompactified.    This article is organized as follows. In the the section II, we review the higher-dimensional dS brane solutions and cosmological solutions in the general six-dimensional Einstein-Maxwell-dilaton theory via dimensional reduction from the dS brane solutions. In the section III, we analyze the stability and cosmological evolutions of the higher-dimensional dS compactifications. In the section IV, we discuss the possible fate of unstable solutions in the higher dimensional theory from the six-dimensional perspectives. In the section V, we shall close this article. ",
        "conclusions": "We discussed the stability of a de Sitter (dS) brane solutions in the higher-dimensional Einstein-Maxwell theory. We confirmed that the dS brane solutions with relatively high expansion rates are unstable against scalar-type linear perturbations with respect to the symmetry of dS spacetime. Such an instability commonly appears in the wide class of compactifications with an external dS spacetime and can be understood as a type of radionic instabilities.  Then, we generalized a relation found in the six-dimensional dS brane solutions Ref. \\cite{ksm}, which has very similar structure to the ordinary laws of thermodynamics, to the case of the higher dimensions as Eq. (\\ref{1st}). In this relation, the area of dS horizon (integrated over two internal dimensions) behaves the thermodynamical entropy. The area of dS horizon is essentially double-valued function of the flux and (one of) the brane tensions. A dynamically unstable solution is also in the thermodynamically unstable branch. The boundary between the thermodynamically stable and unstable branches is given by a curve, where one-to-one map from the plane of model parameters $(\\alpha,\\lambda)$, where $\\alpha$ controls the degree of warp and $\\lambda$ characterizes the curvature of dS spacetime, to the conserved quantities breaks down. These are the generalizations of the results of the case of six dimensions discussed in \\cite{ksm} to higher dimensions.    Then, we discussed the possible cosmological evolutions. We firstly focused on the case $\\alpha=1$, the case of an exactly rugby ball shaped bulk. In this case, during the cosmological evolution, the bulk keeps the rugby-ball type shape. There are two possibilities of the evolutions of unstable cosmological solutions: One possibility is to settle down a stable configuration. The other one is that the compactified two extra-dimensions are decompactified. This strongly depends on the initial condition and if the internal space is initially shrinking, the final configuration is another stable dS brane solution. Then, the size of the internal space in the final configuration usually is smaller than the initial unstable configuration. Furthermore, if initially $\\lambda_{\\rm inst}<\\lambda<\\lambda_{\\rm AdS}$, where $\\lambda_{\\rm inst}$ and $\\lambda_{\\rm AdS}$ are given in Eq. (\\ref{yoko}) and Eq. (\\ref{ads}), the final configuration is also a dS brane solution, while $\\lambda_{\\rm AdS}<\\lambda<\\lambda_{\\rm max}$, the final configuration is an AdS brane solution. This picture can be easily generalized to the case of unstable dS configurations with non-trivial magnetic flux. In this case of the warped bulk $\\alpha<1$, the basic behavior remains the same as the rugby ball case: the initially unstable dS brane solution evolves toward the stable dS/AdS brane solutions (or decompactified). For the observer on the (+)-brane, during the cosmological evolution $\\phi/\\beta_-$ is conserved while for the observer on (-)-brane, $\\phi/\\beta_+$ is conserved. In general, the shape of the bulk of the final stable configuration is different from that of the initial stable one.   Finally, we have discussed the stability of brane cosmological solutions in the six-dimensional Einstein-Maxwell-dilaton theory Eq. (\\ref{6d_action}), including the Nishino-Sezgin (Salam-Sezgin), gauged supergravity as the special limit. This class of six-dimensional Einstein-Maxwell-dilaton theory has equivalent structure to the $(D+2)$-dimensional Einstein-Maxwell theory and the number of dimensions $D$ affects the dilatonic coupling in the reduced six-dimensional theory. Thus, if the seed dS brane solution is unstable in $(D+2)$-dimensional spacetime, the corresponding cosmological solution in six dimensions is also unstable: if the final configuration is another dS brane solution in $(D+2)$ dimensions, this also may be true in six dimensions. The cosmological evolution in higher dimensions may be seen as a process of the transition from the initial cosmological inflation to the current dark energy dominated Universe from the four-dimensional perspectives."
    },
    "0801/0801.1825.txt": {
        "abstract": "We present results from an extensive survey of 64 cavities in the X-ray halos of clusters, groups and normal elliptical galaxies. We show that the evolution of the size of the cavities as they rise in the X-ray atmosphere is inconsistent with the standard model of adiabatic expansion of purely hydrodynamic models. We also note that the majority of the observed bubbles should have already been shredded apart by Rayleigh-Taylor and Richtmyer-Meshkov instabilities if they were of purely hydrodynamic nature. Instead we find that the data agrees much better with a model where the cavities are magnetically dominated and inflated by a current-dominated magneto-hydrodynamic jet model, recently developed by \\citet{LiMHD1} and \\citet{LiMHD2}. We conduct complex Monte-Carlo simulations of the cavity detection process including incompleteness effects to reproduce the cavity sample's characteristics. We find that the current-dominated model agrees within $1\\sigma$, whereas the other models can be excluded at $>5\\sigma$ confidence. To bring hydrodynamic models into better agreement, cavities would have to be continuously inflated. However, these assessments are dependent on our correct understanding of the detectability of cavities in X-ray atmospheres, and will await confirmation when automated cavity detection tools become available in the future. Our results have considerable impact on the energy budget associated with active galactic nucleus feedback. ",
        "introduction": "One of the current most active areas of research in X-ray astronomy is aimed at solving the problem of feedback in elliptical galaxies, groups and clusters of galaxies. All of these systems suffer from the same unsolved problem that the gas in their cores has cooling times significantly shorter than their age, suggesting that the gas should already have cooled and formed stars \\citep{McNamaraAGNReview}. This problem, historically referred to as the ``cooling flow problem'', points toward the existence of a heat source that is preventing the gas from cooling effectively in the center. Due to a lack of better understanding, this effect is generally referred to with the generic term ``feedback''.  Many possible feedback mechanisms have been proposed in recent years. The most basic idea is the conversion of gravitational energy into thermal energy, as gas falls into the center towards a deeper gravitational potential. However, this mechanism has been shown to only be significant in elliptical galaxies with a rather steep gravitational potential, such as NGC~6482 \\citep{PonmanNGC6482}, and does not solve the cooling flow problem.  Another heat source that has often been suggested is feedback from supernovae. If the gas cools sufficiently, stars will form, which will then produce supernovae, which in turn will heat the ambient medium \\citep{BinneyCoolevolution}. While this effect is certainly important, the amount of energy released can only be sufficient to offset cooling in low-luminosity elliptical galaxies, but does not even come close to the radiative losses in groups or cluster, or even high-luminosity ellipticals. Even in normal elliptical galaxies, supernova feedback models have been shown to need fine-tuning to produce gas halos that resemble X-ray observations \\citep{MathewsReview}.  Instead, the central active galactic nucleus (AGN) has emerged as the most popular heat source in the last few years, though this still remains an active matter of debate in the community. \\citet{McNamaraAGNReview} give a recent review of AGN heating in clusters. {\\it Chandra} and {\\it XMM-Newton} observations \\citep[e.g.][]{McNamaraHydraA} have revealed the existence of giant cavities in clusters, which are often associated with radio emission from ongoing jet activity of the central AGN, actively inflating these ``bubbles''. This direct link has also been observed on the smaller group \\citep[e.g.][]{MoritaHCG62} and galaxy scales \\citep[e.g.][]{Fin01,DiehlGallery,DiehlAGN,DiehlEntropy}. AGN outbursts have been shown to have potentially more than sufficient energy to offset cooling in a large fraction of clusters \\citep{BirzanCavities}, and even in normal elliptical galaxies and groups \\citep{BestRadioloudAGN, BestAGNcooling}.  However, while the AGN has {\\it potentially} sufficient energy to offset cooling, the detailed dynamics of how the released energy can actually be transferred into thermal energy of the gas is poorly understood. A lot of theoretical work has been undertaken in the purely hydrodynamical regime \\citep[e.g.][]{BrueggenJets} to understand jet dynamics in inflating these cavities, but with limited success so far \\citep[see e.g.][]{VernaleoJetProblems}. The most realistic simulations to date have been conducted by \\citet{HeinzJets}, who employ the dentist's drill model combined with a dynamical cluster initial setup to avoid some of the issues other more idealized simulations have. But it remains to be seen how successful this model is on the less dynamic group and galaxy scale, where AGN heating has recently been found to be important as well \\citep[e.g.][]{DiehlGallery, DiehlAGN, BestRadioloudAGN, BestAGNcooling, FabianAccretionJet, ShurkinNGC1399_4649}  A problem of purely hydrodynamic jet simulations is that they ubiquitously produce a shock at the termination point of the jet, which results in hot, shocked gas surrounding the jet lobe. This is contrary to the vast majority of X-ray observations, though shocks around radio lobes have recently been reported in NGC~3801 \\citep{CrostonShocks} and Centaurus~A \\citep{KraftCentaurusA, KraftCentaurusA2}. Another major issue is that the created cavities often do not morphologically resemble X-ray observations, and get shredded apart quickly. Rayleigh-Taylor and Richtmyer-Meshkov instabilities are expected to grow on the top of the bubble and Kelvin-Helmholtz instabilities develop on the shear flow at the bubble sides, as it rises. Nevertheless, bubbles with rise times on the order of several times the predicted instability timescales have been observed, calling this scenario into question. Something has to suppress the growth of these instabilities. In general, tangential magnetic fields \\citep{DeYoungMHDBubbles}, viscosity \\citep{ReynoldsViscousBubble}, or continuous inflation of the bubble \\citep{SokerRT} have been proposed to suppress the growth of instabilities.  Alternatively, \\citet{LiMHD1} and \\citet{LiMHD2, LiMHD3} have proposed a magnetically dominated jet model instead. Their 3D magneto-hydrodynamic (MHD) simulations advocate a current-carrying jet that evolves purely by injecting magnetic flux and energy into a small volume surrounding the central black hole. The injected magnetic fields are not force-free and interact with the ambient medium to form a self-collimated jet. As the background pressure drops, the jet expands into a wide jet lobe, excavating a bubble in the ambient medium. Mock X-ray observations of these systems show a strong resemblance to observations of cavities in clusters \\citep{DiehlAlaska}. In particular, these cavities generally expand subsonically and appear cooler than the surrounding medium, similar to observations.  This paper addresses the question of whether one can use the bubbles' sizes  as they rise in the intracluster medium as a function of their distance to the cluster center as an indicator for which model may indeed be correct.  The paper is organized as follows: In section \\ref{s.data} we describe the methodology on how we identify the cavities, and extract their main properties. \\S\\ref{s.models} then presents the predictions of the various models that we are considering, and focus on their predictions for bubble sizes as a function of radius. In section \\ref{s.observations} we show how this relates to observations. We then discuss various incompleteness issues associated with the data in section \\ref{s.discussion}, before we try to reproduce the properties of our data set with Monte-Carlo simulations in \\S\\ref{s.montecarlo} and conclude (\\S\\ref{s.conclusions}). ",
        "conclusions": "We present the first systematic analysis of cavity sizes from cavities located in cluster, group and galaxy environments. We summarize predictions from four particular models: three purely hydrodynamic (AD53, AD43 and CIH) and two magnetically dominated models (FML and CDJ, see Table \\ref{t.models}). The two simplest hydrodynamic models AD43 and AD53 assume the bubble is formed a priori close to the center and its interior pressure due to gas pressure from a hot, low density gas. The bubbles then adiabatically expand during the buoyant rise in the intracluster medium. These two models differ only in the use of two different adiabatic indeces $\\Gamma=5/3$ and $4/3$. The CIH model also assumes adiabatic expansion, but rather continuously injects energy into the cavities. This requires the jet to stay connected to the cavities as they rise, which may be challenging to achieve in reality.  The first magnetically dominated model (FML) assumes that the magnetic field configuration consists of a dipole magnetic field generated by randomly oriented magnetic flux loops. We show that the predictions from this model are identical to the $\\Gamma=4/3$ hydrodynamic model, which makes a distinction of these models from cavity sizes alone unfeasible. Thus, one cannot possibly use cavity sizes to determine the fraction of pressure supplied by magnetic loops threaded through the bubble surface. The last model that we consider is a magnetically dominated model in which the jet carries a current (CDJ). In this model, the bubble size is set by the point where the interior magnetic pressure due to the flowing current is matched by the outside pressure.  Our analysis of the 64-cavity sample shows that the bubble sizes are much bigger than predicted from the adiabatic and the magnetic dipole model, and that they are more consistent with the continuously inflated hydrodynamic model and the current-dominated jet models instead. In particular we find that the inferred currents in these systems scale almost linearly with the core radius of the cluster and the central black hole mass.  We conduct Monte-Carlo simulations of the bubble detection process including incompleteness issues, which we also discuss in length. We show that the magnetically dominated model is always favored compared to the other four models. For the case that the initial bubble size is linked linearly to the core radius, the current-dominated model is consistent with the data within one standard deviation. Even with fine-tuning, we are unable to bring the other four models into quantitative agreement with observations, and never get closer than a $\\sim 5\\sigma$ offset. Any viable future jet model has to at least mimic the behavior of the CDJ model in terms of the size evolution of bubbles.  In order for the purely hydrodynamic model to work, we also need a mechanism to efficiently suppress instabilities. Viscosity \\citep{ReynoldsViscousBubble} or magnetic draping \\citep{RuszkowskiTangledFields} may provide this added stability. Continuous inflation may suppress the growth of Rayleigh-Taylor instabilities on top of the bubble \\citep{SokerRT}. A fine-tuned jet model continuously inflating cavities even at large radii may also result in a faster size evolution of the cavities. However, bringing such a model into agreement with constraints from both X-ray observations and old particle ages inside the cavities will be challenging.  We stress the fact that we do not consider our results as final. One of our key assumptions is that the bubbles are always in pressure equilibrium with the ambient gas. We base this assumption on the general lack of of shocked gas surrounding the bubbles. If this is incorrect and the bubbles are overpressured closer to the center, they would expand faster than the simple AD53 or AD43 models that we have considered here. To be a bit more quantitative, let us take a fiducial observed bubble of size $R_b$ at $r=2\\,r_c$ in a cluster with $\\beta=0.5$ and assume that it is currently in pressure equilibrium. If we further take the steeper slope of the CDJ evolution as a given, we can reproject the bubble size back to the center: $R_{b,0}=R_b\\,[1+(r/r_c)^2]^{-3/8}\\approx 0.55\\,R_b$. Since we know its pressure and volume now, as well as its size before, we determine the the bubble's internal pressure $p_{b,0}$ it must have had to start with: $p_{b,0}=p_b\\, (V_b/V_{b,0})^{4/3}=p_0\\, [1+(r/r_c)^2]^{-3/4} \\,(R_b/R_{b,0})^4$, assuming $\\Gamma=4/3$. This yields an overpressure of $p_{b,0}/{p_0}\\approx 3.34$. Such an overpressure would have driven a shock that should have been easily detectable. A systematic search for such shocks in clusters and galaxies is needed to assess whether all bubbles are overpressured close to the center, but beyond the scope of this paper. However, for the overpressure model to work, we must also explain how all clusters overpressure their bubbles in a very similar fashion to stay in agreement with the correlations discovered in this paper. So far, only Centaurus~A \\citep{KraftCentaurusA,KraftCentaurusA2} and NGC~3801 \\citep{CrostonShocks} show evidence of shocks surrounding their radio lobes. In addition, an initially overpressured bubble would have adjusted its bubble size on the order of a sound crossing time for the bubble, which is much shorter than the observed bubble ages.  We also do emphasize that the current assessment of incompleteness effects may be oversimplified and influence these results. In order to properly address this important issue, an automated bubble detection tool is desperately needed. Incompleteness has a considerable impact on the energy budget associated with AGN feedback, and has to be taken into account in further studies of the subject. We also suggest that the current analysis be repeated with a larger and more homogenous sample, as the predictions of the models will be much easier to distinguish from each other.  We provide a framework to test not only currently available models, but that can easily be extended to future models of jet-lobe systems. We encourage modelers to find theoretical considerations predicting cavity sizes as a function of distance to cluster centers, or to extract this information directly from their numerical simulations."
    },
    "0801/0801.4613_arXiv.txt": {
        "abstract": "We calculate the radiation spectrum and its time variability of the black hole accretion disk-corona system based on the three-dimensional magnetohydrodynamic simulation. In explaining the spectral properties of active galactic nuclei (AGNs), it is often assumed that they consist of a geometrically thin, optically thick disk and hot, optically thin corona surrounding the thin disk. As for a model of the corona, we adopt the simulation data of three-dimensional, non-radiative MHD accretion flows calculated by Kato and coworkers, while for a thin disk we assume a standard type disk.  We perform Monte Carlo radiative transfer simulations in the corona, taking into account the Compton scattering of soft photons from the thin disk by hot thermal electrons and coronal irradiation heating of the thin disk, which emits blackbody radiation. By adjusting the density parameter of the MHD coronal flow, we can produce the emergent spectra which are consistent with those of typical Seyfert galaxies.  Moreover, we find rapid time variability in X-ray emission spectra, originating from the density fluctuation produced by the magnetorotational instability in the MHD corona. The features of reflection component including iron fluorescent line emission are also briefly discussed. ",
        "introduction": "Thanks to the rapid progress in the observational studies in recent years, our understanding of the radiation properties of black hole accretion flows have been really deepened.  Accreting black holes, such as active galactic nuclei (AGNs) and black hole binaries (BHBs) during their very-high spectral state [state with luminosities around a few tenths of the Eddington limit ($L_{\\rm Edd}$)], show the radiation spectra dominated by two components; the thermal bump in UV/soft X-ray band and the power-law emission with a spectral index of $\\alpha\\sim 1$ in the X-ray band (possibly with a high energy cutoff around MeV; for the multiwavelength spectrum of a typical AGN, see Reynolds et al. 1997). These components are often explained by the disk-corona model (Liang \\& Price 1975; Bisnovatyi-Kogan \\& Blinnikov 1977; Haardt \\& Maraschi 1991, 1993).  In this model, the accretion flow consists of geometrically thin, optically thick accretion disk whose structure is studied by Shakura \\& Sunyaev (1973), and hot, optically thin corona surrounding the disk.  The thermal bump is believed as the thermal emission from the optically thick disk (see Kishimoto et al. 2005 for its observational implication), and the power-law component is interpreted to be formed by photons which are emitted from the disk and Compton up-scattered by hot electrons in the corona (see Thorne \\& Price 1975; Shapiro et al. 1976 but in the context of the inner hot accretion flow model).  Such a structure that the hot gas (i.e. the corona) coexists with the cool gas (i.e. the disk) is justified by the existence of the reflection component observed in the X-ray spectra  (Pounds et al. 1990; see also Guilbert \\& Rees 1988; Lightman and White 1988). This reflection component is accompanied by the iron fluorescent line emission broadened relativistically (for reviews, see Mushotzky et al. 1993; Fabian et al. 2000; Reynolds \\& Nowak 2003).  On the other hand, our theoretical understanding of black hole accretion has also made a rapid progress.  Now the dynamics of accretion flows is understood in terms of magnetohydrodynamics (MHD), since Balbus \\& Hawley (1991) rediscovered the magnetorotational instability (MRI) as the fundamental mechanism of angular momentum transfer in accretion disks. Detailed dynamical features of MHD accretion flows have been investigated via global three-dimensional numerical simulations (Matsumoto 1999; Stone \\& Pringle 2001; Machida et al. 2001; Hawley \\& Krolik 2001, 2002; Machida \\& Matsumoto 2003; Armitage \\& Reynolds 2003; De Villiers et al. 2003; Gammie et al. 2003; Igumenshchev et al. 2003). Such magnetically dominated accretion flows as those simulated in these calculations are radiatively inefficient accretion flows (RIAFs) whose mass accretion rates are much smaller than the critical value of $\\sim L_{\\rm Edd}$ (see, e.g., Mineshige et al. 1995). In such flows, the dissipated energy would not be radiated away efficiently, because of their low density, and be advected inward to the central black hole (Ichimaru 1977; Narayan \\& Yi 1994; Abramowicz et al. 1995; Kato et al. 1998).  Although the detailed structures and behavior of magnetohydrodynamical accretion flows are shown by numerical methods, it is still unclear that if these simulational results can be applied in realistic situations in the universe.  In recent years, the researches concerning the observational properties like the spectra and their time variabilities expected from numerically simulated accretion flows have been extensively performed.  Since the RIAF model is believed to fit the emergent spectrum of Sagittarius ${\\rm A}^*$, some authors have calculated the time dependent radiation spectra predicted from simulated MHD accretion flows with the aim to reproduce the observed behavior of Sgr ${\\rm A}^*$ (Hawley \\& Balbus 2002; Goldstone et al. 2005; Ohsuga et al. 2005; Moscibrodzka et al. 2007).  However, as for the accretion flows with moderately high mass accretion rate, which are considered to be a good model for AGNs and BHBs showing the thermal bump and the power-law emission in their spectra, no attempts have been made so far to calculate the spectra predicted from MHD simulations taking into account realistic radiation processes.  As noted in the above, in order to reproduce such spectra, the simulated accretion flows should have two components: geometrically thin cool disk and optically thin hot corona.  However, no simulation data which incorporate such two-component effects are available.  In this study we calculate for the first time the emergent spectra of two-component accretion flows based on the three-dimensional MHD simulation by Kato et al. (2004, hereafter KMS04).  Instead of fully solving the dynamics of two-component accretion flows, we adopt the simulation result of RIAF-like MHD accretion flow for the optically thin corona, and assume that the optically thick, geometrically thin disk is embedded in the corona with mutual interaction through radiation.  Other interactions, such as mass evaporation/condensation or heat conduction, are neglected, for simplicity (Meyer \\& Meyer-Hofmeister 1994; Meyer et al. 2000; Liu et al. 2002, 2007).  The disk is emitting soft photons with thermal spectrum, and those photons are up-scattered by hot electrons in the corona.  After being scattered in the corona, a part of upscattered photons return to the disk and are thermalized there, thereby heating the accretion disk.  The disk emission is, hence, enhanced by this returning process. In this paper we perform three-dimensional Monte Carlo radiative transfer simulations to properly calculate  such radiation processes and the emergent spectra.  The plan of this paper is as follows.  We show the detail of the model and method used in our calculation in \\S 2.  The results are presented in \\S 3, and compare them with the spectra observed in typical AGNs and discuss the similar points and discrepancies between them in \\S 4. In \\S 5 we summarize our study. ",
        "conclusions": "We have calculated the emergent spectra and their time variabilities predicted based on the disk-corona model, in which a cold standard disk at the equatorial plane is sandwiched by a hot coronal flow. As for the structure and dynamics of the corona we use the three-dimensional MHD simulation data by Kato (2004).  In this section we discuss our results especially in the context of AGNs.  We have shown that our black hole disk-corona system can reproduce the power-law X-ray emission with the photon index $\\alpha\\sim 1-2$ by adjusting the density parameter properly.  The power-law indices and the cutoff energy scales of the spectra are roughly in agreement with the observed spectra of Seyfert galaxies.  Moreover, we find significant variability of the power-law X-ray emission. The power-law X-ray emission flux predicted from our model changes by a few tens of percent on timescales of the orbital period near the last stable orbit, which is about $10^3(M/10^8M_{\\odot})~{\\rm sec}$, while the power-law index nor the cutoff energy do not change considerably. This variability comes purely from the fluctuation of the coronal flow around $r\\sim 20r_S$, where the scattering optical depth of the coronal flow attains its largest value in its structure.  This fluctuation is driven by the turbulence as a result of MRI, and its amplitude is large enough to explain the observed X-ray variability in Seyfert galaxies (e.g. Miniutti et al. 2007).  In UV/soft X-ray band, however, we cannot see the bump-like structure which is often found in the spectra of AGNs. The same problem occurred in other disk-corona models (Shimura et al. 1995; Liu et al. 2003). This bump is supposed to be the thermal radiation coming directly from the optically thick disk, while in those disk-corona models the disk is wholly covered with corona. In order to reproduce the UV/soft X-ray bump with a model, the coronal structure may need to be patchy or locally concentrated.  In the present study we do not consider the viewing angle dependence of the spectra.  If we observe this accretion flow system face-on, then the scattering optical depth along the line of sight would be smaller than in the case observed with non-zero viewing angle and the optically thick component should be observed more clearly.   Iron K$\\alpha$ fluorescent line emissions would occur when the disk is irradiated by X-ray from the corona.  The profile of this line emission would be broadened due to relativistic effects (see Reynolds \\& Nowak 2003 for a review).  The detailed structure of this line profile is determined by the line emissivity distribution on the disk, which depends on the spectra of local X-ray irradiation from the corona.  Recently {\\it Suzaku} has observed the broadened iron line profile from MCG-6-30-15 in detail and they implied that the emissivity profile should be as steep as $F_{\\rm line}(r)\\propto r^{-4.4}$ (Miniutti et al. 2007).  This is not easy to understand, since standard disks are known to produce continuum flux $F_{\\rm cont}(r)\\propto r^{-3}$.  Some authors proposed that this is a manifestation of Kerr black hole effects (e.g. Wilms et al. 2001). Such a steep emissivity profile can be reproduced by thermal Comptonization in the corona, if the irradiating corona has the temperature and density profiles which are arising inward, as was pointed out by Kawanaka et al. (2005). This is because the fraction of high energy photons increases inward so that the number of irradiating photons which are capable of producing iron fluorescent line emissions should increase inward more rapidly than $r^{-3}$. As for the MHD corona model which we adopted in this study, however, the temperature and density have flatter profiles (see KMS04), so such a steep line emissivity profile is not expected as long as we consider only Compton upscattering to be the X-ray emission process in this corona.  One of the way of producing such a steep emissivity profile so as to produces the iron line profile as is observed is to introduce moderate radiative cooling in the simulation (Mineshige et al. 2002). Radiative cooling can be efficient in dense region which is generally located in the inner part of the disk. Then the density profile near the event horizon of the coronal flow can be steeper and Compton up-scattering would be more efficient. As a result, hard X-ray photons irradiating the disk would increase and iron fluorescence line photons would become more efficient. Another possibility is that the transient X-ray emission accompanying with magnetic reconnection flares in the localized region of the corona would contribute to the iron line emissions from the disk, which is not taken into account in our present calculation. Such small-scale X-ray sources are often assumed when explaining the constant reflection component including iron line emissions (Miniutti \\& Fabian 2004; Malzac et al. 2006; Nayakshin 2007). If we analyze the local electron heating (or nonthermal acceleration) around magnetic reconnection flares and hard X-ray emission processes in the MHD coronal flow simulation, then we can know the local X-ray irradiation onto the disk, which may lead us to understanding the observed broad iron line profiles and their time variabilities from the microphysics of X-ray emission processes.  We can check the consistency of the assumption that the MHD coronal flow which we adopt in the calculation is cooled dominantly by advection, by comparing the heating rate ($\\sim GMm_p/(2r)\\Omega$) and the cooling rate ($\\sim$Comptonization luminosity per electron). According to the result of our calculation, the innermost region ($0<r\\lesssim 20r_{\\rm S}$) is not cooled efficiently by inverse Compton scattering.  However, in the outer region ($r\\gtrsim 20r_{\\rm S}$) local heating rate and cooling rate are comparable, which means that the assumption of an advection-dominated flow is only marginally justified. Generally when Compton cooling becomes efficient the coronal temperature will become lower, and then Compton upscattering will be inefficient.  This means that the power-law X-ray emission would not be generated so much from that region. As far as we consider the spectrum and energy flux in X-ray and higher energy band, to which the photons from the innermost region mainly contribute, this inconsistency would hardly affect the results.  Finally, we should mention the interaction between the corona and the underlying disk.  In the transition zone between the hot corona and the cold disk where density and temperature abruptly change, the heat conduction and the mass evaporation would be important as the mass and energy exchanging processes (Meyer \\& Meyer-Hofmeister 1994; Liu et al. 2002).  Evaporation of photospheric material was actually shown to be essential in the context of solar flares (see e.g. Yokoyama \\& Shibata 1994).  By including these effects in the simulation we will be able to obtain a more realistic model of the coronal structure and dynamics.  To what extent the disk-corona structure can extend to the inner region is important when we consider the relativistic skewing of the iron line emission profile because most of the observations of iron line profiles imply that the iron fluorescence line photons are emitted from around/inside the last stable orbit.  The detailed analysis of such disk-corona interactions and their observational implications are left as a future work.  \\bigskip We would like to thank Masayuki Umemura, Toshihiro Kawaguchi and Ken Ohsuga for helpful comments and discussions.  We are also grateful to an anonymous referee for his/her valuable comments, which helped us to improve the manuscript in a great deal. The numerical calculations were carried out on Altix3700 BX2 at YITP in Kyoto University.  This work is supported in part by Research Fellowship of the Japan Society for the Promotion of Science for Young Scientists (N. K.)."
    },
    "0801/0801.1170_arXiv.txt": {
        "abstract": "We examine the central-galaxy luminosity -- host-halo mass relation for 54 Brightest Group Galaxies (BGGs) and 92 Brightest Cluster Galaxies (BCGs) at $z<0.1$ and present the first measurement of this relation for a sample of known BCGs at $0.1<z<0.8$ ($\\bar{z}\\sim0.3$). At $z<0.1$ we find $L_K\\propto M_{200}^{0.24\\pm0.08}$ for the BCGs and the early-type BGGs in groups with extended X-ray emission and $L_K\\propto M_{200}^{0.11\\pm0.10}$ for the BCGs alone.  At $0.1<z<0.8$ we find $L_K\\propto M_{200}^{0.28\\pm0.11}$.  We conclude that there is no evidence for evolution in this relationship between $z<0.1$ and $z<0.8$: BCG growth appears to still be limited by the timescale for dynamical friction at these earlier times, not proceeding according to the predictions of current semi-analytic models. ",
        "introduction": "\\label{intro} Brightest cluster galaxies (BCGs) are the most massive galaxies known in the Universe and are observed to have remarkably small dispersion in their absolute magnitudes compared to other early-type cluster galaxies (e.g. \\citealt{aragon-salamanca98,cm98}).  Their position at the centre of clusters suggests their evolution and their unique properties are linked to their special environment (e.g. \\citealt{edge91,vonderlinden06}).  However, it is still not clear exactly which processes drive their evolution.  Of central importance in this context is the mass assembly history of these most massive galaxies and how that depends on their environment. Semi-analytic models based on merger trees from the Millennium N-body simulation predict that, while the stars in BCGs form at high-redshift, the mass of the galaxy only assembles recently (doubling in mass since $z\\sim1$; \\citealt{delucia06}), with the most massive galaxies in the most massive clusters having undergone the most mass assembly.  Observations of BCGs agree that their properties are consistent with having undergone mergers (e.g. \\citealt{boylan-kolchin06,bernardi07}) and BCGs in the most massive clusters do appear to have undergone more stellar mass evolution than those in less massive clusters, as shown by their surface brightness profiles \\citep{brough05}.  However, that evolution must have predominantly occurred before $z\\sim1$ as shown by the uniformity in their absolute magnitudes \\citep{brough02}, and their steep metallicity gradients \\citep{brough07}.  An important clue in this context lies in the central galaxy luminosity--host-halo mass (L$_{\\rm c}$M$_{\\rm h}$) relation, which shows a scaling between BCG mass and that of its host halo. \\cite{lin04} found that $L_K\\propto M_{200}^{0.26\\pm0.04}$ for 93 BCGs at $z<0.1$, \\cite{popesso07} found $L_r\\propto M_{200}^{0.33\\pm0.04}$ for 217 BCGs at $0.01<z<0.25$ in the Sloan Digital Sky Survey (SDSS), and \\cite{hansen07} found $\\langle L_i\\rangle \\propto M_{200}^{0.30\\pm0.01}$ for the SDSS MaxBCG cluster sample ($0.1<z<0.3$). Studies of the L$_{\\rm c}$M$_{\\rm h}$-relation have not ventured beyond $z\\sim0.3$, so it is not clear whether this relationship evolves with cosmic time.  The L$_{\\rm c}$M$_{\\rm h}$-relation is also a natural output from analyses of cosmological simulations: Halo Occupation Distribution functions (HOD; e.g. \\citealt{vandenbosch05,zheng07}) and Conditional Luminosity Functions (CLF; e.g. Cooray \\& Milosavljevi\\'{c} 2005; henceforth CM05).  These statistically assign numbers of galaxies or the luminosity distribution of galaxies in a given dark matter halo, distinguishing between central and satellite galaxies and enabling a direct comparison with observations.  A related approach is the application of a direct relationship between the mass of a dark matter halo or sub-halo and the observed galaxy luminosity (e.g. \\citealt{vale06,vale07}; henceforth VO06, VO07).  CM05 predict $\\langle L\\rangle \\propto M_{200}^{<0.3}$ above halo masses of $4\\times10^{13}h^{-1}M_{\\odot}$ as the timescale on which dynamical friction causes satellite galaxies to fall in to the centre exceeds the age of the host system halo. VO06 predict $\\langle L_{K} \\rangle \\propto M_{100}^{0.28}$ as a result of the hierarchical formation of structure in the Universe, with a prescription for sub-halo mass loss.  We interpret the L$_{\\rm c}$M$_{\\rm h}$-relation in this paradigm: that it results from the BCG growing in step with its host halo.  \\cite{yang05} analysed the HOD of galaxy groups in the 2-degree Field Galaxy Redshift Survey ($0.01<z<0.2$) and found $\\langle L_{cen,b_J}\\rangle \\propto M_h^{0.25}$.  More recently, \\cite{yang07} determined the CLF of groups ($0.01<z<0.2$) in the SDSS and found $\\langle L_{cen,r}\\rangle \\propto M_h^{0.17}$.  At higher redshifts, \\cite{zheng07} analysed the HOD using the two-point correlation function of the DEEP2 galaxy redshift survey (${\\bar z}\\sim1$) in comparison to the SDSS (${\\bar z}\\sim0$) and determined the L$_{\\rm c}$M$_{\\rm h}$-relation and its evolution.  They found the VO06 model to be a good fit to their data at $z\\sim0$ and $z\\sim1$ suggesting that this relationship has not evolved.  \\cite{white07} found $\\langle L_{cen,B}\\rangle \\propto M_h^{0.36}$ at $z\\sim0.5$ from their HOD analysis of $\\sim2L_{\\star}$ red galaxies in the NOAO Deep Wide Field Survey.  These studies, however, do not directly identify central galaxies.  In this Letter, we examine the central-galaxy luminosity -- host halo mass relation for a sample of known BCGs in groups and clusters at $z<0.1$ and, for the first time, the relationship for known BCGs in clusters at redshifts $0.1<z<0.8$.  This will enable us to directly determine the evolution in this relationship for known BCGs and to compare to the HOD analyses at these redshifts. It will also enable us to examine the \\cite{delucia06} prediction of these galaxies having doubled in mass since $z\\sim1$.  In Section~\\ref{Data} we introduce our sample and then present our comparison of BGG and BCG luminosities with host system mass in Section~\\ref{Results}.  We discuss the implications of our results in Section~\\ref{Discussion} and draw our conclusions in Section~\\ref{Conclusions}.  Throughout this paper we assume $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m=0.3$ and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "\\label{Conclusions} We have examined the BGG and BCG luminosity -- host halo mass relation at: $z<0.1$ and, at higher redshifts than previously measured: $0.1<z<0.8$.  We conclude that:  \\begin{itemize}  \\item BCGs follow a relationship with the mass of their host cluster. This relationship does not evolve from $z\\sim1$ to $z\\sim0$.  \\item Our data are consistent with the predictions of two illustrative HOD models for the relationship of central galaxy luminosity with host halo mass. %  \\item The BCG -- halo-mass relation appears to be controlled by the timescale for dynamical friction even at $0.1<z<0.8$.  \\end{itemize}"
    },
    "0801/0801.4563_arXiv.txt": {
        "abstract": "% Galaxies are seen from different viewing angles and their \\emph{observed} properties change as a function of viewing angle.  In many circumstances we would rather know the \\emph{intrinsic} properties of galaxies --- those properties that do not depend on viewing angle.  For a large sample of galaxies it is possible to recover the intrinsic properties of galaxies, statistically, by looking for correlations with galaxy inclination, and then applying a correction to remove those dependencies. Studying the intrinsic properties of galaxies can give a different impression of the galaxy population and help avoid the mistake of connecting observed properties to quantities that don't depend on inclination like halo mass. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.2116_arXiv.txt": {
        "abstract": "Inflation and moduli stabilisation mechanisms work well independently, and many string-motivated supergravity models have been proposed for them. However a complete theory will contain both, and there will be (gravitational) interactions between the two sectors. These give corrections to the inflaton potential,  which generically ruin inflation. This holds true even for fine-tuned moduli stabilisation schemes. Following a suggestion by~\\cite{anaXW}, we show that a viable combined model can be obtained if it is the K\\\"ahler functions ($G= K+\\ln |W|^2$) of the two sectors that are added, rather than the superpotentials (as is usually done). Interaction between the two sectors does still impose some restrictions on the moduli stabilisation mechanism, which are derived. Significantly, we find that the (post-inflation) moduli stabilisation scale no longer needs to be above the inflationary energy scale. \\vskip 0.1in  \\noindent {\\bf Keywords:} inflation, cosmology of theories beyond the SM ",
        "introduction": "Many attempts have been made to implement inflation in extensions of the standard model, although to date there is still no model that is truly convincing.  Supersymmetric (SUSY) theories appear to be more promising. They include numerous moduli fields, i.e.\\ scalar fields which in the supersymmetric limit have an exactly flat potential, as is required for slow-roll inflation.  Any one of these moduli fields could play the role of the inflaton field.  As a concrete example we will consider $F$-term hybrid inflation in this work. In the SUSY limit it  has a flat direction, but when extended to include gravity the situation is less rosy. The large energy density during inflation breaks SUSY spontaneously, and supergravity (SUGRA) effects lift the flatness of the moduli potential. This is the infamous $\\eta$-problem~\\cite{copeland,dine}.  Furthermore, the particular form of the SUGRA potential means that all other, non-inflationary sectors of the full theory will couple to the inflation sector. The coupling will be small, in the models we consider it is only of gravitational strength, but it can nevertheless have large effects.  This is a generic problem for all inflation models, and as we will see, this small coupling between the different sectors frequently kills an otherwise good model. As a specific example, we will study the effects of a modulus stabilisation sector on 4D ${\\mathcal N}=1$ SUGRA $F$-term hybrid inflation. All the other, non-inflationary moduli fields must be fixed during inflation, and so a full SUGRA theory must include additional physics to do this.  For this we will consider  KKLT-like~\\cite{KKLT} and KL-like~\\cite{KL2} moduli stabilisation schemes.  As we will see, the moduli sector gives rise to additional --- and quite generically fatal --- corrections to the inflaton potential. This raises the questions of whether the original SUSY hybrid inflation model can actually be embedded in a full, realistic theory, and if so, are its original predictions valid?  For the answer to both these questions to be yes, the coupling between the two sectors must somehow be minimal, so that neither the moduli corrections to the inflation potential, nor the inflaton corrections to the moduli stabilisation potential ruin the model. As we will show, it is possible, but non-trivial, to achieve this.  There are of course many other models of inflation, which offer alternative approaches to the issue of moduli-inflation coupling. For example, in modular inflation models the modulus field itself is the inflaton~\\cite{modular}. In a sense, the coupling is maximal --- nevertheless successful (fine-tuned)  models have been constructed~\\cite{rt}.  In brane inflation models  the inflaton potential arises from brane interactions, and depends explicitly on the volume modulus.  Stabilising the modulus field then inevitably gives a curvature correction to the inflation potential~\\cite{KKLMMT}. However explicit examples have been constructed where, for fine-tuned parameters, the corrections to $\\eta$ cancel to a high degree, allowing inflation~\\cite{baumann,krause}.  In contrast to the above models, our strategy is to decouple the inflation and modulus sectors as much as possible.  One advantage of this is that it also allows us to decouple the scale of inflation from the gravitino mass scale. At the cost of tuning, it is then possible to have the gravitino in the phenomenologically favoured TeV range without the need for low scale inflation.   The $\\eta$-problem is a common feature of SUGRA inflation models.  To illustrate it, consider a canonically normalised inflaton field with $K= |\\phi|^2$.  The inflationary potential is of the form $V \\sim \\e^K V_* \\sim V_*(1+ |\\phi|^2 + \\cdots)$, with $V_*$ the nearly constant energy density driving inflation.  It follows that the slow-roll parameter $\\eta  = V''/V$ is of order unity, and slow-roll inflation does not occur.  To avoid this conclusion one can fine-tune the model such that the coefficient of the $|\\phi|^2$-term in the potential cancels.  More elegantly perhaps, one can try to achieve the same using a symmetry.  An example of the latter approach is the (accidental) Heisenberg symmetry of the K\\\"ahler potential in $D$-term hybrid inflation~\\cite{gaillard}.  In this paper we avoid the above $\\eta$-problem by using a shift symmetry for the inflaton, $\\phi \\to \\phi + a$, which leaves the K\\\"ahler potential invariant~\\cite{modular,shift2}.  Since the inflaton field $\\Re (\\phi)$ no longer appears explicitly in the K\\\"ahler potential, the large mass corrections to the inflaton field are avoided.  However, the shift symmetry does not kill all the corrections to the inflaton potential.  In the presence of moduli fields $\\eta$- (and $\\epsilon$-) problems appear again.  As a concrete example, consider the case of a single modulus field $T$. If moduli fields are present, they need to be fixed during inflation.   The modulus potential typically has a local minimum at finite field value separated by a barrier from the global minimum at infinity.   The classic example is the KKLT potential~\\cite{KKLT}.  To assure that the modulus does not run away during inflation the barrier should be large. This is the case if the modulus mass is large $m_T > H_*$, with $H_*$ being the Hubble constant during inflation~\\cite{KL2}.  Now since the moduli stabilisation mechanism breaks SUSY, there are soft corrections to the inflaton potential, typically of $\\Or(m_{3/2} H_*)$.  The flatness of the inflaton potential is lost unless the gravitino mass is sufficiently small $m_{3/2} < H_*$.  The problem with this is that one cannot tune the gravitino mass arbitrarily:  in a generic, KKLT-like potential $m_{3/2} \\sim m_T$, and a small gravitino mass is at odds with keeping the modulus fixed. It is therefore difficult to embed inflation in such a scheme.  A solution to the above moduli problem put forward by Kallosh and Linde (henceforth denoted by KL)~\\cite{KL}  is to fine-tune the modulus potential so that $m_{3/2} \\ll m_T$. Then if the Hubble constant during inflation is between these two mass scales, the modulus remains fixed while the soft  corrections to the inflaton mass are small.  Such a set-up has the additional advantage that the gravitino mass can be in the TeV range without the need for low scale inflation.   KL gave an explicit realisation of this idea using a racetrack potential for the modulus.  All problems then appear to be solved, but this is deceiving.  Although the moduli corrections are small after inflation thanks to the fine-tuning in the KL set-up, this is not necessarily true during inflation.  During inflation the modulus field $T$ is slightly displaced from its post-inflationary minimum, disrupting the minute fine-tuning of the potential, with potentially serious consequences.  Indeed, as we will show, in $F$-term hybrid inflation the effects of the modulus displacement are substantial, resulting in $\\eta \\approx -3$ and ruining inflation. The need to include the dynamics of the modulus field during inflation was previously noted in \\cite{mfi,lalak}.  In this paper we will study $F$-term hybrid inflation, which serves to illustrate all the observations made above. It is a multi-field model of inflation, consisting of the inflaton field, and two oppositely charged waterfall fields which are responsible for ending inflation. When combined with a KKLT modulus sector, the corrections to both the inflaton and the waterfall field potentials are large.  Although the mass correction to the inflaton can be protected by a shift symmetry, this is not the case for the waterfall fields, and as a result there is generally no graceful exit from inflation.  Tuning the modulus sector, as in the KL set-up, can reduce these corrections to a harmless size.  However all of this is under the assumption that the modulus $T$ is fixed during inflation.  Taking the modulus dynamics into account we find that even in the fine-tuned KL-stabilisation scheme the corrections are not harmless after all. On the contrary, they prevent inflation from working.  In all previous studies of the effect of the moduli sector on inflation~\\cite{mfi,lalak,mdi,riotto}, the K\\\"ahler and superpotentials of the modulus and inflaton sectors were simply added to get the combined theory, i.e.\\ take $W_{\\rm total} = W_{\\rm inf} + W_{\\rm  mod}$ to get the full superpotential.  In this paper we instead multiply the superpotentials: $W_{\\rm total} = W_{\\rm inf} W_{\\rm  mod}$, as proposed by Ach\\'ucarro and Sousa~\\cite{anaXW}. As we will show, this greatly reduces the moduli corrections. Indeed $F$-term hybrid inflation combined with KL, or even KKLT,  in this way can give a viable inflation model.  Although multiplying superpotentials may sound odd at first, it is natural in a supergravity formulation in terms of the K\\\"ahler function $G = K +\\ln |W|^2$.  Any supersymmetric theory only depends on the K\\\"ahler- and  superpotential through the combination $G$, suggesting that it is the only significant quantity. Adding the K\\\"ahler functions of the two sectors is equivalent to adding their K\\\"ahler potentials and multiplying their superpotentials.  Adding K\\\"ahler functions has the nice property that a SUSY critical point of the modulus sector is automatically a SUSY critical point of the full theory as well~\\cite{anaXW,binXW} --- this feature is at the heart of the reduced moduli corrections.  In the limit of a small gravitino mass, all the corrections to the inflaton potential are small, including those due to the dynamics of the modulus field during inflation.  The resulting inflationary model thus gives similar inflationary predictions to the usual $F$-term hybrid inflation in the absence of a modulus sector.  Although there are still some constraints on the model parameters,  we want to stress that successful inflation is achieved without the need for fine-tuning --- this is in contrast to  most other combined inflaton-moduli models. A notable feature of the model is that it is possible for the vacuum modulus mass to be smaller than the Hubble scale during inflation, without the modulus running off to infinity.  This paper is organised as follows.  In the next section we provide the relevant background material.  We start with a short review  of standard $F$-term hybrid inflation, both in a SUSY and SUGRA theories. This is followed by a concise discussion of moduli stabilisation in KKLT- and KL-style schemes.  In section~\\ref{s:add} we discuss the resulting model when the two sectors are combined by adding superpotentials.  As we will see, even in the fine-tuned KL set-up this does not give a working model.  In section~\\ref{s:mult} we combine the modulus and inflaton sectors by their multiplying superpotentials, or equivalently by adding their K\\\"ahler-functions. The modulus corrections to the inflaton potential now are under control, and for a certain range of parameters we get successful inflation. The parameter range for which the standard $F$-term hybrid inflation predictions apply is determined in section~\\ref{s:inf}.  We end with some concluding remarks.  Throughout this article we will work in units with $M_{\\rm pl} = 1/\\sqrt{8\\pi G_N} =1$. ",
        "conclusions": "The flatness of the inflationary potential in SUGRA models is typically spoilt by corrections coming from supersymmetry breaking. Ironically enough, the vacuum energy which drives inflation breaks SUSY spontaneously, and so gives soft corrections to the inflaton; this is the well-known $\\eta$-problem. Introducing a shift symmetry for the inflaton will protect the inflation sector from itself, and remove the problem. However there will still be corrections coming from other sectors of the full theory, which can also disrupt inflation. In this paper we studied the effects of a moduli stabilisation sector on a $F$-term SUGRA hybrid inflation model.  We considered both a KKLT-like moduli stabilisation scheme, in which there is only one scale in the potential so $m_T \\sim m_{3/2}$, as well as a fine-tuned two-scale KL-like set-up with $m_T \\gg m_{3/2}$. In the KKLT set-up, requiring the modulus to be fixed during inflation raises the scale of the modulus potential, and as a result the soft corrections to both the inflaton slope and the waterfall field masses are too large for inflation to take place.  This problem is circumvented in the KL set-up where the gravitino mass, and consequently the corrections to the inflationary potential, can be tuned arbitrarily small.  One would be inclined to conclude that KL moduli stabilisation can be combined almost effortlessly with inflation.  But this is not true.  The above conclusions only hold in the limit that the modulus field remains fixed during inflation.  Although this seems like a good approximation, as the displacement of the modulus minimum during inflation is indeed small, the correction to the flat inflaton potential is nevertheless large.  In fact, it gives $\\eta \\approx -3$, and thus no slow-roll inflation.  This analysis shows that it is important to take the dynamics of all fields during inflation into account, otherwise crucial effects may be missed.  We have proposed a way to solve all of the above problems, and successfully combine $F$-term hybrid inflation with moduli stabilisation.  The idea is to combine the modulus and inflaton sectors not by adding their respective superpotentials, as is usually done, but by adding their respective K\\\"ahler functions $G = K + \\ln|W|^2$ instead.  Adding  K\\\"ahler functions corresponds to adding K\\\"ahler potentials and multiplying superpotentials.  This way of combining sectors greatly reduces their interactions.  In particular, for the case of combining inflation with a modulus sector, it greatly reduces the displacement of the modulus during inflation. Consequently the correction to the inflationary potential is harmlessly small.  For the fine-tuned two-scale KL set-up, or for a one-scale KKLT set-up with a fine-tuned mass scale, the corrections to the inflaton slope and waterfall masses are small as well.  Hence we indeed succeeded in constructing a successful model of inflation in the presence of moduli.  Even when multiplying superpotentials, there are still some constraints on the moduli sector parameters for viable inflation.  The graviton mass should be small enough to suppress the moduli corrections during inflation.  The modulus mass needs to be heavy and non-tachyonic during inflation to remain stabilised.  Finally the loop potential should be dominated by the contribution of the waterfall fields rather than by the modulus contribution. Nevertheless, there is still a large region of gravitino and modulus mass scales for which inflation works, and the inflationary predictions are nearly indistinguishable from the global SUSY model in the absence of moduli fields.  \\ack We are both grateful to K. Sousa and particularly A. Ach\\'ucarro for useful discussions, inspiration and for the suggestion that multiplication of superpotentials could be natural and helpful. We also thank N. Bevis, C. Burgess and J. Rocher for a useful comments.  SCD thanks the Netherlands Organisiation for Scientific Research (NWO) for financial support."
    },
    "0801/0801.0439_arXiv.txt": {
        "abstract": "We study galaxies that host both nuclear star clusters and active galactic nuclei (AGN) implying the presence of a massive black hole. We select a sample of 176 galaxies with previously detected nuclear star clusters that range from ellipticals to late-type spirals.  We search for AGN in this sample using optical spectroscopy and archival radio and X-ray data.  We find galaxies of all Hubble types and with a wide range of masses ($10^{9}-10^{11}$~M$_\\odot$) hosting both AGN and nuclear star clusters.  From the optical spectra, we classify 10\\% of the galaxies as AGN and an additional 15\\% as composite, indicating a mix of AGN and star-formation spectra. The fraction of nucleated galaxies with AGN increases strongly as a function of galaxy and nuclear star cluster mass.  For galaxies with both a NC and a black hole, we find that the masses of these two objects are quite similar.  However, non-detections of black holes in Local Group nuclear star clusters show that not all clusters host black holes of similar masses.  We discuss the implications of our results for the formation of nuclear star clusters and massive black holes. ",
        "introduction": "Nuclear star clusters (NCs) are massive star clusters coincident with the photocenters of galaxies. They are very common and have been found in $\\sim$75\\% of local late-type spirals \\citep{boker02} and Virgo dwarf elliptical galaxies \\citep{cote06}.  Their size is similar to that of globular clusters \\citep{boker04a}, but NCs are 1-2 orders of magnitude brighter and more massive \\citep{walcher05}.  Also, unlike most globular clusters, they have extended star formation histories \\citep{walcher06,rossa06} and complex morphologies \\citep{seth06}.  The luminosity of nuclear star clusters correlates with galaxy luminosity in both ellipticals \\citep{lotz01,graham03,cote06} and spirals \\citep{carollo98,boker04a}.  Recently, it has been shown that the masses of NCs follow scaling relationships with galaxy mass ($M_{gal}$), bulge velocity dispersion ($\\sigma$), and S\\'ersic index \\citep{ferrarese06,wehner06,rossa06,graham07}.  These scaling relations are very similar to those seen for massive black holes (MBHs\\footnote{We use massive black holes to refer to all non-stellar-mass black holes, including those normally referred to as intermediate mass and as super-massive black holes.}), appearing to extend those relations to lower masses.  Thus, \\citet{wehner06} and \\citet{ferrarese06} suggest that there may be a single scaling relation linking the mass of a central massive object (CMO; either a NC or MBH) to the large-scale properties of the galaxy. The existence of an $M_{CMO}-\\sigma$ or $M_{CMO}-M_{gal}$ relation suggests that the formation of the CMO is linked in some way to the evolution of the galaxy.  Theoretically, these scaling relationships can be understood in multiple ways.  They could be created by feedback from either NCs or MBHs regulating the star formation in the galaxy as a whole \\citep[e.g.,][]{mclaughlin06b}, or alternatively, they could simply result from gas accretion onto the nucleus in proportion to the galaxies' mass \\citep[e.g.,][]{li07}.  Despite this interesting connection between NCs and MBHs, and their link to galaxy formation and evolution, no systematic study of the overlap between these classes of objects exists. A handful of objects, including the Milky Way, are already known to host both NCs and MBHs (see \\S\\ref{litsec}).  \\citet{ferrarese06} and \\citet{wehner06} show there is a rough transition at galaxy masses of $\\sim$10$^{10}$~M$_\\odot$ (and corresponding CMO mass of $\\sim$10$^7$~M$_\\odot$), above which galaxies typically host MBHs, and below which galaxies have NCs.  While there is good evidence that more massive galaxies do not in fact host NCs \\citep{cote06}, it remains unclear how common MBHs are in lower mass galaxies \\citep[e.g.,][]{greene07}.  Recent theoretical work shows that MBHs could form from stellar mergers in a young, dense cluster environment \\citep{miller04,portegieszwart04}, and a direct link between NC and MBH formation may therefore exist.   We present a systematic study of the overlap between NCs and MBHs aimed at better understanding the relation between the two types of objects and the formation mechanism of CMOs in general. Starting with a sample of galaxies with known NCs, we search for active galactic nuclei (AGN) that are powered by accretion onto an MBH. This study gives a lower limit on the number of systems with MBHs, since quiescent and heavily obscured MBHs will not be detected as AGN.  We begin by describing our sample of galaxies with NCs, drawn from several different catalogs (\\S\\ref{samplesec}). Using optical spectra and radio and X-ray data, we examine our sample galaxies for evidence of AGN activity (\\S\\ref{agnsec}).  We then review galaxies for which detections of both AGN or MBHs and NCs exist in the literature (\\S\\ref{litsec}). We discuss the demographics of galaxies with AGN and NCs, the relative masses of these CMOs in galaxies where they co-exist, and the implications of this study for CMO formation, in \\S\\ref{dissec}.  We conclude and discuss future work in \\S\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We have assembled a sample of 176 galaxies known to host nuclear star clusters to study the relationship between nuclear star clusters and massive black holes.  We then use optical spectroscopy and radio and X-ray data to look for AGN in this sample.  For 75 galaxies with available optical spectra, we classify 7 galaxies as AGN and 11 as composite objects.  X-ray catalogs from {\\it Chandra, ROSAT}, and {\\it XMM} provide detections of 22 galaxies in the sample, 4 of which have well-localized sources with X-ray luminosities indicating they are likely AGN. Lastly, we have assembled previously published results for $9$ galaxies that indicate they have both NCs and MBHs/AGN; three are included in our sample.  From this work we conclude that galaxies with both NCs and MBHs are relatively common.  In addition, by examining the objects in our sample we find that: \\begin{enumerate} \\item Galaxies that host both NCs and AGN/MBHs span all Hubble types, have magnitudes as faint as $M_B \\sim -16$, and masses as small as 10$^9$~M$_\\odot$.  For galaxies in our sample with available spectra, $\\gtrsim 10$\\% appear to host both NCs and MBHs. \\item Galaxies with NCs have AGN detection fractions that increase strongly with increasing galaxy and NC mass, consistent with previous studies of the general galaxy population.  Variation of the AGN fraction with Hubble type in our sample is also consistent with the full Palomar survey \\citep{ho97b}.  This suggests that nuclear clusters do not play a strong role in promoting or limiting the activity of massive black holes. \\item In four galaxies with NCs and measured black hole masses the ratios of the MBH mass to the NC mass are between 0.1 and 1.  The luminosity of the AGN in our sample of galaxies with NCs are also consistent with these mass ratios.  However, the non-detection and low masses of MBHs in Local Group NCs suggest a much wider range of MBH-to-NC mass ratios. \\item For galaxies of the same mass, NCs in late-type spiral galaxies are typically an order of magnitude less massive than those in elliptical galaxies.  This result will be discussed further in a forthcoming paper. \\end{enumerate}  Although the data presented here provide compelling evidence for a significant overlap between the MBH and NC population, the evidence for the presence of an active MBH in many of the individual objects is quite weak. A number of observations could strengthen this evidence, including high-resolution {\\it Chandra} observations or {\\it Spitzer} mid-IR spectroscopy \\citep[e.g.,][]{satyapal07}.  Furthermore, less than half of the galaxies in our sample have available optical spectroscopy -- obtaining spectroscopy for more is desirable and will help fill gaps in parameter space (e.g., the lack of spectra of low-mass early type spiral).  Finally, detailed studies of objects with both NCs and MBHs will be required to improve our understanding of the connections between these central massive objects.  Acknowledgments: The authors thank Christy Tremonti for providing significant assistance with the SDSS spectra.  We also acknowledge valuable advice received from Margaret Geller, Julianne Dalcanton, David Helfand, Beth Willman, Andrew West, Nelson Caldwell, Andr\\'es Jord\\'an, Bob Becker, Ramesh Narayan, Martin Elvis, and Giuseppina Fabbiano.  Anil Seth gratefully acknowledges the support of the CfA Postdoctoral Fellowship.  Marcel Ag\\\"ueros is supported by an NSF Astronomy and Astrophysics Postdoctoral Fellowship under award AST-0602099.  This work would not have been possible without use of the Sloan Digital Sky Survey ({\\tt http://www.sdss.org/}), the NASA Extragalactic Database ({\\tt http://nedwww.ipac.caltech.edu/}), HyperLeda ({\\tt http://leda.univ-lyon1.fr/}), and NASA's High Energy Astrophysics Science Archive Research Center ({\\tt http://heasarc.gsfc.nasa.gov/})."
    },
    "0801/0801.0864_arXiv.txt": {
        "abstract": "{On October 24, 2007 the periodic comet \\object{17P/Holmes}  underwent an astonishing outburst that increased its apparent  total     brightness from magnitude $V\\sim17$ up to $V\\sim2.5$  in  roughly two days.  In this contribution we report on Wendelstein $0.8\\,$m telescope (WST) photometric observations  of  the   early   evolution    stages    of the  outburst.} {We   studied   the evolution of the                structure   morphology, its   kinematic,  and    provided  an estimate of the ejected dust mass.} {We  analized   $126$  images of the comet in the $BVRI$ photometric bands spread      between     26   October,   2007   and  November   20,   2007. The   bright    comet    core  appeared well separated from that one of a quickly  expanding dust cloud in all the data, and the bulk of the latter was contained in the field of view of our instrument during the days soon after   the   outburst, allowing precise estimates both of the separation velocities        of        the       two         luminous    baricenters, and        of           the            expansion    velocity  of the dust cloud. The ejected   dust  mass   was   derived    on     the   base   of differential photometry on background stars occulted by the moving cloud.} {The two   cores      were        moving         apart       from     each other        at        a            relative       projected     constant velocity of $(9.87\\,\\pm\\,0.07)\\,$arcsec/day ($0.135\\,\\pm\\,0.001$ km/sec). In               the           inner          regions       of        the dust              cloud              we         observed     a     linear increase in size at  a mean constant velocity   of $(14.6\\,\\pm\\,0.3)\\,$ arcsec/day                                ($0.200\\,\\pm\\,0.004\\,$ km/sec). Evidence        of          a   radial  velocity                gradient in      the         expanding       cloud       was      also       found. Our     estimate       for  the expanding coma's mass was of the order of $10^{-2}-1\\,$ comet's mass    implying a significant disintegration event. } { We   interpreted our observations in the context of an explosive scenario which   was   more   probably   originated  by  some internal instability processes,        rather    than    an   impact with an asteroidal  body. Due to the peculiar characteristics of this  event, further  observations and investigations are necessary in order   to  enlight the nature of the physical           processes          that       determined            it. } ",
        "introduction": "\\label{sect:introduction} In this    Letter,    we     concentrate on the analysis of the outburst of   comet     \\object{17P/Holmes},     occurred   on   October 24, 2007 (Santana 2007). We  provide an overview analysis of the early  evolution phases of the surrounding  cometary environment, in  the  period between October   $26$,   $2007$     and   November   $20$, $2007$   considering the global structure of the expanding material, its kinematic properties and inferring the ejected dust mass.  Comet    brightenings  are    well documented in the  literature, and in  general  they   can be associated with  either evident  cometary  nuclear fragmentation or not (see e.g. Sekanina et al. 2002   and  references therein).   Despite    their    recurrence, the  magnitude   and  the characteristics of   the   outbursts occurring to comet \\object{17P/Holmes} largely   justify       its    longstanding reputation in the annals of astronomy (Barnard 1896). The periodic comet \\object{17P/Holmes}  was  discovered  on November 6, 1892             by E. Holmes   in  London,  during an   outstanding   brightness   increase, followed       by         another similar event   on    January 16, 1893. When        discovered    by      Mr. Holmes        the           comet was         around         $5$      months      past         perihelion $(\\rm T\\,=\\,$June 13). On       October       24,      2007            around $6$ months after the last perihelion   passage                       $(\\rm T\\,=\\,$May 4, 2007)                             the                             comet underwent          a          similar              phenomenon        to those observed more than   one               hundred              years ago. The heliocentric distance of the           object        at        the      time    of the two major events was around            $\\rm \\Delta_{sun}\\,=\\,$2.39 AU         and $\\rm \\Delta_{sun}\\,=\\,$2.44 AU   in  1892    and      2007 respectively, and                 the                    orbital inclination remained substantially            unaltered        during                  this period                                      ($\\rm i\\,\\sim\\,20^{\\circ}$). This  Letter   is    structured    in    the    following    way:    in Sect.~\\ref{sect:observations},  we  present     the   observations   we acquired, and the data reduction method. In Sect.~\\ref{sect:morphology}, we discuss the    morphological     evolution     of     the  expanding structure. In  Sect.~\\ref{sect:cinematic}, we calculated the separation velocities  between the comet    and    the   center of the dust cloud, as well as the expansion velocity   of     the                    cloud. In Sect.~\\ref{sect:mass},   we  estimated   the   ejected   dust   mass. Finally in Sect.~\\ref{sect:conclusions},  we  sum   up     our  results. ",
        "conclusions": "\\label{sect:conclusions}  In    this  Letter we analized the early phases of the outburst that comet  \\object{17P/Holmes} underwent   on    October 24, 2007.    We acquired    $BVRI$ photometric      images      at       the      Wendelstein $0.8\\,$m telescope, between 26/10/2007 and 20/11/2007.  We   observed         a        spherically simmetric              dust               cloud                         moving away    from      the comet nucleus with a mean projected constant velocity of $(9.87\\,\\pm\\,0.07)$    arcsec/day   ($0.135\\,\\pm\\,0.001$ Km/sec),  while   the dust       cloud      was     expanding     with    a   mean constant velocity of $(14.6\\,\\pm\\,0.3)$       arcsec/day         ($0.200\\,\\pm\\,0.004\\,$ Km/sec). These         results        are        in       agreement     with the ones obtained       during        past        outbursts      of      this    comet. Evidence   of  a gradient in the expansion velocity of the dust cloud was also found   with the velocity increasing towards the external regions.    Finally, performing     differential     photometry   on      background stars occulted by      the   moving  cloud,  and assuming a uniform and homogeneous spherical mass distribution,     we      derived      for     the coma's mass a value of $10^{12}-10^{14}$kg,          around                   $10^{-2}-1$ comet's mass. We    interpreted  our observations in the context of an explosive event, probably caused by some internal instability processes,  rather than  by    an asteroidal         body's             impact. Various               mechanisms   have     been    proposed    in  the   past to explain    comets'    splitting,    involving    tidal,  thermal and    rotational  forces (Sekanina 1997).   These        processes      could have  played an important    role      for the         event        discussed here, although  some  specific  characteristics allow to consider it as more peculiar. For example, generally  the separation velocities of the  splitting  components are of the order of a few $\\rm m/s$, while   in this case we found a projected relative velocity around $2$ orders of magnitude  larger. The outburst  itself represents the largest apparent brightness increase ever observed for a  comet. In    conclusion,    we     underline  the importance to consider other            observational  results   in   order to accurately characterize this             event               and        to     provide   a more insight view on         its           enigmatic and still not well  understood  nature.    {\\bf"
    },
    "0801/0801.4440_arXiv.txt": {
        "abstract": "Cosmic shear offers a remarkbly clean way to measure the equation of state of the Universe and its evolution. Resolution over a wide field is paramount, and Antarctica offers unique possibilities in this respect. There is an order of magnitude gain in speed over temperate sites, or a factor three in surface density. This means that PILOT outperforms much larger telescopes elsewhere, and can compete with the proposed DUNE space mission. ",
        "introduction": "Weak lensing has been widely identified as the most promising route to measuring the evolution of the equation of state of the Universe, and hence to understanding the nature of dark energy. A large number of huge lensing surveys are planned: Panstarrs, LSST, The Dark Enery Survey on the CTIO 4-m, KIDS on VST, HyperSuprimeCam for Subaru; and the DUNE and SNAP satellites.  For weak lensing, resolution is paramount, because the galaxies used to measure the effect are $<1''$ in size, and must be at least partially resolved. Antarctica offers unique possibilities for wide-field, high-resolution imaging, so even modest sized telescopes may offer world-beating performance. This paper is a preliminary investigation into the possibilities for measuring weak lensing with PILOT - the Pathfinder for an International Large Optical Telescope, a 2-m class optical/infrared telescope proposed for Dome C on the Antarctic Plateau with first light at the end of 2012.  We assume a 2.4m f/10 telescope, with a $0.75deg^2$ CCD camera, with image scale $\\sim 0.1''$/pixel. ",
        "conclusions": ""
    },
    "0801/0801.4781.txt": {
        "abstract": "We present new HST optical imagery as well as new UV and IR spectroscopic data obtained with the Hubble and Spitzer Space Telescopes, respectively, of the halo planetary nebula DdDm-1. For the first time we present a resolved image of this object which indicates that the morphology of DdDm-1 can be described as two orthogonal elliptical components in the central part surrounded by an extended halo. The extent of the emission is somewhat larger than was previously reported in the literature. We combine the spectral data with our own previously published optical measurements to derive nebular abundances of He, C, N, O, Ne, Si, S, Cl, Ar, and Fe. Our abundance determinations include the use of the newly developed program ELSA for obtaining abundances directly from emission line strengths along with detailed photoionization models to render a robust set of abundances for this object. The metallicity, as gauged by oxygen, is found to be 0.46~dex below the solar value, confirming DdDm-1's status as a halo PN. In addition, we find that Si and Fe are markedly underabundant, suggesting their depletion onto dust. The very low (but uncertain) C/O ratio suggests that the chemistry of the nebula should be consistent with an oxygen-rich environment. We find that the sulfur abundance of DdDm-1 is only slightly below the value expected based upon the normal lockstep behavior between S and O observed in H~II regions and blue compact galaxies. The central star effective temperature and luminosity are estimated to be 55,000~K and 1000~L$_{\\odot}$, respectively, implying an initial progenitor mass of $<$1~M$_{\\odot}$. Finally, we report on a new radial velocity determination from echelle observations. ",
        "introduction": "Planetary nebulae (PNe) have long served as useful tracers of galactic interstellar abundances of elements such as O, Ne, S, Cl, and Ar, alpha elements whose abundances are expected to evolve in lockstep due to their common synthesis sites in massive stars. In a recent study of S, Cl, and Ar abundances in over 80 type~II PNe located mainly in the northern Galactic hemisphere, \\citet{HKB04} discovered that while Ne and O abundances track each other closely, S abundances in a large fraction of objects fall significantly below the values expected for their O abundances and exhibit a large amount of scatter as well. These authors also showed that the same pattern is present in data in the large southern survey by \\citet{KB94}.  \\citet{HKB04} suggested that this tendency of PNe to display a S deficit, a situation they dubbed the ``sulfur anomaly'', was probably due to the underestimation for many objects of the ionization correction factor used to account for unobserved ions of S when determining the total gas-phase elemental abundance from the optically observable ions of S$^+$ and S$^{+2}$. However, measurements of S$^{+3}$ abundances by \\citet{D03} and the results from numerous ISO projects summarized in \\citet{PBS06} using the [S~IV] 10.5$\\mu$m line appear to rule out this explanation, although it is probably too early to draw definite conclusions. Instead, \\citet{PBS06} suggested that the sulfur anomaly could result from the depletion of S onto dust due to the formation of sulfides such as MgS and FeS. Sulfide formation is expected to occur most readily in carbon rich environments, i.e., in objects where C/O $>$ 1.  The goal of the project described in this paper is to determine and study the gas phase abundances of the halo PN DdDm-1, an object for which we present new optical imagery as well as spectroscopic measurements of important UV and IR emission lines. DdDm-1 was chosen primarily because of the large amount of data both new (presented here) and previously published for this object.  Of the many catalogued PNe in our Galaxy, only about 12 are believed to be located in the halo.  Abundance studies of these objects can be used to determine the chemical composition of the halo and can provide additional insight into the evolution of this region of the Milky Way. These objects have consistently been shown to possess sub-solar levels of metals [\\citet{TPRP81}, \\citet{PTPR92}, \\citet{HHM97}, \\citet{D03}]. DdDm-1 has been included in many large abundance surveys but has only been studied in detail in a few instances, e.g. \\citet{BC84}, \\citet{D03}, and \\citet{WCP05}. Here we focus on it exclusively.  In this paper we present new optical HST imagery, new UV data taken with the FOS on HST, and new IR data taken with the IRS on the Spitzer Space Telescope. The imagery allows us for the first time to resolve the nebula and study its morphology. Then combining the spectral data with our own previously published optical data, we compute the abundances of He, C, N, O, Ne, Si, S, Cl, Ar, and Fe using both empirical and numerical methods to ensure robust results. For the empirical method we employ the program ELSA \\citep{J07}; the program Cloudy \\citep{Fer98} is used for the numerical method. The numerical approach also allows us to infer the effective temperature and luminosity of the central star. Since the Spitzer data allow us to determine the abundance of S$^{+3}$, the primary unobservable ion when only optical measurements are available, we evaluate the accuracy of the sulfur ionization correction factor for DdDm-1. With a new sulfur abundance in hand, we then assess the S deficit of DdDm-1, as well as check on the consistency of the deficit's magnitude with the value of C/O derived from our new UV measurements. Finally, we present new echelle data for DdDm-1 and determine its radial velocity.  The paper is organized as follows. In Section 2, we discuss our observations; \\textsection 3 contains a description of our procedure for computing abundances; in \\textsection 4, we discuss our results and we give a summary and conclusions in \\textsection 5. ",
        "conclusions": "We have reported on new IR and UV spectra of the halo planetary nebula DdDm-1 obtained with the Spitzer Space Telescope with the IRS and Hubble Space Telescope with the FOS, respectively. By combining these new data with existing optical measurements, the nebular abundances of He, C, N, O, Ne, Si, S, Cl, Ar, and Fe were determined. The abundance analysis included the computation of detailed photoionization models which made use of an input stellar atmosphere whose characteristics were consistent with a central star of low metallicity. We also present new echelle data for DdDm-1. We have found the following:  \\begin{enumerate}  \\item We present what we believe to be the first resolved image of DdDm-1. Reconstructed imagery taken with the HST in 1993 indicated that the morphology of DdDm-1 can be described as two orthogonal elliptical components in the central part surrounded by an extended halo. The extent of the emission is somewhat larger than was previously reported in the literature.  \\item We present new UV, IR, and echelle data, where the first two sets were acquired using HST and Spitzer, respectively.  \\item Our new determinations of the Si and Fe abundances indicate that their levels are far below those expected from the metallicity of DdDm-1, possibly indicating that their gas phase abundances have been depleted by dust formation.  \\item We determine that C/O $<$ 1, although our C/H abundance is uncertain. However, a C/O ratio below unity suggests that the chemistry of the nebula is O-rich in character. Thus, sulfides should be absent any dust that has formed in the environment of DdDm-1.  \\item We find that the abundance of S$^{+3}$ which we determined directly from our new [S~IV] 10.5$\\mu$m measurement agrees closely with another modern one by \\citet{D03} and is consistent with the level predicted by the value of the ICF obtained when only optical lines of [S~II] and [S~III] are used. The small abundance of S$^{+3}$ that we derive indicates that DdDm-1 is a relatively low excitation nebula when compared with other PNe.  \\item Our total gas-phase S abundance for DdDm-1 is consistent with the value expected from its O abundance, under the assumption of lockstep behavior between these two elements. Thus DdDm-1 has a negligible S deficit. On the other hand, if the large S deficits observed in other PNe are indeed due to sulfide formation in a C-rich environment as others have suggested, then the small S deficit of DdDm-1 is entirely consistent with its C-poor properties.  \\item For the central star we find that T$_{eff}$=55,000~K and L=1000~L$_{\\odot}$. Comparing the star with theoretical model tracks suggests that it is a He-burning object with a mass of less than 1~M$_{\\odot}$.  \\item The heliocentric radial velocity of DdDm-1 is -310.3 $\\pm$ 1.4 km s$^{-1}$.  \\end{enumerate}"
    },
    "0801/0801.1115_arXiv.txt": {
        "abstract": "Type 2 quasars are luminous active galactic nuclei (AGN) whose central regions are obscured by large amounts of gas and dust. In this paper, we present a catalog of type 2 quasars from the Sloan Digital Sky Survey (SDSS), selected based on their optical emission lines. The catalog contains 887 objects with redshifts $z<0.83$; this is six times larger than the previous version and is by far the largest sample of type 2 quasars in the literature. We derive the \\oiiil {} luminosity function for $10^{8.3} L_{\\odot} < L_{\\rm [OIII]} < 10^{10} L_{\\odot}$ (corresponding to intrinsic luminosities up to $M[2500{\\rm \\AA}]\\simeq -28$ mag or bolometric luminosities up to $4\\times 10^{47}$ erg s$^{-1}$). This luminosity function provides robust lower limits to the actual space density of obscured quasars, due to our selection criteria, the details of the spectroscopic target selection, as well as other effects. We derive the equivalent luminosity function for the complete sample of type 1 (unobscured) quasars and determine the ratio of type 2 to type 1 quasar number densities. Our data constrain this ratio to be at least $\\sim 1.5:1$ for $10^{8.3} L_\\odot <  L_{\\rm [OIII]} < 10^{9.5} L_\\odot$ at $z<0.3$, and at least $\\sim 1.2:1$ for $L_{\\rm [OIII]} \\sim 10^{10} L_\\odot$ at $0.3<z<0.83$. Type 2 quasars are at least as abundant as type 1 quasars in the relatively nearby Universe ($z\\la 0.8$) for the highest luminosities. ",
        "introduction": "\\label{sec:intr}  Active galactic nuclei (AGN) can be classified as type 1 (unobscured) or type 2 (obscured) based on the presence or absence of broad hydrogen and helium emission lines in their optical spectra. Unification models of AGN attribute the distinction between these two types to a difference in the observer's viewing angle to a nucleus surrounded by non-isotropic obscuring material \\citep{anto93, urry95}. These models are well-established for low-luminosity, nearby AGN. However, their applicability to high-luminosity AGN (i.e., quasars, classically defined to be sources with bolometric luminosities greater than $10^{45}$ erg s$^{-1}$) has long been controversial. Moreover, while it has long been known that obscured AGN dominate the low-luminosity population in the local Universe \\citep{oste88, salz89, huch92, hao05b, simp05}, the situation is less clear for high-luminosity AGN. This is in part because these objects are much rarer and more difficult to sample, since the AGN luminosity function decreases steeply with luminosity.  In this paper, we present a catalog of 887 type 2 quasars from the Sloan Digital Sky Survey (SDSS; \\citealt{york00}), the largest sample of type 2 quasars in the literature to date. They are selected based on their optical emission lines, have redshifts $z< 0.83$ and \\oiiil {} emission line luminosities extending to $\\mbox{\\loiii} \\ga 10^{10} L_\\odot$ (corresponding to intrinsic UV luminosity $M_{2500}\\la -28$ mag). Multi-wavelength observations of the most luminous objects from a previous version of the sample (\\citealt{zaka03}, hereafter Paper I) have confirmed that they have infrared luminosities up to and above $10^{46}$ erg s$^{-1}$, have spectral energy distributions expected of type 2 quasars \\citep{zaka04, vign04, ptak06, vign06, zaka08}, and contain type 1 quasars in their centers revealed by polarimetric measurements \\citep{zaka05, zaka06}.  With this updated catalog drawn from roughly three times as much SDSS data, we can now sample the high-luminosity AGN population in sufficiently large numbers to draw quantitative conclusions. First, we derive the \\oiiil {} luminosity function (LF) of type 2 quasars. Then, by directly comparing the space densities of type 2 and type 1 sources, we place robust lower limits on the fraction of obscured quasars in the local Universe as a function of \\oiiil {} luminosity. Studying the space densities of different types of AGN provides strong constraints on the simplest AGN unification scenario as well as its modifications. Moreover, quantifying the obscured quasar population is essential for many applications, such as relating the present mass density of local black holes to the accretion history of the entire AGN population (e.g., \\citealt{solt82,yu02,marc04}), understanding the origin of the cosmic X-ray background (e.g., \\citealt{coma95,gill07}), and studying the effects of luminosity on AGN structure (e.g., \\citealt{lawr91,urry95,hopk06}).  Our results are complementary to those derived from obscured quasars selected from hard X-ray surveys \\citep{ueda03, szok04, barg05, mark05, trei06, beck06, sazo07} and mid-infrared color selection \\citep{lacy05, ster05, mart06, poll07}.  For example, up to 20\\% of hard X-ray selected AGN do not show any emission lines in their optical spectra \\citep{rigb06}, and therefore would not be included in our sample. On the other hand, our sample would include Compton-thick objects that are missed in X-ray surveys, as long as their optical spectra meet our selection criteria.  The paper is structured as follows. In \\S \\ref{sec:samp}, we discuss our selection of type 2 quasars based on their optical emission lines and in \\S\\ref{sec:type2_lf}, we determine the \\oiiil {} luminosity function from this sample. In \\S\\ref{sec:type1_lf}, we determine the equivalent luminosity function from a complete sample of type 1 quasars and in \\S\\ref{sec:type12}, we determine the ratio of type 2 to type 1 quasars using the derived luminosity functions. We discuss caveats and implications of these results in \\S\\ref{sec:discussion} and we conclude in \\S\\ref{sec:conclusions}. We adopt a `concordance' cosmology, $h = 0.7$, $\\Omega_{\\rm M}=0.3$, and $\\Omega_\\Lambda = 0.7$. We identify emission lines using air wavelengths, identify objects with J2000 coordinates, and use asinh magnitudes \\citep{lupt99} corrected for Galactic extinction \\citep{schl98}. We often express luminosities in units of solar luminosities $L_\\odot = 3.826 \\times 10^{33}$ erg s$^{-1}$. ",
        "conclusions": "\\label{sec:conclusions}  In this paper, we present the largest sample of type 2 quasars to date. Table~\\ref{tab:qso2} lists the 887 optically-selected sources with redshifts $z\\lesssim 0.83$. These objects are selected from the spectroscopic database of the SDSS on the basis of their emission line properties. Candidate sources are required to have no broad (FWHM$>1100$ km/sec) components in their permitted emission lines. To distinguish type 2 AGN from star-forming galaxies, we use the standard line diagnostic diagrams involving [OIII]/H$\\beta$, [NII]/H$\\alpha$ and [SII]/H$\\alpha$ line ratios (at redshifts $z<0.3$) or require other signs of a hidden AGN, such as the presence of the [NeV]3346,3426 emission lines (at redshifts $0.3<z<0.8$). We place a lower limit on the [OIII]5007 line luminosity of $10^{8.3}L_{\\odot}$, ensuring that the bolometric luminosities of these objects are above the classical Seyfert/quasar separation of $10^{45}$ erg s$^{-1}$.  For this sample, we calculate the \\oiiil\\ luminosity function using the 1/\\vmax\\ method (Fig.~\\ref{fig:phitable_hao}). The selection function that we calculate for each object involves taking into account four different spectroscopic target selection algorithms, of which two (Low-z QSO and High-z QSO) are color-based, one (Galaxy) is based on optical morphology and one (Serendipity FIRST) is based on the radio properties. By comparing how well each algorithm performs at selecting type 2 quasars, we find that color-based algorithms are rather ineffective and select no more than 20\\% of objects. Indeed, there is no well-defined region of the color-color space where type 2 quasars concentrate, even when only narrow redshift ranges are considered. The success of the SDSS in selecting a large number of type 2 quasars is due to the unprecedented size of the spectroscopic survey and the multitude of different target selection algorithms which allow for serendipitous objects.  We extend the \\oiiil\\ AGN luminosity function to luminosities about 2 orders of magnitude higher than was previously done, up to $L_{\\rmoiii}\\simeq 10^{10}L_{\\odot}$. This value corresponds to an intrinsic UV luminosity of $M_{2500}=-28$ mag and a bolometric luminosity of $4\\times 10^{47}$ erg s$^{-1}$.  We also derive the \\oiiil\\ luminosity function for a complete sample of 8,003 $z<0.83$ type 1 quasars taken from the SDSS quasar catalog and find it to be in excellent agreement with other measurements. We can then directly compare the luminosity functions of type 1 and type 2 quasars (Fig.~\\ref{fig:phitable_qso12_union_zbin}) and constrain the type 2 quasar fraction as a function of luminosity (Fig.~\\ref{fig:type2_frac}). We argue that the type 2 quasar luminosity function and the type 2 quasar fraction that we derive are robust lower limits. The main reasons are that (i) there may be a significant number of type 2 quasars that do not meet our spectroscopic selection criteria, and (ii) narrow lines are more extincted in type 2 quasars than they are in type 1 quasars, biasing the type 2/type 1 ratio at a given luminosity to lower values.  Objects with different \\oiiil {} luminosities and different redshifts suffer from different selection biases. Our best data are at low redshifts and relatively low luminosities ($z<0.3$ and $L_{\\rmoiii}<10^9L_{\\odot}$) and at high redshifts and relatively high luminosities (\\zc\\ and $L_{\\rmoiii}>10^{9.5}L_{\\odot}$). In these regimes we find that type 2 quasars are more abundant than type 1 quasars, with the type 2/type 1 ratios of 1.5:1 and 1.2:1, correspondingly."
    },
    "0801/0801.3440_arXiv.txt": {
        "abstract": " ",
        "introduction": "Recently, there has been a surge of interest in models where the Standard Model (SM) or the Minimal Supersymmetric Standard Model (MSSM) communicates with a partially hidden sector via either $Z'$ or Higgs interactions  \\cite{HV,W,wells,PW,PW2,HV2}.   These Hidden Valley or Higgs Portal models provide a stimulating and consistent alternative to the usual model building assumption of a desert above the weak scale.   Higgs-sector and $Z'$ interactions between the hidden sector and the SM states are special in that they involve gauge-invariant operators of dimension $d_{O} \\leq 4$, and thus can be induced by physics at arbitrarily high scales with unsuppressed couplings.  In the case of a $Z'$ the interactions can either occur directly with SM states if they are charged under the $U(1)'$ or, possibly more interestingly, indirectly due to a kinetic-mixing term, $\\ep F_{Y}^{\\mu\\nu} F'_{\\mu\\nu}$, between hypercharge and the new $U(1)$, in which case $\\ep$, and thus the effective size of the SM-hidden sector interaction, can be suppressed \\cite{kineticmixing, HV, HV2}. On the other hand, in the case of the Higgs-sector interactions of interest to us here, couplings of the form $|H|^2 s^2$ involving the SM or MSSM Higgs states and new SM gauge singlet states can be large, especially in the situation where the TeV-scale theory UV-completes not far above the weak scale to a strongly interacting theory with light composite states.  It is interesting to ask whether such models lead to new dark matter candidates with qualitatively different phenomenology.  In this paper we argue that dark matter communicating with a supersymmeterized SM purely via Higgs-sector interactions (the Higgs Portal) leads to new and unusual features.\\footnote{Other works which consider aspects of dark matter phenomenology in the context of Hidden Valley or Higgs Portal models are contained in Ref.~\\cite{gmsbdm,hiddendm}, while earlier related studies are contained in Ref.~\\cite{gaugesingletdm}. } First, as we will show, the thermal relic abundance in such a scenario can be consistent with the measured density of dark matter for masses as high as $\\sim 30\\tev$, much larger than are usually considered (while also being consistent with the upper bound on the mass of thermal relic dark matter derived from unitarity~\\cite{GK}).   Second, for dark matter masses above $\\sim 1\\tev$ non-perturbative Sommerfeld corrections \\cite{sommer} to the low-velocity annihilation rate are large.  Several authors have recently recognised the potential importance of these corrections to the dark matter relic density calculations \\cite{gunion, hisano, profumo, strumia}, which lead to enhanced annihilation rates in the case of attractive interactions.  Even more importantly, as we will argue in detail in a companion paper \\cite{companion}, these corrections have the potential to greatly enhance the indirect annihilation signals by factors of up to $10^5$ beyond those predicted without consideration of the Sommerfeld factor, potentially leading to a significant change in the optimal search strategy.  As well as providing examples in which the dark matter particle is beyond the kinematic reach of the Large Hadron Collider (LHC) but is potentially detectable by indirect and direct dark matter searches, the models presented here are independently motivated by the desire to raise the MSSM upper bound on the lightest Higgs mass, and so relax the current tension with the LEP2 Higgs-mass exclusion limit.  It is also interesting that our models may be given a UV completion in so-called ``Fat-Higgs\" models \\cite{Harnik:2003rs}\\footnote{Other models in a similar class to the Fat Higgs model are discussed in Ref.~\\cite{uvcomplete}.} in which some TeV-scale states are composites of the underlying strong-coupling dynamics.  This UV completion is consistent with both collider constraints and aesthetic requirements such as gauge coupling unification.  This completion is discussed in detail in Section~\\ref{fathiggs}.  Furthermore, the existence of partially hidden (secluded) sectors is common in models that attempt to embed the SM within a larger structure.   Well studied examples include higher-rank GUT models, such as those based upon $E_6$ \\cite{gut}, and supersymmetry breaking models, in particular the messenger sectors of gauge-mediated supersymmetry breaking models~\\cite{gmsbdm}.  More recently, it has been argued that secluded or hidden sectors in the form of Randall-Sundrum-like warped ``throats'' \\cite{throats} are a ubiquitous feature of the landscape of string compactifications \\cite{ubiquitous}, thus implying that there is not an insignificant probability that a hidden or secluded throat with a mass scale close to the weak scale exists.  In fact, as argued by Patt and Wilczek \\cite{PW}, the scales in sectors interacting by Higgs portal interactions are commonly tied together.  Naturally, if our dark matter candidate is to be the dominant component of the cosmological dark matter, we must ensure that the usual neutralino dark matter candidate of the MSSM leads either to a subdominant relic density or is unstable.  In the case in which $R$-parity is conserved and a neutralino is the lightest supersymmetric particle (LSP), the thermally generated abundance of such a state is in many models well below the measured dark matter density. In particular, wino-like or higgsino-like LSPs annihilate very efficiently, leading to subdominant abundances~\\cite{dmreview}. Coannihilations with other supersymmetric states can also deplete the neutralino abundance in many models~\\cite{GS}. Alternatively, instead of being a neutralino, the LSP could be a different supersymmetric state, such as a gravitino. Within the context of gauge-mediated supersymmetry breaking, for example, the LSP is typically a light gravitino which constitutes only a very small fraction of the cosmological dark matter abundance. On the other hand, if there exist $R$-parity violating interactions, then the LSP will be unstable thus evading this issue entirely.\\footnote{A late-decaying LSP may even be beneficial in that it can correct the BBN prediction for the $^6$Li to $^7$Li ratio \\cite{lithium}.}  Turning to the structure of our paper, in Section~\\ref{models} we introduce our models and explain how they are a modified form of the so-called Minimal Non-minimal Supersymmetric Standard Model (MNSSM), while in Section~\\ref{som} we give a brief introduction to the physics of the Sommerfeld enhancement that plays an important role in our calculations. In Section~\\ref{reliccalc} we summarize the calculation of the relevant dark matter annihilation cross section including the Sommerfeld enhancement and present our results for the relic density.  In Section~\\ref{detection} we briefly discuss the direct and indirect detection of our dark matter candidate, leaving a more detailed study for a companion paper~\\cite{companion}.   Section~\\ref{fathiggs}, in which we demonstrate that our models may be given a UV completion in so-called ``Fat-Higgs\" models where the states are composites of underlying strongly coupled dynamics, is somewhat outside the main development of our paper and may be skipped by readers only interested in dark matter phenomenology. Finally, our conclusions are given in Section~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}   In this article, we have discussed models in which a very heavy (3-30 TeV) dark matter candidate is present. In particular, we have focused on models motivated by Higgs Portal and Hidden Valley models, in which the dark matter (and the rest of the partially hidden sector) interacts with the Standard Model and its superpartners only through Higgs interactions.  Dark matter annihilations in this scenario are considerably enhanced by non-perturbative contributions known as the Sommerfeld effect. Through this enhancement, dark matter particles with masses well above the electroweak scale can be produced thermally in the early universe with an abundance consistent with the measured density of dark matter. The dark matter particle in this scenario, although well beyond the reach of the Large Hadron Collider, is still potentially detectable by direct and indirect dark matter experiments. Although we leave the details of this to future work~\\cite{companion}, we point out that Sommerfeld corrections can dramatically enhance the dark matter annihilation rate in low velocity dispersion environments, such as dwarf spheriodal galaxies, thus considerably improving the prospects for indirect dark matter searches.  The particular model we study, which adds two extra SM singlet states to the MSSM spectrum, is independently motivated by the desire to raise the upper bound on the lightest higgs mass, thus lessening the LEP fine-tuning constraints.  We also showed that our model may be given a UV completion in the form of the previously considered  ``Fat Higgs'' models, where the singlet states are composites arising from strong hidden-sector dynamics.  Thus, Higgs Portal Dark Matter provides an example of an attractive and motivated alternative to conventional MSSM neutralino dark matter which is less fine-tuned and may be tested by current and future indirect detection experiments.  \\bigskip \\bigskip  {\\bf Acknowledgements} JMR and SMW are partially supported by the EC Network 6th Framework Programme Research and Training Network ``Quest for Unification\" (MRTN-CT-2004-503369) and by the EU FP6 Marie Curie Research and Training Network ``UniverseNet\" (MPRN-CT-2006-035863). DC is supported by the Science and Technology Facilities Council. DH is supported by the United States Department of Energy and NASA grant NAG5-10842. Fermilab is operated by the Fermi Research Alliance, LLC under Contract No. DE-AC02-07CH11359 with the United States Department of Energy. We would like to thank Markus Ahlers, Joe Silk, and Tim Tait for useful discussions."
    },
    "0801/0801.3676_arXiv.txt": {
        "abstract": "We investigate cosmologies where the accelerated expansion of the Universe is driven by a field with an anisotropic equation of state. We model such scenarios within the Bianchi I framework, introducing two skewness parameters to quantify the deviation of pressure from isotropy. We study the dynamics of the background expansion in these models. A special case of anisotropic cosmological constant is analyzed in detail. The anisotropic expansion is then confronted with  the redshift and angular distribution of the supernovae type Ia. In addition, we investigate the effects on the cosmic microwave background (CMB) anisotropies for which the main signature appears to be a quadrupole contribution. We find that the two skewness parameters can be very well constrained. Tightest bounds are imposed by the CMB quadrupole, but there are anisotropic models which avoid this bound completely. Within these bounds, the anisotropy can be beneficial as a potential explanation of various anomalous cosmological observations, especially in the CMB at the largest angles. We also consider the dynamics of linear perturbations in these models. The covariant approach is used to derive the general evolution equations for cosmological perturbations taking into account imperfect sources in an anisotropic background. The implications for the galaxy formation are then studied. These results might help to make contact between the observed anomalies in CMB and large scale structure and fundamental theories exhibiting Lorentz violation. ",
        "introduction": "There is a remarkable amount of observational evidence that the large-scale structure of the Universe resembles nearly the simplest and most symmetric imaginable system. Hence, it is well approximated by the flat geometry version of the Friedmann-Lema\\^{i}tre models which are both homogeneous and isotropic. The fluctuations in the cosmic microwave background (CMB) are a crucial probe of these properties, and therefore the recent detections of some unexpected features in the CMB temperature anisotropies have raised a lot of speculations about the need to reconsider some of the basic cosmological assumptions. A hemispherical asymmetry has been reported \\cite{Eriksen:2003db}. The angular correlation spectrum seems to be lacking power at the largest scales \\cite{Hinshaw:2006ia}. The alignment of the quadrupole and octupole (the so called Axis of Evil \\cite{Land:2005ad}) could also be seem as an extra-ordinary and unlikely result of statistically isotropic perturbations, even without taking into account that these multipoles happen also to be aligned to some extent with the dipole and with the equinox. The Axis of Evil does not show any correlation with the lack of angular power \\cite{Rakic:2007ve}. Given the {a \\it posteriori} nature of these considerations, the cosmic variance (and the inevitable arbitrariness in any statistics), the statistical significance of all these anomalies is indeed a debatable issue \\cite{Magueijo:2006we,Land:2006bn,Copi:2005ff,Hajian:2006ud}.  Various possible cosmological effects have been proposed as possible explanations for those anomalous features. For instance, to introduce a preferred axis, one has to generate some isotropy breaking. According to whether this occurs at an early time or at late times, one may classify the models into those in which an originally isotropic CMB fluctuation is distorted on its way to us \\cite{Gordon:2005ai} and into those whose statistical anisotropy is imprinted in the primordial fluctuations \\cite{Armendariz-Picon:2005jh}. A primordial imprint has been suggested to be left by isotropization happening during inflation \\cite{Gumrukcuoglu:2006xj,Ackerman:2007nb}, parity violating couplings  \\cite{Alexander:2006mt} vector field gradients \\cite{Armendariz-Picon:2007nr} or non-standard spinors \\cite{Bohmer:2007ut}. On the other hand, the apparent alignments could also be caused by local effects in our neighborhood. Notice, however, that the Axis of Evil has been argued to probably persist after foreground removal \\cite{deOliveira-Costa:2006zj}. Local voids have been though suggested as its possible origin \\cite{Inoue:2006rd} (and quite interestingly such voids could explain the present days acceleration too \\cite{Enqvist:2006cg}).  In the case CMB is distorted during the late acceleration, the signatures of anisotropy would automatically be seen at the smallest multipoles of the CMB, since the perturbation wavelengths corresponding to these angles enter inside the horizon at the same epoch that the dark energy dominance begins. Therefore one would have one coincidence problem less. The apparent statistical anisotropy has been associated with dark energy in the form of vector fields \\cite{Armendariz-Picon:2004pm,Boehmer:2007qa,Libanov:2007mq,us}, shear viscous fluids \\cite{Koivisto:2005mm} and elastic solids \\cite{Battye:2006mb}. See also \\cite{o1,o2,o3,o4} for other investigations. In general, if dark energy is something else than the cosmological constant $\\Lambda$-term, one expects it to have anisotropic stresses at least at the perturbative level. This is also a generic prediction of modified gravity theories \\cite{val1,val2,Koivisto:2006ie,Koivisto:2005yc,clifton,skordis,amar}. Moreover, it could be used to distinguish among scalar field theories in which the possible nonminimal coupling occurs only at the perfect fluid matter sector \\cite{doug1,doug2,Koivisto:2005nr,luca}. Such perturbative shear stresses have in general been attempted to be constrained through parameters describing directly the difference of the metric variables \\cite{Caldwell:2007cw} and the viscous properties of generalized fluids \\cite{Ichiki:2007vn,Mota:2007sz,brook1}. Unfortunately, the perturbative anisotropic stress could escape detection, when its influences are restricted to so large scales that they are only seen through the amplitude of the first few multipoles of the CMB. In spite of that, in many specific models the Jeans scale of the (effective) dark energy is small enough for dark energy to form small-wavelength perturbations which might be probed by e.g. weak lensing experiments \\cite{amendola,uzan,bruck,nunes}. Moreover, a {\\it statistical} anisotropy due to dark energy would leave specific signatures which are not described alone by the amplitude of the CMB angular power spectrum $\\sim \\langle a_{lm},a^*_{lm} \\rangle$, where $a_{lm}$ are the coefficients of the spherical expansion of the anisotropy. Instead, correlations between different multipoles $l$ are predicted, and also $m$-dependent patterns (whereas statistical isotropy would predict, on average, the same value regardless of $m$). The theoretical prediction from statistical anisotropy to the CMB is thus the nondiagonality of the matrix $\\langle a_{lm}, a^*_{lm} \\rangle$.  Here we continue a previous study which investigated the implications and origins of an anisotropic acceleration \\cite{Koivisto:2007bp}. There we studied exact anisotropic but homogeneous solutions of the cosmological equations. We now extend such investigation and generalise to the case where inhomogeneities are present and introduce a covariant framework for the study. The Bianchi I model is sufficiently simple to allow semi-analytical calculations, and it captures the basic effect present also in more complicated anisotropic models by featuring their important common property, direction-dependent expansion rates. Such metric is useful to describe  magnetic fields \\cite{Barrow:1997sy} or their Yang-Mills generalizations \\cite{Barrow:2005df}. As is well known, the anisotropy due to such fields tends to decay. Usually the constraints inferred on the Bianchi type I models have indeed been for cases in which the universe may initially be anisotropic but will then isotropize. Solutions which do not isotropize would require anisotropic matter sources, which in addition should dilute slower than radiation or dust. There has been no evidence that such cosmological sources would exist. However, as the universe is presently believed to be dominated by some source with unexplained negative pressure, a question we would like to ask is whether this pressure could also be anisotropic. This is of interest even without the apparent hints of statistical isotropy at the CMB.  After the study of Barrow on constraining anisotropic stresses in the late universe \\cite{Barrow:1997sy}, we note that very recently there has been also interest in employing Bianchi I to study the CMB. Campanelli {\\it et al}  \\cite {Campanelli:2006vb} describe magnetic fields present at the last scattering epoch resulting in an ellipsoidal universe. Similar effects have been proposed due to moving dark energy \\cite{BeltranJimenez:2007ai}. Rodrigues has suggested that an anisotropic cosmological constant \\cite{Rodrigues:2007ny} could result from an infrared noncommutative property of the spacetime. Longo has claimed that handedness of spiral galaxies is anisotropic  \\cite{Longo:2007gr} and that this, (together with the Axis of Evil) could be a result of large-scale magnetic fields \\cite{Longo:2007pc}. If true, this would be a strong further motivation for our study since our model could unify magnetic and dark energy fields by describing Yang-Mills field or the mutually coupled system. In addition to these approaches, the anisotropic expansion rates have been exploited in inflationary considerations \\cite{Bohmer:2007ut,Gumrukcuoglu:2006xj,Ackerman:2007nb,barrow1,barrow2}. Then the appearance of anomalous features can be consistent, but does not necessitate, the existence of anisotropic sources. Inflation can smooth out initial anisotropies and inhomogeneities, but if the universe did not undergo too many e-folds of inflation \\cite{burgess,easson}, some traces of such possible anisotropies could now be detectable at the largest scales. On the other hand, if one assumes isotropic initial conditions for the inflation, anisotropies could originate during inflation from a nonstandard energy source. This could be the same as dark energy.  We attack the problem at three fronts: phenomenologically, theoretically and observationally. In section \\ref{basics} we present our phenomenological description of anisotropic sources in terms of the skewness parameters $\\delta$ and $\\gamma$, and derive the basic equations describing the model as a dynamical system. We then check the generic asymptotic evolution of the universe in some simple cases. This is a unifying description which can be used for any particular model of inflation or dark energy. Especially we contemplate the possibility of generalizing the usual $\\Lambda$ term in such a way that it would correspond to a constant vacuum energy that may exert anisotropic pressure. Possibilities of unifying inflation and dark energy arise therefore naturally. Section \\ref{bounds} is devoted to the study of the implications of these models to cosmological observations. Explicit constraints are derived from the luminosity distance-redshift relationship of the supernovae of type Ia (SNIa) and the amplitude of the quadrupole anisotropy in the CMB. We find that the additional parameters we introduced, $\\delta$ and $\\gamma$, are well constrained. Finally, the last part of this article considers the inhomogeneities in the models considered here. In Section \\ref{covariant} we briefly introduce our notation for the covariant formalism and then present the general equations describing anisotropic cosmologies. Then we apply these general results into some interesting cases. In section \\ref{structure} we study then large scale structure formation under various specific assumptions about the model. We conclude in section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We investigated the possibility that the accelerated expansion of our universe, both during the early inflation and the present dark energy domination, is driven by a component with an anisotropic equation of state. The motivation for this investigation comes from the frequent appearance vector fields and possible Lorentz violations in fundamental theories, from the need to the test the robustness of the basic assumptions of cosmology, and from the hints of statistical anisotropy of our universe that several observations seem to suggest.  From the examination of the background dynamics we learned that  \\begin{itemize}  \\item There exist (anisotropic) scaling solutions even without coupling between the components.  \\item An anisotropic generalization of the $\\Lambda$-term has nontrivial dynamics given if its isotropic pressure is time-varying.  \\item The quadropole constraint on the anisotropy may be avoided, and the present SNIa data allows large anisotropy.  \\end{itemize}  From the analysis of system of perturbations the following was found  \\begin{itemize}  \\item Even in the limit of weak anisotropy, the linearized perturbations will be direction-dependent.  \\item The linearized anisotropy is the sum of FLRW perturbations and Bianchi I background only for perfect-fluid cosmology.  \\item Either in the limit of vanishing entropy or in the limit of homogeneous field the almost FLRW treatment is not valid.  \\end{itemize}  Let us elaborate on these and other results.  We started by presenting a phenomenological description of a such an anisotropy in terms of the skewness parameters $\\delta$ and $\\gamma$, and derived the basic equations describing the model as a dynamical system, and obtained the generic asymptotic evolution of the universe in some simple cases. Several scaling solutions were found both within the matter dominated epoch as well as during an accelerated expansion phase (with the summary of fixed points presented in the Table \\ref{tab1}). This seems to be interesting, since while in the FLRW universe it has been proven difficult to address the coincidence problem by finding a model entering from a matter-dominated scaling solution to an accelerating scaling solution, allowing for the presence of three expansion rates now opens up the possibility of describing a universe entering from a perfect fluid dominated scaling to an anisotropically accelerating scaling era, which might eventually help to understand the coincidence problem. In that framework, it is easier to find models with natural parameter values avoiding the fine-tuning problems of cosmological constant or scalar field quintessence \\cite{Boehmer:2007qa,Koivisto:2007bp,Jimenez:2008au}.  We also contemplated upon the possibility of an anisotropic generalization of the cosmological constant constant term. Interestingly, we found that a model with a constant density and constant anisotropic pressures ($\\rho = \\rho_0$, $\\delta = \\delta_0$, $\\gamma = \\gamma_0$) is indeed a consistent covariant modification of the Einstein's cosmological term which could generate a viable expansion history (with small enough $\\delta_0$ and $\\gamma_0$) that exhibits nontrivial anisotropic properties. We mentioned noncommutative properties of spacetime as a possible origin of anisotropies in a cosmological $\\Lambda$-term. One could consider different motivations by writing the vacuum energy as an integral of the zero-point energies over the $\\bf{k}$ up to some cut-off (summed over every possible field). If there were any isotropy-breaking physical principle which allowed only wavevectors with specific directions, the vacuum energy would acquire nonzero $\\delta$ and $\\gamma$. More specifically, if the cut-off depended not only on the magnitude $k$ of the wavevector ${\\bf{k}}$, but also on its direction (i.e. some directions would be excluded from the integral or contribute with less weight to the vacuum energy) this could explain both the smallness of the observed cosmological constant and the apparent anisotropic features in the cosmological data.  The amount of possible anisotropy, as described by the skewness parameters, can be constrained by the CMB and SNIa observations. We found that the present CMB data constraints can be much tighter than the constraints from the anisotropies in the luminosity-distance -redshift relationship of SNIa. However, there are classes of models where the CMB bounds from the background quadrupole anisotropy can be avoided altogether while there is significant anisotropy at late times that could be observed in the late SNIa distributions and which could leave statistically anisotropic imprints at the large-angle CMB. We point out, however, that in more general models, e.g. with time-varying $\\delta$ and $\\gamma$, one could allow more anisotropy. It is in principle possible for an arbitrarily anisotropic expansion to escape detection from CMB (considering only the effects from the background), as long as the expansion rates evolve in such a way that $e_z=e_y=0$. In other words, the (background) quadrupole vanishes, if each scale factor has expanded - no matter how anisotropically - the same amount since the last scattering.  We also considered the inhomogeneous cosmology of these models. The covariant formalism was employed to derive the general evolution equations for the perturbations in a universe with background anisotropies, i.e. imperfect fluid terms and shear of the metric. Our formalism and equations thus provide a tool to study fully consistently and covariantly the fundamentally anisotropic properties of matter inhomogeneities in cosmology, which may be hidden in the standard description of the overdensity scalar perturbation. One may witness the anisotropy by monitoring the solutions of the individual equations of the system (\\ref{d_evol2}), while summing their gradients yields the scalar quantity which has a straightforward interpretation as the overdensity in the standard FLRW cosmology, which at small scales does not depend on the gauge choice.  The implications of anisotropic sources for the formation of structures were discussed in several limiting cases. In particular, in the cases  examined, the deviations from isotropy are small enough that one may treat all the imperfect terms as first order perturbations. We noticed that this gives a result which is not obtainable by a straightforward generalization of the usual FLRW equations. If one considers the evolution equation of matter perturbations (in its closed form (\\ref{d_evol2})), one notices that while the metric shear would enter only at the second order, the imperfect fluid terms are present at the first order level. Thus, even when the Bianchi universe would be close to the FLRW model, one cannot consider the anisotropies in the perturbed Bianchi universe as a sum of FLRW perturbations and the small anisotropies of the exact Bianchi model. This is seen from the fact that when an imperfect source is present, its perturbations (\\ref{a_source}) can break explicitly the statistical isotropy already at the linear level. In general, the perturbation structure of these models allows a rich array of possibilities. The anisotropic properties seem to usually extend to the smaller scales also, but this feature is model-dependent and the constraints it implies must be studied on a case-by-case basis.  The present SNIa data allows anisotropic acceleration, but future SNAP data could set things straight about the skewness of dark energy and so of its nature. Such possibility would open a completely new window not only on the nature of the CMB anomalies but also into high energy physics models beyond the usual isotropic candidates of inflation or dark energy models such as the inflaton, the quintessence scalar fields or even an isotropic cosmological constant. Hence, an anisotropic equation of state may not only successfully unify the early inflation with the present days acceleration, but might also be the culprit for both the apparent problem of cosmic coincidence as well as the large-angle anomalies in the CMB.   The detailed signatures of the anisotropies, either originating from the early or the late accelerating epoch of the universe, to the full sky maps of the CMB should be determined: could these signatures rule some or all of the models out, or could they in fact be identified with the observed anomalies? Investigations along these lines are ongoing."
    },
    "0801/0801.4921_arXiv.txt": {
        "abstract": "We study the perturbative behaviour of charged topological-AdS black holes. We calculate both analytically and numerically the quasi-normal modes of the  electromagnetic and gravitational perturbations. Keeping the charge-to-mass ratio constant, we show that there is a second-order phase transition at a critical temperature at which the mass of the black hole vanishes. We pay special attention to the purely dissipative modes appearing in the spectrum as they behave singularly at the critical point. ",
        "introduction": "It is well known that if a black hole is initially perturbed, the surrounding geometry will start vibrating into quasi-normal oscillation modes, whose frequencies and decay times depend only on the intrinsic features of the black hole itself, being insensitive to the details of the initial perturbation. The radiation associated with these modes is expected to be seen with gravitational wave detectors in the coming years, giving valuable information on the properties of black holes. For these reasons, quasi-normal modes (QNMs) of black holes in asymptotically flat spacetimes have been extensively studied (for reviews, see~\\cite{KS,N}).  The Anti-de Sitter - conformal field theory (AdS/CFT) correspondence has led to an intensive investigation of black hole QNMs in asymptotically AdS spacetimes. Quasi-normal modes in AdS spacetime were first computed for a conformally invariant scalar field, whose asymptotic behaviour is similar to flat spacetime~\\cite{CM}.  Subsequently, motivated by the AdS/CFT correspondence, Horowitz and Hubeny made a systematic computation of QNMs for scalar perturbations of Schwarzschild-AdS (S-AdS) spacetimes~\\cite{HH}.  Their work was extended to electromagnetic and gravitational perturbations of S-AdS black holes in~\\cite{CL}. The study of scalar perturbations was further extended to the case of Reissner-Nordstr\\\"om-AdS (RN-AdS) black holes in~\\cite{WLA}. Finally, the QNMs of  scalar, electromagnetic and gravitational perturbations of RN-AdS black holes were presented in~\\cite{kokkotas} using the results of~\\cite{MM}.  The QNMs of AdS black holes have an interpretation in terms of the dual conformal field theory (CFT) \\cite{cft}. According to the AdS/CFT correspondence, a large static black hole in AdS corresponds to an (approximately) thermal state in the CFT. Perturbing the black hole corresponds to perturbing this thermal state, and the decay of the perturbation describes the return to thermal equilibrium. So we obtain a prediction for the thermalization timescale in the strongly coupled CFT. In ref.~\\cite{HH} it was shown that the QNMs for the scalar perturbations of large Schwarzschild-AdS black holes scaled with the temperature and it was argued that the perturbed system in the dual description will approach to thermal equilibrium of the boundary conformal field theory. However, when the black hole size is comparable to the AdS length scale there is a clear departure from this behaviour. It was then conjectured that this behaviour may be connected with a Hawking-Page phase transition \\cite{phase1, phase2} which occurs when the temperature lowers.  These results were further confirmed in \\cite{CL,WLA,kokkotas,Konoplya}. However, the behaviour of QNMs for small black holes is still poorly understood. Another interesting finding of the electromagnetic and gravitational perturbations is that  purely dissipative modes appear in the spectrum which are pure imaginary QNMs. In such perturbed classical backgrounds the presence of these dissipative modes indicate that the boundary theory reaches thermal equilibrium with no oscillations. It was shown in \\cite{CL} that for axial perturbations of Schwarzschild-AdS black holes these highly damped modes scale as the inverse of the black hole radius and this behaviour persisted in the case of Reissner-Nordstr\\\"{o}m-AdS black holes \\cite{kokkotas}.  It was observed in \\cite{Martinez:2004nb} that  for small black holes as we lower the temperature to a critical value there is a phase transition of a vacuum topological black hole towards a hairy black hole (MTZ). This claim was supported in \\cite{Koutsoumbas:2006xj} by calculating the QNMs of electromagnetic perturbations of the MTZ black hole and topological black holes. It was found that there is a change in the slope of the QNMs as we decrease the value of the horizon radius below a critical value, and this change has been attributed to the phase transition.  It was also observed in \\cite{Koutsoumbas:2006xj} that for small black holes  the quasi-normal frequencies converge toward the imaginary axis, i.e., their real part decreases and after the first few quasi-normal frequencies, it vanishes, indicating that for black hole radius smaller than the AdS length scale  there are only a finite number of QNMs. It was shown that the finite number of such modes for small horizons is due to the existence of bound states behind the horizon, which is an unobservable region. Further evidence that the behaviour of QNMs may provide indications of a phase transition was provided in \\cite{Shen:2007xk}.   Topological black holes~\\cite{mann}-\\cite{Birmingham:2006zx}, having hyperbolic horizons, introduce new features not present in spherical black holes. In~\\cite{Emparan:1999gf} it was shown that hyperbolic black holes can be described as thermal Rindler states of the dual conformal field theory in flat space. It was also found that, for small topological black holes of the size of the AdS length scale, the entropy at strong coupling is larger than the entropy obtained from field theory at lowest perturbation order at weak coupling. One possible explanation put forward was that there is a phase transition of topological black hole to vacuum AdS space at a critical temperature, which however was not observed in~\\cite{Emparan:1999gf} (see also~\\cite{Birmingham, Emparan:1998pf}).   In this work we make a detailed study of electromagnetic and gravitational perturbations of charged topological black holes (CTBH) in AdS space. Studying these perturbations of the background geometry, we show that the second-order phase transition observed at a critical temperature in~\\cite{Koutsoumbas:2006xj} occurs in more general configurations including charge. We calculate both analytically and numerically the QNMs of axial and in some cases the polar perturbations of large and small black holes. We find that for large black holes the QNMs irrespectively of the value of the charge, exhibit a negative slope. However, if the value of the black hole radius is smaller than the AdS length scale and the charge is small, the propagating QNMs (whose frequencies have a non-vanishing real part) of both axial and polar perturbations are finite with a positive slope. We attribute this behaviour to a second order phase transition which occurs as the temperature approaches a critical value (or the black hole radius is approaching the length scale of the AdS space). We also find that, in the case of small black holes, if we increase the charge, the number of propagating QNMs is again infinite and a part of positive slope coexists with the negative slope frequencies. This observation indicates that in the case of small black holes, the charge plays the r$\\hat{o}$le of an order parameter and according to the AdS/CFT correspondence we expect that the dual boundary theory is described by a thermal state with coexisting phases.   We  also study the purely dissipative modes appearing in the spectrum, both analytically and numerically of electromagnetic and gravitational perturbations of the charged topological black holes. For large black holes the purely dissipative modes of both axial and polar perturbations scale linearly with temperature. Also the intermediate black holes to a good approximation depend linearly on the temperature. This is the expected behaviour for large and intermediate black holes and they agree with the normal QNMs behaviour discussed in~\\cite{HH}. If the black hole radius is smaller than the AdS length scale then we find a clear departure from linearity with temperature. If the charge is small the purely dissipative modes scale with the temperature according to $a+b/(T-T_0)$ where $a, b$ are constants. Then, for a fixed charge to mass ratio we observed an infinite change of slope at $T=T_{0}$ signaling a second order phase transition. As we increase the charge the temperature dependence changes drastically.  The paper is organized as follows. In section \\ref{sec2} we review the basic properties of the topological black holes and we discuss their thermodynamics. In section \\ref{sec3} we present the analytical calculations of the QNMs of electromagnetic and gravitational perturbations of the CTBH. In section \\ref{sec4} we study numerically their behaviour and in section \\ref{sec5} we discuss in detail the purely dissipative modes appearing in the spectrum. Finally, section \\ref{sec6} contains our summary. ",
        "conclusions": "\\label{sec6}  We have studied the perturbative behaviour of the charged topological black holes. We have calculated both analytically and numerically the QNMs of electromagnetic and gravitational perturbations of these black holes.  For large black holes we found analytically that the axial $Z_{2}$ QNMs are purely dissipative, depend on the charge and scale as the inverse of the black hole horizon. The $Z_{1}$ axial modes are proportional to the black hole horizon but analytical expressions cannot be obtained in general for non-zero charge. For zero charge the potentials for both axial and polar modes reduce to electromagnetic potential and the wave function can be written in terms of the Heun function, leading to a semi-analytic expressions of the QNMs.  For small black holes, at the critical point with zero charge and mass the wave equation simplifies and the QNMs can be explicitly calculated. For small changes around the critical point, the real part of the $Z_{2}$ modes increases above the critical point giving a positive slope, whereas below the critical point it gives a negative slope.  Above the critical point, for these modes there is a critical value of $\\xi$ below which there are only purely dissipative modes. Below the critical point, there are no purely dissipative modes for any value of $\\xi$. The $Z_{1}$ modes exhibit the opposite behaviour. Similar behaviour is exhibited by the polar modes.  These results are also supported by  numerical investigations of the QNMs. The numerical results show clearly a change of slope of QNMs around a critical temperature for all kinds of perturbations. We found that the purely dissipative modes scale linearly with temperature for large black holes, while for small horizons they scale according to $a+b/(T-T_{0})$. Then, for a fixed charge to mass ratio we observed an infinite change of slope at $T=T_{0}$ signaling a second order phase transition.  The numerical results show also an interesting dependence of the modes on the charge of the black hole. For small horizons and small charge the number of propagating ($\\Re\\omega \\ne 0$) QNMs is finite, while as the charge increases, positive slope frequencies coexist with frequencies of negative slope and the number of propagating QNMs is again infinite. As the charge increases, a drastic change in the temperature dependence occurs.  \\newpage"
    },
    "0801/0801.1503_arXiv.txt": {
        "abstract": "{ I provide an incomplete inventory of the astronomical variability that will be found by next--generation time--domain astronomical surveys.  These phenomena span the distance range from near--Earth satellites to the farthest Gamma Ray Bursts.  The surveys that detect these transients will issue alerts to the greater astronomical community; this decision process must be extremely robust to avoid a slew of ``false'' alerts, and to maintain the community's trust in the surveys.  I review the functionality required of both the surveys and the telescope networks that will be following them up, and the role of VOEvents in this process.  Finally, I offer some ideas about object and event classification, which will be explored more thoroughly by other articles in these proceedings.  } ",
        "introduction": "Next generation surveys such as Pan-STARRS\\footnote{http://pan-starrs.ifa.hawaii.edu} and LSST\\footnote{http://www.lsst.org} promise to open up the time domain of astronomical variability to the general community as a service. This will allow the global study of all on--going phenomena in real--time, enabling both the small aperture amateur astronomer and the large aperture professional.  It also places great responsibility on the surveys themselves to provide a reliable stream of information.  The scope of the planned data release is unprecedented; 10$^{5-6}$ transient ``alerts'' are predicted to be generated on a nightly basis by the LSST alone.  This will place a huge burden on follow--up networks in the very near future.  Undertaking follow--up of these alerts could easily consume {\\it all} of the available global telescope resources unless intelligent decisions are made about which events to focus on.  The volume of the data streams will preclude this decision from being made by a human.  To enable intelligent, autonomous follow--up systems, the VOEvent protocol has been developed as a means of automatically conveying information between astronomical resources.  To take advantage of this stream, Heterogeneous Telescope Networks (HTNs) are being implemented to undertake follow--up of alerts.  This purpose of this conference is to address the practical implementation of this marriage.  This paper will address the types of astronomical variability the surveys will have to sort through, the types of alerts that the surveys will generate, and present various ideas on the classification of these alerts.  I will emphasize that the surveys must do more than generate alerts; they must also listen to the follow--up results of the community or risk retaining outdated classifications based upon their own limited data. ",
        "conclusions": "I have detailed some of the interplay between surveys and follow--up networks in time--domain astronomy, highlighting their potential roles and responsibilities.  While the concept of data federation using VOEvent's citation mechanism may work, it may also lead to different inferences by different resources.  An alternative is to broker the evolving ``truth'' regarding an event by trusted agents in the system. These agents have the potential to fully realize (or scuttle) the successful interplay between surveys and follow--up HTNs.  The surveys and agents each play a role in the ultimate classification of each event.  I emphasize that the surveys should also {\\it listen} to the network, and decide if and how to allow these external resources to influence their internal event classifications.  Finally, the need for accurate and precise data models is commensurate with the data flow that will be compared against them.  These models should be built sooner rather than later using currently existing data sets, to ensure that the promise of the first several years of these surveys is not lost to faulty or inaccurate alerts."
    },
    "0801/0801.4618_arXiv.txt": {
        "abstract": "{In these proceedings I discuss various extragalactic surveys which will be undertaken over the next few years and which will be complementary to any H{\\sc i} and/or continuum surveys with the SKA-pathfinder telescopes. I concentrate on the near-infrared public surveys which will be undertaken with the Visible and Infrared Survey Telscope for Astronomy (VISTA), and in particular the VISTA Deep Extragalactic Observations (VIDEO) survey which will provide the ideal data set to combine with any deep SKA-pathfinder observations of the extragalactic sky. After highlighting the links that the SKA pathfinders have with the various VISTA surveys, I briefly describe an approved large area survey to be carried out with the {\\sl Herschel Space Observatory} which has a large scientific overlap with the SKA pathfinder telescopes.}  \\FullConference{From planets to dark energy: the modern radio universe\\\\ October 1-5 2007\\\\ University of Manchester, Manchester, UK}  \\begin{document} ",
        "introduction": "The Square Kilometre Array promises to be one of the most influential telescopes of any era in astronomy. In the area of galaxy formation and evolution, it will be able to map out the neutral Hydrogen content of galaxies up to and possibly beyond $z \\sim 1.5$. However, it will be insensitive to the underlysing stellar populations and various physical mechanisms that are associated with galaxies which do not emit at radio wavelengths. Thus, even in the age of the SKA complementary multi-wavelength surveys will still be paramount to further our understanding of the Universe. Furthermore, the various SKA pathfinder telescopes will also benefit greatly from multi-wavelength surveys which are being or will be carried out concurrently with the deep radio surveys.  Such surveys are already underway or being planned, and many more will undoubtedly be initiated in the coming years. In these proceedings I focus on the near-infrared surveys which will be carried out by the ESO-Visible and Infrared Survey Telescope for Astronomy (VISTA), to begin operation in 2008 at Paranal, Chile, in particular the second deepest tier of the extragalactic public surveys to be undertaken on this telescope, namely the VISTA Deep Extragalactic Observations (VIDEO) survey.  I also highlight a recently approved large-area survey to be conducted with the {\\sl Herschel Space Observatory} which I believe will be of benefit and benefit from observations with the various SKA pathfinder telescopes. ",
        "conclusions": ""
    },
    "0801/0801.1517_arXiv.txt": {
        "abstract": "In the synchrotron radiation model, the polarization property depends on both the configuration of the magnetic field and the geometry of the visible emitting region. Some peculiar behaviors in the X-ray afterglows of {\\it Swift} gamma-ray bursts (GRBs), such as energetic flares and the plateau followed by a sharp drop, might by highly linearly-polarized because the outflows powering these behaviors may be Poynting-flux dominated. Furthermore, the broken-down of the symmetry of the visible emitting region may be hiding in current X-ray data and will give rise to interesting polarization signatures. In this work we focus on the polarization accompanying the very early sharp decline of GRB X-ray afterglows. We show that strong polarization evolution is possible in both the high latitude emission model and the dying central engine model which are used to interpret this sharp X-ray decline. It is thus not easy to efficiently probe the physical origin of the very early X-ray sharp decline with future polarimetry. Strong polarization evolution is also possible in the decline phase of X-ray flares and in the shallow decline phase of X-ray light curves characterized by chromatic X-ray VS. Optical breaks. An {\\it XRT}-like detector but with polarization capability on board a {\\em Swift}-like satellite would be suitable to test our predictions. ",
        "introduction": "\\label{sec:Observation} The successful launch of the {\\it Swift} satellite in Nov 2004 opened a new window to reveal what has been happening in the early afterglow phase of gamma-ray bursts (GRBs). As summarized in Zhang et al. (2006) and Nousek et al. (2006), in a canonical {\\it Swift} GRB X-ray afterglow lightcurve some interesting features are emerging (detected in a good fraction of but not all bursts), including a very early sharp decline (phase-I), a shallow decline or even plateau, (phase-II), preceding the conventional or so-called ``normal\" decay phase (phase-III \\& -IV), and the energetic X-ray flares (phase-V). Various modifications of the standard GRB afterglow model have been put forward to explain the observations (see M\\'esz\\'aros 2006; Piran \\& Fan 2007; Zhang 2007 for recent reviews). Although the underlying physical processes are not very clear, it is widely suggested that the X-ray Phases I, II and V may be relevant to prolonged activities of the central engine while the simultaneous optical afterglow is dominated by GRB forward shock emission (Fan \\& Wei 2005; Zhang et al. 2006; Ghisellini et al. 2007; Panaitescu 2007). If so, the early time UV/optical and X-ray polarization behaviors should be independent and may be very different from each other. In some models, the outflows powering energetic flares and X-ray plateaus followed by a sharp drop are Poynting-flux dominated (e.g., Gao \\& Fan 2006; Giannios 2006; Troja et al. 2007), thus high-linear polarization of these synchrotron radiation photons is expected (Fan, Zhang \\& Proga 2005).  Except the configuration of the magnetic field, the geometry of the emitting area can influence the polarization property of the synchrotron radiation as well. The broken-down of the symmetry of the visible emitting region may be hiding in the peculiar X-ray data and can give rise to interesting polarization signatures. Motivated by this idea, in section 2 we calculate the polarization property of Phase-I and show that strong evolution is likely. In section 3, we argue that similar phenomena are also expected in the decline phase of the X-ray flares and possibly in the late time part of Phase-II, especially those associated with a chromatic break in optical band. We also discuss the detection prospect there. In section 4, we summarize our work with some discussions. ",
        "conclusions": "Many surprises in GRB afterglows, mainly in the X-ray band, have been brought since the successful launch of {\\it Swift} in 2004. Though not clearly understood yet, the (re-)activity of GRB central engine is likely to play an important role on producing afterglows, as speculated by Katz, Piran \\& Sari (1998). The different temporal behaviors in the UV/optical and X-ray afterglow favor that the peculiar X-ray emission traces the prolonged activity of the central engine while the optical afterglow is dominated by the forward shock emission \\cite{FW05,GG07}. If true, the polarization behaviors of the early X-ray and optical afterglow may be independent and different. At such an early time, the linear polarization of the optical forward shock emission is expected to be very small \\cite{GL99} unless the interstellar medium is magnetized \\cite{GK03}, because the Lorentz factor of the outflow is so large that the visible emitting region is well within the cone of the ejecta and is thus symmetric. But in the X-ray band, interesting polarization signatures may be present. For example, Fan et al. (2005) argued that some X-ray flares might have a high linear polarization because these new outflows are likely to be Poynting-flux dominated. Such an argument is also applicable to the X-ray plateau followed by a sharp drop as detected in GRB 070110.  In this work, we focus on the polarization properties of the very early sharp decline of X-ray afterglows (Phase-I). In the high latitude emission model, the later the emission, the larger the angular $\\theta$ (see Figure \\ref{fig:Cartoon} for details). So at late time, the symmetry of the emitting ring is broken when part of the emitting ring is out of the cone of the ejecta. Strong polarization evolution thus emerges (see Section 2.1). For comparison, results from the dying central engine model are a bit more complicated. If the outflow powering the fast decaying X-ray tail emission always has a Lorentz factor $\\Gamma>1/(\\theta_{\\rm j}-\\theta_{\\rm v})$, no strong polarization evolution is expected unless the physical composition of the outflow changed significantly. However, if the Lorentz factor drops rapidly to satisfy $\\Gamma<1/(\\theta_{\\rm j}-\\theta_{\\rm v})$, then polarization evolution is evident. In this scenario, if the magnetic field involved is shock-generated and is thus random, then strong polarization evolution is expected; while if the magnetic field involved is brought from the central engine and is thus large-scale ordered, then weak polarization evolution is expected due to the significant polarization background (the degree is $\\sim 0.5$). One of our goals to calculate the polarization of Phase-I is to see whether it's possible, via X-ray polarimetry, to distinguish between the high latitude emission model and the dying central engine model. This goal is not easy to fulfill in view of the above comparison. However, our calculation does support that strong polarization evolution is likely to accompany Phase-I, the decline phase of X-ray flares (Phase-V) and possibly the late part of Phase-II. Fruitful novel features are expected to appear in GRB X-ray afterglow polarimetry, though detection of them is beyond the capability of current monitors. We have witnessed the breakthrough made by {\\it Swift} XRT and we are likely to experience it once more in the coming X-ray polarimetry era. One byproduct of the future X-ray polarimetry is to probe the cosmological birefringence effect that arises in some quantum gravity models (Fan, Wei \\& Xu 2007 and the references therein).  Finally we'd like to point out that the optical flares in a few bursts (e.g., B\\\"oer et al. 2006)  may also have a central engine origin (Wei, Yan \\& Fan 2006). Interesting polarization signatures are also expected."
    },
    "0801/0801.4568_arXiv.txt": {
        "abstract": "In the CDM cosmological framework structures grow from merging with smaller structures.  Merging should have observable effects on galaxies including destroying disks and creating spheroids.  This proceeding aims to give a brief overview of how mergers occur in cosmological simulations. In this regard it is important to understand that dark matter halo mergers are not galaxy mergers; a theory of galaxy formation is necessary to connect the two.  Mergers of galaxies in hydrodynamical simulations show a stronger dependence on mass than halo mergers in N-body simulations.    If one knows how to connect galaxies to dark matter halos then the halo merger rate can be converted into a galaxy merger rate.  When this is done it becomes clear that major mergers are many times more common in more massive galaxies offering a possible explanation of why Hubble type depends on galaxy mass. ",
        "introduction": "Structure growth via mergers is one of the main predictions of CDM type cosmologies.  However, what is predicted is the merger rates of dark matter halos, which are not directly observable.  Using dark matter halo merger rates to predict galaxy merger rates requires a theory of galaxy formation or at least a model of how galaxies populate dark matter halos.  In a similar way, what can actually be observed are close galaxy pairs, disturbed galaxies, or morphological differences between galaxies, all of which can only be indirectly tied to galaxy mergers using theoretical models. Thus connecting theory to observations poses a number of difficulties which are often not given enough attention.  In particular the halo merger rate is often used as an indicator of galaxy merger rates.   If galaxy mass scaled linearly with dark matter halo mass then this could possibly be true. But differences in the shapes of the galaxy stellar mass and halo mass functions imply that galaxy formation is much less efficient in low and high mass halos.  Thus we should expect that galaxy merger statistics should differ from halo merging statistics. ",
        "conclusions": ""
    },
    "0801/0801.2737.txt": {
        "abstract": "% context heading (optional) %{} leave it empty if necessary {Cosmic rays are confined in the atmospheres of galaxy clusters and, therefore, they can play a crucial role in the heating of their cool cores.} %, resulting in a alternative to mechanical heating processes of cluster cores.} % aims heading (mandatory) {%Aims. We discuss here the thermal and non-thermal features of a model of cosmic ray heating of clusters' cores that can provide a solution to the cooling-flow problems. To this aim, we generalize a model originally proposed by Colafrancesco, Dar \\& DeRujula (2004) and we show that our model predicts specific correlations between the thermal and non-thermal properties of galaxy clusters and enables various observational tests.} % methods heading (mandatory) {%Method. The model reproduces the observed temperature distribution in clusters by using an energy balance condition in which the X-ray energy emitted by clusters is supplied, in a quasi-steady state, by the hadronic cosmic rays, which act as ``warming rays'' (WRs). The temperature profile of the intracluster (IC) gas is strictly correlated with the pressure distribution of the WRs and, consequently, with the non-thermal emission (radio, hard X-ray and gamma-ray) induced by the interaction of the WRs with the IC gas and the IC magnetic field.} % results heading (mandatory) {%Results. The temperature distribution of the IC gas in both cool-core and non cool-core clusters is successfully predicted from the measured IC plasma density distribution. Under this contraint, the WR model is also able to reproduce the thermal and non-thermal pressure distribution in clusters, as well as their radial entropy distribution, as shown by the analysis of three clusters studied in details: Perseus, A2199 and Hydra. The WR model provides other observable features of galaxy clusters: a correlation of the pressure ratio (WRs to thermal IC gas) with the inner cluster temperature $(P_{WR}/P_{th}) \\sim (kT_{inner})^{-2/3}$, a correlation of the gamma-ray luminosity with the inner cluster temperature $L_{\\gamma} \\sim (kT_{inner})^{4/3}$, a substantial number of cool-core clusters observable with the GLAST-LAT experiment, a surface brightness of radio halos in cool-core clusters that recovers the observed one, a hard X-ray ICS emission from cool-core clusters that is systematically lower than the observed limits and yet observable with the next generation high-sensitivity and spatial resolution HXR experiments like Simbol-X.} % conclusions heading (optional), leave it empty if necessary {%Conclusions. The specific theoretical properties and the multi-frequency distribution of the e.m. signals predicted in the WR model render it quite different from the other models so far proposed for the heating of clusters' cool-cores. Such differences make it possible to prove or disprove our model as an explanation of the cooling-flow problems on the basis of multi-frequency observations of galaxy clusters.} ",
        "introduction": "\\label{sec.intro}  Theoretical description of the intra-cluster (IC) plasma in the central regions of many clusters predicts that it is cooler than in their outskirts, because it radiates X-rays at such a rate that the plasma cooling time is much shorter than the age of the cluster. However, such rapid cooling is not observed and the central-temperature depression in cluster's cores is not as deep as expected on the basis of the plasma's cooling rate: the observed central temperature $T_{\\rm inner}$ settles, in fact, at a fraction $\\sim 1/3 -1/2$ of the outer temperature $T_{\\rm outer}$ (see, e.g. McNamara 1997,  Peterson et al. 2003, Piffaretti et al. 2005, Donahue et al. 2005, see also Bregman 2004 for a review). The lack of cooling gas with $T<T_{inner}$ in cluster cores is also supported by the absence of line emission corresponding to gas below $T_{\\rm inner}$, as shown by high spatial resolution imaging with Chandra (e.g. McNamara et al. 2000; Fabian 2000; Blanton et al. 2001; Allen et al. 2001; Lewis et al. 2002; Blanton et al. 2003) and high spectral resolution measurements with the XMM-Newton RGS instrument (Peterson et al. 2001, Tamura et al. 2001, Kaastra et al. 2001; Kahn et al. 2002; Peterson et al. 2002) and EPIC instrument (B\\\"ohringer et al.~2001, 2002; Molendi \\& Pizzolato 2001; Matsushita et al.~2002). These are referred here as the cooling--flow (CF) problems.  It has been widely recognized that some form of heating is necessary to quench the cooling flow and form a warm--core in the IC gas structure. Available sources of heating in clusters cores are provided by AGN jets and lobes, pressure waves, buoyant bubbles and cavities, intra-cluster shock waves, leptonic and hadronic cosmic-rays.\\\\ % The presence of a central radio source with jet and lobes, the mechanical heating provided by pressure waves produced by the jetted AGNs, the formation of buoyant bubbles filled with energetic plasma, the thermal conduction from the hot outer layers of the clusters, may certainly alleviate the CF problems (see e.g., B\\\"ohringer et al. 2002; Churazov et al. 2002; Ruszkowski \\& Begelman 2002, Fabian 2004, Voit \\& Donahue 2005, Br\\\"uggen \\& Kaiser 2002, Ruszkowski et al. 2004, Vernaleo \\& Reynolds 2006, Reynolds et al. 2005).\\\\ % Some difficulties with solutions along these lines, however, might stand out: % a first difficulty is that not all cooling flow clusters contain a powerful radio source at their center (e.g., the cluster RXJ0820.9+0752; Bayer-Kim et al.~2002, and the clusters A1650 and A2244, Donahue et al. 2005), nor a large amount of spherically distributed buoyant bubbles out to the cluster's core boundary. % Another difficulty is that the bremsstrahlung cooling rate is proportional to $n_{\\rm e}^2\\,T^{1/2}$, with $n_{\\rm e}$ the plasma's electron number density\\footnote{Hydrogen in the plasma is fully ionized, and we are using here $n_{\\rm p}\\approx n_{\\rm e}$.} and $T$ its local temperature: a suitable conduction and/or heating and pressure-building mechanism must somehow adapt itself to this behaviour, particularly as a function of $n_{\\rm e}(r)$, which varies by orders of magnitude along the clusters' radii. % Other difficulties stand out with the spatial distribution of the energy deposited by AGN jets and cavities: on one hand, the energy released from AGN jets and cavities seems to be sufficient to quench cooling flow and increase the central gas entropy (e.g., Sijacki \\& Springel 2006); on the other hand, it is not yet clear if such energy can be properly distributed to reproduce the observed cool-core temperature profiles (see, e.g., discussion by Vernaleo \\& Reynolds 2005, 2006, Heinz et al. 2006, Vernaleo \\& Reynolds 2007). Small duty cicles and a higher fraction of AGNs in cluster cores (see Bird et al. 2007) might however weaken the difficulties related to the AGN heating scenarios. In addition, there seems to be a small AGN activity in groups (Dwarakanath \\& Nath 2006).\\\\ % Sound waves and shocks do not seem to be efficient in quenching cooling flows (see, e.g., Fujita \\& Suzuki 2006). %  Irrespective of the nature of the heating source, it seems that a general property of the heating agent -- which derives from the smoothness and from the similarity of the heating distribution required to quench cooling in several clusters -- is that they must be spatially distributed in the whole cluster core. % In fact, it has been shown (Fabian \\& Sanders 2006, Fabian et al. 2006, Sanders \\& Fabian 2007) that the non-thermal pressure in the Perseus cluster (one of the best studied clusters with cool cores and mini radio-halos) is quite similar (slightly steeper) to the thermal one.\\\\ % In this context, it has been shown that a CR population in clusters that is radially distributed as the radial profile of the IC gas, is able to recover the observed radial structure of the IC gas temperature in several clusters (Colafrancesco, Dar \\& DeRujula 2004). % % The CRs in this last theory are not confined to the disk of the Galaxy, but permeate a much larger halo, being produced by \"cannonballs\" travelling for kiloparsecs in the interstellar medium of galaxies, decelerating by knocking out its constituents, which are thereby accelerated to become CRs\\footnote{When referring to CRs in CF clusters, we use the expressions ``CRs'' and ``WRs'' as synonyms.} (see, e.g., DeR\\'ujula 2004 for a review, and references therein).\\\\ % Beyond the debate on the cannonball picture (see, e.g., Hillas 2006; Dar \\& DeR\\'ujula 2006) and its relevance for the CR origin, we stress here that any mechanism which produces radial distributions of WRs similar to the profile of the IC gas can provide quenching of cooling flows by the action of WRs heating.  In this paper we consider and justify, on the basis of a consistent picture of CR diffusion in cluster atmospheres (see Appendix A), that the WR number density $n_{WR}(r)$ can be, in fact, distributed with a radial profile that is similar to that of the IC gas, and follows the general form $n_{WR} \\sim n^{\\alpha}_{e}$, i.e. it is proportional to a power $\\alpha$ of the IC gas density $n_e(r)$. % The total amount of WRs in our model and the slope $\\alpha$ of the $n_{WR}(r)$ profile are left as free parameters that will be subsequently constrained by fitting the temperature profile of the IC gas in each specific cluster core. The resulting CR density profile will be further tested against the entropy and pressure profiles in the cluster cores, the spectrum and surface brightness profile of the cluster radio halos and the future observations in the hard X-ray and gamma-ray bands.\\\\ % We will first show that the presence of WRs in cluster cores are able, in fact, to heat the IC gas and reproduce their radial temperature profiles. % We will then start from this constraint to discuss the consequences that such a population of WR hadrons (protons), produced and stored in the atmospheres of galaxy clusters, has for the gamma-ray emission, for the hard X-ray (HXR) emission and for the presence of diffuse radio emission of both cool-core and non cool-core clusters. % We will finally show that it is possible to use these predictions, in the framework of a multi-wavelength observational strategy, to set constraints on the amount and spatial distribution of WRs in clusters, and therefore disentangle between the WR model discussed here and other models for the onset of warm-cores in galaxy clusters.\\\\ % The structure of the paper is the following: in Sect.2 we delineate the theoretical framework that is able to model the temperature structure of the cluster cool cores and determine the main parameters (i.e., the radial distribution and the density of WRs) that fit the cluster temperature profiles. In Sect.3 we will also discuss the density and pressure distribution of WRs in the clusters' atmospheres that is consistent with the temperature profile. % In Sect.4 we will apply our procedure to a sample of ten well studied clusters. These clusters have been chosen according to the following criteria: i) 4 clusters which have cool cores and diffuse radio emission (these clusters certainly have a population of CRs in their centers); ii) 2 clusters with radio halos (i.e. with CRs in their atmospheres) but without cool cores, in order to compare their properties with the radio-active, cool-core clusters; iii) 4 clusters with cool cores but with no evident central diffuse radio emission in order to compare these non radio--active clusters with the previous ones. % In Sect.4 we also study the thermal and non-thermal pressure structure of the selected clusters and we will show that the pressure ratio $P_{WR}/P_{th}$ found at the cluster center correlates with the central gas temperature $T_{inner}$ providing evidence for the basic regulation mechanism induced by WRs. We will discuss the predictions of our model for some specific clusters (like Perseus, A2199 and Hydra) for which detailed thermal and non-thermal pressure information is available as well as information on their radial entropy distribution. % In Sect.5 we will discuss the expected gamma-ray luminosity and flux of the cluster sample here considered that is  correlated with the cool-core temperature structure predicted in our model. % In Sects.6 and 7 we will present analogous considerations for the expected radio and HXR emission that are constrained by the cluster temperature structure. % We will discuss the differences between our model and other scenarios for the cool-core heating in Sect.8 and we will resume our main conclusions in Sect.9.\\\\ % Throughout the paper, we use a flat, vacuum--dominated cosmological model with $\\Omega_m = 0.3$, $\\Omega_{\\Lambda} = 0.7$ and $h = 0.7$. ",
        "conclusions": "\\label{sec.conclusions}  The WR heating model that we presented in this paper is able to reproduce the temperature, the pressure and the entropy structure of both cool-core and non cool-core clusters and it provides results that are consistent with all the available evidence of non-thermal phenomena emerging from the cores of clusters, e.g. diffuse radio emission in the form of mini halos and diffuse radio halos, HXR emission limits obtained with BeppoSAX, Chandra and XMM, and gamma-ray limits obtained with EGRET.  This model provides a theoretical description of the physics of cool cores that is directly related to the presence of several observable consequences of the presence of WRs in clusters' atmospheres. The main conclusions of our work are: \\begin{itemize} \\item the presence of WRs in cluster atmosphere provides a simple solution for the temperature, pressure and entropy profiles of the thermal IC plasma and of the non-thermal plasma observed in the inner regions of several clusters studied in this paper. \\item The WR distribution produces a substantial emission of gamma-rays by their hadronic $pp \\to \\pi^0 \\to \\gamma \\gamma$ interations with the IC gas. A large fraction (6 clusters out of 10 considered in this study) have gamma-ray fluxes at $E > 100$ MeV detectable with the GLAST-LAT experiment. The gamma-ray luminosity of the cool-core and non-cool core clusters correlates with the cluster's inner temperature, and thus provides a specific scaling behaviour of our WR model that can be verified with the next coming GLAST observations. \\item The WRs distribution can reproduce the radio emission of all the cool-core clusters with mini-halos that we studied (A2199, Perseus, \\rxj) except for A2390 (whose mini radio-halo flux could be affected by residual radio emission from the lobes of the central cD radio galaxy). \\item The WRs distribution produces ICS emission with fluxes well below the HXR limits of A2199, Coma and A2163 provided by BeppoSAX and well below the not-thermal emission detected by Chandra and XMM in the central regions of Perseus. This means that ICS emission from WRs is not the explanation of the emission of these clusters in excess with respect to their thermal bremsstrahlung radiation. \\item The distribution of WRs in clusters directly relates the thermal IC gas temperature, pressure and entropy distribution in cool-core clusters (observable in the X-ray energy band) with non-thermal diffuse emission features observable at radio (synchrotron), HXRs (ICS), gamma-rays (mainly from $\\pi^0 \\to \\gamma \\gamma$ decay) and microwaves (by Sunyaev --Zel'dovich effect). We predict that the expected levels of these non-thermal emission features in the WR model will be testable with the next-coming experiments in the HXR band (Simbol-X), in the gamma-rays (GLAST-LAT) and in the microwaves (SPT). These experiments can determine the amount of WRs which are present in galaxy clusters. \\item These observable predictions makes the WR model entirely testable by using a multi-frequency observational strategy that links, therefore, thermal and non-thermal phenomena in clusters cores. \\item The specific theoretical and observational features of the WR model render it quite different from other models so far proposed for the heating of cool-cores in galaxy clusters. Such peculiar differences make it possible to prove or disprove the WR model and, in general, the properties of any model that has been proposed as an explanation of the cooling-flow problems. \\end{itemize}"
    },
    "0801/0801.0457_arXiv.txt": {
        "abstract": "{} {The origin of the observed variability of the gas-phase D/H ratio in the local interstellar medium is still debated, and in particular the role of deuterium depletion onto dust grains. Here we extend the study of the relationship between deuterium and  titanium, a refractory species and tracer of elemental depletion, and explore other relationships.} {We have acquired high resolution spectra for nine early-type stars using the VLT/UVES spectrograph, and detected the absorption lines of interstellar TiII. Using a weighted orthogonal distance regression (ODR) code and a special method to treat non symmetric errors, we compare the TiII columns with the corresponding HI, DI and also OI columns. In parallel we perform the same comparisons for available FeII data.} {We find a significant correlation between TiII/HI and D/H in our data set, and, when combined with published results, we confirm and better constrain the previously established trends and extend the trends to low HI columns.  We exclude uncertainties in HI and OI columns as the main contributor to the derived metals-deuterium correlations by showing that the TiII/HI ratio is positively correlated with DI/OI. We find a similar correlation between FeII/HI and DI/OI.  The TiII gradients are similar or slightly smaller than for FeII, while one would expect larger variations on the basis of the higher condensation temperature of titanium. However we argue that ionisation effects introduce biases that affect iron and not titanium and may explain the gradient similarity.  We find a less significant negative correlation between the TiII/DI ratio and the hydrogen column, possibly a sign of different evaporation of D and metals according to the cloud properties. More TiII absorption data along very low H column lines-of-sight would be useful to improve the correlation statistics. } {} ",
        "introduction": "The measurement of the present-day ratio of deuterium (D) to hydrogen (H) in the Galaxy places important constraints on its chemical evolution and has been the goal of intense work, from the Copernicus era (Rogerson and York, 1973) to HST (eg. Linsky, 1998), IMAPS (Jenkins et al., 1999, Sonneborn et al., 2000) and FUSE (e.g. Moos et al., 2002). Measurements inside the Local Bubble (LB), the volume of tenuous gas surrounding the Sun (e.g. Snowden et al., 1998, Welsh et al., 1999, Lallement et al., 2003) appear to be consistent with a single value for D/H (Moos et al., 2002, H{\\'e}brard \\& Moos, 2003, Oliveira et al., 2003), but other measurements (e.g. Jenkins et al., 1999, Sonneborn et al., 2000,  Hoopes et al., 2003, Friedman et al., 2006, H{\\'e}brard et al., 2005, Oliveira \\& H{\\'e}brard, 2006) suggest variations of the interstellar D/H ratio at larger distances, i.e. beyond $\\simeq$ 150-200 parsecs. Quantitatively, within and at the periphery of the LB (log N(HI) $\\leq$ 19.3 cm$^{-2}$), the D/H ratio is of the order of 15 ppm (see Figure 1 of Linsky et al., 2006), while lines-of-sight measurements with higher N(HI) show a strong variability, with maximum values of $\\simeq$ 23 ppm. For a number of lines-of-sight (LOS) with log N(HI) $\\leq$ 20.5 cm$^{-2}$, D/H is significantly lower, $\\simeq$ 7 ppm.  According to Draine (2004) and Linsky et al. (2006) (see also Linsky 2007), the wide range in the gas-phase D/H ratios is due to different amounts of deuterium depletion onto interstellar grains. Theoretical support for this interpretation is the preferential adsorption of deuterium onto interstellar grains, in particular PAHs (Jura, 1982, Draine, 2004). Observational support comes from the correlations between D/H and both Fe and Si depletions, as well as the correlations with the rotational temperatures of H$_{2}$ for lines-of-sight possessing molecular gas (Linsky et al., 2006). The D/H trends quoted above may be explained by the depletion hypothesis if one hypothesizes a recent heating of the LB and especially recent shocks in star-forming regions at its periphery. This interpretation has important consequences and implies that the total D/H ratio including deuterium in the gas and in the dust is at least 23 parts per million of hydrogen, which poses a challenge for models of galactic chemical evolution.  On the other hand, Oliveira \\& H{\\'e}brard (2006) show that high D/H values are also found along distant lines-of-sight with associated high interstellar columns. At high N(HI) the selection bias against high D/H values (due to dominance of dense, cold, and thus strongly depleted regions) should not lead to any variability (Oliveira et al., 2006). Moreover, D/O measurements do not fully support this depletion scenario, since the DI/OI ratio is found to be more constant than DI/HI (H\\'ebrard et al., 2005). Interestingly the LB D/H ratio derived from the D and O measurements is only of the order of 13 ppm (H{\\'e}brard \\& Moos 2003). These authors suggest that the uncertainties in the measurement of HI in the data play an important role in the subsequently derived relationships. Two quotients with the same denominator remain naturally correlated if there are measurement errors on this denominator. FeII/HI  vs. DI/HI or FeII/OI vs. DI/OI relationships could result from this effect.  Infall of external unprocessed (or less processed) gas also creates D/H variability, and indeed a continuous infall is a general property in the Galaxy (Romano et al., 2006). Support for the existence of a significant infall into the Sun's neighborhood comes from D, O, N, and $^{3}$He abundance values  found in the local interstellar cloud that surrounds the Sun. Geiss et al. (2002) interpret these abundances as being due to the local mixture of processed galactic gas and very weakly processed (Magellanic type) gas of intergalactic origin having fallen onto the disk. Such an infall could have occurred recently, leading to a present-day incomplete mixing. On the other hand, the wide D/H range implies a non realistic infall strength and a very weak mixing after the infall, as argued by Linsky et al. (2006).  It is therefore mandatory to investigate in more detail the depletion mechanisms and the associated deuterium behaviour. Prochaska et al. (2005) made an important step by showing a positive correlation between the abundance of titanium, a refractory species, and that of D/H. Here we present results we have obtained in the southern hemisphere with the Ultraviolet and Visual Echelle Spectrograph (UVES) at the ESO Very Large Telescope (VLT), and we add nine new TiII measurements towards D/H targets. In parallel to our observations and analysis, Ellison et al. (2007) also gathered and analysed new VLT-UVES  and Keck Observatory HIRES data. We have combined our results with their results and those from Prochaska et al. For targets in common, we have compared the two results. We have extended the correlative study, in particular to lower column-densities. We have used sophisticated methods to treat asymmetric errors in both correlated quantities.  Titanium is a particularly well suited element for the study of a potential correlation between depletion and deuterium abundance. It has a high depletion which is linked to its high condensation temperature. \\TiII\\ is optically accessible in the near ultraviolet via the $\\lambda\\lambda$ 3229.199,3242.994,3383.768 \\AA\\ absorption lines (wavelengths from Morton, 1991), and interstellar titanium has been observed in absorption since the 1970's (e.g. Wallerstein and Goldsmith, 1974, Stokes, 1978, Hobbs, 1984). Moreover, as already argued by these authors, \\TiII\\ is the dominant ionization state in HI regions due to the nearly exact coincidence of the ionization potentials of \\TiII\\ and HI. Also, \\TiII\\ is coupled to HI via charge-exchange reactions (Field and Steigman, 1971). The observed close N(\\TiII) vs. N(HI) relationship, in contrast with the more dispersed N(NaI)/N(HI) or N(CaII)/N(HI) relations (see e.g. Figure 7 of Welsh et al., 1997), results from those physical properties.  The detection of the interstellar TiII lines requires high spectral resolution (R $>$ 50,000) and appropriate optical systems and detectors adapted to the near ultraviolet wavelength range. Prochaska et al. (2005) used HIRES at the Keck Observatory and observed absorption towards 7 northern hemisphere stars. Due to this limited number of targets, it is important to extend their study to other D/H targets. Moreover, the observed trend is strongly dependent on the measurement towards Feige 110, and indeed H\\'ebrard et al. (2005) have shown that the column of HI is probably significantly underestimated for this star. They also quote the contradicting case of $\\gamma^{2}$Vel,  a line-of-sight with high D/H and low titanium abundance. Finally, while they covered a  range of logN(HI) from 19.85 and 21.05 cm$^{-2}$, it is important to extend the study to lines-of-sight with lower HI columns, because there is a strong D/H variability for logN(H) $\\leq$ 20 cm$^{-2}$ and no titanium column density data are currently available within this range.     Section 2 describes the observations, data analysis and TiII column determination. In section 3 we perform linear comparison fits between TiII and other species, using our data in combination with other published data. We perform corresponding correlative studies for available FeII data. Section 4 discusses the results.  \\begin{figure*} \\centering \\includegraphics[width=16cm,angle=0]{8820fig1.eps} \\caption{Normalised spectra around the TiII 3384 \\AA\\ (solid lines) and 3242 \\AA\\ (dashed lines) transitions for our target stars}. The abscissa is the heliocentric velocity in km/s. \\label{spectra} \\end{figure*} ",
        "conclusions": "We have observed nine southern hemisphere D/H target stars at high spectral resolution and derived the column-densities of interstellar ionised titanium towards those targets. A comparison between the D/H ratio and the titanium depletion reveals a clear correlation and reinforces the result of Prochaska et al. (2005) and the more recent result of Ellison et al. (2007). Our work extends the correlation over a broader range of interstellar HI columns, namely more than two orders of magnitude, showing that the trend extends down to log N(HI)=19.0 cm$^{-2}$. We also make use of the orthogonal distance regression method for our linear fits, in order to take into account errors in quantities both in abscissa and ordinate.  One star is far out of the correlation, $\\mu$Col. This target is already strongly discrepant in correlations shown by Linsky et al (2006) between D/H and iron and silicon abundances. We have shown evidence here for a non-interstellar origin of the broadest absorption feature and note that excluding this absorption would bring the star back to the main trend.  We have compared both the Ti and Fe depletions with the DI/HI ratio, but also with the DI/OI ratio, when available. For both metals there is a clear relationship also with DI/OI, which demonstrates that the metal abundance- D/H correlation is not primarily due to uncertainties in HI column measurements, and that D/O is also correlated with metal depletions in the ISM.  The ionized titanium correlation gradient is found to be similar to or slightly smaller than the corresponding gradient for ionized iron. Ellison et al (2007) have emphasized that given its high condensation temperature one would expect the titanium gradient to be the highest. We have argued here that ionization effects do affect FeII/HI and result in an overestimate of the gradient, while the titanium abundance is unaffected due to the nearly exact coincidence between TiII ad HI ionization potentials. Thus we believe that the slopes measured for both metals are not a concern with respect to the depletion hypothesis as the source of deuterium abundance variations. Further work is needed to correct for the ionization and obtain a more reliable comparison between the different metal-deuterium relationships.   There is some (although weak) evidence that the TiII/DI ratio decreases towards high column densities. This may signify that D is preferentially released into the ISM (when compared to titanium) in the dense regions that are predominant along distant lines-of-sight. It is important to investigate further this potential relationship, because, if confirmed, it would be a first indication of the different sources of interstellar D and metals. Recording more titanium data along shorter lines-of-sight is essential. TiII, that does not suffer ionisation biases, is an ideal tracer here.  It remains to explain why very high D/H ratios do not have corresponding very high D/Os, as already argued by  H{\\'e}brard et al (2005).  This is reinforced here by our findings (section 3.2 and Figure 9a) that the three high D/O values found by Oliveira et al (2006) may be overestimated. Additionally, Oliveira and H{\\'e}brard (2006) derive a high D/H for a very distant sightline while one would expect in this case a very strong depletion. Certainly some of the physical processes are not yet understood. On the other hand, since very discrepant points appear in the correlation with D/H, but none with D/O, the high D/H ratios may be a consequence of underestimated H columns.  Assuming proportionality between metals and D depletion, one can derive a mean O/H ratio from the ratio of the Ti/H vs. D/H and Ti/H vs. D/O slopes (or similarly for Fe). Using TiII one finds (Table 3) 6.05 10$^{-8}$/1.35 10$^{-4}$ = 448 ppm using our targets and the Prochaska et al. targets, and 6.46 10$^{-8}$/1.29 10$^{-4}$ = 501 ppm using all targets.  Using FeII one finds 1.49/4.96 10$^{3}$= 300 ppm. There must be some significance to this difference that remains to be elucidated. These two values are respectively higher and lower than the mean value of 347 ppm derived by Cartledge et al. (2004) within 800 parsecs. It is also interesting to compare with the range of solar photospheric O/H values recently proposed by Grevesse et al. (2007). The value derived here from the TiII data, 480-500 ppm, is within their quoted range of 407-513 ppm, while the value derived from FeII, 300 ppm, is significantly below. These calculations suggest that the slopes derived from the linear fits using TiII are meaningful.  Is there something to learn about the local ISM history from the D/H abundance pattern? The large-scale event that has given rise to the Gould belt of enhanced star formation is about 60 Myrs old (see e.g. Perrot \\& Grenier, 2003). Two main classes of scenarios have been proposed for the belt origin: the interaction between a giant external cloud and the disk on one hand (Comeron and Tora, 1994, Olano, 2001) and a strong explosive event on the other (Olano, 1982). Because the age of the belt is smaller than the mixing time of gases (a few hundreds Myrs, see e.g. simulations by De Avillez and Breitschwerdt, 2004), in the former case of cloud collision the belt may have left imprints in the form of incomplete mixing between disk and \"fresh\" gas.  In the latter case there must be dust evaporation at the shocked front regions of the belt but no gas mixing. A link between the Gould belt expansion and the D/H distribution has been suggested because highly variable D/H regions shown by Linsky et al (2006, Fig. 1) are located at distances (inferred from the H columns) that roughly correspond to the front of the expanding Gould belt of young stars and supernovae (Lallement, 2007). Whatever the origin of the belt, the subsequent and observed expansion produces star formation and supernovae at the periphery of the belt, i.e. evaporation of dust in the shocked and heated regions of its periphery, which could explain the observed variability of D/H at intermediated columns (10$^{19.5-20.5}$ cm$^{-2}$). The difference between the two scenarios for the belt origin is that in the cloud collision case some deuterium abundance variability due to gas mixing must add to abundance variation due to grain destruction, and there must be a larger amount of less astrated gas within the Gould belt as compared with outside the belt. There are no signs of astration variations that follow from the present correlative analysis, in agreement with the constancy of O/H up to large distances (Andr{\\'e} et al., 2003). If D and metal abundance variability has some link with the belt, then the absence of signs of differential astration at columns lower than 10$^{21}$ cm$^{-2}$ seems to favor the explosive event scenario.   Independantly of a potential effect of the Gould belt on the D/H distribution, our results provide a support for the ideas, originally suggested by Draine et al. (2004) and Linsky et al. (2006), that time-dependent deuterium depletion onto dust grains is responsible for the variability of the D/H ratio, and for the low values of D/H observed for many lines of sight beyond the Local Bubble. On the other hand, it remains to explain why very high D/H ratios do not correspondingly present very high D/O values. This is important for the precise determination of the present galactic D/H ratio."
    },
    "0801/0801.0985_arXiv.txt": {
        "abstract": "We review observations of extended regions of radio emission in clusters; these include diffuse emission in `relics', and the large central regions commonly referred to as `halos'. The spectral observations, as well as Faraday rotation measurements of background and cluster radio sources, provide the main evidence for large-scale intracluster magnetic fields and significant densities of relativistic electrons. Implications from these observations on acceleration mechanisms of these electrons are reviewed, including turbulent and shock acceleration, and also the origin of some of the electrons in collisions of relativistic protons by ambient protons in the (thermal) gas. Improved knowledge of non-thermal phenomena in clusters requires more extensive and detailed radio measurements; we briefly review prospects for future observations. ",
        "introduction": "\\label{Introduction}  In the last 10 years, the improved capabilities (sensitivity, spectral and spatial resolution) of multi-wavelength telescopes have allowed us to study in detail the formation and evolution of the largest gravitationally bound systems in the Universe, i.e. galaxy clusters. Following the hierarchical scenario of structure formation, massive clusters form through episodic mergers of smaller mass units (groups and poor clusters) and through the continuous accretion of field galaxies.  It has now been proven that major cluster mergers, with their huge release of gravitational binding energy ($\\sim 10^{64}$ ergs), deeply affect the physical properties of the different components of clusters, i.e. the temperature, metallicity and density distribution of the thermal intracluster medium (ICM) emitting in X-rays \\citep[e.g.][]{Buote02, Sauvageot05, Ferrari06a, Kapferer06, Markevitch07}, the global dynamics and spatial distribution of galaxies \\citep[e.g.][]{Girardi02, Ferrari03b, Ferrari05, Maurogordato07}, as well as their star-formation rate \\citep[e.g.][]{Gavazzi03, Poggianti04, Ferrari06b}. The typical signatures that allow to identify merging clusters from optical and X-ray observations are: a) substructures in the X-ray and optical surface densities \\citep[see][and references therein]{Buote02}, b) non-Gaussian radial velocity distributions of cluster members \\citep[see][and references therein]{Girardi02}, c) patchy ICM temperature, pressure, entropy and metallicity maps \\citep[e.g.][]{Finoguenov05, Kapferer06}, d) sharp X-ray surface brightness discontinuities, accompanied by jumps in gas and temperature \\citep[``cold fronts'', see, e.g.,][]{Markevitch00}, e) absent or disturbed cooling-cores \\citep[e.g.][]{Markevitch99}, f) larger core radii compared to (nearly) relaxed clusters \\citep[][]{Jones99}. There are also indications that recent merging events lead to a depletion of the nearest cluster neighbours \\citep[e.g.][]{Schuecker99}. Additionally, deep radio observations have revealed the presence of diffuse and extended ($\\sim$ 1 Mpc) radio sources in about 50 merging clusters. Their radio emission is not related to a particular cluster member, but rather to the presence of relativistic electrons (Lorentz factor $\\gamma \\gg$ 1000) and weak magnetic fields ($\\mu$G) in the intracluster space.  In this review, we focus on radio observations of this non-thermal component in galaxy clusters. We outline our current knowledge on the presence of non-thermal processes in the intracluster gas, and their physical connections with the thermodynamical evolution of large-scale structure. The relevance of the study of extended cluster radio emission for cosmology is pointed out. On smaller scales, there are only few indications of the possible presence of extended radio emission in galaxy groups. These systems host diffuse $\\sim$ 1 keV gas called intragroup medium (IGM) \\citep[see, e.g.,][]{Mulchaey96}. Radio and hard X-ray emission possibly related to a non-thermal component of the IGM has been recently pointed out by \\citet{Delain06} and \\citet{Nakazawa07} respectively. The existence of diffuse radio sources in galaxy groups has indeed to be tested with observations of higher sensitivity.  Radio observations of the emission from individual radio galaxies are not treated here. For a discussion on cluster radio galaxies see the reviews by \\citet{Feretti02a} and \\citet{Feretti07}. X-ray observations and simulations of the non-thermal component in clusters are reviewed by \\citet{Rephaeli08} - Chapter 5, and \\citet{Dolag08} - Chapter 15, this volume. The adopted cosmology is $\\Lambda$CDM (${\\rm H}_0$=70 km ${\\rm s}^{-1} {\\rm Mpc}^{-1}$, $\\Omega_{\\rm m} = 0.3$, $\\Omega_{\\Lambda} = 0.7$). ",
        "conclusions": ""
    },
    "0801/0801.2178_arXiv.txt": {
        "abstract": "{The Sun has recently been predicted to be an extended source of gamma-ray emission, produced by inverse-Compton (IC) scattering of cosmic-ray (CR) electrons on the solar radiation field. The emission was predicted to be extended and a confusing foreground for the diffuse extragalactic background even at large angular distances from the Sun. The solar disk is also expected to be a steady gamma-ray source. While these emissions are expected to be readily detectable in the future by GLAST, the situation for available EGRET data is more challenging. } {  The theory of gamma-ray emission from IC scattering on the solar radiation field by Galactic CR electrons is given in detail. This is used as the basis for detection and model verification using EGRET data.   } {  We present a detailed study of the solar emission using the EGRET database, accounting for the effect of the emission from 3C 279, the moon, and other sources, which interfere with the solar emission. The analysis was performed for 2 energy ranges, above 300 MeV and for 100-300 MeV, as well as for the combination to improve the detection statistics. The technique was tested on the moon signal, with our results consistent with previous work. } {Analyzing the EGRET database, we find evidence of emission from the solar disk and its halo.   The observations are compared with our model for the extended emission. The spectrum of the solar disk emission and the spectrum of the extended emission have been obtained. The spectrum of the moon is also given.} { The observed intensity distribution and the flux are consistent with the predicted model of IC gamma-rays from the halo around the Sun.} ",
        "introduction": "Solar quiet gamma-ray astronomy started playing a significant role in the early 90's thanks to the EGRET mission. \\citet{hudson} pointed out the importance of gamma-ray emission from the quiet Sun as an interesting possibility for highly sensitive instruments such as EGRET. \\citet{seckel} estimated that the flux of gamma rays produced by cosmic-ray interactions on the solar surface would be detectable by EGRET. This emission is the result of particle cascades initiated in the solar atmosphere by Galactic cosmic rays. Meanwhile \\citet{thompson97} detected in the EGRET data the gamma rays produced by cosmic ray interactions with the lunar surface. Although they expected similar interactions on the Sun, the EGRET data they analyzed did not show any excess of gamma rays consistent with the position of the Sun. They obtained an upper limit of the flux above 100 MeV of 2 $\\times$ 10$^{-7}$cm$^{-2}$s$^{-1}$ . \\citet{fairbairn}, analyzing the EGRET data of the solar occultation of 3C279 in 1991, found a photon flux of about 6 $\\times$ 10$^{-7}$ cm$^{-2}$s$^{-1}$ above 100 MeV in the direction of the occulted source, which they used to put limits on axion models; however they did not consider the extended solar emission.  Inverse Compton scattering of cosmic-ray electrons on the solar photon halo around the Sun has been estimated to be an important source of gamma-ray emission  \\citep{orlando,moskalenko} \\footnote{The same process from nearby stars has been studied in \\citet{orlando} and \\citet{orlandoc}}. In our previous work, we predicted this to be comparable to the diffuse background even at large angular distance from the Sun. The formalism has now been improved and more accurate estimates have been obtained. The anisotropic scattering formulation and the solar modulation have been implemented, in order to give a more precise model for the EGRET data. In this work we explain in detail our model of the extended emission produced via inverse Compton scattering of cosmic-ray electrons on the solar radiation field.  A first report of our detection with EGRET data was given in \\citet{orlandob}. Here we describe the full analysis including the spectrum of disk and extended components.  Our result is very promising for the forthcoming GLAST mission which certainly will be able to detect the emission even at larger angular separation from the Sun. Moreover, the extended emission from the Sun has to be taken into account since it can be strong enough to be a confusing background for Galactic and extragalactic diffuse emission studies. ",
        "conclusions": "In this paper the theory of gamma-ray emission from IC scattering of solar radiation field by Galactic CR electrons has been given in detail. Analyzing the EGRET database, we find evidence for emission from the Sun and its vicinity. For all models the expected values are in good agreement with the data, with rather big errors bars. This also means that, because of the limited sensitivity of EGRET, it is not possible to prove which model better describes the data. The spectrum of the moon has also been derived.   Since the inverse Compton emission from the heliosphere is extended, it contributes to the whole sky foreground and has to be taken into account for diffuse background studies. Moreover, since the emission depends on the electron spectrum and its modulation in the heliosphere, observations in different directions from the Sun can be used to determine the electron spectrum at different position, even very close to the Sun. To distinguish the modulation models (eg. with GLAST) an accuracy of $\\sim$10$\\%$ would be required. With the point source sensitivity of GLAST around 10$^{-7}$cm$^{-2}$s$^{-1}$ per day above 100 MeV, it will be possible to detect daily this emission, when the Sun is not close to the Galactic plane. In additional a precise model of the extended emission, such as that presented here, would be required in searches for exotic effects such as in \\citet{fairbairn}"
    },
    "0801/0801.2452_arXiv.txt": {
        "abstract": "Through photoionization modeling, constraints on the physical conditions of three $z$ $\\sim$ 1.7 single-cloud weak \\ion{Mg}{2} systems ($W_r$(2796) $\\leq$ 0.3\\AA) are derived.  Constraints are provided by high resolution $R$ $=$ 45,000, high signal-to-noise spectra of the three quasars HE~0141$-$3932, HE~0429$-$4091, and HE~2243$-$6031 which we have obtained from the European Southern Observatory (ESO) archive of Very Large Telescope (VLT), Ultraviolet and Visual Echelle Spectrograph (UVES). Results are as follows.  \\begin{description}  \\item[(1)] The single-cloud weak \\ion{Mg}{2} absorption in the three $z$ $\\sim$ 1.7 systems is produced by clouds with ionization parameters of $-$3.8 $<$ $\\log U$ \\lsim\\ $-$2.0 and sizes of 1 -- 100~pc.  \\item[(2)] In addition to the low-ionization phase \\ion{Mg}{2} clouds, all systems need an additional 1 -- 3 high-ionization phase \\ion{C}{4} clouds within 100 \\kms\\ of the \\ion{Mg}{2} component.  The ionization parameters of the \\ion{C}{4} phases range from $-$1.9 $<$ $\\log U$ $<$ $-$1.0, with sizes of tens of parsecs to kiloparsecs. In all cases at least one \\ion{C}{4} cloud is centered at the same velocity (within 5 \\kms\\ in our systems) as the \\ion{Mg}{2} clouds.  \\item[(3)] Two of the three single-cloud weak \\ion{Mg}{2} absorbers have near-solar or super-solar metallicities, if we assume a solar abundance pattern. Although such large metallicities have been found for $z$ $<$ 1 weak \\ion{Mg}{2} absorbers, these are the first high metallicities derived for such systems at higher redshifts.  These strong constraints were possible because of the specific shapes of the \\lya\\ profiles in these cases.  All single-cloud weak \\ion{Mg}{2} absorbers may have high metallicities, but in some cases the kinematic spread of the \\ion{C}{4} cloud contributions to \\lya\\ do not allow such a determination.  \\item[(4)] Two of the three weak \\ion{Mg}{2} systems also need additional low-metallicity, broad \\lya\\ absorption lines, offset in velocity from the metal-line absorption, in order to reproduce the full \\lya\\ profile.  \\item[(5)] Metallicity in single-cloud weak \\ion{Mg}{2} systems are more than an order of magnitude larger than those in Damped \\lya\\ systems at $z$ $\\sim$ 1.7.  In fact, there appears to be a gradual decrease in metallicity with increasing \\nhi, from these, the most metal-rich \\lya\\ forest clouds, to Lyman limit systems, to sub-DLAs, and finally to the DLAs. Weak \\ion{Mg}{2} absorbers could be near local pockets in which star-formation has occurred, but where there is little gas to dilute the metals that are dispersed into the region, resulting in their very high metallicities.  We speculate that DLAs may be subject to the opposite effect, where a large dilution of metals produced in the vicinity will occur, leading to a small metallicity.  \\end{description} ",
        "introduction": "\\label{sec:intro}  Many single-cloud weak \\ion{Mg}{2} absorbers ($W_r$(2796) $\\leq$ 0.3\\AA) at 0.4 $<$ $z$ $<$ 1.0 have been found to have solar or even super-solar metallicities \\citep{rig02,cha03}. This is quite puzzling in view of the fact that they tend to be found at distances of 40 -- 100~kpc from star-forming galaxies \\citep{chu05,mil06}, and not near detected sites of star formation.  These objects have a cross-section similar to the absorption cross section of luminous galaxies, thus they occupy a significant volume of the universe \\citep{rig02}. Because weak \\ion{Mg}{2} systems only contain a small mass, they do not trace the cosmic metal abundance, however, it is still important to understand local star-formation activities in the full range of environments over cosmic time.  \\citet{nar07} presented the results of a search for weak \\ion{Mg}{2} absorption in VLT/UVES quasar spectra over a cumulative redshift path of 77.3 at 0.4 $<$ $z$ $<$ 2.4.  A total of 116 \\ion{Mg}{2} absorbers with 0.02 $<$ $W_r(2796)$ $<$ 0.3~\\AA\\ were detected, with $\\sim$ 60\\% of these having just a single-cloud, narrow component in \\ion{Mg}{2}. It has been argued that a large fraction of the population of multiple-cloud weak \\ion{Mg}{2} absorbers (i.e., absorbers with multiple-clouds, which are resolved by high-resolutional ($R$ $\\sim$ 40,000) spectroscopies) have a different physical origin from the single-cloud weak \\ion{Mg}{2} absorbers \\citep{rig02,mas05,din05}. \\citet{nar07} found that the redshift path density, $dN/dz$, of weak \\ion{Mg}{2} absorbers increases from $z$ $=$ 2.4 to $z$ $=$ 1.2, where it peaks, and subsequently declines until the present.  The paper of \\citet{nar07} followed a weak \\ion{Mg}{2} survey of \\citet{lyn06}, which used a small subset of high quality VLT/UVES quasar spectra.  From that survey, \\citet{lyn07} selected the nine $z$ $>$ 1.4 weak \\ion{Mg}{2} absorbers and applied photo-ionization models.  Of these, four were single-cloud weak \\ion{Mg}{2} absorbers with coverage of the corresponding \\lya\\ transitions so that metallicities could be constrained.  In general, \\citet{lyn07} found that the physical conditions of the $z$ $>$ 1.4 weak \\ion{Mg}{2} absorbers were similar to those of the same class of absorbers at lower redshift, suggesting that the same production mechanisms are responsible over this time period.  However, in contrast to the situation at low redshift \\citep{rig02,cha03}, \\citet{lyn07} were only able to derive lower limits to the metallicity, ranging from 1/100$^{th}$ -- 1/10$^{th}$ of the solar value.  Single-cloud weak \\ion{Mg}{2} absorbers, in addition to the cloud that produced the \\ion{Mg}{2} absorption, have separate, lower density \\ion{C}{4} clouds that are spread over tens to hundreds of \\kms\\ \\citep{cha03,lyn07}.  The reason that only lower limits could be derived for the metallicity in these cases was because of the contributions to the \\lya\\ profile from the kinematically spread \\ion{C}{4} clouds.  On the other hand, \\citet{lyn07} also suggested that there could be a build-up of metals in the population of single-cloud weak \\ion{Mg}{2} absorbers from $z$ $\\sim$ 2 to $z$ $\\sim$ 0.4.  In view of the large spread of metallicity constraints at any redshift, a sample of only four $z$ $>$ 1.4 single-cloud weak \\ion{Mg}{2} absorbers was not sufficient to resolve the issue of whether metallicity evolution occurs in the population.  We therefore identified three additional $z$ $>$ 1.4 candidates for photo-ionization modeling from the larger survey of \\citet{nar07}. These systems, the $z$ $=$ 1.78169 absorber toward HE~0141$-$3932, the $z$ $=$ 1.68079 absorber toward HE~0429$-$4091, and the $z$ $=$ 1.75570 absorber toward HE~2243$-$6031, have high $S/N$-ratio VLT/UVES coverage of key constraining metal-line transitions, as well as \\lya. We apply Cloudy \\citep{fer98} photoionization models to constrain the ionization parameters and metallicities of the phases\\footnote{Phases are defined as regions of gas with similar metallicity, temperature, and volume density.} of gas that are required to fit the data.  Our focus is on the question of the metallicity evolution of the single-cloud weak \\ion{Mg}{2} absorbers, and on comparing the result to the metallicity evolution of other classes of quasar absorption line systems.  We will begin in \\S~\\ref{sec:data} with a summary of the data that we have used to constrain the properties of the three $z$ $\\sim$ 1.7 single-cloud weak \\ion{Mg}{2} absorbers.  Next, in \\S~\\ref{sec:photo}, we describe our procedure for photo-ionization modeling with Cloudy, which is based on comparing synthetic model profiles to the data. \\S~\\ref{sec:systems} presents the detailed results for the three absorbers.  Finally, in \\S~\\ref{sec:discussion}, we summarize those results, compare the metallicities of single-cloud weak \\ion{Mg}{2} absorbers to DLAs, to sub-DLAs, and to Lyman limit absorbers, and discuss the implications of this comparison.  Throughout this paper, we use a cosmology with $H_{0}$=72 \\kms Mpc$^{-1}$, $\\Omega_{m}$=0.3, and $\\Omega_{\\Lambda}$=0.7. ",
        "conclusions": "\\label{sec:discussion}  In this study, we applied photoionization models to constrain the physical properties of three single-cloud weak \\ion{Mg}{2} absorption systems at $z$ $\\sim$ 1.7.  Along with results presented in \\citet{lyn07}, these complete a sample of the seven $z$ $>$ 1.5 single-cloud weak \\ion{Mg}{2} absorbers found in the VLT archive \\citep{nar07} for which metallicity constraints could be derived because of simultaneous coverage of various metal lines and \\lya.  We start this final section by summarizing the model constraints derived for the three $z$ $\\sim$ 1.7 single-cloud weak \\ion{Mg}{2} absorbers studied in this paper (\\S~\\ref{sec:summary}). We then proceed to compare our results to the other four single-cloud weak \\ion{Mg}{2} systems at $z$ $\\sim$ 1.7 from \\citet{lyn07} in \\S~\\ref{sec:comp_weakMgII}. In this section, we also compare to lower redshift single-cloud weak \\ion{Mg}{2} systems from the literature to investigate evolutional trends. We address the high metallicity of these systems in the context of DLAs (whose metallicities are usually smaller than the solar value by a factor of 10 -- 30) in \\S~\\ref{sec:comp_DLA}, sub-DLAs in \\S~\\ref{sec:comp_subDLA}, and strong and multiple-cloud weak \\ion{Mg}{2} systems (the strong ones are usually associated with bright ($L$ $>$ 0.05$L^{*}$) galaxies within 40$h^{-1}$ kpc; e.g., \\citealt{ber91}) in \\S~\\ref{sec:comp_strongMgII}, respectively.  We finally discuss a possible correlation between metallicity and total hydrogen column density in \\S~\\ref{sec:metal}.   \\subsection{Summary of Results} \\label{sec:summary}  System~1 can be fit with a single low-ionization \\ion{Mg}{2} cloud and three higher ionization clouds (two of them offset in velocity from the \\ion{Mg}{2}) that give rise to \\ion{C}{4} absorption.  The low ionization cloud has $\\log Z$ $>$ $-$0.8 and two of the high ionization clouds must have metallicities at least this high as well. Two offset low-metallicity clouds, producing only \\lya\\ absorption appear to be clustered with this system.  System~2 is also fit with one \\ion{Mg}{2} and three \\ion{C}{4} clouds, though all four clouds have similar, relatively high ionization parameters ($\\log U$ $\\sim$ $-$1.9).  The metallicity of the \\ion{Mg}{2} cannot be well-constrained directly from comparison to the \\lya\\ because the kinematical spread of the \\ion{C}{4} lines give rise to a large \\lya\\ equivalent width, however near solar metallicities are possible, and even favored for a solar abundance pattern.  Also, the blueward \\ion{C}{4} cloud will over-produce absorption in \\lya\\ unless its metallicity is $\\log Z$ $>$ $-$1.1.  A separate, offset low-metallicity component was again needed to fit the red portion of the \\lya\\ profile.  System~3 has the best metallicity constraint among the seven single-cloud weak \\ion{Mg}{2} absorbers at $z$ $>$ 1.5 because of the absence of offset \\ion{C}{4} components, thereby keeping the \\lya\\ profile very narrow.  Two phases are still required, with clouds of different ionization parameters centered at the same velocity that produce both the narrower low-ionization transitions, and the broader high-ionization transitions.  Both the low- and the high-ionization phases are constrained to have solar or super-solar metallicities.  The most striking thing about these model results, is that in all three $z$ $\\sim$ 1.7 systems, the metallicity of at least one phase of gas is constrained to be greater than one tenth the solar value.  The three systems all have at least two phases of gas: the low ionization phase that arises in a layer of gas $\\sim$1 -- 100~pc thick with a density of 0.001 -- 0.1~\\cmmm, and the high ionization phase that comes from a larger region ($\\sim$ 0.1 -- 10~kpc) with a lower density.  For these systems, even metallicities of the larger high-ionization phase regions range from one tenth solar up to solar. We will return to a discussion of the surprising issue of how regions with \\lognhi\\ $<$ 15 (i.e., over a million times less than the threshold for star formation) can be enriched in metals to the solar value.   \\subsection{Comparison to Other Single-Cloud Weak \\ion{Mg}{2} Systems} \\label{sec:comp_weakMgII}  The properties of the three $z$ $\\sim$ 1.7 single-cloud weak \\ion{Mg}{2} systems, derived from the observed profiles, are compared to those of 23 other single-cloud weak \\ion{Mg}{2} absorbers in Figure~\\ref{fig5}, both at similar and at lower redshifts. Figure~\\ref{fig5} shows that there is no significant evolution in the observed \\ion{Mg}{2}~$\\lambda$2796 profiles.  Figure~\\ref{fig6} presents quantities derived from photoionization models, for our three absorbers and for others taken from the literature, using the same methods as we have used here \\citep{rig02,cha03,din05,lyn07}.  The specific constraints that are displayed in Figures~\\ref{fig5} and \\ref{fig6} are also listed in Table~\\ref{tab7}, with the relevant references.  \\citet{lyn07} derived similar constraints for the phase structure of four different single-cloud weak \\ion{Mg}{2} absorbers at $z$ $\\sim$ 1.7.  They also placed constraints on the metallicities of those absorbers, of $\\log Z$ $>$ $-$1.5, $>$ $-$1.5, $>$ $-$1.0, and $>$ $-$2.0 for the four different low-ionization phases.  They note that the metallicities could be significantly higher than these lower limits, however, these absorbers (by chance) tended to have larger contributions to the \\lya\\ absorption from offset \\ion{C}{4} clouds and from \\lya-only clouds.  Results from modeling lower redshift (0.4 $<$ $z$ $<$ 1.4) single-cloud weak \\ion{Mg}{2} systems also typically yield a two-phase structure, with similar densities.  Metallicity constraints are again often limited because of the difficulty in separating the contributions of the two phases.  However, of the 11 cases of single-cloud weak \\ion{Mg}{2} absorbers for which some metallicity constraint could be obtained, 2 cases require the metallicity of the low-ionization phase to be be solar or even super-solar \\citep{cha03}. Furthermore, a total of 7/11 of the low redshift cases required a metallicity of at least one tenth of the solar value.  We note that the metallicities are often quite likely to be higher than these strict lower limits because they are derived assuming that none of the \\lya\\ absorption comes from the high-ionization clouds.  We conclude that some fraction of weak \\ion{Mg}{2} systems have been demonstrated to have solar or super-solar metallicity, at low redshift (2 systems at \\zabs\\ = 0.8181 and 0.9056; \\citealt{cha03}) and even at $z$ $\\sim$ 1.7 (1 system at \\zabs\\ = 1.7557; this paper), which corresponds to an age of the Universe of 3.7~Gyr.  Values constrained to be greater than 1/10th the solar value are even more common. Furthermore, it is likely that in many cases for which lower limits on metallicity have been derived the actual value is much higher.  The metallicity constraint that we have derived for the \\zabs\\ = 1.7557 absorber in this paper is high because this system is a rare case that does not have offset \\ion{C}{4} absorption that contributes to the \\lya\\ absorption.  In other ways there is no reason to think it is unusual.  The same applied for the low redshift weak \\ion{Mg}{2} systems with the highest metallicity constraints \\citep{cha03}.  We therefore know that single-cloud weak \\ion{Mg}{2} absorbers at both low and high redshift definitely have metallicities of at least 1/10th solar, but that some and probably many have solar or even supersolar metallicities.  Thus even before $z=1.5$ star formation must have polluted certain environments with metal-rich gas.  Since \\citet{lyn07} could not constrain any $z$ $>$ 1.5 single-cloud weak \\ion{Mg}{2} absorbers to have close to solar metallicity, and since several lower redshift solar metallicity cases were known, they tentatively suggested that there might be a gradual build-up of metals in the single-cloud weak \\ion{Mg}{2} absorber population.  In this paper, we have shown there is at least one strong counter-example, i.e. of $z$ $>$ 1.5 solar metallicity cases, and therefore there appears to be no metallicity evolution in the population.  This is not to say that the population has a narrow range of physical properties: there is a large spread of cloud densities and metallicities at all redshifts (see Figure~\\ref{fig6}).  However, it does suggest that common processes are at work to create the metals, and that the metals are mixed into fairly similar surrounding environments.  The existence of high metallicity (greater than solar) compact intergalactic clouds ($\\sim 100$~pc) has also been demonstrated by \\citet{sch07}, who derived robust lower limits on metallicities of \\ion{C}{4} absorbers using similar photoionization modeling techniques.  \\citet{sch07} find that the number density of this population of high metallicity \\ion{C}{4} absorbers is of the same order as that of single-cloud weak \\ion{Mg}{2} absorbers, and suggests that the former might evolve from the latter as material expands. More observations allowing direct connections between these populations will be of great interest.  In fact, two of our three $z$ $>$ 1.5 weak \\ion{Mg}{2} systems have \\ion{C}{4} phases with metallicities comparable to those derived by \\citet{sch07} for the high metallicity \\ion{C}{4} cloud population.  \\subsection{Comparison to Damped \\lya\\ Systems} \\label{sec:comp_DLA}  Damped \\lya\\ systems (DLAs), whose neutral hydrogen column densities are greater than \\lognhi\\ $=$ 20.3, dominate the neutral gas in the Universe at high redshift \\citep[e.g.,][]{lan95,sto00,per03}. By identifying metal absorption lines that correspond to high redshift DLAs, it is possible to constrain the global metal abundances and trace the evolution of metallicity in neutral gas from $z$ $=$ 5 to 0 \\citep[e.g.,][]{pet94,pro00}.  The metallicity evolution, measured from DLAs, however, has two big problems: (i) the global mean metallicity is less than 1/6th of the solar abundance at $z$ $>$ 1 \\citep[e.g.,][]{pet99,pro00}, and is still less than the solar by a factor of $>$ 4 if the metallicity redshift relationship is extrapolated down to $z$ $\\sim$ 0 using a least-squares linear fit \\citep{pro03,kul05,per06b} \\footnote{At $z$ $<$ 1, there are only a few cases of metallicity measurements in DLAs using high quality spectra, and in these cases the metallicity is evaluated to be $\\sim$ 1/10th of the solar value \\citep{var00,pet00}}, and (ii) the metallicity is also much smaller than the theoretically expected value from the star forming activity at higher redshift \\citep[e.g.,][]{mad98,pet99,wol03}.  This discrepancy (i.e., ``missing-metals problem'') remains an unresolved problem.  There are at least two possible classes of ideas proposed for the origin of this discrepancy: (i) a large amount of dust in metal-rich DLAs obscures background quasars, which prevent us from detecting DLAs with higher metallicities \\citep[e.g.,][]{boi98,fal93,vla05}, and (ii) the DLA region is sampling a lower metallicity than is found in the star-forming parts of the galaxy \\citep{ell05a,wol03}.  The dust obscuration scenario is supported by the small abundances of depleted elements such as Cr and Fe, relative to undepleted elements like Zn, in DLA systems \\citep[e.g.,][]{pet97,kha04}. \\citet{vla05} estimated that $\\sim$30~\\%\\ to 50\\%\\ DLAs are missed as a consequence of obscuration.  In absorbers expected to have a lower dust abundance, sub-DLA systems (super-LLSs) with \\lognhi\\ $=$ 19.0 -- 20.3, a higher metallicity has been measured \\citep[e.g.,][]{per06a,pro06,pet00,kul07}. Although \\citet{ell01} did not see any significant differences in DLAs toward radio-selected quasars, compared to those toward optically selected quasars at $z$ $>$ 2, the dust-depletion effect is expected to be more important at lower redshift.  Alternatively, \\citet{ell05a} and \\citet{che05} suggest that there could be an metallicity gradient as a function of a distance from the galactic center, with DLAs produced at a larger impact parameter than the stellar emission.  DLAs that are produced in the host galaxies of gamma-ray bursts tend to have higher metallicities \\citep[e.g.,][]{djo04}, which could be more representative of typical star-forming regions.  Also, small regions with DLA column densities could be ejected from galaxies by superwinds or as a part of AGN outflows \\citep{mac99,ham97,gab06,bou06}.  Figure~\\ref{fig7} summarizes the metallicities derived for DLAs at $z$ $=$ 0.4 to 4.8 from \\citet{pro03} and \\citet{kul05}, and compares to the values for single-cloud weak \\ion{Mg}{2} absorbers.  It is striking that weak \\ion{Mg}{2} absorbers, despite having drastically smaller \\nhi\\ values, have considerably higher metallicities than the DLAs, both at low redshift and at $z$ $\\sim$ 1.7.  Because they are not near regions with current star formation, the dust-to-gas ratio in single-cloud weak \\ion{Mg}{2} absorbers is expected to be very small.  Even for DLAs, it has been demonstrated that there is no bias against observing high metallicities due to dust obscuration of their background quasars \\cite{ell01,ell05b}.  Even if the dust-to-gas ratio is similar to those in DLAs, their total dust column density should be much smaller than those in DLAs, because of their lower gas column densities.  Thus, they would not be subject to the metallicity bias that was more likely to have affected DLAs.  Thus the weak \\ion{Mg}{2} absorbers provide an opportunity to see some high metallicity regions of the universe at high redshift, though they are not the same type of regions that would produce high metallicity DLA absorption.  Relating to the idea that DLAs sample low metallicity parts of galaxies, the weak \\ion{Mg}{2} absorbers, despite their low column densities, must somehow sample higher metallicity regions which are not closely related to luminous galaxies.  This provides clear evidence for metallicity inhomogeneities, which because of the size constraints for single-cloud weak \\ion{Mg}{2} absorbers, are on very small scales (parsecs to hundreds of parsecs).   \\subsection{Comparison to sub-DLA Absorbers} \\label{sec:comp_subDLA}  Sub-DLA absorbers, with \\lognhi\\ $=$ 19.0 -- 20.3, have been found to have systematically higher metallicities than DLAs with the \\nhi-weighted value larger by 0.5 -- 0.8 dex, at 0.6 $<$ $z$ $<$ 3.2 \\citep{kul07}.  Several sub-DLA systems (super-LLSs) at 0.5 $<$ $z$ $<$ 2.0 have been found with metallicities that are solar or super-solar \\citep[e.g.,][]{per06a,pro06}.  Even small numbers of such high metallicity systems contribute significantly to the total metal mass density at high redshift \\citep{pro06,kul07}.  The \\citet{kul07} mean values of metallicity for sub-DLAs, determined from [Zn/H], are reproduced in our Figure~\\ref{fig7} from their Table~1.  Although there is a range of metallicities at each redshift, there is a clear increase in the metallicity of sub-DLAs from $\\langle z \\rangle$ $=$ 1.5 to $\\langle z \\rangle$ $=$ 0.9.  This is consistent with the global star formation history in the universe over this period, and the metallicity values are much more in line with expectations from chemical evolution models than those for DLAs \\citep{kul07,som01}.  Some single-cloud weak \\ion{Mg}{2} absorbers, with \\lognhi\\ $\\sim$ 15 -- 16, also have solar or super-solar metallicities, both at $\\langle z \\rangle$ $=$ 0.9 and at $\\langle z \\rangle$ $=$ 1.5.  There is considerable overlap between the metallicity constraints for sub-DLAs and single-cloud weak \\ion{Mg}{2} absorbers at low and intermediate redshifts.  However, it is important to note that many of the weak \\ion{Mg}{2} metallicities are only lower limits, and it is quite likely that the values are higher at least in some cases.  The reason that we cannot derive higher metallicities for those systems could simply be the separate \\ion{C}{4} clouds that happen to overlap in velocity. It is also important to note that for $z$ $\\sim$ $1.7$, the redshift of the absorbers we have studied in this paper, there are not many sub-DLA systems with super-solar metallicity, but all single-cloud weak \\ion{Mg}{2} absorbers potentially could have solar- or super-solar metallicity because we can place only lower limits on metallicity. Although two super-solar metallicity sub-DLAs were investigated by \\citet{pro06}, they were selected for study from a much larger sample because of their extremely strong \\ion{Zn}{2} absorption. However, much larger data samples of sub-DLAs and weak \\ion{Mg}{2} systems will be required before concluding whether there are any differences of metallicity between these two categories.  Another apparent difference between the metallicities of single-cloud weak \\ion{Mg}{2} absorbers and sub-DLAs, is that there is no clear increase in the metallicity of the single-cloud weak \\ion{Mg}{2} absorbers over the range 0.4 $<$ $z$ $<$ 1.7.  This type of absorption is apparently produced by the same types of regions present over this redshift range, though they are known to be less common at low redshifts and at $z$ $>$ 2 \\citep{nar05,nar07,lyn06,lyn07}.  The sample size of weak \\ion{Mg}{2} systems is now much smaller than those for DLAs and Sub-DLAs, and only a few cases have been confirmed to have super-solar metallicities. It is quite important to increase the sample size of weak \\ion{Mg}{2} absorbers, particularly to include other systems with only one \\ion{C}{4} component such that a strong constraint on metallicity can be derived.  At $z < 1.5$ this will involve high resolution observations with {\\it HST}/COS or {\\it HST}/STIS of quasars for which high-resolution optical spectra are also available.   \\subsection{Comparison to Strong \\ion{Mg}{2} and Multiple-Cloud Weak \\ion{Mg}{2} Absorbers} \\label{sec:comp_strongMgII}  Intermediate in \\lognhi\\ between single-cloud weak \\ion{Mg}{2} absorption and sub-DLA's are Lyman limit and partial Lyman limit systems, with \\lognhi\\ $=$ 16 -- 19.  These populations loosely correspond to the bulk of the strong \\ion{Mg}{2} absorber population, with $W_r$(2796) $>$ 0.3~{\\AA}, and to multiple-cloud weak \\ion{Mg}{2} absorbers.  Fewer of these systems have been studied in detail to date, but it is still possible to derive good constraints on metallicities in many cases, using photoionization modeling techniques the same as those employed here.  The ten cases of strong and multiple-cloud weak \\ion{Mg}{2} absorbers at 0.7 $<$ $z$ $<$ 1.9 \\citep{mas05,zon04,lyn07,pro99,din03}, shown on Figure~\\ref{fig7}, have metallicities ranging from $-$1.5 $<$ $\\log Z$ $<$ 0.6.  These are consistent with the ranges for both sub-DLAs and single-cloud weak \\ion{Mg}{2} absorbers, but tend to be higher than those for DLAs. Without larger samples, we cannot distinguish differences or evolutionary trends, but again we note that many of these are measurements of the metallicity, while for the single-cloud weak \\ion{Mg}{2} absorbers the low metallicity values are all lower limits.   \\subsection{Metallicity vs. \\nhi} \\label{sec:metal}  In summary, Figure~\\ref{fig7} shows that it is the lowest \\nhi\\ systems, that appear to have the highest metallicities, both at $z$ $<$ 1 and at $z$ $\\sim$ 1.7.  The data are even consistent with a gradual increase in metallicity with decreasing \\nhi, though there is a huge spread at any given value. Such a trend has already been pointed out using compiled samples of DLAs and sub-DLAs \\citep{boi98,ake05,mei06,kha07}, extending down to sub-DLAs with \\lognhi\\ $\\sim$ 19 \\citep{per03}. \\citet{yor06} also suggested that a similar trend continued to strong \\ion{Mg}{2} systems with smaller hydrogen column densities.  The individual systems plotted on Figure~\\ref{fig7}, however, are only the systems with detected \\ion{Mg}{2} absorption.  In particular, for single-cloud weak \\ion{Mg}{2} absorbers we are sure to be missing objects with metallicities significantly less than those of the systems that we do detect.  In fact, those lower metallicity objects are part of the much larger \\lya\\ forest population. Figure~\\ref{fig7} also shows the typical metallicities of \\lya\\ forest clouds, both those with \\lognhi\\ $>$ 14.5, which have $\\log Z$ \\lsim\\ $-$2.0 \\citep{son96,cow95,tyt95}, and those with lower column densities, which have a mean metallicity of $\\log Z $ \\lsim\\ $-$3.2 \\citep{cow98,lu98}.  These much lower values for the \\lya\\ forest overlap with the DLA metallicities at $z$ $>$2.5, but are smaller than those for low redshift DLAs.  The fact that the high metallicity population of single-cloud weak \\ion{Mg}{2} absorbers represents only a small fraction of all \\lya\\ forest clouds does not diminish their significance.  The single-cloud weak \\ion{Mg}{2} absorbers at $z$ $\\sim$ 1 would account for 25 -- 100\\%\\ of the \\lya\\ forest clouds with 15.8 $<$ \\lognhi\\ $<$ 16.8 \\citep{rig02}.  More importantly, the cross-section of single-cloud weak \\ion{Mg}{2} absorbers in the plane of the sky is comparable to that of galaxies at $z$ $\\sim$ 1, considering the regions of galaxies $\\lsim$30 -- 40~kpc that produce Lyman limit absorption.  Thus there are significant regions of the universe covered by these mysterious near-solar or super-solar metallicity objects with small \\lognhi.  As described above, one explanation proposed for the larger metallicities for sub-DLAs than for DLAs is the dust obscuration of quasars that have the highest metallicity DLAs in their foregrounds. Although there may be some debate about whether this is a plausible explanation of that trend, it is not likely to explain the continued increase in metallicity toward the lowest \\lognhi\\ weak \\ion{Mg}{2} absorbers.  There is not likely to be a large selection effect due to dust obscuration for even the sub-DLA systems.  It seems more plausible that different types of absorbers are sampling different types of regions in and around galaxies, which can have large variations in metallicities even on small scales.  Different types of galaxies, e.g. star-forming versus quiescent, will have different fractions of area covered by regions of the different types. By comparing to the mass-metallicity relation seen in star-forming galaxies at local universe \\citep{tre04} and higher redshift \\citep{sav05,erb06}, \\citet{yor06} and \\cite{kha07} proposed that the DLAs are associated with low-mass ($<$ $10^9$ $M_{\\odot}$) galaxies, while sub-DLAs and weaker LLSs are probably the systems that arise in massive spiral/elliptical galaxies.  The question remains: how do the single-cloud weak \\ion{Mg}{2} absorbers develop such high metallicities even at redshifts as high as $z$ $\\sim$ 1.7?  We know that the lines of sight that pass through these objects have \\nhi\\ five or six orders of magnitude below the star formation threshold.  We also know that most single-cloud weak \\ion{Mg}{2} absorbers are not located very close to luminous galaxies, but that they tend to be in the vicinity \\citep{chu05,mil06}. Outflowing gas from dwarf galaxies in the local universe \\citep[e.g.,][]{sto04} or from Lyman break galaxies at higher redshift \\citep[e.g.,][]{ade05} could partially contribute to the high metallicities in single-cloud weak \\ion{Mg}{2} systems.  \\citet{sch07} have argued that high metallicity intergalactic \\ion{C}{4} clouds may be ejected as small clumps in superwinds from star-forming galaxies. However, at least some of the very high metallicity single-cloud weak \\ion{Mg}{2} absorbers show signs of \\ion{Fe}{2} line detections (e.g. our System~1 and several systems in \\citealt{rig02}), although sometimes it is detected with only a few $\\sigma$ level. For those systems, it has been stated that they are not $\\alpha$-enhanced, and that \"in situ\" star formation, including the less energetic Type~Ia SNe that increase the iron to magnesium ratio, is responsible for enriching the gas \\citep{rig02}. Here, readers should be aware that it is not the gas along the weak \\ion{Mg}{2} absorber line of sight that has formed stars, since it has much too low a column density.  It must instead be a nearby molecular cloud region that was actively star forming more than 1 billion years before the time from which we observe the absorber.  That region would itself produce DLA absorption if a line of sight passed through it. It is known that many DLAs at $z$ $<$ 1 are associated only with dwarf or low surface brightness galaxies \\citep{rao03}, and that in at least one case no galaxy is detected in deep narrow-band H$\\alpha$ images \\citep{bou01}.  So then why do DLAs not have high metallicities as well?  And how do we explain the general trend of increasing metallicity with decreasing \\nhi?  The reason could be simple.  Large quantities of metals can be produced in star forming regions in a variety of environments, inside and outside of galaxies.  However, the metals will be dispersed into widely different types of regions, depending upon the environment. Within giant galaxies molecular cloud regions are typically surrounded by a significant volume of high density, high column density gas.  So the metallicity is diminished by a large dilution effect as the ejected metals are spread through the galaxy region.  In an intergalactic star-forming region, which is metal-polluted by {\\it in-situ} star formation or has been ejected from a nearby dwarf galaxy, the clouds may not have a significant volume of high column density gas around it, thus the dilution could be minimal.  This would especially apply in a shallow potential well, where much of the surrounding gas is ejected by the early Type~II SNe, and the remnant low density gas remaining is greatly enhanced by the metals produced by later Type~Ia's which encounter a relatively sparse surrounding medium.  The same type of dilution effect could lead to DLAs having lower metallicity than sub-DLAs.  \\citet{pro06} suggests that the super-solar metallicity sub-DLAs that they observed at $z$ $\\sim$ 1.7, extreme members of the sub-DLA population, are related to regions affected by the feedback from actively star-forming regions.  The metals would be spread to larger lower density sub-DLA regions by energetic processes.  Obviously, the metallicity of a region is determined by a combination of the amount of star formation activity in its vicinity and the overall density of its gas.  In the case of the single-cloud weak \\ion{Mg}{2} absorbers, we conclude that the density effect dominates, leading to unexpected patches of high metallicity outside of galaxies. In the contrasting case of a DLA, there is just too much gas around it to dilute the metals coming through it, and a low metallicity often results."
    },
    "0801/0801.2387_arXiv.txt": {
        "abstract": "We present the results of a survey for stellar and substellar companions to 82 young stars in the nearby OB association Upper Scorpius. This survey used nonredundant aperture-mask interferometry to achieve typical contrast limits of $\\Delta$$K\\sim$5-6 at the diffraction limit, revealing 12 new binary companions that lay below the detection limits of traditional high-resolution imaging; we also summarize a complementary snapshot imaging survey that discovered 7 directly resolved companions. The overall frequency of binary companions ($\\sim$35$^{+5}_{-4}$\\% at separations of 6-435 AU) appears to be equivalent to field stars of similar mass, but companions could be more common among lower-mass stars than for the field. The companion mass function has statistically significant differences compared to several suggested mass functions for the field, and we suggest an alternate log-normal parameterization of the mass-function. Our survey limits encompass the entire brown dwarf mass range, but we only detected a single companion that might be a brown dwarf; this deficit resembles the so-called ``brown dwarf desert'' that has been observed by radial-velocity planet searches. Finally, our survey's deep detection limits extend into the top of the planetary mass function, reaching 8-12 $M_{Jup}$ for half of our sample. We have not identified any planetary companions at high confidence ($\\ga$99.5\\%), but we have identified four candidate companions at lower confidence ($\\ga$97.5\\%) that merit additional followup to confirm or disprove their existence. ",
        "introduction": "The detection and characterization of low-mass companions has become one of the highest priorities of the astronomical community. Radial velocity surveys have discovered over 200 extrasolar planetary companions over the past decade, and both RV surveys and coronagraphic imaging surveys have discovered an abundance of stellar-mass companions (e.g. Marcy \\& Butler 2000; McCarthy \\& Zuckerman 2004; Metchev 2005; Johnson et al. 2006; Naef et al. 2007). However, very few brown dwarf companions have been identified, an unexpected result given that the observational signatures of more massive companions are far larger than those of planetary-mass companions and that free-floating brown dwarfs are very common (Kirkpatrick et al. 2000; Luhman et al. 2003; Chiu et al. 2006; Slesnick et al. 2006a, 2006b). This dearth of companions between the stellar and planetary mass regimes is popularly known as the ``brown dwarf desert''. The existence and extent of the brown dwarf desert can provide key constraints on star and planet formation since it represents the extreme mass limit of both processes.  If the stellar-mass binary companions of solar-mass stars are drawn from the Initial Mass Function (IMF; e.g. Kroupa 1995) or formed via some other process that preferentially forms low-mass companions (e.g. Duquennoy \\& Mayor 1991, hereafter DM91), then brown dwarf companions should be common unless another process inhibits their formation or dynamically strips them. However, if stellar companions are formed via the fragmentation of a protostellar core, then there are no a priori expectations that brown dwarfs should form. Indeed, even if fragmentation can form an extremely inequal-mass pair, the long collapse timescale for low-mass objects might lead to their preferential photoevaporation by the higher-mass, more luminous companion.  It is also unclear whether brown dwarfs could form via planetary formation processes. Radial velocity surveys suggest that the giant planetary mass function is well-fit by a power law, $dN/dm\\propto$$M^{-1.05}$, for masses of $\\sim$1-10 $M_{Jup}$ (Marcy et al. 2005). If this power law extends to higher masses, there should be as many ``planetary'' companions with masses of 10-25 $M_{Jup}$ as with masses of 4-10 $M_{Jup}$ or 1.6-4.0 $M_{Jup}$. An absence of these companions suggests either that the function is not a power law or that the power law is truncated by some limit. For example, submillimeter disk surveys suggest that protoplanetary disks have a mean mass of $\\sim$5 $M_{Jup}$ by the age of 1-2 Myr (Andrews \\& Williams 2005), with a small fraction ($\\sim$5\\%) having masses of $\\sim$30-100 $M_{Jup}$. Unless massive planets are formed very early or efficiently accrete the entire disk mass, this could impose an upper cutoff on the distribution of planetary masses.  The brown dwarf desert has been studied mostly at very small or very large separations. The radial velocity exoplanet surveys that have proven so successful over the past decade should have detected any brown dwarfs within their outer separation limit ($\\sim$3-5 AU), and they have set very low upper limits on the frequency of close brown dwarf companions to solar mass stars ($<$1\\%; Marcy \\& Butler 2000; Grether \\& Lineweaver 2006). Similarly, high-resolution coronagraphic imaging surveys have demonstrated sufficient sensitivity to identify brown dwarf companions at typical separation limits of $>$50 AU (e.g. Gizis et al. 2001; Neuh\\\"auser et al. 2003; McCarthy \\& Zuckerman 2004; Neuh\\\"auser \\& Guenther 2004; Metchev 2005). They have measured frequencies which are low, but somewhat inconsistent (and perhaps not anomalously low; 1$\\pm$1\\% by McCarthy \\& Zuckerman, compared to 6.8$^{+8.3}_{-4.9}$\\% by Metchev and 18$\\pm$14\\% by Gizis et al.). A survey for wide companions to high-mass (2-8 $M_{\\sun}$) stars in Upper Sco by Kouwenhoven et al. (2007) found a relatively low frequency for brown dwarf companions, 0.5$\\pm$0.5\\% at separations of 130-520 AU. Finally, there have been an intriguing sample of candidate planetary-mass companions identified at large separations (e.g. Chauvin et al. 2004; Neuh\\\"auser et al. 2005), but both their mass and formation mechanism are still uncertain and their frequency is still unconstrained (e.g. Masciadri et al. 2005; Kraus et al. 2006; Biller et al. 2007; Ahmic et al. 2007).  However, these surveys don't study the actual separation range where most giant planets and binary companions are expected to form. Most giant planets at small orbital radii ($\\la$5 AU) are thought to have migrated inward, so their mass distribution may not match that of their more distant brethren. The binary formation process may also be different for small separations ($\\la$10 AU), with H$_2$ dissociation softening the equation of state and leading to enhanced fragmentation over that expected for larger length scales (Whitworth \\& Stamatellos 2006). Similarly, giant planets aren't expected to form at very large radii ($\\ga$30 AU) since the formation timescale is too long, and the frequency of wide binary companions may differ significantly from those of closer binaries (e.g. Kraus \\& Hillenbrand 2007a) since the fragmentation occurs on a length scale that is several orders of magnitude larger.  Ideally, the desert should be studied at the separation range where giant planets and most binaries are thought to form ($\\sim$5-30 AU; Lissauer \\& Stephenson 2007; DM91), but this has been impossible using existing techniques. For example, theoretical models (Chabrier et al. 2000) suggest that a 50 $M_{Jup}$ brown dwarf located 15 AU (1-2$\\arcsec$) from a nearby field star will have a contrast ratio of $\\Delta$$K\\sim$10-15 magnitudes at a separation of only $\\sim$1\\arcsec. The contrast problem could have been addressed by observing young stars since their substellar companions would be intrinsically more luminous ($\\Delta$$K<$5 mag), but most young stars are further away, so the separations are even smaller (0.1-0.2$\\arcsec$; $\\sim$$\\lambda$$/D$). Sensitivities near the diffraction limit have traditionally been far too shallow to detect such companions. However, new advances in high-resolution imaging techniques are now opening up this critical regime; our survey will use one such technique, non-redundant aperture mask interferometry.  The technique of non-redundant aperture masking has been well-established as a means of achieving the full diffraction limit of a single telescope (e.g. Nakajima et al. 1989; Tuthill et al. 2000). The reason for the technique's success over direct imaging is that the calibration is independent of structure of the wavefront over scales larger than a single sub-aperture, but it still preserves the angular resolution of the full aperture. This technique, when applied to seeing-limited observations, requires observations to be taken in a speckle mode with sub-apertures of diameter smaller than the atmospheric coherence length, limiting the technique to objects brighter than about $m_H=5$. The use of adaptive optics allows for longer integration times and larger sub-apertures, extending the technique to much fainter targets.  Published detections have been able to recover astrometrically discovered binary systems with contrast ratios of 3:1 at 0.6 $\\lambda$$/D$ and 100:1 at $\\lambda$$/D$ (Pravdo et al. 2006; Lloyd et al. 2006; Ireland et al. 2008) using total  observation times of $\\sim$10 minutes. The inner limit of companion detectability at high contrast is $\\lambda/2B_L$, where $B_L$ is the longest baseline in the mask (typically 80-95\\% of the aperture diameter). Typical closure phase errors are such that aperture masking can unveil high contrast companions at separations ~5 times closer than direct imaging in both $H$ and $K$ bands.  In this paper, we describe an aperture-mask interferometry and direct imaging survey to detect stellar and substellar companions to young stars in the nearby OB association Upper Scorpius. This survey directly studies the age and separation range corresponding to the peak of planet formation, offering the first glimpse of the brown dwarf desert in this critical range of parameter space. In Section 2, we describe our survey sample, and in Section 3, we describe the observations and data analysis techniques. In Section 4, we summarize the results of our survey. In Section 5, we combine these results with previous binary surveys to place constraints on the stellar binary frequency, mass function, and separation distribution, and in Section 6, we consider constraints on the corresponding parameters for the planetary population. Finally, in Section 7, we discuss the implications of our survey for the extent and aridity of the brown dwarf desert. ",
        "conclusions": "We present the results of a survey for stellar and substellar companions to 82 young stars in the nearby OB association Upper Scorpius. This survey used nonredundant aperture-mask interferometry to achieve typical contrast limits of $\\Delta$$K\\sim$5-6 at the diffraction limit, revealing 12 new companions that lay below the detection limits of traditional high-resolution imaging; we also summarize a complementary snapshot imaging survey that discovered 7 directly resolved companions. The overall frequency of binary companions ($\\sim$33$^{+5}_{-4}$\\% at separations of 6-435 AU, including companions reported in the literature) appears to be similar to field stars of similar mass, but the companion mass function appears to be more biased toward equal-mass companions than the equivalent mass function in the field. This result could indicate an environmental or dynamical effect, but our number statistics are not yet sufficient to place strong constraints on its nature.  Our survey limits encompass the entire brown dwarf mass range and we detected two companions with $q\\le$0.1, a number which is consistent with a flat mass ratio distribution. However, both of these companions have mass ratios near 0.1 and only one has a mass which might fall below the substellar boundary, so we hesitate to rule out the existence of any deficit that might denote a brown dwarf ``desert''. Our survey's deep detection limits also extend into the top of the planetary mass function; we have not identified any planetary companions at high confidence ($\\ga$99.5\\%), but we have identified four candidate companions at lower confidence ($\\ga$97.5\\%) that merit additional followup to confirm or disprove their existence. The lack of planets within the brown dwarf mass range also is not a significant proof of the existence of a desert.  Finally, we note that our survey results are extremely encouraging with respect to the potential for future discoveries. We are currently extending our survey efforts to the Taurus-Auriga star forming region and to several nearby moving groups, and this expansion of our sample should make any conclusions much more robust. Our ability to precisely measure astrometry for close ($\\sim$2-3 AU) binary systems could also allow us to measure dynamical masses for many young stars on a timescale of $\\la$5 yr. Finally, achieving similar detection limits for planetary-mass companions in Taurus-Auriga and the nearby moving groups will significantly enhance our limits on the properties of young planets; a similar null detection for our full sample would significantly rule out the canonical values for the planetary distribution function, confirming either that these values are wrong or that evolutionary models significantly overestimate the luminosity (and detectability) of young planets."
    },
    "0801/0801.4724_arXiv.txt": {
        "abstract": "{We compare the small-scale features visible in the Ne~{\\sc{viii}} Doppler-shift map of an equatorial coronal hole (CH)  as observed by SUMER with the small-scale structures of the magnetic field as constructed from a simultaneous photospheric magnetogram by a potential magnetic-field extrapolation. The combined data set is analysed with respect to the small-scale flows of coronal matter, which means that the Ne~{\\sc{viii}} Doppler-shift used as tracer of the plasma flow is investigated in close connection with the ambient magnetic field. Some small closed-field regions in this largely open CH are also found in the coronal volume considered. The Doppler-shift patterns are found to be clearly linked with the field topology. ",
        "introduction": "In this study we present observations of an equatorial coronal hole (CH) and investigate the magnetic field structures in the source regions of the fast solar wind. We directly compare the Dopplergrams obtained by the SUMER (Solar Ultraviolet Measurements of Emitted Radiation) instrument, yielding the plasma flow velocity, with the photospheric magnetograms at the bottom of the CH, which via magnetic field extrapolation gives us the coronal magnetic field. In the case studied here we thus combine spectroscopic data with the three dimensional (3-D) magnetic field, an approach providing a clearer physical picture of the plasma conditions and flow pattern prevailing in the corona. This CH was studied before by \\cite{xia03} and \\cite{xia2003}. Therefore we refer the reader to their papers for further details.  The new aspect here is the 3-D magnetic field, which is constructed from magnetograms and extrapolated to all heights above the photosphere. We use the Greens function method to compute the current-free magnetic field from the scalar potential, \\begin{eqnarray} \\Phi({\\bf r}) = -\\frac{1}{2 \\pi} \\int_{\\partial \\Omega} B_z({\\bf r^{'}}) \\, \\frac{d \\sigma^{'}}{|{\\bf r}-{\\bf r^{'}}|}, \\;\\; {\\bf B} =  \\nabla \\Phi, \\end{eqnarray} where $B_z$ is the line-of-sight (LOS) photospheric magnetic field, $\\partial \\Omega$ denotes the bottom boundary surface (photosphere) of the computational box, $d \\sigma = dx \\, dy$, ${\\bf r}=\\sqrt{x^2+y^2+z^2}$, and ${\\bf B}({\\bf r})$ is the 3-D magnetic field. {\\tw Details of the extrapolation technique for CHs and the quiet Sun have been described in earlier work by \\cite{solanki2004}.} Potential fields can be reconstructed from the LOS component of the photospheric magnetic field alone. The corresponding measurements are available, e.g., from the MDI (Michelson Doppler Imager) LOS magnetograph data on SOHO.  For the reconstruction of the dynamic coronal magnetic field, such as in twisted loops, it is usually necessary to include the currents, which often are assumed to be parallel to the magnetic field. These force-free magnetic fields are mathematically more difficult to model, and one needs for their reconstruction additional observational input data, such as the photospheric magnetic field vector \\citep{wiegelmann2004}, or optical images of the coronal plasma structures. This method has been used by \\cite{marsch2004}, who carried out a comprehensive study of fields and flows in active regions (ARs) associated with sunspots. \\begin{figure*} \\setlength{\\unitlength}{1.0cm} \\begin{picture}(12,6) \\put(-0.7,0){\\includegraphics[width=12cm]{fig1.jpg}} \\end{picture} \\parbox{5.5cm}{\\vspace{-4.5cm}\\caption{Magnetic structures in a coronal hole. The gray-coding shows the field strength in the photosphere. The black line gives roughly the boundary of the coronal hole. {\\tw The field of view for SUMER is marked as a red rectangle in both panels.} The magnetic field was constructed from a MDI magnetogram. Left figure: mostly closed loops at various scales. Only closed magnetic field lines with $B \\ge 30$~G are shown. Right figure: Only open fields with large photospheric values, $B \\ge 100$~G. The open flux is bundled in narrow uniform filaments and originates in stronger fields concentrated at small-scale footpoints. The flux tubes expand as they extend into the corona. The average expansion factor (ratio of the maxium to the minimum of the flux tube area) is $28 \\pm 11$.} \\label{fig1}} \\end{figure*} ",
        "conclusions": "The present study has emphasized the need to analyze solar ultraviolet images and spectrogramms or Doppler-shift maps in close connection with the coronal magnetic field, which can routinely be constructed and extrapolated to the outer corona from photospheric magnetograms. The field is the key player in the low-beta corona in constraining plasma flow and guiding the nascent solar wind outflow through open coronal funnels, which according to Fig. \\ref{fig3}  have a more complicated geometry at low heights (below about $10$\\arcsec) than assumed in the models considered by \\cite{marsch1997} or \\cite{hackenberg2000}. The flux tube geometry in the transition region may be of importance to the initial acceleration of the solar wind (see, e.g., \\citeauthor{2002ApJ...566..562L} \\citeyear{2002ApJ...566..562L}).   The results in Fig. \\ref{fig2} are of central importance to appreciate the key role played by the coronal magnetic field. They corroborate previous findings of SUMER, and complement the results published by \\cite{xia2003} and in the thesis of \\cite{xia03}. The Ne~{\\sc{viii}} Doppler-shift maps show a close relationship with the magnetic carpet structure and funnels inside the CH. Largest blue shifts with speeds up to 20~km~s$^{-1}$ (darkest blue patches) are associated with those regions where intense open magnetic fields of a uniform polarity are concentrated. In contrast, the plasma confined in small loops in the CH does not reveal any significant flow."
    },
    "0801/0801.1101_arXiv.txt": {
        "abstract": "Motivated by the recently discovered class of faint (10$^{34}$--10$^{35}$ \\eps) X-ray transients in the Galactic Center region, we investigate the 2--10 keV properties of classical and recurrent novae.  Existing data are consistent with the idea that all classical novae are transient X-ray sources with durations of months to years and peak luminosities in the 10$^{34}$--10$^{35}$ \\eps\\ range.  This makes classical novae a viable candidate class for the faint Galactic Center transients.  We estimate the rate of classical novae within a 15 arcmin radius region centered on the Galactic Center (roughly the field of view of \\xmm\\ observations centered on Sgr A*) to be $\\sim$0.1 per year.  Therefore, it is plausible that some of the Galactic Center transients that have been announced to date are unrecognized classical novae.  The continuing monitoring of the Galactic Center region carried out by \\chandra\\ and \\xmm\\ may therefore provide a new method to detect classical novae in this crowded and obscured region, where optical surveys are not, and can never hope to be, effective. Therefore, X-ray monitoring may provide the best means of testing the completeness of the current understanding of the nova populations. ",
        "introduction": "Recently, several groups have reported their detections of relatively faint X-ray transients in the \\chandra\\ and \\xmm\\ observations of the Galactic Center region \\citep{Pea2005,Sea2005,Mea2005}. These authors conclude that these transients are collectively located near the Galactic Center, based on their absorbing columns and their sky distribution, although no direct distance measurements are available. With this assumption, the inferred luminosities of these transients are in the 10$^{34}$--10$^{35}$ \\eps\\ range.  The authors of these studies claim that such a luminosity is too high for cataclysmic variables (CVs), semi-detached binaries in which the accreting object is a white dwarf. Instead, they argue for neutron star or black hole accretors based solely on the luminosity.  However, the Galactic Center transients are sub-luminous compared to the known transient populations of black hole or neutron star binaries \\citep{Sea2005,Mea2005}, requiring a new population (see, e.g., \\citealt{KW2006}).  The accretion driven X-ray luminosities of CVs are indeed insufficient to explain the Galactic Center transients.  Non-magnetic CV X-ray luminosities are in the range 10$^{30}$--10$^{32}$ \\eps, with the highest value being 3$\\times 10^{32}$ \\eps\\ for the old nova, V603~Aql \\citep{Bea2005}. Magnetic CVs, the intermediate polars (IPs) in particular, are more luminous in 2--10 keV X-rays, with estimated luminosities often exceeding 10$^{33}$ \\eps\\ \\citep{SRea2006}.  However, since the highest luminosity recorded for an IP is 1.3$\\times 10^{34}$ \\eps\\ during the outbursts of the unusual IP (and another old nova), GK~Per \\citep{Hea2004}, IPs are also not likely candidates for the Galactic Center X-ray transients.  However, the above discussion is incomplete because it is limited to the accretion driven X-ray luminosities of CVs.  In reality, CVs can generate higher X-ray luminosities through nuclear fusion, which is a more efficient source of energy than accretion onto a white dwarf.  Indeed, classical novae have been known to emit 2--10 keV X-rays at luminosities exceeding 10$^{34}$ \\eps.  We present below a summary of X-ray properties of classical as well as recurrent novae. ",
        "conclusions": "Classical and recurrent novae are a known class of transient X-ray sources that reach luminosities in the 10$^{34}$ -- 10$^{35}$ \\eps\\ range.  The shell X-ray phase of novae may last months to several years, although they probably soften as they age, gradually making them less conspicuous above 2 keV.  Novae have the right spectral and temporal characteristics to explain some of the faint Galactic Center transients that have been detected with \\chandra\\ and with \\xmm\\ in recent years.  If the existing literature accurately reflects the rate of X-ray transients near the Galactic Center, then the known population of classical (and recurrent) novae are probably a small, but not negligible, contributor to the overall transient population.  \\citet{Mea2005} have argued that dynamical processes may lead to a high space density of X-ray binaries within 1 pc of the Galactic Center. That is, the concentration of X-ray emitters is forcing considerations of a new population of objects not seen elsewhere in the Galaxy. We propose that any such studies include white dwarf binaries, since Galactic Center specific processes could produce an additional population of novae beyond disk and bulge novae that are currently known. The combination of X-ray monitoring and population synthesis may represent our best hope of obtaining a complete picture of nova populations in the Galaxy, because optical surveys cannot possibly be effective in the Galactic Center region.  Note that optical surveys fail to detect the majority of Galactic novae overall --- compare the inferred Galactic nova rate of $\\sim$30 yr$^{-1}$ to the actual rate of discovery ($\\sim$6 yr$^{-1}$).  \\chandra\\ and \\xmm\\ have been monitoring the Galactic Center region more or less regularly over the last $\\sim$8 years.  Even at 0.1 novae per year within 15 arcmin of the Galactic Center, it is probable that a nova will be detected as a 2--10 keV X-ray transient soon, if one hasn't been already.  A concurrent infrared monitoring campaign will be required, however, to prove beyond a reasonable doubt that a particular X-ray transient is due to a classical nova.    \\clearpage"
    },
    "0801/0801.1930_arXiv.txt": {
        "abstract": "Red-sequence galaxies record the history of terminated star-formation in the Universe and can thus provide important clues to the mechanisms responsible for this termination.  We construct composite samples of published cluster and field galaxy photometry in order to study the build-up of galaxies on the red-sequence, as parameterised by the dwarf-to-giant ratio (DGR).  We find that the DGR in clusters is higher than that of the field at all redshifts, implying that the faint end of the red-sequence was established first in clusters.  We find that the DGR evolves with redshift for both samples, consistent with the ``down-sizing'' picture of star formation.  We examine the predictions of semi-analytic models for the DGR and find that neither the magnitude of its environmental dependence nor its evolution is correctly predicted in the models.  Red-sequence DGRs are consistently too high in the models, the most likely explanation being that the strangulation mechanism used to remove hot gas from satellite galaxies is too efficient.  Finally we present a simple toy model including a threshold mass, below which galaxies are not strangled, and show that this can predict the observed evolution of the field DGR. ",
        "introduction": "Red-sequence galaxies are an important population for understanding galaxy formation and evolution in general.  They represent systems with very little or no on-going star-formation and thus are a unique tracer of the past activity of galaxies.  \\citet{De-Lucia:2004xa} found a deficit of faint red-sequence galaxies in clusters at z$\\sim$0.8 relative to local clusters.  This would suggest that star formation ended earlier for the most luminous/massive galaxies at high redshift and would support the ``down-sizing'' picture first proposed by \\citet{Cowie:1996xw}, in which the termination of star-formation progresses from the most massive to the least massive galaxies as the universe ages.  Similar results for clusters spanning a range of redshifts have been found by \\citet{Tanaka:2005mk}, \\citet{de-Lucia:2007li}, \\citet{Stott:2007wc}, \\citet{Gilbank:2007rq} and others.  Similarly, surveys of field galaxies have attempted to characterise the evolution of the red-sequence luminosity function.  The relative contributions of passive evolution/termination of star-formation and number density evolution/dry merging are still open questions (e.g., \\citealt{Bell:2004lb},  \\citealt{Cimatti:2006vz}).   Red-sequence galaxies can thus place constraints on galaxy formation prescriptions in theoretical models.  Recent observations in the local universe using SDSS data \\citep{Baldry:2006bq,Weinmann:2006cn} have found that the fractions of satellite galaxies which are red are too high in current semi-analytic models.  With the recent abundance of surveys targeting red-sequence galaxies both in the cluster and the field, the time is now ripe to draw together these works to build a picture of the build-up of galaxies on the red-sequence as a function of environment and to confront these with predictions of galaxy formation models. ",
        "conclusions": "Previous works have suggested that there appears to be little or no evolution in the bright end of the RSLF, both in clusters and the field (e.g., \\citealt{Gilbank:2007rq}; \\citealt{Scarlata:2007oz}).  If this is the case, the evolution of the DGR can be attributed to the build-up of faint galaxies on the red-sequence with increasing cosmic time.  This implies that star-formation is terminated in giant galaxies first and later in dwarf galaxies, in agreement with the ``down-sizing'' picture of star-formation. Furthermore, the lower value of DGR in the field compared with clusters implies that star-formation was completed in cluster galaxies before field galaxies.  Performing a similar analysis on the B06 semi-analytic model predicts a non-evolving DGR out to z$=$1, with little difference between clusters and the field.  Such a result suggests that faint galaxies move onto the red-sequence too efficiently in the model, with this build-up of the faint end of the red-sequence already having been completed by the earliest epoch considered here (z$\\sim$1).  The methods available in the models for the quenching of star-formation include: feedback from supernovae or AGN and hot halo gas removal. This latter process only applies to galaxies which are satellites (i.e. not the most massive, central object) in their dark matter halo.  The relative importance of the various quenching mechanisms are summarised in fig.~4 of B06.  At the bright end of the red-sequence, most of the galaxies are the central object in their dark matter halo and thus feedback from AGN or supernova winds are required to turn off star-formation.  At the faint end (in the dwarf regime), red-sequence galaxies are predominantly satellites. We have verified that this is still the case for all the redshifts we consider here.  In most galaxy formation models (e.g., B06, \\citealt{De-Lucia:2006sa}), when any galaxy enters a more massive dark matter halo and becomes a satellite, it instantly has all of its hot gas removed in a process generally referred to as `strangulation` \\citep{Larson:1980wj,Balogh:2000yv}. Thus, once it has consumed the cold gas already present in its disk, its fuel can no longer be replenished by hot gas from the halo and its star-formation is terminated.  This prescription may be too efficient.  Recent simulations by \\citet{McCarthy:2007mr}  found that the removal of the initial hot galactic halo gas does not occur instantaneously, but likely takes $\\sim$ 1 Gyr to remove the bulk of it, and that as much as 30\\% can still remain after 10 Gyr.  In addition, the models of \\cite{Dekel:2006ii}  predict the existence of a critical halo mass ($\\sim10^{12}$ M$_\\odot$) for haloes, below which strangulation may not occur.  Thus, we expect additional parameters involving strangulation efficiency/timescale and/or a threshold halo mass for strangulation to be important in tuning the models to reproduce the observations presented in Fig.~1.  To illustrate the effect of a threshold halo mass for strangulation, we use the following toy model to predict the evolution of the field DGR.  Field surveys do not just probe isolated (i.e. non-group) galaxies.  Indeed, they are dominated by groups, especially toward lower redshift as cosmic structure grows. Thus it is group-scale physics which is important in determining the properties of galaxies in field surveys.  We assume further that galaxies in haloes below some (group-size) threshold mass, $M_{th}$, have no red-sequence satellites, i.e. DGR $=0$.  We assume that haloes above $M_{th}$ have a DGR given by that of the observed cluster population at the same redshift (from the best-fit cluster line in Fig.~1).  We consider $M_{th} = (10^{12}, 10^{12.5}, 10^{13}$) M$_\\odot$ and measure the fractional contribution of haloes in the appropriate mass range from the Millennium Simulation. We overplot the expected evolution of the field DGR for the different threshold masses in Fig.~1.  It can be seen that both the $M_{th} = 10^{12}$ and $10^{12.5}$ M$_\\odot$ models give good agreement with the observed field DGR. In order to reproduce the observations so well, our model requires both an increasing contribution to the field by groups due to growth of structure {\\it and} an intrinsic evolution of the DGR in haloes above $M_{th}$.  If either of these factors is removed (e.g. if a constant group fraction at all redshifts or a constant cluster DGR at all redshifts is assumed) then the predicted evolution of the field DGR is too flat with respect to the observations.  If groups do not exhibit exactly the same DGR as clusters, but instead are intermediate between clusters and the field (this is currently an open question), then this would effectively shift our predicted lines down and a lower $M_{th}$ would then reproduce the observations.  Of course, a scenario in which the timescale for the removal of hot gas is much longer may fit the observations equally well and this or some alternative combination of threshold mass/timescale cannot be ruled out.  Indeed, the observed evolution of the cluster DGR we have had to include in our model may reflect the effect of a strangulation timescale.  Disentangling the effects of a timescale and threshold mass for the quenching of satellite galaxies is a complex observational problem.  The colour difference between central and satellite red-sequence galaxies may be a useful test of models.  For example, in the B06 prescription, it is predicted that satellite red-sequence galaxies are significantly redder than central red-sequence galaxies.  \\citet{van-den-Bosch:2007da} recently measured the average colours of satellite and central galaxies in the local universe from an approximately halo mass-selected sample and found that the former were indeed on average redder.  They also showed that measuring halo masses from statistical samples for $\\sim$10$^{12}$ M$_\\odot$ systems is not possible with high completeness and thus, if there is a threshold mass for strangulation around this limit, such studies will require detailed targeted observations of groups (e.g., \\citealt{Wilman:2005rd}).  An alternative approach to get at the timescale for strangulation is the study of the star formation histories of individual galaxies undergoing transformation (e.g., \\citealt{Moran:2007zk}).  It is possible that the observations in Fig.~1, coupled with new strangulation prescriptions in the semi-analytic models may improve constraints on the likely values of the threshold mass and timescale for strangulation."
    },
    "0801/0801.3104_arXiv.txt": {
        "abstract": "We present limits on planetary companions to pulsating white dwarf stars. A subset of these stars exhibit extreme stability in the period and phase of some of their pulsation modes; a planet can be detected around such a star by searching for periodic variations in the arrival time of these pulsations. We present limits on companions greater than a few Jupiter masses around a sample of 15 white dwarf stars as part of an on-going survey. One star shows a variation in arrival time consistent with a 2\\,$\\mathrm{M_J}$ planet in a 4.5 year orbit. We discuss other possible explanations for the observed signal and conclude that a planet is the most plausible explanation based on the data available. ",
        "introduction": "}   All main-sequence stars with mass less than about 8\\,M$_{\\odot}$ will end their lives as white dwarf stars (WDs). As such, WDs are a fossil record of star formation in the Galaxy from earliest times to just a few hundred million years ago. WDs offer a window into the ultimate fate of planetary systems, including our own solar system, and whether planets can survive the final stages of stellar evolution. The properties of a WD are relatively insensitive to the mass of the progenitor: the mass distribution of isolated WDs is narrowly distributed in a peak 0.1\\,M$_{\\odot}$ wide around a mean mass of 0.59\\,M$_{\\odot}$ \\citep{Kepler07}. Surveys of WDs can therefore search a wide range of main sequence progenitor masses and the low luminosity of a WD means any companion planets can be potentially followed up with current direct detection technology in the mid-infrared.  \\citet{Livio84} considered the fate of a planet engulfed in the envelope of a red giant star. They determined that below a certain mass the planet would be evaporated and destroyed, but larger objects would accrete material and spiral in toward the stellar core. They predicted that the end state of these systems would be a tight binary consisting of a white dwarf and a brown dwarf and suggested this mechanism might explain the origin of cataclysmic variable (CV) systems.   Planets that are not engulfed, and sufficiently far from the stellar surface that tidal drag is small, will drift outwards to conserve angular momentum \\citep[as described in][]{Jeans24}. \\citet{Duncan98} investigated the stability of the outer solar system when the sun undergoes mass loss as a red giant. They found that the system was stable on timescales of at least 10 billion years for reasonable amounts of mass loss, but for larger amounts typical of WDs formed from more massive stars, the planets' orbits became unstable on timescales of $\\lesssim 10^{8}$ years. \\citet{Debes02} looked what would happen if the orbits of two planets became unstable and determined that if orbits crossed, the most likely result was that one planet would be scattered into a shorter period orbit, while the other would be boosted into a longer orbit or ejected from the system. For planets scattered inward, the extreme flux from a newborn white dwarf would strip the outer atmospheric layers. \\citet{Villaver07} estimated that a 2\\,$\\mathrm{M_J}$ planet 1.8\\,AU from a young white dwarf would lose half its mass in this manner. The picture drawn by this brief survey of theory is a population of planets in long period orbits around WDs, with a number of objects scattered closer to the star, or in very tight binaries.  Searches for sub-stellar companions to WDs have concentrated on exploiting the lower contrast between star and companion, especially in the infrared \\citep[e.g.][]{Probst83, Burleigh02, Farihi05survey, Friedrich06, Debes06, Mullally07}. Although a couple of brown dwarf stars have been found with this approach \\citep{Zuckerman92, Farihi04}, as yet no direct detections of planets have been claimed. \\citet{Silvotti07}, using the same timing method discussed here, report the detection of a planetary mass companion in a 1.7\\,AU orbit around an extreme horizontal branch sdB star.  Pulsating WDs allow the possibility of searching for the presence of companion planets as changes in the observed arrival time of the  pulsations. Hydrogen atmosphere (DA) WDs pulsate in an instability strip approximately 1200\\,K wide near 12,000\\,K and are known as DAVs (or ZZ~Ceti stars). The pulsations are non-radial g-modes with periods of order 100--1500s and amplitudes of a few percent.  The pulsation properties of DAVs vary with temperature. Those near the hot end of the strip tend to have a smaller number of shorter period, lower amplitude modes with sinusoidal lightcurves and are known as hDAVs, while those nearer the red edge show more, larger amplitude modes and asymmetric pulse shapes. The change in pulsation properties is most likely due to increasing depth of the convection layer near the surface of the star \\citep[][and subsequent articles]{Brickhill83}.  \\citet{Stover80} first noted that modes on hDAVs often exhibited an impressive stability in the period and phase of pulsation. \\citet{Kepler05a} measured the rate of period change, $\\dot{P}$, of one mode in the hDAV \\object[WD0912+354]{G117-B15A} to be $3.57(82) \\times 10^{-15}$, while \\citet{Mukadam03} constrained $\\dot{P}$ of a mode in \\object[WD0133-116]{R548} to $\\le 5.5(1.9) \\times 10^{-15}$ and \\citet{ODonoghue87} constrained $\\dot{P}$ of \\object[WD1425-811]{L19-2} to $<3.0 \\times 10^{-14}$. Together with pulsars, these objects are the most stable astrophysical clocks known.   Measurements of $\\dot{P}$ require datasets of between 10 and 30 years, but the investment in time yields a suitable scientific reward. Measurement of $\\dot{P}$ provides a rare opportunity to directly test models of structure and composition of the core of a star \\citep{Kepler05a}, constrain the current rate of change of the gravitational constant \\citep{Benvenuto04}, as well as provide useful constraints on the mass of the hypothesized axion or other super-symmetric particles \\citep{Isern92, Corsico01, Bischoff-Kim07b}  If a planet is in orbit around a star, the star's distance from the Sun will change periodically as it orbits the center of mass of the planetary system. If the star is a stable pulsator like a hDAV, this will cause a periodic change in the observed arrival time of the otherwise stable pulsations compared to that expected based on the assumption of a constant period. The change in arrival time, $\\tau$, is given by \\begin{equation} \\tau = \\frac{a_p m_p \\sin{i}}{M_* c} \\end{equation}  \\noindent where $a_p$ is the semi-major orbital axis of the planet, $m_p$ is the planet mass, $M_*$ is the mass of the white dwarf, $c$ is the speed of light and $i$ is the inclination of the orbit to the line of sight. In common with astrometric methods, the sensitivity increases with the orbital separation, making long period planets easier to detect given data sets with sufficiently long baselines.  In 2003 we commenced a pilot survey of a small number of DAVs in the hope of detecting the signal of a companion planet. We present here a progress report of the first 3--4 years of observations on 12 objects, as well as presenting limits around 3 more objects based partly on archival data stretching as far back as 1970. For one object we find a signal consistent with a planetary companion. Further observations are necessary to confirm the nature of this system. For our other objects, we can constrain the presence of planets down to a few Jupiter masses at 5\\,AU, with more stringent limits for stars with archival data. ",
        "conclusions": "We present our results on an on-going survey for planets around 15 pulsating white dwarf stars. Our survey data spans 3--4 years with archival data on some stars stretching back to 1970. We are already sensitive to planets down to a few Jupiter masses at distances of approximately 5\\,AU. For one star, we observe a curvature in the O-C diagram consistent with a planet in a 4.5 year orbit. Further observations are necessary to span a full orbit of this candidate object. If confirmed, this will be the first planet discovered around a WD, and together with the planet discovered around an sdB, suggests that planets regularly survive the death of their parent star, and that WDs will be fruitful targets for planet searches."
    },
    "0801/0801.4381_arXiv.txt": {
        "abstract": "{} {We have compiled one of the largest normal-galaxy samples ever to probe \\x\\ luminosity function evolution separately for early and late-type systems.} {We selected 207 normal galaxies up to redshift $z\\sim 1.4$, with data from four major \\chandra\\ \\x\\ surveys, namely the \\chandra\\ deep fields (north, south and extended) and XBootes, and a combination of \\x\\ and optical criteria. We used template spectral energy-distribution fitting to obtain separate early- and late-type sub-samples, made up of 101 and 106 systems, respectively.  For the full sample, as well as the two sub-samples, we obtained luminosity functions using both a non-parametric and a parametric, maximum-likelihood method.} {For the full sample, the non-parametric method strongly suggests luminosity evolution with redshift. The maximum-likelihood estimate shows that this evolution follows $\\sim (1+z)^{k_{\\rm total}}$, $k_{\\rm total}=2.2\\pm0.3$. For the late-type sub-sample, we obtained $k_{\\rm late}=2.4^{+1.0}_{-2.0}$. We detected no significant evolution in the early-type sub-sample. The distributions of early and late-type systems with redshift show that late types dominate at $z\\ga 0.5$ and hence drive the observed evolution for the total sample. } {Our results support previous results in \\x\\ and other wavebands, which suggests luminosity evolution with $k=2-3$. } ",
        "introduction": "Non-AGN dominated (\\lq\\lq normal\\rq\\rq) galaxies have intrinsically weak \\x\\ luminosities and correspondingly faint fluxes. Thus, it is only with the most recent, new-generation X-ray missions, \\chandra\\ and \\xmm, that such galaxies have been detected in cosmologically significant redshifts ($z>0$). The first such \\x\\ selected sample was obtained with the \\chandra\\ Deep Fields North and South \\citep[CDF-N, CDF-S;][]{2003AJ....126..539A,2002ApJS..139..369G}, reaching fluxes of $f(0.5-2.0 \\ {\\rm keV}) \\sim 10^{-17}$~\\funits \\citep[see also][]{2003AJ....126..575H}.  Based on these observations, \\citet{2004ApJ...607..721N} obtained an \\x\\ luminosity function (XLF) of normal galaxies, characterised by luminosity evolution that follows $\\sim (1+z)^{2.7}$.  A population of rapidly evolving, star-forming galaxies has been detected in recent years in other wavebands \\citep[e.g.][and references therein]{2004ApJ...615..209H}.  However, observations in the \\x\\ band provide unique insight into physical processes, complementing information obtained from other wavelength regions.  In the most massive and luminous early-type galaxies, the X-ray emission is dominated by the hot interstellar medium (ISM) at $kT\\sim 1$~keV, with a smaller fraction contributed by low-mass X-ray binaries (LMXBs) associated with the old stellar population. However, for fainter \\x\\ systems, the evidence suggests that LMXBs may well be the dominant X-ray emitting component \\citep[e.g.][and references therein]{2000ApJ...544L.101S,2003ApJ...586..826K}.  On the other hand, the X-ray emission in late-type galaxies is mainly due to a mixture of low- and high-mass X-ray binaries (HMXBs) with a contribution from hot gas ($kT\\sim 1$~keV) \\citep[see][for a comprehensive review]{2006ARA&A..44..323F}.  Understanding how these populations of binary stars and, consequently, their \\x\\ luminosity evolve with time is obviously closely linked to XLF normal-galaxy evolution. Considering LMXBs and HMXBs as the overall dominant component in \\x\\ emission from normal galaxies, \\citet{2001ApJ...559L..97G} adopt a semi-empirical approach to link \\x-binary lifetimes with star-formation rates (SFRs) in a cosmological context.  They show that evolving SFRs significantly affect \\x\\ binary populations, hence, integrated \\x\\ galactic emission, with the possibility of significant evolution in \\x\\ luminosities even in relatively low redshifts $\\approxlt 1$ \\citep[see also][]{1998ApJ...504L..31W}. They predict different galaxy evolution rates at \\x\\ wavelengths, as compared to other wavelength regions, which depend on the star-formation history of the Universe, as well as evolutionary timescales for LMXBs and HMXBs. Conversely, if the \\x\\ evolution of normal galaxies is known, one may obtain insight into properties such as the characteristic timescales of LMXBs and HMXBs. For instance, if HMXBs tracing the instantaneous SFR dominate the integrated \\x\\ emission, the total \\x\\ luminosity is expected to follow the star-formation history of the Universe as observed in the optical and IR bands. On the contrary, LMXBs have much longer evolutionary timescales, and the integrated \\x\\ luminosity of normal galaxies would present a time delay of the order of $\\sim 1$~Gyr compared to optical and IR observations. Constraining the number density and evolution of LMXBs is also relevant to the LISA gravitational wave mission, as this will be sensitive to gravitational radiation from binaries with periods shorter than 4 hours. Such sources are primarily LMXBs. Determining LMXB evolutionary timescales can thus provide information on the number of expected LISA detections.  As the relative contribution of LMXBs and HMXBs to the integrated \\x\\ luminosity of normal galaxies closely depends on galaxy type, it is imperative to disentangle XLF behaviour among different galaxy types. Unfortunately, because of the scarcity of \\x-detected normal galaxies, it is not surprising that very few results have been obtained.  Using moderate size samples, \\citet{2005MNRAS.360..782G} and \\citet{2006MNRAS.367.1017G} calculated XLFs separately for early-type/absorption-line and late-type/emission-line galaxies. By comparing the predictions of the latter XLF with observed normal-galaxy number counts, \\citet{2006ApJ...641L.101G} detected luminosity evolution proportional to $\\sim (1+z)^{2.7}$ for {\\it late} types, driven by sources with \\lxo~$>-2$. \\citet{2007MNRAS.377..203G} then compared the predictions of this XLF with observed counts for galaxies selected by exploitation of the tight \\x-IR correlation and detected luminosity evolution proportional to $\\sim (1+z)^{2.4}$ for their sample of star-forming galaxies. Recently, \\citet{arXiv:0706.1791v1} have used the GOODS survey to obtain 40 early-type and 46 late-type galaxies up to a redshift $z\\sim 1.2$. Their XLFs suggest luminosity evolution proportional to $\\sim (1+z)^{1.6}$ for early types and $\\sim (1+z)^{2.3}$ for late types.  In this paper we aim to substantially increase the numbers of \\x\\ detected normal galaxies, to be able to investigate XLF evolution separately for different galaxy types. We compiled one of the largest normal galaxy samples ever, and, for the first time, probed redshift evolution directly and independently for early and late-type systems.  The structure of this paper is as follows: In \\scr{sec:sel} we describe galaxy selection for our sample.  In \\scr{sec:lf} we present our XLFs, obtained using two different methods. We present our results in \\scr{sec:res}, and discuss them in \\scr{sec:disc}. We conclude with predictions related to future observations and missions in \\scr{sec:fut}. ",
        "conclusions": "\\label{sec:disc} \\subsection{Comparison with previous results} To probe XLF evolution for different galaxy types, large numbers are necessary. \\citet{2005MNRAS.360..782G} used 46 galaxies to construct XLFs for emission- and absorption-line systems. However, their results were for a single redshift bin $z<0.22$. \\citet{2006MNRAS.367.1017G} had 67 galaxies, likewise limited to a single redshift bin $z<0.2$. Although \\citet{2004ApJ...607..721N} had a large sample of 210 galaxies, they constructed no galaxy-type specific sub-samples. It is thus the first time that XLFs were calculated directly for early and late-type normal galaxies. The large size of our total sample allows splitting into three redshift bins that, in turn, strongly suggest evolution with redshift. After repeating the procedure separately for early and late-type systems, we detected no evolution for early types and strong evolution for late types.  Our results are mostly in good agreement with related work in the literature. The results for 36 normal galaxies ($z=0.01 \\rightarrow 0.3$) by \\citet{2006ApJ...644..829K} are shown by stars in \\fr{fig:xlf-total.eps}. The results from \\citet{2004ApJ...607..721N} for two redshift bins are also shown. When compared to our results, these show an apparent systematic shift to higher luminosities, which might suggest AGN contamination. Note, however, that for the sake of clarity we do not show errors from \\citet{2004ApJ...607..721N}, which would bring their points and, in particular, the apparently markedly discrepant point at $(\\log L_X, \\, \\log \\, \\phi) \\sim (40.2,-1.7)$ into better agreement with ours.  In Figs.~\\ref{fig:xlf-early.eps} and \\ref{fig:xlf-late.eps} we also show the results for early- and late-type galaxies by \\citet{2006MNRAS.367.1017G}. We constructed the XLF curves by using the function parameters quoted by these authors, but also imposed luminosity evolution following the $k$ values we found in the present paper. For early types we see that our largest sample cannot be adequately parametrised by their XLF, which shows too steep an exponential cutoff. The late-type XLF shows broad agreement with our non-parametric bins, but it is clearly not a good representation of our data. The discrepancies, especially for early types, may partly stem from these authors using a selection criterion \\lxol2, thus obtaining fewer luminous sources.   \\subsection{Evolution} The fact that late-type galaxies are driving the observed evolution of the total sample can be understood if we look at the redshift distributions of the two galaxy types.  The histograms in \\fr{fig:Nz.eps} show the observed distributions: Early-types dominate in the lowest redshift bins, whilst late-types dominate at intermediate redshifts. A Kolmogorov-Smirnov (KS) test shows that the two distributions differ significantly. The probability that the value of the KS statistic $D$ obtained may be due to chance alone is very small, $p=0.006$.  The observed trend is also corroborated by theoretical distributions, shown by curves in the same figure, which we construct using our ML fit parameters together with the area curve information of our data. Qualitatively, this also agrees with the results of \\citet{2005ApJ...625..621B}. These authors used a large sample from the Great Observatories Origins Deep Survey (GOODS) fields to probe mass assembly of morphologically distinct normal galaxies. They found an increasing proportion of early-type systems with decreasing redshift from $z\\sim 1$ to 0, which is similar to what we see in our redshift distribution.  Furthermore, we found that it is the late types with highest \\lxo\\ that are mostly responsible for the observed trend. In \\fr{fig:Nz.eps} we see that numbers for late-types appear to rise after $z=0.5$. Out of 53 late-type systems at $z>0.5$, 39, or 74\\%, have $\\log (f_X/f_R) > -2$. This finding agrees with the claim by \\citet{2006MNRAS.367.1017G} that their deduced evolution stems primarily from such sources.  It also provides {\\it a posteriori} justification of our upper-limit choice ($\\log (f_X/f_R) \\le -1$) in the selection criteria for the present sample.  We cannot exclude that this observation is, at least partially, a consequence of a selection effect. Given the steep relation $L_X\\propto L_B^{1.5}$ for spirals \\citep{2001ApJS..137..139S}, we might expect to preferentially detect intrinsically brighter galaxies at higher redshift, which would lead to a selection against low $f_X/f_R$ values. Even so, that we see bright {\\it late}-types may still be significant. As \\citet{2006A&A...454..447T} have shown, at lower redshift, bright galaxies with \\lxo~$>-2$ are mostly early-type.  It is encouraging that our luminosity evolution estimates are in good agreement with results by other authors.  Taken together, \\citet{2004ApJ...607..721N}, \\citet{2004ApJ...615..209H}, \\citet{2005A&A...440...23R}, \\citet{2006ApJ...641L.101G}, \\citet{2007MNRAS.377..203G}, \\citet{arXiv:0706.1791v1}, as well as the present paper, lead to a broad consensus that {\\it the data are consistent with luminosity evolution with an index $k=2-3$.}  In particular, \\citet{2004ApJ...615..209H} investigates luminosity function evolution for a sample of radio-selected star-forming galaxies by combining constraints from the global SFR density evolution with those from the 1.4 GHz radio source counts at submillijansky levels. He finds luminosity evolution $\\sim (1+z)^{2.7\\pm0.6}$ for star-forming galaxies. Given the tight correlation between radio, far infrared, and \\x\\ luminosity functions \\citep{2003A&A...399...39R}, the agreement of this result with our result for late-type systems is particularly significant for two reasons. Firstly, it supports our main result that late-type systems are driving the overall redshift evolution of normal galaxies in the local Universe. Secondly, as our normal galaxy selection criteria are independent of those for 1.4 GHz sources, it suggests that any AGN contamination is unlikely to have had significant impact on our XLFs.  \\citet{2007MNRAS.377..203G} obtained essentially the same evolution index for late-types, $k_{\\rm late}=2.4$, by employing methods largely independent of ours. Because they used \\x-IR correlations which hold for normal galaxies, but not for AGN, their sample is virtually guaranteed to be free of AGN contamination. Although their sample is $\\sim 4$ times smaller than ours, the agreement in the evolution index is, once more, significant. \\citet{2005A&A...440...23R} showed that the blue galaxy luminosity function for the 25,000-strong spiral galaxy sample of \\citet{2003A&A...401...73W} requires a pure luminosity evolution index $\\la 3$. This is also consistent with our results, under the assumption that blue galaxies are largely the same population as the late-type systems in our sample. \\citet{arXiv:0706.1791v1} jointly fitted low and high-redshift XLFs, obtaining pure-luminosity evolution with $k_{\\rm early} = 1.6$ and $k_{\\rm late} = 2.3$. These authors used Bayesian techniques and a Monte Carlo Markov chain analysis. The agreement, at least for late-types, is thus particularly encouraging.  Comparisons with results for red galaxies are less straightforward. Assuming that our early-type galaxies are the same population as red galaxies in optical surveys, with no significant dust contamination, that, unlike \\citet{arXiv:0706.1791v1}, we do not measure significant evolution for early-type galaxies, disagrees at face value with the observed evolution in optical wavebands \\citep{2007ApJ...654..858B,2003A&A...401...73W}. However, this is consistent with models that predict a fast decline in luminosities at optical wavelengths, with \\x\\ luminosities remaining high for $\\ga 1$ Gyr \\citep{2006IAUS..230..417E}, in systems with a significant contribution from LMXBs. \\citet{1998ApJ...504L..31W} predict that the combined effect of LMXBs and HMXBs may lead to a \\lq twin peak\\rq\\ in the evolving XLF of normal galaxies; due to HMXBs, the first peak would be expected to occur close to the SFR peak at $z\\approx 1.5$. The second peak would be delayed until $z\\approx 0.5-1$, due to delayed turn-on of LMXBs. To first order, our early-type systems are dominated by LMXBs, which could mean that we may approximately be witnessing this second peak. However, this approximation may be inadequate, even to first order, due to the caveat that the most luminous systems are dominated by the hot ISM. Similarly, our late-type systems are dominated by HMXBs, whose \\x\\ luminosity has peaked at high $z$, and we are witnessing their dimming at lower redshifts. Indeed, an evolution index $k\\sim 2.4$ is consistent with \\lq Peak-M\\rq\\ models, which are characterised by longer evolutionary timescales for LMXBs \\citep[][Table 2]{2001ApJ...559L..97G}.  \\subsection{Comparison with AGN} It is not clear whether this scenario is consistent with results for the XLF evolution of AGN-dominated systems. \\citet{2003ApJ...598..886U} and \\citet{2005ApJ...635..864L} found that luminosity-dependent density evolution explains their XLFs best. They showed that XLFs for luminous AGN peak at higher redshifts than for less luminous ones. By assuming that SMBH growth is closely linked to starburst activity, \\citet{2003ApJ...598..886U} linked this to normal-galaxy evolution in a first-order scenario, where luminous AGN once lived in what later became early-type galaxies. Their AGN activity peaked at high $z\\approx 2$, following strong starbursts, and decreased rapidly after $z\\la 2$. On the other hand, galaxies that hosted less luminous AGN at high $z$ have smaller spheroids, i.e. are of later type. Their star-formation peaks at lower $z$, and so does their AGN activity, peaking at $z\\approx 0.6-0.7$.  We did not specifically investigate luminosity-dependent density evolution in this paper. As explained in the previous section, our results supporting pure-luminosity evolution are corroborated by a number of different authors using a variety of data and techniques. This would then cast doubt on a general, unified evolutionary scheme, encompassing both AGN and normal galaxies.  Even so, it is possible to make our results broadly consistent with these AGN evolutionary models, at least in a qualitative sense. Thus early types that, presumably, correspond to bright AGN/high star-formation at high redshift would be expected to show declining luminosities at the redshifts we probe. That we detect no such evolution is already explained by the possibility of delayed LMXB evolution. For late types, we do detect declining luminosities. As the AGN peak redshift is within the range probed by our sample, our selection criteria naturally exclude any AGN, instead picking galaxies whose AGN is either switching off or makes a very low contribution to the total luminosity.       In this paper we have shown the importance of increasing numbers of \\x-luminous normal galaxies for probing XLF behaviour as a function of both redshift and galaxy type. We selected 207 normal galaxies (101 early-type and 106 late-type) from the three \\chandra\\ deep fields and XBootes. Our major results are: \\begin{enumerate} \\item Both methods of XLF estimation, parametric and non-parametric suggest that the full galaxy sample is consistent with pure luminosity evolution, which is driven exclusively by evolution of late-type systems. \\item The evolution index $k$, where $L^{\\ast}\\equiv L^{\\ast}_0 (1+z)^k$, is $2.2\\pm0.3$ for the full sample, \\aer{2.4}{+1.0}{-2.0} for the late-type sample, and \\aer{-0.7}{+1.4}{-1.6} for the early-type sample. \\item Our results agree broadly with results from other work. \\item We estimate that, with a 1 Ms exposure with {\\it XEUS}, counts from normal galaxies will overtake AGN at faint fluxes and will allow probing XLF evolution to high-redshift. \\end{enumerate}"
    },
    "0801/0801.4348_arXiv.txt": {
        "abstract": "We present the latest results on CIB fluctuations from early epochs from deep Spitzer data. The results show the existence of significant CIB fluctuations at the IRAC wavelengths (3.6 to 8 $\\mu$m) which remain after removing galaxies down to very faint levels. These fluctuations must arise from populations with a significant clustering component, but only low levels of the shot noise. There are no correlations between the source-subtracted IRAC maps and the corresponding fields observed with the {\\it HST} ACS at optical wavelengths. Taken together, these data imply that 1) the sources producing the CIB fluctuations are individually faint with $S_\\nu<$ a few nJy at 3.6 and 4.5 $\\mu$m; 2) have different clustering pattern than the more recent galaxy populations; 3) are located within the first 0.7 Gyr (unless these fluctuations can somehow be produced by - so far unobserved - local galaxies of extremely low luminosity and with the unusual for local populations clustering pattern), 4) produce contribution to the net CIB flux of at least 1-2 nW/m$^2$/sr at 3.6 and 4.5 $\\mu$m and must have mass-to-light ratio significantly below the present-day populations, and 5) they have angular density of $\\sim$ a few per arcsec$^2$ and are in the confusion of the present day instruments, but can be individually observable with {\\it JWST}. ",
        "introduction": "The cosmic infrared background (CIB) is a repository of emissions throughout the entire history of the Universe. The recent years have seen significant progress in CIB studies, both in identifying and/or constraining its mean level (isotropic component) and fluctuations (see Kashlinsky 2005 for a recent review). The CIB contains emissions also from objects inaccessible to current (or even future) telescopic studies and can, therefore, provide unique information on the history of the Universe at very early times. One particularly important example of such objects, of particular reference to this conference, concerns Population III stars (hereafter Pop~III), the still elusive zero-metallicity stars expected to have been much more massive than the present stellar populations (see Bromm \\& Larson 2004 for a recent review). Herebelow I will use the term \"era of the first stars\", or \"Pop~III era\", with the understanding that the actual era may be composed of objects of various nature from purely zero-metallicity stars, to low- metallicity stars to even possibly mini-quasars whose contribution to the CIB is driven by energy released by gravitational accretion, as opposed to stellar nucleosynthesis.  Extensive numerical investigations of collapse and fragmentation of the first objects forming out of density fluctuations specified by the standard $\\Lambda$CDM model suggest that Pop III stars were quite massive and lived at $z> 10$, well within the first Gyr of the Universe's evolution. If predominantly massive, they are expected to have left a significant level of diffuse radiation redshifted today into the IR, and it has been suggested that the CIB contains a detectable contribution from Pop~III in the near-IR, manifest in both its mean level and its anisotropies (e.g. Bond et al 1986, Santos et al 2002, Salvaterra \\& Ferrara 2003, Cooray et al 2004, Kashlinsky et al 2004).  In the past several years a group of us (Kashlinsky, Arendt, Mather \\& Moseley 2005, 2007a,b,c - hereafter KAMM1, KAMM2, KAMM3, KAMM4) have used deep-integration {\\it Spitzer} data to measure the CIB fluctuations component arising from early populations. These provide first observational insights into the global evolution of the Universe at early cosmic epochs. Our measurements revealed significant CIB fluctuations at the IRAC wavelengths (3.6 to 8 $\\mu$m) which remain after removing galaxies down to very faint levels (KAMM1, KAMM2). These fluctuations must arise from populations that have a significant clustering component, but only low levels of the shot noise (KAMM3). Furthermore, there are no correlations between the source-subtracted IRAC maps and the corresponding fields observed with the {\\it HST} ACS at optical wavelengths (KAMM4). Taken together, these data imply that 1) the sources producing the CIB fluctuations have a very different clustering pattern than galaxies at intermediate redshifts and are individually faint with $S_\\nu<$ a few nJy at 3.6 and 4.5 $\\mu$m; 2) are located within the first $\\simeq 0.7$ Gyr (unless these fluctuations can somehow be produced by - so far unobserved - local galaxies of extremely low luminosity and with the unusual for local populations clustering pattern), 3) they produce contribution to the net CIB flux of at least 1-2 nW/m$^2$/sr at 3.6 and 4.5 $\\mu$m and must have mass-to-light ratio significantly below the present-day populations, and 4) their angular density is $\\sim$(a few) arcsec$^{-2}$, so they are in the confusion of the current instruments, but can be individually observable with {\\it JWST}.  Below, I will discuss the latest measurements of the fluctuations in the CIB by our team (Kashlinsky, Arendt, Mather \\& Moseley - KAMM) and explore their implications for the nature of the sources contributing to these anisotrpies, specifically in the era of the first stars. Following the Introduction, Sec. 2 reviews the current measurements at both near-IR (IRAC) and optical (ACS) wavelengths and Sec. 3 discusses the nature of the cosmological populations producing these CIB anisotropies. ",
        "conclusions": ""
    },
    "0801/0801.2933_arXiv.txt": {
        "abstract": "% The opacity of a spiral disk due to dust absorption influences every measurement we make of it in the UV and optical. Two separate techniques directly measure the total absorption by dust in the disk: calibrated distant galaxy counts and overlapping galaxy pairs. The main results from both so far are a semi-transparent disk with more opaque arms, and a relation between surface brightness and disk opacity. In the Spitzer era, SED models of spiral disks add a new perspective on the role of dust in spiral disks. Combined with the overall opacity from galaxy counts, they yield a typical optical depth of the dusty ISM clouds: 0.4 that implies a size of $\\sim$ 60 pc. Work on galaxy counts is currently ongoing on the ACS fields of M51, M101 and M81. Occulting galaxies offer the possibility of probing the history of disk opacity from higher redshift pairs. Evolution in disk opacity could influence distance measurements (SN1a, Tully-Fisher relation). Here, we present first results from spectroscopically selected occulting pairs in the SDSS. The redshift range for this sample is limited, but does offer a first insight into disk opacity evolution as well as a reference for higher redshift measurements. Spiral disk opacity has not undergone significant evolution since z=0.2. HST imaging would help disentangle the effects of spiral arms in these pairs. Many more mixed-morphology types are being identified in SDSS by the GalaxyZoo project.  The occulting galaxy technique can be pushed to a redshift of 1 using many pairs identified in the imaging campaigns with HST (DEEP2, GEMS, GOODS, COSMOS). ",
        "introduction": "The number of distant galaxies seen in an {\\em HST} image through the face-on foreground spiral is a direct indication of its opacity, after proper calibration using artificial galaxy counts \\citep[the ``Synthetic Field Method\", SFM,][]{Gonzalez98,Holwerda05a}. We have obtained calibrated galaxy counts for a sample of 32 deep {\\em HST/WFPC2} fields. The main results from the disk opacity study are: (1) most of the disks are semi-transparent \\citep[][Figure \\ref{f:ra}]{Holwerda05b}, (2) arms are more opaque \\citep{Holwerda05b}, (3) as are brighter sections of the disk \\citep{Holwerda05d}. We did not find a relation between HI profiles and disk opacity but this can be better explored on a single disk \\citep{Holwerda05c}. The optimal distance of the foreground disk is $\\sim$ 10 Mpc, the compromise between crowding effects and solid angle constraints \\citep{Gonzalez03,Holwerda05e}.   \\begin{figure} \\centering \\includegraphics[width=\\textwidth]{Holwerda_p42_f2.eps}% \\caption{The relation between the optical depth from the SED model [$\\tau_m$(SED)], and the observed optical depth from the number of distant galaxies [$\\tau$(SFM)]. Model C from \\protect\\cite{Natta84} is fit through these. Cloud optical depth is 0.4, and  more than a single cloud along the line-of-sight is needed, especially for optically thick disks. Figure from \\protect\\cite{Holwerda07a}.} \\label{f:tc} \\end{figure}    The sample from \\cite{Holwerda05b} has an overlap of 12 galaxies with {\\em SINGS}; we compared the disk opacity from number of distant galaxies to the dust surface density obtained from an SED model \\citep{Draine07a}. Expressed as optical depths, the relation between the {\\it observed} optical depth ($\\tilde{\\tau}$) and the optical depth inferred from the SED dust mass ($\\tau_m$) is solely a function of the typical cloud optical depth ($\\tau_c$): $\\tilde{\\tau}/\\tau_m = (1-e^{-\\tau_c})/\\tau_c$ \\citep{Natta84}. Figure \\ref{f:tc} shows the values for $\\tilde{\\tau}$ and $\\tau_m$ and the fit of $\\tau_c$=0.4. This implies that most of the disks are made up of small, optically thin, cold ISM structures with more than one cloud along the line-of-sight, especially in optically thick disks \\citep{Holwerda07a}. From their distribution in the E[I-L] color map, it becomes clear that the number of distant galaxies does not drop exclusively as a result of grand spiral opaque structures but that the unresolved dusty ISM disk is equally important \\citep{Holwerda07b}.    \\begin{figure} \\centering \\includegraphics[width=\\textwidth]{Holwerda_p42_f3.ps}% \\caption{\\label{f:raz}Radial opacity profile for the local (lower), z=0-0.1 (middle) and z=0.1-0.2 (top) spirals from occulting pairs. Open circles are disk sections for the local pairs. Figure from \\protect\\cite{Holwerda07c}.} \\end{figure} ",
        "conclusions": ""
    },
    "0801/0801.2428_arXiv.txt": {
        "abstract": "We investigate the linear growth and vertical structure of the magnetorotational instability (MRI) in weakly ionised, stratified protoplanetary discs. The magnetic field is initially vertical and dust grains are assumed to be well mixed with the gas over the entire vertical dimension of the disc. For simplicity, all the grains are assumed to have the same radius ($a = 0.1$, $1$ or $3\\  \\mu$m) and constitute a constant fraction (1 \\%) of the total mass of the gas. Solutions are obtained at representative radial locations ($R = 5$ and $10 $ AU) from the central protostar for a minimum-mass solar nebula model and different choices of the initial magnetic field strength, configuration of the diffusivity tensor and grain sizes.  We find that when no grain are present, or they are $\\gtrsim 1\\ \\mu$m in radius, the midplane of the disc remains magnetically coupled for field strengths up to a few gauss at both radii. In contrast, when a population of small grains ($a = 0.1 \\mu$m) is mixed with the gas, the section of the disc within two tidal scaleheights from the midplane is magnetically inactive and only magnetic fields weaker than $\\sim 50$ mG can effectively couple to the fluid. At 5 AU, Ohmic diffusion dominates for $z/H \\lesssim 1$ when the field is relatively weak ($B \\lesssim$ a few milligauss), irrespective of the properties of the grain population. Conversely, at 10 AU this diffusion term is unimportant in all the scenarios studied here. High above the midplane ($z/H \\gtrsim 5$), ambipolar diffusion is severe and prevents the field from coupling to the gas for all $B$. Hall diffusion is dominant for a wide range of field strengths at both radii when dust grains are present.  The growth rate, wavenumber and range of magnetic field strengths for which MRI-unstable modes exist are all drastically diminished when dust grains are present, particularly when they are small ($a \\sim 0.1 \\mu$m). In fact, MRI perturbations grow at $5$ AU (10 AU) for $B \\lesssim 160$ mG (130 mG) when $3 \\ \\mu$m grains are mixed with the gas. This upper limit on the field strength is reduced to only $\\sim 16$ mG (10 mG) when the grain size is reduced to $0.1 \\ \\mu$m. In contrast, when the grains are assumed to have settled, MRI unstable modes are found for $B \\lesssim 800$ mG at 5 AU and 250 mG at 10 AU (Salmeron \\& Wardle 2005). Similarly, as the typical size of the dust grains diminishes, the vertical extent of the dead zone increases, as expected. For $0.1 \\ \\mu$m grains, the disk is magnetically inactive within two scaleheights of the midplane at both radii, but perturbations grow over the entire section of the disk for grain sizes of $1 \\ \\mu$m or larger. When dust grains are mixed with the gas, perturbations that incorporate Hall diffusion grow faster, and are active over a more extended cross section of the disc,  than those obtained under the ambipolar diffusion approximation.  We conclude that in protoplanetary discs, the magnetic field is able to couple to the gas and shear over a wide range of fluid conditions even when small dust grains are well mixed with the gas. Despite the low magnetic coupling, MRI modes grow for an extended range of magnetic field strengths and Hall diffusion largely determines the properties of the perturbations in the inner regions of the disc. ",
        "introduction": "\\label{sec4:intro} Magnetic fields may regulate the `disc accretion phase' of star formation by providing means of transporting away the excess angular momentum of the disc, enabling matter to accrete. The more generally relevant mechanisms associated with this are MHD turbulence induced by the magnetorotational instability (MRI; Balbus \\& Hawley 1991, 1998) and outflows driven centrifugally from the disc surfaces (Blandford \\& Payne 1982, Wardle \\& K\\\"onigl 1993; see also the review by K\\\"onigl \\& Pudritz 2000). These processes are, in turn, thought to play key roles in the dynamics and evolution of astrophysical accretion discs. Magnetically driven turbulence is likely to impact disc chemistry  (e.g. Semenov, Wiebe \\& Henning 2006, Ilgner \\& Nelson 2006) as well as the properties -- and evolution -- of dust grains mixed with the gas (e.g. Turner et al. 2006). Magnetic resonances have been shown to modify the net tidal torque exerted by the disc on a forming planet and thus, alter the speed and direction of the planet's migration through the disc (Terquem 2003; Johnson, Goodman \\& Menou 2006). Finally, the magnetically inactive (dead) zones (Gammie 1996; Wardle 1997)  are not only the regions where planets are thought to form, but they have also been invoked  as a possible mechanism to stop their inward migration (e.g. Matsumura \\& Pudritz 2005).  In protoplanetary discs, however, the magnetic diffusivity can be high enough to limit -- or even suppress -- these processes. The specific role magnetic fields are able to play in these environments is, therefore, largely determined by the degree of coupling between the field and the neutral gas. A critical parameter for this analysis is the ionisation fraction of the fluid, which reflects the equilibrium between ionisation and recombination processes taking place in the disc. Ionisation processes outside the disc innermost sections ($R \\gtrsim 0.1$ AU) are non-thermal, driven by interstellar cosmic rays, X-rays emitted by the magnetically active protostar and radioactive decay (Hayashi 1981; Glassgold, Najita \\& Igea 1997; Igea \\& Glassgold 1999; Fromang, Terquem \\& Balbus 2002). On the other hand, free electrons are lost through recombination processes which, in general, take place both in the gas phase and on grain surfaces (e.g. Nishi, Nakano \\& Umebayashi 1991).  Dust grains affect the level of magnetic coupling in protoplanetary discs when they are well mixed with the gas (e.g. in relatively early stages of accretion and/or when turbulence prevents them from settling towards the midplane). They do so in two ways. First, they reduce the ionisation fraction by providing additional pathways for electrons and ions to recombine on their surfaces. Second, charged dust particles can become important species in high density regions (Umebayashi \\& Nakano 1990; Nishi, Nakano \\& Umebayashi 1991). For example, at $1$ AU in a disc where $0.1 \\ \\mu$m grains are present, positively charged particles are the most abundant ionised species within two scaleheights from the midplane (see Fig.~6 of Wardle 2007; hereafter W07). As these particles generally have large cross sections, collisions with the neutrals are important and they become decoupled (or partially decoupled) to the magnetic field at densities for which smaller species, typically ions and electrons, would still be well tied to it.  Both of these mechanisms act to lower the conductivity of the fluid, especially near the midplane where the density is high and ionisation processes are inefficient. In the minimum-mass solar nebula disc (Hayashi 1981, Hayashi et el. 1985), for example, X-rays are completely attenuated below $z/H \\sim 1.7$ (see Fig.~$1$ of Salmeron \\& Wardle 2005; hereafter SW05). As a result, in the disc inner sections the magnetic coupling may be insufficient to provide adequate link between the neutral particles and the field. Moreover, recent calculations for a minimum-mass solar nebula disk exposed to X-ray and cosmic ray ionising fluxes (W07) indicate that near the surface (above 3 - 4 tidal scaleheights) the magnetic diffusivity can also be severe, even though in these regions the ionising flux is strongest and the electron fraction is significantly larger than it is in the disc interior. This effect results from the strong decline in the number density of charged particles high above the midplane. These considerations suggest that magnetic activity in the inner regions of weakly ionised discs may well be confined to intermediate heights above the midplane ($z/H \\sim 2 - 4 $).  It is clear that the level of magnetic diffusion is strongly dependent on the presence, and size distribution, of dust particles suspended in the gas phase. In fact, once dust grains have settled, the ionisation fraction may be enough to produce adequate magnetic coupling over the entire vertical extension of the disk even at $R \\lesssim 1$ AU (W07). A realistic study of the properties of the MRI in these discs must, therefore, incorporate a consistent treatment of dust dynamics and evolution (unless they are assumed to have settled, a good approximation to model relatively late accretion stages, as was the case in SW03 and SW05). This analysis is further complicated because dust grains have complex spatial and size distributions (e.g. Mathis, Rumpl \\& Nordsieck 1977, Umebayashi \\& Nakano 1990, D'Alessio et al. 2006), determined by the competing action of processes involving sticking, shattering, coagulation, growth (and/or sublimation) of ice mantles and settling to the midplane (e.g. Weidenschilling \\& Cuzzi 1993). Previous results (SW03, SW05) highlight also the importance of incorporating in these studies all three diffusion mechanisms between the magnetic field and the neutral fluid (namely, the Ohmic, Hall and Ambipolar diffusivities), as Hall diffusion largely determines the growth and structure of MRI perturbations, particularly in the disc inner regions (e.g. within distances of the order of a few AU from the central protostar; SW05). This is implemented here via a diffusivity tensor formalism (Cowling 1957, Norman \\& Heyvaertz 1985, Nakano \\& Umebayashi 1986, Wardle \\& Ng 1999, W07).  Dust grains can affect the structure and dynamics of accretion discs via two additional mechanisms: Dust opacity can modify the radiative transfer within the disc -- which, in turn, can dramatically alter its structure -- and dust particles may become dynamically important if their abundance is sufficiently high. In this study both effects are small because the disc is vertically isothermal and grains constitute only a small fraction of the mass of the gas (see below).   In this paper we study the vertical structure and linear growth of the MRI in a disc where dust grains are well mixed with the gas over its entire vertical dimension. Results are presented for two representative distances ($R = 5$ and $10$ AU) from the central protostar. For simplicity, we assume that all particles have the same radius -- $a = 0.1$, $1$ or $3 \\ \\mu$m -- and constitute $1$\\% of the total mass of the gas, a typical assumption in studies of molecular clouds (Umebayashi \\& Nakano 1990). This fraction is constant with height, which means that we have also assumed that no sedimentation has occurred. Although this is a very simplified picture, the results illustrate the importance of dust particles in the delicate ionisation equilibrium of discs, and consequently, on their magnetic activity.  The paper is organised as follows. The adopted disc model is described in section \\ref{sec4:cond}. Section \\ref{sec4:formulation} briefly summarises the formulation and methodology, which are based on SW05. We refer the reader to that study -- and references therein -- for further details. This section also includes a discussion of the typical dependence of the components of the diffusivity tensor and magnetic coupling with height with (and without) grains for the two radial positions of interest. Section \\ref{sec4:mri} then presents the vertical structure and linear growth of unstable MRI modes at these radii and compares solutions incorporating different diffusion mechanisms and assumptions regarding the presence, and size, of dust grains. These results, and possible implications for the dynamics and evolution of low conductivity discs, are discussed in section \\ref{sec4:discussion}. Our main conclusions are summarised in section \\ref{sec:summary}. ",
        "conclusions": "\\label{sec4:discussion}  In this paper we have examined illustrative examples of the impact of dust grains in the magnetic activity of protoplanetary discs and, in particular, in the linear growth and vertical structure of MRI perturbations. Solutions were computed for $R = 10$ and $5$ AU assuming a single size grain population of radius $a = 0.1$, 1 or 3 $\\mu$m, and constituting 1 \\% of the total mass of the gas, is well mixed with the gas phase over the entire vertical extent of the disc. This fraction is independent of height, so we have also assumed that the grains have not sedimented towards the midplane. Our results indicate that the perturbations' wavenumber and growth rate are significantly reduced when grains are present. Furthermore, the magnetically inactive -- dead -- zone, which was practically non-existent when grains were settled, extends to $z/H \\sim 3$ at either radii when 0.1 $\\mu$m-sized grains are considered. At 10 AU (5 AU), unstable perturbations were found in this case for $B \\lesssim 10$ mG (16 mG), a much reduced range compared with the strengths for which they exist when no grains are involved (250 and 795 mG, respectively). This maximum field strength corresponds well to the equipartition field at % the height at which the perturbations peak, as expected (e.g. $z/H \\approx 3.7$ at 10 AU; see lower right panel of Fig.~\\ref{fig:grains1}).  These results illustrate the impact of dust particles on the dynamics -- and evolution -- of low conductivity discs. They can also be used to estimate the maximum magnetic field strength to support magnetic activity at $1$ AU. The magnetic coupling at this radius, in a disc including $0.1\\  \\mu$m grains, is too weak below $z/H \\sim 3.5$ to allow the field to sufficiently couple with the gas (W07). Assuming that the maximum field strength for the MRI to grow is also of the order of the equipartition field at about this height, we can roughly estimate that the MRI should be active for $B \\lesssim 400$ mG at $1$ AU. This is a smaller range than the several gauss for which unstable modes exist when dust grains are not present (SW05). However, MRI modes still grow in this case for a wide range of field strengths.  The results just presented were obtained assuming that all grains have the same size and are well mixed with the gas at all $z$. More realistic spatial, and size, distributions must incorporate the effects of dust dynamics and evolution within the disc. Observations of the mid -- and far -- infrared spectra of discs have provided credible indications that dust properties in discs are indeed different from those of particles in diffuse clouds (e.g. Van Dishoeck 2004, D'Alessio et al. 2006 and references therein). Two aspects of this dust evolution have been clearly identified. First, dust grains coagulate from $0.1 \\ \\mu$m to $\\sim 1$ mm particles. Second, (silicate) material becomes crystallised. It is believed that this crystallisation occurs in the disc, given that crystalline silicates are absent from the interstellar medium. Furthermore, the presence of this material at radial locations where the temperature is too low to produce them, suggests significant radial mixing takes place as well (e.g. Van Dishoeck 2004 and references therein). Finally, simulations of dust dynamics and evolution also suggest that in quiescent environments, the grains tend to settle and agglomerate into bigger particles (e.g. Weidenschilling \\& Cuzzi 1993) and efficiently coagulate and grow icy mantles (Ossenkopf 1993). All these processes modify the surface area of dust grains and impact the recombination rate on their surfaces and the way they drift in response to magnetic stresses. It is also expected that a residual population of small grains would remain (e.g. Nomura et al. 2006) even when the mean grain size maybe relatively large ($a \\gtrsim 1 \\mu$m). This is an important consideration because these small grains tend to carry a significant fraction of the grain charge (S07). The effect of this settling in the spectral energy distribution (and optical appearance) of protostellar discs has been investigated in recent studies (e.g. Dullemond \\& Dominik 2004).  How quickly, and to what height, dust particles are able to settle is an important, and largely unanswered, question. According to Nakagawa, Nakazawa \\& Hayashi (1981), the mass fraction of $\\sim 1$ - $10 \\ \\mu$m grains well mixed with the gas, diminishes from $\\sim 10^{-1}$ to $10^{-4}$ in a timescale of about $2 \\ee 3 $ to $10^{5}$ years. Moreover, although the timescale for dust grains to sediment  all the way to the midplane may exceed the lifetime of the disc, they may be able to settle within a few scaleheights from the midplane in a shorter timescale (Dullemond \\& Dominik 2004). This is complicated even more by the expectation that the transition between sections were dust grains are well mixed with the gas, and those completely depleted of them, occurs gradually (Dullemond \\& Dominik 2004).   MHD turbulence may itself be an important factor for the settling of dust particles. It may, in particular, produce sufficient vertical stirring to prevent settling below a certain height (Dullemond \\& Dominik 2004, Carballido et al. 2005, Turner et al. 2006). However, this is contingent on the disc being able to generate and sustain MHD turbulence in the vertical sections where the dust is present. This is not guaranteed, even if turbulence exists in other regions, as dust grains efficiently reduce the ionisation fraction (and magnetic coupling) of the gas. As a result, the efficiency -- and even the viability -- of MHD turbulence in the presence of dust grains, is an important topic that merits careful investigation.    We have explored in this paper the linear growth and vertical structure of MRI unstable modes at two representative radii ($R = 5$ and 10 AU) in protoplanetary discs.  Dust grains are assumed to be well mixed with the fluid over the entire section of the disc and, for simplicity, are taken to have the same radius ($a = 3$, 1 or 0.1 $\\mu$m). They constitute a constant fraction (1\\%) of the total mass of the gas. These solutions are compared with those arrived at assuming that the grains have settled to the midplane of the disc (SW05). We have also explored which disc sections are expected to be magnetically coupled and the dominant diffusion mechanism as a function of height and the strength of the magnetic field, which is initially vertical.  Our models use a minimum-mass solar nebula disc (Hayashi 1981, Hayashi et al. 1985) and incorporate all three diffusion mechanisms between the magnetic field and the neutral gas: Ohmic, Hall and Ambipolar. The diffusivity components are a function of height (as is the density) and are obtained using the method described in W07, to which we refer the reader for details. Essentially, this formalism uses a chemical reaction scheme similar to that of Nishi, Nakano \\& Umebayashi (1991), but it incorporates higher dust-grain charge states that are likely to occur in discs on account of their larger gas density and temperature in relation to those of molecular clouds. Our calculations also include a realistic ionisation profile, with the main ionising sources being cosmic rays, X-rays and (to a lesser extent) radioactive decay.  Solutions were obtained at the two radii of interest for different grain sizes and configurations of the diffusivity tensor as a function of the magnetic field strength. We refer the reader to SW03 and SW05 for further details of the integration procedure. The main findings of this study are detailed below.  \\subsubsection*{Magnetic diffusivity} \\begin{enumerate} \\item When no grain are present, or they are $\\gtrsim 1\\ \\mu$m in radius, the midplane of the disc remains magnetically coupled for field strengths up to a few gauss at both radii. \\item In contrast, when a population of small grains ($a = 0.1 \\mu$m) is mixed with the gas, the section of the disc below $z/H \\sim 2$ ($z/H \\sim 2.5$) is magnetically inactive at $R = 10$ AU (5 AU). Only magnetic fields weaker than 25 mG (50 mG) can couple to the gas. \\item At 5 AU, Ohmic diffusion dominates for $z/H \\lesssim 1.2$ when the field is relatively weak ($B \\lesssim$ a few milligauss), irrespective of the properties of the grain population. Conversely, at 10 AU this diffusion term is unimportant in all the scenarios studied here. \\item High above the midplane ($z/H \\gtrsim 4.5 - 5$, depending on the specific model), ambipolar diffusion is severe and prevents the field from coupling to the gas for all $B$. This is consistent with previous results by W07. \\item Hall diffusion is dominant for a wide range of field strengths and grain sizes at both radii (see Figs.~\\ref{fig:contour_10AU} and \\ref{fig:contour_5AU}). \\end{enumerate}  \\subsubsection*{Magnetorotational instability} \\begin{enumerate} \\item The growth rate, wavenumber and range of magnetic field strengths for which unstable modes exist are all drastically diminished when dust grains are present, particularly when they are small ($a \\sim 0.1 \\ \\mu$m; see Figs.~\\ref{fig:grains1} and \\ref{fig:grains1_1}). \\item In all cases that involve dust grains, perturbations that incorporate Hall diffusion grow faster than those obtained under the ambipolar diffusion approximation. \\item At 10 AU, unstable MRI modes grow for $B \\lesssim 80$ mG (10 mG) when the grain size is $1 \\ \\mu$m ($0.1 \\ \\mu$m), a much reduced range compared with the $\\sim 250$ mG for which they exist when ions and electrons are the only charge carriers (SW05). When the grains are relatively large ($a = 1$ and $3 \\ \\mu$m), Hall diffusion controls the structure the modes, which grow even at the midplane. In contrast, when the grains are small ($a = 0.1 \\ \\mu$m), the perturbations grow only for $z/H \\gtrsim 3$ and are shaped mainly by ambipolar diffusion. \\item At 5 AU, MRI perturbations exist for $B \\lesssim 80$ mG (16 mG) when the grains are $1 \\ \\mu$m ($0.1 \\ \\mu$m) in size. For comparison, the upper limit when no grains are present is $\\sim 800$ mG (SW05). These modes are shaped largely by ambipolar diffusion, the dominant mechanism at the height where they peak. \\end{enumerate}  We conclude that in protoplanetary discs, the magnetic field is able to couple to the gas and shear over a wide range of fluid conditions even when small dust grains are well mixed with the gas. Despite the low magnetic coupling, MRI modes grow for an extended range of magnetic field strengths and Hall diffusion largely determines the properties of the perturbations in the inner regions of the disc."
    },
    "0801/0801.1700_arXiv.txt": {
        "abstract": "We present infrared observations of 66 starburst galaxies over a wide range of oxygen abundances, to measure how metallicity affects their dust properties. The data include imaging and spectroscopy from the {\\it Spitzer} Space Telescope, supplemented by groundbased near-infrared imaging.  We confirm a strong correlation of aromatic emission with metallicity, with a threshold at a metallicity [$12+{\\rm log(O/H)}$] $\\sim$ 8.  The large scatter in both the metallicity and radiation hardness dependence of this behavior implies that it is not due to a single effect, but to some combination. We show that the far-infrared color temperature of the large dust grains increases towards lower metallicity, peaking at a metallicity of 8 before turning over.  We compute dust masses and compare them to HI masses from the literature to derive the gas to dust ratio, which increases by nearly 3 orders of magnitude between solar metallicity and a metallicity of 8, below which it flattens out. The abrupt change in aromatic emission at mid-infrared wavelengths thus appears to be reflected in the far-infrared properties, indicating that metallicity changes affect the composition of the full range of dust grain sizes that dominate the infrared emission. In addition, we find that the ratio L(8~\\micron)/L(TIR), important for calibrating 24~\\micron\\ measurements of high redshift galaxies, increases slightly as the metallicity decreases from $\\sim$ solar to $\\sim$ 50\\% of solar, and then decreases by an order of magnitude with further decreases in metallicity.  Although the great majority of galaxies show similar patterns of behavior as described above, there are three exceptions, SBS~0335-052E, Haro~11, and SHOC~391. Their infrared SEDs are dominated energetically by the mid-IR near 24~\\micron\\ rather than by the $60 - 200$~\\micron\\ region. In addition, they have very weak near infrared outputs and their SEDs are dominated by emission by dust at wavelengths as short as 1.8~\\micron. The latter behavior indicates that the dominant star forming episodes in them are extremely young. The component of the ISM responsible for the usual far infrared emission appears either to be missing, or inefficiently heated, in these three galaxies. ",
        "introduction": "\\label{sec:introduction}  One of the early discoveries of infrared astronomy was the substantial infrared output of star forming galaxies \\citep{kleinmann70}.  It quickly became apparent that the infrared emission was ubiquitous and represented a substantial portion of the bolometric output associated with recent star formation, since most of the ultraviolet and optical emission of the young stars is absorbed and re-emitted in the infrared \\citep{rieke72,rieke79,soifer87}.  Since then, this phenomenon has been studied extensively (as summarized in many reviews, e.g., \\citet{telesco88,moorwood96,sanders96, kennicutt98, genzel00, sauvage05, lagache05}). To first order, this work has established that the overall properties of the best-studied star-forming and infrared-luminous galaxies are dependent on two parameters: age and luminosity.  Thus, the spectral energy distributions have been fitted with templates that vary primarily with luminosity (e.g., \\citet{devriendt99,chary01,dale02}), and stellar population synthesis models demonstrate how the characteristics of the galaxy evolve with the age of the dominant star-forming episode (e.g., \\citet{rieke93, leitherer95, engelbracht98}).  The initial studies of infrared galaxies were largely confined to relatively luminous, massive, and hence metal-rich examples. Indications that metal-poor low-luminosity galaxies might behave differently were found in the extensive mid-infrared spectroscopic survey by \\citet{roche91}. With improvements in sensitivity and sophistication of mid and far infrared instrumentation, it has become apparent that metallicity - and the accompanying changes in the hardness of the UV radiation field - constitutes a third critical parameter influencing the overall infrared properties of star-forming galaxies. The most dramatic dependence applies to the mid-IR aromatic features (sometimes termed \"PAH\" features).  In high metallicity galaxies, these features are quite similar in strength and other properties \\citep{roche91,dale07}.  Spectra obtained with the Infrared Space Observatory (ISO) showed that they are significantly weaker in a number of low-metallicity galaxies \\citep{thuan99, madden00, galliano03, galliano05}.  Observations with {\\it Spitzer} have confirmed this trend and indicated that there is a fairly sharp transition in the relative aromatic feature strength near a metallicity of $12 + {\\rm log(O/H)} = 8$ \\citep{houck04,engelbracht05, wu06, madden06}.  Is the change in aromatic feature characteristics a direct result of the low metallicity, or is the fundamental effect due to the increased hardness of the interstellar UV radiation field in low metallicity environments?  Is the change confined to the aromatic carriers, or do other components of the interstellar dust undergo changes also?  With the sensitivity of {\\it Spitzer}, it is possible to explore these questions. In this paper, we describe infrared measurements of a sample of 66 star-forming galaxies over a very wide range of oxygen abundances (which we treat as measurements of the total metallicities), $7.1 \\le 12 + {\\rm log (O/H)} \\le 8.85$. We find systematic changes in the behavior of the far-infrared color temperatures and in the dust masses near the same metallicity where the aromatic feature strength appears to change. Thus, near this metallicity, there may be a general change in dust properties over the full range of sizes and compositions that contribute to the observed infrared emission.  In addition to providing insights to the composition and behavior of the interstellar dust, an understanding of galaxy behavior with metallicity is necessary to interpret infrared observations of galaxies at high redshift, where their metallicities are observed to be lower (by up to a factor of order two) than locally \\citep{liang04, mouhcine06, rupke07}.  Since our study includes this range in the overall context of infrared behavior with metallicity, we can compare with the observed infrared properties at high redshift. We find evidence for an increase in L(8~\\micron)/L(TIR) with decreasing metallicity down to $\\sim$ half solar. This effect would result in overestimates of L(TIR) based on observations at 24~\\micron\\ of galaxies at z $\\sim$ 2, consistent with the bias reported by Rigby et al.\\ (2007) in studies of luminous galaxies near this redshift.  Our discussion begins in \\S\\ref{sec:data} by describing the sample, how the data were reduced, and how the photometric and spectroscopic measurements were made.  We describe basic properties of the galaxies in \\S\\ref{sec:results}, including spectral energy distributions (SEDs) and luminosities. \\S\\ref{sec:metaleffects} describes the effects of metallicity on the galaxy infrared properties. The paper is summarized in \\S\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We present new infrared images and spectra, both from the {\\it Spitzer} Space Telescope and from the ground, for a sample of 66 star-forming galaxies.  The sample spans a wide range of metallicities, from the lowest known in a star-forming galaxy to near solar.  The imaging covers the range from 1~\\micron\\ to 180~\\micron, while the spectra cover 5~\\micron\\ to 40~\\micron. These observations represent the first detections of dust emission from some of the lowest metallicity star-forming galaxies known, including the current record holder SBS~0335-052W.  We perform photometry on the images to compute spectral energy distributions (SEDs).  We demonstrate that, with a few exceptions, the SEDs of the galaxies are very similar in the near infrared, where they are dominated by stellar emission, and the far infrared, where they are dominated by emission from dust in thermal equilibrium.  The transition from stellar to dust emission occurs around 4.5~\\micron, with little scatter for galaxies with a metallicity [$12 + \\rm {log(O/H)}$] above 8.  The scatter in this transition wavelength increases considerably at low metallicities.  The SEDs exhibit a strong metallicity dependence in the mid infrared, largely due to changes in the strength of the aromatic features.  We use a variety of simple models to derive the fraction of emission due to stars in the mid infrared, particularly in the IRAC bands at 4.5~\\micron, 5.8~\\micron, and 8.0~\\micron.  While this fraction is a relatively constant 70\\% at 4.5~\\micron, it has a strong dependence on metallicity at 8.0~\\micron, where it ranges from 4\\% in metal-rich galaxies to 30\\% in the lowest-metallicity galaxies, reflecting the lower dust content and weak aromatics in these galaxies.  We confirm previous evidence for a substantial reduction in aromatic feature strengths below $12 + {\\rm log(O/H)} \\sim 8.2$, and also with increased radiation hardness. The scatter in the aromatic EWs against both metallicity and radiation hardness implies that this reduction is not controlled by a single parameter, but probably by some combination of effects.  We compute dust properties (temperature and mass) for each galaxy.  We find an anticorrelation between dust temperature and metallicity, with equilibrium dust temperatures of $\\sim23$~K near solar metallicity up to 40~K at low metallicity $\\sim$ 8, and then falling temperatures with further reductions in metallicity.  The derived dust masses span over 8 orders of magnitude, from one-tenth of a solar mass to over 50~million solar masses. They exhibit a very steep dependence on metallicity, as $\\sim$ Z$^{-2.5}$ down to Z $\\sim$ 8 but have a much weaker dependence for Z $<$ 8. The change in dust behavior in terms of aromatics, far infrared color temperature, and dust/gas mass ratio all near Z = 8 indicates that there near this metallicity there is a general modification of all components of the interstellar dust that dominate the infrared emission.  We show that the ratio L(8~\\micron)/L(TIR) increases with decreasing metallicity from solar to about 50\\% solar and then decreases with further reductions in Z.  This behavior is important for interpretation of 24~\\micron\\ measurements of star forming galaxies at redshifts z $\\sim$ 2, since the signals for them are dominated by aromatic emission and it is likely that they have lower metallicity than is typical of local template galaxies.  We find 3 galaxies, SBS~0335-052E, Haro~11, and SHOC~391, that have anomalous far infrared spectral energy distributions, with weak emission near 70~\\micron\\ and an SED that is dominated energetically by the mid-IR near 24~\\micron.  In addition, they have weak stellar outputs in the near infrared and are dominated by dust emission down to wavelengths as short as 2~\\micron. This latter behavior indicates that they are the sites of very young dominant star forming episodes.  Their metallicities tend to be low, but not different from other galaxies with behavior much more similar to the majority of infrared-active galaxies.  It appears that the dust responsible for the far infrared emission in most galaxies is either absent, or not efficiently heated in these three objects. They are interesting targets for further study since they represent an extreme state of the ISM."
    },
    "0801/0801.1193_arXiv.txt": {
        "abstract": "{Galaxy metallicities have been measured to redshift z$\\sim$2 by gas-phase oxygen abundances of the interstellar medium using the R$_{23}$ and N2 methods. Galaxy stellar metallicities provide crucial data for chemical evolution models but have not been assessed reliably much outside the local Universe.} {We determine the iron-abundance, stellar metallicity of star-forming galaxies at redshift z$\\sim$2, homogeneously-selected and observed as part of the Galaxy Mass Assembly ultra-deep Spectroscopic Survey (GMASS).} {We compute the equivalent width (EW) of a rest-frame mid-ultraviolet (mid-UV), photospheric absorption-line index, the 1978~\\AA\\ index, found to vary monotonically with stellar metallicity by Rix, Pettini and collaborators (R04), in model star-forming galaxy (SFG) spectra created using the theoretical massive star models of Pauldrach and coworkers, and the evolutionary population synthesis code Starburst99. The 1978~\\AA\\ index is sensitive to Fe III transitions and measures the iron-abundance, stellar metallicity. To accurately determine the 1978~\\AA\\ index EW, we normalise and combine 75 SFG spectra from the GMASS survey to produce a spectrum corresponding to a total integration time 1652.5 hours (and a signal-to-noise ratio $\\sim$100 for our~1.5~\\AA~binning) of FORS2 spectroscopic observations at the Very Large Telescope.} {We measure a iron-abundance, stellar metallicity of log$(Z/Z_\\odot) = -0.574\\pm0.159$ for our spectrum representative of a galaxy of stellar mass 9.4$\\times$10$^{9}$~M$_\\odot$ assuming a Chabrier initial mass function (IMF). We find that the R04 model SFG spectrum for log$(Z/Z_\\odot) = -0.699$ solar metallicity provides the best description of our GMASS coadded spectrum. For similar galaxy stellar mass, our stellar metallicity is $\\sim0.25$ dex lower than the oxygen-abundance, gas-phase metallicity quantified by Erb and collaborators (E06) for UV-selected star-forming galaxies at $z=2$.} {We measure the iron-abundance, stellar metallicity of star-forming galaxies at redshift z$\\sim$2 by analysing the 1978~\\AA\\ index in a spectrum created by combining 75 galaxy spectra from the GMASS survey. We find that our measurement is $\\sim$0.25 dex lower than the oxygen-abundance gas-phase metallicity at similar values of galaxy stellar mass. We conclude that we are witnessing the establishment of a light-element overabundance in galaxies as they are being formed at redshift z$\\sim$2. Our measurements are indeed reminiscent of the $\\alpha$-element enhancement seen in the likely progenitors of these starburst galaxies at low-redshift, i.e. galactic bulges and early-type galaxies.} ",
        "introduction": "Galaxies are believed to form by the merging of dark matter haloes, the isolated collapse of a gas cloud and subsequent gas accretion, or a combination of these scenarios. Photometric and spectroscopic surveys measure the rate of galaxy mass assembly and the variation in galaxy star formation rates with redshift \\citep{lil96,mad96,mau04,cim04}.  Galaxies are unlikely to evolve as closed boxes but exchange gas, metals and stars with their environment. Following the early work of \\cite{lar74} cosmological models propose that mass outflows and winds driven by supernovae and AGN activity suppress star formation (via so-called ``feedback'') increasingly effectively in galaxies of increasingly lower mass, and enrich the interstellar medium (ISM) with products of stellar nucleosynthesis \\protect\\citep{spr03,ben03,del04,cro06,bow06,brt07,fin07}. Such winds are required for example to prevent the overproduction of galactic low-mass stars and describe the observed metal-enrichment history of the Universe. In the local Universe and to a redshift of z$\\sim$2 a ``mass-metallicity relation'' is observed such that more massive star-forming galaxies are more metal-rich~\\citep{leq79,tre04,sav05,erb06,lee06}. \\cite{tre04} (hereinafter T04) found a tight correlation between the gas-phase metallicity and stellar mass of 53,400 star-forming galaxies of redshift 0.005$ < $z$ < $0.25, using imaging and spectroscopic data from the Sloan Digital Sky Survey (SDSS). Their relation has a 1$\\sigma$ scatter from the median of only 0.1 dex and is linear for stellar masses 8.5 $<$ log(M/M$_\\odot$) $<$ 10.5. At higher redshift~\\citet{fin07} were able to more accurately reproduce the gas phase mass-metallicity relation at z$\\sim$2~using ``momentum-driven winds'' \\citep{mur05,opp06} instead of constant velocity winds traditionally proposed. ``Momentum-driven winds'' remove metals into the intergalactic medium by radiation pressure as photons are absorbed and scattered by dust, and supernovae and galactic winds. Unlike thermal energy, momentum cannot be radiated away and ``momentum-driven winds'' can transport metals to significant distances outside the galaxy.  Measuring galaxy metallicity as a function of redshift probes the rate of chemical enrichment. In the local Universe metallicity is measured using both spectroscopic absorption and emission lines. \\cite{gal05} found good agreement between stellar and gas-phase metallicities using SDSS galaxy spectra for which both absorption and emission lines were measurable. Faber and collaborators at Lick Observatory pioneered measurement of luminosity-weighted stellar age and metallicity of nearby early-type galaxies by developing the ``Lick system'' of spectroscopic absorption-line indices \\citep{bur84,fab85,gon93,gor93,guy94,wor97,tra98} for the rest-frame optical wavelength range $\\sim$4000$-$6200\\AA, and the single-burst evolutionary population synthesis models of \\cite{cla94}. The ``Lick system'' was further developed by \\cite{tra00} and Thomas, Maraston and collaborators \\citep{tho03,tho04} to account for the effects of light-element overabundance ratios on Lick/IDS index measurements using respectively the empirical relations of \\cite{tri95}, and the theoretical models of \\cite{mar05} and calibrations of \\cite{kor05}.  At high-redshift emission lines are more readily detected than absorption lines and galaxy metallicity has been measured using primarily the oxygen-abundance, gas-phase metallicity of the ISM of star-forming galaxies, using the R$_{23}$ \\citep{pag79} and N2 \\citep{sto94,rai00} methods \\citep{kza99,kko00,cli01,lil03,kob04,sha04,mai05,erb06,lam06,mai06}. At z$>$1 the gas metallicity of the ISM has been constrained using the~CIV$\\lambda$1550~\\citep{meh02} and MgII and FeII absorption lines \\citep{sav04}. Pettini and collaborators have measured the chemical abundances of the ISM using rest-frame UV absorption lines in a spectrum of the gravitationally-lensed Lyman-$\\alpha$ galaxy MS 1512-cB58 at redshift z$\\sim$2.73 \\citep{pet02} and damped Lyman-$\\alpha$ systems \\citep{pet02a}. Stellar metallicity data are required to fully test chemical evolution models. \\cite{dem04} provided a first constraint of stellar metallicity for a small sample of very massive ($>10^{11}M_\\odot$) $z\\sim2$ galaxies from the K20 survey \\citep{cim02,ema04}. They showed tentative evidence for solar or supersolar metallicity but with large uncertainties due to limitations in the achieved S/N ratio. \\protect\\cite{rix04} compared the iron-abundance, stellar metallicity for MS 1512-cB58 at z=2.73, and Q1307-BM1163 at z=1.411, with previously determined light-element abundances (magnesium, silicon and sulphur) and iron-abundance gas-phase metallicities of the ISM. For MS 1512-cB58, the iron-abundance, stellar metallicity is a factor of two higher than the light-element abundance gas-phase metallicities. The iron-abundance gas-phase metallicity is a factor of four lower than the light-element abundance metallicities: this is attributed to dust depletion of the ISM and the time delay in release of products of supernovae Ia \\protect\\citep{pet02}. For Q1307-BM1163, the iron-abundance, stellar metallicity is a factor of two higher than the oxygen-abundance, gas-phase metallicity measured using the N2 method. \\cite{erb06} compared two rest-frame mid-UV spectra (obtained for four of the six bins of stellar mass examined in their mass-metallicity relation analysis i.e. two lowest and two highest of six bins of mass) with star-forming galaxy spectra generated using the Starburst99 \\citep{lei99} model. \\cite{erb06} found good agreement between the gas-phase metallicity oxygen abundance determined using the N2 method and the metallicity of the model spectrum that best described the rest-frame mid-UV spectra as judged by eye.  We present the first robust measurement of iron-abundance, stellar metallicity for mid-IR selected star-forming galaxies (SFGs) at z$\\sim$2. We measure a photospheric absorption-line index, defined by \\citet{rix04} (hereinafter R04) in the rest-frame mid-ultraviolet galaxy spectrum (mid-UV) that is dominated by the light of massive stars. R04 modelled star-forming galaxy spectra using the evolutionary population synthesis models Starburst99 (\\citet{lei99}; hereinafter L99) and non-local thermal equilibrium (NLTE) stellar atmosphere models of OB stars \\citep{pau01}. Using stellar atmosphere models, the wavelength coverage of theoretical star-forming galaxy (SFG) models was extended redward to $\\sim$2100~\\AA. Star-forming galaxy model spectra were generated for five values of metallicity, 0.05,~0.2,~0.4,~1.0,~and~2.0 solar, assuming a continuous star formation history and a Salpeter IMF. R04 found that the equivalent width (EW) of an absorption-line system, centred on 1978~\\AA~within an Fe~III transition blend, varied monotonically with metallicity, after 100 Myr of star formation. They defined the EW of the ``1978~\\AA~index'', EW(1978), to have limits of 1935~\\AA~and 2020~\\AA~avoiding the~Al~III $\\lambda$$\\lambda$1855, 1863 interstellar lines, the nebular emission line C~III$]\\lambda$1909 and other weaker lines.  We measure the 1978~\\AA~index in a spectrum created by adding 75 star-forming galaxy spectra from the Galaxy Mass Assembly ultradeep Spectroscopic Survey (GMASS) data corresponding to over 1652.5 hours of integration time with FORS2 at the VLT. Given the difficulty of measurement of the 1978~\\AA~index and the limited range of galaxy stellar masses, we do not analyse spectra as a function of stellar mass. It would indeed be challenging to achieve this measurement for massive galaxies alone, because the most massive galaxies are fainter in the rest-frame mid-UV (e.g. \\citealt{ema04,kon06}) and they are of course rarer (cf.~Figure \\ref{hist}).  In Section \\ref{data} we provide a brief description of the GMASS survey. In Sections \\ref{selection} and \\ref{coadd}, we describe respectively, the selection of, and the combination of GMASS star-forming galaxy spectra. In Section \\ref{metallicity} we outline the measurement of the 1978~\\AA~index and its error, and contrast our observed data with the R04 star-forming galaxy model spectra. In Section \\ref{discuss} the 1978~\\AA~index measurement is compared with predictions of galaxy stellar and gas metallicity from the cosmological simulations of \\citet{fin07}, and the oxygen-abundance gas-phase metallicities of \\citet{erb06} and we briefly discuss our results and conclusions.  We use AB photometric magnitudes, and assume a WMAP cosmology with $\\Omega_\\Lambda, \\Omega_M = 0.73, 0.27$, and $h = $H$_0$[km s$^{-1}$ Mpc$^{-1}$]$/100=0.71$ \\citep{spe03}. We assume a Chabrier initial mass function (IMF) \\protect\\citep{cha03} from 0.1 to 100 $M_\\odot$. Whether the IMF is universal between galaxies and at different redshift is a subject of ongoing debate. The empirical measurement of the Galactic IMF is not well determined and the uncertainty increases for the individual kinematical components (disk, bulge and thick disk) of the Galaxy \\protect\\citep{cha03}, and to high-redshift. For a review of the observational constraints provided by measurement of the mass-to-light ratios of nearby galaxies, the redshift evolution of the Fundamental Plane and galaxy cluster, element-abundances, see \\protect\\cite{ren05}. In this paper, the Chabrier IMF is chosen to provide a reference for all data shown and our conclusions are independent of this choice. ",
        "conclusions": "We have measured the iron-abundance, stellar metallicity of star-forming galaxies at redshift z$\\sim$2, for a spectrum created by adding together 75 star-forming galaxy spectra from the GMASS survey. The iron-abundance, stellar metallicity is determined by measuring the EW of a rest-frame mid-UV, photospheric absorption-line index, the 1978~\\AA\\ index, defined by \\cite{rix04} using the theoretical massive star models of \\cite{pau01}, and the evolutionary population synthesis code Starburst99. We measure an iron-abundance, stellar metallicity of 0.267 $_{-0.082}^{+0.118}$ solar (log$(Z/Z_\\odot) = -0.574\\pm0.159$). This is lower by $\\sim$0.25 dex than the oxygen-abundance, gas-phase metallicity measured by \\citet{erb06} for similar galaxy stellar mass and redshift. At least part of this difference may be the result of different systematic errors between the two estimates although both ourselves and \\cite{erb06} find that systematic errors are very unlikely to reproduce our observed data results. We postulate that our measurements are reminiscent of the $\\alpha$-element enhancement seen in the likely progenitors of these starburst galaxies, i.e. galactic bulges and early-type galaxies."
    },
    "0801/0801.1470_arXiv.txt": {
        "abstract": "We report on time-dependent axisymmetric simulations of an X-ray excited flow from a parsec-scale, rotating, cold torus around an active galactic nucleus. Our simulations account for radiative heating and cooling and radiation pressure force. The simulations follow the development of a broad bi-conical outflow induced mainly by X-ray heating. We compute synthetic spectra predicted by our simulations. The wind characteristics and the spectra support the hypothesis that a rotationally supported torus can serve as the source of a wind which is responsible for the warm absorber gas observed in the X-ray spectra of many Seyfert galaxies. ",
        "introduction": "Approximately $50\\%$ of low-red-shift AGNs exhibit spectra rich in atomic absorption features in the X-ray band. These spectra, referred to as warm absorbers, contain numerous lines and bound-free edges from ions of intermediate-Z elements, and indicate blue-shifts of the absorbing material of $\\simeq 10^2 - 10^3$ km s$^{-1}$. The term refers to the fact that the X-ray absorbing gas has an electron temperature and ionization degree which are intermediate between that expected for neutral and fully ionized material. Warm absorber gas has been known for more than two decades,  beginning with observations using the  {\\it Einstein} Observatory satellite \\cite{Halpern84}, although the limited spectral resolution of early X-ray experiments hampered a detailed study of its properties.  Recent X-ray telescopes, {\\it Chandra} and XMM-{\\it Newton}, provide unprecedented spectral resolution up to $\\sim$10 keV, and have revealed much of the complexity of warm absorber spectra. The ubiquity of X-ray warm absorbers and the implication of outflow in objects which are likely powered by accretion, combine to make warm absorbers of great interest for the understanding of active galactic nuclei.  While observationally the reality of the X-ray absorbing gas is firmly established, the source of this gas remains uncertain. If the wind originates from an accretion disk \\newline \\cite{Konigl94,Murray95}, then warm absorber gas can occupy a vast spatial region from the inner accretion disk up to and beyond the broad line region.  In addition, it is likely to have a speed which is comparable to the escape velocity in the region where the wind originates. If so, the line widths favor a wind originating far outside the inner accretion disk, at  $\\sim 1\\, {\\rm pc}$, for an AGN containing a $10^6 M_\\odot$ black hole.  The natural source of gas at this distance is the molecular torus which is responsible for obscuring the broad line region in Seyfert 2 galaxies, and which is thought to exist in some form in many low and intermediate luminosity AGN \\cite{AntonucciMiller86}. The existence of an outflow from the torus has been suggested by \\cite{KrolikBegelman86, KrolikBegelman88}, and as the source of warm absorber flows by \\cite{KrolikKriss1,KrolikKriss2}.  Direct interferometric observations of this dusty structure have recently become available: mid-infrared high spatial resolution studies of the active nucleus of NGC 1068,  revealed a dust torus - like structure that is 2.1 pc thick and 3.4 pc in diameter ~\\cite{JaffeNATUR}. Evidence of at least two temperature components has been established: the first one is extended and warm (320 K) and conceals the second, which is compact and hot ($>800 K$). It is natural to assume that the latter component can be attributed to the inner part of the torus, heated by the radiation from the compact nucleus. Although such evidence is most solid for nearby active galaxies it is likely that the same obscuring torus paradigm may apply in quasars whose central regions are heavily obscured by gas and dust (Type II quasars) \\cite{Zakamska06}.  The work done so far in modeling torus warm absorber flows employ either one-dimensional models e.g. \\cite{ChelNet05} or two dimensional models which omit rotational forces and also employ an artificial shape for the torus, which is adopted as a boundary for the flow \\cite{BalsKrolik}. These calculations support the idea that the outflow originates  at $r\\sim 1\\,$ pc - the estimate supported both from virial arguments and also deduced from the warm absorber line widths. The content of the warm absorber gas is highly ionized cosmically abundant  heavy elements. This gas can occupy a region from 0.01 pc to $\\sim\\,$10 pc from the nucleus. The column densities are in the range $10^{21} - 10^{24} \\, {\\rm cm}^{-2}$. These estimates are consistent with constraints on the position of the warm absorber gas, which are deduced from the absence of correlated variability of the line equivalent widths to changes of the continuum \\cite{Net03, Beh03}.  The goal of the current work is to show that  a state-of-the-art  hydrodynamic model of the evaporative  torus wind possesses key characteristics of the the warm absorber. To approach this goal, in this paper we present results of models containing the following ingredients: 1) Time-dependent axisymmentric numerical simulations that take into account rotational forces , X-ray heating and cooling and radiation pressure both in spectral lines and continuum; 2) Exploration of the resulting distributions of density, velocity etc.; and 3) A calculation of the synthetic X-ray absorption spectra using the density and velocity distributions as input. Although calculating the spectra is crucial to check the relevance of the particular hydrodynamic solution to the warm absorber phenomenon,  the primary goal of the work reported here is the outflow dynamics . ",
        "conclusions": ""
    },
    "0801/0801.4255_arXiv.txt": {
        "abstract": "The inversion of a gravitational lens system is, as is well known, plagued by the so-called mass-sheet degeneracy: one can always rescale the density distribution of the lens and add a constant-density mass-sheet such that the, also properly rescaled, source plane is projected onto the same observed images. For strong lensing systems, it is often claimed that this degeneracy is broken as soon as two or more sources at different redshifts are available. This is definitely true in the strict sense that it is then impossible to add a constant-density mass-sheet to the rescaled density of the lens without affecting the resulting images. However, often one can easily construct a more general mass distribution -- instead of a constant-density sheet of mass -- which gives rise to the same effect: a uniform scaling of the sources involved without affecting the observed images. We show that this can be achieved by adding one or more circularly symmetric mass distributions, each with its own center of symmetry, to the rescaled mass distribution of the original lens. As it uses circularly symmetric distributions, this procedure can lead to the introduction of ring shaped features in the mass distribution of the lens. In this paper, we show explicitly how degenerate inversions for a given strong lensing system can be constructed. It then becomes clear that many constraints are needed to effectively break this degeneracy. ",
        "introduction": "Being an ill-posed problem, it is not surprising that gravitational lens inversion is plagued by degeneracies. Several classes of degeneracies were first identified by \\citet{Gorenstein} and were later reinterpreted by \\citet{Saha2000} in terms of changes in the arrival-time surface. Of these, the most widely known degeneracy is usually called the mass-sheet degeneracy, although recently the more correct name of steepness degeneracy has been suggested \\citep{SahaSteepness}. This degeneracy was first mentioned in the context of the strong lensing system Q0957+561 \\citep{FalcoMassSheet}, but has also been studied in weak lensing systems (e.g. \\citet{SchneiderMassSheet}) and even in the context of microlensing \\citep{Paczynski}.  In strong lensing systems, it is often claimed that the presence of two sources at different redshifts suffices to break the mass-sheet degeneracy (e.g. \\cite{Abdelsalam2}). While it is definitely true that a constant-density mass-sheet can no longer be used to construct degenerate solutions in such a case, we show in this paper how a more general mass distribution, to be added to the original lens mass distribution, can be constructed which gives rise to a similar type of degenerate solutions. In particular, the alternative name of steepness degeneracy is still applicable since the construction of degenerate solutions still requires the original mass distribution to be rescaled. Although the method described below is quite straightforward, the authors, having performed a thorough literature study, are not aware of this result being published before.  In section \\ref{sec:formalism} the necessary equations from the gravitational lensing formalism will be reviewed. These equations will be used to derive the mass-sheet degeneracy in section \\ref{sec:masssheet} and to explain the extension to multiple redshifts in section \\ref{sec:extension}. Finally, in section \\ref{sec:conc}, the results and their implications are discussed. ",
        "conclusions": "\\label{sec:conc}  In this article we have presented a straightforward extension of the mass-sheet degeneracy to multiple redshifts. Although there is no actual sheet of mass involved, the procedure and effect of the degeneracy are the same. In both cases, the existing mass distribution is rescaled and a component is added, leading to a rescaling of the sources. This justifies placing the degeneracy described above in the same category as the mass-sheet or steepness degeneracy. Although we illustrated the construction of degenerate solutions using a strong lensing example, it is clear that weak lensing studies can be affected as well. In particular, if weak lensing data alone leads to solutions prone to the mass-sheet degeneracy, it will not suffice to add strong lensing data of a single source to break the degeneracy if strong and weak lensing regions do not overlap. Otherwise, a similar construction as above is possible, using one critical density in the weak lensing region and another in the strong lensing region.  However, as was shown by \\citet{Bradac2004}, using weak lensing data it is in principle possible to break the mass-sheet degeneracy -- including the extension described here -- if individual source redshifts are available and if sources with a rather high distortion are included. Another way to break the degeneracy is to add information about the magnification, for example by using source number statistics \\citep{Broadhurst1995} or Type Ia supernovae observations \\citep{Holz}. Additional information about stellar dynamics in the gravitational lens can also help to break the degeneracy \\citep{Koopmans2006}. Since the mass-sheet degeneracy rescales the time delay surface, time delay measurements can be used to break it as well.  This degeneracy seems to explain what we observed during tests of our non-parametric inversion algorithm \\citep{Liesenborgs2}. Although an example using two sources was used to avoid the mass-sheet degeneracy, we found that our algorithm did not succeed in finding the correct source sizes. However, the general shape and features of the reconstructed mass map did resemble closely the true mass distribution, which was used to create the images that served as input for the inversion routine. It is now clear that images of two sources at different redshifts do not provide enough constraints to firmly establish the scale of these sources.  The procedure outlined above still allows much freedom in the interpolation scheme, and therefore in the precise shape of $\\Sigma_{\\rm gen}$. It is clear, however, that in general it will always be necessary to introduce substructure in $\\Sigma_{\\rm gen}$, which may lead to the presence of substructure in the new mass map. The presence of such substructure will limit the range in which the parameter $\\lambda$ may lie: if too much substructure is introduced, additional images will be predicted and the created mass distribution will no longer be compatible with the observed situation. However, by using several versions of $\\Sigma_{\\rm gen}$, each with a different center of symmetry, the amount of substructure that needs to be introduced to obtain a specific scaling can be limited.  It is straightforward to extend the procedure to more than two sources, but at some point it will no longer be possible to construct such degenerate solutions. How many sources are needed will depend on a number of factors. The number of the images and their sizes play an important role, as this determines how easily their annuli will overlap and therefore how difficult it will be to find appropriate scaling centers. The precise redshifts of the sources will determine how much substructure needs to be introduced, since the redshifts will determine if the corresponding critical densities lie close to each other or differ a lot. If the latter is the case, automatically more substructure will need to be introduced, allowing only a limited range of $\\lambda$ values to avoid the prediction of extra images.  Although at a first glance the mass-sheet or steepness degeneracy is easily broken, this article shows that it is in fact a lot more difficult to do so and a relatively large number of constraints may be needed. The examples above also suggest that circularly symmetric features should be distrusted, as they are easily introduced in degenerate solutions."
    },
    "0801/0801.2190_arXiv.txt": {
        "abstract": "{Recognizing the transport of Energetic Neutral Atoms (ENA), from their place of birth to Earth orbit, has become an important issue in light of the forthcoming launch of the NASA SMEX mission IBEX, which is devoted to imaging of the heliospheric interface by in-situ detection of ENAs. } {We investigate the modifications of both energy of survival probability of the hydrogen ENA (H ENA) detectable by IBEX (0.01 -- 6~keV), between the termination shock and Earth orbit. We take into account the influence of the variable and anisotropic solar wind and of solar EUV radiation. } {Energy changes of the atoms are calculated by numerical simulations of the orbits of H ENA between $\\sim 100$~AU from the Sun and Earth orbit, taking into account solar gravity and Lyman-$\\alpha$ radiation pressure, which is variable in time and depends on the radial velocity of the atom. To calculate the survival probabilities of the atoms against ionization, a detailed observation-based 3D and time-dependent model of H ENA ionization is constructed, and with the use of this model the probabilities of survival of the atoms are calculated by numerical integration along the previously-calculated orbits.} {Due to radiation pressure, H ENA reach the Earth orbit practically without energy and direction change, apart from the atoms of energy lower than 0.1~keV, during high solar activity. The survival probability of H ENA increases from just $\\sim 2\\%$ for the slowest detectable ENA at solar minimum to $\\sim 80$\\% for the fastest ENA. For a given energy at Earth orbit we expect fluctuations in the survival probability of amplitude between $\\sim 20$ percent at 0.01~keV to just a few percent at 6~keV and a modulation of survival probability as a function of the location at Earth orbit, ecliptic latitude of the arrival direction, and phase of solar cycle with an amplitude of a few dozen percent for 0.1~keV atoms at solar minimum to a few percent for 6~keV atoms at solar maximum. } {With appropriate account of local transport effects IBEX should be able to discover departures from symmetry in the flux of H ENA from the heliospheric interface at a level of a few percent.} ",
        "introduction": "Hydrogen Energetic Neutral Atoms (H ENA) provide important information about the physical state of plasma in remote parts of the heliosphere, especially in the heliospheric interface \\citep{gruntman:92b, hsieh_etal:92a, gruntman:97, gruntman_etal:01a, scherer_fahr:03a, fahr_scherer:04a, heerikhuisen_etal:07a, fahr_etal:07a, sternal_etal:08a}. Measurements of heliospheric H ENA are planned by the NASA SMEX mission IBEX \\citep{mccomas_etal:04a, mccomas_etal:05a, mccomas_etal:06}. IBEX will be a spin-stabilized Earth satellite observing Energetic Neutral Atoms (ENA) in the 0.01 -- 6~keV energy band, with sensors looking perpendicularly to the spin axis directed at the Sun. The prime target of the mission will be imaging of the heliospheric interface via neutral atoms and the goal of an individual observation will be to obtain the kinematic parameters of the locally-registered atom beyond the termination shock (TS). To that end, understanding of the processes affecting the atoms between their source region and the detector is needed. In particular, a control on the modifications of the kinematic parameters of the atoms between the source region and the detector and a grasp on loss processes along the way are necessary. Some insights into this problem were presented by \\citet{bzowski_tarnopolski:06a}.  Since the ENA gas inside the TS is collisionless at all energies, then -- from Boltzmann equation -- the transport of ENA from TS to the inner heliosphere is governed solely by dynamic effects and gain and loss processes. The dynamic effects include heliocentric acceleration by a combined action of the solar gravity and the opposing solar radiation pressure, whose strength depends on the radial velocity of the incoming atom. The losses include all known ionization processes operating in the heliosphere: charge exchange, EUV ionization, and electron impact \\citep{rucinski_etal:96a}. For inwardly-traveling ENA, the gains (considered by IBEX as unwanted foreground) include charge exchange between CME, CIR, and pickup ions on the one hand, and all kinds of neutral populations present inside TS on the other hand, as well as production of ENA in the plasma environment of the spacecraft. In this paper, they are neglected because they will constitute a small addition to the net signal \\citep{fahr_etal:07a}.  The paper discusses modifications of kinematic parameters and survival probabilities of H ENA, suitable for detection by IBEX, traveling from the heliospheric termination shock (TS) to 1~AU. First, we discuss the radiation-pressure force affecting the atoms and the resulting trajectories of individual H ENA reaching perihelia of their trajectories at Earth orbit. Then we briefly present the ionization processes destroying the incoming H ENA will be briefly presented, providing a longer discussion in the Appendix. Subsequently, we show the survival probabilities of the H ENA, with emphasis on asymmetries and anisotropies that result from local conditions in the solar wind. The discussion is based, as much as possible, on available solar-wind and solar-radiation output data and does not consider measurement technicalities, such as detection efficiencies, geometric factors, proper motion of the spacecraft etc., which are left to a future study. Emphasis is placed on atoms at the energies expected in the heliosheath region, i.e. up to 0.2~keV; faster atoms are not, however, disregarded, since the studies mentioned earlier predict a non-vanishing production  of H ENA to the upper limit of the IBEX sensitivity threshold. Attention is drawn to solar-cycle-related effects, which include changes in the solar-wind anisotropy and net flux of solar Lyman-$\\alpha$ pressure, and to fluctuations on a timescale of the solar rotation period.  The term ``source region'' frequently used in the paper should be understood as an origin of the H ENA, sufficiently distant from the Sun for its dynamical influence on the atoms to be negligible. It was arbitrarily set to be $\\sim 100$~AU from the Sun. An atom that is unaffected by any radiation pressure and at infinity travels at 12~km/s, at 1~AU will have accelerated to 43.8~km/s, while an atom that travels with the velocity 12~km/s at 100~AU from the Sun, at 1~AU will have accelerated to 43.6~km/s  and the difference between the velocities is then just 0.2~km/s; when radiation pressure is added and larger velocities are adopted, this difference will decrease further.  Dynamical conclusions discussed in Sect.~2 are valid for all atoms reaching perihelion at 1~AU from the Sun irrespective of exact location in the 3D space; the conclusions about the survival probability, however, presented in Sect.~3 were obtained for H ENA reaching perihelia of their trajectories at 1~AU in the ecliptic plane, which is referred to as ``at Earth orbit''.  \\begin{figure*} \\centering \\includegraphics[width=8cm]{9393fig1a.eps} \\includegraphics[width=8cm]{9393fig1b.eps} \\caption{Radiation pressure acting on H atoms, expressed as a factor $\\mu$ of compensation of the solar gravity force, shown as a function of their radial velocity with respect to the Sun for solar minimum and maximum (left-hand panel), and wavelength- and disk-integrated monthly-averaged solar Lyman-$\\alpha$ flux based on the SOLAR 2000 model (right-hand panel). The profiles of radiation-pressure factor $\\mu$ are shown according to the model defined in Eq.(\\ref{e3}) for the net flux $I_{\\mathrm{min}} = 3.7 \\cdot 10^{11}$ cm$^{-2}$~s$^{-1}$ (solar min.) and $I_{\\mathrm{max}} = 6.0 \\cdot 10^{11}$ cm$^{-2}$~s$^{-1}$ (solar max.).} \\label{lyaPlot} \\end{figure*} \\begin{figure*} \\centering \\includegraphics[width=8cm]{9393fig2a.eps} \\includegraphics[width=8cm]{9393fig2b.eps} \\caption{Instantaneous radiation pressure acting on H ENA during 81 days before perihelion located at 1~AU, expressed as the gravity-compensation factor $\\mu$ and shown for a few energies corresponding to the lowermost IBEX energy channels and equal to 0.01~keV (solid), 0.015~keV (dots), 0.029~keV (dash-dot), 0.056~keV (dashed), 0.107~keV (dash-dot-dot), 0.208~keV (lowermost solid lines), for solar minimum (left-hand panel), and maximum conditions (right-hand panel). The perihelion velocities correspond, respectively, to 43.8, 53.6, 74.5, 103.6, 143.2, and 199.6~km/s.  } \\label{radPresPlot} \\end{figure*} ",
        "conclusions": "We investigate modifications of H ENA energy at Earth orbit with respect to their energy at 100~AU from the Sun and their probabilities of survival against losses induced by local heliospheric ionization processes. The range of energy at Earth orbit that we have studied corresponds to the sensitivity range of the planned NASA SMEX mission IBEX, with emphasis on the H ENA energies expected from the inner heliosheath (mostly 0.1 -- 0.2~keV, up to $\\sim 1$~keV). Realistic, observation-based models of both ionization processes and radiation pressure have been used. The ionization model reproduces the time-dependent spatial anisotropy of the solar wind, and the radiation-pressure model takes into account the dependence of $\\mu$ on both the radial velocities of the atoms and on the heliolatitude (Figs.~\\ref{lyaPlot}, \\ref{swprofPl}, \\ref{betaElLati}, \\ref{omniPl}, and \\ref{photoIonPl} and Eqs~(\\ref{equ2}), (\\ref{e3}), and (\\ref{e4})).  Based on the results of our numerical simulations,  we propose that, due to radiation pressure, H ENA reach Earth orbit without a significant changes in energy and direction apart from the atoms with energies below 0.1~keV during high solar activity (Figs.~\\ref{sourceEnClipped}, \\ref{deflectionPlot}). The probability of H ENA survival is a function of the solar-cycle phase, the ecliptic longitude of the location at Earth orbit, and the ecliptic latitude of the approach direction. Especially important in this respect is the latitudinal range of the slow solar-wind region (Figs.~\\ref{wionAniso} and \\ref{absWeightsPl}). For all cases considered, the time of flight from the source region at 100~AU to Earth orbit is equal to the force-free travel time, to an accuracy of 2 solar rotation periods, down to 1~AU energies $\\sim 0.05$~keV.  The survival probability of H ENA declines rapidly with decreasing energy at Earth orbit from $\\sim 0.8$ for $E_{\\mathrm{1\\,AU}} = 6$~ keV to $\\sim 0.4$ for $E_{\\mathrm{1\\,AU}} = 0.2$~keV and $\\sim 0.3$ for $E_{\\mathrm{1\\,AU}} = 0.1$~keV for H ENA approach directions within the ecliptic plane (Fig.~\\ref{radPressRole}). A yearly modulation of the survival probability of H ENA originating in the polar regions, is expected due to the yearly movement of the  Earth-bound detector in heliolatitude. The yearly amplitude of the maximum-to-minimum probability ratio decreases with increasing  energy at Earth orbit; for $E_{\\mathrm{1\\,AU}} \\simeq 0.1$~keV the ratio is equal to $\\sim 1.15$, and for 0.2 keV to $\\sim 1.1$ (Fig.~\\ref{yearlyModulation}). The courses of the modulation for the north and south poles are in antiphase.  Apart from a short time interval at solar maximum, one also expects an appreciable modulation of the survival probability as a function of heliolatitude (Figs.~\\ref{wionAniso} and \\ref{absWeightsPl}). During solar minimum in the case of H ENA with $E_{\\mathrm{1\\,AU}} = 0.1$~keV the amplitude of this modulation should not exceed 1.6, 1.3 for $E_{\\mathrm{1\\,AU}} = 0.2$~keV and 1.1 for $E_{\\mathrm{1\\,AU}} = 6$~keV. Appreciable left-right and north-south asymmetries are expected too, changing with position at Earth orbit (Figs.~\\ref{wionAniso} and \\ref{absWeightsPl}), as well as a distinct correlation of the survival probability profiles with the latitude structure of solar wind. The amplitude of the latter modulation decreases with widening of the slow solar wind region (Fig.~\\ref{absWeightsPl}).  Because of fluctuations of the solar wind and of solar EUV radiation at time scales shorter than solar rotation period fluctuations of the H ENA survival probability at Earth orbit are expected. Based on the solar wind observations at 1~AU and on the simulations discussed in the paper a phenomenology function expressing the amplitude of these fluctuations with respect to the survival probability as a function of the H ENA energy at Earth has been found (Eq.(\\ref{fluctuEq})): it decreases with an increase of energy and becomes lower than 15\\% for $E_{\\mathrm{1\\,AU}}\\simeq 0.1$~keV (Fig.~\\ref{fluctuPl}).  We have also studied the contributions of various factors affecting the losses of H ENA traveling from the source region to Earth orbit. The uncertainty in the charge-exchange cross-section yields uncertainty in the survival probability approximately of a few percent, much lower than the amplitude of the probability fluctuations due to fluctuations in the solar wind. The strength of radiation pressure affects the survival probability on the level of $\\sim 20$\\% for $E_{\\mathrm{1\\,AU}}\\simeq0.1$~keV and $\\sim 15$\\% for $E_{\\mathrm{1\\,AU}}\\simeq0.2$~keV. Photoionization contributes 20\\% and 15\\% of the losses at the energies $E_{\\rm 1\\,AU} \\simeq 0.1$~keV and $E_{\\rm 1\\,AU} \\simeq  0.2$~keV, respectively, and far more for approach directions from high latitudes. Electron ionization for energies above 0.1~keV contributes less than 10\\% of the losses (Fig.~\\ref{miscIonRole}).  All modulation factors discussed increase appreciably -- although to varying degrees -- with a decrease in H ENA energy at Earth orbit. They should, however, disappear during solar maximum, apart from those related to solar-wind fluctuations, when the large-scale structure of the solar wind is spherically symmetric.  The analysis presented in this paper has neglected the observational effect of a proper motion of the detector, which will be the subject of future studies. Based on results obtained thus far it is concluded that the most important elements needed for a successful interpretation of H ENA measurements, such as those planned for the forthcoming IBEX mission, are monitoring of the local solar wind, the latutudinal structure of the solar wind, and the solar Lyman-$\\alpha$ flux. The solar EUV flux in the spectral range responsible for photoionization of hydrogen is also important, and especially its dependence on heliolatitude. Taking into account the local effects of H ENA transport from source to Earth orbit, we expect that IBEX should be able to discover departures from symmetry in the heliospheric interface of approximately a few percent.  \\begin{appendix}"
    },
    "0801/0801.2159_arXiv.txt": {
        "abstract": "We analyze a high-resolution spectrum of a microlensed G-dwarf in the Galactic bulge, acquired when the star was magnified by a factor of 110. We measure a spectroscopic temperature, derived from the wings of the Balmer lines, that is the same as the photometric temperature, derived using the color determined by standard microlensing techniques. We measure [Fe/H]=$0.36 \\pm 0.18$, which places this star at the upper end of the Bulge giant metallicity distribution.  In particular, this star is more metal-rich than any bulge M giant with high-resolution abundances.  We find that the abundance ratios of alpha and iron-peak elements are similar to those of Bulge giants with the same metallicity.  For the first time, we measure the abundances of K and Zn for a star in the Bulge. The [K/Mg] ratio is similar to the value measured in the halo and the disk, suggesting that K production closely tracks $\\alpha$ production. The [Cu/Fe] and [Zn/Fe] ratios support the theory that those elements are produced in Type II SNe, rather than Type Ia SNe. We also measured the first C and N abundances in the Bulge that have not been affected by first dredge-up. The [C/Fe] and [N/Fe] ratios are close to solar, in agreement with the hypothesis that giants experience only canonical mixing. ",
        "introduction": "}    The Galactic Bulge underwent an intense burst of star formation early in the formation of the Galaxy, leading to a very different stellar population and chemical evolution history than found in the Milky Way disk or halo \\citep[e.g.,][]{ortolani:95,zoccali:03,mcwilliam:94}. In particular, massive stars may dominate the pollution at almost all metallicities, leading to unique abundance patterns \\citep[e.g.,][]{lecureur:07}. Similar events are thought to mark the formation of other galactic spheroids, making the Bulge stellar population a template for interpreting extragalactic observations. As a result of its unique formation history in the Galaxy, the Bulge has been the subject of intensive study.  The detection of RR Lyrae stars \\citep{baade:46} first indicated that the Bulge contained old stars. With deeper photometry, the main sequence turnoff (MSTO) of the Bulge was detected. \\citet{terndrup:88} found a mean age of 11-14 Gyr for stars in Baade's Window, with a negligible fraction of stars with ages $<$ 5 Gyr. Photometry reaching more than 2 magnitudes below the MSTO with {\\it Hubble Space Telescope} confirmed the generally old nature of the Bulge \\citep{ortolani:95,holtzman:98}, although \\citet{feltzing:00} included a reminder that a young metal-rich population has similar MSTO colors and luminosities to an older, more metal-poor population. Therefore, deriving ages reliably from photometry of the MSTO requires adequate knowledge of the Bulge metallicity distribution function (MDF), specifically for the MSTO stars. However, because of the faintness of those stars, the measurement of the Bulge MDF has historically relied on giants.  Since the discovery of both M giants and RR Lyr stars, it has been known that the Bulge giants span a wide range in metallicity. Low-dispersion spectra provided the first quantitative measure of the MDF \\citep{whitford:83,rich:88}.  \\citet{sadler:96} measured indices from low-dispersion spectra of 268 K giants (both red clump stars and first ascent giants) to derive a mean metallicity\\footnote{We adopt the usual spectroscopic notation that [A/B] $\\equiv$ {\\rm log}$_{\\rm 10}$(N$_{\\rm A}$/N$_{\\rm B}$)$_{\\star}$ -- {\\rm log}$_{\\rm 10}$(N$_{\\rm A}$/N$_{\\rm B}$)$_{\\odot}$}) $\\langle{\\rm [Fe/H]}\\rangle =-0.11$, with a dispersion of 0.46 dex. Recalibration by \\citet{fulbright:06} based on high-resolution spectra of 15 stars in common with the \\citet{sadler:96} sample reduced the mean metallicity to $-0.22$.  \\citet{ramirez:00} measured $\\langle{\\rm [Fe/H]}\\rangle =-0.21$ from low-dispersion near-infrared spectra for 72 M giants in the inner Bulge. The good agreement between \\citet{ramirez:00} result and the recalibrated \\citet{sadler:96} result is somewhat surprising. At the bright end of the giant branch, only metal-rich first ascent giants become M giants. However, both K giants and M giants become red clump stars, and lower luminosity metal-rich giants are K stars as well. Therefore, the inclusion of red clump stars  and fainter giants in the \\citet{sadler:96} sample make the biases in their sample more similar to those of the \\citet{ramirez:00} study.  \\citet{zoccali:03} measured both the MDF of the bulge and the age of the stars using deep photometry of the Bulge in the optical and near-infrared wavelengths. They constructed the M$_{K}$ and (V-K)$_0$ color magnitude diagram for a low-reddening window at ($l,b$)=(0.277, $-$6.167). Because the slope of the red giant branch (RGB) in these colors depends on the metallicity, RGB stars with different metallicities fall on different parts of the color-magnitude diagram. They could therefore use their photometry to derive the MDF of 503 giants by comparing the positions of bright (M$_K < -4.5$) RGB stars with globular cluster fiducials of known metallicity. After correcting for small biases in their MDF caused by their magnitude cutoff, they find an MDF with a peak at [M/H]$=-0.1$, a sharp cutoff at [M/H]$=-0.2$ and few stars with [M/H]$<-1$. Adopting the metallicities derived from the giants for the MSTO dwarfs, they estimated that the Bulge is coeval with the halo and argued that the lack of stars above the prominent MSTO of the Bulge ruled out a significantly younger population.  These measurements of the Bulge giant MDF can be improved by metallicity measurements from high-dispersion measurements of many stars in several fields throughout the Bulge. Recently, \\citet{lecureur:08} and \\citet{zoccali:08}  have obtained a total of $\\sim 1000$ K giant spectra at (R$\\sim$20,000) in 4 Bulge windows. They confirm the metal-rich nature of the Bulge.    With such a metal-rich population having been reached so quickly after star formation began, the Bulge is expected to have a distinct chemical evolution compared to other Galactic populations, such as the halo or the disk because, for example, the contributions of longer-lived polluters, such as Type Ia supernovae (SNe) or low-mass AGB stars, should be small. Indeed, \\citet{mcwilliam:94} measured high [$\\alpha$/Fe] ratios in giants, in particular, high [Mg/Fe] for [Fe/H] values up to solar, arguing for little Type Ia SN contribution of Fe compared to the thick or thin disks. \\citet{fulbright:07} confirmed the overall enhancement in [Mg/Fe] and strengthened the conclusion of \\citet{mcwilliam:94} that the other $\\alpha$ abundance ratios do not track [Mg/Fe] exactly. [O/Fe], [Si/Fe], [Ca/Fe] and [Ti/Fe] begin to decrease around [Fe/H]=0, while [Mg/Fe] does so at supersolar [Fe/H].  \\citet{fulbright:07} suggested that metallicity-dependent Type II SN yields could explain the different behaviors of the $\\alpha$ elements. \\citet{mcwilliam:07} explained the low [O/Mg] ratios in metal-rich Bulge giants through a different metallicity-dependent mechanism: Wolf-Rayet winds leading to less effective O production in metal-rich massive stars. By [Fe/H]$\\sim$0.2, all the [$\\alpha$/Fe] ratios have begun to decline, probably indicating the introduction of large amounts of Fe from Type Ia SNe \\citep[e.g.,][]{cunha:06}.  Several studies have looked at the abundances of the light elements Na and Al in the Bulge, two elements whose production should depend on the metallicity of the massive stars that exploded as Type II SNe. \\citet{cunha:06} measured Na in 7 K and M giants and found the predicted increase in [Na/Fe] and [Na/O] in the most metal-rich stars. \\citet{lecureur:07} found supersolar [Al/Fe] at all metallicities and, for [Fe/H] $>$0, enhanced [Na/Fe] compared to the ratios in disk stars. Interestingly, at higher metallicities, the scatter in [Al/Fe] and [Na/Fe] increased and became larger than could be explained by observational errors. Interpreting the Na and Al abundances as the result of Type II SN production may be problematic. The surface abundances of Al and Na have been shown to be increased by large amounts of internal mixing in metal-poor globular cluster stars \\citep[e.g.][]{shetrone:96}, where the products of  proton-capture reactions deep inside the star are mixed up to the surface, leading to enhancements in these two elements. However, \\citet{lecureur:07} argued that the C and N abundances in the giants they studied were consistent with only mild mixing and, therefore, that the high Na and Al had to be due to the overall chemical evolution of the Bulge. \\citet{cunha:06} also found evidence for mild mixing in giants, affecting C and N, but not O, Na, or Al.  Finally, there is little information on the neutron-capture elements in the Bulge. The absorption lines for these elements are concentrated in the blue part of the optical spectrum, where the crowding from Fe, CN, and other lines is severe. Near-IR spectra have essentially no lines of these elements. While there are a few lines of Ba in the red, these lines in metal-rich giants are so saturated that reliable measurements are very difficult. As a result, only the neutron-capture element Eu has published results so far. \\citet{mcwilliam:94} found [Eu/Fe]$>$0 in Bulge giants, which is likely because of the production of Eu in the r-process. Measuring additional neutron-capture element abundances would test this idea, because the r-process is better at making some elements (e.g., Eu) than others (e.g., Ba).  Our knowledge of the metallicity and abundance ratios of Bulge stars has generally been confined to the bright giants, which are usually the only ones accessible to high-resolution observations.  But studying the dwarfs has several advantages.  Their abundances of the elements such as C and N are unaffected by dredge-up processes on the giant branch. We can measure elements, such as S and Zn, that not measured in giants, because the hotter temperatures of the dwarfs decrease the strength of CN and increase the strength of certain atomic lines. In addition, it is critical to know the metallicity of stars at the MSTO to accurately measure the ages from their color and luminosity. Finally, individual ages can be assigned to dwarf stars near the turnoff.  The advent of large surveys to identify and follow-up microlensing events, such as the Microlensing Observations in Astrophsycis (MOA) collaboration\\footnote{http://www.massey.ac.nz/$\\sim$iabond/alert/alert.html}, the Optical Gravitational Lens Experiment\\footnote{http://www.astrouw.edu.pl/$\\sim$~ogle/ogle3/ews/ews.html} (OGLE), the Microlensing Follow Up Network\\footnote{http://www.astronomy.ohio-state.edu/$\\sim$microfun/} ($\\mu$FUN) and the Probing Lensing Anomalies Network\\footnote{http://planet.iap.fr/} (PLANET), provides an opportunity to study otherwise unobservable Bulge dwarfs.  During high-magnification microlensing events, it is possible to obtain high-resolution, high signal-to-noise ratio spectra of faint stars with a huge savings in observing time: a factor $A\\times 10^{0.4\\Delta m}$ where $A$ is the magnification and $\\Delta m$ is the number of magnitudes below sky of the unmagnified star.  In \\citet{johnson:07} , we reported the detailed abundances for a highly-magnified Bulge dwarf, \\ogle, which from a 15 minute exposure at magnification $A\\sim 135$, was shown to be one of the most metal-rich stars ever observed.  It also provided the first measurements of S and Cu in the Bulge.  Here we present a high-resolution spectrum of the Bulge G-dwarf MOA-2006-BLG-99S, taken at magnification $A=110$.   Finally, we respond the challenge: ``Ask not what microlensing can do for stellar spectroscopy -- ask what stellar spectroscopy can do for microlensing.'' There is one important way that the spectroscopic study of bulge dwarfs can benefit microlensing. Whenever a source approaches or transits a ``caustic'' (line of infinite magnification) caused by the lens, one can measure $\\rho$, the ratio of the angular source radius to angular Einstein radius $\\rho=\\theta_*/\\theta_{\\rm E}$, from the microlens lightcurve. Then $\\theta_*$ is inferred from the dereddened color and magnitude of the source to yield $\\theta_{\\rm E}=\\theta_*/\\rho$, which in turn provides important constraints on the lens properties. Because spectroscopy is not normally available for these microlensed sources, the dereddened color and magnitude are estimated by comparing the source position on an instrumental color-magnitude diagram with that of the clump and then assuming that the Bulge clump is similar to the local clump as measured by {\\it Hipparcos} \\citep{ob03262}.  This procedure undoubtedly suffers some statistical errors and could suffer systematic errors as well.  For example, the Bulge clump may have a different color from the local one.  High-resolution spectra of an ensemble of microlensed bulge sources will test this procedure for both statistical and systematic errors. For OGLE-2006-BLG-265S, the standard microlensing procedure yielded $(V-I)_0=0.63\\pm 0.05$, whereas \\citet{johnson:07} obtained $(V-I)_0=0.705\\pm 0.04$ from high-resolution spectroscopy.  This difference hints at a possible discrepancy, but only by repeating this procedure on a number of dwarfs can this be confirmed. ",
        "conclusions": "The results for \\moaspaces demonstrate the unique information that can be obtained from high-resolution spectra of microlensed dwarfs in the Bulge. Because microlensing is not biased in metallicity, we can measure the MDF of the Bulge main-sequence stars when sufficient number of highly magnified Bulge dwarfs have been observed with high-dispersion spectrographs. However, the two dwarfs we have studied so far both have [Fe/H]$> 0.30$, which makes them more metal-rich than any of the M giants observed at high-resolution. If the temperatures and metallicities measured for the two dwarfs are accurate, then these stars are younger than the bulk of the Bulge population.  For many elements, the abundance ratios we measure in \\moaspaces are not surprising. The iron-peak and $\\alpha$-elements (O, Mg, Si, Ca, and Ti) fall on the trend observed in the Bulge giants. The solar values for the [$\\alpha$/Fe] ratios suggest that Type Ia SN ejecta are responsible for some of the Fe, which would be expected if these metal-rich stars were part of a younger population in the Bulge. The [C/Fe] and [N/Fe] ratios are also $\\sim$ 0, suggesting that AGB stars have also begun to contribute these light elements to offset the contribution of Fe from Type Ia SNe. It is therefore unexpected to find [Ba/Fe]$\\sim$-$0.5$. Such low [Ba/Fe] values, when found in the halo, imply that only Type II SNe and the r-process have polluted the gas, and not the s-process from AGB stars.  The C and N abundances we measure in \\moaspaces are unaffected by any mixing or dilution and therefore are the result of the chemical evolution of the Bulge. Comparing the total amount of C+N to that measured in Bulge giants, we conclude that only mild mixing, as expected in theoretical models, has occurred in the Bulge. This agrees with the assertions by \\citet{lecureur:07}, \\citet{zoccali:06}, and \\citet{mcwilliam:07} that the O, Na and Al abundances in Bulge giants should not have been affected by mixing. The [Na/Fe] and [Al/Fe] ratios for \\moaspaces fall on the lower end of the values seen in Bulge giants of this metallicity. The [Na/Fe] and [Al/Fe] ratios in giants are marked above all else by large scatter and more abundance ratios in dwarfs are needed before a comparison of scatter can be made.  We can use our measurements of the abundances of K, Cu and Zn, rare for Bulge stars, to exploit the different star formation history of the Bulge to study the nucleosynthesis sites of these elements. The supersolar [Cu/Fe] and subsolar [Zn/Fe] values agree with the model of \\citet{bisterzo:05} and the production of these two elements predominantly in Type II SNe. The chemical evolution of K is more difficult to understand, because the [K/Mg] value in the Bulge is similar to that in the metal-poor halo, although the production of K should be enhanced in metal-rich SNe.  Finally, we note that the standard microlensing technique of estimating the intrinsic color and magnitude of the source star by using the offset from the position of the red clump can be tested by deriving a temperature and gravity from the spectrum. For \\moas, the photometric and spectroscopic temperature agree very well, although the agreement was considerably less good for \\ogle.  In summary, the abundances in \\moaspaces provide unique information about the formation and evolution of the Bulge. Larger samples of dwarfs observed in this way would allow the derivation of the dwarf MDF and the trends in abundance ratios over a wide range of metallicity. By taking advantage of microlensing events in the Bulge, we can achieve this goal with a modest amount of telescope time."
    },
    "0801/0801.2968_arXiv.txt": {
        "abstract": "The ability to map the cosmological expansion has developed enormously, spurred by the turning point one decade ago of the discovery of cosmic acceleration.  The standard model of cosmology has shifted from a matter dominated, standard gravity, decelerating expansion to the present search for the origin of acceleration in the cosmic expansion.  We present a wide ranging review of the tools, challenges, and physical interpretations. The tools include direct measures of cosmic scales through Type Ia supernova luminosity distances, and angular distance scales of baryon acoustic oscillation and cosmic microwave background density perturbations, as well as indirect probes such as the effect of cosmic expansion on the growth of matter density fluctuations.  Accurate mapping of the expansion requires understanding of systematic uncertainties in both the measurements and the theoretical framework, but the result will give important clues to the nature of the physics behind accelerating expansion and to the fate of the universe. ",
        "introduction": "\\label{sec:intro}  A century ago our picture of the cosmos was of a small, young, and static universe.  Today we have a far grander and richer universe to inhabit, one that carries information on the strongest and weakest forces in nature, whose history runs from singularities and densities and temperatures far beyond our terrestrial and laboratory access to the vacuum and temperatures near absolute zero.  Understanding our universe relies on a wide range of physics fields including thermodynamics, classical and quantum field theory, particle physics, and gravitation.  Perhaps most amazing is that while our universe is huge, it is finite in well defined ways, and can be encompassed and comprehended.  The visible universe extends for nearly an equal number of orders of magnitude above the size of the Earth as the proton lies below.  The number of particles is of order $10^{80}$, large but not unbounded, and the age of the universe is some 14 billion years, only three times the age of the Earth.  And the universe is simple in ways we have no right to expect: it is essentially electrically neutral and its large scale dynamics is governed by gravity and no other forces of nature, it is nearly in thermal equilibrium for most of its history, and the geometry of space is maximally symmetric.  A few characteristics hold where we might be tempted to say, as Alfonso X ``The Wise'' did in the 13th century: ``Had I been present at the creation of the world, I should have recommended something simpler''.   The universe has about a billion times more entropy per baryon than expected, and in a related sense has a greatly unequal ratio of matter to antimatter. The dynamics is neither kinetic energy dominated nor potential energy dominated but apparently perfectly balanced, giving the critical energy density and a flat spatial geometry.  But the most important cosmological discovery of the 20th century was that the maximal symmetry does not extend to spacetime; that is, we do not live in a steady state universe unchanging in time.  Discovery of the cosmic expansion of space in the 1910s and 1920s and that the evolution arose from a hot, dense, early state called the Big Bang in the 1960s gave rise to modern cosmology as a exemplar of, window on, and laboratory for physics.  Ten years ago the discovery of the {\\it acceleration\\/} of that expansion revolutionized cosmology and a great array of overlapping fields of physics.  This article addresses our current knowledge of the cosmic expansion and our prospects for exploring it in detail. In particular, we focus on the recent epoch of acceleration and the astrophysical tools for mapping the cosmic expansion history to reveal the nature of new physics beyond our present standard model.  In the remainder of this section we discuss how the expansion of our universe impacts fundamental questions of the origin and fate of the cosmos and everything in it, and its intimate relation with the nature of gravity, as a force itself and reflecting on unification with quantum theory.  We present the effects of acceleration in \\S\\ref{sec:eff}, but do not go into the root causes of it, which are highly speculative at this time; see, e.g., the focus issue on dark energy in \\citet{grgissue}. Techniques for directly mapping the expansion appear in \\S\\ref{sec:dist}, while \\S\\ref{sec:gro} briefly mentions indirect effects of acceleration through the growth of structure in the universe.  Obtaining an accurate map is crucial to our understanding, and \\S\\ref{sec:sys} focuses on challenges imposed by systematic uncertainties within the theoretical interpretation and data analysis.  For future measurements these may be the key issues in advancing our knowledge.  Future prospects for mapping the cosmic expansion from the Big Bang to the final fate are outlined in \\S\\ref{sec:fut}, and \\S\\ref{sec:concl} summarizes and concludes.   \\subsection{The dynamic universe}  In the 1910s the frequency shift of spectral lines from astronomical sources with respect to laboratory measurements were observed by Slipher and the dynamics of space was found by Einstein within his theory of general relativity.  While Einstein, together with de Sitter, acted to counteract the expansion of space and retain a static state through introduction of the cosmological constant, other researchers in the 1920s such as Friedmann and Lema\\^{\\i}tre calculated the expansion history of a universe containing matter and components with pressure, and Weyl recognized that a physical expansion scaling linearly with distance occurred naturally. The observations of Hubble in 1929 established the expanding universe as the basis of cosmology.  From the 1930s to 1980s, astronomical observations of the behavior of sources at different redshifts, and the increasing corroboration of redshift as directly related to distance, confirmed the picture of an expanding universe.  (For a collection of some important early papers, see \\citet{BernsteinJ}.)  Figure~\\ref{fig:history} shows the evolution of our capabilities to map the cosmic expansion and the subsequent understanding achieved.  Comparison of theoretical calculations and observations with respect to primordial nucleosynthesis and the cosmic microwave background (CMB) radiation clarified the model as one arising from a hot, dense, near singular state given the name Big Bang.  However, estimates of the matter density (and even more so other component contributions, save for radiation measured directly in the CMB) remained somewhat vague, and in fact little different from Friedmann's 1922 work. Data could not make a definitive statement as to whether the matter and energy density amounted to less than, equal to, or possibly greater than the critical density needed for a spatially flat, asymptotically static expansion.  This all changed dramatically in 1998 with the discovery of acceleration of the cosmic expansion, by two independent groups mapping the distance-redshift relation of Type Ia supernovae \\citep{riess98,perl99}. The expansion history behavior jumped from being somewhere between a critical, matter dominated universe and an open, spatial curvature dominated universe (approaching an empty universe akin to the Minkowski space of merely special relativity), and instead shifted toward one with similarities to de Sitter space governed by a cosmological constant. What was so revolutionary was not that such behavior merely changes the quantitative aspects of the expansion, illustrated in Figure~\\ref{fig:history}, but differs qualitatively and fundamentally, breaking the relation between geometry and destiny.   \\begin{figure}[!htb] \\begin{center} \\psfig{file=history1.ps,height=3.in} \\psfig{file=history2.ps,height=3.in} \\psfig{file=history3.ps,height=3.in} \\caption{Mapping the expansion history has been a major theme of cosmology, revealing the constituents and nature of our universe. From early models with diverse properties (top left panel), scientists narrowed in to a region shown by the green box corresponding to Big Bang models with a long matter dominated epoch.  By the 1980s cosmologists believed the universe corresponded to an Einstein-de Sitter universe with critical matter density, or perhaps some model more toward the empty universe curve (top right panel).  The remarkable discovery in 1998 of accelerated expansion showed that the correct model must be similar to the solid black curve, partaking of characteristics of the de Sitter universe with a cosmological constant $\\Lambda$.  The challenge now is to further improve our cosmic mapping ability to zoom in on the narrow green triangular region, revealing the nature of the new physics behind acceleration (bottom panel).  (See \\citet{Linprl} for the specific models.) } \\label{fig:history} \\end{center} \\end{figure}   \\subsection{Geometry and destiny} \\label{sec:geom}  For a non-accelerating universe described on large scales by a smooth, homogeneous and isotropic, i.e.\\ Friedmann-Robertson-Walker model, geometry and destiny are inextricably linked.  Whether the spatial curvature is positive, zero (flat), or negative correlates one-to-one with whether the expansion is closed (bounded), critical (asymptotically static), or open (unbounded).  However, the Einstein field equations governing the cosmic expansion involve not only the amount of energy density but the full energy-momentum contribution.  So there is a loophole to escape destiny.  The Friedmann equations (the Einstein equations in a homogeneous, isotropic model) can be written as \\beq \\dot a^2+k-8\\pi G\\rho a^2/3\\equiv \\dot a^2+V_{\\rm eff}(a)=0, \\eeq where $a$ is the expansion scale factor, $k$ is the spatial curvature, and $\\rho$ is the total energy density.  (We consider the effective potential $V_{\\rm eff}$ below.) From the continuity (or conservation of energy-momentum) equation, the total energy density $\\rho$ will evolve with expansion as $\\rho\\sim a^{-3(1+w)}$, where $w=p/\\rho$ is the pressure to energy density, or equation of state, ratio (easily generalized for a time varying ratio). So $\\rho a^2\\sim a^{-(1+3w)}$.  If $w>-1/3$ then the sum of the $k$ and $\\rho a^2$ terms ranges from $-\\infty$ to $k$ for all values of $a$, i.e.\\ the entire expansion history.  Thus, whether $\\dot a^2$ ever reaches zero and hence whether there is a maximum value of $a$ or expansion continues without halt is wholly governed by the value of the spatial curvature $k$: is it positive or negative.  Geometry controls density.  However, if $w<-1/3$ then the $\\rho a^2$ term grows with $a$ and eventually dominates as the expansion factor grows large.  This strongly negative contribution can overcome a positive $k$, so this breaks the link between geometry and destiny.  One can get a visual appreciation for this by considering the $k$ and $\\rho a^2$ terms together as making up an effective potential energy (see, e.g., \\citet{araa}), such as standardly used in physics analysis to determine whether a system with a certain kinetic energy is bound or not.  Here, the analog of kinetic energy is $\\dot a^2$, and we want to know whether the minimum kinetic energy, $\\dot a^2=0$, hits the potential energy curve for any value of $a$, indicating a maximum expansion factor.  Figure~\\ref{fig:veff} for the effective potential $V_{\\rm eff}$ vs.\\ $a$ then illustrates the conditions discussed above.  If $w>-1/3$ then the effective potential ranges from $-\\infty$ to $k$ for all values of $a$, i.e.\\ the entire expansion history.  Thus, whether the effective potential crosses the minimum kinetic energy value of zero and hence whether there is a maximum value of $a$ or expansion continues without halt is wholly governed by the value of the spatial curvature $k$: geometry controls density.   \\begin{figure}[!htb] \\begin{center} \\psfig{file=veff.ps,width=3.4in} \\caption{Effective potentials for universes classified according to their geometry (sign of $k$) and whether the cosmic expansion accelerates ($w<-1/3$) or not.  Models crossing the dotted zero line do not expand forever.  For $w>-1/3$ the geometry determines the destiny: whether models with physical kinetic energy, hence $V_{\\rm eff}\\le0$, can achieve $a\\to\\infty$, i.e.\\ expand forever.  By contrast, for accelerating universes the destiny is eternal expansion regardless of geometry -- unless the energy density $\\rho$ can go negative. } \\label{fig:veff} \\end{center} \\end{figure}   However, if $w<-1/3$ then the second term of the effective potential eventually dominates as the expansion factor grows large.  This strongly negative contribution can overcome a positive $k$, so one can in fact have an eternal, positive curvature universe.  Conversely, if the energy density itself goes negative (e.g.\\ due to a negative cosmological constant) then a finite, negative curvature universe is possible.  To determine the crucial quantity, the value of $w$, we need to take into account the entire energy-momentum not just the energy density. From the other Friedmann equation, \\beq \\frac{\\ddot a}{a}=-\\frac{4\\pi G}{3}(\\rho+3p)=-\\frac{4\\pi G}{3}\\rho\\,(1+3w), \\eeq we see that the condition $w<-1/3$ is precisely the condition for accelerated expansion, $\\ddot a>0$.  When one of the components of the universe has sufficiently negative pressure and contributes substantial energy density, such so-called dark energy can cause the total equation of state to drop below $-1/3$ and cause acceleration.  Destiny then becomes unhinged from geometry, and we must understand the nature of dark energy to predict the fate of the universe.  We discuss this further in \\S\\ref{sec:fate}.   \\subsection{Acceleration}  Einstein's equivalence principle states that acceleration is curvature is gravity.  While space may (or may not) be flat, {\\it spacetime\\/} curvature is nonzero in an expanding universe.  Rather than viewing gravity as a force acting at a distance, we can view it as the curvature of spacetime and particles follow paths of least action (geodesics) in this curved spacetime.  An intuitive way of seeing these deep connections is given in Figure~\\ref{fig:equiv}.  \\begin{figure}[!htb] \\begin{center} \\psfig{file=eqvprinc.ps,width=3.4in} \\caption{A photon experiences a frequency shift from acceleration, gravity, or spacetime curvature.  Here the spacetime diagram illustrates a Pound-Rebka experiment of a photon propagating from one atom to another at height $h$ above it.  A gravitational field leads to a gravitational redshift $z=gh/c^2$, while equivalently a uniform acceleration $g$ over a time $t$ leads to a velocity $v=gt$ and a Doppler shift $z=v/c=gt/c=gh/c^2$. Since a frequency change is equivalent to a change of time intervals, the horizontal, parallel line segments representing the emitted photon pulse period $t_0$ and received period $t'$ cannot be equal.  This means the photon propagation lines (wavy diagonal lines), which should each be at $45^\\circ$ in a flat spacetime diagram, must not in fact be parallel.  Thus acceleration, gravitation, and spacetime curvature are equivalent. } \\label{fig:equiv} \\end{center} \\end{figure}   Thus, mapping the expansion history and its acceleration is equivalent to mapping spacetime or to exploring the gravitational nature of the universe.   \\subsection{Revolution in physics}  Is the discovery and further exploration of cosmic acceleration truly revolutionary?  Let us ask what is our current understanding of the universe through the Standard Model: baryons, photons, neutrinos, etc.\\ make up roughly 4\\% of the energy density of the universe.  Understanding 4\\% is like one letter out of the alphabet, e.g.\\ reading  \\begin{quote} \\fon{The nature} o\\fon{f the physics} o\\fon{f the accelerating universe is the premier mystery f}o\\fon{r this generati}o\\fon{n} o\\fon{f physics. The challenge} o\\fon{f a new, d}o\\fon{minant, str}o\\fon{ngly negative pressure c}o\\fon{mp}o\\fon{nent} o\\fon{f the universe is cl}o\\fon{sely tied with the f}o\\fon{undati}o\\fon{ns} o\\fon{f quantum the}o\\fon{ry and the nature} o\\fon{f the vacuum; it is central t}o \\fon{the the}o\\fon{ry} o\\fon{f gravitati}o\\fon{n and the nature} o\\fon{f spacetime, p}o\\fon{ssibly even the number} o\\fon{f dimensi}o\\fon{ns.  The behavi}o\\fon{r} o\\fon{f dark energy g}o\\fon{verns the fate} o\\fon{f the universe.} \\end{quote}  We may hope to soon detect and eventually characterize dark matter. What will be our understanding once we have succeeded with this characterization and employed it to extend the Standard Model to a new theory of high energy physics such as supersymmetry?  Then we will comprehend 25\\% of the energy density of the universe, like all the vowels (with y) of the alphabet:  \\begin{quote} \\fon{Th}e \\fon{n}a\\fon{t}u\\fon{r}e o\\fon{f th}e \\fon{ph}y\\fon{s}i\\fon{cs} o\\fon{f th}e a\\fon{cc}e\\fon{l}e\\fon{r}a\\fon{t}i\\fon{ng} u\\fon{n}i\\fon{v}e\\fon{rs}e i\\fon{s th}e \\fon{pr}e\\fon{m}ie\\fon{r m}y\\fon{st}e\\fon{r}y \\fon{f}o\\fon{r th}i\\fon{s g}e\\fon{n}e\\fon{r}a\\fon{t}io\\fon{n} o\\fon{f ph}y\\fon{s}i\\fon{cs.  Th}e \\fon{ch}a\\fon{ll}e\\fon{ng}e o\\fon{f} a \\fon{n}e\\fon{w, d}o\\fon{m}i\\fon{n}a\\fon{nt, str}o\\fon{ngl}y \\fon{n}e\\fon{g}a\\fon{t}i\\fon{v}e \\fon{pr}e\\fon{ss}u\\fon{r}e \\fon{c}o\\fon{mp}o\\fon{n}e\\fon{nt} o\\fon{f th}e u\\fon{n}i\\fon{v}e\\fon{rs}e i\\fon{s cl}o\\fon{s}e\\fon{l}y \\fon{t}ie\\fon{d w}i\\fon{th th}e \\fon{f}ou\\fon{nd}a\\fon{t}io\\fon{ns} o\\fon{f q}ua\\fon{nt}u\\fon{m th}eo\\fon{r}y a\\fon{nd th}e \\fon{n}a\\fon{t}u\\fon{r}e o\\fon{f th}e \\fon{v}a\\fon{c}uu\\fon{m;} i\\fon{t} i\\fon{s c}e\\fon{ntr}a\\fon{l t}o \\fon{th}e \\fon{th}eo\\fon{r}y o\\fon{f gr}a\\fon{v}i\\fon{t}a\\fon{t}io\\fon{n} a\\fon{nd th}e \\fon{n}a\\fon{t}u\\fon{r}e o\\fon{f sp}a\\fon{c}e\\fon{t}i\\fon{m}e\\fon{, p}o\\fon{ss}i\\fon{bl}y e\\fon{v}e\\fon{n th}e \\fon{n}u\\fon{mb}e\\fon{r} o\\fon{f d}i\\fon{m}e\\fon{ns}io\\fon{ns.  Th}e \\fon{b}e\\fon{h}a\\fon{v}io\\fon{r} o\\fon{f d}a\\fon{rk} e\\fon{n}e\\fon{rg}y \\fon{g}o\\fon{v}e\\fon{rns th}e \\fon{f}a\\fon{t}e o\\fon{f th}e u\\fon{n}i\\fon{v}e\\fon{rs}e\\fon{.} \\end{quote}  It is apparent that for true understanding we will need to know the nature of dark energy.  Only then will we have a complete and comprehensible picture:  \\begin{quote} The nature of the physics of the accelerating universe is the premier mystery for this generation of physics.  The challenge of a new, dominant, strongly negative pressure component of the universe is closely tied with the foundations of quantum theory and the nature of the vacuum; it is central to the theory of gravitation and the nature of spacetime, possibly even the number of dimensions.  The behavior of dark energy governs the fate of the universe. \\end{quote}  Understanding cosmic acceleration undeniably will extend the frontiers of physics and rewrite the textbooks.  It is not excessive to suggest that we are faced with a revolution in our understanding of nature as profound as the mystery of blackbody radiation a century ago.  Blackbody radiation taught us the existence of photons and the structure of atoms; dark energy may lead us to the existence of quantum gravity and the structure of the vacuum. ",
        "conclusions": "\\label{sec:concl}  Mapping the cosmological expansion is a key endeavour in our quest for understanding physics -- gravitation and other forces, spacetime, the vacuum -- and the origin, evolution, and future of the universe.  The Big Bang model of a hot, dense, early universe expanding and adiabatically cooling, and forming structure through gravitational instability from primordial seed perturbations, is remarkably successful and simple. The concordance cosmology is (close to) spatially flat and 13.7 billion years old.  This clear statement represents a substantial advance of our knowledge over a decade ago.  The discovery of the acceleration of the cosmic expansion created a renaissance of cosmological exploration, offering the hope of connecting quantum physics and gravitation, extra dimensions and the nature of spacetime, while severing the bonds between geometry and destiny.  This opens up completely two premier questions in science: the origin and the fate of the universe.  To understand the nature of the gravitationally repulsive dark energy pulling the universe apart, we must map the cosmological expansion in greater detail and accuracy than ever envisioned.  We have shown that a considerable part of the optimal approach for mapping the expansion history is set purely by the innate cosmological dependences. Observational programs must follow these foundations as basic science requirements: in particular the need for a wide range of redshifts, $z\\approx0-2$ and robust anchoring to either low or high redshift. Our capabilities for measuring the expansion through direct geometric probes are increasing, and techniques are continuing to develop. There are no short cuts -- detailed design of successful surveys works within this framework with the purpose to minimize systematic uncertainties in the measurements.  Every technique has systematic uncertainties, appearing in multiple guises. While observational systematics are most familiar, arising even in purely geometric techniques, equally important are issues in the data analysis, e.g.\\ combining heterogeneous data sets, and in the theoretical interpretation, e.g.\\ susceptibility to biases from assumptions of high redshift behavior or of non-expansion physics (perturbation growth behavior, coupling, etc.).  We have given several concrete examples of the effects of systematic uncertainties for various probes.  These provide a cautionary tale for survey design, as we are now entered into the systematics dominated data era -- and may well soon approach the theory systematics era.  We cannot rely on assumptions that any part of the dark sector is simple and ignorable while we concentrate on another aspect.  For true progress, we emphasized the role of complementarity in building from robust, clean answers to more complex investigations. Probes employing growth of structure give windows on both expansion per se and gravitational laws, and we commented that the excitement of testing general relativity is equally matched with the challenge of creating a framework for analysis.  Acceleration of the cosmic expansion heralds a revolution in physics, if we can characterize and understand it.  To comprehend this new aspect of the universe, we must map the expansion not only at recent epochs, but encompass the early universe.  Once we have garnered sufficient understanding, the prize is answering the question of the fate of the universe, now unbound from the question of the cosmic geometry.  The good news is that we likely have at least 24 billion years to do so!   \\ack  I thank the Aspen Center for Physics for a superb working environment, and Los Alamos National Laboratory and the Santa Fe 07 workshop, University of Heidelberg, and the Dark Cosmology Centre and Niels Bohr Summer Institute, for hospitality. I gratefully acknowledge contributions by Georg Robbers, Ramon Miquel, and David Rubin of Figures~\\ref{fig:dlsss}, \\ref{fig:dustprior}, and \\ref{fig:doomsday} respectively.  Colleagues too numerous to mention helped form my thinking on this wide ranging topic, but I will single out Bob Wagoner for laying the foundations. This work has been supported in part by the Director, Office of Science, Department of Energy under grant DE-AC02-05CH11231."
    },
    "0801/0801.3125_arXiv.txt": {
        "abstract": "Over the last few years, several interesting observations were obtained with the help of solar Adaptive Optics (AO). In this paper, few observations made using the solar AO are enlightened and briefly discussed. A list of disadvantages with the current AO system are presented. With telescopes larger than 1.5m are expected during the next decade, there is a need to develop the existing AO technologies for large aperture telescopes. Some aspects of this development are highlighted. Finally, the recent AO developments in India are also presented. ",
        "introduction": "Sun provides an unique opportunity for greater understanding of physical processes happening in stellar atmospheres. The solar atmosphere is structured to very small scales which are dynamic in nature. Understanding the nature of these small-scale structures and dynamics may provide important clues to some puzzling and unanswered questions. It is now very well realised that the small-scale dynamics can influence the large-scale structures (DeRosa, 2005). In hydrodynamic condition, two important scales determine the structuring of the solar atmosphere: (i) Pressure scale height, and (ii) Photon mean free path. Both these scales correspond to a value of 70km (i.e., about 0.\"1) at the photosphere. However, smaller magnetic structures are seen in realistic numerical magneto-hydrodynamic simulations (Stein \\& Nordlund, 2006; Schussler \\& Vogler, 2006). Grid sizes used in these simulations are smaller than 0.\"1 and future simulations are aimed to achieve a grid size of 0.\"02 or better.  Few years before, resolving structures on the Sun to better than an arcsecond for longer durations were not possible even at good sites due to the Earth's atmosphere. The scenario changed with the advancement of the Adaptive Optics (AO) technology. Hence, modern telescopes thrive on improving the spatial, spectral, and temporal resolutions using an AO system. It may be stated that most of the recent large solar ground based telescope (aperture $>=$ 50cm) designs have an integrated AO system. The realisation of the day-time AO technology lagged behind the night-time due to the difficulties caused by the extended nature of the source as well as the poor seeing conditions prevail during the day-time (Rimmele, 2004; Keller, 2005).  The first solar AO system (Acton, 1992 \\& 1993) was developed at Lockheed and tested at the Dunn Solar Telescope (DST). This system worked with high contrast features like solar pores and hence severely limited the scientific use. The development of the correlating Shack-Hartmann wavefront sensor, for the low-order AO system at DST (Rimmele, 2000),  has changed the perspective of the solar AO. All solar AO systems currently in operation are based on the correlating Shack-Hartmann sensor. Following Keller (2005), Table~I lists the successful AO systems at different solar observatories around the globe. Figure~1 shows an example Sunspot image observed under good seeing condition using a high-order AO (HOAO) system operated at DST. A small portions of the bottom penumbral region is blown-up to show the high spatial resolution achieved with this observation.  \\begin{figure}[t] \\hspace*{0.2in} \\includegraphics[scale=0.22]{fig1a.ps} \\includegraphics[scale=0.22]{fig1b.ps} \\caption{Speckle reconstructed sunspot observed at the DST using the high-order AO (HOAO) system. Courtesy: F. Woeger, T. Rimmele, M. Komsa, C. Berst, S. Fletcher, \\& S. Hegwer: NSO/AURA/NSF.} \\end{figure}  \\begin{table}[h] {\\bf Table 1: A list of Successful AO Systems.} \\\\ {\\tiny \\begin{tabular}{|ccccccccc|} \\hline & & & & & & & & \\\\ AO & Act- & Sub- & Camera & Frame & Hardware & DM & Ref. & First  \\\\ & uator & aperture & & Rate & & & & Light \\\\ & & & & & & & & \\\\ \\hline & & & & & & & & \\\\ 76cm & 57 & 19 &  ---   & Analog & --- &  ---    & Acton & 1986 \\\\ DST & & & & & & & 1992 & \\\\ Lockheed & & & & & & & & \\\\ & & & & & & & & \\\\ 76cm&  97 & 24 &---& $<$1.6kHz & 24 & Xenitics & Rimmele & 1998 \\\\ DST & & & & & DSPs &Inc. & 2000 & \\\\ LOAO & & & & & & & & \\\\ & & & & & & & & \\\\ 48cm  &  19 & 19 & Dalsa & 955Hz & 566MHz & AOPTIX & Scharmer &  1999 \\\\ SVST & & & CA-D6& & Alpha &Tech. Inc. & 2000 & \\\\ & & & & & & & & \\\\ 76cm & 97 & 76 & Custom & 2.5kHz & 40 & Xenitics&Rimmele &\t2002 \\\\ DST & & & built & & DSPs &Inc. &2003 & \\\\ HOAO & & & CMOS & & & & & \\\\ & & & & & & & & \\\\ 70cm & 35 & 36 & Dalsa & 955Hz & 8X900MHz & Laplacian & von der Luhe & 2002 \\\\ VTT & & & CA-D6 & & Sun &Optics &2003 & \\\\ KAOS & & & & & & & & \\\\ & & & & & & & & \\\\ 1.5m & 37 & 120-200 & Dalsa & 955Hz & 1GHz &\tOkotech  & Keller &\t2002 \\\\ McMath & & & CA-D6 & & Pentium & &2003 & \\\\ Low Cost & & & & & & & & \\\\ & & & & & & & & \\\\ 97cm & 37 & 37 & Dalsa & 955Hz & 1.45GHz & AOPTIX & Scharmer & 2003 \\\\ SST & & & CA-D6 & &Athlon &Tech. Inc. &2003 & \\\\ & & & & & & & & \\\\ 65cm & 97 & 76 & Custom & 2.5kHz & 40 & Xenitics & Denker & 2004 \\\\ BBSO & & & Built & &DSPs &Inc. &2007 & \\\\ HOAO & & & CMOS & & & & & \\\\ & & & & & & & & \\\\ \\hline \\end{tabular} } \\end{table} ",
        "conclusions": ""
    },
    "0801/0801.4828_arXiv.txt": {
        "abstract": "Among more than 200 extrasolar planet candidates discovered to date, there is no known planet orbiting around normal binary stars.  In this paper, we demonstrate that microlensing is a technique that can detect such planets.  Microlensing discoveries of these planets are possible because the planet and host binary stars produce perturbations at a common region around center of mass of the binary stars and thus the signatures of both planet and binary can be detected in the light curves of high-magnification microlensing events. The ranges of the planetary and binary separations of systems for optimal detection vary depending on the planet mass.  For a Jupiter-mass planet, we find that high detection efficiency is expected for planets located in the range of $\\sim$ 1 AU -- 5 AU from the binary stars which are separated by $\\sim$ 0.15 AU -- 0.5 AU. ",
        "introduction": "Since the first discovery by \\citet{wolszczan92}, more than 200 extrasolar planet candidates have been discovered. Among the currently known extrasolar planet-hosting stars, approximately 20\\% are members of binaries or multiple systems \\citep{haghighipour06}. Nearly all of the planets in these systems revolve around only one of the stars. The only exception is the planetary system PSR B1620-26 \\citep{lyne88, backer93}, but this system consists of not normal stars but a pulsar and a white dwarf.   Besides the orbital configuration of the known planets in binaries where the planet revolves around one of the widely separated binary members, a planet can stay in a stable orbit if the planet is located well outside the binary stars and thus seeing the two stars as approximately a single object. However, such planets, which were imagined as a planet with two simultaneously rising suns in the {\\it Star Wars} saga (Tatooine), have not been seen in planet searches using the radial velocity technique because the search programs avoid short-period binary stars. The chance to detect such planets via the transit technique is also very low because these planets tend to have wide orbits.   Planets in binary stellar systems can be detected by using the microlensing technique. Such a possibility was first mentioned by \\citet{bennett99}.  They reported a candidate planet orbiting around a binary system from the observation and light-curve modeling of a microlensing event MACHO-97-BLG-41. However, their interpretation of this event is very special in the sense that the source trajectory is precisely aligned with the line connecting the planet and the center of mass of the binary. Besides, the interpretation of the signal was rejected because a rotating binary-lens model provides an excellent fit to the observed data \\citep{albrow00}. \\citet{lee08} also investigated the lensing signal of planets in binary systems. However, their analysis is restricted to planets orbiting around one of the binary members. In this paper, we investigate the feasibility of detecting planets with two simultaneously rising suns by using the microlensing technique.       \\begin{figure}[ht] \\epsscale{1.2} \\plotone{f1.eps} \\caption{\\label{fig:one} Lensing geometry of a planet revolving around close binary stars. The big dashed circle  represents the Einstein ring of a mass corresponding to the total mass of the binary stars. The angle $\\phi$ denotes the position angle of the planet as measured counterclockwise from the binary axis and $d_b$ and $d_p$ denote the projected separations between the binary stars and between the planet and the binary center of mass, respectively. The right panel shows the enlarged view of the region around the binary stars. The figure drawn in red color represents the lensing caustic produced by the lens system. }\\end{figure}     \\begin{figure*}[ht] \\epsscale{0.7} \\caption{\\label{fig:two} Lensing magnification patterns of triple lens systems composed of a planet revolving around binary stars. The labels above and on the right represent the projected separations between binary stars ($d_b$) and between the planet and the center of mass of the binary ($d_p$), respectively. In each map, the coordinates are centered at the position of the center of mass of the binary stars and the $x$-axis is aligned with the binary axis. The heavier star is on the right.  The planets have a common position angle of $\\phi= 45^\\circ$ as measured counterclockwise from the binary axis (see Figure~\\ref{fig:one}).  The grey-scale is drawn such that brighter tone represents the region of higher magnification. The figures drawn in solid curves represent the lensing caustics. The presented patterns are for a Saturn-mass planet revolving around binary stars with a total mass of $0.5\\ M_\\odot$ and a companion/primary mass ratio of 0.5. The light curves resulting from the source trajectories marked by straight lines with arrows in the individual panels are presented in the corresponding panels of Figure~\\ref{fig:three} (thick solid curves in black color).  The panels blocked by thick black solid lines show the cases where the planet dominates the perturbation pattern, while the panels blocked by thick white solid lines show the opposite cases where the binary stars dominate the pattern. }\\end{figure*} ",
        "conclusions": "We investigated the feasibility of detecting planets orbiting close binary stars by using the microlensing technique.  We presented the channel for which one can identify the signatures of such planetary systems.  We illustrated the signatures for various planet-binary configurations and explained the tendencies of the signatures.  We also estimated the ranges of the binary and planetary separations for which the microlensing method is efficient in detecting these planetary systems.  Discovering planets orbiting around close binary stars is difficult by using the radial velocity or transit technique.  Therefore, microlensing is an important technique that can provide information about these planetary systems, which is very important for our understanding of planet formation."
    },
    "0801/0801.1120_arXiv.txt": {
        "abstract": "We consider axisymmetric relativistic jets with a toroidal magnetic field and an ultrarelativistic equation of state, with the goal of studying the lateral structure of jets whose pressure is matched to the pressure of the medium through which they propagate. We find all self-similar steady-state solutions of the relativistic MHD equations for this setup. One of the solutions is the case of a parabolic jet being accelerated by the pressure gradient as it propagates through a medium with pressure declining as $p(z)\\propto z^{-2}$. As the jet material expands due to internal pressure gradients, it runs into the ambient medium resulting in a pile-up of material along the jet boundary. The magnetic field acts to produce a magnetic pinch along the axis of the jet. The relative importance of the two components -- the boundary pile-up and the magnetic pinch -- depends on the magnetization parameter of the jet and on the total energy it carries relative to the ambient pressure. Such jets can be in a lateral pressure equilibrium only if their opening angle $\\theta_j$ at distance $z$ is smaller than about $1/\\gamma$, where $\\gamma$ is the characteristic bulk Lorentz-factor at this distance; otherwise, different parts of the jet cannot maintain causal contact. We construct maps of optically thin synchrotron emission from our self-similar models to illustrate the observational features of such jets. We suggest that the boundary pile-up may be the reason for the limb-brightening of the sub-parsec jet of M87. We find that if the synchrotron emissivity falls with the distance from the jet axis (as in the magnetically pinched case), the polarization fraction rises toward the edge, as seen in 3C273 and Mkn501. Projection effects and the emissivity pattern of the jet have a strong effect on the observed polarization signal, so the interpretation of the polarization data in terms of the geometry of magnetic fields is rather uncertain. For example, jets with toroidal magnetic fields display the `spine-sheath' polarization angle pattern seen in some BL Lac objects. ",
        "introduction": "\\label{sec_intro}  Accretion onto black holes is frequently accompanied by collimated outflows. Although many of the physical properties of these flows proved difficult to measure, strong evidence exists that they are relativistic in a variety of environments. This evidence includes the one-sided appearance of jets and observations of apparent superluminal motions in Galactic black-hole binaries \\citep{mira99} and in active galactic nuclei (AGN; e.g., \\citealt{jors05}), as well as the rapid gamma-ray variability in AGN \\citep{dond95, ahar05}.  Jets are also magnetized. The polarization fraction of synchrotron radiation from AGN jets approaches its theoretical maximum value in some objects (e.g., \\citealt{zava05, list05, push05}), indicating that the magnetic field is highly ordered. The maps of rotation measure of 3C273 \\citep{asad02, zava05} and of several BL Lac objects \\citep{gabu04} show strong gradients of the rotation measure across the jet which are expected if the toroidal component dominates. The polarization position angles indicate toroidal fields in more than half of all BL Lac objects \\citep{gabu00, list05, push05}, but such observations may be subject to projection effects and relativistic aberration \\citep{pari03, lyut05}. The strength of the magnetic field is poorly constrained, so its true dynamical importance cannot be evaluated from observations.  The current theoretical paradigm is that most relativistic jets are probably accelerated by large-scale electromagnetic stresses \\citep{blan82, li92, cont94}. Acceleration is rapid until the flow passes the fast magnetosonic surface and attains rough equipartition between kinetic energy flux and Poynting flux. Further increase in the jet Lorentz-factor tends to be more gradual, but eventually it is possible for most of the Poynting flux to be converted to kinetic energy \\citep{bege94, vlah03a, vlah03b, besk06, komi07}. Ultimately, the flow is confined laterally by either the inertia of the disk at the base of the jet, in which case the collimating force is transmitted along the jet by magnetic tension, or the pressure of the ambient medium along its walls \\citep{bege95}.  In this paper we investigate the interior of laminar ultrarelativistic jets whose boundary pressures are matched to their surroundings. The flux-freezing condition implies that the poloidal component of the magnetic field decays faster than the toroidal component in an expanding jet \\citep{bege84}, and we consider the structure of the jet far enough from its initial acceleration site that the toroidal component dominates. \\citet{cont95} showed formally that at large distances down the jet the poloidal magnetic field is dynamically unimportant and discussed the critical points of such solutions.  As forces exterior to the jet gradually collimate it, curving the outermost streamlines slightly toward the axis, some of the material in the jet interior runs into the wall and piles up. Thus, jets may develop an edge-brightened or `hollow-cone' structure, with much of the jet energy, momentum and emissivity concentrated into a narrow sheath along the interface between the jet and its surroundings. The pinching effect due to magnetic tension has the opposite effect, tending to focus jet energy toward the axis \\citep{usty95, usty00, ston00, komi07}. We find that for some functional dependencies of the ambient medium pressure $p_a(z)$ such jets can develop self-similar structures that admit a wide range of prescribed entropy and toroidal magnetic field distributions, and we investigate the interplay of pressure and magnetic field for the self-similar solutions.  In Section \\ref{sec_solutions} we describe the basic equations, our assumptions and the family of self-similar solutions that we find. In Section \\ref{sec_sync} we discuss the appearance of the synchrotron radiation from our solutions. We conclude in Section \\ref{sec_conclusions}. In Section \\ref{sec_solutions}, unless otherwise noted we use units in which $c=1$ and the magnetic field is redefined to eliminate factors of $4 \\pi$ from Maxwell's equations. The magnetic field in Gaussian units is $\\sqrt{4 \\pi}$ times our magnetic field ${\\bf B}$ and the current (and its density) in Gaussian units is $1/\\sqrt{4\\pi}$ times our current (and density). Our Greek indices run from 0 to 3 and the Latin ones from 1 to 3, and summation is assumed over repeating indices (using the Minkowski metric for Greek indices). We use Gaussian and CGS units in Section \\ref{sec_sync}. ",
        "conclusions": "\\label{sec_conclusions}  We have considered axisymmetric relativistic MHD jets with an ultrarelativistic equation of state and toroidal magnetic fields. We have then found all self-similar solutions of relativistic MHD equations in this setup assuming that the jet propagates in a pressure equilibrium with the surrounding matter. There are three geometrically self-similar solutions which occur in the medium with pressure scaling as $P(z)\\propto z^0$ (cylindrical jet), $\\propto z^{-2}$ (parabolic jet) and $\\propto z^{-4}$ (conical jet). The lateral structure of the jet is determined by the balance of the radial pressure gradient, the tension of the toroidal magnetic field lines and the centrifugal force. The last term is relevant only when the streamlines are curved. Our models represent a special case of a more general situation in which interactions between the jet and the medium lead to shocked layers \\citep{brom07}; in our case, solutions with no shocks are possible.  While the self-similarity condition is quite restrictive, in that only a limited set of geometries can be studied, the lateral structure of the parabolic jet displays some generic features which we investigate in detail. The pressure gradients act to push the jet material against the walls of the jet, whereas the magnetic field tension acts to pile the material closer to the axis. The relative importance of the two effects is determined by the physical parameters of the jet, such as its magnetization $\\beta_B$ and the total energy. The pinched core component is stronger for a higher value of $\\beta_B$, whereas the pile-up component is stronger for a jet with a higher Lorentz-factor (or higher total energy, to put it another way).  Our models describe jets whose Lorentz-factors are smaller than the critical value $\\gamma_{\\rm cr}\\simeq 1/\\theta_j$, where $\\theta_j$ is the opening angle of the jet. We find that for larger Lorentz-factors $\\gamma\\gg\\gamma_{\\rm cr}$ either the formal solution of the lateral equilibrium equation does not exist at all, or it exists in the form of a thin pile-up of the material along the edge of the jet. The condition $\\gamma<\\gamma_{\\rm cr}$ is equivalent to requiring that the jet maintains causal contact in the lateral direction as it propagates.  Edge pile-up is a generic feature of a jet with curved streamlines and it might explain edge-brightening of the jet in M87 in the acceleration region \\citep{ly07, kova07}. We also find that even purely toroidal magnetic fields can produce the `spine-sheath' polarization structure \\citep{push05} when projection effects are taken into account. We find that for core-dominated jets the polarization fraction is expected to be near zero on the projected jet axis and rise toward the edge, as observed in 3C273 \\citep{zava05} and Mkn501 (B1652+398; \\citealt{push05}). The orientation of the polarization position angle is determined not just by the geometric projection effects, but also by the emissivity profile. For emissivity rising toward the axis, the polarization perpendicular to the axis may become the dominant orientation even for purely toroidal magnetic fields. These effects significantly complicate the interpretation of the observed polarization position angle in terms of the intrinsic direction of the magnetic field."
    },
    "0801/0801.1250_arXiv.txt": {
        "abstract": "{Observations with the {{\\it Hinode\\/}} spectro-polarimeter have revealed strong horizontal internetwork magnetic fields in the quiet solar photosphere.} {We aim at interpreting the observations by means of results from numerical simulations. } {Radiative MHD simulations of dynamo action by near-surface convection are analyzed with respect to the relation between vertical and horizontal magnetic field components. } {The dynamo-generated fields show a clear dominance of the horizontal field in the height range where the spectral lines used for the observations are formed. The ratio between the averaged horizontal and vertical field components is consistent with the values derived from the observations. This behavior results from the intermittent nature of the dynamo field with polarity mixing on small scales in the surface layers. } {Our results provide further evidence that local near-surface dynamo action contributes significantly to the solar internetwork fields.} ",
        "introduction": "The ubiquitous existence of small-scale `internetwork' magnetic fields of mixed polarity in the so-called quiet solar photosphere is strongly indicated by various observational diagnostics \\citep[e.g.,][and further references therein]{Khomenko:etal:2003, Lites:Socas-Navarro:2004, Trujillo:etal:2004}.  Recent high-resolution space-borne observations with the spectropolarimeter of the Solar Optical Telescope aboard the {\\it Hinode\\/} satellite have considerably strengthened the case for internetwork fields and, furthermore, have revealed that the measured internetwork flux is dominated by strongly inclined, almost horizontal magnetic field \\citep{Lites:etal:2007, Orozco-Suarez:etal:2007}. Considerable amounts of highly time-dependent horizontal magnetic flux have also been found in ground-based observations with lower spatial resolution \\citep{Harvey:etal:2007}.  The ubiquity of the small-scale, mixed-polarity internetwork field suggests a local origin of at least a significant part of the measured flux.  Recently, we have demonstrated by means of radiative magneto-convection simulations that local dynamo action by near-surface convective flows is a possible source for the internetwork flux \\citep{Voegler:Schuessler:2007}. Here we show that the spatial structure of the dynamo-generated field provides a natural explanation for the observed dominance of the horizontal field component in the middle photosphere. ",
        "conclusions": "A direct comparison of the simulation results with the observations is not possible because the individual values for the averaged fields in the simulation are both (by about a factor three) smaller in the relevant height range than the values for the `apparent' derived from the observation, the discrepancy probably becoming even more severe when the actual spatial resolution of the observations is taken into account. This is not particularly surprising since the magnetic Reynolds number of the simulation is still orders of magnitude smaller than the actual value in the corresponding solar layers, so that the saturation level of our simulation is probably considerably below the level to be expected for the real Sun. In fact, a preliminary simulation with a roughly doubled Reynolds number of about 5000 shows an increase of the magnetic energy by a factor of about 1.7 with respect to the case shown here. Interestingly, it turns out that the optical depth profiles of the averaged field strengths in this case can be very well approximated by just multiplying the curves shown in Fig.~\\ref{fig:tauprof} by $\\sqrt{1.7}$, meaning that the main difference between the simulations reduces to a simple scale factor for the field strength. Together with the fact that the observed internetwork fields are predominantly weak and their energy is significantly smaller than the kinetic energy of the convective motions, this suggests that the general nature of the dynamo-generated field may not be significantly different from the case shown here, apart from a higher overall amplitude. In particular, we expect that the ratio of the average horizontal and vertical field is not strongly affected by the amplitude of the dynamo-generated field. Of course, these assertions need to be demonstrated by further simulations with higher Reynolds numbers.  The clear dominance of the horizontal field in the mid photosphere seems to be a rather specific property of the strongly intermittent field generated by near-surface turbulent dynamo action. Accordingly, the dynamo simulation of \\citet[][with a closed bottom boundary and a local treatment of radiative transfer]{Abbett:2007} also exhibits strong horizontal field in the photospheric layers. On the other hand, models with an imposed net vertical flux \\citep[e.g.,][]{Voegler:etal:2005} or our recent simulations of the decay of a granulation-scale mixed-polarity field (at magnetic Reynolds numbers below the threshold of dynamo action) do not show this behavior; in these cases, the intricate small-scale mixing of polarities that is characteristic for the dynamo does not dominate the field structure. The rapid decay with height of such a dynamo-generated field is also consistent with the apparent lack of strong horizontal field in the chromosphere \\citep{Harvey:etal:2007}.  What are the alternatives to near-surface dynamo action?  `Shredding' of pre-existing magnetic flux (remnants of bipolar magnetic regions) cannot explain the large amount of observed horizontal flux since the turbulent cascade does not lead to an accumulation of energy (and generation of a spectral maximum) at small scales. On the other hand, such a behavior is typical for turbulent dynamo action. Flux emergence from the deeper convection zone in the form of granule-sized small bipoles would have to proceed such a high rate in order to maintain the ubiquitous strong horizontal fields that it probably would not have gone undetected in the past \\citep[see, however,][]{Centeno:etal:2007}. The sporadic appearance of horizontal internetwork fields (HIFs) described by \\citet{Lites:etal:1996} and interpreted as small-scale flux emergence events seems to be unsufficient to explain the ubiquitous horizontal field now found with {\\it Hinode}. On the other hand, flux recycling of an overall background flux by granulation probably represents a significant source of horizontal field in network and plage regions. Emergence of extended horizontal field strands in granules as observed by \\cite{Ishikawa:etal:2007} is not seen in local dynamo simulations.  In the real Sun, probably all three sources, i.e., dynamo, shredded fields, and small-scale flux emergence from deeper layers, contribute to the internetwork flux in unknown amounts.  In any case, the strong horizontal fields in the quiet photosphere inferred by the observations indicate that the source of these fields at the solar surface is a mixed-polarity field whose energy is mostly contained in those spatial scales where the dynamo-generated flux resides.  Therefore, the observational results obtained with the {\\it Hinode\\/ SP} together with the analysis presented here provides strong indication that surface dynamo action represents a significant source for the internetwork field in the solar photosphere."
    },
    "0801/0801.4643_arXiv.txt": {
        "abstract": "{This is the fourth in a series of papers that deal with angular momentum transport by internal gravity waves in stellar interiors.} {Here, we want to examine the potential role of waves in other evolutionary phases than the main sequence.} {We study the evolution of a $3\\,M_\\odot$ Population\\,I model from the pre-main sequence to the early-AGB phase and examine whether waves can lead to angular momentum redistribution and/or element diffusion at the external convection zone boundary.} {We find that, although waves produced by the surface convection zone can be ignored safely for such a star during the main sequence, it is not the case for later evolutionary stages. In particular, angular momentum transport by internal waves could be quite important at the end of the sub-giant branch and during the early-AGB phase. Wave-induced mixing of chemicals is expected during the early-AGB phase.} {} ",
        "introduction": "In recent years, several authors studied the impact of internal gravity waves (IGWs) in a variety of main sequence stars. These waves were initially invoked as a source of mixing in stellar interiors in low-mass stars with an extended surface convection zone (Press~1981; Garc\\'\\i a L\\'opez \\& Spruit~1991; Schatzman~1993; Montalb\\'an~1994) and also as an efficient process in the synchronization of massive binary stars (Goldreich \\& Nicholson~1989). More recently, it was suggested that IGWs may play a role in braking the solar core (Schatzman~1993; Zahn, Talon, \\& Matias~1997; Kumar \\& Quataert~1997). This idea was confirmed first in static models (Talon, Kumar, \\& Zahn~2002) and recently in the complete evolution of solar-mass models, evolved all the way from the pre-main sequence to 4.6\\,Gy (Charbonnel \\& Talon~2005).  All these authors find that IGWs are easily excited, and a similar conclusion is reached by studies of convection on top of a stably stratified layer in 2--D and 3--D hydrodynamic numerical simulations (\\eg Hurlburt et al.~1986,~1994; Andersen~1994; Nordlund et al.~1996; Kiraga et al.~2003; Dintrans et al.~2005; Rogers \\& Glatzmeier~2005a,b). These waves have a strong impact on stellar evolution, especially through their effect on the rotation profile. Through differential filtering, IGWs indeed play a major role in the redistribution of angular momentum in stars,  which determines the extent and magnitude of rotation-induced mixing (Talon \\& Charbonnel~2005, hereafter TC05).  In the spirit of applying a unique set of physical principles to stellar evolution, one needs to assess the impact of such waves on stars of all masses and at various evolutionary stages. It is our purpose in this series of papers to consistently examine the full Hertzsprung-Russel diagram (HRD) for determining when such waves are efficiently emitted and how they can affect stars when their complete rotational history is being considered.  In Talon \\& Charbonnel~(2003, hereafter Paper~I; see also Talon \\& Charbonnel 1998), we showed how the appearance of IGWs in solar-metallicity main sequence stars with an effective temperature $T_{\\rm eff} \\apprle 6700~{\\rm K}$ (i.e., when the surface convection zone becomes substantial) can explain the existence of the lithium dip\\footnote{The so-called lithium dip, observed in all galactic open clusters older than $\\approx 200~{\\rm Gyr}$, as well as in the field, refers to a strong lithium depletion in a narrow region of $\\approx 300~{\\rm K}$ in effective temperature, centered around $T_{\\rm eff} = 6700~{\\rm K}$ (e.g. Wallerstein et al.~1965; Boesgaard \\& Tripicco~1986; Balachandran~1995).} in stars undergoing rotational mixing.  In Talon \\& Charbonnel~(2004, hereafter Paper~II), we examined the IGW generation in Population~II main sequence stars. We showed that, along the lithium plateau\\footnote{The lithium plateau refers to the constant Li abundance of Population~II dwarf stars above $T_{\\rm eff} \\approx 5500~{\\rm K}$ (e.g. Spite \\& Spite~1982).}, the net angular momentum luminosity of IGWs is constant and high enough to enforce quasi solid-body rotation similar to that of the Sun in these stars. We proposed that this behavior could play a major role in explaining the constancy of the lithium abundance in the oldest dwarf stars of our Galaxy (see also a discussion in Charbonnel \\& Primas~2005).  Here, we wish to look at other evolutionary stages, especially the pre-main sequence (PMS) and the more advanced stages for intermediate-mass stars. We focus in particular on stars in which IGWs generation by the surface convection zone is limited on the main sequence, \\ie, Pop\\,I star originating from the left side of the Li dip. To do so, we follow the evolution of a $3\\,M_\\odot$, $Z=0.02$ star from the PMS up to the early asymptotic giant branch (early-AGB)\\footnote{The peculiar case of thermally pulsing AGB (TP-AGB) stars will be presented elsewhere.}. We estimate at which stages waves are efficiently generated in the outer convection zone. This is a prerequisite for evaluating their impact on the rotation profile of the corresponding stellar models. We also look at the existence of a shear layer oscillation (or SLO) as a direct source of turbulence and mixing at the convective boundary.  We begin in \\S\\,\\ref{sec:popI} with a description of the evolution of relevant characteristics for our $3\\,M_\\odot$, $Z=0.02$ star. In \\S\\,\\ref{sec:formalism}, we recall the main aspects of the formalism we used to evaluate the impact of IGWs on stellar evolution. The following sections are devoted to discussing the pre-main sequence (\\S\\,\\ref{sec:pms}), the main sequence and the sub-giant branch (\\S\\,\\ref{sec:ms}), the red-giant branch (\\S\\,\\ref{sec:giants}), and the clump and the early-AGB (\\S\\,\\ref{sec:AGB}).  \\begin{figure}[t] \\centerline{ \\psfig{figure=8620fig1.eps,height=6.5cm,angle=-90} } \\caption{Evolution of our $3\\,M_\\odot$ Pop\\,I model in the HR diagram, from the PMS to the early-AGB. Numbers correspond to selected evolutionary points that are listed in Table~\\ref{tab:popI}. For clarity, the PMS is drawn with a dashed line, and the corresponding models are identified by squares, while the star symbol corresponds to the appearance of the convective core. The circles correspond to evolutionary points on the main sequence, sub-giant, and first ascent giant branches, while triangles refer to evolutionary points on the clump and on the early-AGB. \\label{fig:HRI}} \\end{figure}  \\begin{figure}[t] \\centerline{ \\psfig{figure=8620fig2.ps,height=6.5cm,angle=-90} } \\caption{Evolution of the temperature at the boundaries of convection zones for the evolutionary points selected for the $3\\,M_\\odot$ Pop\\,I star and given in Table~\\ref{tab:popI}. Core and surface (i.e., $T_{\\rm eff}$) temperatures are also shown. The evolutionary phases are indicated. \\label{fig:zc3I}} \\end{figure}  \\begin{figure*}[t] \\centerline{ \\psfig{figure=8620fig3.ps,height=7.5cm,angle=-90} } \\caption{Evolution of the surface convection zone properties for the $3\\,M_\\odot$ Pop\\,I evolutionary points from Table~\\ref{tab:popI}. It shows the thermal conductivity $K_T$ (in ${\\rm cm^2\\,s^{-1}}$) at the convective boundary, the convective flux $F_c=\\rho v_c^3$ (in ${\\rm g\\,cm^{-3}}$) taken $0.5~H_P$ into the convection zone, the mass of the convection zone, the spherical harmonic degree corresponding to the convective scale $\\ell_c$, the convective frequency $\\nu_c$ (in $\\mu$Hz), the typical convective timescale $\\tau_L=\\alpha H_P/v_c$, as well as the total energy luminosity (in ${\\rm g\\,cm^2\\,s^{-3}}$) prior to any damping in waves (this is for waves with $\\nu < 3.5\\,\\mu {\\rm Hz}$ and $1<\\ell<60$). \\label{fig:zc2I}} \\end{figure*} ",
        "conclusions": "In this paper, we examined when and how internal gravity waves generated by surface convection zones are expected to have an impact on stellar models that only have shallow surface convection zones during the main sequence. We find that several evolutionary stages can experience angular momentum redistribution by IGWs: \\begin{itemize} \\item During the pre-main sequence, just as the star leaves the Hayashi track, waves are efficiently excited and filtered. While this has no impact on light element abundances and only has a limited effect on fast rotators (see \\S\\,\\ref{sec:pms}), it could strongly modify the main-sequence rotation profile of slow rotators; \\item At the end of the sub-giant branch (\\ie, at the boundary of the Hayashi track), angular momentum redistribution is once more quite efficient and could totally change the angular momentum profile that has been established during the main sequence and the sub-giant phase contraction; \\item During the early-AGB, IGW generation again becomes substantial: mixing can be expected to occur both at the convective boundary, due to the presence of an SLO and at the hydrogen-burning-shell boundary, due to the enhanced damping of waves caused by the increase in the Brunt-V\\\"ais\\\"al\\\"a frequency. \\end{itemize} Stellar models including meridional circulation, shear turbulence, and wave-induced mixing are underway to quantify this impact (Charbonnel \\& Talon, in preparation)."
    },
    "0801/0801.2779_arXiv.txt": {
        "abstract": "We present a 1.1~mm wavelength imaging survey covering 0.3~deg$^2$ in the COSMOS field. These data, obtained with the AzTEC continuum camera on the James Clerk Maxwell Telescope (JCMT), were centred on a prominent large-scale structure over-density which includes a rich X-ray cluster at $z \\approx 0.73$. A total of 50 millimetre galaxy candidates, with a significance ranging from 3.5--8.5$\\sigma$, are extracted from the central 0.15~deg$^2$ area which has a uniform sensitivity of $\\sim$1.3~mJy/beam. Sixteen sources are detected with $S/N \\ge 4.5$, where the expected false-detection rate is zero, of which a surprisingly large number (9) have intrinsic (de-boosted) fluxes $\\ge 5$~mJy at 1.1~mm. Assuming the emission is dominated by radiation from dust, heated by a massive population of young, optically-obscured stars, then these bright AzTEC sources have FIR luminosities $> 6\\times 10^{12}~\\rm L_{\\odot}$ and star formation-rates $> 1100~\\rm M_{\\odot}/yr$. Two of these nine bright AzTEC sources are found towards the extreme peripheral region of the X-ray cluster, whilst the remainder are distributed across the larger-scale over-density. We describe the AzTEC data reduction pipeline, the source-extraction algorithm, and the characterisation of the source catalogue, including the completeness, flux de-boosting correction, false-detection rate and the source positional uncertainty, through an extensive set of Monte-Carlo simulations. We conclude with a preliminary comparison, via a stacked analysis, of the overlapping MIPS 24~\\micron~data and radio data with this AzTEC map of the COSMOS field. ",
        "introduction": "\\label{sec:int} A decade after the discovery of a population of extremely luminous, high-redshift dust-obscured galaxies detected by their sub-millimetre and millimetre wavelength emission \\citep{smail97,hughes98,barger98}, over 200 sub-mm/mm galaxies (hereafter SMGs) have been detected with signal to noise ratio $\\ge 4$ in blank field surveys \\citep[e.g.,][]{borys03,greve04,laurent05,coppin06} and in surveys towards moderate redshift clusters designed to probe the faintest SMGs via lensing \\citep[e.g.,][]{smail98,smail02,chapman02}. Their high FIR luminosities (L$_{FIR}\\sim 10^{12-13}~\\rm L_\\odot$) and inferred star formation rates \\citep[SFR $\\gg 100~\\rm M_\\odot/yr$,][]{smail97,hughes98,barger98} suggest that these galaxies are high-redshift analogs to the local ULIRG population \\citep{sanders96}, and that they may be the progenitors of the massive elliptical population observed locally.  Until recently, the relatively modest mapping speeds of SCUBA \\citep[850~\\micron,][]{holland99} on the 15-m James Clerk Maxwell Telescope (JCMT), MAMBO \\citep[1.2~mm,][]{kreysa98} on the Institut de Radio Astronomie Millimetrique (IRAM) 30-m telescope and Bolocam \\citep[1.1~mm,][]{glenn98,haig04} on the 10-m Caltech Submillimeter Observatory (CSO), have restricted SMG surveys to $< 300$~arc-min$^2$ in size, limiting our understanding of the brightest, rarest SMGs and resulting in wide variations in the derived number counts as a result of small number statistics and cosmic variance \\citep[e.g.,][]{chapman02,smail02,scott02,borys03}. With new emphasis on large ($> 300$~arc-min$^2$) sub-mm/mm blank field surveys, \\citep{greve04,laurent05,mortier05,bertoldi07}, an accurate characterisation of the bright end of the SMG number counts and the mean properties of the SMG population is now becoming possible \\citep[e.g.,][]{coppin06}.  We surveyed a 0.15~deg$^2$ region within the COSMOS field \\citep{scoville07a} with uniform sensitivity at 1.l~mm with the AzTEC camera \\citep{wilson08} on the JCMT. The AzTEC survey field is centred on a prominent large-scale structure as traced by the galaxy density \\citep{scoville07b}, including a massive galaxy cluster at $z=0.73$ (Figure 1). This AzTEC map has no overlap with the MAMBO/COSMOS survey \\citep{bertoldi07} and only a small amount of overlap with the shallower Bolocam survey (J. Aguirre, private communication). Both MAMBO and Bolocam surveys cover a low galaxy-density region of the COSMOS field, whilst our new AzTEC observations are designed to examine the impact of massive large-scale foreground structures on SMG surveys in order to provide a measure of the importance of cosmic variance in the observed source-density at millimetre wavelengths.  In this paper we present the AzTEC mm survey of the COSMOS field, including the data reduction and source catalogue. Since this is the first in a series of papers describing the surveys completed by AzTEC on the JCMT, we provide an extensive description of the data analysis pipeline used to extract sources from AzTEC maps.  The JCMT observations, pointing, and calibration strategy are described in \\S~\\ref{sec:obs}. A detailed description of the data reduction algorithm is given in \\S~\\ref{sec:dat}.  In \\S~\\ref{sec:scat}, we present the AzTEC map and source catalogue, followed by a discussion of simulations used to determine flux boosting, false-detection rate, completeness, and source positional uncertainty in the map in \\S~\\ref{sec:sim}. A preliminary comparison of the mm sources to the radio and MIPS 24~\\micron~populations is made in \\S~\\ref{sec:multi_compare}, and we discuss the contribution of AzTEC sources to the Cosmic Infrared Background in \\S~\\ref{sec:cib}.  The large number of bright SMGs identified in the AzTEC/COSMOS field strongly suggests a bias in the number density introduced by the known large-scale structure that is present in the map. A detailed treatment of this analysis is beyond the scope of this paper and is deferred to Paper II (Austermann et al., in prep.).  The multi-wavelength imaging data from the HST/ACS, \\textit{Spitzer} IRAC and MIPS, as well as deep radio imaging from the VLA is particularly valuable for identifying and studying the nature of the SMGs identified by AzTEC. We will present a complete study of the multi-wavelength properties of the SMGs detected in the COSMOS field in Paper III.  We assume a flat $\\Lambda$CDM cosmology with $\\Omega_M = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $H_0 = 73$~km~s$^{-1}$~Mpc$^{-1}$ throughout.  \\begin{figure*} \\begin{center} \\includegraphics[width=17.5cm]{figure1col.eps} \\caption{{\\bf Left:} The galaxy density map from \\citet{scoville07b}, with the boundaries of the AzTEC, Bolocam, and MAMBO mm surveys within the COSMOS field indicated. The location of the $z = 0.73$ cluster environment is identified by the dashed circle. {\\bf Right:} The AzTEC/COSMOS map with $\\ge 3.5 \\sigma$ source candidates identified by circles with diameters equal to twice the AzTEC FWHM on the JCMT. The map has been trimmed to the ``75\\% coverage region'' and has an average rms noise level of 1.3~mJy/beam and an area of 0.15~deg$^2$. The signal map has been Wiener filtered for optimal identification of sources as described in \\S~\\ref{sec:wiener}.} \\label{fig:survey_map} \\end{center} \\end{figure*} ",
        "conclusions": "\\label{sec:con} We have imaged a 0.15~deg$^2$ region within the COSMOS field with AzTEC, a new mm-wavelength camera, with uniform sensitivity of 1.3~mJy/beam at 1.1~mm. We have identified 50 source candidates in the AzTEC/COSMOS map with signal to noise ratio $\\ge$ 3.5, 16 of which are detected with $S/N \\ge 4.5$, where the expected number of false-detections is zero. Seven of the $\\ge 5\\sigma$ source candidates have been followed up and confirmed with SMA imaging \\citep{younger07}. The sources are spread throughout the field, with only 3 located in the $z = 0.73$ cluster environment. Our catalogue is 50\\% complete at an intrinsic flux density of 4~mJy, and is 100\\% complete at 7~mJy.  The positional uncertainty of these AzTEC sources due to random and confusion noise is determined through simulations which show that sources with $S/N \\ge 3.5$ have $\\ge$ 80\\% probability of being detected within 4.5\\arcsec of their true location.  Comparing our $\\ge 3.5\\sigma$ source candidate list with the radio source catalogue of \\citet{schinnerer07}, we find that the fraction of AzTEC sources with potential radio counterparts is 36\\% and is consistent with that found in the SCUBA/SHADES survey \\citep{ivison07} at similar flux levels. From averaging the AzTEC map flux at the locations of the radio and MIPS~24~\\micron~\\citep{sanders07} source positions, we statistically detect the faint mm emission (below our detection threshold) of radio and MIPS~24~\\micron~sources and thereby demonstrate that errors in the mean astrometry of our map arising from the pointing model are small ($<2$\\arcsec). Estimates of the resolved fraction of the millimetre CIB due to these radio and mid-IR galaxy populations is 7 $\\pm$ 1\\% and 21 $\\pm$ 3\\% respectively.  The AzTEC/COSMOS field samples a region of high galaxy over-density compared to the regions imaged with MAMBO and Bolocam, and our AzTEC/COSMOS map contains a large number of very bright mm sources (9 with corrected flux density $\\ge 5$~mJy, where 4--5 are expected for an unbiased field). We will present a complete analysis of the relationship between the SMG population and the foreground galaxy population in Paper II of this series.  The availability of extensive high quality multi-wavelength data from the radio to the X-ray makes the follow-up analysis of the detected sources readily possible and will allow us to study the nature of these sources. A full analysis of the multi-wavelength properties of the sources detected in this survey will be presented in Paper III."
    },
    "0801/0801.0898_arXiv.txt": {
        "abstract": "On 10 March 2001 the active region NOAA 9368 produced an unusually impulsive solar flare in close proximity to the solar limb. This flare has previously been studied in great detail, with observations classifying it as a type 1 white-light flare with a very hard spectrum in hard X-rays. The flare was also associated with a type II radio burst and coronal mass ejection. The flare emission characteristics appeared to closely correspond with previous instances of seismic emission from acoustically active flares. Using standard local helioseismic methods, we identified the seismic signatures produced by the flare that, to date, is the least energetic (in soft X-rays) of the flares known to have generated a detectable acoustic transient. Holographic analysis of the flare shows a compact acoustic source strongly correlated with the impulsive hard X-ray, visible continuum, and radio emission. Time-distance diagrams of the seismic waves emanating from the flare region also show faint signatures, mainly in the eastern sector of the active region. The strong spatial coincidence between the seismic source and the impulsive visible continuum emission reinforces the theory that a substantial component of the seismic emission seen is a result of sudden heating of the low photosphere associated with the observed visible continuum emission. Furthermore, the low-altitude magnetic loop structure inferred from potential--field extrapolations in the flaring region suggests that there is a significant inverse correlation between the seismicity of a flare and the height of the magnetic loops that conduct the particle beams from the corona. ",
        "introduction": "\\label{S-Introduction}  Recent developments in the study of flare acoustic emissions \\citep{donea05,donea06,moradi07,Martinez-Oliveros2006} have bolstered the view that seismic emission from flares offers major new insights both into flare physics and helioseismology, ranging from a greatly improved understanding of flare dynamics and kinematics, to an understanding how seismic emission is generated differently by turbulence in magnetic subphotospheres from inside the quite Sun.  Sunquakes emanate from compact sources that take only a small fraction of the energy released by flares. The surface manifestation of these sources appears as circular (or near-circular) waves propagating outward from the solar surface, approximately 20\\,--\\,60 minutes after the impulsive phase of the flare. \\citet{donea05} considered the possibility that relatively weak flares might be able to produce sunquakes and that acoustically active flares may indeed be more common than previously thought. This was confirmed soon after by \\citet{betal2006} following comprehensive helioseismic observations of flares using helioseismic holography and the data from Michelson Doppler Imager (MDI) onboard the Solar and Heliospheric Observatory (SOHO).  While the majority of the flare acoustic transients discovered to date have been released by the more energetic X-class flares, recently however, a number of strong acoustic emissions from M-class flares have been discovered. \\citet{donea06} analyzed the helioseismic properties of the strong seismic transient produced by the M9.5 class flare of 9 September 2001, and very recently,  \\citet{Martinez-Oliveros2006} performed a comprehensive electromagnetic and acoustic analysis of the M7.4-class flare of 14 August 2004 -- the smallest flare known to have produced a detectable acoustic transient prior to the discovery of the sunquake reported in this paper. Their findings, along with those of \\citet{donea05} and \\citet{moradi07}, have identified a number of distinct observational characteristics that distinguish acoustically active flares from others:  \\begin{enumerate} \\item The sites of seismic emission generally coincide spatially with impulsive hard X-ray (HXR) and microwave (MW) emissions, suggesting a relation to thick-target heating of the chromosphere by energetic particles. \\item The sites of seismic emission similarly coincide spatially with impulsive continuum emission, suggesting acoustic emission associated with extra heating and ionization of the low photosphere. \\item The seismicity of the active region appears to be closely related to the heights of the coronal magnetic loops that conduct high-energy particles. \\end{enumerate}  In this paper, we will examine the last of the above characteristics -- introduced by \\citet{Martinez-Oliveros2006} -- in greater detail, with further evidence from the flare of 10 March 2001. The evidence we will provide reinforces the theory that shorter coronal loops are more likely to be conducive to a more rapid injection of trapped, high-energy electrons into the chromosphere at their footpoints. This enhances the magnitude and suddenness of the chromospheric heating that gives a rise to the intense visible continuum emission seen in all acoustically active flares. This mechanism appears to be a prospective source of the energy required to drive a powerful acoustic transient into the solar interior. ",
        "conclusions": "The standard flare scenario divides the flare process into a number of phases. In this scenario, the flare particles are accelerated to relativistic or super-relativistic velocities in the corona and injected into magnetic field loops whose footpoints are in active-region chromospheres. Inevitably, some particles are going to be trapped in the coronal magnetic field, while others, those in the magnetic loss cone, will precipitate directly into the chromosphere. Eventually the majority of the trapped particles are either scattered into the loss cone and precipitated, or thermalized (or both) by thermal plasma in the magnetic loop. In the case of the very sudden and impulsive flare of 10 March 2001, the hypothesis is that acceleration and injection of particles into the magnetic loop occurred in a short period of time \\citep{uddinetal2004}.  This kind of phenomenon can be described using the trapping and injection model proposed by \\citet{aschwanden2004}. In this model the rate of precipitation of charged particles into the chromosphere is controlled by the relaxation time of the system. The aperture angle of the loss cone changes with time, significantly opening as the magnetic field collapses to a more potential configuration. It is important to note that in this model, the time of acceleration and injection of the particles into the magnetic field are almost the same and relatively short compared with the precipitation and trapping time. It is fair to assume that if the relaxation time is short, the aperture of the loss cone also will change rapidly, allowing more particles to reach the chromosphere in a short period of time. This depends on efficient scattering of high-energy electrons into the expanded loss cone, which is greatly enhanced by chromospheric ablation of thermal plasma into short, low-lying loops.  Rapid evacuation of trapped electrons is suggested by observations of a rapid decay in non-thermal MW emission.  As a general rule, thermalization of particles in a magnetic trap is small compared to losses due to precipitation. Hence, high-energy electrons evacuated from the coronal loop in this way contribute to HXR bremsstrahlung emission substantially as well as their counterparts that were initially injected into the loss cone.   A much more complex model of particle precipitation, in which processes such as non-thermal excitation and ionization of hydrogen atoms, and non-thermal plasma heating (coulomb and ohmic) is explored by \\citet{ah1986}, \\citet{zk1993}, and \\citet{zharkova2007}. Interestingly, the later show the ohmic heating of the corona by the electron beams is so effective that the corresponding particle-induced downward propagating shocks are almost depleted of energy, leaving very little energy to reach the photosphere and induce any kind of seismic activity. Perhaps this is the explanation for why we did not see any seismic sources at the location of the $\\mathrm H\\alpha$ K2 kernel in Figure~\\ref{egression}. However, we also want emphasize here the possibility that photospheric heating also contributes to flare acoustic emission.  As the whole flaring process occurred relatively rapidly (and given the highly-impulsive properties of the 10 March 2001 flare, it is not unreasonable to assume so), the solar chromosphere was heated quite suddenly. As we have already seen, the multi-wavelength emissions of the flare indicate this; furthermore the strong spatial and temporal correspondence between the different types of emissions point to radiative back-warming playing a significant role in the heating mechanism. This conclusion was in fact drawn by both \\citet{lietal2005} and \\citet{ding03} to explain the origin of the continuum feature of the 10 March 2001 flare in terms of an ``electron-beam-heated flare model'', with chromospheric radiative back-warming suspected being the chief heating agent, originating in the temperature-minimum region.  The above conclusions, when viewed in conjunction with those of \\citet{donea05}, \\citet{donea06}, \\citet{moradi07}, and \\citet{Martinez-Oliveros2006}, provide direct evidence of flare acoustic emission being driven, in part, by heating of the low photosphere. The basic principle here is that the chromospheric radiation further heats up the photosphere, with the result being of an optically-thick $\\mathrm H^{-}$ bound-free absorption, which then introduces a pressure transient directly to the underlying medium. The photospheric heating hypothesis is well supported by our observations and previous ones -- which all indicate that instances of flare seismic emission have been characterized by a close spatial correspondence between the seismic emission and sudden WL emission during the impulsive phase of the flare. Radiative fluxes characteristic of WL emission seen in all acoustically active flares, if emitted downward from the chromosphere (as well as upward), are probably sufficient to heat the photosphere a few percent within a few seconds of the onset of the incoming radiative flux. This process, described by \\citet{machado1989}, is called ``back-warming''.  According to rough models described by \\citet{donea06} and \\citet{moradi07}, such heating (if applied suddenly) should cause a pressure transient in the heated layer that drives a seismic transient whose energy flux is of the order of those estimated for acoustically active flares. The energy invested into the seismic transient is in proportion to \\begin{equation} \\varepsilon \\sim (\\frac{\\Delta I_c}{I_c})^2 \\end{equation} a fraction of order $\\Delta I_c/I_c$ times the radiative energy suddenly emitted by the flare. We therefore expect acoustic emission due to photospheric heating to be inefficient in flares whose WL signatures are weak, diffuse, or not very sudden, and this is consistent with examples we have encountered to date.  In conclusion, it must be stated that to date, no one, single mechanism can fully explain the mechanics of flare acoustics and their observational signatures - because to do so would be a gross over-simplification of the problem. What these results have shown is that the study of flare mechanics, as well as helioseismology, would greatly benefit from the development of detailed models of solar flare-induced seismic emission -- from the corona, down to the photosphere, including modeling of realistic active region subphotospheres. Credible models would need to include realistic subphotospheric thermal anomalies to represent penumbral and perhaps umbral subphotospheres, and realistic photospheric magnetic fields extrapolated to depths of a few hundred kilometers beneath the photosphere. The later should also include an account for the highly inclined magnetic fields that characterize sunspot penumbra -- where a significant majority of the seismic emissions observed to date have been detected.   \\thanks{The authors would like to sincerely thank Profs. Paul Cally, Markus Aschwanden, Valentina Zharkova, Drs. Charlie Lindsey and Wahab Uddin, and Diana Besliu-Ionescu for their helpful and interesting comments which contributed directly to the development and/or improvement of this article.}"
    },
    "0801/0801.2315_arXiv.txt": {
        "abstract": "% The yellow hypergiants are found in a stage between the massive Red Supergiants and the Wolf-Rayet stars. This review\\footnote{This is an updated and slightly expanded version of a Keynote Talk given at ``Biggest, Baddest, Coolest Stars'' (ASP Conf Series) eds. D. Luttermoser, B. Smith, and R. Stencel} addresses current issues concerning the evolution of massive stars, concentrating on the transitional post-Red Supergiant phase.  Few yellow hypergiants are known and even fewer show direct evidence for having evolved off the Red Supergiant branch. Indeed, only two such rare objects with clear evidence for having gone through of a previous mass losing phase are known, IRC +10420 and HD 179821. We will review their properties, discuss recent results employing near-infrared interferometry, integral field spectroscopy and polarimetry.  Finally, their real-time evolution is discussed. ",
        "introduction": "A large variety of evolved massive stars are present in the upper parts of the HR diagram. Amongst others, we find the hot, mass losing Wolf-Rayet (WR) stars (Crowther, 2007), the unstable Luminous Blue Variables (LBVs, Humphreys \\& Davidson 1994), the enigmatic B[e] supergiants (Zickgraf et al. 1986), the massive Red Supergiants (RSG, Massey et al. 2008) and in between these phases, the Yellow Hypergiants (YHGs, de Jager 1998). An HR diagram with LBVs, Yellow Hypergiants and Red Supergiants is shown in Figure 1.    \\begin{figure} \\plotfiddle{./oudmaijer_1.ps}{7.cm}{90}{50}{50}{180}{0} \\caption{HR diagram displaying Luminous Blue variables (filled circles, taken from Smith et al. 2004), yellow hypergiants (asterisks) at their highest observed temperatures (de Jager 1998), and Red Supergiants (triangles) from Levesque \\& Massey (2005). Many Galactic YHG data are taken from de Jager (1998), HD 179821 is from from Reddy \\& Hrivnak (1999). A large number of YHGs are now known from detailed studies of massive Galactic clusters. F15 is reported by Figer et al. (2006), Star 49 is presented by Davies et al. (2007a), and data for W4, W8a, W12a, W16a, W32 and W265 in Westerlund 1 are taken from Clark et al (2005).  The dashed line denotes the Humphreys-Davidson limit.} \\end{figure}  A commonly cited evolutionary scenario linking most of these classes of object can be summarized as follows (e.g. Meynet \\& Maeder 2003). After a massive star with a mass of more than about 50 $M_{\\odot}$ moves off the main sequence it evolves straight to the WR phase via the B supergiant and LBV phases. In special circumstances it can then explode as a Supernova giving rise to a Gamma Ray Burst (e.g. Meynet \\& Maeder 2007). Lower mass stars with masses typically larger than 10 $M_{\\odot}$ evolve all the way from the blue to the red to become a Red Supergiant, during which phase they can lose half their mass via a strong, long lasting mass loss phase. When the star evolves off the RSG branch it can evolve via the YHG phase to the LBV phase to eventually become a WR star.  It should be mentioned however that it is still not clear whether any LBVs have evolved off the RSG (e.g. Lamers et al. 2001).  The situation proves even less settled for the B[e] supergiants, it is possible that they constitute a rapidly rotating sub-sample of the objects that evolve through the LBV phase.  In this paper, we focus on the evolutionary phase that immediately follows the Red Supergiant stage, the post-RSG phase.  This term would of course apply to all objects, including the WR stars, that have visited the Red Supergiant branch at some point during their evolution. For the purposes of this paper, we will confine ourselves to the first stop in their journey from the RSG towards the blue. A quick glance at the HR diagram brings us then immediately to the Yellow Hypergiants. They are located on the red side of the LBVs and the apparent empty region in the HR diagram where the temperatures are between $\\sim$7000 and 10 000 K. As outlined in Figure 1, the LBVs vary in temperature between $\\sim$9000 K up to the S Dor instability strip, for which higher luminosities correspond to higher temperatures. The YHGs are located on the red side of this region, which has historically been referred to as the Yellow Void (e.g. de Jager and Nieuwenhuijzen 1997), and have not been observed to cross it bluewards. Given that the LBVs are found on the blue side of the Yellow Void and also do not seem to cross it, while the larger number of known LBVs and YHGs seem to fill in the ``void'' since the term was introduced, it could be argued that ``Yellow Wall'' is perhaps a more appropriate term.  Yellow hypergiant stars are not in a quiet phase of their evolution. Observationally they have been found to exhibit explosive events and to undergo multiple blue- to red movements and vice versa in the HR diagram. After the stars evolve from the RSG branch and reach temperatures around 7000 K, their envelopes become unstable and a large mass loss ensues (e.g. Stothers \\& Chin 2001). A large enough mass loss rate can give rise to an optically thick wind, or in other words, the central star is surrounded by a cool pseudo-photosphere. After the eruption has occurred, the motion of the star in the HR diagram reverses and the object now follows a redward loop returning to the RSG branch. After the wind clears, the star will be on a blueward loop again until it can ``bounce'' another time against the Yellow Void.  This so-called ``bouncing against the Yellow Void'' has been observed in several instances. A compelling observation is presented by de Jager \\& Nieuwenhuijzen in 1997. They observed HR 8752 and found that it had undergone at least two such ``bounces'' in the preceding 30 years. In both cases, the bounce was associated with an episode of increased mass loss. Just before the mass loss event, the star had reached its maximum temperature, to decrease thereafter as expected for a newly formed thick pseudo-photosphere.  The duration of these episodes is around 10 years for HR 8752. This is shorter than what has been observed for the object Var A in M33. Humphreys et al. (2006), observed a spectral type variation for this object from F to M and back to F again over a period of 45 years, and present evidence that the high mass loss and subsequent clearing of the pseudo-photosphere are responsible for the redward and blueward motions in much the same manner as explained above. Similar events, but at much smaller timescales have been observed for the famous object $\\rho$ Cas. Lobel et al. (2003) describe the latest outburst of the star during which its temperature decreased from 7000 K to less than 4000 K in a matter of a few hundred days, to increase back to almost 6000 K in roughly the same amount of time.  Once a sufficient amount of mass has been ejected, and it is not yet clear how much, the stellar envelope can survive the instabilities and the star can evolve through the void to re-appear as an LBV. But this may not apply to all YHGs as speculated by Smith et al. (2004). Inspection of the HR diagram in Figure 1 reveals that there may be a gap in luminosity around log$(L/L_{\\odot})$ between 5.6..5.8 where there are apparently no LBV counterparts to the YHGs. Smith et al. propose that for stars within this luminosity range, their pseudo-photospheres will have even higher opacities due to the bi-stability mechanism (e.g. Vink et al. 1999). Subsequently they will not be visible as LBVs because the pseudo-photospheres effectively result in a YHG spectrum for a longer time. Instead, the objects that can cross the Yellow Void would only re-appear as the much hotter Ofpe/WN9 stars or B[e] stars.         \\begin{figure}  \\plotfiddle{./oudmaijer_2.ps}{8.cm}{0}{42}{42}{-150}{-50} \\caption{Infrared colours of the  K-G-F-A super- and hypergiants listed in Table 2 of de Jager 1998. The upper panel shows the {\\it K $-$ [25]} colour in magnitudes (17 objects), that of the lower panel the {\\it K $-$ [8]} colour of 19 stars. See text for details. }  \\end{figure} ",
        "conclusions": ""
    },
    "0801/0801.3288_arXiv.txt": {
        "abstract": "We find that the classical nova V723\\,Cas (1995) is still an active X-ray source more than 12 years after outburst and analyze seven X-ray observations carried out with \\swift\\ between 2006 January 31 and 2007 December 3. The average count rate is $0.022\\,\\pm\\,0.01$\\,cts\\,s$^{-1}$ but the source is variable within a factor of two of the mean and does not show any signs of turning off. We present supporting optical observations which show that between 2001 and 2006 an underlying hot source was present with steadily increasing temperature. In order to confirm that the X-ray emission is from V723\\,Cas, we extract a \\rosat\\ observation taken in 1990 and find that there was no X-ray source at the position of the nova. The \\swift\\ XRT spectra resemble those of the Super Soft X-ray binary Sources (SSS) which is confirmed by \\rxte\\ survey data which show no X-ray emission above 2\\,keV between 1996 and 2007. Using blackbody fits we constrain the effective temperature to between $T_{\\rm eff}=(2.8-3.8)\\times10^5$\\,K and a bolometric luminosity $\\gtrsim 5\\times10^{36}$\\,erg\\,s$^{-1}$ and caution that luminosities from blackbodies are generally overestimated and temperatures underestimated. We discuss a number of possible explanations for the continuing X-ray activity, including the intriguing possibility of steady hydrogen burning due to renewed accretion. ",
        "introduction": "\\begin{deluxetable}{ccccc} \\tablecaption{\\label{ferats}Journal of Optical observations} \\tablenum{1} \\tablehead{\\colhead{Date} & \\colhead{UT Time } & \\colhead{$t_{\\rm exp}$ (s)} & \\colhead{Obs.\\tablenotemark{a}} & \\colhead{[Fe\\,{\\sc x}]/[Fe\\,{\\sc vii}]\\tablenotemark{b}}\\\\ \\colhead{} & \\colhead{(hh:mm:ss)} & \\colhead{} & \\colhead{} & \\colhead{} } \\startdata 2001 Jan 19 & 03:57:21  &  300  &  SO 90'' & 0.02 \\\\ 2003 Sep 13 &  09:39:37  &  600  &  MMT & 0.62 \\\\ 2004 Jun 26 &  10:17:09  &  600  &  SO 90'' & 0.81 \\\\ 2006 Nov 9  &  02:42:59  &  600  &  SO 90'' & 2.33 \\enddata \\tablenotetext{a}{Steward Observatory Bok 90 inch telescope and the Multiple Mirror Telescope.} \\tablenotetext{b}{Ratio of the [Fe\\,{\\sc x}] (6373\\AA) to [Fe\\,{\\sc vii}] (6087\\AA) lines.}  \\end{deluxetable}  \\begin{figure*}[!ht] \\resizebox{\\hsize}{!}{\\includegraphics{optfe}} \\caption{\\label{optspec}Optical spectra of V723\\,Cas taken between January 2001 and November 2006, focusing on the $[$Fe\\,{\\sc x}$]$ and $[$Fe\\,{\\sc vii}$]$ emission lines at 6376\\,\\AA\\ and 6089\\,\\AA, respectively (see Table~\\ref{ferats}). The ratio of these lines is continuously increasing, suggesting that there is a hot source ionizing increasingly tenuous ejecta.} \\end{figure*}   Classical Novae (CNe) are the observable events caused by thermonuclear runaways on the surfaces of white dwarfs (WDs), fueled by material accreted from a companion star. Material dredged up from below the WD surface is mixed with the accreted material and violently ejected. The outburst continues until either nuclear burning has converted all the remaining hydrogen on the WD surface to helium or it has been ejected from off the WD. The bolometric luminosity at the peak of the explosion is near or exceeds the Eddington luminosity, $\\sim 1\\times10^{38}$\\,erg\\,s$^{-1}$ for a 1-M$_\\odot$ WD \\citep[see for example][]{schwarz_lmc}. Initially, the radiative energy output of the nova peaks in the optical, since the binary is surrounded by expanding ejecta that are not transparent to the high-energy radiation produced on the surface of the WD by nuclear burning. As the ejecta expand and thin (and the remaining material on the WD burns hydrostatically in a shell), the photosphere recedes to inner and hotter layers. The peak of the spectral energy distribution then shifts to higher energies \\citep{gallstar78}.  The X-ray bright phase of the evolution has only been observed sporadically, with the first systematic observations obtained with \\rosat\\ \\citep[e.g.,][]{kr02}. The X-ray spectrum during this phase was found to sometimes resemble that of the class of Super Soft X-ray Binary Sources \\citep[SSS;][]{heuvel} such as \\object[Cal 83]{Cal\\,83} in the LMC \\citep[e.g.,][]{lanz04}. This phase is now called the SSS phase. V1974 Cyg was the first nova to be followed in X-rays from before the SSS phase until decline. \\cite{krautt96} reported a slow rise and fast decline of the X-ray brightness after the peak had been reached. The first spectral modeling of V1974 Cyg was carried out by \\cite{krautt96} using blackbody fits, but \\cite{balm98} reanalyzed the spectra using LTE atmosphere models, because blackbodies did not fit the spectra. It is further known that luminosities and temperatures derived from blackbody fits are inconsistent with realistic physical conditions, with a tendency to overestimate the luminosity and thus underestimate the temperature \\citep{kahab}. The LTE models presented by \\cite{balm98} were an attempt to account for complex absorption patterns that change the shape of the spectrum. Although only 24 spectral bins were available, an attempt was made to determine a large number of parameters (temperature, surface gravity, bolometric luminosity, neutral hydrogen column density, and the abundances of carbon, oxygen, and neon). Further, LTE and a static WD atmosphere were assumed. The combination of a small ratio of independent data points to adjustable parameters and the limitations of current atmosphere models call for caution when drawing conclusions from these models. The effects of the expansion, especially during the early phases of the evolution, have been investigated with the PHOENIX code by \\cite{petz05}.  In addition to the SSS spectrum \\cite{krautt96} found a hard component that was fit with a MEKAL model by \\cite{balm98} and thus was assumed to be an optically thin plasma that is radiatively cooling. This phenomenon has been observed in a number of novae \\citep[e.g.,][]{swnovae}, the most recent example being RS\\,Oph \\citep[e.g.,][]{ness_rsoph}.  A systematic search of all \\rosat\\ observations of novae including the \\rosat\\ All Sky Survey and the archives of pointed \\rosat\\ observations was carried out by \\cite{orio01}. Of a sample of 108 CNe and recurrent novae (RNe) that exploded within the last $\\sim 100$ years, very few SSS spectra were found implying that the SSS phase must be short. A recent survey of all novae observed with \\swift\\ also resulted in few detections of CNe in the SSS phase \\citep[e.g.,][]{swnovae}. Few Galactic (e.g., \\citealt{orio01,swnovae} and references within) and extra-galactic \\citep{pietsch05,pietsch06} novae have had the duration of this phase determined. In the extra-galactic surveys candidates for novae in the SSS phase can be confused with supernova remnants, foreground neutron stars, or black hole transients \\citep{orio06}.  SSSs are believed to be WDs with accretion rates high enough to sustain steady nuclear burning \\citep{heuvel,kahab}. Stable shell burning can cause the WD to grow in mass via accretion with a buildup of hydrogen deficient material. Thus, the SSSs are strong candidates for single-degenerate supernovae Ia (SN\\,Ia) progenitors \\citep{starrf04}. In general, CNe are not considered to be SN\\,Ia progenitors because the outburst is believed to eject more material than is accreted.  While theoretical models predict novae to be active for 10s to 100s of years \\citep[e.g.,][]{st89,sala06}, all novae observed so far in X-rays have turned off after much less than a decade \\citep{orio01,swnovae}, implying mass must be lost during the outburst by some mechanism \\citep[e.g.,][]{macdonald85}. The typical observed outburst duration is of order one year with the longest observed Galactic nova in outburst being GQ\\,Mus at $\\sim$9 years \\citep{oegel93,shanley95}. The outburst duration depends on the nuclear burning rate, the amount of hydrogen left on the WD after the explosion, and the rate of any ongoing mass loss. Unfortunately, all these properties including the WD mass, the most critical driver, are poorly known.  A different approach was proposed by \\cite{greiner03} who found from the available data that systems with shorter orbital periods display long durations of supersoft X-ray phases, while long-period systems show very short or no SSS phase at all. They speculate that shorter periods may be related to a higher mass transfer rate (e.g., by increased irradiation) and thus to the amount of material accreted before the explosion.  In this paper we present \\swift\\ observations in X-rays of the slow nova V723\\,Cas (\\S\\ref{back}). In \\S\\ref{obs} \\& \\S\\ref{analysis} we describe the observations and our analysis and address the possible future X-ray evolution in \\S\\ref{disc}.\\\\ ",
        "conclusions": "\\label{conc}  Although other CNe have been observed to be active for several years, V723\\,Cas is now confirmed to be the longest CN observed in outburst and as of December 2007 shows no indication it is turning off. While the observed count rate was variable, the spectral hardness remained unchanged. From blackbody fits we find a temperature range from 2.8$\\times$10$^5$ to 3.8$\\times$10$^5$\\,K and the luminosity greater than 5$\\times$$10^{36}$\\,erg\\,s$^{-1}$. We caution that these numbers have limited physical meaning, as it is well known that blackbody fits lead to overestimated luminosities and thus underestimated temperatures \\citep{kahab}. Nevertheless, the high temperatures (especially if the numbers are underestimated) indicate that nuclear burning is still continuing. V723\\,Cas may be burning the remaining hydrogen left on the WD surface after the initial explosion. If so, the X-ray emission will decline once the fuel is exhausted. It is unknown when this event will happen but we plan on continuing our \\swift\\ monitoring program to detect a possible X-ray turn off and to further investigate the changes in the X-ray count rate.  There is also the intriguing possibility that the nova is now fed via renewed accretion similar to that observed in the prototype SSS, Cal\\,83. If so, V723 Cas may have evolved into a permanent SSS. Evidence that points to this possibility is the exceptional length of time in outburst and the long orbital period similar to that in Cal\\,83 and Cal\\,87. Note also that \\cite{greiner03} found that long-period systems show shorter SSS phases than short-period systems. V723\\,Cas does not fit into this picture, suggesting that this is a different situation. Establishing the WD mass and determining whether accretion has re-established will assist in understanding the fate of V723 Cas. Schwarz et al. (in preparation) are presently obtaining optical spectra during the different phases of the orbit to determine whether the V723\\,Cas system is dominated by an illuminated subgiant secondary, a re-established accretion disk, or some mixture of both. Any orbital variations in the lines may also be used to provide mass estimates of the system components."
    },
    "0801/0801.3241_arXiv.txt": {
        "abstract": "We present a detailed analysis of the properties of tidally stripped material from disrupting substructure haloes or \\emph{subhaloes} in a sample of high resolution cosmological $N$-body host haloes ranging from galaxy- to cluster-mass scales. We focus on devising methods to recover the infall mass and infall eccentricity of subhaloes from the properties of their tidally stripped material (i.e. tidal streams). Our analysis reveals that there is a relation between the scatter of stream particles about the best-fit debris plane and the infall mass of the progenitor subhalo. This allows us to reconstruct the infall mass from the spread of its tidal debris in space. We also find that the spread in radial velocities of the debris material (as measured by an observer located at the centre of the host) correlates with the infall eccentricity of the subhalo, which allows us to reconstruct its orbital parameters. We devise an automated method to identify leading and trailing arms that can, in principle at least, be applied to observations of stellar streams from satellite galaxies. This method is based on the energy distribution of material in the tidal stream. Using this method, we show that the mass associated with leading and trailing arms differ. While our analysis indicates that tidal streams can be used to recover certain properties of their progenitor subhaloes (and consequently satellites), we do not find strong correlations between host halo properties and stream properties. This likely reflects the complicated relationship between the stream and the host, which in a cosmological context is characterised by a complex mass accretion history, an asymmetric mass distribution and the abundance of substructure. Finally, we confirm that the so-called ``backsplash'' subhalo population is present not only in galaxy cluster haloes but also in galaxy haloes. The orbits of backsplash subhaloes brought them inside the virial radius of their host at some earlier time, but they now reside in its outskirts at the present-day, beyond the virial radius. Both backsplash and bound subhaloes experience similar mass loss, but the contribution of the backsplash subhaloes to the overall tidal debris field is negligible. ",
        "introduction": "The currently favoured model for cosmological structure formation is the $\\Lambda$CDM model. This model assumes that we live in a spatially flat Universe whose matter content is dominated by non-baryonic Cold Dark Matter (CDM) and whose present-day expansion rate is accelerating, driven by Dark Energy ($\\Lambda$). In the context of this model, the formation and evolution of cosmic structure proceeds in a ``bottom-up'' or hierarchical manner. Low mass gravitationally bound structures merge to form progressively more massive structures, and it is through such a merging hierarchy that galaxies, groups and clusters form \\cite[e.g.][]{Davis.etal.85, Springel.Frenk.White.06}.  The tidal disruption of satellite galaxies around our Galaxy and others is characteristic of the hierarchical merging scenario. As a satellite galaxy orbits within the gravitational potential of its more massive host, it is subject to a tidal field that may vary in both space and time. The gravitational force acting on the satellite strips a stream of tidal debris from it, and in some cases tidal forces may be sufficient to lead to the complete disruption of the system. This tidal debris tends to form two distinct arms -- a \\emph{leading} arm ahead of the satellite and a \\emph{trailing} arm following the satellite.  The process by which a satellite galaxy suffers mass loss and the subsequent formation of a stream of tidal debris (hereafter referred to as a \\emph{tidal stream}) has been studied extensively using numerical simulations \\cite[e.g.][]{Johnston.98, Ibata.etal.01, Helmi.04}. A number of these studies have followed the orbital evolution of an individual satellite galaxy with a known mass profile, realised with many particles, in a static analytic and non-cosmological\\footnote{By ``non-cosmological'' we mean that the systems were evolved in isolation, in the absence of the background expansion of the Universe, without infall and accretion, and without a population of substructures.} host potential \\citep[e.g.][and many others]{Johnston.98, Ibata.etal.01, Helmi.04}. This approach allows for well-defined test cases to be investigated and for a systematic survey of orbital and structural parameters to be carried out. For example, \\cite{Helmi.04} showed that the width of the tidal debris in the plane perpendicular to its orbit increases when the host halo potential is flattened rather than spherical.  In this regard, it is important to note that stream properties are sensitive to both the internal properties of its progenitor satellite and the details of the orbit that it follows.  For example, we expect the length and width of a stream to tightly correlate with the age of the stream and the mass/internal velocity dispersion of its progenitor \\cite[e.g.][]{Johnston.Hernquist.Bolte.96, Johnston.98,Johnston.Sackett.Bullock.01,Penarrubia.etal.06}. Understanding what role the progenitor satellite and its orbit plays in shaping stream properties is therefore vital if we wish to use streams to address more fundamental problems, such as the nature of the dark matter.\\\\  While previous studies have provided important insights into the formation and evolution of tidal streams and their dependence on the properties of both satellites and host dark matter haloes, none have considered tidal streams in cosmological host haloes\\footnote{We note the work of \\citet[][]{Penarrubia.etal.06}, which employs a semi-analytical approach to study tidal streams in an evolving host potential whose time variation is based on results of cosmological simulations.}. Cosmological haloes have complex mass accretion histories and mass distributions that can at best be crudely approximated by fitting formulae, and the interplay between a halo's asphericity, its clumpiness, and how its mass grows in time may be non-trivial. How this complexity may affect the results of previous studies, which were based on analytic potentials and approximations to halo growth, is difficult to say. For example, interactions between the stream and substructures lead to the dynamical heating, causing an increase in internal velocity dispersion and a subsequent broadening of the stream \\citep{Moore.etal.99, Ibata.etal.02, Johnston.Spergel.Haydn.02}.  Therefore, while studying tidal streams in cosmological haloes is a challenging problem, such an approach is necessary if our results are to reflect the complexity of nature (assuming the validity of the $\\Lambda$CDM model, of course!).  This paper is the first of in a series that will examine the formation and evolution of tidal streams in \\emph{cosmological} dark matter haloes, and determine how the properties of these tidal streams depend on the properties of the satellite galaxies from which they originate and on the host dark matter haloes in which they orbit. We use high resolution cosmological simulations of individual galaxy- and cluster-mass dark matter haloes -- the \\emph{host haloes} -- that allow us to track \\emph{in detail} the evolution of a large population of their substructures or \\emph{satellite galaxies} over multiple orbits\\footnote{For the purposes of our study, we treat substructure haloes and satellite galaxies as interchangeable, insofar as both suffer mass loss and the tidally stripped material is a dynamical tracer of the host potential. However, we note that the correspondence between dark matter substructures and luminous satellite galaxies is not a straightforward one \\citep[][]{Gao.DeLucia.etal.04}.}. We seek to answer the questions we consider to be key to the study of tidal streams;  \\begin{itemize} \\item \\emph{Can observations of tidal streams be used to infer the properties of their parent satellite?} \\item \\emph{Can tidal streams reveal properties of the underlying dark matter distribution of the host halo?} \\end{itemize}  The present paper primarily focuses on the first question and can be summarised as follows. In \\Sec{sec:simulations} we describe the fully self-consistent cosmological simulations of structure formation in a \\LCDM\\ universe that we have used in our study. The suite of host haloes (consisting of eight cluster and one Milky Way sized dark matter halo) are presented in \\Sec{sec:hosts}. In \\Sec{sec:subhaloes} we shift attention to the formation and evolution of debris fields. In this regard, \\Sec{sec:subhaloes} deals with mass loss from subhaloes and the contribution of so-called ``backsplash'' satellites. These subhaloes are found outside the virial region of the host at the present day, but their orbits took them inside the virial radius at earlier times. We show that they exist not only in cluster-sized haloes as found by e.g. \\cite{Gill.etal.05.3} and \\cite{Moore.Diemand.Stadel.04}, but also in our galactic halo. In \\Sec{sec:streams_sat} we examine how well satellite properties such as its original mass and the eccentricity of its orbit can be reproduced using the whole debris field of the corresponding satellite. Furthermore, we present in \\Sec{sec:trailing/leading} a method to distinguish between trailing and leading arms in an automated fashion and we investigate properties of the streams such as their mass, shape and velocity dispersion. We pay particular interest to the energy distribution of debris, since it can be used to distinguish between leading and trailing arm observationally\\footnote{It is important to remember that our tidal streams consist of ``particles'' of dark matter, whereas an observer detects stellar streams. However, the physical processes involved in the formation of both dark matter and stellar streams are identical, and so we can think of them interchangeably.}. Finally, we summarise our results in \\Sec{sec:summary} and comment on their significance in the context of future papers, which will try to establish the missing link between the properties of tidal streams and the host haloes only briefly touched upon in this study. ",
        "conclusions": "\\label{sec:summary}   In this paper, which is the first in a series, we have investigated the tidal disruption of subhaloes (or \\emph{satellite galaxies}) orbiting in cosmological dark matter host haloes. We used nine high resolution cosmological $N$-body simulations of individual galaxy- and cluster-mass haloes, each of which contain in excess of $\\sim 1$ million particles within the virial radius, and within which a few hundred subhaloes can be reliably resolved. The time sampling of our simulations is such that we can track subhalo orbits in detail, which allows us to both determine orbital planes for subhaloes and to follow the mass loss they suffer as a function of time.  In the first part of the paper (i.e. \\Sec{sec:hosts}) we concluded that, while all of our host haloes are sufficiently relaxed, they exhibit a range of \\emph{integral} properties.  For example, they show a spread in formation times, mass assembly histories and shapes. These are precisely the kind of host properties that we expect to affect the properties of tidal debris field stripped from orbiting satellites.  The second part of the paper (i.e. Sections~\\ref{sec:subhaloes}--\\ref{sec:trailing/leading}) focused on the subhaloes and their tidal disruption. Our main conclusions may be summarised as follows;  \\begin{itemize}  \\item We predict that ``backsplash'' satellite galaxies should be found not only in the outskirts of galaxy clusters but also around galaxies.  \\item Backsplash subhaloes contribute as much as 13\\% of the mass of the total tidal debris field within the host. This suggests that the region outside the virial radius of the host galaxy should be considered also when observing the debris of satellites and searching for possible corresponding satellite galaxies.  \\item Tidal stream properties correlate well with progenitor satellite properties. For example, we find a relationship between the scatter about the (best-fit) debris plane and the infall mass of the satellite, as well as the spread in radial velocity and the infall eccentricity.  \\item We devised a method to separate leading and trailing debris arm in cosmological simulations.  \\item The masses of leading and trailing arms differ.  \\item We do not find differences in the velocity dispersions of leading and trailing arms.  \\item We present a method that uses energy distributions to separate leading and trailing arms and that can, in principle, be employed by observers.  \\item The formation of the tidal debris field is extremely complex in ``live'' host haloes such that there appears to be little correlation between debris properties and host halo properties (age, mass assembly history, dynamical state, shape).  \\end{itemize}  We set out to address two central questions key to understanding tidal streams; \\begin{itemize} \\item \\emph{Can observations of tidal streams be used to infer the properties of their parent satellite?} \\item \\emph{Can tidal streams reveal properties of the underlying dark matter distribution of the host halo?} \\end{itemize}  These have proven to be very weighty questions to address, and so we focused on the first question in this paper. As the bullet points above indicate, the answer would appear to be yes. For example, there is a clear correlation between mass loss and orbital infall eccentricity, and between deviation of the tidal debris field from its orbital plane and the infall mass of the progenitor satellite. Addressing the second question is less straightforward. We have examined our data for correlations between stream properties and host properties as summarised Table~\\ref{t:haloprop}, but evidence for correlations with host mass accretion history, dynamical state and shape are extremely weak at best.  Indeed, the only suggestion of a correlation that we find is between the mass ratio of leading to trailing arm and the virial mass of the host; while all our cluster mass haloes have more mass in the leading debris arm, the relation is reversed in the galaxy mass halo.  Further, it is interesting to note that the galaxy halo G1 appears to be distinct from the eight galaxy cluster haloes (with the possible exception of C1) in terms of how its bulk properties relate to its stream properties. However, it is not clear whether this is a systematic effect, given that G1 is the only realisation of a galactic type halo. For example, the differences are not so striking if one compares G1 with C1. We are in the process of building a sample of galaxy haloes to supplement G1.\\\\  We now conclude with some general remarks. Previous studies of individual satellites disrupting in analytic host potentials using controlled $N$-body experiments have hinted at simple and direct links between, for example, the flattening of the host potential and the dispersion of the debris field \\citep[e.g., ][]{Ibata.etal.03, Helmi.04, Penarrubia.etal.06}. However, our experience suggests that the situation is vastly more complicated in ``live'' host haloes, where the host undergoes a complex mass accretion history, contains a wealth of substructure and cannot be easily modeled as a simple ellipsoid. \\citet[][]{Penarrubia.etal.06} pointed out that all present-day observable stream properties can only constrain the present mass distribution of the host halo, independent of its past evolution. We would argue strongly that even present-day properties of the host halo will be difficult to infer from stream properties, given the complexity of interactions responsible for the emergence of debris fields. There are multiple processes driving these effects and so while we note tentative correlations, we require further simulations to help us understand what is happening. This will represent the focus of the following paper in this series. We point out that this situation is akin to the findings of \\citet[][]{Gill.etal.04.2}; they presented a detailed study of the dynamics of satellite galaxies in similar host haloes and were also unable to establish correlations to the host properties.  We close with a cautionary note: this paper is \\textit{not} to be understood as a definitive quantitative analysis of tidal streams in dark matter haloes. Rather it shows that certain (unexpected) correlations between the debris field and satellite properties exist. Our study therefore represents the first step towards our understanding of whether or not tidal streams from disrupting satellite galaxies can be used to deduce robustly properties of their host dark matter haloes and of the orbits of the satellites."
    },
    "0801/0801.4657_arXiv.txt": {
        "abstract": "We report the discovery of a giant,  loop-like stellar structure around the edge-on spiral galaxy NGC\\,4013. This arcing feature extends 6$\\arcmin$ ($\\sim$ 26 kpc in projected distance) northeast from the center and 3$\\arcmin$ ($\\simeq 12$ kpc) from the disk plane; likely related features are also apparent on the southwest side of the disk, extending to 4$\\arcmin$ ($\\sim$ 17 kpc). The detection of this low surface-brightness ($\\mu_{\\rm R}= 27.0^{+0.3}_{-0.2}$ \\magsqarcsec) structure  is independently confirmed in three separate datasets from three different telescopes.  Although its true three dimensional geometry is unknown,  the sky- projected morphology of this structure displays a match with  the theoretical predictions for the edge-on, projected view  of a stellar tidal streams of a dwarf satellite moving in a low inclined ($\\simeq 25^\\circ$), nearly circular orbit. Using the recent model of the Monoceros tidal stream in the Milky Way by Pe\\~narrubia et al. as template, we find that the progenitor system may have been a galaxy with an initial mass $6\\times 10^8 M_\\odot$, of which current position and final fate is unknown. According to this simulation, the tidal stream may be approximately  $\\sim 2.8$ Gyr of age.  Our results demonstrate that NGC\\,4013, previously considered a prototypical isolated disk galaxy in spite of having one of the most prominent \\hone warps detected thus far, may have in fact suffered a recent minor merger. This discovery highlights that undisturbed disks at high surface brightness levels in the optical but warped in \\hone maps may in fact reveal complex signatures of recent accretion events in deep photometric surveys. ",
        "introduction": "In the last decade, the study of the formation and evolution of the Milky Way (MW) has been revolutionized by the first generation of wide-field, digital imaging surveys. The resulting extensive photometric databases have revealed for the first time the existence of spectacular stellar tidal streams \\citep[e.g., that from the Sagittarius dwarf galaxy;][]{Ibata2001a,MD2001,Majewski2003} as well as large stellar sub-structures in the halo \\citep{Newberg2002,RP2004,Juric2005}, interpreted as the fossils of the hierarchical formation of our Galaxy. The discovery of the Monoceros tidal stream \\citep{Newberg2002,Yanny2003}, located close to the Galactic plane outside the MW disk \\citep[as well as similar structures seen around the M31 disk,][]{Ibata2005}, indicates that mergers might play a relevant role in the formation of the outer regions of spiral disks \\citep{Jorge2006} like that of the MW.  Moreover, inside-out disk formation from continual accretion of in-falling material is an observed feature in cold dark matter ($\\Lambda$CDM) simulations of the growth of structures on galaxy- sized scales \\citep[e.g.,][]{Abadi2003}, and is now also a common feature of galactic chemical evolution models \\citep[e.g.,][]{ALC01, CMR01} attempting to explain trends in disk chemical abundance patterns. These various results provide clear evidence that the destruction of satellite galaxies plays a relevant role not only in the formation of MW-like spiral galaxies generally, but for their disks as well as their halos.  Furthermore, these results suggest that the stellar mass assembly of the MW disk, and disks in general, likely continues actively to the present epoch.  The accretion of satellite galaxies on low-inclination orbits may be an important formation mechanism of galactic disks in the $\\Lambda$CDM paradigm \\citep[e.g.,][]{Abadi2003}. The existence of low-inclined tidal streams in the outer regions of our Galaxy \\citep[e.g., the Monoceros or Triangulum/Andromeda tidal streams,][]{Yanny2003,Majewski2003,Grillmair2006} supports this prediction, although the external origin of these over-densities is still a matter of debate, mainly due to the severe extinction hindering the exploration of low-latitude areas. Recently, for example, Kazantzidis et al. (2008) and Younger et al. (2008) have suggested that, in a $\\Lambda$CDM context, the bombardment by dark matter sub-halos of an initially cold disk formed at $z\\simeq 1$ would also yield the formation of low-surface brightness substructures ($\\mu \\simgreat 25$  \\magsqarcsec) with filamentary shape that may locally resemble tidal streams in phase-space. Clearly, additional observational data (e.g. detailed chemical abundances) obtained along the Monoceros stream are needed to distinguish between both theoretical scenarios.  The search for extragalactic analogies to MW minor mergers is required not only to (1) show that the MW is not unusual in this respect, but to (2) obtain externally-viewed ``snapshots'' of different phases of such interactions, (3) explore the range of possible mass/orbit combinations for such activity, and (4) estimate the fractional contribution of accreted mass and the mass spectrum of such events in the life of MW-like systems, an issue that remains unresolved \\citep{Majewski1999}.  In this way, a systematic survey of tidal streams around other, nearby disk galaxies can provide new constraints and insights on the hierarchical formation and structure of MW-like galaxies beyond the previously limited views afforded by our own Galaxy and M31 \\citep{Ibata2007}. A need to build up statistical information on the number and distribution of tidal streams in galaxies is driven by the availability of state-of-the-art, high resolution cosmological simulations that offer, for the the first time, the opportunity to use these observations to probe the theoretical predictions of galaxy formation in the framework of the $\\Lambda$CDM paradigm \\citep[e.g.,][]{Bullock2005}.   Promising galaxies in the hunt for extragalactic tidal streams are those displaying outstanding asymmetries in optical or \\hone images. It has been long suggested that these perturbations are a result of gravitational interaction with nearby companions. In some cases, apparently isolated galaxies exhibit morphologies more commonly associated with interacting systems. Recently, the prototypical isolated, warped disk galaxy, NGC\\,5907, was found to be surrounded by a spectacular stellar tidal stream \\citep{Shang1998,Martinez-Delgado2008}, but one with no obvious companion (which may have been completely disrupted).  In the case of both the MW and M31 there are also obvious disk warps --- both stellar and gaseous --- but no immediately obvious perturber, though small, nearby, likely tidally disrupting companion satellites have been suggested as a cause in each case --- \\citep{Sato1986,Ibata1998,Bailin2003,Bailin2004}. These examples suggest that spirals with warped disks but no large nearby companions, may have undergone {\\em minor} mergers in the last Gyrs.  Under this prepense, we have been observing disk galaxies to faint surface brightnesses in the search for evidence of minor mergers.  We \\citep{Martinez-Delgado2008} have previously reported our observations of the debris of a minor merger around the NGC\\,5907 disk system.  Here we report similar work on the edge-on disk, NGC\\,4013. NGC\\,4013 is one of the 62 luminous ($M_{B} <-16.9$) members of the Ursa Major cluster of galaxies, a nearby, late-type dominated, and low mass galaxy cluster. According to the HYPERLEDA database, the mean heliocentric radial velocity of NGC\\,4013 (836 \\kms) places it at a distance of 14.6 Mpc. It is also a relatively isolated system, with the two nearest, slightly more luminous, cluster neighbours (NGC\\,4051 and NGC\\,3938) at $\\sim170/250$ kpc projected distance away. Classified as an Sb galaxy with a maximum observed rotational velocity of 195 \\kms \\citep{Bottema1996}, an extinction corrected total absolute magnitude of -20.1 \\mabs $_{B{\\rm band}}$ \\citep{Verheijen2001}, an optical scale-length of 2.8\\,kpc (40\\arcsec) \\citep{vdk1982} and an isophote 25 mag arcsec$^{-2}$ radius of 5.2\\arcmin \\citep{rc3} , NGC\\,4013 is very similar to the Milky Way.  NGC\\,4013 is, moreover, famous for its prodigiously warped \\hone disk, with a line of nodes close to parallel with that of the line of sight and with one of the largest warp angles observed ($\\sim 25\\deg$).  On one side the warp extends out to about 8\\,kpc (2$\\arcmin$) from the nominal plane \\citep{Bottema1987,Bottema1995,Bottema1996}. The box/peanut shaped bulge of NGC\\,4013 with its 'X'-like morphology indicates the presence of a bar being observed edge-on \\citep{Patsis2006}.   Here we report the optical detection of a faint, low-surface brightness, loop-like structure that appears to be part of a giant, low-inclination stellar tidal stream of a disrupted dwarf satellite looking very similar to the expected form of the Monoceros tidal stream. This looping structure suggests the likelihood of there having been a previous interaction of NGC\\,4013 with a low-mass companion. Given the fact that NGC\\,4013's very prominently warped \\hone layer matches in shape and orientation those of the apparent tidal loop makes this galaxy system a very compelling example of the likely link between disk warps and mergers, and an interesting case study for models of warp formation. ",
        "conclusions": "\\label{discussion}  Our deep images of NGC\\,4013 show a complex stellar halo, including a coherent, stream-like structure inclined $\\sim\\!25\\deg$ with respect to the galaxy disk. We have independently confirmed the presence of these low surface brightness features with three data sets from three different telescopes. In addition, if one knows what to look for, the loop is also barely visible on the Digitized Sky Survey 2 (DSS2) blue plate and on the archival $4.5\\mu$ Spitzer Infrared Telescope Facility image. The existence of the loop structure is therefore without question. The discovered structures are evidently debris material produced by the tidal disruption of a low-mass galaxy companion to NGC 4013: Such features are expected from tidal disruption of a companion galaxy \\citep{Johnston2001} and similar features are also observed in NGC\\,5907 \\citep{Shang1998,Martinez-Delgado2008}, NGC\\,3310 \\citep{Wehner2005}, and a handful of other external galaxies \\citep[summarized in][Table 1]{Martinez-Delgado2008}.  The sky-projected geometry of this structure is in fact  similar to the edge-on, external perspective of a tidal stream that results from the tidal disruption of a single dwarf galaxy satellite moving on a low-inclined orbit  predicted by different theoretical simulations (e.g. the Monoceros tidal stream: Pe\\~narrubia et al. 2005;Canis Major over-density: Martin et al. 2005, their Fig. 14). For comparison purpose,  we will use here as  template the N-body model constructed by Pe\\~narrubia et al. (2005)  to reproduce the geometry and kinematics of the Monoceros tidal stream. This model simulates the accretion of a low-mass satellite (6 $\\times$10$^{8}$ M$_{\\sun}$) on a nearly circular (e=0.10$\\pm$0.05), low inclined ($i\\!\\sim\\!25^{\\circ}$) orbit in the last 3\\,Gyr and provides a remarkable good match to the projected, edge-on view of the stream wrapping NGC 4013 (see Fig.~\\ref{fig:FoS}). Considering that NGC\\,4013 has a smaller disk scale-length, $h_R\\!=\\!2.8$ kpc, and a lower maximum circular velocity, $V_{c,{\\rm max}}\\!=\\!195$ \\kms \\citep{Bottema1996} than the Milky Way, we can apply dynamical considerations to estimate the stream age from the N-body model, so that ${\\rm Age}/{\\rm Age}_{Mon}= h_R/h_{R,MW}\\times  V_{c,{\\rm max}}/V_{c,{\\rm max},MW}$. For the Milky Way we have that $h_R=3.5$ kpc and $V_{c,{\\rm max}}=220$ km/s, thus ${\\rm Age}/{\\rm Age}_{Mon}\\approx 0.9$. Since some pieces of the Monoceros stream are ${\\rm Age}_{Mon}=3$ Gyr old, we estimate that the stream surrounding NGC\\,4013 may be as old as $\\simeq 2.8$ Gyr. However, although the model provides a reasonable qualitative resemblance to the morphology of NGC 4013, it cannot fully explain the observed structure. For example, while the observed tidal structure in Fig. 3  is clearly  more prominent on its Eastern side of the galaxy, the model shows loops of roughly equal brightness on both sides (as well as various other features that not detected in our deep image). Unfortunately, with no kinematic data and only the projected geometry of the stream as constraint, the orbit and mass of the progenitor system may be degenerated, and in that respect the estimates presented here must be considered illustrative only.  Although the progenitor system of the Monoceros and NGC 4013 tidal streams might have shared similar orbits and initial masses, our estimates reveal that the surface brightness of the latter is approximately $\\sim$2--3 magnitudes higher than the value reported for this low galactic latitude tidal debris candidate ($\\mu_{\\rm B}\\sim 30.0 $\\magsqarcsec: J. de Jong, private communication). A possible explanation might be that the progenitor of the NGC\\,4013 stream was more luminous than that of the Monoceros stream. Since both models have the same initial mass, this hypothesis would imply that the mass-to-light ratio of the NGC\\,4013 stream progenitor was considerably lower than that of the Monoceros stream. However, the N-body models used by Pe\\~narrubia et al. (2005) assume that the stellar and dark matter components of the progenitor system have the same spatial distribution (i.e. mass-follows-light models). In a $\\Lambda$CDM cosmogony, stars only populate the inner-most regions of dwarf galaxies (Strigari et al. 2007, Pe\\~narrubia et al. 2008a) and can only be tidally stripped after most of the dark matter halo beyond the luminous radius has been lost (Pe\\~narrubia et al. 2008b), which introduces an additional parameter: the spatial segregation of the stellar component with respect to the surrounding dark matter halo of the progenitor galaxy, which is fundamental to determine the survival time of dwarf galaxies and the properties of their associated tidal streams.  For example, the more deeply embedded within the halo a stellar component is, the thinner, brighter and colder will be its associated tidal stream. Therefore, for $\\Lambda$CDM- motivated models, the halo size, the stellar spatial segregation and the total luminosity of a dwarf galaxy are all parameters that influence the resulting surface brightness of an associated tidal stream but which,  unfortunately, cannot be constrained in absence of additional data (e.g. kinematics).   Whereas the detection of the stream around NGC\\,4013 cannot shed light on the nature of the Milky Way stellar over-densities (including on the controversial origin of the Monoceros stream; see Sec. 1), it does suggest that the accretion of satellite galaxies on low-inclination, low-eccentric orbit may be relatively common during the hierarchical formation of spiral galaxies. In fact, their occurrence in the known sample of stellar tidal streams in the Local Volume \\citep[$D\\!<\\!15$\\,Mpc; see][]{Martinez-Delgado2008,Pohlen2004} is $\\gtsim 20\\%$. In this context, the presence of giant disks around spiral galaxies that extend out to several scale-lengths of the inner disk (e.g.~the extended disk found by \\cite{Ibata2005} around M31 is detectable out to a projected radius of 80\\,kpc with the present instrumentation) might also be signature of the accretion of massive satellite galaxies moving on low-inclination, low-eccentricity orbits \\citep{Jorge2006}, adding support to the scenario that the disks of spiral galaxies may grow inside out as a result of this type of accretion events.   The detection of galactic warps in spiral galaxies that present signatures of recent accretion events is quickly increasing with the availability of wide-field, deep photometric surveys. Although the origin of galactic warps is still unknown, it was early proposed that the gravitational perturbations induced by satellite galaxies may cause the formation of warps (Burke 1986, although see Hunter \\& Toomre 1969 for counter-arguments). Since those early studies, several alternative scenarios have been proposed in the literature. For example, it has been shown that the interaction of a stellar disk with the surrounding dark matter halo may also induces the formation of warps \\citep{Sparke1988,Binney1998}, as well as bending instabilities \\citep{Revaz2004}, intergalactic accretion flows onto the disk \\citep{Lopez2002} and cosmic in-fall \\citep{Shen2006}.  A crucial fact that has been often invoked to dismiss the tidal origin of galactic warps was the existence of warped spiral galaxies in apparent isolation \\citep{Sancisi1976}. NGC\\,4013, with its prominent, and rather symmetrical \\hone warp, along with a really unassailable indication of a past merger event, provides an important counter-example to the above argument; indeed, NGC\\,4013 may well be the Rosetta Stone for warp theories, since morphologically the warp and the merger debris seem so closely aligned.  With NGC\\,5907 and NGC\\,4013 we have now two examples of apparently isolated galaxies with significantly warped gaseous disks, which at the same time show evidence for an ongoing tidal disruption of a dwarf companion. The connection between the warp of the MW and its satellite galaxies has been already explored in some recent studies. For example, \\cite{Bailin2003} and \\cite{Weinberg2006} study the effects of the Sagittarius dwarf and the Magellanic Clouds on the MW disk, respectively. Interestingly, the existence of warped galaxies that have suffered a recent accretion of satellite galaxies on low-inclination, low-eccentricity orbits (like the MW and NGC\\,4013) suggests that these kind of accretion events may induce the formation of warps more efficiently than the perturbations from satellite galaxies moving on highly eccentric, nearly-polar orbits. This is a possibility worth to be further investigated.  With the growing numbers of examples showing a connection between warped disks and evidence of mergers, even for disk systems with no obvious companions, we may be close to resolving the mystery of warps, and in a way be consistent with prevailing $\\Lambda$CDM models of galaxy formation."
    },
    "0801/0801.3077_arXiv.txt": {
        "abstract": "Low-frequency radio observations of neutral hydrogen during and before the epoch of cosmic reionization will provide hundreds of quasi-independent source planes, each of precisely known redshift, if a resolution of $\\sim 1$ arcminutes or better can be attained.  These planes can be used to reconstruct the projected mass distribution of foreground material. A wide-area survey of 21 cm lensing would provide very sensitive constraints on cosmological parameters, in particular on dark energy. These are up to 20 times tighter than the constraints obtainable from comparably sized, very deep surveys of galaxy lensing although the best constraints come from combining data of the two types. Any radio telescope capable of mapping the 21cm brightness temperature with good frequency resolution ($\\sim $ 0.05 MHz) over a band of width $\\simgt$~10 MHz should be able to make mass maps of high quality.   If the reionization epoch is at $z\\simlt 9$ very large amounts of cosmological information will be accessible.  The planned Square Kilometer Array (SKA) should be capable of mapping the mass with a resolution of a few arcminutes depending on the reionization history of the universe and how successfully foreground sources can be subtracted.  The Low-Frequency Array (LOFAR) will be able to measure an accurate matter power spectrum if the same conditions are met. ",
        "introduction": "The pregalactic 21 cm radiation originates from a period when most of the hydrogen in the universe was neutral before most of the galaxies formed and emitted ionizing radiation.  This is often called the dark ages.  At $z\\sim 1100$  the universe cooled to the extent where nearly all of the hydrogen became neutral. We know this from the spectrum of the CMB.  Some time between $z=6$ and $z=30$ most of the hydrogen in the universe reionized.   We know this from the lack of sufficient Lyman-$\\alpha$ absorption in high redshift quasar spectra.  The 21 cm line comes from the hyperfine splitting of the ground state of hydrogen (coupling between the spin of the electron and proton).  It is redshifted to several meters when it reaches us.  There will be many statistically independent regions of 21~cm emission at different redshifts (or frequencies) at a single position on the sky.  Gravitational lensing will distort the images of these independent sources in a coherent way.  By ``stacking'' up the 21~cm maps in different frequencies the lensing distortion can be separated from the intrinsic structure in the 21~cm emission \\cite{ZandZ2006,metcalf&white2006}.  There are several telescopes being planned and built to observe the 21~cm radiation from the epoch of reionization.   Among these are the 21 Centimeter Array (21CMA, formerly known as PAST)\\footnote{http://21cma.bao.ac.cn/}, the Mileura Widefield Array (MWA) Low Frequency Demonstrator (LFD)\\footnote{ttp://www.haystack.mit.edu/ast/arrays/mwa/}, the core array of LOFAR (Low Frequency Array)\\footnote{www.lofar.org} and the core of SKA (Square Kilometer Array)\\footnote{www.skatelescope.org/}. Only LOFAR and SKA will have sufficient collecting area to be relevant for gravitational lensing observations as presently conceived. ",
        "conclusions": "Much will depend on future instrument design and the as yet unknown characteristics of the 21 cm absorption and emission, particularly around the epoch of reionization.  If reionization happens unexpectedly early there may not be a large enough range of redshift within the observed frequency range.  Despite this, the planned specifications for SKA might enable us to make high fidelity maps of the matter distribution all the way back to $z\\simeq 10$ and, if enough area can be surveyed, very good statistical information would be possible enabling very high precision measurements of cosmological parameters.  Realistic upgrades to the collecting area and array size of the next generation of telescopes would greatly improve their sensitivity to lensing."
    },
    "0801/0801.1558_arXiv.txt": {
        "abstract": "Recent N-body simulations have shown that the assembly history of galactic halos depend on the density of large-scale environment. It implies that the galaxy properties like age and size of bulge may also vary with the surrounding large-scale structures, which are characterized by the tidal shear as well as the density. By using a sample of $15,882$ well-resolved nearby galaxies from the Tully Catalog and the real space tidal field reconstructed from the 2Mass Redshift Survey (2MRS), we investigate the dependence of galaxy morphological type on the shear of large-scale environment where the galaxies are embedded. We first calculate the large scale dimensionless overdensities ($\\delta$) and the large-scale ellipticities ($e$) of the regions where the Tully galaxies are located and classify the Tully galaxies according to their morphological types and create subsamples selected at similar value of $\\delta$ but span different ranges in $e$. We calculate the mean ellipticity, $\\langle e\\rangle$,  averaged over each subsample and find a signal of variation of $\\langle e\\rangle$ with galaxy morphological type: For the case of $0.5\\le\\delta\\le1.0$, the ellipticals are found to be preferentially located in the regions with low ellipticity. For the case of $-0.3\\le\\delta\\le 0.1$, the latest-type spirals are found to be preferentially located in the regions with high ellipticity. The null hypothesis that the mean ellipticities of the regions where the ellipticals and the latest type spirals are located are same as the global mean ellipticity averaged over all types is rejected at $3\\sigma$ level when $-0.3\\le\\delta\\le 0.1$. Yet, no signal of galaxy-shear correlation is found in the highly overdense/underdense regions. The observed trend suggests that the formation epochs of galactic halos might be a function not only of halo mass and large-scale density but also of large-scale shear. Since the statistical significance of the overall trend is low, it will require a sample of at least $100,000$ galaxies to verify the existence of this correlation. ",
        "introduction": "It has been known for long that the physical properties of galaxies such as morphological type, color, luminosity, spin parameter, star formation rate, concentration parameter and so on are functions of their environments \\citep{dre80,pos-gel84,whi-etal93,lew-etal02,gom-etal03,got-etal03,roj-etal05, kue-ryd05,bla-etal05,ber-etal06,cho-etal07,par-etal07}. Most previous works have largely focused on the galaxies located in the highly dense regions and usually quantified the environmental dependence of galaxy properties in terms of cross-correlations with local density on small scale, which is  believed to be established by environment-dependent processes like galaxy-galaxy interaction.  The currently popular $\\Lambda$CDM model predicts that the galaxy properties are correlated with not only the small scale but also the large-scale environment. Here the small scale and the large-scale environments represent the surrounding regions smaller and larger than $5h^{-1}$Mpc, respectively. Recent high-resolution N-body simulations of a $\\Lambda$CDM cosmology have demonstrated that the assembly history of galactic halos of a given mass change with the density field smoothed on sufficiently large scale \\citep{gao-etal05}. Since the galaxy content depend strongly on assembly history of its host halo \\citep{spr-etal05,cro-etal07}, this numerical finding suggests that the intrinsic properties of galaxies may be correlated with the density of large-scale environments.  In fact, several groups have already found observational evidences for the existence of cross-correlations between galaxy properties and large-scale tidal shear. \\citet{nav-etal04} found that the spin axes of spiral galaxies near the Local Supercluster lie preferentially on the Supergalactic Plane. \\citet{tru-etal06} also showed by analyzing recent observational data from large galaxy surveys that the spin axes of void galaxies tend to be inclined to the void surfaces, keeping the initial memory of tidally induce alignments. \\citet{pan-bha06} have shown that the galaxy luminosity and colour are dependent on the filamentarity of large-scale environment. More recently, \\citet{pan-bha08} have also presented that the star formation rate of galaxies are cross-correlated with the filamentarity of the surrounding large-scale structures. \\citet{her-etal07} found a one-to-one correspondence between the value of the galaxy spin parameter and the morphological type, while \\citet{her-cer06} and \\citet{cer-her08} have noted that the spin parameter of the observed galaxies is strongly correlated with their color and chemical abundance.  The existence of cross-correlations between galaxy properties and density of large-scale environment implies that the galaxy content still have the memory of the initial conditions of the Lagrangian regions where the host halos formed. In the scenario of the biased galaxy formation, the galaxy sites correspond to the high peaks of the linear density field \\citep{kai84,bar-etal86}. But, since the initial density peaks are not spherical, they should be characterized not only by the peak heights but also by their ellipticity and prolateness. Hence, if the galaxy properties are still linked to the large-scale density, then it is likely that they are also linked to the large-scale ellipticity and prolaticity.  Our goal here is to test this statistical link. Instead of the ellipticity of the large-scale density field, however, we consider the ellipticity of the large-scale potential field (i.e., the tidal shear). Since  the potential field is smoother than the density field, we believe that the large-scale potential field reflects more directly the linear conditions. Furthermore, the tidal shear field, defined as the second derivative of the potential field, has a large-scale coherence, resulting in highly anisotropic spatial distribution of galaxies at present epoch \\citep{bon-etal96}. Therefore, the tidal shear field is more readily measurable from the large-scale spatial distribution of present galaxies.  In previous approaches, however, the correlation between galaxy properties and the shear of large-scale structure was measured in redshift space. Since the redshift distortion effect could contaminate the measurements of the filamentarity of the large scale structure, it will be desirable to measure it in real space. Moreover, true as it is that in the linear regime the density and the tidal shear at a given region are mutually independent, it is likely that the two quantities have developed cross-correlations in the subsequent non-linear regime. Therefore, to find an independent relationship between galaxy properties and large-scale tidal shear, one has to first account for the density-shear correlations and remove its effect.  We attempt here to measure observationally an independent relationship between the morphological types of nearby large galaxies and the real-space tidal shear reconstructed from all sky survey. The organization of this paper is as follows:  In \\S 2, the observational data are described. In \\S 3, the shear of environment is defined and its mean value averaged over each galaxy subsample at similar density is measured as a function of galaxy morphological type. In \\S 4, the results are discussed and a final conclusion is drawn. ",
        "conclusions": "By measuring observationally the mean ellipticities of large-scale structures where the nearby galaxies are embedded as a function of galaxy morphological types, we have tentatively detected an independent relationship between galaxy morphology and shear of environment: In the mildly overdense environment with large scale dimensionless overdensity of $0.05\\le\\delta\\le1.06$, the ellipticals are found to prefer the low-shear region; In the mildly underdense environment with large scale density of $-0.3\\le\\delta\\le-0.1$, the latest-type spirals (Scd-Sm) are found to prefer the high-shear region. No correlation signal, however, is found either in the highly overdense ($\\delta> 1$) or in the highly underdense environment ($\\delta<-0.5$).  This observational results may be explained by the dependence of halo formation epochs on the shear of large-scale environment. It has been known that the formation epochs of galactic halos depend not only on their mass but also on the density of large-scale environment \\citep{gao-etal05}. Here we argue that there is some observational evidence to support the hypothesis that the formation epochs of halos also depend on the shear of the large-scale environment. In the high-shear environment the halos formed relatively early without growing at late time due to strong tidal disruption from the surrounding matter distribution. Thus, the late-type spirals are likely to be located in this high-shear environment. Meanwhile in the low-shear environment the halos formed relatively recently since there is no strong tidal distribution, having grown at late times through hierarchical merging. Hence, the ellipticals are likely to be found in this low-shear environment. In the highly overdense regions, however, even in case that there is tidal disruption from the surrounding matter distribution, the other stronger effects of environmental processes should mask the shear-dependence of galaxy properties. Likewise in the highly underdense regions, the tidal field is too weak to produce any significant galaxy-shear correlation. It will be interesting to test against N-body simulations whether or not the formation epochs of galactic halos with similar mass at similar density is a function of the ellipticity of large-scale dark matter distribution. Our future work is in this direction.  It should be noted here that our results suffer from low statistical significance due to the small sample size. We have seen only a marginal effect of the large-scale shear for the latest-type spirals (Scd-Sm). Meanwhile no signal has been found for the Sb-Sc galaxies even though both types of galaxies occupy similar regions. To remove the effect of the stronger galaxy-density correlations, we had to constrain the overdensities of galactic regions, which resulted in large errors.  The null hypothesis that there is no galaxy-shear correlation over the six bins is found to be rejected at only the $10\\%$ level due to large errors. Given our result, it will require a sample of more than $100,000$ galaxies to test the hypothesis at $>90\\%$ confidence level. This line of investigation will provide new insight into the formation and evolution of galaxies in a filamentary cosmic web."
    },
    "0801/0801.1072_arXiv.txt": {
        "abstract": "In this paper  we present the results of two detailed N-body simulations of the interaction of a sample of four massive globular clusters in the inner region of a triaxial galaxy. A full merging of the clusters takes place, leading to a slowly evolving cluster which is quite similar to observed Nuclear Clusters. Actually, both the density and the velocity dispersion profiles match qualitatively, and quantitatively after scaling, with observed features of many nucleated galaxies. In the case of dense initial clusters, the merger remnant shows a density profile more concentrated than that of the progenitors, with a central density higher than the sum of the central progenitors central densities. These findings support the idea that a massive Nuclear Cluster may have formed in early phases of the mother galaxy evolution and lead to to the formation of a nucleus, which, in many  galaxies, has indeed a luminosity profile similar to that of an extended King model. A correlation with galactic nuclear activity is suggested. ",
        "introduction": "In the first paper of this series devoted to the study of the interaction of globular clusters (GCs) in triaxial galaxies, we analyzed the face-on collision between two GCs moving on quasi-radial orbits in the central galactic region, with the aim, also, to understand how effective is the tidal distortion. Our first finding was that the face-on collision is practically uneffective respect to the role of the external tidal field. Actually, the tidal erosion has been shown to be such to destroy {\\it loose} GCs (initial King concentration parameter $c<1$) in few passages through the galactic centre, while \\textit{tight} cluster ($c\\geq 1.6$) keeps bound a substantial amount of their mass up to the complete orbital decaying. At this latter regard, another important result is that the orbital energy dissipation due to the tidal interaction is of the same order of that caused by dynamical friction. At the light of these results, and given that dynamical friction was shown to be important in segregating massive GCs in triaxial potentials (\\citealt*{pesce}; \\citealt{capdol93}; \\citealt{cv05}), a natural further step in the present program of investigation is to study the possible merging of a set of GCs decayed in the central galactic region to see whether a sort of \\lq super star cluster \\rq (SSC) results from the merging, and what its morphological and dynamical characteristics could be.  It is quite ascertained the existence of very bright ($10^7\\div 10^8 $L$_\\odot$) clusters of stars in the central region of galaxies across the Hubble sequence thanks to VLT and HST observations \\citep{walch05,vandermarel,wehner}. On the basis of integrated colours and of the estimated $M/L$ ratios, these clusters are usually thought to be young or, at least, to contain a significant population of young stars. In any case, that of the age is a controversal point, because, for instance, the nuclear star clusters (NCs) in M82 show evidence for mass segregation despite their spectra are well fitted by stellar population synthesis models with ages $10$--$50$ Myr \\citep{McCr}. \\citet{bok04} fitted analytical models to Hubble Space Telescope images of 39 NCs in order to determine their effective radii after correction for the instrumental point-spread function. They compared the luminosities and sizes of NCs to those of other ellipsoidal stellar systems, in particular the Milky Way globular clusters, finding for NCs a narrow size distribution statistically indistinguishable from that of Galactic GCs, even though the NCs are, on average, 4 mag brighter than the old GCs. They discuss some possible interpretations of the similarity among NCs and Galactic GCs and, from a comparison of NCs luminosities with various properties of their host galaxies, find that more luminous galaxies harbor more luminous NCs. It remains unclear whether this correlation reflects the influence of galaxy size, mass, and/or star formation rate.  However, \\citet{rossa06}, by means of spectroscopic HST/STIS data, derived NCs ages for 19 galaxies and found that they range from 10 Myr to 10 Gyr, with a non negligible presence of old clusters. A comparable result has been deduced also for 9 NCs of very late-type, bulge-less spirals, by \\citet{walch06} by means of a high-resolution spectroscopic survey with VLT/UVES.  \\citet{genz03}, in their study of the stellar cuspy distribution around the Supermassive Black Hole (SBH) in the Galactic center, found that the K-band luminosity function of the local NC (within 9'' of Sgr A*) resembles that of the large-scale Galactic bulge except for showing an excess of stars at $K_s \\leq 14$. It fits with population synthesis models of an old, metal-rich stellar population with a contribution from young, early, and late-type stars at the bright end. In the central arcsecond, they argued that a stellar merger model is the most appealing explanation. These stars may thus be ``super blue-stragglers'', formed and ``rejuvenated'' through mergers of lower mass stars in the very dense  ($\\geq 10^8 \\msol$ pc$^{-3}$) environment of the cusp.  Another intriguing piece of the puzzle of the structure of SSCs is given by \\citet{baumg03}. Through a comparison between the observational data on the kinematical structure of the very bright cluster G1 in M31 ---obtained with the HST WFP Camera 2 and Space Telescope Imaging Spectrograph instruments--- and their results of dynamical simulations carried out using the special purpose computer GRAPE-6, they obtained very good fits when starting simulations with initial conditions extracted from the end product of a previous simulation concerning a merging between two pre-existent star clusters. The merging explains observed features without the need to invoke the presence of an intermediate-mass black hole in the center of G1.  Not many simulations have been presented in the literature that study the possible formation scenarios for SSCs. Among these, we remind those by \\citet{fell02,fell05} finding that star clusters aggregates, like the ones found in the Antennae or Stephan Quintet, are very likely the merger progenitors of SSCs. % Interestingly enough, these authors also find that the resulting SSC is a stable and bound object, whose density profile is well fitted by a King profile, even if the mass loss of the merger product occurs through every perigalacticon passage. ",
        "conclusions": "In this paper we studied the modes of interaction of a sample of few (4) globular clusters whose orbital motion is limited to the core of the galactic triaxial field, because they have experienced a significant and rapid dissipation of the orbital energy by dynamical friction acting on their large initial masses. We studied 2 sets of 4 GCs of different initial conditions, characterised by a higher (case 1) and lower (case 2) central density. Dynamical friction was properly taken into account by the triaxial generalization of the classical Chandrasekhar formula \\citep{pesce} and it is the main cause of orbital energy dissipation up to the beginning of the merging process, when the individual cluster sizes are comparable with the orbital size. Since this stage on, the main responsible of the residual orbital energy loss is the tidal torque. The merger is completed in $\\sim 18$ galactic core crossing times, i.e. in a time much shorter than the Hubble time and, in any case, small compared to the total orbital decaying time.  The merged system keeps some of the characteristics of the preexisting objects, attaining a relaxed structure which has, in the case 2, a mildly triaxial shape (inherited by the environment) and a halo which is diffusing in the external field assuming its phase space properties. The case 1 results into a merger configuration which conserves better the spherical simmetry of the \\lq progenitors\\rq\\ and whose radial density profile is more concentrated to the center than that expected on the basis of the mere space superposition of the 4 progenitors. This is, likely, the consequence of some violent relaxation of the merged cluster to the state of a quasi-stationary Super Star Cluster, whose stellar density distribution maintains a King shape and is very similar to Nuclear Clusters observed at the center of many galaxies. Even in the limits of our simulations (both on the number of merging objects and on the time extension of the integration), we may infer from a comparison of our results with the characteristics (surface brightness profile, integrated light and velocity dispersion) of the nuclei of many galaxies \\citep{bok02, geha, walch05, cote06} that such nuclei may have actually formed by the merging of few tens (when merging progenitors are quite compact) up to few hundreds (when progenitors are looser) massive GCs decayed orbitally into the inner galactic regions.  As remarkable result, the radial profile of the line-of-sight velocity dispersion of the merger remnant shows the same minimum at the galactic center as actually found by observations of a sample of Virgo cluster nucleated dEs \\citep{geha} and of late-type spirals \\citep{walch05}. Finally, an important theoretical result is that high densities ($\\gtrsim 10^7 \\msol$ pc$^{-3}$), critical to allow a significant accretion onto a massive object, may be reached by mean of a merging of less than $50$ clusters of the type considered in case 1. This may have important implications on the role of NCs in supplying the innermost galactic activity."
    },
    "0801/0801.3846_arXiv.txt": {
        "abstract": "{ We describe the design and implementation of an autonomous adaptive software agent that addresses the practical problem of observing undersampled, periodic, time-varying phenomena using a network of HTN-compliant robotic telescopes. The algorithm governing the behaviour of the agent uses an optimal geometric sampling technique to cover the period range of interest, but additionally implements proactive behaviour that maximises the optimality of the dataset in the face of an uncertain and changing operating environment. ",
        "introduction": "\\label{Section:Intro}    The eSTAR project \\citep{allan04_estar_spie, allan07_estar_htn_iii} is an agent-based software system that aims to establish an intelligent robotic telescope network for efficient and automated observing. The system comprises \\emph{user agents} that run an observing programme on behalf of an astronomer, and \\emph{embedded agents} that provide the interface to observational resources such as telescopes. One role of eSTAR is to provide the software infrastructure for Robonet-1.0 \\citep[e.g.][]{bode04robonet_canonical, mottram06_high_level_robonet}, a network of three autonomous 2-m robotic telescopes strategically located across the globe. eSTAR implements the HTN\\footnote{Heterogeneous Telescope Networks} protocol \\citep{allan06_htn_protocol_standard} using the XML dialect RTML \\citep{pennypacker02_rtml2.1, hessman06_rtml3}. The Robonet-1.0 network thus provides a full reference implementation of an HTN-compliant environment.    This paper describes the design and implementation of a user agent that applies adaptive scheduling to optimise an observing run. Adaptive scheduling was described in an earlier paper \\citep{saunders06_agent_metrics} as ``techniques... to allow an autonomous software entity to calculate an idealised observing strategy, while remaining able to respond flexibly to scheduling constraints outside of its immediate control.'' In particular, the adaptive scheduling agent (ASA) implements the optimal geometric sampling strategy for period searching described in \\citet{saunders06_optimal_placement} to address the practical problem of observing undersampled, periodic, time-varying phenomena.    An example of the type of time-domain astronomy that this agent was designed to perform is the surveying of variability in star-forming regions. Especially in the youngest clusters, a large fraction of the cluster members are T-Tauri stars that can exhibit significant variability due to rotational modulation of features at the stellar surface. However the range of periods among cluster members is large, ranging from a few hours to many days. Additionally, aliasing, particularly as a consequence of diurnal sampling, is usually a problem for data obtained in the ``classical paradigm'', where an on-site observer performs observations using a single telescope. A robotic network spread across longitude can break this aliasing pattern, but datasets are typically undersampled with respect to shorter periods in the range of interest. Given a limited number of observations, the optimal geometric sampling technique determines the best time to make individual observations in order to achieve similar sensitivity to periods across the full range of interest. Of particular importance is the property that as long as the set of gaps between observations remains unchanged, observations may be arbitrarily reordered \\citep{saunders06_optimal_placement}. This property is fully exploited by the adaptive algorithm.    Unfortunately, observations are not guaranteed. In the dispatch model of telescope scheduling, users request observations but it is up to the telescope scheduler to decide whether such requests are honoured \\citep{fraser06_robonet_scheduling}. Observations may still fail even when submitted to an entirely cooperative scheduler, due to the possibility of telescope downtime or inclement weather. Thus any system seeking to reliably implement an optimal geometric sampling technique must in practice deal with such observation uncertainty. ",
        "conclusions": "We have built an adaptive scheduling agent that implements the optimal geometric sampling strategy for period searching described by \\citet{saunders06_optimal_placement} in the context of an HTN-compliant network of telescopes. The algorithm automatically penalises the agent for choosing too many short gaps, but is similarly hostile to the chasing of large gaps when there are shorter, more valuable alternatives. In this way a balance point is determined that unambiguously defines the best option for the agent at any point in the run. The deceptively simple rule sequence for finding the optimality implicitly utilises the reordering property of optimal sequences, maximises the agent's cautiousness in the face of uncertainty, and provides a computationally cheap and scalable way to unambiguously calculate the best choice of action at any point. It is also stable --- a snapshot of the run at any point in time is optimal for that subset of points, and minimises the effects of observation failure on the sequence. In the case of perfect observing behaviour (no failures), the optimal sequence is naturally recovered. Importantly, the degree of failure exhibited by the network can also change dynamically without adversely affecting the algorithm, because it makes no assumptions about the stability of the network, and makes no attempt to model the behaviour of the telescopes on which it relies."
    },
    "0801/0801.1905_arXiv.txt": {
        "abstract": "We show an interesting correlation between the surface brightness and temperature structure of the relaxed clusters \\rxj1720 ~ and \\ms1455 , hosting a pair of cold fronts, and their central core--halo radio source. We discuss the possibility that the origin of this diffuse radio emission may be strictly connected with the gas sloshing mechanism suggested to explain the formation of cold fronts in non major merging clusters. We show that the radiative lifetime of the relativistic electrons is much shorter than the timescale on which they can be transported from the central galaxy up to the radius of the outermost cold front. This strongly indicates that the observed diffuse radio emission is likely produced by electrons re--accelerated via some kind of turbulence generated within the cluster volume limited by the cold fronts during the gas sloshing. ",
        "introduction": "\\label{par:intro}  A number of cool core clusters of galaxies host very peculiar diffuse radio sources at their center characterized by a {\\it core--halo} (CH herafter) structure (e.g. PKS\\,0745--191;  \\citealt{1991MNRAS.250..737B}). These objects show an amorphous steep spectrum halo which surrounds a bright core coincident with the dominant galaxy, with no jets on the kpc scale. Even though the connection between CHs and cool core clusters is well established, the origin of these sources remains uncertain. \\cite{1995ApJ...451..125S} argued that they might be produced by the interaction between the radio source and the thermal gas cooling at the cluster center, i.e. in the form of disruption of the jets by the high--pressure ambient medium and buoyant effects on the radio plasma from the disrupted jets. In this Letter we show an interesting spatial correlation between the CH emission and X-ray features in two galaxy clusters \\rxj1720 ~ ($z=0.159$)  and \\ms1455 ~ ($z=0.258$), both showing a pair of cold fronts (CF hereafter) in their core. We suggest a connection between the origin of the CH structure and the same gas sloshing mechanism responsible for the formation of the CFs. We use $H_0=70$~km~s$^{-1}$~kpc$^{-1}$, $\\Omega=0.27$, and $\\Lambda=0.73$. ",
        "conclusions": "\\label{par:discussion}  Using high resolution simulations differet authors showed that CFs in the core of relaxed galaxy clusters may be the result of gas sloshing induced by a minor merger with a small dark matter halo occurred within the last few Gyrs  (see e.g \\citealt{2006ApJ...650..102A} --AM hereafter-- and references therein). The only condition for the CF to form is that the cluster has a steep entropy profile (as observed in cool core clusters) and the dark matter sub--halo has a non--zero impact parameter. When the sub--halo flies through the main cluster, it produces a gravitational disturbance that pushes the cool gas away from its initial position. This makes the gas slosh and two or more CFs form on opposite side with respect to the cluster center. AM noticed that, when the gas is displaced from the center for the first time, it does not fall back radially, and the subsequent CF(s) is (are) not exactly concentric but combines into a spiral pattern (see 1.9--2.1 Gyrs panels in Fig.~7 of AM). As shown in \\S~\\ref{par:image}, both the CF pairs in \\rxj1720 ~ and \\ms1455 , which are clearly two distinct non concentric structures, seem to combine into a spiral pattern very similarly to what predicted by AM. These structures are remarkably similar to the simulated ones in the surface brightness map (compare left panels of Figs.~\\ref{fig:ratio_rxj1720} and \\ref{fig:fig:ratio_ms1455} with upper--left panel of Fig.~19 of AM) as well as in the temperature map (compare right panels of Figs.~\\ref{fig:ratio_rxj1720} and \\ref{fig:fig:ratio_ms1455} with lower--left panel of Fig.~19 of AM). Thus the X-ray observations of these two clusters seems to confirm the prediction of AM for the gas sloshing.  What appear to be very interesting is that both clusters host a CH radio source. These sources seem to be well confined (in projection) within the CF pairs of the respective cluster. Furthermore, we noticed a spatial correlation between the radio emission and the X-ray spiral structure which might suggest that the radio emitting electrons are trapped into the same ICM flows that produced the CFs. Given the lack of jets on the kpc scale, the relativistic electrons need to be transported by some other mechanism from the central galaxy up to the radii of the outermost CFs. We might speculate that the relativistic electrons injected by the central AGN, and trapped into the ICM magnetic field, are transported at the CFs radii by the motion of the low entropy thermal gas induced by the gas sloshing. In fact, the AM simulations show that the CFs are formed by the continuous circulation of lower entropy gas that comes from the inner core region and falls back soon after (see Fig~7 of AM). However, AM show that the speed of this gas circulation is relatively low with an upper limit of $v_{max}=500~{\\rm km/s}$. In both clusters the outermost CF is at a radial distance of $\\Delta r \\approx 150~{\\rm kpc}$. Assuming  radial transport, we estimate that the minimum time needed for the relativistic electrons to be displaced up to the position of the outermost front is $t_{min}\\ge 3\\times 10^{8}\\, {\\rm yr}$. This time is much larger than the radiative lifetime of relativistic electrons which suffer energy losses mainly due to Inverse Compton with the cosmic microwave background photons ($t\\approx 10^{7-8}/(1+z)^4\\, {\\rm yr}$; see e.g. \\citealt{2001MNRAS.320..365B}). This simple calculation suggests that the observed radio emission might be produced by re--acceleration of a population of relic electrons, injected into the ICM by a previous activity of the central AGN. Due to the subsonic nature of the gas sloshing, we can exclude that the re--acceleration may be driven by  thermal shock. An alternative and plausible mechanisms may be the re--acceleration by MHD turbulence, as suggested to explain the cluster mini--halos (\\citealt{2002A&A...386..456G}) and giant radio halos (e.g. \\citealt{2004MNRAS.350.1174B}).  These authors claim that, to be effective, this re--acceleration mechanism requires only a modest level (few per cent) of turbulence. In fact, 2D simulations with much higher resolutions have shown that the gas sloshing could automatically create turbulence in a core (e.g. \\citealt{2004ApJ...612L...9F}).  It is worth noting that the CH/CFs correlation described in this Letter is not limited to these two clusters but seems to be present also in other clusters hosting a CH, as e.g. 2A\\,0335+096 (\\citealt{2003ApJ...596..190M}), PKS\\,0745--191 (\\citealt{1991MNRAS.250..737B}), and A\\,2052 (\\citealt{1993ApJ...416...51Z}). These results calls for further and deeper investigation of the radio--X-ray connection in relaxed clusters with CFs, which will be addressed in forthcoming papers."
    },
    "0801/0801.2698_arXiv.txt": {
        "abstract": "{We present the analysis of single epoch long slit spectra of the three brightest images of the gravitationally lensed system J1131-1231. These spectra provide one of the clearest observational evidence for differential micro-lensing of broad emission lines (BELs) in a gravitationally lensed quasar. The micro-lensing effect enables us: (1) to confirm that the width of the emission lines is anti-correlated to the size of the emitting region; (2) to show that the bulk of \\FeII~is emitted in the outer parts of the Broad Line Region (BLR) while another fraction of \\FeII~is produced in a compact region; (3) to derive interesting informations on the origin of the narrow intrinsic \\MgII\\, absorption doublet observed in that system.}  \\resumen {}   \\begin{document} ",
        "introduction": "J1131-1231 is one of the nearest gravitationally lensed AGN. The lensing galaxy ($z_l = 0.295$) splits the light rays from the source ($z_s = 0.66$) into four macro-images: three bright images (A-B-C) separated on the sky by typically 1$\\arcsec$ and a fainter component (D) located at 3.6$\\arcsec$ from A (Sluse et al. \\citealp{SLU03}). The present study focuses on the lensed images A-B-C. The lensing effects in a system like J1131-1231 occur for two different regimes. First, the lensing galaxy produces four {\\it resolved macro-images} of the background source. Second, the compact masses in that galaxy (typically $10^{-6} < M < 10^6\\,M_\\odot $) split each macro-lensed image into multiple of {\\it unresolved micro-images} separated by a few micro-arcseconds. Because gravitational lensing magnifies the source, each (unresolved) macro-image $i$ is amplified by a factor $M_i$ associated with macro-lensing and by a factor $\\mu_i$ due to micro-lensing. The typical length scale for micro-lensing (in the quasar plane) is the Einstein radius $R_E$ of the micro-lens: \\begin{equation} \\label{equ:micro} R_E= \\sqrt{\\frac{4GM}{c^2}\\frac{D_{ls}D_{os}}{D_{ol}}}=14.3 \\sqrt{\\frac{M h^{-1}}{M_{\\odot}}}\\,{\\rm light}-{\\rm days}, \\end{equation} \\noindent where $D_{os}$, $D_{ls}$, $D_{ol}$ are the angular-size distances between observer and source (os), lens and source (ls) and observer and lens (ol) (see e.g. Wambsganss~\\citealp{WAM06} for a review on micro-lensing). Micro-lensing acts as a magnifying glass which amplifies regions of the source on scales smaller or equal to the micro-lens Einstein radius. For J1131-1231, a solar mass Einstein radius has a size similar to the usually assumed size of the BLR, meaning that regions as large as the BLR can be amplified by micro-lensing. Because micro-lensing occurs independently in each macro-lensed images, we can track for spectral differences between multiple images of a lensed quasar. Such differences reveal the selective micro-amplification of small regions of the source. Unfortunately, spectral differences between macro-lensed images are not only due to micro-lensing. The time delay between the macro-lensed images can also induce spectral differences between images (because each lensed image is a snapshot of the source at a different epoch). Hopefully, the time delay between images A-B-C of J1131-1231 is less than a few days (e.g. Saha et al.~\\citealp{SAH06}) and thus can be neglected. Differential reddening between lensed images due to the lensing galaxy might also exist. This effect seems however to be negligible in the case of J1131-1231 (Sluse et al.~\\citealp{SLU06}, ~\\citealp{SLU07}).  After a brief description of the data (Sect.~\\ref{sec:data}), we interpret the main spectral differences observed between images A-B-C of J1131-1231 in the micro-lensing framework (Sect.~\\ref{sec:results}). More results can be found in Sluse et al. (\\citealp{SLU07}; SLU07). ",
        "conclusions": ""
    },
    "0801/0801.1187_arXiv.txt": {
        "abstract": "{}{% To understand the origin of the \\lya\\ line profile of the prototypical Lyman break galaxy (LBG) MS 1512--cB58 and other similar objects. To attempt a consistent fit of \\lya\\ and other UV properties of this galaxy. To understand the main physical parameter(s) responsible for the large variation of \\lya\\ line strengths and profiles observed in LBGs.} {% Using our 3D \\lya\\ radiation transfer code (Verhamme \\etal\\ 2006), we compute the radiation transfer of \\lya\\ and UV continuum photons including dust. Observational constraints on the neutral gas (column density, kinematics, etc.) are taken from other analysis of this object.} {% The observed \\lya\\ profile of MS 1512--cB58 is reproduced for the first time taking radiation transfer and all observational constraints into account. The observed absorption profile is found to result naturally from the observed amount of dust and the relatively high HI column density \\nh. Radiation transfer effects and suppresion by dust transform a strong intrinsic \\lya\\ emission with EW(\\lya) $\\ga$ 60 \\AA\\ into the observed faint superposed \\lya\\ emission peak. We propose that the vast majority of LBGs have intrinsically EW(\\lya) $\\sim$ 60--80 \\AA\\ or larger, and that the main physical parameter responsible for the observed variety of \\lya\\ strengths and profiles in LBGs is \\nh\\ and the accompanying variation of the dust content. Observed EW(\\lya) distributions, \\lya\\ luminosity functions, and related quantities must therefore be corrected for radiation transfer and dust effects. The implications from our scenario on the duty-cycle of \\lya\\ emitters are also discussed. }{} ",
        "introduction": "\\label{s_intro} Discovered serendipitously by \\citet{Yee96}, the Lyman Break Galaxy (LBG) MS 1512--cB58 (hereafter cB58) located at $z\\sim 2.73$ is strongly gravitationally magnified by the foreground cluster MS 1512+36. It is by far the best studied LBG, which has been observed at many wavelengths \\citep[see][]{Elli96,Fray97,Bech97,Naka97,Pett00,Tepl00,Bake01,Sawi01, Sava02,Pett02,Bake04,Tepl04}. In particular, rest-frame UV \\citep{Pett00,Pett02}, optical \\citep{Tepl00,Tepl04}, and radio \\citep{Bake01, Bake04} observations have allowed to draw a fairly global picture of this star-forming galaxy and its interstellar medium (ISM): it is young ($\\sim$ 50--300 Myr), massive ($\\sim 10^{10}$ \\msun), with a sub-solar metallicity ($Z \\sim 2/5$ \\zsun), a star formation rate (SFR $\\sim 24$ \\msunyr) typical of LBGs, and it is surrounded by outflowing material at $\\vexp \\sim 250$ \\kms.  Thanks to its apparent brightness the rest-UV spectra of this galaxy are of comparable quality as the best studied local starbursts. This has in particular made it possible to detect numerous stellar and interstellar (IS) lines, to study their kinematics, and to use them for detailed UV line fits to constrain its stellar content, age, star formation history etc. \\citep[e.g.][]{Elli96,Pett00,deMell00}. Apart from its brightness, magnified by a factor $\\sim$ 30 thanks to gravitational lensing \\citep{Seitz98},  the properties of cB58 are empirically found to be typical of LBGs \\citep{Shapley03}. More precisely, in the sample of \\citet{Shapley03} cB58 is found in the quartile ($\\sim 25$ \\%) of LBGs with the strongest \\lya\\ absorption. As such a representative and given the quality and amout of observational data available, cB58 is of particular interest.  Our objectives are to examine if the observed \\lya\\ line profile and strength of cB58 can be understood, and if one is able to obtain consistent diagnostics from the \\lya\\ line and from the other UV and broad band features used in earlier papers. Such an attempt is carried out here for the first time with a \\lya\\ radiation transfer code \\citep{Verh06}. Other LBGs with varying strengths of \\lya\\ emission will be analysed in a subsequent paper \\citep{Verh07}. In particular we wish to elucidate what causes the broad \\lya\\ absorption and very weak superposed \\lya\\ emission in cB58 and in other LBGs showing similar \\lya\\ spectra \\citep[cf.][]{Shapley03}. Similarly we want to understand what physical parameter(s) distinguish the different subtypes of LBGs showing large variations especially in \\lya. Finally, fitting observed starburst spectra with detailed \\lya\\ radiation transfer modeling should in general provide insight for the use of \\lya\\ as a diagnostic in distant galaxies.  The remainder of the paper is structured as follows. The main observational properties of cB58 are summarised in Sect.\\ \\ref{s_obs}. The principles of our simulations and fit method are described in Sect.\\ \\ref{s_fit}. In Sect.\\ \\ref{s_results} we discuss the implications from our results on the understanding of \\lya\\ emission in LBGs, and on the properties of these galaxies. Our main conclusion are summarised in  Sect.\\ \\ref{s_conclude}. ",
        "conclusions": "\\label{s_conclude} A 3D \\lya\\ and UV continuum radiation transfer code \\citep{Verh06} has for the first time been applied to the prototypical Lyman break galaxy MS 1512--cB58 at $z=2.7$. Since this is one of, if not the best studied LBG, and since it is part of the quartile of LBGs showing predominantly \\lya\\ in absorption \\citep[cf.][]{Shapley03}, a detailed \\lya\\ line profile analysis including radiation transfer and dust effects is of great interest.  Three different geometries were explored for the material surrounding the central starburst; a spherically symmetric expanding shell with velocity \\vexp, a foreground slab moving towards the observer, and a geometry mimicing a shell with a lower velocity in the back. All available observational constraints were used (see Table \\ref{tab_pars}). In particular these include: the velocity \\vexp\\ of the foreground material, its Doppler parameter $b$, column density \\nh, and the observed extinction.  The width of the intrinsic \\lya\\ emission was taken from the observed FWHM of \\ha. The two free parameters were: the equivalent width of the intrinsic \\lya\\ emission, and the velocity \\vback\\ of the background slab (if applicable).  Our radiation transfer calculations have confirmed what could be expected from our earlier expanding shell simulations and from the lower-than-usual observed velocity shift between the peak of ``remnant'' \\lya\\ emission and the blue IS absorption lines (cf.\\ Sect.\\ \\ref{s_ism}); the \\lya\\ line profile of cB58 cannot be reproduced by an expanding shell with isotropic constant velocity \\vexp. Considering, however, a lower shell velocity in the back allowed us to obtain excellent fits to the observed \\lya\\ line profile.  Interestingly, the input spectrum requires a fairly strong \\lya\\ emission, with a lower limit of EW(\\lya) $\\ga 60$ \\AA. Radiation transfer and the suppresion of photons by dust are responsible for transforming such an input spectrum into the observed \\lya\\ absorption profile with a superposed faint emission peak. In this way the observed \\lya\\ line strength and profile can be reconciled with the strong intrinsic \\lya\\ emission expected from the approximately constant star-formation history of cB58 derived from earlier detailed UV spectral analysis \\citep[cf.][]{Pett00,deMell00}.  In fact we suggest that cB58 and most other LBGs have intrinsically EW(\\lya) $\\sim$ 60--80 \\AA\\ or larger, and that the main physical parameter responsible for the observed variety of \\lya\\ strengths and profiles in LBGs is \\nh\\ and the accompanying variation of the dust content (see Sect.\\ \\ref{s_discuss}).  This explains not only the absorption-dominated object cB58, but also observed correlations between E(B-V) and EW(\\lya) in LBGs, the absence of bright LBGs with strong \\lya\\ emission, and other correlations.  Among the implications from our work are that observed EW(\\lya) distributions and \\lya\\ luminosity functions must be corrected for radiation transfer and dust effects. Furthermore relatively short duty cycles, suggested earlier in the literature, are not required to explain the variations observed between different LBG types. Our proposed unifying scenario will be detailed further and subjected to additional tests in a subsequent publication on a sampe of LBGs with strong \\lya\\ emission \\citep{Verh07}. Providing a clearer picture of the physical links between the observables including \\lya\\ line strength and profile, and the star-formation, host galaxy, and outflow properties and evolution of LBGs and LAEs remains an objective for the near future."
    },
    "0801/0801.3182_arXiv.txt": {
        "abstract": "We discuss results from numerical simulations of star cluster formation in the turbulent interstellar medium (ISM). The thermodynamic behavior of the star-forming gas plays a crucial role in fragmentation and determines the stellar mass function as well as the dynamic properties of the nascent stellar cluster. This holds for star formation in molecular clouds in the solar neighborhood as well as for  the formation of the very first stars in the early universe.  The thermodynamic state of the ISM is a result of the balance between heating and cooling processes, which in turn are determined by atomic and molecular physics and by chemical abundances. Features in the effective equation of state of the gas, such as a transition from a cooling to a heating regime, define a characteristic mass scale for fragmentation and so set the peak of the initial mass function of stars (IMF).  As it is  based on  fundamental physical quantities and constants,  this is an attractive approach to explain the apparent universality of the IMF in the solar neighborhood as well as the transition from purely primordial high-mass star formation to the more normal low-mass mode observed today. ",
        "introduction": "Identifying the physical processes that determine the masses of stars and their statistical distribution, the initial mass function (IMF), is a fundamental problem in star-formation research. It is central to much of modern astrophysics, with implications ranging from cosmic re-ionisation and the formation of the first galaxies, over the evolution and structure of our own Milky Way, down to the build-up of planets and planetary systems.  Near the Sun the number density of stars as a function of mass has a peak at a characteristic stellar mass of a few tenths of a solar mass, below which it declines steeply, and for masses above one solar mass it follows a power-law with an exponent $dN/d{\\log}m \\propto m^{-1.3}$. Within a radius of several kpc % this distribution shows surprisingly little variation (Salpeter 1955; Scalo 1998; Kroupa 2001; Kroupa 2002; Chabrier 2003).  This has prompted the suggestion that the distribution of stellar masses at birth is a truly universal function, which often is referred to as the Salpeter IMF, although note that the original Salpeter (1955) estimate was a pure power-law fit without characteristic mass scale.   The initial conditions in star forming regions can vary considerably, even in the solar vicinity. If the IMF were to depend on the initial conditions, there would be no reason for it to be universal. Therefore a derivation of the characteristic stellar mass that is based on fundamental atomic and molecular physics would be highly desirable. In this proceedings contribution we argue that indeed the thermodynamic properties of the star-forming cloud material determine the characteristic mass scale for fragmentation and subsequent stellar birth.   The thermodynamic state of interstellar gas is a result of the balance between heating and cooling processes, which in turn are determined by fundamental atomic and molecular physics and by chemical abundances. The derivation of a characteristic stellar mass can thus be based on quantities and constants that depend solely on the chemical abundances in a molecular cloud. It also explains why deviations from the ``standard'' mode of star formation are likely to occur under extreme environmental conditions such as those occurring in the early universe or in circum-nuclear starburst regions. ",
        "conclusions": "\\label{conclusions}  In this proceedings paper we have discussed several studies, where the thermodynamic behavior of the interstellar medium plays a crucial role in fragmentation and subsequent stellar birth. These examples range from star formation in the solar neighborhood at the present day to the formation of the very first stars in the early universe, and support the idea that the distribution of stellar masses depends, at least in part, on the thermodynamic state of the star-forming gas. Dips in the effective EOS, such that the relation between temperature $T$ and density $\\rho$ changes from decreasing $T$ with increasing   $\\rho$ to increasing $T$ with increasing $\\rho$, i.e.\\ the transition from a cooling ($\\gamma < 1$) to a heating ($\\gamma >1$) regime, define a characteristic mass scale for fragmentation.  The thermodynamic state of interstellar gas is a result of the balance between heating and cooling processes, which in turn are determined by fundamental atomic and molecular physics and by chemical abundances. The derivation of a characteristic stellar mass can thus be based on quantities and constants that depend solely on the chemical abundances of the star forming gas. This is an attractive feature explaining the apparent universality of the IMF in the solar neighborhood as well as the transition from purely primordial high-mass star formation to the more normal low-mass mode observed today. Clearly more work needs to be done to investigate the validity of this hypothesis."
    },
    "0801/0801.3461_arXiv.txt": {
        "abstract": "We present a \\chandra\\ study of the hot ISM in the giant elliptical galaxy \\src. In common with other group-centred ellipticals, its temperature profile rises with radius in the outer parts of the galaxy, from $\\sim$0.7~keV at 2~kpc to $\\sim$0.9~keV by 20~kpc. However, within the central $\\sim$2~kpc the trend reverses and the temperature peaks at $\\sim$1.1~keV within the innermost 200~pc. Under the assumption of hydrostatic equilibrium, we demonstrate that the central temperature spike arises due to the gravitational influence of a quiescent central super-massive black hole. We constrain the black hole mass (\\mbh) to $(3.35^{+0.67}_{-0.95})\\times 10^9$\\msun\\ (90\\% confidence), in good agreement with stellar kinematics measurements. This is the first direct measurement of \\mbh\\ based on studies of hydrostatic X-ray emitting gas, which are sensitive to the most massive black holes, and is a crucial validation of both mass-determination techniques. This agreement clearly demonstrates the gas must be close to hydrostatic, even in the very centre of the galaxy, which is consistent with the lack of morphological disturbances in the X-ray image. \\src\\ is now one of only a handful of galaxies for which \\mbh\\ has been measured by more than one method. At larger radii, we were able to decompose the gravitating mass profile into stellar and dark matter (DM) components. Unless one accounts for the DM, a standard Virial analysis of the stars dramatically over-estimates the stellar mass of the galaxy. We find the measured J-band stellar mass-to-light ratio, $1.37\\pm0.10 M_\\odot L_\\odot^{-1}$, is in good agreement with simple stellar population model calculations for this object. ",
        "introduction": "\\label{sect_introduction} It is becoming increasingly clear that supermassive black holes (SMBHs) in the centres of galaxies are intimately involved in the evolution of their hosts. Models of galaxy formation suggest that energy injection from active galactic nucleus (AGN) outbursts, in particular during hierarchical assembly, can have a significant impact on the structure of the evolving galaxy \\citep[\\eg][]{silk98a,dimatteo05a}. In galaxy groups and clusters, cavities in the hot intra-cluster medium (ICM) are a direct, observable manifestation of an episodic, AGN-driven feedback process acting on the ICM \\citep[\\eg][]{birzan04a}. Such feedback has been implicated in global heating of the gas \\citep[\\eg][]{donahue06a}, redistributing metals through the ICM \\citep[\\eg][]{mathews04a} and invoked as a possible solution to the ``cooling paradox'' \\citep[][and references therein]{mathews03a}.  Evidence for SMBHs has been found in increasingly large numbers of non-active galactic nuclei, leading to the picture of ubiquitous SMBHs in % stellar spheroids which has begun to emerge \\citep[][for a review]{ferrarese05a}. The black hole mass (\\mbh) has been found to correlate with properties of the host, such as luminosity or mass of the bulge \\citep{kormendy95a,magorrian98a,mclure02a,marconi03a}, the central light concentration \\citep{graham01a}, or the central velocity dispersion (\\sigmac) of the galaxy \\citep{gebhardt00b,ferrarese00b,tremaine02a}. These observational relationships are starting to become a key constraint on galaxy-formation models \\citep[\\eg][]{robertson06a,granato04a}. There remain however, concerns over the exact shape of the relations. In particular, different authors have found different slopes for the  \\mbh-\\sigmac\\ relation \\citep{tremaine02a,ferrarese05a}, and there are suggestions it may deviate from a simple powerlaw relation \\citep{wyithe06a}. Possible disagreement in the masses of the largest SMBHs inferred from the \\mbh-\\sigmac\\ and \\mbh-bulge luminosity relations \\citep{lauer07a} raise questions over which relation is more fundamental. Clearly, further progress requires a careful assessment of possible systematic biases in the black hole mass measurements, especially in the high-mass regime.  A few independent techniques have evolved to measure the masses of SMBHs, each relying on its own set of simplifying assumptions. In galaxies with quiescent nuclei, the presence of the black hole is measured by its gravitational influence on the kinematics of either the stars or, if present, a central ionized gas disk \\citep[\\eg][]{gebhardt03a,macchetto97a}. Since gas disks are not found in every galaxy, stellar kinematical studies, based on measuring the line-of-sight velocity dispersion profile are more generally applicable. Such studies are, however, complicated by a strong degeneracy between orbital structure and the gravitating mass profile, particularly in slowly-rotating galaxies \\citep{binney82}. Although measuring high-order moments of the velocity dispersion profile allows the degeneracy to be partially broken \\citep{vandermarel93a}, the solution is not, in general, unique \\citep{valluri04a}. To obtain interesting constraints on the SMBH mass, therefore, additional constraints such as a restricted range of allowed orbits or smoothly-varying phase-space are required \\citep[\\eg][]{magorrian98a,gebhardt03a}. If these assumptions are unrepresentative, however, \\citeauthor{valluri04a} pointed out that the systematic errors in the \\mbh\\ determination can exceed the statistical errors.  Given the good agreement in the shape of the \\mbh-\\sigmac\\ relation measured using masses obtained with different methods \\citep[\\eg][]{gebhardt00c,ferrarese01a}, it seems unlikely that SMBH masses determined from state-of-the-art ``3-integral'' methods \\citep{vandermarel98a,gebhardt03a} are {\\em on average} in error. Nevertheless, a proper assessment of the systematics inherent in the mass determination, by comparing masses {\\em for the same black hole} measured using different techniques is of considerable importance. To date there are only a handful of galaxies for which recent stellar kinematical \\mbh\\ measurements have been compared with at least one reliable, independent technique. Stellar and gas dynamical estimates of \\mbh\\ have been found to agree reasonably in at least two systems, NGC\\thin 3379 and Cen\\thin A \\citep{shapiro06a,marconi06a}, although there is evidence of non-circular motions in NGC\\thin 3379, which can lead to a systematically mis-estimated \\mbh. Similar agreement has been found for M\\thin 81 \\citep{kormendy04a}, although the stellar kinematics measurement was based on ``2-integral'' models, which are potentially unreliable as they place constraints on the orbital structure which appear unrealistic, such as assuming the orbits are described entirely by only two integrals of motion \\citep[for more discussion see][]{ferrarese05a}. Non-circular gas motions may explain strong discrepancies between \\mbh\\ measured by these different techniques in the galaxies IC\\thin 1459 and NGC\\thin 4335 \\citep{cappellari02a,verdoeskleijn02a}. Despite this concern, the \\mbh-\\sigmac\\ relation is usually constructed by combining samples of galaxies in which \\mbh\\ has been determined from stellar kinematics and samples for which gas dynamics has been used \\citep[\\eg][]{tremaine02a}, and there is little evidence that the sub-samples obey appreciably different relations.   Other methods of determining \\mbh\\ are available in galaxies containing active galactic nuclei (AGNs), such as measuring the motions of masing gas \\citep{miyoshi95a} or ``reverberation mapping'', using the variability of the  AGN \\citep[\\eg][]{peterson04a}. In two cases \\mbh\\ determined from reverberation mapping has been compared to stellar kinematics results \\citep{davies06a,onken07a}, but since the uncertainties in the former technique lead to masses accurate only to within a factor of $\\sim$3 \\citep{onken04a}, these measurements do not provide good constraints on the validity of the stellar models.      A new method to identify SMBHs in giant elliptical galaxies was suggested by \\citet{brighenti99c}, who pointed out that the gravitational influence of a central black hole should have an impact on the distribution of the characteristically hot interstellar medium (ISM). The ISM should gradually flow inwards as it loses energy from the emission of thermal X-rays but, since the local cooling time far exceeds the local free-fall timescale, the flow is highly subsonic and the gas remains close to hydrostatic equilibrium \\citep{mathews03a}. In the powerful gravitational field close to the black hole, the inflowing gas should be sufficiently compressed to cause a strong temperature peak, which could be used as a black hole diagnostic. However, detections of central temperature peaks in galaxies are rare \\citep{humphrey04b,humphrey05a,humphrey06a} and this hypothesis has, therefore, yet to be verified observationally. In part this reflects the difficulty in obtaining precise temperature constraints from X-ray spectroscopy on scales of a few hundred parsecs. AGN-induced disturbances are also common in the centres of early-type galaxies \\citep{jones02a,birzan04a} and these may disturb the gas enough to destroy many such hot cores.   In this paper, we present a detailed \\chandra\\ study of \\src, a nearby, very relaxed and X-ray bright galaxy which is arguably the most promising candidate in which to search for this effect. Based on stellar kinematical studies, \\citet{gebhardt03a} reported a black hole mass of $(2.0^{+0.4}_{-0.6})\\times 10^9$\\msun\\ (1-$\\sigma$ errors), placing it at the upper end of the \\mbh-\\sigmac\\ relation. A modest archival \\chandra\\ dataset has already revealed a remarkably relaxed X-ray morphology \\citep[implying hydrostatic equilibrium is a good approximation:][]{buote95a} and a temperature profile which peaks in the centre \\citep{humphrey06a,randall03}. The object is faint in the radio \\citep{condon02a}, indicating that it does not host a powerful AGN, and there is little evidence of morphological disturbances indicating AGN-ISM interaction (\\S~\\ref{sect_image}). On a large scale, excellent agreement between the mass determined from globular cluster kinematics and X-ray methods \\citep{bridges06a} also provides good support for hydrostatic equilibrium. We present deep new data which enable us to obtain interesting temperature constraints on scales as small as $\\sim$150~pc, which is crucial given the $\\sim$60~pc sphere of influence of the SMBH on the gas\\footnote{Defined as $GM_{BH}/c_s^2$, where $c_s$ is the central gas sound-speed (in the absence of a black hole), $\\sim 490 km\\ s^{-1}$ from our data. This is slightly smaller than the sphere of influence on the stars, $GM_{BH}/\\sigma_*^2$, which is $\\sim$100~pc for the measured $\\sigma_*=385 km\\ s^{-1}$ \\citep{gebhardt03a}.}. For the first time in {\\em any} system we are able to place constraints on \\mbh\\ from hydrostatic X-ray gas.  We adopted a nominal distance of 15.6~Mpc for \\src, based on the I-band SBF distance modulus estimate of \\citet{tonry01}, which was corrected by -0.16 to account for recent revisions to the Cepheid zero-point. At this distance, 1\\arcsec\\ corresponds to 76~pc. Unless otherwise stated, all error-bars represent 90\\%\\ confidence limits. ",
        "conclusions": "\\label{sect_discussion} Using the assumptions of hydrostatic equilibrium, spherical symmetry, a single-phase ISM and that the stellar mass follows the optical light, we have obtained an estimate for the mass of the supermassive black hole at the centre of \\src\\footnote{We note that, technically we also require the entropy profile not to behave pathologically at small radii. However, since hydrostatic equilbrium requires dS/dr $>0$, the most discrepant behaviour we might expect to see in S is a precipitous drop. In such a case, this would tend to {\\em reduce} a central peak in the model temperature profile, making the evidence for a supermassive black hole even stronger.}. This is the first time that the mass of a supermassive black hole has been determined from the distribution of hydrostatic, X-ray emitting gas. In principle, this method can also be applied to other giant elliptical galaxies with relaxed X-ray morphologies, allowing the most massive black holes to be studied.   We placed a lower limit on the black hole mass of $1.6\\times 10^9$\\msun, at 99\\%\\ confidence. Our measured mass agrees (within $\\sim$2-$\\sigma$) with that determined from stellar kinematical studies, $(2.0^{+0.4}_{-0.6})\\times 10^9$\\msun\\ (1-$\\sigma$ error), based on a different set of assumptions. This makes \\src\\ one of only a handful of galaxies for which \\mbh\\ has been measured {\\em via} multiple complementary methods. Although the overall agreement in the shape of the \\mbh-\\sigmac\\ relation when measured using different techniques \\citep[\\eg][]{gebhardt00c} provides good evidence that the masses from stellar kinematics are not {\\em on average} in error, the systematic uncertainties on any individual measurement of \\mbh\\ are more difficult to assess. In the few systems where the mass can be inferred from both gas and stellar dynamics, the results have been mixed, possibly on account of non-circular gas motions (\\S~\\ref{sect_introduction}). Therefore, the good agreement we found between the measurements of \\mbh\\ from stellar dynamics and hydrostatic gas analysis provides crucial, independent verification that stellar dynamics can be used to determine reliable SMBH masses, and effectively rules out pathological orbital structure for the central part of \\src.  Our results provide strong support not only for the stellar kinematical results but also the single-phase, hydrostatic equilibrium approximation we adopted. Firstly, the X-ray image shows a remarkably symmetrical morphology; we do not confirm the evidence of significant asymmetries previously reported with shallower data \\citep{randall03,shurkin07a}. Secondly, the hydrostatic mass model actually fits the data remarkably well, despite the complicated temperature and density profiles. Thirdly, we find good agreement between \\mbh\\ determined {\\em via} both optical and X-ray techniques.  As our derived mass profile was extended to larger radii, we found excellent agreement with stellar mass profile within $\\sim$3~kpc. At even larger radii, the mass determination from the kinematics of globular clusters agreed very well with our results \\citep{bridges06a}. Taken together, these indicate that the ISM \\src\\ must be very close to hydrostatic equilibrium at all scales, or there is a remarkable conspiracy in the data.  In a recent provocative paper, \\citet{diehl07a} analysed the X-ray morphology of a sample of early-type galaxies and concluded that, since the majority show clear evidence of morphological disturbances in their cores, probably all elliptical galaxies are out of hydrostatic equilbrium. Quite apart from the obviously false syllogism that, since there exists disturbed (non-hydrostatic) ISM in many early-type galaxies, {\\em all} early-type galaxies (including those without clear disturbances) are non-hydrostatic, \\src\\ represents a clear counter-example to this argument. This is unsurprising, since  cosmological simulations of galaxy clusters have shown that the existence of asymmetries typically does not translate to mass errors of more than 25\\%  \\citep{tsai94a,evrard96a,nagai07a}, even in systems that are manifestly more disturbed than \\src. Comparing optical and X-ray mass estimates for two systems (NGC\\thin 1399 and M87) that are much more disturbed than \\src, a recent study by \\citet{churazov07a} reaches a similar conclusion. Hence, even when low-level asymmetrical features are eventually detected with higher quality data, as they must, there is no evidence from either theoretical or observational studies that this implies a systematic error in the derived mass much larger than the other systematic errors we addressed in this paper.     With \\chandra\\ we were unable to resolve the sphere of influence ($\\sim$60~pc) of the SMBH on the gas although, since the emissivity is proportional to the square of the gas density (which rises steeply towards the centre), we were able to probe the gas density within a factor $\\sim$2--3 of it. Provided hydrostatic equilibrium holds, resolving the sphere of influence is not strictly necessary to place constraints on the central mass if one takes care to compute reliable average densities and temperatures within each radial bin. Ideally, however, X-ray imaging spectroscopy at scales smaller than this scale is desirable so that the SMBH mass is not determined primarily by a single data-point. With such data one could also search for morphological disturbances at the smallest scales, which would indicate local deviations from hydrostatic equilbrium. Unfortunately, such imaging is beyond the capabilities of \\chandra\\ and will have to wait until the era of higher spatial resolution X-ray instruments. Nonetheless, within the central $\\sim$200~pc, the dynamical time-scale is only $8\\times 10^5$~yr, as compared to a cooling time $\\sim 2\\times 10^7$~yr so that, unless the gas is being constantly stirred up, we would expect it to relax quickly into a state approximating hydrostatic equilibrium. Although ISM disturbances due to AGN activity appear common in early-type galaxies \\citep[\\eg][]{jones02a}, these are typically seen on large scales, comparable to any radio lobes \\citep[which, although very weak in \\src, are on $\\sim$kpc scales:][]{stanger86a}. Given that the dynamical time-scale is shortest in the central parts of the galaxy, \\src\\ would have to be in a very unusual state for AGN activity to stir up the gas within only the central $\\sim$100~pc but leave the X-ray morphology outside this range undisturbed. Although the current X-ray data do not allow us completely to rule this out, turbulence or bulk motion would provide additional support against gravity \\citep[\\eg][]{nagai07a}, lowering the central temperature for a given mass and entropy, and leading to an underestimate of \\mbh. Since a deviation from this assumption would lead to poorer agreement with the optical data, our results indicate that the gas is very close to hydrostatic.    In addition to being close to hydrostatic equilibrium, our results indicate that the gas must also be close to single phase. We tested this explicitly by adding an extra {\\bf vapec} spectral component in each annulus. We found that the fit did not improve significantly ($\\Delta \\chi^2$=9 for 30 fewer degrees of freedom) and, excepting the innermost bin, the extra component was negligible in each annulus, supporting the single-phase assumption. Although we cannot rule out multi-phase gas within the innermost $\\sim$200~pc, it is difficult to understand why genuinely multi-phase gas would be confined only to this region. Fitting a single temperature model to data where multi-temperature models are actually required in general should lead us to underestimate the mass \\citep[\\eg][]{allen98}, so, if anything, the effect of multi-phase gas would be to make our measured \\mbh\\ more discrepant with the optical data (implying the gas is probably single-phase). It is worth noting that our models do indicate a steeply-rising temperature profile within the innermost bin and in such cases, two-temperature spectral models are often statistically better descriptions of the data than single-temperature models \\citep[\\eg][]{buote00c}. However, one can still obtain reliable mass profiles by interpreting a single-temperature fit to the data provided one takes care to compare its parameters to a suitably averaged model, as we did in the present work \\citep{gastaldello07a}.   One of the intriguing results from our analysis is an explanation of the central temperature spike seen  in this system. As we show in Fig~\\ref{fig_temp_profile}, if we omit the central black hole, our model temperature profile is much less centrally-peaked than the real data. This confirms the predictions of \\citet{brighenti99c}, who suggested that such features should be prevalent when \\chandra\\ temperature profiles of elliptical galaxies hosting massive black holes are investigated. This detection, therefore, opens the possibility of using the same technique to measure the central black hole mass in other, morphologically relaxed, early-type galaxies. Since significant X-ray halos are most often associated with the most massive galaxies, this will make such a method particularly sensitive to the most massive black holes. However, even for a black hole as massive as $\\sim 3\\times 10^9$\\msun, the sphere of influence is only $\\sim$50~pc, requiring precise temperature constraints on a comparable scale. Within the limitations of current instrumentation, this effectively restricts the analysis to very X-ray bright galaxies within $\\sim$20~Mpc, severely limiting the number of systems in which we might hope to see it. With the  improved spatial resolution promised by proposed future  X-ray satellites,  such as the \\genx\\ mission \\citep{windhorst06a}, though, such measurements should become routine.     In practice, the presence of a central temperature spike is not actually a sufficient diagnostic for a massive central black hole. As is clear from Fig~\\ref{fig_temp_profile}, a mild central increase in the temperature can arise as a consequence of the centrally-peaked stellar mass distribution. In fact, if the stellar mass was significantly more cuspy (for example in a ``powerlaw'' galaxy, rather than the ``cored'' \\src), we would have seen a significant central temperature rise even without a black hole.  More importantly, the mass profile alone does not determine the shape of the temperature profile, but instead it is the interplay of the mass and entropy profiles. The shape of the latter is set by the balance between heating and cooling, which is not completely understood. In general, the flatter the central entropy profile, the more centrally-peaked is the temperature profile, so quantitative differences in the thermal histories of individual galaxies would explain in part the paucity of such temperature spikes. In many objects, this is compounded by the presence of AGN-driven disturbances \\citep[\\eg][]{birzan04a,osullivan05a}, which tend to stir up the gas, disturbing hydrostatic equilibrium and mixing hot, central gas with cooler gas from larger radii. Combined with the observational difficulty of the measurement, these effects may explain the few reported galaxies exhibiting this feature \\citep[\\eg][]{humphrey06a,humphrey04b}, in contrast to the ubiquity of black holes.  In addition to the black hole mass estimate, our method also allowed us to determine directly the total stellar mass of the system, without making any assumptions regarding the stellar mass-to-light ratio, and requiring only that stellar mass follows light. We therefore obtained {\\em both} data-points (albeit not completely independently) to place on a \\mbh\\ {\\em  versus} stellar mass relation. We estimated that the mass of the SMBH is $\\sim$1.7\\%\\ times the mass in stars of the galaxy. It is interesting to compare this to the relations between \\mbh\\ and bulge mass (\\mbul) found by \\citet{marconi03a}. For a bulge mass of $\\sim 2\\times 10^{11}$\\msun, those authors would predict a ratio of \\mbh\\ to \\mbul\\ of $\\sim$0.2\\%, which is discrepant at almost an order of magnitude from our measurement (larger than the expected scatter). \\citeauthor{marconi03a} estimated \\mbul\\ by inserting the stellar velocity dispersion and effective radius into the Virial relation, which actually gives an estimate of the {\\em total} gravitating matter within, roughly, the effective radius. For the case of \\src, which they include in their sample, they adopted an effective radius of 7.5~kpc (corrected to our adopted distance) and derived \\mbul $\\simeq7.8\\times 10^{11}$\\msun. This is considerably higher than the {\\em stellar} mass we measured within the galaxy and, not unexpectedly, falls far short of the $\\sim3\\times 10^{13}$\\msun\\ Virial mass of the entire system \\citep{humphrey06a}. Instead, this is approximately the total gravitating mass within $\\sim$19~kpc, of which dark matter comprises almost $\\sim$75\\%. Clearly this illustrates the difficulty in interpreting Virial theorem estimates for \\mbul, especially at the most massive end where galaxies sit in the centres of group-like dark matter halos. In order to gain insight into the SMBH-bulge relation, a full decomposition of the mass profile into stellar and dark components is, therefore, essential.  Finally, based on the mean stellar population age and abundance estimated from our Lick index analysis \\citep{humphrey06a}, we can estimate the expected mass-to-light (M/\\lj) ratio for different simple stellar population models, assuming all stars were born in a swift burst of star-formation. Using the models of \\citet{maraston05a}, the J-band ratio is $1.86\\pm0.21$\\mpl, assuming a \\citet{kroupa01a} IMF, and depending only very weakly on the metallicity of the stars. As an alternative, if we adopt version 2 of the PEGASE code \\citep{fioc97a}, the predicted J-band M/L ratio is $1.55\\pm0.16$\\mpl, indicating the level of systematic uncertainty in such population models. These results compare favourably with our observed stellar M/\\lj\\ ratio ($1.37\\pm 0.10$\\mpl), although the \\citeauthor{maraston05a} models are slightly discrepant. We note that our measured M/\\lj\\ is completely inconsistent with the predictions for a Salpeter IMF ($2.7\\pm0.3$\\mpl\\ for both sets of models), as we found in \\citet{humphrey06a}, unless the age for the stellar population is dramatically over-estimated. In that case it would need to be closer to $\\sim$7~Gyr than the 13$\\pm$2~Gyr we measured."
    },
    "0801/0801.1008_arXiv.txt": {
        "abstract": "The Warm-Hot Intergalactic Medium (WHIM) is thought to contribute about $40-50$~\\% to the baryonic budget at the present evolution stage of the universe. The observed large scale structure is likely to be due to gravitational growth of density fluctuations in the post-inflation era. The evolving cosmic web is governed by non-linear gravitational growth of the initially weak density fluctuations in the dark energy dominated cosmology. Non-linear structure formation, accretion and merging processes, star forming and AGN activity produce gas shocks in the WHIM. Shock waves are converting a fraction of the gravitation power to thermal and non-thermal emission of baryonic/leptonic matter. They provide the most likely way to power the luminous matter in the WHIM. The plasma shocks in the WHIM are expected to be collisionless. Collisionless shocks produce a highly non-equilibrium state with anisotropic temperatures and a large differences in ion and electron temperatures. We discuss the ion and electron heating by the collisionless shocks and then review the plasma processes responsible for  the Coulomb equilibration and collisional ionisation equilibrium of oxygen ions in the WHIM. MHD-turbulence produced by the strong collisionless shocks could provide a sizeable non-thermal contribution to the observed Doppler parameter of the UV line spectra of the WHIM. ",
        "introduction": "\\label{Introduction} Cosmological simulations of the large-scale structure (LSS) predict that about $40-50$~\\% of baryons at epoch $z < 2$ could reside in the Warm-Hot Intergalactic Medium (WHIM) with temperatures $10^5-10^7$ K at moderate overdensities $\\delta < 100$ \\citep{CenO99,Dave_ea01,Fang_ea02}. The WHIM heating is due to shocks driven by gravitationally accelerated flows in the LSS structure formation scenario (e.g. \\citealt{Kang_ea07}). Numerical simulations predict the observational signatures of the web gas as a function of redshift. The simulations account for feedback interactions between galaxies and the intergalactic medium, and demonstrate that the X-ray and ultraviolet \\ion{O}{vi}, \\ion{O}{vii} and \\ion{O}{viii} lines and the \\ion{H}{i} Lyman alpha line are good tracers of low-density cosmic web filamentary structures (e.g. \\citealt{Hellsten_ea98,Tripp_ea00,Cen_ea01,Furlanetto_ea04}). Intervening metal absorption systems of highly ionised C, N, O, Ne in the soft X-ray spectra of bright Active Galactic Nuclei (AGN) were suggested to be tracer of the WHIM. The predicted distribution of ion column densities in the WHIM absorbers is steep enough to provide  only a few systems with $N_{\\mbox{O\\,{\\scriptsize VII}}} > 10^{15}$~cm$^{-2}$  along an arbitrary chosen line of sight (e.g. \\citealt{Fang_ea02}). Therefore, the detection of the WHIM is particularly difficult and requires very sensitive UV and X-ray detectors, both for absorption and for emission processes (e.g. \\citealt{Lehner_ea07,Nicastro_ea05,Kaastra_ea06,Takei_ea07} and \\citealt{richter2008,durret2008} - Chapters 3 and 4, this volume). Future projects and namely {\\sl Cosmic Origin Spectrograph} ({\\sl COS}), the {\\sl X-Ray Evolving Universe Spectrometer} ({\\sl XEUS}), {\\sl Constellation-X} and the {\\sl Diffuse Intergalactic Oxygen Surveyor} ({\\sl DIOS}) will increase the signal-to-noise ratio in the spectra allowing to study weak systems with $N_{\\mbox{H\\,{\\scriptsize I}}} < 10^{12.5}$~cm$^{-2}$ and $N_{\\mbox{O\\,{\\scriptsize VII}}} < 10^{15}$~cm$^{-2}$. Simulations of spectra of the broad Ly$\\alpha$ absorption lines and the highly ionised oxygen lines in the weak systems require thorough modelling of physical condition in the plasma (see e.g. \\citealt{Mewe90,PaerelsK03,Kawahara_ea06}). We discuss in this paper the heating and equilibration processes in the shocked WHIM plasma affecting the spectral simulations. A discussion of collisionless shock physics relevant to cosmological shocks in the WHIM can be found in \\citealt{bykov2008} - Chapter 7, this volume. To this end, in this paper we discuss first some specific features of collisionless shock heating of ions of different charge states providing highly non-equilibrium initial states just behind the magnetic ramp region that relaxes to an equilibrium state through Coulomb collisions, and the relation of the processes to simulations of emission/absorption spectra of the WHIM and observational data analysis. ",
        "conclusions": "We discussed some specific features of the WHIM heating processes by collisionless plasma shocks driven by gravitationally accelerated flows in the LSS structure formation scenario.  $\\bullet$ Collisionless shock heating of ions in the WHIM results in a highly non-equilibrium initial state with a strongly anisotropic quasi-Maxwellian ion distribution just behind a very thin magnetic ramp region. The ion temperatures  just behind the shock depend on the ion atomic weight, the charge state and the shock magnetic field inclination. The ion temperatures could decline from a simple linear scaling with the ion mass providing an excessive heating of heavy ions.  $\\bullet$ The ion and electron temperatures relax to the equilibrium state through Coulomb collisions in a layer of the depth $N_{\\rm eq} \\sim 10^{16} v_{\\rm 7}^4$~cm$^{-2}$ behind a shock of velocity $v_{\\rm 7}$.   $\\bullet$ Strong collisionless plasma shocks with an efficient Fermi acceleration of energetic particles can generate strong MHD waves in the downstream region that will result in a non-thermal broadening of the emission/absorption lines in the observed WHIM spectra."
    },
    "0801/0801.1604_arXiv.txt": {
        "abstract": "Neutrino telescopes are moving steadily toward the goal of detecting astrophysical neutrinos from the most powerful galactic and extragalactic sources. Here we describe analysis methods to search for high energy point-like neutrino sources using detectors deep in the ice or sea. We simulate an ideal cubic kilometer detector based on real world performance of existing detectors such as AMANDA, IceCube, and ANTARES. An unbinned likelihood ratio method is applied, making use of the point spread function and energy distribution of simulated neutrino signal events to separate them from the background of atmospheric neutrinos produced by cosmic ray showers.  The unbinned point source analyses are shown to perform better than binned searches and, depending on the source spectral index, the use of energy information is shown to improve discovery potential by almost a factor of two. ",
        "introduction": "\\label{Sec1}  With the construction of IceCube at the South Pole \\cite{IceCube} and of ANTARES in the Mediterranean Sea \\cite{ANTARES}, together with existing R\\&D programs for a cubic kilometer array at these latitudes \\cite{km3net,NEMO}, neutrino astronomy is entering a very promising era. IceCube, when complete in 2011, will consist of up to 80 strings on a hexagonal grid with 124 m spacing, with each string holding 60 optical modules vertically spaced by 17 m. The instrumented part of the strings is deployed between 1450 and 2450 m deep in the ice. The experiment profits from experience acquired with the AMANDA detector, taking data since 1996 and completed in 2000 with 19 strings containing 677 total optical modules between 1500 and 2000 m below the ice surface. Studies on the optimal configuration for a cubic kilometer array in the Mediterranean are ongoing, and the amount of photomultipliers (PMTs) generally considered is somewhat larger than for IceCube, around 6000 \\cite{distefano} and up to about 9000 \\cite{Carr}.  The community is refining methods to detect low statistics signals amongst large backgrounds. The expected background from atmospheric neutrinos in a cubic kilometer detector is of the order of 50 000 upgoing events per year after selection criteria guaranteeing good angular resolution and rejection of misreconstructed cosmic ray muons.  Point-like signals of few events need to be singled out among this large number of background events. Two features distinguish signal from the background: \\begin{itemize} \\item The angular distribution. The signal would cluster around the direction of the neutrino source (assumed here to be point-like) with a spread depending on the detector angular resolution. Angular resolution is limited by detector geometry and by the propagation characteristics of light in the medium, specifically by photon scattering and absorption. The pointing accuracy for astrophysical sources is also limited by the kinematic angle between the parent neutrino and the muon. \\smallskip \\item The energy distribution. The differential energy spectrum of the signal expected from Fermi acceleration mechanisms is close to E$^{-2}$, harder than that of atmospheric neutrinos due to the showering process in the atmosphere. The differential spectrum of atmospheric neutrinos approximately follows a power law of E$^{-3.7}$ above 100 GeV. \\end{itemize} Other signatures may be used, including time dependencies of emissions measured in other detectors such as correlations with gamma ray bursts or TeV gamma ray flares.  However, we focus on steady emissions of neutrinos with time. The methods we have implemented exploit the two features listed above. We show that unbinned methods based on the likelihood ratio hypothesis test perform better than methods based on angular bins, and we show that the introduction of energy dependent information, e.g. the number of hit PMTs, helps discriminate signal and allows an energy spectrum reconstruction even when few events are detected on top of the background. Other unbinned methods have been studied and developed by other authors \\cite{till2,aart,finley,aguilar}.  Sec.~\\ref{Sec2} describes the simulation we use to generate realistic samples of signal and background events in a cubic kilometer detector. Sec.~\\ref{Sec3} describes the unbinned method based on a likelihood ratio analysis, comparing a signal plus background hypothesis to a background only hypothesis. This method has been applied to IceCube data for 9 strings and to 2005-6 data of the AMANDA-II detector ~\\cite{chadicrc07,jimicrc07}. In Sec.~\\ref{Sec4} we describe the performance of the methods and show results in terms of discovery potential. We also emphasize the importance of using energy related information to increase the discovery potential and show the ability to determine the spectral index of the neutrino source. The method is then compared to a search using angular bins, more traditionally applied in neutrino astronomy \\cite{macro,amanda5,SuperKamiokande}. ",
        "conclusions": "\\label{Sec5}  We have shown that likelihood methods improve discovery potential using differences in both the angular distribution of events and energy spectra between the background and signal hypotheses.  Furthermore, it has been shown that 10\\% less flux is required compared to methods using predefined angular search bins. The improvement is more significant when an energy related variable is used to distinguish hard astrophysical neutrino spectra from the softer spectrum of atmospheric neutrinos. In this case, for E$^{-2}$ energy spectra only about half of the flux is required to achieve 5$\\sigma$ compared to binned methods."
    },
    "0801/0801.1118_arXiv.txt": {
        "abstract": "{Asymptotic Giant Branch (AGB) stars are typified by strong dust-driven, molecular outflows. For long, it was believed that the molecular setup of the circumstellar envelope of AGB stars is primarily determined by the atmospheric C/O ratio. However, recent observations of molecules such as HCN, SiO, and SO reveal gas-phase abundances higher than predicted by thermodynamic equilibrium (TE) models. UV-photon initiated dissociation in the outer envelope or non-equilibrium formation by the effect of shocks in the inner envelope may be the origin of the anomolous abundances.} {We aim at detecting \\emph{(i)} a group of `parent' molecules (CO, SiO, HCN, CS), predicted by the non-equilibrium study of Cherchneff (2006) to form with almost constant abundances independent of the C/O ratio and the stellar evolutionary stage on the Asymptotic Giant Branch (AGB), and \\emph{(ii)} few molecules, such as SiS and SO, which are sensitive to the O- or C-rich nature of the star.} {Several low and high excitation rotational transitions of key molecules are observed at mm and sub-mm wavelengths with JCMT and APEX in four AGB stars: the oxygen-rich Mira \\object{WX Psc}, the S star \\object{W Aql}, and the two carbon stars \\object{V Cyg} and \\object{II~Lup}. A critical density analysis is performed to determine the formation region of the high-excitation molecular lines.} {We detect the four `parent' molecules in all four objects, implying that, indeed, these chemical species form whatever the stage of evolution on the AGB. High-excitation lines of SiS are also detected in three stars with APEX, whereas SO is only detected in the oxygen-rich star \\object{WX Psc}. } { This is the first multi-molecular observational proof that periodically shocked layers above the photosphere of AGB stars show some chemical homogeneity, whatever the photospheric C/O ratio and stage of evolution of the star. } ",
        "introduction": "\\label{Introduction} Circumstellar envelopes of Asymptotic Giant Branch stars (AGBs) have long been known to be efficient sites of molecule formation.  While the outer layers of such envelopes experience penetration of interstellar UV photons and cosmic rays resulting in a fast ion-molecule chemistry, the deepest layers are dominated by a non-equilibrium chemistry due to the passage of shocks generated by stellar pulsation. Dust forms in those inner gas layers, still bound to the star, and grains couple to the gas to accelerate it, thereby generating stellar wind and mass loss phenomena. The described processes greatly modify the abundances established by the equilibrium chemistry in the dense, hot photosphere \\citep{Tsuji1973A&A....23..411T}.  For a long time, the gas chemical composition was believed to be dominated entirely by the C/O ratio of the photosphere. A C/O ratio greater than one implied that all the oxygen was tied in CO, leading to an oxygen-free chemistry, whereas a C/O ratio less than one meant that no carbon bearing molecules apart from CO could ever form in an oxygen-rich (O-rich) environment. This picture, based essentially on thermal equilibrium considerations applied to the gas, has been first disproved by the detection of SiO at millimeter (mm) wavelength in carbon-rich (C-rich) AGBs \\citep{Bujarrabal1994A&A...285..247B}. As for O-rich AGBs, CO$_2$ infrared (IR) transition lines were detected in various objects with the Short-Wavelength Spectrometer (SWS) onboard the Infrared Space Observatory (ISO) \\citep[e.g.,][]{Justtaonont1996A&A...315L.217J, ryde1998Ap&SS.255..301R}.  Theoretical modeling describing the chemistry in the inner wind of the extreme carbon star \\object{IRC+10216} showed that the formation of SiO was due primarily to hydroxyl OH reaction with atomic silicon close to the photosphere as a result of shock activity and therefore non-equilibrium chemistry \\citep{Willacy1998A&A...330..676W}.  Later on, \\citet{Duari1999A&A...341L..47D} showed that CO$_2$ formation in the O-rich Mira \\object{IK Tau} results from the reaction of OH radicals with CO in the shocked regions, implying again that non-equilibirum chemistry was paramount to the formation of C-bearing species in O-rich Miras. It was then recently proposed that the inner wind of AGBs shows a striking homogeneity in chemical composition, despite their photospheric C/O ratio and stage of stellar evolution \\citep{Cherchneff2006A&A...456.1001C}. In particular, \\citet{Cherchneff2006A&A...456.1001C} showed that when taking shock chemistry into account, molecules such as SiO, HCN and CS are present in comparable amount in the inner layers of M, S, and C AGBs, whereas specific molecules (e.g.\\ SO and HS for O-rich Miras and C$_2$H$_2$ for carbon stars) are typical for O-rich or C-rich chemistries.  In this letter, we present observations carried out with the JCMT and the APEX telescope of four AGBs: one O-rich, \\object{WX Psc}, one S star ($\\equiv$ C/O $\\approx$ 1), \\object{W Aql}, and two carbon stars, \\object{II Lup} and \\object{V Cyg}. We focus on the detection of (sub)mm transitions of CO, SiO, HCN, CS, SiS and SO in order to confirm or disprove the above hypothesis and to check for homogeneity in AGB winds. ",
        "conclusions": "  \\vspace{-1.5ex} \\begin{enumerate} \\item The observations reported in this letter confirm the status of `parent' molecules for CO, SiO, HCN, and CS in AGB stars, whose observed molecular lines form close to the star in support of the theoretical predictions of \\citet{Cherchneff2006A&A...456.1001C}. The excitation analysis suggests that SiO is emitted closest to the star, followed by HCN and CS. \\item High-excitation lines of SiS are detected in all stars, implying that SiS too forms close to the star, whatever the stage of stellar evolution. However, a thorough line analysis is necessary to prove or disprove that SiS is more abundant in C stars than in O-rich Miras, as predicted by \\citet{Cherchneff2006A&A...456.1001C}. \\item SO appears to be typical of O-rich AGBs only. With only two excitation lines detected, no information can be drawn on the locus of SO formation. However, the SO($10_{11}$-$10_{10}$) line has an estimated velocity of only 10\\,km/s in \\object{WX Psc} when the wind terminal velocity for this object is 19.2\\,km/s. This fact suggests that the line excitation occurs close to star. In any case, the inner chemistry of O-rich AGB envelopes appears to be as rich, if not richer, than that of C-rich stars.  \\end{enumerate}  A detailed line analysis of the present data coupled to further observational campaigns with JCMT and APEX are planned to corroborate these results."
    },
    "0801/0801.3092_arXiv.txt": {
        "abstract": "{The  development of parallel  supercomputers allows today the  detailed study  of the  collapse and  the  fragmentation of prestellar   cores    with   increasingly   accurate   numerical simulations.    Thanks  to   the   advances  in   sub-millimeter observations, a wide range of observed initial conditions enable us to study the different  modes of low-mass star formation. The challenge for the simulations  is to reproduce the observational results.}  {Two main numerical  methods, namely AMR and SPH, are widely used  to simulate the  collapse and the  fragmentation of prestellar cores. We compare thoroughly these two methods within their standard framework.}  {We use  the AMR code RAMSES and the SPH code DRAGON.   Our physical model is as  simple as possible, and  consists  of  an  isothermal  sphere  rotating  around  the $z$-axis.  We  first study the conservation  of angular momentum as a function of the  resolution.  Then, we explore a wide range of   simulation  parameters  to   study  the   fragmentation  of prestellar cores.}  {There seems to be a convergence between the two  methods, provided  resolution in  each case  is sufficient. Resolution criteria  adapted to our physical cases,  in terms of resolution per  Jeans mass, for  an accurate description  of the formation  of protostellar  cores are  deduced from  the present study.   This  convergence is  encouraging  for  future work  in simulations   of   low-mass   star  formation,   providing   the aforementioned  criteria  are  fulfilled.}  {} ",
        "introduction": "Star formation is  known for being the place  of extreme variations in length and density scales. Although  it is established that stars form in dense cores, the non-linear evolution makes it difficult to perform accurate  calculations of  the  collapse and  the  fragmentation of  a prestellar core. The star formation  process is the outcome of complex gas dynamics involving non-linear interactions of gravity, turbulence, magnetic  field and  radiation.   Early theoretical  pioneer works  by \\cite{Larson_1969},  \\cite{Penston_1969} or \\cite{Shu_1977}  are among the  many illustrations of  the high  complexity of  the gravitational collapse.  Recently,  \\cite{Klein_2007} pointed out  that developing a theory for low-mass star formation remains one of the most elusive and important  goals  of   theoretical  astrophysics.   The  computational challenge stems  from the  fact that star  formation occurs  in clouds over  many  orders  of   magnitude  in  spatial  and  density  scales. Following  the gravitational  collapse while  resolving  precisely the Jeans length, which scales as $\\lambda_\\mathrm{J} \\propto \\rho^{-1/2}$ for  an   isothermal  gas,  is   a  major  difficulty   for  numerical simulations.  Different  approaches  are  used   to  study  star  formation  through numerical simulations and include  more and more detailed physics. One key question resides  in the validation of the  numerical methods used to study low-mass star  formation.  Nowadays, two completely different numerical method are used with sufficient accuracy:  \\begin{enumerate} \\item AMR: Adaptive Mesh Refinement method for Eulerian grids \\item SPH:  Smoothed Particle Hydrodynamics method  for a Lagrangian approach. \\end{enumerate}  No systematic  comparison between the  two methods has been  done with low-mass  star formation  calculations.  However,  a lot  of numerical works  have  been  carried  out  and  some of  them  are  common  test calculations for  convergence testing and  intercode comparisons.  The most famous  model was first calculated by  \\cite{Boss_1979} and since then, it  has been  recalculated by several  authors with  even higher spatial           resolution           \\cite[e.g][]{bate_burkert-1997, Truelove_1998,Kitsionas_2002,Arreaga_2007}.   The   SPH  approach  has generated a  lot of  detailed investigations on  the influence  of the number          of           particles          and          neighbors \\citep{Lombardi_1999,rasio-1999,Attwood_2007},    and   criteria   for numerical  convergence   have  been  extracted   from  these  studies. \\cite{Nelson_2006} performed a large investigation of the influence of these parameters on disk  fragmentation, and concluded that the better the  resolution  the   later  the  fragmentation.   \\cite{Dehnen_2001} investigated  the optimal gravitational  force softening  necessary in three-dimensional $N$-body  codes.  \\cite{bate_burkert-1997} provide a minimum resolution criterion for SPH calculations with self-gravity to accurately model fragmentation.  Less studies have been performed with AMR    since    AMR   codes    have    become   available    recently. \\cite{Truelove_1997} give an empirical  criterion for the Jeans length resolution   in   AMR  calculations   to   avoid  spurious   numerical fragmentation.  There   are  not  much of    direct  comparison  between SPH  and  AMR calculations. Comparison  in  the context of  cosmological simulations has   been done through the Santa   Barbara Cluster Comparison Project \\citep{Santa_Barbara_1999}.        \\cite{Fromang_2006}  compares quite successfully AMR hydrodynamical collapse calculations with the ones of \\cite{Hosking_2004}, using the SPH method.\\\\  In the present paper, we  compare thoroughly the two approaches in the context  of low-mass prestellar  core formation.   We stress  that the main goal of  this paper is to investigate  whether convergence can be achieved between the two methods.  We have conducted calculations over a wide range of numerical resolution parameters, in order to study the dependency  of  angular  momentum  conservation and  fragmentation  on physical and numerical initial  conditions.  We then derive resolution criteria necessary to describe accurately prestellar core formation.    The paper  is organized  as follows: in  \\S2 we briefly  introduce our collapse  model.   In \\S3,  we  present the  two  codes  used for  our comparative study as well as  our initial numerical conditions and the criteria we fulfill to resolve gravitational collapse.  The problem of angular momentum conservation  is examined in detail \\S4.   In \\S5, we tackle  the fragmentation  issue  and explore  the  dependency of  the results  on the numerical  parameters. First,  we study  the numerical convergence of AMR and  SPH calculations separately.  Then, we compare the respective converged calculations.  This convergence study is done for  different test  cases. In  section 6  we conclude  this  paper by deriving for each method  empirical required numerical criteria for an accurate  description  of  gravitational  collapse  and  fragmentation adapted to our test cases.\\\\  The convention in this paper  is to call \"particles\" the SPH particles and \"cells\" the AMR cells in order to avoid confusion. ",
        "conclusions": "We have investigated the effect of numerical resolution in AMR and SPH calculations on the collapse and the fragmentation of rotating cores.  We show  that we  reach good convergence  between AMR and  SPH methods provided one  uses sufficient numerical  resources.  First, we  take a simple model to study local angular momentum conservation. The initial study shows that  local angular momentum is better  conserved with the AMR  approach for  equivalent computational  needs, whereas  SPH gives better dynamical  times.  As shown in Fig.   \\ref{homorot}a, a smaller number of  particles in standard  SPH calculations leads to  bad local angular  momentum conservation.  Numerical  torques on  the rotational axis are  accumulated for denser  particles, whereas our  model should remain   axisymmetric.    In   AMR   calculations,  a   poor   initial computational domain resolution  (i.e.  $\\ell_\\mathrm{min} < 6$) leads to unphysical  transfer of  gravitational energy to  rotational energy (see Fig.   \\ref{homorot}b).  A  significant loss of  angular momentum will affect  fragmentation since  less rotational support  can balance gravitational   collapse.   The   smallest  parameter   set   for  SPH calculations  required to  go through  gravitational  collapse without significant loss of angular momentum  corresponds to a number of $\\sim 530$   particles   per   Jeans    mass   at   the   critical   density $\\rho_\\mathrm{c}$,  i.e.  5  particles  per Jeans  {\\it length}.   The equivalent  minimum  resolution  criterion  for  AMR  calculations  is $\\ell_\\mathrm{min} \\ga 6$ and $N_\\mathrm{J}=4$.  Then we  investigate fragmentation issues for  three different initial condition.   For  the least  prone  to  fragment case  ($\\alpha=0.65$, $\\beta=0.04$), we show  that AMR and SPH methods  give similar results when  the respective  Jeans resolution  criteria are  fulfilled. These results  agree  with  the  semi-analytical criteria  on  fragmentation derived  by \\cite{Tsuribe_Inutsuka_1999}  for the  isothermal collapse ($\\alpha  \\la 0.55  - 0.65\\,  \\beta$). We  study extensively  the case $\\alpha=0.5$, $\\beta=0.04$.  We first reach good agreement between AMR calculations    with   the   parameters    $\\ell_\\mathrm{min}=6$   and $N_\\mathrm{J}=15$  and  SPH  calculations  with  $N_\\mathrm{p}=5\\times 10^5$  and $N_\\mathrm{N}=50$,  i.e.  $\\sim  5370$ particles  per Jeans mass at critical density $\\rho_\\mathrm{c}$.  These parameter sets seem to  be  a   lower  resolution  limit  for  dense   core  collapse  and fragmentation SPH and AMR calculations  in order to get good agreement in both time  and space scales.  Using a lower  number of particles or number  of points  per  Jeans  length will  lead  to inaccurate  early fragmentation due to numerical effects.  Initially, we compare the two converged  calculations and,  for  this specific  case,  we find  good agreement  between the  two  methods (see  Fig \\ref{comp_a050}).   The price to pay in computer time,  however, is larger with the SPH method for the fragmentation study, due  to not using sink particles.  In the case  of  low  thermal  support,  the  dynamic  quickly  becomes  very non-linear  and numerical  convergence of  the simulation  can  not be achieved  as easily  as  for higher  thermal  support.  A  statistical analysis over a large number of  simulations would be needed to see if converging results  can be obtained in term  of fragment distributions (in mass, size ...).  The   two   approaches   show   good   agreement   for   the   general pictures. Details  are better resolved  in AMR calculations  thanks to the refinement  method based  on the local  Jeans length,  whereas the resolution  deteriorates with  increasing density  with  standard SPH. Numerical  calculations  of   protostellar  collapse  should  thus  be conducted with  great care, with  a detailed examination  of numerical resolution. The  present work  can be used  to assess the  validity of numerical tools to study star formation."
    },
    "0801/0801.4167_arXiv.txt": {
        "abstract": "New solar wind data from the \\textit{Voyager}~1 and \\textit{Voyager}~2 spacecraft, together with the \\textit{SOHO} SWAN measurements of the direction that neutral hydrogen enters into the inner heliosheath and neutral helium measurements provided by multiple observations are expected to provide more reliable constraints on the ionization ratio of the local interstellar medium (LISM) and the direction and magnitude of the interstellar magnetic field (ISMF). In this paper we use currently the most sophisticated numerical model of the heliospheric interface, which is based on an MHD treatment of the ion flow and kinetic modeling of neutral particles, to analyze an ISMF-induced asymmetry of the heliosphere in the presence of the interplanetary magnetic field and neutral particles. It is shown that secondary hydrogen atoms modify the LISM properties leading to its shock-free deceleration at the heliopause. We determine the deflection of hydrogen atoms from their original trajectory in the unperturbed LISM and show that it occurs not only in the plane defined by the ISMF and LISM velocity vectors, but also, to a lesser extent, perpendicular to this plane. We also consider the possibility of using 2--3 kHz radio emission data to further constrain the ISMF direction. ",
        "introduction": "The interaction of the solar wind (SW) with the local interstellar medium (LISM) determines the structure of the outer heliosphere. The mathematical problem of this interaction should articulate a physical model that takes into account the fundamental processes occurring in the SW plasma when it collides with a weakly ionized LISM. The model should be accompanied by boundary conditions at large distances from the Sun, in the unperturbed LISM, and at 1~AU, where the SW properties are monitored by a fleet of spacecraft, such as \\textit{ACE}, \\textit{WIND}, \\textit{Ulysses}, \\textit{STEREO}, etc. Ultimately, the solution to the problem of providing the boundary conditions on an inner surface must be based on the use of solar observations (see, e.~g., Wang \\& Sheeley~2006). The properties of the partially ionized LISM generally cannot be measured \\emph{in situ}. For this reason, we are obliged to rely on numerical modeling, which is capable of providing parameter-dependent results for the quantities measured in the inner heliosheath.  The direction and magnitude of the LISM velocity $\\mathbf{V}_\\infty$ are determined from the observations of He atoms (see M\\\"obius et al.~2004 and references therein), which can propagate easily into the SW region within 1~AU from the Sun. The temperature of H ions and atoms in the LISM is usually assumed to be the same as that of neutral He. The number densities of the LISM H atoms ($n_{\\mathrm{H}\\infty}$) and ions ($n_\\infty$), and even the ionization ratio of the hydrogen plasma, are not well know. Constraints on the densities can be derived from the ratio of neutral H and He line densities, which is 14.7 \\citep{Dupuis,Wolff}, and the abundance ratio between H and He as chemical elements in the LISM, which is 10 \\citep{Vallerga96}. Several MHD calculations that include neutral particles self-consistently \\citep{Izmod05,Jacob06,Pogo07} use $n_\\mathrm{H\\infty}$ in the range 0.15 to 0.2~cm$^{-3}$, which yields a neutral H density of about 0.1~cm$^{-3}$ at the SW termination shock (TS) -- a value consistent with the Ulysses pick-up ion (PUI) measurements \\citep{Gloeckler}.  The strength and direction of the interstellar magnetic field (ISMF) are not determined with sufficient accuracy. According to starlight polarization measurements in the solar neighborhood \\citep{Tinbergen}, the ISMF vector $\\mathbf{B}_\\infty$ belongs to the Galactic plane. The theory for the origin of the Local Bubble suggested by \\citet{Cox03} requires strong magnetic fields of $4$--$5\\,\\mu \\mathrm{G}$ nearly parallel to the LISM velocity. Recent measurements performed by the Solar Wind ANisotropy (SWAN) experiment on the \\textit{SOHO} spacecraft determined the direction that neutral hydrogen enters the inner heliosheath \\citep{Lallement05}. This direction and the initial direction of the LISM velocity (the angle between these is $4^\\circ \\pm 1^\\circ$) define a so-called hydrogen deflection plane (HDP). The deflection of neutral H flow occurs through charge exchange in the regions where there is a difference between the velocities of H atoms and ions. Since it is assumed that the velocities of H atoms and He atoms coincide with that of H ions and because the proton density in the SW is substantially smaller than that at the LISM side of the heliopause (HP), the major deflection takes place in the outer heliosheath -- the region of compressed plasma on the interstellar side of the~HP. Thus, it is reasonable to assume that the deflection of the neutral H flow is caused by an asymmetry in the LISM plasma distribution just outside the HP. This suggestion of \\citet{Lallement05} was confirmed by \\citet{Izmod05} who simulated the SW--LISM interaction with an ISMF vector $\\mathbf{B}_\\infty$ lying in the HDP at $45^\\circ$ to $\\mathbf{V}_\\infty$. The asymmetry in the plasma distribution caused by the obliquity of the ISMF (see Pogorelov \\& Matsuda~1998; Ratkiewicz et al.~1998; Pogorelov et al.~2004, 2006) does indeed result in an H-atom deflection angle close to that determined by \\citet{Lallement05}. If we ignore the interplanetary magnetic field (IMF), as it was done by \\citet{Izmod05}, the plane defined by the vectors $\\mathbf{B}_\\infty$ and $\\mathbf{V}_\\infty$ ($B$-$V$ plane) is the symmetry plane for the heliosphere. As a result, the deflection of the neutral H flow in the direction perpendicular to the $B$-$V$ plane should on the average be zero. However, as shown by \\citet{Pozank06} and \\citet{Pogo07}, this is not precisely so, since the presence of IMF eliminates the plane of symmetry and therefore allows an H-atom deflection perpendicular for the $B$-$V$ plane.  With V2 crossing the TS at a distance of about 84~AU (according to the announcement of the \\textit{Voyager} science team at the AGU Fall Meeting 2007), which is 10~AU closer to the Sun than V1 did in December~2004, it is interesting to analyze the asymmetry of the heliosphere using a fully three-dimensional numerical model based on ideal MHD description of the interacting SW and LISM ions that incorporates both the IMF and ISMF \\citep{Pozaog04,Pozaog06,Pogo07}, together with a self-consistent kinetic description for H-atoms \\citep{Jacob06,Jacob07}. These two very different sets of equations are coupled via source terms that describe momentum and energy transfer between atoms and ions through resonant charge exchange. Only on the basis of a self-consistently coupled model can we determine conclusively whether the ISMF can make the HP sufficiently asymmetric, so that the TS is substantially closer to the Sun in the V2 direction than in the V1 direction. This letter addresses the above issues by analyzing the distribution of charged and neutral particles in the heliosphere for the case when the $B$-$V$ plane is parallel to the HDP determined by \\citet{Lallement05}. We also discuss the possibility of using 2--3 kHz radio emission data \\citep{Kurth03} to further constrain the ISMF orientation. ",
        "conclusions": "New \\textit{in situ} and remote observations give us a unique opportunity to constrain previously undetermined properties of the LISM and of the ISMF threading it. We applied our MHD-kinetic model to determine the deflection of H-atoms from their original direction in the unperturbed LISM. We showed that the absence of a heliospheric symmetry plane, which is due to the presence of IMF and, very likely, the regions of slow and fast SW, can introduce an angle between the $B$-$V$ plane and the resulting HDP. However, this angle is smaller than a possible $15^\\circ$-angle between these planes determined by the accuracy of measuring the H-atom flow direction. This factor, as well as physical conditions for 2--3 kHz radio emission, can be used to adjust the orientation of the ISMF in such a way that the distribution of radio emission sources observed by the \\textit{Voyager} spacecraft will be approximately aligned with the Galactic plane. The angle of about $30^\\circ$ between $\\mathbf{B}_\\infty$ and $\\mathbf{V}_\\infty$ is sufficient to introduce a $4^\\circ$-degree average deflection of the H-atom flow from its original direction in the unperturbed LISM."
    },
    "0801/0801.4484_arXiv.txt": {
        "abstract": "The Alpha Magnetic Spectrometer (AMS) to be installed on the International Space Station (ISS) will be equipped with a proximity Ring Imaging Cherenkov (RICH) detector for measuring the velocity and electric charge of the charged cosmic particles. This detector will contribute to the high level of redundancy required for AMS as well as to the rejection of albedo particles. Charge separation up to iron and a velocity resolution of the order of 0.1\\% for singly charged particles are expected. A RICH protoptype consisting of a detection matrix with 96 photomultiplier units, a segment of a conical mirror and samples of the radiator materials was built and its performance was evaluated. Results from the last test beam performed with ion fragments resulting from the collision of a 158\\,GeV/c/nucleon primary beam of indium ions (CERN SPS) on a lead target are reported. The large amount of collected data allowed to test and characterize different aerogel samples and the sodium fluoride radiator. In addition, the reflectivity of the mirror was evaluated. The data analysis confirms the design goals. \\vspace{-0.5cm} ",
        "introduction": "AMS~\\cite{bib:ams} (Alpha Magnetic Spectrometer) is a precision spectrometer designed to search for cosmic antimatter, dark matter and to study the relative abundance of elements and isotopic composition of the primary cosmic rays. It will be installed in the International Space Station (ISS), where it will operate at least for three years.  The spectrometer  will be capable of measuring the rigidity ($R\\equiv pc/ |Z| e$), the charge ($Z$), the velocity ($\\beta$) and the energy ($E$) of cosmic rays within a geometrical acceptance of $ \\sim$0.5\\,m$^2$sr. Fig. \\ref{fig:ams} shows a schematic view of the AMS spectrometer. The system is composed of several subdectors: Transition Radiation Detector (TRD), Time-of-Flight (TOF), Silicon Tracker (STD), Anticoincidence Counters (ACC), superconducting magnet, Ring Imaging \\CK\\ detector and electromagnetic calorimeter (ECAL). \\begin{figure}[htb] \\begin{center} \\vspace{0.2cm} \\epsfig{file=AMS2.DET.3D.May03-newII.ps,width=0.71\\linewidth,clip=,bbllx=0,bblly=0,bburx=612,bbury=550} \\vspace{-0.5cm} \\caption{A whole view of the AMS Spectrometer.} \\end{center} \\label{fig:ams} \\vspace{-1.0cm} \\end{figure} ",
        "conclusions": "AMS-02 will be equipped with a proximity focusing RICH enabling velocity measurements with a resolution of about 0.1\\% and extending the charge measurements up to the iron element. Velocity reconstruction is made with a likelihood method. Charge reconstruction is made in an event-by-event basis. Evaluation of both algorithms on real data taken with in-beam tests at CERN, in October 2003 was done. The detector design was validated and a refractive index 1.05 aerogel was chosen for the radiator, fulfilling both the demand for a large light yield and a good velocity resolution. The RICH detetector is being constructed and its assembling to the AMS complete setup is foreseen for 2008."
    },
    "0801/0801.2431_arXiv.txt": {
        "abstract": "The acceleration of the universe can be explained either through dark energy or through the modification of gravity on large scales. In this paper we investigate modified gravity models and compare their observable predictions with dark energy models. Modifications of general relativity are expected to be scale-independent on super-horizon scales and scale-dependent on sub-horizon scales. For scale-independent modifications, utilizing the conservation of the curvature scalar and a parameterized post-Newtonian formulation of cosmological perturbations, we derive results for large scale structure growth, weak gravitational lensing, and cosmic microwave background anisotropy. For scale-dependent modifications, inspired by recent $f(R)$ theories we introduce a parameterization for the gravitational coupling $G$ and the post-Newtonian parameter $\\gamma$. These parameterizations provide a convenient formalism for testing general relativity. However, we find that if dark energy is generalized to include both entropy and shear stress perturbations, and the dynamics of dark energy is unknown a priori, then modified gravity cannot in general be distinguished from dark energy using cosmological linear perturbations. ",
        "introduction": "\\label{sec:intro}  Cosmic acceleration has been revealed by measurements of the redshift-distance relation $\\chi(z)$ where $\\chi$ is the comoving radial distance.  The Hubble expansion rate follows from $H(z)=(d\\chi/dz)^{-1}$ (in units where $c=1$). This determination assumes only that the observable universe is adequately described by the Robertson-Walker metric, an assertion that is testable empirically \\cite{Hogg04,LuHellaby07} and does not imply the validity of general relativity (GR).  The inference of dark energy follows once the Einstein field equations of general relativity are imposed on the Robertson-Walker metric yielding the Friedmann equations. These equations imply that a stress-energy-momentum component with negative pressure is needed to explain cosmic acceleration. This substance may be vacuum energy (i.e., a cosmological constant, giving rise to the $\\Lambda$CDM model) or a scalar field \\cite{Copeland06}. The dark energy equation of state for uniform expansion, $p(\\rho)$, can be determined from measurement of $\\rho(z)$, which itself follows from $H(z)$ combined with the first Friedmann equation. Measuring $w(z)=p/\\rho$ is the primary goal of dark energy experiments \\cite{DETF}.  Another possibility is that general relativity requires modification on large distance scales and at late times in the universe.  In this case cosmic acceleration would arise not from dark energy as a substance but rather from the dynamics of modified gravity. Modified gravity is not particularly attractive theoretically, but the observed cosmic acceleration is so surprising that all plausible explanations should be considered.  Observations of the cosmic expansion history cannot distinguish dark energy from modified gravity \\cite{fRexphist,Song06}.  Testing gravity requires exploring the evolution of spatial inhomogeneity (e.g.\\ \\cite{ZLBD07, inhom} and references therein).  Modified gravity theories must pass tests within the solar system and in relativistic binaries \\cite{Will06}.  They are expected to show significant departures from general relativity only on cosmological distance scales. The combination of cosmic microwave background anisotropy, weak gravitational lensing, and the growth of clustering of dark matter and galaxies provides an opportunity to discriminate between dark energy and modified gravity.  Performing cosmological tests of modified gravity requires a set of predictions.  There are two approaches to generating these predictions. In the first, a theory is specified by its Lagrangian (or other fundamental description) which provides the equations of motion for both homogeneous expansion and cosmological perturbations.  A class of theories can be specified by giving a Lagrangian with free parameters, e.g. $f(R)$ theories where $R$ is the Ricci scalar \\cite{Song06}.  A second approach is inspired by the parameterized post-Newtonian framework for solar system tests \\cite{ThorneWill71}. Here one begins with the solution of the gravitational field equations (i.e., the metric) instead of the Lagrangian.  Several authors have recently adopted this framework or a similar one \\cite{HuSawicki07b,Caldwell07,Amendola07,Amin07,Linder07}. The difficulty here is to find a good ``Newtonian'' description in cosmology to which one adds ``Post-Newtonian'' parameters. Metric perturbations in the scalar sector governing the growth of cosmic structure are characterized by two spatial scalar fields, $\\Phi$ (the Newtonian potential) and $\\Psi$ (the spatial curvature potential). Even if we introduce the relationship $\\Psi=\\gamma\\Phi$ with Eddington parameter $\\gamma$ \\cite{Eddington22}, there remains one unknown function of space and time. In the solar system case, by contrast, $\\Phi=-GM/r$ is known to provide an excellent approximation to planetary dynamics.  On solar system scales, and even within galaxies, one can use test-particle orbits to determine $\\Phi$, and light rays to determine $\\Phi+\\Psi$.  This comparison yields impressive limits on $|\\gamma-1|$ in the solar system \\cite{Will06}.  However, on a scale of several kpc, gravitational lensing combined with stellar dynamics in elliptical galaxies yields a current best result $|\\gamma-1|=0.02\\pm0.07$ \\cite{Bolton06}.  At the scale of Gpc where dark energy appears to drive accelerated expansion, there are no longer any bound test-particle orbits to measure $\\Phi$, so a different approach, based on cosmological perturbation theory, is needed.  Previous work in the cosmological parameterized post-Newtonian framework has either assumed that some of the Einstein field equations remain valid with modified gravity \\cite{Caldwell07} or has examined the dynamics of individual theories, e.g.\\ \\cite{Dore07}.  Neither approach is ideal. One would prefer to sample all possible theories in a broad class, and for each theory to constrain the potentials by a consistency condition that does not assume general relativity or any particular modification thereof.  Such a consistency condition was found recently in Ref.\\ \\cite{Bertschinger06} for the long-wavelength perturbations of a Robertson-Walker spacetime.  This result was derived assuming that gravity is described by a classical four-dimensional metric theory having a well-defined infrared limit (i.e., the theory is well behaved for very long wavelength perturbations).  For practical application, assumptions must also be made about the background spatial curvature and entropy perturbations.  Assuming an inflationary or equivalent origin of perturbations, long-wavelength isentropic perturbations are imprinted in the spatial curvature on a flat background.  In general relativity, these spatial curvature fluctuations, represented by the gauge-invariant $\\zeta$ field of Bardeen et al.\\ \\cite{BST83} or the ${\\cal R}$ field of Lyth \\cite{Lyth85}, are time-independent in the long-wavelength limit. Ref.\\ \\cite{Bertschinger06} presented a derivation of the conserved curvature perturbation (calling it $\\kappa$) making no assumption about the field equations except that they have a well-defined infrared limit. Physically this means that curvature perturbations are small and that all waves propagate causally.  In what follows, the curvature perturbation introduced in Ref. \\cite{Bertschinger06} will be called $\\zeta$ although its definition differs from that of Ref. \\cite{BST83}. For long wavelength perturbations on a flat background, ${\\cal R}=\\zeta$.  In the long-wavelength limit all cosmological perturbations factorize into functions of time multiplying the curvature perturbation or its spatial derivatives.  This is true in GR and in modified gravity theories that are well-behaved in the infrared limit. One might naively expect this factorization to hold only on scales larger than the Hubble length. In general relativity, however, signals in the scalar sector propagate at the speed of sound, not the speed of light \\cite{Bertschinger96}, leading to conservation of $\\zeta$ on scales larger than the Jeans length.  To satisfy solar system tests, modified gravity theories for cosmic acceleration must introduce a length scale $L_G$ below which general relativity is recovered. This length scale might be associated, for example, with the dynamics of new scalar degrees of freedom.  The value of this length scale, compared with the size of the systems investigated, plays a crucial role in characterizing the behavior of modified gravity theories.  If $L_G$ is much smaller than the length scales over which linear cosmological structure formation is measured (e.g., $L_G=1$ Mpc), then the factorization of cosmological perturbations on scales larger than the Jeans length remains valid.  We denote this case scale-independent modified gravity.  These theories are like GR in that the curvature perturbation is conserved for the relevant length scales.  This condition yields a great simplification of the dynamics, reducing cosmological perturbations to quadratures.  In GR, waves propagating in the scalar sector travel only at the speed of sound, so that scale-dependence of transfer functions arises only below the Jeans length.  Modified gravity theories, however, typically have additional fields supporting waves that travel at the speed of light.  In this case, $L_G$ is the Hubble length and the factorization of cosmological perturbations no longer holds.  Theories of this type are called scale-dependent modified gravity theories.  Now two quantities, $\\gamma$ and the gravitational coupling $G_\\Phi$ (the generalization of Newton's constant in the Poisson equation for $\\Phi$), are needed to characterize gravity, and both will vary with length scale as well as with time.  Even such complicated models can still be approximated by parameterizations, as we will discuss below. Ref. \\cite{HuSawicki07b} found a way of bridging super- and sub-horizon modifications to GR. Our parametrization, while not as general, is simpler because it only involves a few free parameters and no free functions.  Because we do not start with a Lagrangian, we cannot explain cosmic acceleration.  We take the cosmic expansion history as given from observations. Rather than providing a complete theory of modified gravity, we provide a framework for observational tests of gravity in cosmology.  This paper is organized as follows.  Section \\ref{sec:long} describes the curvature perturbation and its use to build scale-independent modified gravity theories.  Section \\ref{sec:observe} works out the growth of structure on sub-horizon scales, shows that the Poisson equation is modified, and derives results for cosmic microwave background anisotropy and weak gravitational lensing for scale-independent modified gravity theories.  Section \\ref{sec:f(R)} then examines the sub-horizon behavior of a currently popular class of theories known as $f(R)$ models and shows that they are scale-dependent.  Section \\ref{sec:shear} considers the alternative hypothesis that dark energy has a peculiar stress tensor while gravity is governed by GR. Finally, results are summarized and conclusions are presented in Section \\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  Parameterizing modified gravity theories in cosmology is much more difficult than parameterizing post-Newtonian gravity in the solar system because the gravitational potentials $\\Phi$ and $\\Psi$ in the GR limit are not static Coulomb potentials in cosmology.  In GR, on scales larger than the Jeans or nonlinear length scales, $\\Psi=\\Phi$ factorizes into a product of a time-dependent growth function and a spatially-varying curvature perturbation.  If this factorization persists in modified gravity theories, then a simple post-Newtonian parameterization can be obtained.  This is the approach we introduced in Section \\ref{sec:long}.  It has the practical virtue of yielding easily calculated predictions for all observables in the linear regime, including the growth of matter clustering, peculiar velocities, microwave background anisotropy, and weak gravitational lensing.  Introducing the Eddington parameterization $\\Psi=\\gamma\\Phi$, where $\\gamma$ depends on time but not on space, we showed that structure grows faster (gravity is stronger) in models with $\\gamma<1$. However, the shape of the transfer functions (i.e., their dependence on spatial wavenumber) for these scale-independent models is unchanged compared with GR.  Unfortunately, realistic modified gravity theories such as the $f(R)$ models have scale-dependent effects and can no longer be described by only one post-Newtonian quantity $\\gamma$.  Instead, the strength of gravity, described by $G_\\Phi$ in the Poisson equation, can vary independently of $\\gamma$.  Even so, because galaxy clustering and dynamics depends on $G_\\Phi$ but not on $\\gamma$, while weak lensing depends on $\\gamma$, observations could, in principle, measure modified gravity parameters assuming there is no dark energy.  If dark energy is complex enough to require two additional fields to characterize its stress tensor (e.g., shear stress potential and entropy), then it appears that there is enough information in the dark energy model to account for any $\\Phi$ and $\\Psi$ without modifying gravity.  One cannot prove gravity is modified unless one can account for all significant contributions to the stress-energy tensor.  Thus, our hope to describe all modified gravity models with two parameters, yielding predictions measurably different from all dark energy models, has not been realized.  Distinguishing modified gravity from dark energy will require making additional assumptions.  Nevertheless, the $(\\beta,s)$ parameterization of scale-independent modified gravity presented in Section \\ref{sec:long}, and the $(\\alpha_1,\\alpha_2,\\beta_1,\\beta_2)$ parameterization of scale-dependent modified gravity models presented in Section \\ref{sec:scaledep}, are still useful for characterizing observational data.  If measurements of galaxy clustering, peculiar velocities, and weak lensing are all consistent with $\\beta=0$ and $\\alpha_1=\\alpha_2=\\beta_1=\\beta_2=0$, for example, then modified gravity and exotic dark energy models can both be excluded.  If measurements require nonzero parameters, however, dark energy and modified gravity remain viable explanations until additional assumptions are made to distinguish them,e.g.\\ restriction of the Lagrangian to a particular form.  A generic prediction of modified gravity theories in cosmology is that the gravitational coupling $G_\\Phi$ in the Poisson equation should vary with time and with length scale.  Departures from GR could be important not only in the linear regime of cosmological perturbations but perhaps also in the nonlinear regime (albeit on scales much larger than the solar system).  Nonlinear effects may allow modified gravity to be distinguished from exotic dark energy, assuming that the dark energy fluctuations are small.  For this reason it would be valuable to perform N-body simulations of structure formation using variable $G_\\Phi$, extending previous work \\cite{Nbody} to the scale-independent and scale-dependent modified gravity models discussed in the current paper."
    },
    "0801/0801.0602_arXiv.txt": {
        "abstract": "A simple explicit example of a Roberts--type dynamo is given in which the $\\alpha$--effect of mean--field electrodynamics exists in spite of point--wise vanishing kinetic helicity of the fluid flow. In this way it is shown that $\\alpha$--effect dynamos do not necessarily require non--zero kinetic helicity. A mean--field theory of Roberts--type dynamos is established within the framework of the second--order correlation approximation. In addition numerical solutions of the original dynamo equations are given, that are independent of any approximation of that kind. Both theory and numerical results demonstrate the possibility of dynamo action in the absence of kinetic helicity.  Key words: Mean--field dynamo action, $\\alpha$--effect, modified Roberts dynamo ",
        "introduction": "\\label{intro}  The essential breakthrough in the understanding of the origin of the large--scale magnetic fields of cosmic objects came with the development of mean--field dynamo theory. A central component of this theory is the $\\alpha$--effect, that is, a mean electromotive force with a component parallel or antiparallel to the mean magnetic field in a turbulently moving electrically conducting fluid. The $\\alpha$--effect, which occurs naturally in inhomogeneous turbulence on a rotating body, is a crucial element of the dynamo mechanisms proposed and widely accepted for the Sun and planets, for other stellar objects and even for galaxies; see, e.g., \\cite{krauseetal80,ruedigeretal04,brandenburgetal05}.  In the majority of investigations this effect has been merely calculated in the so-called second--order correlation approximation (SOCA), or first--order smoothing approximation (FOSA), which ignores all contributions of higher than second order in the turbulent part of the fluid velocity. Moreover in many cases attention has been focussed on the high--conductivity limit only, which can be roughly characterized by short correlation times of the turbulent motion in comparison to the magnetic--field decay time for a turbulent eddy. Under these circumstances the $\\alpha$--effect is closely connected with some average over the kinetic helicity of the turbulent motion, that is, of $\\bu \\cdot \\bomega$, where $\\bu$ means the turbulent part of the fluid velocity, $\\bomega$ the corresponding vorticity, $\\bomega = \\bnab \\x \\bu$; see, e.g., \\cite{krauseetal80,ruedigeretal04,raedler00b,brandenburgetal05,raedleretal07}. More precisely, the coefficient $\\alpha$ for isotropic turbulence turns out to be equal to $- \\frac{1}{3} \\int_0^\\infty \\ol{\\bu (\\bx, t) \\cdot \\bomega (\\bx, t - \\tau)} \\, \\dd \\tau$, often expressed in the form $- \\frac{1}{3} \\ol{\\bu \\cdot \\bomega} \\, \\tau_c$ with equal arguments of $\\bu$ and $\\bomega$ and some appropriate time $\\tau_c$. Here and in what follows overbars indicate averages. For anisotropic turbulence the trace of the $\\alpha$--tensor is just equal to three times this value of $\\alpha$.  These findings have been sometimes overinterpreted in the sense that mean--field dynamos, or even dynamos at all, might not work without kinetic helicity of the fluid motion, that is, if $\\bu \\cdot \\bomega$ vanishes. There exist however several counter--examples.  Firstly, the mean-field dynamo theory offers also dynamo mechanisms without $\\alpha$--effect, for which the average of the kinetic helicity may well be equal to zero. For instance, a combination of the so--called $\\bOmega \\x \\bJ$--effect, which may occur even in the case of homogeneous turbulence in a rotating body, with shear associated with differential rotation, can act as a dynamo \\cite{raedler69b,raedler70,roberts72,raedler76,stix76b,raedler80,moffattetal82,raedler86,raedleretal03}. The possibility of a dynamo due to turbulence influenced by large-scale shear and the shear flow itself \\cite{rogachevskiietal03,rogachevskiietal04,rogachevskiietal06b} is still under debate \\cite{brandenburg05b,raedleretal06,ruedigeretal06,BRRK07}.  Secondly, in dynamo theory beyond the mean--field concept a number of examples of dynamos without kinetic helicity are known. The dynamo proposed by Herzenberg \\cite{herzenberg58}, usually considered as the first existence proof for homogeneous dynamos at all, works without kinetic helicity. Likewise in the dynamo models of Gailitis working with an axisymmetric meridional circulation in cylindrical or spherical geometry \\cite{gailitis70,gailitis93,gailitis93b} there is no kinetic helicity. We mention here further the dynamo in a layer with hexagonal convection cells proposed by Zheligovsky and Galloway \\cite{zheligovskyetal98}, in which the kinetic helicity is zero everywhere.  Thirdly, it is known since the studies by Kazantsev \\cite{kazantsev68} that small--scale dynamos may work without kinetic helicity. This fact has meanwhile been confirmed by many other investigations; see, e.g., \\cite{brandenburgetal05}.  Let us return to the mean--field concept but restrict ourselves, for the sake of simplicity, to homogeneous and statistically steady turbulence. We stay first with the second--order correlation approximation but relax the restriction to the high--conductivity limit. As is known from early studies, e.g. \\cite{krauseetal71b,krauseetal80,moffattetal82}, the crucial parameter for the $\\alpha$--effect is then no longer the mean kinetic helicity $\\ol{\\bu \\cdot \\bomega}$. The $\\alpha$-effect is in general determined by the function $h (\\bxi, \\tau) = \\ol{\\bu (\\bx, t) \\cdot \\bomega (\\bx + \\bxi, t + \\tau)}$ or, what is equivalent, by the kinetic helicity spectrum, defined as its Fourier transform with respect to $\\bxi$ and $\\tau$. In the high--conductivity limit this brings us back to the above--mentioned results. In the low conductivity limit, that is, large correlation times in comparison to the magnetic--field decay time for a turbulent eddy, the $\\alpha$--effect is closely connected with the quantity $\\ol{\\bu \\cdot \\bpsi}$, where $\\bpsi$ is the vector potential of $\\bu$, that is $\\bnab \\x \\bpsi = \\bu$ with $\\bnab \\cdot \\bpsi = 0$. For isotropic turbulence the coefficient $\\alpha$ is then equal to $ - \\frac{1}{3 \\eta} \\, \\ol{\\bu \\cdot \\bpsi}$, where $\\eta$ means the magnetic diffusivity of the fluid \\cite{krauseetal71b,krauseetal80,raedler00b,raedleretal07,suretal07}. In the anisotropic case the trace of the $\\alpha$--tensor is again equal to three times this value of $\\alpha$. In general $\\ol{\\bu \\cdot \\bomega}$ does not vanish in this limit but it is without interest for the $\\alpha$--effect. Both $\\ol{\\bu \\cdot \\bomega}$ and $\\ol{\\bu \\cdot \\bpsi}$ are quantities, which, if non--zero, indicate deviations of the turbulence from reflectional symmetry. In the second--order correlation approximation the $\\bOmega \\x \\bJ$--effect, too, is completely determined by the kinetic helicity spectrum \\cite{moffattetal82}.  Beyond the second--order correlation approximation the $\\alpha$--effect is no longer determined by the kinetic helicity spectrum alone. In an example given by Gilbert et al.\\ \\cite{gilbertetal88} an $\\alpha$--effect (with zero trace of the $\\alpha$--tensor) occurs in a non-steady flow even with zero kinetic helicity spectrum. Note that small-scale dynamos may also work with zero kinetic helicity spectrum, but they do not produce large scale fields.  Interesting simple dynamo models have been proposed by G.\\ O.\\ Roberts \\cite{robertsgo70,robertsgo72}. In his second paper \\cite{robertsgo72} he considered dynamos due to steady fluid flows which are periodic in two Cartesian coordinates, say $x$ and $y$, but independent of the third one, $z$. When speaking of a ``Roberts dynamo\" in what follows we refer always to the first flow pattern envisaged there [equation (5.1), figure 1]. This dynamo played a central role in designing the Karlsruhe dynamo experiment \\cite{muelleretal00,stieglitzetal01,muelleretal02,stieglitzetal02}. It can be easily interpreted within the mean--field concept. In that sense it occurs as an $\\alpha$--effect dynamo with an anisotropic $\\alpha$-effect. The flow pattern proposed by Roberts, and in a sense realized in the \\mbox{Karlsruhe} experiment, shows non--zero kinetic helicity. In this paper we want to demonstrate that this kind of dynamo works also with a slightly modified flow pattern in which the kinetic helicity is exactly equal to zero everywhere (but the kinetic helicity spectrum remains non--zero).  In section II we present a simple mean--field theory of Roberts dynamos at the level of the second--order correlation approximation. In section III we report on numerical results that apply independently of this approximation. Finally, in section IV some conclusions are discussed. ",
        "conclusions": "We have examined a modified version of the Roberts dynamo with a steady fluid flow as it was considered in the context of the \\mbox{Karlsruhe} dynamo experiment. By contrast to what might be suggested by simple explanations of this dynamo the necessary deviation of the flow from reflectional symmetry is not adequately described by the kinetic helicity $\\bu \\cdot \\bomega$ of the fluid flow or its average $\\ol{\\bu \\cdot \\bomega}$. Dynamo action with steady fluid flow requires in truth a deviation which is indicated by a non--zero value of the quantity $\\ol{\\bu \\cdot \\bpsi}$, where $\\bpsi$ is the vector potential of $\\bu$. Even if the kinetic helicity is equal to zero everywhere, dynamo action occurs if only the quantity $\\ol{\\bu \\cdot \\bpsi}$, the helicity of the vector potential of the flow, is unequal to zero. This fact is not surprising in light of the general findings of mean--field electrodynamics.  We stress that our result does not mean that the quantity $\\ol{\\bu \\cdot \\bpsi}$ is in general more fundamental for the $\\alpha$--effect and its dynamo action than $\\ol{\\bu \\cdot \\bomega}$. In cases with steady fluid flow indeed $\\ol{\\bu \\cdot \\bpsi}$ is crucial. However, with flows varying rapidly in time, the crucial quantity is $\\ol{\\bu \\cdot \\bomega}$.  Both the mean kinetic helicity $\\ol{\\bu \\cdot \\bomega}$ and the quantity $\\ol{\\bu \\cdot \\bpsi}$ are determined by the helicity spectrum of the flow. That is, the Roberts dynamo does not work with zero kinetic helicity spectrum. In this sense the situation is different from that in the interesting example of an $\\alpha$--effect dynamo given by Gilbert et al.\\ \\cite{gilbertetal88}, which works with a very special non--steady flow for which the kinetic helicity spectrum is indeed zero.  Looking back at the Karlsruhe experiment we may state that a modified (technically perhaps more complex) version without kinetic helicity of the fluid flow should show the dynamo effect, too. As table~\\ref{tab1} exemplifies, however, the kinetic helicity can well reduce the dynamo threshold. In that sense we may modify the title of the paper of Gilbert et al.\\ and say: Helicity is unnecessary {\\it for Roberts dynamos} -- but it helps!  \\bigskip  {\\bf Acknowledgement} The results reported in this paper have been obtained during stays of KHR at NORDITA. He is grateful for its hospitality. AB thanks the Kavli Institute for Theoretical Physics for hospitality. This research was supported in part by the National Science Foundation under Grant No.\\ PHY05-51164."
    },
    "0801/0801.4601_arXiv.txt": {
        "abstract": "{} {We analyse the point-symmetric planetary nebula \\object{NGC\\,6309} in terms of its three-dimensional structure and of internal variations of the physical conditions to deduce the physical processes involved in its formation.} {We used VLA-D $\\lambda3.6$-cm continuum, ground-based, and HST-archive imaging as well as long slit high- and low-dispersion spectroscopy.} {The low-dispersion spectra indicate a high excitation nebula, with low to medium variations of its internal physical conditions ($10,600\\,{\\rm K}     \\lesssim T_{\\rm e}$[\\ion{O}{iii}] $\\lesssim 10,900\\,{\\rm K}$; $10,100\\,{\\rm K}     \\lesssim T_{\\rm e}$[\\ion{N}{ii}]  $\\lesssim 11,800\\,{\\rm K}$; $1440\\,{\\rm cm}^{-3} \\lesssim N_{\\rm e}$[\\ion{S}{ii}]  $\\lesssim 4000\\,{\\rm cm}^{-3}$; $1700\\,{\\rm cm}^{-3} \\lesssim N_{\\rm e}$[\\ion{Cl}{iii}]$\\lesssim 2600\\,{\\rm cm}^{-3}$; $1000\\,{\\rm cm}^{-3} \\lesssim N_{\\rm e}$[\\ion{Ar}{iv}] $\\lesssim 1700\\,{\\rm cm}^{-3}$). The radio continuum emission indicates a mean electron density of $\\simeq1900$\\,cm$^{-3}$, emission measure of $5.1\\times10^5$\\,pc\\,cm$^{-6}$, and an ionised mass $M$(\\ion{H}{ii})$\\simeq0.07\\,$M$_\\odot$. In the optical images, the point-symmetric knots show a lack of [\\ion{N}{ii}] emission as compared with similar features previously known in other PNe. A rich internal structure of the central region is seen in the HST images, resembling a deformed torus. Long slit high-dispersion spectra reveal a complex kinematics in the central region, with internal expansion velocities ranging from $\\simeq20$ to 30\\,km\\,s$^{-1}$. In addition, the spectral line profiles from the external regions of NGC\\,6309 indicate expanding lobes ($\\simeq40$\\,km\\,s$^{-1}$) as those generally found in bipolar nebulae. Finally, we have found evidence for the presence of a faint halo, possibly related to the envelope of the AGB-star progenitor.} {Our data indicate that NGC\\,6309 is a quadrupolar nebula with two pairs of bipolar lobes whose axes are oriented PA=${40\\degr}$ and PA=${76\\degr}$. Equatorial and polar velocities for these two pairs of lobes are 29 and 86\\,km\\,s$^{-1}$ for the bipolar system at PA=${40\\degr}$ and 25 and 75\\,km\\,$s^{-1}$ for the bipolar system at PA=${76\\degr}$. There is also a central torus that is expanding at 25\\,km\\,s$^{-1}$. Kinematical age for all these structures is around 3700 to 4000\\,yr. We conclude that NGC\\,6309 was formed by a set of well-collimated bipolar outflows (jets), which were ejected in the initial stages of its formation as a planetary nebula. These jets carved the bipolar lobes in the previous AGB wind and their remnants are now observed as the point-symmetric knots tracing the edges of the lobes.} ",
        "introduction": "In spite of all the work that up to now has been done to understand the origin of planetary nebulae (PNe) morphologies and their evolution (e.g., Kwok, Purton \\& FitzGerald \\cite{kwo78}; Kahn \\& West \\cite{kah85}; Balick \\cite{bal87}; Icke \\cite{ick88}; Mellema \\cite{mel95}; Perinotto et al. \\cite{per04}; Rijkhorst, Mellema \\& Icke \\cite{rij05}; Sch\\\"onberner et al. \\cite{sch05}; Sch\\\"onberner et al. \\cite{sch07}), some morphological structures still remain as unsolved problems. In particular, the so-called point-symmetry (Stanghellini, Corradi \\& Schwarz, \\cite{sta93}; Gon\\c calves et al. \\cite{gon03}) appears enigmatic. Some observational studies about point-symmetric PNe (e.g. Miranda \\& Solf \\cite{mir92}; L\\'opez, Meaburn \\& Palmer \\cite{lop93}) relate the formation of these objects to collimated outflows from a precessing central source known as bipolar rotating episodic jets (BRETs, see L\\'opez, V\\'azquez \\& Rodr\\'{\\i}guez \\cite{lop95}). However, in some cases, there are not definite proofs about the jet nature of these features as they are not confirmed by spectroscopic studies (e.g. V\\'azquez et al. \\cite{vaz99a}, V\\'azquez et al.\\cite{vaz02}). Although some theoretical models have intended to explain point-symmetric PNe (e.g., Cliffe et al. \\cite{cli95}; Livio \\& Pringle \\cite{liv96}, \\cite{liv97}; Garc\\'{\\i}a-Segura \\& L\\'opez \\cite{gar00}; Rijkhorst, Icke \\& Mellema \\cite{rij04}), the origin and shaping of this kind of objects remains a puzzle.  NGC 6309 is a PN whose morphology strongly suggests a BRET scenario for its origin, as can be seen in the H$\\alpha$+[\\ion{N}{ii}] and [\\ion{O}{iii}] images by Schwarz, Corradi \\& Melnick (\\cite{sch92}). It has two prominent, point-symmetric `arms' formed by pairs of condensations in addition to a bright internal elliptical structure. Some previous studies on \\object{NGC\\,6309} have determined an expansion velocity $V_{\\rm exp}$[\\ion{O}{iii}]=34\\,km\\,s$^{-1}$ (Sabbadin \\cite{sab84}), as well as mean physical parameters and total abundances (G\\'orny et al. \\cite{gor04}), namely, electron density and temperature $N_{\\rm e}$[\\ion{S}{ii}]=2600\\,cm$^{-3}$, $T_{\\rm e}$[\\ion{N}{ii}]=12,097\\,K, $T_{\\rm e}$[\\ion{O}{iii}]=11,845\\,K; and abundance ratios He/H=0.1, N/H=8.20, O/H=8.64, Ne/H=7.82, and S/H=6.49 (log H=12, except for He/H ratio). Armour \\& Kingsburgh (\\cite{arm01}) found similar values. NGC\\,6309 has also been included in the list of PNe with low ionisation structures (LIS) by Gon\\c calves et al. (\\cite{gon03}). Finally, the central star of \\object{NGC\\,6309} has been classified as a ``weak emission line star'' by G\\'orny et al. (\\cite{gor04}).  In spite of this suggestive morphology, an internal kinematic study of NGC\\,6309, as well as a full analysis of the physical conditions in the different nebular regions, has not yet been done. In this paper, we have carried out a radio-optical study of NGC\\,6309, including radio continuum mapping, optical imaging, and long-slit optical spectroscopy, in both, high- and low-dispersion. We discuss our results on the morphology, kinematics, physical conditions, ionic abundances, and nature of the gas emission to explore the possible mechanisms involved in the formation of NGC\\,6309. ",
        "conclusions": "We have carried out an analysis of the planetary nebulae NGC\\,6309, based on ground-based and space-based imaging, high- and low-dispersion spectroscopy, as well as VLA-D radio continuum. We summarize the main conclusions of this work as follows.  NGC\\,6309 can be described as a quadrupolar PN formed by a bright central torus, two systems of bipolar lobes oriented at different directions, and point-symmetric knots that trace the edges of the lobes. The torus expands at 25\\,km\\,s$^{-1}$, whereas the polar expansion velocity of the lobes is 75\\,km\\,s$^{-1}$ for a first bipolar system at PA 76\\degr, and 86\\,km\\,s$^{-1}$ for another bipolar system at PA 40\\degr. Assuming a distance of 2\\,kpc for the nebula, the kinematic ages of the structures ranges from 3700 to 4000\\,yr, suggesting that they have been formed in a short time span. The knots at the edges of the lobes suggest that the lobes have been formed by rapidly precessing bipolar jets that carved cavities in the previous red giant envelope. In addition to these structures, we detect a circular halo surrounding the torus and bipolar lobes, which probably corresponds to the envelope ejected in the AGB phase by the central star. There is also a conelike structure embedded in one of the lobes and with its base on the torus.  We also study internal variations of the physical conditions and chemical abundances in NGC\\,6309. The low-dispersion spectra indicate a high-excitation nebula, with low to medium variations of its internal physical conditions ($10,600\\,{\\rm K}     \\lesssim T_{\\rm e}$[\\ion{O}{iii}] $\\lesssim 10,900\\,{\\rm K}$; $10,100\\,{\\rm K}     \\lesssim T_{\\rm e}$[\\ion{N}{ii}]  $\\lesssim 11,800\\,{\\rm K}$; $1440\\,{\\rm cm}^{-3} \\lesssim N_{\\rm e}$[\\ion{S}{ii}]  $\\lesssim 4000\\,{\\rm cm}^{-3}$; $1700\\,{\\rm cm}^{-3} \\lesssim N_{\\rm e}$[\\ion{Cl}{iii}]$\\lesssim 2600\\,{\\rm cm}^{-3}$; $1000\\,{\\rm cm}^{-3} \\lesssim N_{\\rm e}$[\\ion{Ar}{iv}] $\\lesssim 1700\\,{\\rm cm}^{-3}$). The radio continuum emission indicates a mean electron density of $\\simeq1900$\\,cm$^{-3}$; emission measure of $5.1\\times10^5$\\,pc\\,cm$^{-6}$; and an ionised mass $M$(\\ion{H}{ii})$\\simeq0.07\\,$M$_\\odot$. The logarithmic extinction coefficient $c_{\\rm H \\beta}$ range from 0.70 to 0.97 (mean value 0.87), which probably is produced by differences in the internal dust distribution."
    },
    "0801/0801.1953_arXiv.txt": {
        "abstract": "We propose a general framework to scrutinize the performance of semi-analytic codes of galaxy formation. The approach is based on the analysis of the outputs from the model after a series of perturbations in the input parameters controlling the baryonic physics. The perturbations are chosen in a way that they do not change the results in the luminosity function or mass function of the galaxy population.  We apply this approach on a particular semi-analytic model called GalICS. We chose to perturb the parameters controlling the efficiency of star formation and the efficiency of supernova feedback. We keep track of the baryonic and observable properties of the central galaxies in a sample of dark matter halos with masses ranging from $10^{10}$ M$_\\odot$ to $10^{13}$ M$_\\odot$.   We find very different responses depending on the halo mass. For small dark matter halos its central galaxy responds in a highly predictable way to small perturbation in the star formation and feedback efficiency. For massive dark matter halos, minor perturbations in the input parameters can induce large fluctuations on the properties of its central galaxy, at least $\\sim 0.1$ in $B-V$ color or $\\sim 0.5$ mag in $U$ or $r$ filter, in a seemingly random fashion. We quantify this behavior through an objective scalar function we call predictability.  We argue that finding the origin of this behavior needs additional information from other approximations and different semi-analytic codes. Furthermore, the implementation of an scalar objective function, such as the predictability, opens the door to quantitative benchmarking of semi-analytic codes based on its numerical performance. ",
        "introduction": "Hierarchical aggregation seems to be at the heart of galaxy evolution. In a cold dark matter universe, as depicted by numerical  simulations, its structure grows through subsequent mergers and zero fragmentations. The growth and evolution of galaxies, which are thought to use dark matter as scaffolding, is channeled through this hierarchical aggregation, at least for the most massive structures \\citep{LSS_SFW}.  Notwithstanding all the complexity in the process of galaxy formation and evolution, galaxies still are the most basic population unit in the description of large scale structure in the Universe. And still nowadays much work is being invested in galaxy formation to disentangle the influence of the hierarchical context setup by dark matter from the secular baryonic processes on small scales.  Tackling the problem theoretically implies numerical experiments following large structure dynamics and, at the same time, a description of baryonic processes such hydrodynamics and radiative cooling. This is still very challenging in the usual numerical approach that discretizes space and time,  and try to solve a relevant set of equations  to capture the physics \\citep{2002Sci...295...93A,2006AIPC..878....3G}. From the computational point of view it involves achieving an effective resolution spanning at least $5$ orders of magnitude in mass, length and time \\citep{2007arXiv0705.1556N}.  To overcome this barrier the semi-analytic model (hereafter SAM) approach proposes to describe first the non-linear clustering of dark matter on large scales, and describe later the small scale baryonic physics through analytical prescriptions. The connection between the two scales is provided through the dark matter halo, which is the most basic unit of non-linear dark matter structure (see \\cite{2006RPPh...69.3101B} and references therein).  The non linear clustering of dark matter is described through a merger tree, representing the merging history of a given dark matter halo.   The construction methods for merger trees can vary, ranging from Monte-Carlo realizations based on theoretical estimates \\citep{1999MNRAS.305....1S} to the numerical based on N-body simulations \\citep{2006MNRAS.365...11C}, including also hybrid approaches mixing numerical and analytical techniques \\citep{2002MNRAS.333..623T}.   The different analytic implementations of baryonic process span a wide range of philosophical approaches, physical concepts and numerical implementations. Most of them, nonetheless, constructed from observed correlations in our local patch of Universe.  Regardless of the details of these models, what is general to all of them is the underlying merger trees structure complementary with analytic recipes describing the growth of baryonic structure inside the merger trees. The ignorance respect to the physics included in the analytic recipes is usually represented by scalars, which in most of the cases represent efficiencies of physical processes. This implies that a given realization of a semi-analytic run is completely determined by the dark matter input and the baryonic parameters in the simulation.   Most of the work during the last decade was invested in adding and exploring the effect of these parameters on average quantities, especially the luminosity function. The generic parameters to be used have been more or less settled, and most of the models have achieved a good level of internal consistency by reproducing some key observational features. The confidence on the consistency that can be achieved, and the ease to perform a semi-analytic run, have empowered the modelers to select subsamples and make studies about the most massive galaxies or correlations among populations \\citep{2007MNRAS.375....2D,2007arXiv0709.3933H}.   As the complexity and interest in semi analytic techniques grow, two relevant issues must  be addressed in more detail. First, the issue of error  propagation from the incertitude in the input parameters, a factor that might be  important in a hierarchical Universe, where the amplification of small initial errors might be important for the most massive and hierarchical objects. Second, the development of objective ways to compare different types of complexity in semi analytic models. This could allow, for instance, the implementation of simple tests with an objective scalar function to measure the model performance.  The objective of this paper is two-fold: \\begin{itemize} \\item{Propose a methodology to weigh the role of secular baryonic processes in the context of SAMs.} \\item{Propose an objective scalar function that captures the biases and general behavior of semi-analytic models regardless of its detailed implementation.} \\end{itemize}  These two objectives are a result of the same perturbative approach we advocate in this paper. This approach is based on the fact that the only objective information we have to describe the results of a semi-analytic run are the input parameters of the model. In the perturbative approach, we perform semi-analytic runs in the neighborhood of some scalar parameters. This will allow us, as we will show, to get an idea about the limits of our semi-analytic model.  This paper is structured as follows. In Section \\ref{sec:SAM} we describe the structure common to all SAMs and from that point we introduce the concept of perturbations in a semi-analytic model. In Section 3 we introduce the setup for the perturbation experiment of our SAM. We describe in Section 4 the two most relevant qualitative features of the experiment results. We select one of this qualitative results to make a detailed quantitative analysis with three different indices, these results are shown in Section 5. We discuss our results in Section 6. ",
        "conclusions": "We have used $N$-body simulations and a semi-analytic model of galaxy formation to explore the consequences of small perturbations in the input parameters of the semi-analytic model. Specifically we varied the scalar parameters regulating the star formation $\\alpha$ and supernovae feedback efficiency $\\epsilon$. We followed some physical properties for $600$ central to gauge the effect of the perturbations.   This experiment was motivated both by the interest of performing a description of a semi analytic models, making abstraction of the details in the model , and test the significance of this approach to infer the distinctive footprint of the different baryonic processes.  We find that depending on the halo mass, there are two kinds of different of behavior respect to the change in the input parameters. There is a smooth variation of the galactic properties for low mass halos, and a seemingly random unpredictable behavior limited to high mass halos.  We have quantified this behavior through an objective scalar function called  predictability, $P$, defined in Eq.\\ref{predigo}. Values $P\\simeq 1$ mean a rather smooth and predictable response,  while $P\\leq 0$  point a  more unpredictable behavior. The highest predictability is found only in low mass halos while high mass halos present almost always negative values of $P$. This notion of predictability depends on the property we are looking at. In our case the total galactic mass and the total stellar mass showed the general trend of high predictability on the $\\alpha $-$\\epsilon$ plane (upper panels in Fig.\\ref{predictfigure}). While the $(B-V)$ color and the SDSS$_U$ and SDSS$_r$ bands showed almost always negative predictabilities (lower panel in Fig.\\ref{predictfigure}).  Then, we computed the variance $\\sigma_{\\alpha\\epsilon}$ over the $\\alpha$-$\\epsilon$, $\\sigma_{\\alpha\\epsilon}$ plane and weighted it by the predictability over the same landscape. This helped us to estimate the possible variation in a given property after performing a change $\\delta\\alpha-\\delta\\epsilon$ in the input parameters. Using this measure we found that for high mass halos on should expect rather large variations in the galactic properties from a small perturbation in the baryonic parameters (Fig.\\ref{weightedfigure}).  In order to give a scale to the landscape fluctuations for a given galaxy, we compare the landscape variance with the variance in a subset of central galaxies of halos taken from the whole cosmological box simulation. The halos are selected to have similar mass as the parent halo of the galaxies we are studying . The general trend showed that at higher halo masses the variance coming from the modulation of the $\\alpha$ and $\\epsilon$ parameters is on the same order of magnitude as the variation on the galaxy properties over the whole box. The opposite trend is only found for low halo masses  (Fig.\\ref{cosmofigure}).  In this particular case of perturbation related to star formation and supernova feedback, it seems that the quantities that exhibit a dependence on the full star formation history (magnitudes and colours) are the most sensitive to variations of the $\\alpha - \\epsilon$ parameters. For instance, most of the trends we found are not found for the total galaxy mass.  In general, all this evidence seems to point towards a picture where the central galaxies hosted in massive halos, which have grown mainly through mergers, are the most  sensitive to small variations of the baryonic parameters in a way that is comparable of doing a significant variation on the mass of the host halo.   From these results one can expect that if the variation of the others scalar parameters in the model is performed, the predictability $P$, landscape variance $\\sigma_{\\lambda_i}$ and $P$-weighted landscape variance $\\sigma_{\\lambda_i P}$ should be higher than the values quoted in this paper. It is very unlikely that the variations of other parameters could cancel out exactly the influence of the efficiencies $\\alpha$ $\\epsilon$.   There is a hierarchy of causes for this behavior that must be explored. First of all, might this be a signature of the hierarchical build of galaxies? In such a picture it is easy to imagine that mild perturbations at early times might add up to finally yield very different values for very similar initial conditions. This would account for the relatively large values of the variance over the $\\alpha$-$\\epsilon$ plane compared to the intrinsic variances over the whole population of similar halos in the cosmological volume. But probably not for the low predictability values.  Could this be an artifact coming from the semi-analytic models of galaxy formation? In these models, generally the distinction between what is to be considered as the central galaxy depend on which galaxy is the most massive. This is ambiguous when various galaxies inside a dark matter halo have similar masses, in that case the selection of the central galaxy might be subject to noise. This could explain in part the seemingly random landscapes for high mass haloes.  On the last level of the hierarchy, could this be coming from \\emph{our code}? This is impossible to confirm without performing the same kind  of experiment with another fully fledged semi-analytic model. Which take us to the issue of comparison between semi-analytic models of galaxy formation. The predictability, as a meaningful scalar objective function, opens the possibility to measure the biases from different semi-analytic codes. This could allow the comparison of different codes based on its numerical performance, going beyond the rather ill-posed strategy of comparison based on astrophysical performance, i.e. reproducing observations.   Finally, the small perturbations we made on the scalar parameters were constructed  to not have any effect on the mean properties of the galaxies such as  the luminosity function. It means that formally the galaxies we have produced at every perturbation are consistent with the overall galaxy population.  As a consequence of all this, studies making use of selected subpopulations from a wider population generated using semi-analytic models, should bear in mind that this smaller population might not be unique. The dispersion on this subsample of galaxies, coming from the perturbations that can be induced on every parameter in the model, should be explicitly stated. Including that dispersion (in the form of error bars, for instance) seems a necessary condition to make a \\emph{fair use} of semi-analytic models, acknowledging in an explicit manner its limitations on predictability."
    },
    "0801/0801.3167_arXiv.txt": {
        "abstract": "{} {We study the spectral and temporal behavior of the High Mass X-ray Binary \\aof during a `pre-outburst flare' which took place $\\sim5\\,\\mathrm{d}$  before the peak of a normal (type I) outburst in \\augsept. We compare the studied behavior with that observed during the outburst.} {We analyse \\xte observations that monitored \\aof during the outburst. We complete spectral and timing analyses of the data. We study the evolution of the pulse period, present energy-dependent pulse profiles both at the initial pre-outburst flare and close to outburst maximum, and measure how the cyclotron resonance-scattering feature (hereafter CRSF) evolves.} {We present three main results: a constant period $P$=$103.3960(5)\\,\\s$ is measured until periastron passage, followed by a spin-up with a decreasing period derivative of $\\dot{P}$=$(-1.69\\,$$\\pm$${0.04})\\,\\times10^{-8}\\sps $ at MJD 53618, and $P$ remains constant again at the end of the main outburst. The spin-up provides evidence for the existence of an accretion disk during the normal outburst. We measure a CRSF energy of $E_{\\mathrm{cyc}}$$\\sim$$50\\,\\kev$ during the pre-outburst flare, and $E_{\\mathrm{cyc}}$$\\sim$$46\\,\\kev$ during the main outburst. The pulse shape, which varies significantly during both pre-outburst flare and main outburst, evolves strongly with photon energy. }{} ",
        "introduction": "\\label{sect:intro}  The Be/X-ray binary \\aof \\footnote{Referred to as 1A 0535+262 in SIMBAD} is a transient source, characterized by quiescent states with X-ray luminosity $L_{\\mathrm{X}}\\,$$\\lesssim$$10^{36}\\,\\ergs$, interrupted by normal (type I) outbursts, associated in general with periaston passage, when a luminosity $L_{\\mathrm{X}}$$\\sim$$10^{36-37}\\,\\ergs$ is reached, and giant (type II) outbursts for which $L_{\\mathrm{X}}$$>$$10^{37}\\,\\ergs$. \\aof was discovered by \\citet{rosenberg75} during a giant outburst of luminosity level of $L_{(3-7\\kev)}$$\\sim1.2$$\\times10^{37}\\ergs$. Since then, five giant outbursts have been detected in October 1980 \\citep{nagase82}, June 1983 \\citep{sembay90}, March/April 1989 \\citep{makino89}, February 1994 \\citep{finger94_2}, and May/June 2005 \\citep{tueller05}. Following the giant outburst in May/June 2005, two normal outbursts occurred in August/September 2005 \\citep{finger05_1} and December 2005 \\citep{finger05_2}.  The source \\aof is in a binary orbit, with an O9.7IIIe optical companion HDE245770 \\citep{giangrande80}, of orbital period $P_{orb}$$=$$110.3\\,$$\\pm$$0.3\\,\\mathrm{days}$ and eccentricity $e$=$0.47\\,$$\\pm$$0.02\\,$ \\citep{finger94_2}. The estimated distance of the system is $2\\,\\kpc$ \\citep{steele98}. The measured pulse period of the neutron star at MJD 53614.5137 is $P$=$103.39315(5)\\,\\s$ \\citep{caballero07}. During quiescence a spin-down trend of the neutron star has been reported \\citep{finger96,hill07}, and during giant outbursts the neutron star shows a strong spin-up. In the June 1983 giant outburst, a spin-up of $\\dot{\\nu}$$\\sim$$0.6\\times10^{-11}\\,\\hzs$ was measured \\citep{sembay90}. In the February 1994 giant outburst the spin-up reached $\\dot{\\nu}$$\\sim$$1.2\\times10^{-11}\\,\\hzs$ and quasi periodic oscillations were present, confirming the presence of an accretion disk \\citep{finger96}.  The X-ray spectrum of \\aof is typically described by a phenomenological model consisting of a power law with an exponential cutoff. Two absorption like features, interpreted as cyclotron resonance scattering features CRSFs (fundamental and first harmonic), have been observed in the spectrum at $\\sim$$45\\,\\kev$ and $\\sim$$100\\,\\kev$  \\citep{kend94,grove95}. Recent observations during the \\augsept normal outburst with \\inte and \\xte \\citep{kretschmar05,wilson05,caballero07} and \\suzaku \\citep{terada06} have confirmed this result. The centroid energy of the CRSF was measured during the main outburst at different luminosity levels and it was not found to vary with X-ray luminosity.  In this Letter we focus on the \\augsept normal outburst. At the onset of this outburst, a sharp, pre-outburst flare is observed superimposed on the gradual increase towards the peak flux in the main outburst. This pre-outburst flare develops on a time scale of 1 day, short compared to the three-week duration of the main outburst. As can be seen from the light curve measured by \\swift/\\bat , the flare is one of several such flares superimposed on the rising edge of the main outburst (see \\citealt{postnov08}). The pre-outburst flare studied here reaches a PCA flux of $810\\,\\mathrm{counts\\,s^{-1}\\,PCU^{-1}}$, and the maximum flux during the main outburst is $1005\\,\\mathrm{counts\\,s^{-1}\\,PCU^{-1}}$. The spectral and timing behaviours of the source during the pre-outburst flare appear to be different from the ones observed during the main outburst. ",
        "conclusions": "\\label{sect:summ}  In this paper, we present evidence for significant changes in the spectral and timing parameters of the Be/X-ray binary \\aof, during its normal outburst in \\augsept. Our three main results are the following: \\begin{enumerate} \\item The pulse period of the neutron star appears to be constant during the pre-outburst flare, $P$=$103.3960(5)\\,\\s$, and a spin-up starts at periastron, $\\dot{P}$=$(-1.69\\,\\pm{0.04})\\,\\times10^{-8}\\sps$ measured at MJD 53618. The pulse period falls exponentially at the end of the outburst.  \\item The energy-dependent, pulse profiles during the pre-outburst flare are significantly different from those measured for the main outburst.  \\item During the pre-outburst flare, the fundamental cyclotron-line energy centroid reaches $E_{\\mathrm{cyc}}$=$52.0^{+1.6}_{-1.4}\\,\\kev$, significantly higher than for the main outburst, $E_{\\mathrm{cyc}}$=$46.1^{+0.5}_{-0.5}\\,\\kev$. \\end{enumerate}  This is the first observation of a normal outburst that measures a spin-up rate for \\aof, providing evidence that an accretion disk is being detected during a Type I outburst. The measured spin-up $\\dot{P}$=$(-1.69\\,\\pm{0.04})\\,\\times10^{-8}\\,\\sps$, or $\\dot{\\nu}$=$(1.58\\,\\pm{0.04})\\,\\times10^{-12}\\,\\hzs$, is smaller than the spin-up measured during giant outbursts in the past, e.g., $\\dot{\\nu}$$\\sim$0.6$\\times10^{-11}\\,\\hzs$ in June 1983 \\citep{sembay90}, or $\\dot{\\nu}$$\\sim$$1.2\\times10^{-11}\\,\\hzs$ in February 1994 \\citep{finger96}.  The pulse profiles measured during the main outburst are consistent with those during a giant outburst in March/April 1989 albeit at a different luminosity level, $L_{(23-53\\kev)}\\,$$\\sim$$1.26\\,\\times10^{37}\\,\\ergs$ \\citep{kend94}. A comparison between these pulse profiles and those obtained for the pre-outburst flare in this paper suggest that a different mode of accretion takes place during the flare (see \\citealt{postnov08}).  The energy of the CRSF during the main outburst is consistent with the energy measured with \\inte close to the maximum of the main outburst, $E_{\\mathrm{cyc}}$=$45.9\\pm0.3\\,\\kev$ \\citep{caballero07} and with the energy measured with \\suzaku at the end of the main outburst, $E_{\\mathrm{cyc}}$=$46.3^{+1.5}_{-1.3}\\,\\kev$ (using the same Gaussian model for the line, \\citealt{terada06}). From the observed cyclotron energy, the estimated, magnetic field at the site of the X-ray emission during the main outburst, assuming a redshift of z=0.3, is $B$$\\sim$$5.2\\times10^{12}\\G$, and during the pre-outburst flare $B$$\\sim$$5.8\\times10^{12}\\G$. During the main part of the outburst the cyclotron-line energy remains constant within the errors, in spite of the changes in the luminosity. This suggests that the structure of the accretion column is different from that observed in both V~0332 +53 \\citep{mowlavi06} and 4U~0115+63 \\citep{tsygankov06}. The source is likely to be in the sub-Eddington regime, as for Her X-1 \\citep{staubert07} for which it is believed  no shock has formed in the accretion column.   Pre-outburst flares have been observed in other accreting pulsars, such as \\twos \\citep{finger99} and \\exo \\citep{camero05}.  \\citet{hayasaki06} have modeled accretion disks around neutron stars in Be/X-ray binaries. Their SPH simulations reproduce a series of normal outbursts from the accretion disk formed about the neutron star at periastron. In some cases, single flares are found preceding the outburst maximum. We note first of all that the timescale of the flare in their simulations appears to be too long (10\\% of the orbital period) compared to the flare duration observed. The inspection of the \\swift-\\bat light-curve of A0535+26 during the outburst studied here shows that RXTE actually observed only one of a collection of short flares in the rising part of the outburst \\citep{postnov08}. Such flares are not reproduced by the above simulations. This model does not take into account the disk-magnetospheric interaction.  A more plausible interpretation is that the pre-outburst flares could be caused by magnetospheric instabilities between the accretion disk and the neutron-star magnetosphere at the onset of accretion. The instability causes the plasma that has accumulated in the disk-magnetosphere, boundary-layer to rapidly fall onto the neutron star surface close to the magnetic poles, but along different field lines than (quasi)stationary accretion (see \\citealt{postnov08} for details)."
    },
    "0801/0801.2950_arXiv.txt": {
        "abstract": "We study the classical stability of an anisotropic space-time seeded by a spacelike, fixed norm, dynamical vector field in a vacuum-energy-dominated inflationary era. It serves as a model for breaking isotropy during the inflationary era. We find that, for a range of parameters, the linear differential equations for  small perturbations about the background do not have a growing mode.  We also examine the energy of fluctuations about this background in flat-space. If the kinetic terms for the vector field do not take the form of a field strength tensor squared then there is a negative energy mode and the background is unstable. For the case where the kinetic term is of the form of a field strength tensor squared we show that perturbations about the background have positive energy at lowest order. ",
        "introduction": "Inflation has become the standard paradigm for very early Universe cosmology. It supposes an era when the Universe was dominated by vacuum energy. During this era any classical  inhomogeneities are smoothed out by the exponential expansion  of the Universe. The inhomogeneities that we observe today originate in small quantum fluctuations that have their wavelength  redshifted  outside the horizon by the exponential expansion of the Universe.  At some time the Universe  exits the inflationary era and transitions to a radiation dominated era and then eventually to a matter dominated era.  After inflation, the Universe expands slower than the speed at which light travels.  This allows the perturbations with wavelengths that are outside of the horizon to eventually reenter the horizon and generate the inhomogeneities we observe in the microwave background (as well as the large scale structure of the Universe).   Since we have no direct probes of the inflationary era, we are compelled to consider the possibility that some tenets of physics---tenets that are fundamental to our current understanding of the Universe---only reflect post-inflationary developments. One such tenet is that  rotational invariance is not spontaneously (or explicitly) broken.  This symmetry implies angular momentum conservation and is necessary for the classification of elementary particles by their spin ({\\it i.e}., their angular momentum in their rest frame).  Recently the possibility that rotational invariance is broken during the inflationary era has been studied \\cite{Ackerman:2007, Gumrukcuoglu:2007bx, Gumrukcuoglu:2006xj, Pullen:2007tu, Ando:2007hc, Boehmer:2007ut}. In Ref.~\\cite{Ackerman:2007} it was noted that if the breaking of rotational invariance is small there is a very simple and unique signature on the spectrum of density perturbations and hence on the anisotropy of the microwave background radiation.  Ackerman  {\\it et.~al.}~also introduced a simple model for breaking rotational invariance during the inflationary era.  They calculated the influence of the breaking of rotational invariance on the spectrum of density perturbations, verifying the expectations based on their general arguments.   The Lagrange density for the model has the usual Einstein-Hilbert action, a cosmological constant term with vacuum energy density $\\rho_{\\Lambda}$ and the following  terms containing the  four-vector field $u^{\\sigma}$, \\begin{eqnarray} \\label{ulag} {\\cal{L}}_u &=&-\\beta_1 \\nabla^\\mu u^\\sigma \\nabla_\\mu u_\\sigma - \\beta_2 (\\nabla_\\mu u^\\mu)^2 \\nonumber \\\\  &-& \\beta_3 \\nabla^\\mu u^\\sigma \\nabla_\\sigma u_\\mu + \\lambda(u^{\\mu} u_{\\mu} - m^2)\\ . \\end{eqnarray} Here $\\lambda$ is a Lagrange multiplier that enforces the constraint $ g_{\\mu \\nu} u^{\\mu} u^{\\nu}=m^2$. We take $m^2>0$ so the four-vector $u^{\\sigma}$ is spacelike.  This model has a homogeneous but anisotropic background solution to the classical equations of motion. We choose the $x_3$-axis to be aligned along the  four-vector field, \\begin{equation} u^0=0,~~u^1 = u^2=0~~\\text{and}~~u^3 ={m \\over b(t)}, \\label{uvector} \\end{equation} resulting in a space-time metric of the form, \\begin{equation} \\label{coord} {\\rm d}s^2=-{\\rm d}t^2+a(t)^2{\\rm d}{\\bf x}_{\\perp}^2+b(t)^2{\\rm d}x_3^2. \\end{equation} The breaking of rotational invariance by the four-vector field causes the $x_3$-direction to expand at a different rate than the $x_1$ and $x_2$ directions.  Explicitly,\\footnote{It is assumed that the dynamics immediately preceding the inflationary era was rotationally invariant so that $a(0)=b(0)$.} \\begin{equation} a(t)=e^{H_a t},~~~b(t)=e^{H_b t}, \\label{spit1} \\end{equation} where the Hubble parameters are related to the vacuum energy, the scale of rotational invariance breaking, and parameters in the vector Lagrangian by the following equations, \\begin{eqnarray} \\label{backgrounda} &&H_a= {{\\dot a}\\over a} = H_b(1+16 \\pi G  \\beta_1 m^2),\\nonumber  \\\\ && H_b= {{\\dot b}\\over b} ={{\\sqrt{ 8 \\pi G \\rho_{\\Lambda} \\over (1+8 \\pi G  \\beta_1 m^2)(3+32 \\pi G  \\beta_1 m^2)}}} \\,. \\end{eqnarray} Notice that the background solution does not depend on the parameters $\\beta_{2,3}$ and as $m \\rightarrow 0$ the rate of expansion is the same for all spatial dimensions.   The purpose of this paper is to study the  classical stability of this solution.  We first consider the linear differential equations that result from expanding the classical equations of motion to linear order about the background solution discussed above. We find that, for small $G \\beta_i m^2$,  small perturbations do not grow provided that $\\beta_1 >0$ and $\\beta_1+\\beta_2 +\\beta_3 \\ge 0$.   Next we consider the energy of these solutions in flat-space.  Negative energy modes indicate an instability that might not show up in the analysis of the linearized equations of motion.  A special role is played by the case $\\beta_1+\\beta_2 +\\beta_3= 0$ where, in flat-space, the kinetic terms for the vector field take the form of a field strength tensor squared.  In this case, we find that when the energy density is evaluated to quadratic order in  fluctuations about the background, all propagating modes have positive energy. However, when that equality is not satisfied, there is a mode that propagates with negative energy.  Models with vector fields that spontaneously break Lorentz symmetry were first introduced in \\cite{Kostelecky:1989jw}.  Similar fixed-norm timelike Lorentz-violating vector field models have been extensively studied; constraints required by theoretical and observational consistency have been placed on parameters in this theory \\cite{Jacobson:2004ts, Carroll:2004ai, Lim:2004js, Eling:2003rd, Foster:2005dk, Li:2007vz, Seifert:2007fr, Foster:2006az, Eling:2006df, Jacobson:2000xp, Clayton:2001vy, Kostelecky:1989jw, Bluhm:2007bd}. For a review, see \\cite{Jacobson:2008aj}.  Others have classified gauge invariant perturbations in anisotropic scenarios  \\cite{Pereira:2007yy, PhysRevD.34.3570}; however, we choose to work in a particular (nonstandard) gauge for calculational convenience. ",
        "conclusions": "We have studied the small fluctuations about a spatially homogeneous anisotropic inflationary background. The anisotropy was caused by a dynamical four-vector that was constrained to have a spacelike invariant norm.  From the first order stability analysis of the equations of motion we derived the constraints $\\b_1 > 0$ and $\\b_1 + \\b_2 + \\b_3 \\ge 0$ for the parameters in the vector Lagrangian \\eqref{ulag}.  Moreover, we find that all modes have positive energy in flat-space and do not grow with time in the case where the kinetic term for the four-vector corresponds to a field strength tensor squared (\\textit{i.e.}~for $\\b_1 + \\b_2 +\\b_3 =0$).  A negative energy mode propagates if  $\\b_1 + \\b_2 +\\b_3 \\neq 0$, which implies that the background given by Eqs.~\\eqref{uvector}-\\eqref{backgrounda} is unstable when $\\b_1 + \\b_2 + \\b_3 \\ne 0$.   \\subsection*{Post-submission note}  After the posting of this paper, in Ref.~\\cite{Himmetoglu:2008zp}, an instability in the theory with a fixed-norm spacelike vector field with a field strength squared kinetic term was found  when the physical momentum of perturbations is comparable to the Hubble scale. We considered only physical momenta that were much greater than or much less than the Hubble scale.  \\appendix*"
    },
    "0801/0801.3459_arXiv.txt": {
        "abstract": "Analytic models based on spherical and ellipsoidal gravitational collapse have been used to derive the mass functions of dark matter halos and their progenitors (the {\\it conditional} mass function).  The ellipsoidal model generally provides a better match to simulation results, but there has been no simple analytic expression in this model for the conditional mass function that is accurate for small time steps, a limit that is important for generating halo merger trees and computing halo merger rates.  We remedy the situation by deriving accurate analytic formulae for the first-crossing distribution, the conditional mass function, and the halo merger rate in the ellipsoidal collapse model in the limit of small look-back times.  We show that our formulae provide a closer match to the Millennium simulation results than those in the spherical collapse model and the ellipsoidal model of Sheth \\& Tormen (2002). ",
        "introduction": "\\label{intro}  \\cite{PS74} presented an analytical expression for the unconditional mass function of dark matter halos at redshift $z$, $n(M,z)$, based on the spherical collapse model.  This function is closely related to the first-crossing distribution of random walks with a barrier in the excursion set framework \\citep{BCEK91}.  In this framework, the linear over-density computed at a given point in the Lagrangian space fluctuates as a Markovian process when smoothed on successively smaller scales, and a dark matter halo is identified at the point when the random walk of the linear over-density crosses a critical value, or a barrier $\\mathcal{B}(M,z)$.  In the spherical collapse model, this barrier depends only on time and is independent of mass: $\\mathcal{B}=\\delta_c/D(z)$, where $\\delta_c=1.68$ and $D(z)$ is the standard linear growth factor.  Although analytically simple, the spherical collapse model has been found to over-predict the abundance of small halos and under-predict that of massive ones (\\eg, \\citealt{LC94,GB94,T98,ST99}). The reason is mainly that halo collapses are generally triaxial rather than spherical (\\eg, \\citealt{d70,bbks86}).  Based on \\cite{BM96}, \\cite{SMT01} use the ellipsoidal collapse model and obtain fitting functions that provide a closer match to the unconditional halo mass function in N-body simulations. Unlike the spherical collapse model in which the condition for the virialization of a dark matter halo is solely determined by the linear over-density on the scale of the halo mass, the virialization condition in the ellipsoidal collapse model also depends on halo ellipticity and prolateness.  By assuming that a dark matter halo becomes virialized when its third axis collapses, \\cite{SMT01} find a new criterion for the virialization of dark matter halos, which involves all three parameters. They further simplify the virialization condition by fixing the ellipticity and the prolateness at their most likely values for a given over-density, and obtain a fitting formula for the barrier $\\mathcal{B}(M,z)$ that is mass-dependent, in contrast to the constant $\\mathcal{B}(z)$ of the spherical collapse model.  A mass-dependent barrier is commonly referred to as a {\\it moving} barrier, and it is this mass-dependence that suppresses the abundance of small halos while increasing that of massive ones in the ellipsoidal collapse model. Physically, this is because a smaller halo typically has a larger ellipticity and therefore a longer collapsing time.  The relationship between the unconditional mass function and the first crossing distribution associated with barrier-crossing random walks has been extended to obtain the {\\it conditional} mass function of halos \\citep{BCEK91, LC93}.  In this so-called extended Press-Schechter (EPS) model, the conditional mass function $dN(M_1,z_1|M_0, z_0)/dM_1$ gives the average number of progenitor halos (of mass $M_1$ at redshift $z_1$) per unit mass associated with a descendant halo of mass $M_0$ at redshift $z_0$ $\\left(z_1>z_0\\right)$.  Once determined, it can be used to generate merger trees of halos for many applications (e.g., galaxy formation, growth of the central black hole, reionization) using Monte Carlo simulations.  The conditional mass function has a simple analytic form in the constant barrier spherical collapse model \\citep{LC93}.  For a moving barrier (such as the ellipsoidal collapse model), however, exact analytic forms have been found only for the special case of a linear barrier (\\citealt{ST02}, ST02 hereafter); while the same authors have proposed a Taylor-series-like approximation for a general moving barrier.  We find that none of these formulae work well for $z_1-z_0 \\ll 1$, which is important for generating accurate merger trees in most Monte Carlo methods (\\eg, \\citealt{LC93,KW93,SK99,SL99,CLBF00}) and for relating halo merger rates to the conditional mass function (\\S~2 below).  \\cite{FM07} compare the halo merger rates determined from the Millennium simulation (\\citealt{springel05}) with the prediction of the spherical collapse EPS model, finding the latter to overpredict the major merger rates by up to a factor of $\\sim 2$ and underpredict the minor merger rates by up to a factor of $\\sim 5$.  Recently, various other fitting forms for the conditional mass function have been proposed that are calibrated to the results from particular $N$-body simulations, e.g., \\cite{Cole07, Parkinson07,ND07}.  An alternative way of deriving the conditional mass function that does not rely explicitly on fitting to simulations is to solve the integral equation proposed by \\cite{ZH06} (ZH06 hereafter). This method, which is based on the conservation of probability in the excursion set formalism, is accurate but computationally expensive.  In \\S2, we derive almost exact analytic forms for the first crossing distribution, the conditional mass function, and the halo merger rate in the ellipsoidal collapse model in the limit of small look-back times; a limit where earlier work breaks down.  Our method is based on ZH06, but our results are expressed in simple analytic forms.  We compare the predictions of this improved ellipsoidal collapse model with those from the spherical collapse model and the Millennium simulation in \\S\\ref{Millennium}.  We assume the cosmological parameters used in the Millennium simulation: $\\Omega_m=0.25$, $\\Omega_b=0.045$, $h=0.73$, $\\Omega_{\\Lambda}=0.75$, $n=1$, $\\sigma_8=0.9$. ",
        "conclusions": "\\label{summary}  We have derived new analytic formulae for the first-crossing distribution (eq.(\\ref{f_complete})), the conditional mass function (eq.(\\ref{NM})), and the merger rate of dark matter halos (eq.(\\ref{B0})) in the ellipsoidal collapse model in the limit of small look-back times.  Our method is based on solving the first-crossing distribution of random walks with a moving barrier using the exact integral equation of ZH06.  This method results in extra terms in eqs.~(\\ref{f_complete}), (\\ref{NM}), and (\\ref{B0}) that are absent in the spherical collapse model and different from those in the ellipsoidal model of ST02.  Fig.~1 illustrates how these terms correct the discrepancies of ST02 in the conditional mass function for small look-back times $\\Delta z\\la 0.03$.  Eq.~(\\ref{relation}) relates the conditional mass function at small $\\Delta z$ to the halo merger rate.  The halo merger rate of our ellipsoidal collapse model generally agrees better with the Millennium result reported in \\cite{FM07} than that of the spherical collapse model (red vs blue curves in Figs.~\\ref{mr1} and \\ref{mr2}).  The discrepancy between our ellipsoidal collapse model and the Millennium simulation is typically $20\\% \\sim 30\\%$, which is about a factor of two smaller than that of the spherical collapse model.  A comparison between Figs.~\\ref{mr1} and \\ref{mr2} shows a better agreement between model and simulation when the progenitor mass $M_i$ in the analytic formulae is assigned to be the smaller progenitor (option A).  A number of factors not considered in this paper nor in earlier EPS work can contribute to the remaining 20-30\\% discrepancy between the model and simulation.  These include the statistical importance of non-binary mergers, diffuse accretion and tidal stripping of halo mass (such that $M_0\\neq M_1+M_2$), and the impact of a halo's environment on merger rates that can lead to non-Markovian processes in the excursion set model (\\eg ST02, \\citealt{ND07}).  Our conditional mass function in eq.(\\ref{NM}) can be easily incorporated into Monte Carlo simulations to study halo merger histories over a large look-back time.  This is done in a companion paper, in which we compare several existing Monte Carlo algorithms (\\eg, \\citealt{LC93,KW93,SK99,CLBF00}) that are all based on the spherical collapse model, and propose a more accurate method using eq.(\\ref{NM}). Several groups have recently proposed accurately parameterized forms of the conditional mass function based on fits to the Millennium results and incorporated them into different Monte Carlo simulations (\\citealt{Cole07,Parkinson07,ND07}).  As we will show in the companion paper (also see \\citealt{MS07} for a different method), a similar level of accuracy can be achieved using eq.(\\ref{NM}) of this paper and our Monte Carlo method without a priori knowledge of N-body simulations.  We thank James Bullock, Joanne Cohn, Lam Hui, Ravi Sheth, Martin White, and Simon White for useful discussions.  This work is supported in part by NSF grant AST 0407351.  The Millennium Simulation databases used in this paper and the web application providing online access to them were constructed as part of the activities of the German Astrophysical Virtual Observatory."
    },
    "0801/0801.3173_arXiv.txt": {
        "abstract": "The bright galaxy population of the Local Group Analog (LGA) LGG~225 has been imaged with the Galaxy Evolution Explorer ({\\it GALEX}) through its Far- and Near-UV wavebands. A significant fraction  of the group  members   appear to underwent recent/on-going interaction episodes that strongly disturbed overall galaxy morphology. UV-bright regions, sites of intense star formation activity accompanied by intense dust extinction, mark the galaxy outskirts forming irregular structures and tails.  Compared to the Local Group, LGG~225 seems thus to  be experiencing a more intense and active evolutionary phase. ",
        "introduction": "\\label{sec:1}  Poor groups of galaxies represent the defining aggregates of the so called ``field''.  Their importance comes from two facts: {\\it i)} most of galaxies in the {\\em Local} Universe is found in groups, rather than in the cluster environment \\cite{Eke}. {\\it ii)} The transition between galaxy properties typical of field and clusters happens just {\\em at densities characteristic of poor groups}, thus suggesting the existence of re-processing mechanisms driven by the environment richness \\cite{Lewis, Gomez}. Groups show a wide range of properties when adopting a multi-wavelength approach.  As far as X-ray emission is concerned, for instance, those dominated by ellipticals are rich in hot and diffuse intergalactic gas \\cite{Mulch}, while in  spiral-dominated aggregates--- as is the case of our own Local Group--- a low X-ray emission comes from cold gas mainly confined  within fine structures (shells etc.) of single galactic sources (see \\cite{Schim}). Whether a link  exists between  these two different evolutionary outputs in a hierarchical cosmological framework is still matter of debate.  In the group environment we know that  galaxy encounters are extremely efficient in reshaping  member morphology, being stellar velocity dispersion  within galaxies comparable to that of the group as a whole. At the same time, both the merger/accretion rate and the frequency of processes giving  rise to tidal dwarfs are still unknown. In this context, we are currently studying with {\\it GALEX} a sample of Local Groups Analogs which includes so far three systems, namely LGG~93, LGG~127 and LGG~225.  We present here a preliminary analysis of the  LGG~225 group composed of about 10 galaxies (e.g. \\cite{Tul}). We imaged in the Near (NUV) and the Far (FUV) {\\it GALEX} bands the ultraviolet emission of NGC~3447, NGC~3447A,  NGC~3454, NGC~3455 and UGC~6035 aiming at tracing the ongoing  star formation (SF) and its distribution across galaxies and their  intergalactic medium. In order to estimate the paramenters of these SF events we rely on the \\cite{Buz1,Buz2} population synthesis models. ",
        "conclusions": ""
    },
    "0801/0801.1340_arXiv.txt": {
        "abstract": "%  We present recent results from several on-going studies: The first addresses the question of gas-density thresholds for star formation, as probed by the outer disks of normal nearby galaxies. The second concerns the observational evidence for the existence of gravitating non-luminous (GNL) galaxies, as predicted by most recent simulations of galaxy formation in Lambda-CDM cosmologies. We find that (1) If star formation is traced by far-ultraviolet light, then there is no evidence for a threshold to star formation at any gas density so far probed, and (2) there is no evidence for GNL galaxies gravitationally interacting with known optical systems based on the observations (a) that there are no ring galaxies without plausible optically visible intruders, (b) all peculiar galaxies in the Arp Atlas that are bodily distorted have nearby plausibly interacting companions, and (c) there are no convincingly distorted/peculiar galaxies within Karachentsev's sample of more than 1,000 apparently/optically isolated galaxies. ",
        "introduction": "%  Galaxies continue to surprise us, both in their observed structure and by their inferred evolution. Some of these surprises come from moving to new wavelengths and re-observing old friends; while other surprises come staying at familiar wavelengths but looking at the galaxies from a somewhat new perspective or in a slightly revised context.  With the launch of the GALEX satellite into low-Earth orbit on April 28, 2003 it has become possible to image a large number and a wide range of different types of galaxies at largely unexplored wavelengths: in the near and far ultraviolet. Given the high sensitivity of the detectors and the very low sky background (especially in the far ultraviolet channel) it is possible to see features in the ultraviolet out to surface brightness levels unequalled by ground-based optical observations. Moreover, the wide (one-degree) field of view of GALEX also maximizes the possibility of serendipitous discovers, of which there have been many. ",
        "conclusions": " (1) No ring galaxy is being produced from a head-on collision between a spiral and pure dark-matter GNL galaxy (2) No bodily deformed galaxies are the result of collisions and/or near encounters between optical and pure dark-matter GNL galaxies. (3) Optically isolated galaxies show no signs of bodily interactions with pure dark matter GNL galaxies.      \\vfill\\eject"
    },
    "0801/0801.4109_arXiv.txt": {
        "abstract": "#1{\\vskip 7mm \\begin{center}{\\large Abstract}\\par \\smallskip \\begin{minipage}[c]{12cm} \\small #1 \\end{minipage} \\end{center} } \\def\\title#1{\\begin{center}{\\Large\\bf #1}\\end{center}} \\def\\author#1{\\vskip 5mm \\begin{center}{#1}\\end{center}} \\def\\address#1{\\begin{center}{\\it #1}\\end{center}} \\makeatletter \\@ifundefined{lesssim}{\\def\\lesssim{\\mathrel{\\mathpalette\\vereq<}}}{} \\@ifundefined{gtrsim}{\\def\\gtrsim{\\mathrel{\\mathpalette\\vereq>}}}{} \\def\\vereq#1#2{\\lower3pt\\vbox{\\baselineskip1.5pt \\lineskip1.5pt \\ialign{$\\m@th#1\\hfill##\\hfil$\\crcr#2\\crcr\\sim\\crcr}}} \\makeatother  \\begin{document}  \\begin{flushright} IFT-UAM/CSIC-08-05 \\end{flushright}  \\vspace{5mm}  \\title{% A new gravitational wave background from the Big Bang } \\author{% Juan Garc\\'\\i a-Bellido\\footnote{E-mail: juan.garciabellido@uam.es}\\,\\, and\\,\\, Daniel G. Figueroa\\footnote{E-mail: daniel.figueroa@uam.es} } \\address{% Instituto de F\\'\\i sica Te\\'orica CSIC-UAM, \\\\Universidad Aut\\'onoma de Madrid, \\\\ Cantoblanco 28049 Madrid, Spain }  \\vspace*{-2mm}  \\abstract{\\vspace*{2mm} The reheating of the universe after hybrid inflation proceeds through the nucleation and subsequent collision of large concentrations of energy density in the form of bubble-like structures moving at relativistic speeds. This generates a significant fraction of energy in the form of a stochastic background of gravitational waves, whose time evolution is determined by the successive stages of reheating: First, tachyonic preheating makes the amplitude of gravity waves grow exponentially fast. Second, bubble collisions add a new burst of gravitational radiation. Third, turbulent motions finally sets the end of gravitational waves production. From then on, these waves propagate unimpeded to us. We find that the fraction of energy density today in these primordial gravitational waves could be significant for GUT-scale models of inflation, although well beyond the frequency range sensitivity of gravitational wave observatories like LIGO, LISA or BBO. However, low-scale models could still produce a detectable signal at frequencies accessible to BBO or DECIGO. For comparison, we have also computed the analogous % background from some chaotic inflation models and obtained similar results to those % of other groups. The discovery of such a background would open a new observational window into the very early universe, where the details of the process of reheating, i.e. the Big Bang, could be explored. Thus, it could also serve % as a new experimental tool for testing the Inflationary Paradigm. } ",
        "introduction": "Gravitational waves (GW) are ripples in space-time that travel at the speed of light, and whose emission by relativistic bodies represents a robust prediction of General Relativity. % Theoretically, it is expected that the present universe should be permeated by a diffuse background of GW of either an astrophysical or cosmological origin~\\cite{Maggiore}. % Fortunately, these backgrounds have very different spectral signatures that might, in the future, allow gravitational wave observatories like % LIGO~\\cite{LIGO}, LISA~\\cite{LISA}, % BBO~\\cite{BBO} or % DECIGO~\\cite{DECIGO}, to disentangle their origin~\\cite{Maggiore}. Unfortunately, the weakness of gravity will make this task extremely difficult, requiring a very high accuracy in order to distinguish one background from another. %  There are, indeed, a series of constraints on some of these backgrounds, coming % from the % anisotropies in the Cosmic Microwave Background (CMB)~\\cite{Elena}, from Big Bang nucleosynthesis~\\cite{BBN} % or from millisecond pulsar timing~\\cite{pulsar}. Most of these constraints come at very low frequencies, % from $10^{-18}$ Hz to $10^{-8}$ Hz, while present and future GW detectors (will) work at frequencies of order % $10^{-3}-10^3$ Hz. If early universe first order phase transitions~\\cite{KosowskyTurner,Nicolis} or cosmic turbulence~\\cite{Turbulence} occurred around the electro-weak (EW) scale, % GW detectors could have a chance to measure the corresponding associated backgrounds. However, if those processes occurred at the GUT scale, their corresponding backgrounds will go undetected by the actual detectors, since these cannot reach the required sensitivity in the high frequency range of $10^7-10^9$ Hz. There are however recent proposals to cover this range~\\cite{Nishizawa2007,Picasso}, which may become competitive in the not so far future.  Cosmological observations % seem to suggest that something like Inflation must have occurred in the very early universe. % Approximately scale-invariant density perturbations, sourced by quantum fluctuations during inflation, seem to be the most satisfying explanation for the CMB anisotropies. Together with such scalar perturbations one also expects tensor perturbations (GW) to be produced, with an almost scale-free power spectrum~\\cite{Starobinsky}. % The detection of such a background is crucial for early universe cosmology because it would help to determine the absolute energy scale of inflation, a quantity that for the moment is still uncertain, and would open the exploration of physics at very high energies.  However, in the early universe, after inflation, other GWB could have been produced at shorter wavelengths, in a more 'classical' manner rather than sourced by quantum fluctuations. In particular, whenever there are large and fast moving inhomogeneities in a matter distribution, one expects the emission of GW. % At large distances from a source, % the amplitude of the GW is given by % $h_{ij} \\simeq G\\ddot Q_{ij} / c^4r$, with $Q_{ij}$ the quadrupole moment of the mass distribution. % The larger the velocity of the matter distribution, the larger the amplitude of the radiation produced. % However, because of the weakness of gravity, % in order to produce a significant amount of gravitational radiation, it is required % a very relativistic motion (and high density contrasts) in the matter distribution of a source. Fortunatelly, this is indeed believed to be the situation at the end of inflaton, during the conversion of the huge energy density driving inflation into radiation and matter, at the so-called {\\it reheating} of the Universe~\\cite{preheating}, i.e. at the Big Bang.  Note that any background of GW coming from the early universe, if generated below Planck scale, immediately decoupled upon production and, whatever their spectral signatures, they will retain their shape throughout the expansion of the Universe. Thus, the characteristic frequency and shape of the GWB generated at a given time should contain information about the very early state of the Universe in which it was produced. Actually, it is conceivable that, in the not so far future, the detection of these GW backgrounds could be the only way we may have to infer the physical conditions of the Universe at such high energy scales. % However, the same reason that makes GW ideal probes of the early universe $-$ the weakness of gravity $-$ is responsible for the extreme difficulties we have for their detection on Earth. %  In % Refs.~\\cite{GBF,GBFS} we described the stochastic background predicted to arise from reheating after hybrid inflation. Here we will review the various processes involved in the production of such a background. In the future, this background could serve as a new tool to discriminate among different inflationary models, since reheating in each model would give rise to a different GWB with very characteristic spectral features. The details of the dynamics of preheating depend very much on the model and are often very complicated because of the non-linear, non-perturbative and out-of-equilibrium character of the process itself. However, all the cases have in common that only specific resonance bands of the fields suffer an exponential instability, which makes their occupation numbers grow by many orders of magnitude. The shape and size of the spectral bands depend very much on the inflationary model. If one translates this picture into position-space, the highly populated modes correspond to large time-dependent inhomogeneities in the matter distributions which acts, in fact, as a powerful source of GW. For example, in single field chaotic inflation models, the coherent oscillations of the inflaton during preheating generates, via parametric resonance, a population of highly occupied modes that behave like waves of matter. They collide among themselves and their scattering leads to homogenization and local thermal equilibrium. These collisions occur in a highly relativistic and very asymmetric way, being responsible for the generation of a stochastic GWB~\\cite{TkachevGW,JuanGW,EastherLim,EastherLim2,DufauxGW} with a typical frequency today of the order of $10^{7} - 10^{9}$ Hz, corresponding to the present size of the causal horizon at the end of high-scale inflation.  However, there are models like hybrid inflation in which the end of inflation is sudden~\\cite{hybrid} and the conversion into radiation occurs almost instantaneously. Indeed, % hybrid models preheat very violently, via the spinodal instability of the symmetry breaking field that triggers the end of inflation, irrespective of the couplings that this field may have to the rest of matter. Such a process is known as \\textit{tachyonic preheating}~\\cite{tachyonic,symmbreak} and could be responsible for copious production of dark matter particles~\\cite{ester}, lepto and baryogenesis~\\cite{CEWB}, topological defects~\\cite{tachyonic}, primordial magnetic fields~\\cite{magnetic}, etc. In Ref.~\\cite{symmbreak}, it was shown that the process of symmetry breaking in hybrid preheating, proceeds via the nucleation of dense bubble-like structures moving at relativistic speeds, which collide and break up into smaller structures (see Figs.~7 and~8 of Ref.~\\cite{symmbreak}). We conjectured at that time that such collisions would be a very strong source of GW, analogous to the GW production associated with strongly first order phase transitions~\\cite{KosowskyTurner}. As we will show here, this is indeed the case during the nucleation, collision and subsequent rescattering of the initial bubble-like structures produced after hybrid inflation. During the different stages of reheating in this model, gravity waves are generated and amplified until the Universe finally thermalizes and enters into the radiation era of the Standard Model of Cosmology. From that moment until now, % this cosmic GWB will be redshifted as a radiation-like fluid, totally decoupled from any other energy-matter content of the universe, such that today's ratio of energy stored in these GW to that in radiation, could range from $\\Omega_{_{\\rm GW}}h^2\\sim 10^{-8}$, peacked around $f\\sim 10^7$~Hz for the high-scale models, to $\\Omega_{_{\\rm GW}}h^2\\sim 10^{-11}$, peacked around $f\\sim 1$ Hz for the low-scale models.  Finally, since the first paper by Khlebnikov and Tkachev~\\cite{TkachevGW}, studing the GWB produced at reheating after chaotic inflation, there has been some developments. The idea was soon extended to hybrid inflation in Ref.~\\cite{JuanGW}. It was also revisited very recently in Ref.~\\cite{EastherLim,EastherLim2} for the $\\lambda\\phi^4$ and $m^2\\phi^2$ chaotic scenarios, and reanalysed again for hybrid inflation in Refs.~\\cite{GBF,GBFS}, using the new formalism of tachyonic preheating~\\cite{tachyonic,symmbreak}. Because of the increase in computer power of the last few years, we are now able to perform precise simulations of the reheating process in a reasonable time scale. Moreover, understanding of reheating has improved, while gravitational waves detectors are beginning to attain the aimed sensitivity~\\cite{LIGO}. Furthermore, since these cosmic GWBs could serve as a deep probe into the very early universe, we should characterize in the most detailed way the information that we will be able to extract from them. ",
        "conclusions": "To summarize, we have shown that hybrid models are very efficient generators of gravitational waves at preheating, in three well defined stages, first via the tachyonic growth of Higgs modes, whose gradients act as sources of gravity waves; then via the collisions of highly relativistic bubble-like structures with large amounts of energy density, and finally via the turbulent regime (although this effect does not seem to be very significant in the presence of scalar sources), which drives the system towards thermalization. These waves remain decoupled since the moment of their production, and thus the predicted amplitude and shape of the gravitational wave spectrum today can be used as a probe of the reheating period in the very early universe. The characteristic spectrum can be used to distinguish between this stochastic background and others, like those arising from NS-NS and BH-BH coalescence, which are decreasing with frequency, or those arising from inflation, that are flat~\\cite{SKC}. \\\\  We have plotted in Fig.~\\ref{fig5} the sensitivity of planned GW interferometers like LIGO, LISA and BBO, together with the present bounds from CMB anisotropies (GUT inflation), from Big Bang Nucleosynthesis (BBN) and from milisecond pulsars (ms pulsar).  Also shown are the expected stochastic backgrounds of chaotic inflation models like $\\lambda\\phi^4$, both coupled and pure, as well as the predicted background from two different hybrid inflation models, a high-scale model, with $v=10^{-3}M_P$ and $\\lambda\\sim g^2\\sim0.1$, and a low-scale model, with $v=10^{-5}M_P$ and $\\lambda\\sim g^2\\sim 10^{-14}$, corresponding to a rate of expansion $H\\sim 100$ GeV. The high-scale hybrid model produces typically as much gravitational waves from preheating as the chaotic inflation models. The advantage of low-scale hybrid models of inflation is that the background produced is within reach of future GW detectors like BBO~\\cite{BBO}. It is speculated that future high frequency laser interferometers could be sensitive to a GWB in the MHz region~\\cite{Nishizawa2007}, although they are still far from the bound marked with an interrogation sign.  For a high-scale model of inflation, we may never see the predicted GW background coming from preheating, in spite of its large amplitude, because it appears at very high frequencies, where no detector has yet shown to be sufficiently sensitive, unless the spectrum can be extrapolated to lower frequencies, where there are interferometric detectors like BBO which could see a signal. On the other hand, if inflation occured at low scales, even though we will never have a chance to detect the GW produced during inflation in the polarization anisotropies of the CMB, we do expect gravitational waves from preheating to contribute with an important background in sensitive detectors like BBO. The detection and characterization of such a GW background, coming from the complicated and mostly unknown epoch of rehating of the universe, may open a new window into the very early universe, while providing a new test on inflationary cosmology. \\\\"
    },
    "0801/0801.3729_arXiv.txt": {
        "abstract": "We report Cousins $R$-band monitoring of the high-redshift ($z$=4.40) radio quiet quasar Q\\,2203+292 from May 1999 to October 2007. The quasar shows maximum peak-to-peak light curve amplitude of $\\sim$0.3 mag during the time of our monitoring, and $\\sim$0.9 mag when combined with older literature data. The rms of a fit to the light curve with a constant is 0.08 mag and 0.2 mag, respectively. The detected changes are at $\\sim$3-sigma level. The quasar was in a stable state during the recent years and it might have undergone a brightening event in the past. The structure function analysis concluded that the object shows variability properties similar to those of the lower redshift quasars. We set a lower limit to the Q\\,2203+292 broad line region mass of 0.3-0.4 M$_\\odot$. Narrow-band imaging search for redshifted Ly$\\alpha$ from other emission line objects at the same redshift shows no emission line objects in the quasar vicinity. ",
        "introduction": "Many quasars show short-term or/and long-term variability (Ulrich, Maraschi \\& Urry 1997). These changes in the source flux help to constrain the physics and the size of the central engine. Astronomers came to an early understanding that the central engines of the AGNs and QSOs can not be resolved easily, if at all, but that the variability timescales measure the size of the emitting regions (i.e. Blandford \\& McKee 1982). Later on, this became the basis of reverberation studies (for a recent review see Peterson et al. 2004 and the references therein) that could map the innermost broad line regions (hereafter BLR). Nevertheless, the exact variability mechanisms remain unclear.  Typically, the variability is aperiodic, but it does show some dependencies on time lag, luminosity, wavelengths and redshift. For example, the variability amplitude increases with time lag (Vanden Berk et al. 2004), more luminous quasars are less variable (Vanden Berk et al. 2004; de Vries et al. 2005; Giveon et al. 1999), and the variability increases toward the blue part of the spectrum (Vanden Berk et al. 2004; de Vries et al. 2005). It is particularly difficult to constrain the variability versus redshift dependence because of the inevitable biases at high redshift, limiting the quasar luminosity range, the probed time-baseline, etc. of the more distant quasars.  The radio properties of QSOs also seem to be related to the optical variability: the radio-loud ones are relatively more variable that the radio quiet ones, and the blazars show even stronger variability because of beaming, which is quite different than the long-term variability studied in this work. For a more comprehensive review of the QSO variability properties we refer the reader to the summary of Wold, Brotherton \\& Shang (2007). As of now, there is no commonly accepted theory of the QSO variability and the existence of some strange objects, albeit rare ones, complicates the picture even further. A good example is the radio-quiet QSO SDSS\\,J153259.96-003944.1 (Stalin \\& Srianand 2005) who is the prototype of the rare class of weak (or absent) emission-line quasars (WLQs). This object shows strong long-term variability and flat optical spectrum like BL\\,Lacs but there is no radio emission, no optical polarization and no X-ray (Shemmer et al. 2006). These properties may be explained by a deficit of line-emitting gas in the vicinity of the central continuum source, similar to the X-ray weak quasar PHL\\,1811 (Leighly et al. 2007 and the references therein).  The most comprehensive quasar variability studies, especially the ones covering long time scales, are focused either on bright low-redshift ones (Hook et al. 1994; Hawkins 2002), or they study the behavior of the structure function for large samples (de Vries, Becker \\& White 2003; de Vries et al. 2005; Hovatta et al. 2007). The literature is lacking well sampled light curves of individual distant quasars, with the notable exception of Kaspi et al. (2007) with whom we share five objects from our extended program. While challenging, the distant QSOs are also rewarding because the variability can help constrain the size of the QSO accretion disks (Hawkins 2007) at early times. Furthermore, identifying variable high-redshift quasars is a necessary first step for reverberation mapping of their broad line regions (i.e. Kaspi et al. 2007). These arguments motivated us to start an optical monitoring study of QSOs with $z$$\\geq$4.  Here we present a photometric sequence for Q\\,2203+292, which is one of the first detected quasars within that redshift range (Dickinson \\& McCarthy 1987) and it is radio-quiet (Schneider et al. 1992; Schmidt et al. 1995; Omont et al. 1996). Surdej et al. (1993), Crampton, McClure \\& Fletcher (1992) and Kochanek (1993) reported that it is not gravitationally lensed. The only study in the literature on the Q\\,2203+292 variability comes from McCarthy et al. (1988) who concluded from three Lick 3m plates and two plates from the 5m Hale telescope (Longair \\& Gunn 1975; Riley, Longair \\& Gunn 1980) that the quasar did not vary strongly over a period of the 15 years since its discovery. A few additional broadband observations of Q\\,2203+292 are reported elsewhere, but they are in different photometric systems (see Table~\\ref{Table_Literature}).  Turner (1991) calculated for Q\\,2203+292 a mass of the central black hole of (0.69-9.6)$\\times$10$^{8}$ M$_\\odot$ and a mass accretion rate of 1.6-22 M$_\\odot$/yr, depending on the adopted cosmological model. This values are typical for the quasars at the same redshift (Turner 1991; Dietrich \\& Hamann 2004).  \\begin{table} \\caption{Broad band photometry of Q\\,2203+292 since 1987.} \\label{Table_Literature} \\begin{tabular}{c@{  }c@{ }c@{  }c} \\hline Date       & Reference                   & Filter        & Brightness     \\\\ yyyy/mm/dd &                             &               & mag            \\\\ \\hline 1987/09/25 & McCarthy et al. (1988) & $r_{\\rmn{s}}$ & 20.78$\\pm$0.09 \\\\ &                             & $V$           &  22.0$\\pm$0.3  \\\\ &                             & $I$           &  21.1$\\pm$0.3  \\\\ &                             &               &                \\\\ 1988/09/12 & Schneider, Schmidt \\& Gunn   & $r_{\\rmn{4}}$ & 20.88$\\pm$0.05 \\\\ &         (1989)              & $g_{\\rmn{4}}$ & 22.33$\\pm$0.06 \\\\ &                             &               &                \\\\ 1989/09/09 & Crampton et al. (1992)      & $R_C$           &  20.7$\\pm$0.1  \\\\ &                             & $V$           &  22.0$\\pm$0.1  \\\\ &                             &               &                \\\\ 1990/06/21 &    Crampton                 & $R_C$           &  20.4$\\pm$0.1  \\\\ &(private communication)      &               &                \\\\  \\hline \\end{tabular} \\end{table}  \\section[]{Observations and Data Reduction}  \\subsection[]{Observing Strategy and Basic Data Reduction}  We monitored Q\\,2203+292 in $R$-band with a variety of instruments and telescopes. ESO Archive images were also used. The observing log is shown in Table~\\ref{Table_ObsLog}. Typically, the total integration time was split into a few separate frames (as listed in the last column) and the telescope was jittered by a few arcsec between each of them to remove the artifacts caused by the detector's cosmetic defects. All observations were performed in clear, photometric nights. The object was monitored during culmination, whenever possible, to minimize the airmass variation during the observations.  \\begin{table*} \\begin{center} \\begin{minipage}{150mm} \\caption[]{Observing log.} \\label{Table_ObsLog} \\begin{tabular}{c@{   }c@{   }c@{   }c@{   }c@{   }c@{   }c@{   }c@{   }c} \\hline Data & Instrument@Telescope@Site & Pixel Scale & FoV & Airmass & FWHM & Total Integration Time \\\\ yyyy/mm/dd & & [arcsec\\,px$^{-1}$] & [arcmin$^2$] & &[arcsec]&[s]\\\\ \\hline 1999/05/15 & FORS1@VLT@Paranal              & 0.25 & 6.8$\\times$6.8 & 1.76 & 0.4 & 1$\\times$100=100 \\\\ 2003/08/26 & Photometrics\\,AT200A@2m@Rozhen & 0.29 & 5.0$\\times$5.0 & 1.04 & 1.1 & 3$\\times$1200=3600 \\\\ 2004/10/09 & Photometrics\\,AT200A@2m@Rozhen & 0.29 & 5.0$\\times$5.0 & 1.03 & 1.4 & 3$\\times$1200=3600 \\\\ 2005/10/01 & EEV/Marconi\\,42-40@1m@SAO      & 0.27 & 4.6$\\times$4.6 & 1.23 & 1.2 & 3$\\times$200=600 \\\\ 2005/11/04 & VersArray\\,1300B@2m@Rozhen     & 0.26 & 5.7$\\times$5.5 & 1.11 & 1.8 & 9$\\times$300=2700 \\\\ 2005/11/06 & VersArray\\,1300B@2m@Rozhen     & 0.26 & 5.7$\\times$5.5 & 1.07 & 2.0 & 6$\\times$600=3600 \\\\ 2006/06/29 & Photometrics\\,AT200A@2m@Rozhen & 0.29 & 5.0$\\times$5.0 & 1.54 & 1.6 & 10$\\times$180+2$\\times$300=2400 \\\\ 2006/08/18 & Photometrics\\,AT200A@2m@Rozhen & 0.29 & 5.0$\\times$5.0 & 1.04 & 1.4 & 3$\\times$1200=3600 \\\\ 2006/08/19 & Photometrics\\,AT200A@2m@Rozhen & 0.29 & 5.0$\\times$5.0 & 1.02 & 1.4 & 2$\\times$1200=2400 \\\\ 2006/08/25 & FoReRo2@2m@Rozhen              & 0.82 & 7.0$\\times$7.0 & 1.32 & 2.4 & 8$\\times$300=2400 \\\\ 2006/10/23 & EEV/Marconi\\,42-40@1m@SAO      & 0.27 & 4.6$\\times$4.6 & 1.04 & 1.1 & 3$\\times$300=900 \\\\ 2006/10/24 & EEV/Marconi\\,42-40@1m@SAO      & 0.27 & 4.6$\\times$4.6 & 1.09 & 1.0 & 3$\\times$300=900 \\\\ 2007/09/10 & VersArray\\,1300B@2m@Rozhen     & 0.26 & 5.7$\\times$5.5 & 1.14 & 1.5 & 2$\\times$1200=2400 \\\\ 2007/10/04 & VersArray\\,1300B@2m@Rozhen     & 0.26 & 5.7$\\times$5.5 & 1.04 & 1.0 & 3$\\times$1200=3600 \\\\ \\hline \\end{tabular} \\end{minipage} \\end{center} \\end{table*}  \\begin{figure} \\hspace*{0.0cm}\\psfig{file=Fig_1.eps,width=8.3cm,height=7cm} \\caption{Radial flux variation. The difference $R_{Stetson} - R_{inst}$ magnitudes for all detectable stars in the field of the globular clusters NGC\\,7790 and NGC\\,2420, taken with VersArray\\,1300B in direct focus (top) and VersArray\\,512B in red channel of FoReRo2 (bottom) versus the squared normalized distance from the center of the detector $\\rho^2$.} \\label{fig1} \\end{figure}  \\begin{figure} \\hspace*{0cm}\\psfig{file=Fig_2.eps,width=8.0cm,height=9.3cm} \\caption{Difference between the magnitudes of the quasar Q\\,2203$+$292 and the reference stars A and B is shown on the top and middle panels. The difference between A and B reference stars is shown on the bottom panel. The solid line is the linear fit to the data. The first open circle indicates the VLT data and the other open circles are Rozhen Observatory data. The solid circles are the data from the 1m SAO telescope.} \\label{fig2} \\end{figure}  \\begin{figure*} \\hspace*{0.5cm}\\psfig{file=Fig_3.eps,width=13.5cm,height=13.5cm} \\caption{Identification chart with standard stars in the field of Q\\,2203+292. The image was obtained with 2m RCC telescope of NAO, Rozhen on 2004 October 9. The field of view is 5$\\times$5 arcmin$^2$. North is up, and East is to the left.} \\label{fig3} \\end{figure*}   The basic data reduction includes: bias substraction, flat fielding, alignment of individual frames and combination. We used the standard IRAF\\footnote{IRAF is the Image Reduction Analysis and Facility made available to the astronomical community by the National Optical Astronomy Observatories, which are operated by AURA, Inc., under contract with the U.S. National Science Foundation. STSDAS is distributed by the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy (AURA), Inc., under NASA contract NAS5-26555.} routines to perform them.  \\subsection[]{Instrumental Magnitudes and Corrections for Systematical Effects}  We carried out aperture photometry on the combined images. The aperture diameter was set to the size of the FWHM to optimize the signal-to-noise, because the sky contributed comparable flux to the target flux in the wings of the image. This is possible because: (i) the QSO PSF is indistinguishable from the PSF profiles of the stars, and (ii) the PSF variations across the field of view are negligible. Therefore, the selection of the aperture does not have an effect on the relative photometry we obtained on each individual image. Next, the zero points we determined by comparing the instrumental and the standard magnitudes of the calibration field (measured the same way) contain into themselves the aperture corrections. For more details on the photometric calibration see Sec.~\\ref{Sec_Standards}.  The Photometrics\\,AT200A and the VersArray\\,1300B cameras exhibit spatial flux variations (Fig.~\\ref{fig1}). They were removed as described in Markov (2005), correcting the instrumental magnitudes as follows: \\begin{equation} R = c_{\\rho} \\rho^2 + R_{inst}, \\end{equation} where $c_{\\rho}$ is a known coefficient (0.12 for Photometrics\\,AT200A and 0.17 for VersArray\\,1300B) and $\\rho$ is the distance from the center of the detector. This is normalized by the detector half-size, and it varies from $0$ to $\\sqrt{2}$: \\begin{equation} \\rho = \\sqrt{ \\left(1-\\frac{x}{x_c}\\right)^2 +  \\left(1-\\frac{y}{y_c}\\right)^2}. \\end{equation} Here $x$ and $y$ are the coordinates of the object in pixels and $x_c$ and $y_c$ are the coordinates of the central pixel. Note that for the purpose of our differential photometry this correction is minor, because the quasar and the comparison stars were placed at the same position, within the pointing errors, so $\\Delta \\rho^2$$<$0.1, translates into 0.02 mag extra uncertainty in the differential magnitudes Q--A and Q--B (see Fig. 2). An additional source of error, albeit also small, is the jittering between the individual frames, that comprise the individual epochs of our light curve. It was always 3-4 arcsec (except 12 arcsec in one case) which corresponds to ignorable magnitude error. Interestingly, the FoReRo2 (Jockers et al. 2000) shows no spatial effects, as verified from observations of Stetson standards (Stetson 2000) taken during a few different photometric nights.  Differential light curves of the quasar were generated relative to two nearby comparison stars imaged on the same frame. They were selected among the most stable stars in the field (see Sec.~\\ref{Sec_Standards}). The results of the differential photometry are given in Table~\\ref{Table_LightCurves} and Fig.~\\ref{fig2}. The points from SAO (solid circles) are with bigger errors, because of the smaller telescope aperture and the shorter exposure times (Table~\\ref{Table_ObsLog}).  \\begin{table} \\begin{center} \\caption{Differential instrumental magnitudes between the quasar and the comparison stars A ($\\alpha_{2000}$=22:06:07.15, $\\delta_{2000}$=29:30:12.9) and B ($\\alpha_{2000}$=22:06:06.33, $\\delta_{2000}$=29:31:08.0) and between the two comparison stars.} \\label{Table_LightCurves} \\begin{tabular}{cccc} \\hline Date       &     Q$-$A\t   & Q$-$B\t   &\t A$-$B      \\\\ yyyy/mm/dd &     [mag]\t   & [mag]\t   &\t [mag]      \\\\ \\hline 1999/05/15 & 0.48$\\pm$0.05 & 0.97$\\pm$0.04 & 0.49$\\pm$0.03 \\\\ 2003/08/26 & 0.74$\\pm$0.03 & 1.20$\\pm$0.03 & 0.45$\\pm$0.03 \\\\ 2004/10/09 & 0.82$\\pm$0.03 & 1.28$\\pm$0.03 & 0.46$\\pm$0.03 \\\\ 2005/10/01 & 0.76$\\pm$0.06 & 1.18$\\pm$0.06 & 0.42$\\pm$0.03 \\\\ 2005/11/04 & 0.67$\\pm$0.04 & 1.14$\\pm$0.03 & 0.47$\\pm$0.03 \\\\ 2005/11/06 & 0.70$\\pm$0.03 & 1.19$\\pm$0.03 & 0.48$\\pm$0.03 \\\\ 2006/06/29 & 0.73$\\pm$0.04 & 1.21$\\pm$0.04 & 0.48$\\pm$0.03 \\\\ 2006/08/18 & 0.77$\\pm$0.03 & 1.23$\\pm$0.03 & 0.46$\\pm$0.03 \\\\ 2006/08/19 & 0.71$\\pm$0.03 & 1.18$\\pm$0.03 & 0.47$\\pm$0.03 \\\\ 2006/08/25 & 0.61$\\pm$0.04 & 1.06$\\pm$0.04 & 0.45$\\pm$0.03 \\\\ 2006/10/23 & 0.72$\\pm$0.06 & 1.14$\\pm$0.05 & 0.42$\\pm$0.04 \\\\ 2006/10/24 & 0.71$\\pm$0.06 & 1.27$\\pm$0.06 & 0.56$\\pm$0.04 \\\\ 2007/09/10 & 0.63$\\pm$0.03 & 1.07$\\pm$0.03 & 0.44$\\pm$0.02 \\\\ 2007/10/04 & 0.65$\\pm$0.02 & 1.10$\\pm$0.02 & 0.45$\\pm$0.02 \\\\ \\hline \\end{tabular} \\end{center} \\end{table}    \\subsection[]{Photometric Calibration\\label{Sec_Standards}}  The absolute calibration of our instrumental magnitudes includes three steps. First, we tied the FORS1@VLT image to the Landolt (1992) standard field MARK\\,A. Then, we searched for all stars in common between the FORS1@VLT image and eight other images from NAO Rozhen, obtained under photometric conditions until 2006 August and calculated transformations between the individual frames. Finally, using these nine images we derived new magnitudes for a few additional stars, bringing the number of reference stars in the field to 24. Here we consider only stars with rms$\\leq$0.04 mag making sure the calibration is based only on non-variable sources (Table~\\ref{Table_Stdandards} and Fig.~\\ref{fig3}). The weights used for averaging the magnitudes were $\\sigma^{-2}$, where: \\begin{equation} \\sigma = \\sqrt{\\sigma_{zp}^{2} + \\sigma_{inst}^{2}}. \\end{equation} Here $\\sigma_{zp}$ is the error of the zero-point and $\\sigma_{inst}$ is the instrumental magnitude's error. The total uncertainty is dominated by the zero-point errors with a typical value to $\\sim$0.03 mag. The instrumental error of individual measurements attains 0.01 mag at $R_C$$\\sim$20.5 mag level (Table~\\ref{Table_Stdandards}).  \\begin{table} \\caption{Reference stars in the field of the quasar Q\\,2203+292. The quasar magnitude is for the 2004 October 9 image.} \\label{Table_Stdandards} \\begin{center} \\begin{tabular}{c@{  }c@{  }c@{   }c} \\hline ID & \\multicolumn{2}{c}{RA (2000.0) DEC} & R$_C$(err), mag \\\\ \\hline Q & 22:06:02.70 & 29:30:02.0 & 20.51\\,(0.03) \\\\ A & 22:06:07.15 & 29:30:12.9 & 19.69\\,(0.01) \\\\ B & 22:06:06.33 & 29:31:08.0 & 19.23\\,(0.02) \\\\ 1 & 22:06:12.17 & 29:28:26.0 & 20.29\\,(0.03) \\\\ 2 & 22:06:13.47 & 29:28:45.3 & 19.43\\,(0.03) \\\\ 3 & 22:06:00.42 & 29:29:13.7 & 18.58\\,(0.03) \\\\ 4 & 22:06:11.51 & 29:29:17.9 & 19.40\\,(0.04) \\\\ 5 & 22:05:54.95 & 29:29:22.6 & 20.23\\,(0.03) \\\\ 6 & 22:06:08.80 & 29:29:24.1 & 20.03\\,(0.03) \\\\ 7 & 22:05:58.62 & 29:29:26.2 & 18.60\\,(0.02) \\\\ 8 & 22:06:12.79 & 29:29:40.8 & 19.03\\,(0.02) \\\\ 9 & 22:05:55.39 & 29:29:42.9 & 17.72\\,(0.03) \\\\ 10 & 22:05:54.13 & 29:29:50.9 & 20.26\\,(0.03) \\\\ 11 & 22:05:59.95 & 29:29:51.6 & 19.96\\,(0.03) \\\\ 12 & 22:05:55.51 & 29:29:55.0 & 19.99\\,(0.04) \\\\ 13 & 22:05:54.92 & 29:29:58.5 & 19.88\\,(0.03) \\\\ 14 & 22:05:59.73 & 29:30:00.8 & 18.61\\,(0.03) \\\\ 15 & 22:06:02.24 & 29:30:01.1 & 20.45\\,(0.03) \\\\ 16 & 22:06:11.07 & 29:30:04.1 & 19.67\\,(0.03) \\\\ 17 & 22:06:05.93 & 29:30:29.0 & 19.91\\,(0.03) \\\\ 18 & 22:05:54.04 & 29:31:13.9 & 19.19\\,(0.03) \\\\ 19 & 22:05:59.18 & 29:31:19.4 & 20.18\\,(0.03) \\\\ 20 & 22:05:58.68 & 29:31:22.7 & 18.80\\,(0.02) \\\\ 21 & 22:05:56.55 & 29:31:23.1 & 19.08\\,(0.02) \\\\ 22 & 22:06:02.73 & 29:31:37.0 & 19.24\\,(0.02) \\\\ \\hline \\end{tabular} \\end{center} \\end{table}   \\subsection{Literature data and filter transformation\\label{Sec_MagTransf}}  The observations of McCarthy et al. (1988), Schneider et al. (1989) and (Crampton et al. 1992; private communication) provided additional measurements of Q\\,2203$+$292 (Table~\\ref{Table_Literature}). Since some of them were observed in filters other than the $R_C$ filter used by us and we have to compare the luminosity of Q\\,2203+292 with the luminosities of quasars at similar redshift observed in $r^\\prime$, we were forced to derive transformations to the Cousins $R_C$ system.  We calculated $R_C$$-$$r$ for a $z$=4.4 QSO using the Q\\,2203+292 spectrum from Constantin et al. (2002) convolving it with the transmission curves of the standard $R_C$ filter and the other $r$ filters used in the literature studies. To obtain the necessary spectral coverage, we combined it with the bluest part of the Q\\,2203+292 spectrum from McCarthy et al. (1988), shortwards of 1070 $\\AA$. Here and throughout this paper we used the following zero-point fluxes: 3080 Jy for $R_C$ (Bessell 1979), 2810 Jy for $r_s$ system (Djorgovski 1985), 4471 Jy for $r_4$ (Frei \\& Gunn 1994) and 3631 Jy for $r^\\prime$ AB system (Oke \\& Gunn 1983).  We also calculated $R_C$$-$$r$ as a function of redshift $z$ to compare our corrections with the literature, and to derive K-corrections for a comparison of the properties of Q\\,2203+292 and the SDSS QSOs (see Sec.~\\ref{SecPhysProp}). We created our own composite spectrum, combining the Constantin et al. (2002) median spectrum with the bluest part of the McCarthy et al. (1988) shortwards of 1150 $\\AA$ (which contributes very little flux to any of the $r$ filters). Finally, the reddest part of our template ($\\lambda$$>$1450 $\\AA$) comes from the Vanden Berk et al. (2001). Indeed, the composite is dominated by lower redshift ($z$$\\leq$2) quasars. The overall behavior of $R_C$$-$$r$ as a function of $z$ is shown in Fig.~\\ref{fig4}. The bottom panel demonstrates the agreement between our synthetic photometry and the mean $R_C$$-$$r^\\prime$ colour for the quasars with $r^\\prime$ photometry from the SDSS QSO sample (Schneider et al. 2005) after applying the colour equations of Jester et al. (2005; see their table 1). Although these transformations were derived for quasars with $z$$<$2.1, the figure suggests that they can be extrapolated even up to $z$$\\sim$3.3. For higher redshifts, Ly$\\alpha$ enters the filter passband rendering them unusable.  The corrections for Q\\,2203+292 in $R_C$ system are: $R_C$$-$$r_s$=+0.36 mag, $R_C$$-$$r_4$=$-$0.29 mag. We assign to them tentative errors equal to the differences between the $R$$-$$r$ values derived for this redshift from the quasar spectrum and from our own composite spectrum: 0.02 mag for both $R_C$$-$$r_s$ and $R_C$$-$$r_4$. They were added in quadrature to the observational uncertainties of the first two data points of our light curve (Fig.~\\ref{fig5}). Table~\\ref{Table_UnifLightCurge} lists all available photometry for Q\\,2203+292 in the $R_C$ filter, including the corrected literature measurements. Note, that here we combine systematic with random errors, though.    \\begin{figure} \\vspace*{0.cm}\\psfig{file=Fig_4.eps,width=8.6cm,height=6.7cm} \\caption{The $R_C-r$ colours versus the redshift $z$. The upper panel shows colour terms for $r_s$ (dashed line) and $r_4$ (dotted line). The bottom panel shows $R_{C}-r^\\prime$ (solid line) and the behavior of the mean colour of the SDSS QSO sample according to Jester et al. (2005). The synthetic colours were derived from our composite QSO spectrum.} \\label{fig4} \\end{figure}  \\begin{figure} \\hspace*{-0.2cm}\\psfig{file=Fig_5.eps,width=9.8cm,height=6.7cm} \\caption{Light curve of the quasar Q\\,2203+292 in the rest frame. The open circles show the literature data (Table 1) and the filled circles are our measurements.} \\label{fig5} \\end{figure}    \\begin{table} \\begin{center} \\caption[]{Final $R_C$ light curve of Q\\,2203+292. See Sec.~\\ref{Sec_MagTransf} for details. References: 1 McCarthy et al. (1988), 2 Schneider et al. (1989), 3 Crampton et al. (1992), 4 Crampton (private communication), 5 this work.} \\label{Table_UnifLightCurge} \\begin{tabular}{c@{   }c@{   }c@{   }c} \\hline JD-2447000+& $R_C$(error),mag & Ref. \\\\ \\hline 0064.04166 & 21.14 (0.09) & 1 \\\\ 0417.28020 & 20.59 (0.05) & 2 \\\\ 0773.84030 & 20.70 (0.10) & 3 \\\\ 1064.04170 & 20.40 (0.10) & 4 \\\\ 4313.43375 & 20.22 (0.04) & 5 \\\\ 5878.38186 & 20.41 (0.04) & 5 \\\\ 6288.31451 & 20.51 (0.03) & 5 \\\\ 6645.39887 & 20.44 (0.06) & 5 \\\\ 6679.30514 & 20.37 (0.03) & 5 \\\\ 6681.28100 & 20.40 (0.03) & 5 \\\\ 6916.41181 & 20.41 (0.05) & 5 \\\\ 6966.39679 & 20.46 (0.04) & 5 \\\\ 6967.42921 & 20.42 (0.04) & 5 \\\\ 6972.54666 & 20.31 (0.04) & 5 \\\\ 7032.18690 & 20.39 (0.08) & 5 \\\\ 7033.28848 & 20.46 (0.08) & 5 \\\\ 7354.46565 & 20.30 (0.02) & 5 \\\\ 7378.27326 & 20.34 (0.01) & 5 \\\\ \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "\\subsection[]{Variability}  The quasar shows (Fig.~\\ref{fig5}) a brightness increase of $\\sim$0.75 mag at the beginning of the coverage but it is nearly constant later. Due to the gaps in the lightcurve, any non-linear fluctuations (such as flares) cannot be ruled out. We verify the variability properties of the Q\\,2203+292 by means of differential photometry and Monte Carlo simulation.  The differential light curves (Fig.~\\ref{fig2}) of the quasar were generated with respect to the reference stars A and B, two of the most stable stars near the quasar. The rms of the relative light curve A--B is 0.035 mag, and since the two stars have similar magnitudes, their individual errors are $\\sim$0.025 mag. The rms for the quasar light curve with respect to star A (marked as Q--A) is 0.083 and with respect to star B (marked as Q--A) is 0.086 mag. The maximum peak-to-peak variation of the quasar is $\\sim$0.92 mag over the entire monitoring period (Table.~\\ref{Table_UnifLightCurge}) but it is reduced to 0.21 mag if we consider only the Rozhen observations. The rms of all 18 QSO measurements is 0.20 mag, and if only our 14 measurements are considered, it decreases to 0.08 mag.  To test further the variability of Q\\,2203+292, we carried out a Monte Carlo simulation drawing 18 measurements from a constant source with the measured mean magnitude of Q\\,2203+292. Each of these points was generated from a Gaussian distribution with the observational error of the corresponding measurement, so that the artificial datasets more faithfully represent the properties of the real observations. If we consider all data, including the ones from the literature, none of one million simulated data sets exceeded the observed rms However, the colour transformations can be a source of extra uncertainty, so we carried the same simulation only for our 14 measurements to obtained that in 98.5 per cent of the cases the data are inconsistent with a constant source. Excluding the VLT point lowers this probability down to 86.4 per cent.  We conclude that if the colour transformation of the historical observations can be considered reliable, the quasar have undergone a brightening episode in the past but the unaccounted systematic effects stop is from making a strong statement about this. All the literature data consistently deviate from ours in one direction, albeit by different amount, hinting that the variation may be real. Next, our own data show that during the recent years the quasar is in relatively stable state.  The $R_C$ band in the rest-frame of the quasar correspond to the UV flux between 970 and 1420 \\AA, including the Ly$\\alpha$ emission line. To compare the variability properties of Q\\,2203+292 with those of lower redshift quasars, we calculated the structure function S($\\tau$), which is commonly used to characterize the variability of large quasar samples (i.e. Hughes, Aller \\& Aller 1992): \\begin{equation} S(\\tau) =\\langle [m(t)-m(t+\\tau)]^2 \\rangle. \\end{equation} Here, $m(t)$ is the magnitude at time $t$ and $\\tau$ is the time interval between the two measurements in the QSO rest frame. The broken brackets express ensemble average over measurements with the same time interval. The structure function is less sensitive to the inhomogeneity of the observational coverage, and it can be applied to both individual objects and to samples of objects. Note that sometimes in the literature the structure function is defined as a square root of the ensemble average.  The structure function for Q\\,2203+292 is shown in Fig.~\\ref{fig6}. The small number of measurements that form each bin lead to large uncertainties making it difficult to draw definite conclusions. The overall shape of S($\\tau$) for Q\\,2203+292 is similar to that of other QSOs studied in the literature (i.e. Vanden Berk et al. 2004). It is dominated by observational errors for short time intervals and by the intrinsic QSO variability for the longer ones. There is indication that S($\\tau$) may reach a plateau at time scale just above 1 yr. However, quasars are known to vary on a much longer time-scale, so we interpret this as poor sampling. Note that the structure functions of some QSOs may show intrinsic structure in S($\\tau$) that is often interpreted as variability driven by more than one physical mechanism (see Hughes et al. 1992 for examples). We can not exclude that this can be the case with Q\\,2203+292. Further observations over longer time span are necessary to address this question.  \\begin{figure} \\hspace*{0.0cm}\\psfig{file=Fig_6.eps,width=8.7cm,height=8.7cm} \\caption{$R$-band structure function for the quasar Q\\,2203+292 in the rest frame time. The horizontal bars mark the width of the bins used in the ensemble averaging and the vertical ones are the corresponding rms} \\label{fig6} \\end{figure}  A quick comparisons with the literature, typically for $\\tau$$\\sim$0.5-2 yr, shows that Q\\,2203+292 is relatively more variable than the average QSO in the samples of Cristiani et al. (1996) and Wold et al. (2007) and comparable with those of di Clemente et al. (1996) and de Vries et al. (2003).  Recently, Wold et al. (2007) explored the dependence of the quasar variability from the mass of the central black hole. For $M_{BH}$=10$^8$--10$^9$ M$_\\odot$ -- as estimated for Q\\,2203+292 by Turner (1991) -- they obtain $R$-band amplitude of 0.3-0.4 mag, which is similar to the value reported here.   \\subsection[]{Physical properties of Q\\,2203+292\\label{SecPhysProp}}  To calculate the absolute luminosity of Q\\,2203+292, we adopted the following cosmological parameters: $\\Omega_{\\Lambda}$=0.7, $\\Omega_{M}$=0.3, and $H_0$=70 km s$^{-1}$ Mpc$^{-1}$.  The absolute $R$-band magnitude, $M_R$, is related to the apparent magnitude, $R$, by \\begin{equation} M_R=R-A_R-5~log~d_L-2.5~log\\,(1+z)-25+\\Delta\\,R_{k corr}(z) \\end{equation} \\noindent where $A_R$ is Galactic absorption, and $d_L$ is the luminosity distance for a flat Universe.  The K-correction $\\Delta\\,R_{k corr}(z)$ is calculated by convolving a quasar spectrum with a sensitivity curve for a standard $R_{C}$ filter (Fig.~\\ref{fig7}). We used our composite spectrum (see Sec. 2.4) for the SDSS quasars observed in $r^\\prime$ (Schneider et al. 2005) and the Q\\,2203+292 spectrum from Constantin et al. (2002) for our target. In the latter case $\\Delta\\,R_{k corr}$=0.25 mag, 0.06 mag smaller than the value derived for $z$=4.40 from the composite spectrum. Although this is a systematic rather than random error, we added this difference in quadrature to the $M_R$ uncertainty (equal to the rms of all Q\\,2203+292 measurements) to obtain a conservative error estimate. Assuming $A_R$=0.25 mag (Schlegel, Finkbeiner \\& Davis 1998) and using the average $R_C$=20.46 mag we obtained $M_R$=$-$29.39 mag, typical for the SDSS quasar distribution at that redshift (Fig.~\\ref{fig8}).  \\begin{figure} \\vspace{0.0cm} \\hspace*{0.2cm}\\psfig{file=Fig_7.eps,width=8.1cm,height=4.7cm} \\caption{$R_C$-band K-correction as a function of the redshift.} \\label{fig7} \\end{figure}  \\begin{figure} \\hspace*{0.1cm}\\psfig{file=Fig_8.eps,width=8.3cm,height=8.6cm} \\caption{Comparison of the absolute magnitude $M_R$ of Q\\,2203+292 and 46420 SDSS quasars from Schneider et al. (2005). The apparent $R$ magnitudes of the SDSS quasars were calculated from the SDSS $r^\\prime$ magnitudes according to the colour transformations described in Sec.~\\ref{Sec_MagTransf} (see also Fig.~\\ref{fig4}, bottom panel). The bigger dot marks the average absolute luminosity of Q\\,2203+292 and the errorbars are the rms The inset shows a histogram of $M_R$ for 625 quasars at $z$ between 3.9 and 4.9. Again, the location of Q\\,2203+292 measurement is shown with a bigger dot.} \\label{fig8} \\end{figure}  The minimum mass of the emitting gas in the quasar broad line region (BLR) can be estimated following Baldwin et al. (2003), assuming Case B recombination, when all Ly$\\alpha$ photons escape and the electron temperature is $T_e$=20,000 K: \\begin{equation} M_{BLR} = 5.1 (10^{11}/n_e) (L_{{\\rm Ly}\\alpha}/10^{45})~M_{\\odot}, \\end{equation} where $L_{{\\rm Ly}\\alpha}$ is the Ly$\\alpha$ luminosity and $n_e$ is the electron density.  We measured $L_{{\\rm Ly}\\alpha}$ from the spectra of Schneider et al. (1989) and Constantin et al. (2002): 8.6$\\times$10$^{43}$ and 5.6$\\times$10$^{43}$ erg s$^{-1}$, respectively. Naturally, these are only lower limits because of the Ly$\\alpha$ self-absorption. In both cases, we fitted the quasar continuum with a power law $F_{\\nu} \\propto \\nu^\\alpha$, and fixing $\\alpha$=$-$1.0. Assuming $n_e$=10$^{11}$ cm$^{-3}$, we obtain $M_{BLR}$=0.44 and 0.29 M$_\\odot$ for the measurements from the two spectra.   \\subsection{Miscellaneous: Search for associated emission line objects}  Narrow band Ly$\\alpha$ imaging down to 25.5 mag per sq. arcsec (McCarthy et al. 1988) yielded no other emission line sources at the same redshift within 2$\\times$2 arcmin$^2$ from Q\\,2203+292. Thompson, Djorgovski \\& Beckwith (1994) failed to find associated [O{\\sc iii}] emitters in 18.2$\\times$19.4 arcsec$^2$ field centred at the quasar.  We used the Photometrics\\,AT200A camera at the 2m telescope at the Rozhen Observatory to carry out a search for associated emission line objects on 2006 August 19. The observations were obtained through a narrow band (FWHM=32 \\AA) interference filter $IF658$ centred at 6572 \\AA, corresponding to Ly$\\alpha$ at $z$$\\sim$4.4. Our field of view was 5 arcmin$^2$, which is much bigger than that in the earlier studies. The exposure time was 2 h, which was split into six 1200 s exposures. We found no evidence for sources with emission lines falling into the bandpass of our narrow band filter down to a surface brightness level of $\\sim$24.5 mag per sq. arcsec, in agreement with the previous attempts."
    },
    "0801/0801.2613_arXiv.txt": {
        "abstract": "We report a search for the presence of dust in the intra-cluster medium based on the study of statistical reddening of background galaxies. Armed with the Red Sequence Cluster survey data, from which we extracted ({\\em i}) a catalog of 458 clusters with $z_{clust} < 0.5$ and ({\\em ii}) a catalog of $\\sim$ 90,000 galaxies with photometric redshift $0.5 < z_{ph} < 0.8$ and photometric redshift uncertainty $\\delta z_{ph} / (1+z_{ph}) < 0.06$, we have constructed several samples of galaxies according to their projected distances to the cluster centers. No significant color differences [$<$E($B-R_c$)$>$ = $0.005 \\pm 0.008$, and $<$E($V-z'$)$>$ = $0.000 \\pm 0.008$] were found for galaxies background to the clusters, compared to the references. Assuming a Galactic extinction law, we derive an average visual extinction of $<$A$_V$$>$ = $0.004 \\pm 0.010$ towards the inner $1\\times R_{200}$ of clusters. ",
        "introduction": "\\cite{zwi51} was the first one to observe a weak and diffuse intra-cluster light in the Coma cluster, suggesting the presence of material in the intra-cluster medium (ICM). The jump in sensitivity of the CCD detectors later allowed intra-cluster light to be studied more quantitatively (see e.g., \\citealt{ber95,fel02,zib05}). The existence of an intra-cluster stellar population has further been clearly established with the detection of planetary nebulae \\citep{arn96,fel98}, red giant stars \\citep{fer98} and supernovae \\citep{gal03} in the ICM. In parallel, X-ray spectroscopic observations revealed the presence of heavy metals in the ICM \\citep{mit76,mus96}. This ICM matter is commonly believed to originate from stripping and/or disruption of cluster galaxies, and one can then naturally wonder about the existence of dust in the ICM.  It is true that the ICM is a harsh medium, where strong X-ray radiation and gas temperatures of $\\sim$ 10$^7$ K may impose severe restrictions on the presence of dust. However, the survival timescale for dust grains in this environment is still of the order of few 10$^8$ yr \\citep{dra79}. Besides, many mechanisms, such as tidal or ram pressure stripping, galactic winds, mass loss from ICM red giants and supergiants, or possible accretion of primordial dust, may contribute to the release or replenishment of dust in the ICM [see, e.g.,\\cite{pop00}, \\cite{sch05}, \\cite{dom06} for predictions on the efficiency of the different mechanisms].  While there have been many attempts, using various methods, to detect the existence of dust in the ICM (see Table~\\ref{prevsearch}), the debate is still open. The methods are mostly twofold: either based on a direct search for dust thermal emission, typically peaking in the far infrared to submillimeter range, or on an indirect search for extinction and reddening of background sources.  The direct detection of extended thermal emission in the intra-cluster medium would constitute the best evidence for the presence of ICM dust. Theoretical predictions of the IR emission have been made (see e.g., \\citealt{dwe90}, \\citealt{yam05}); taking into account dust injection by galaxies, sputtering, and heating by the hot ICM plasma, they yield a mean dust temperature of $\\sim 20 - 30$ K, and dust-to-gas ratio $\\sim 10^{-6}$. \\cite{ann93} could not detect significant emission from a submillimeter search towards 11 clusters with cooling flows, up to a limit of $\\sim 10^8$ M$_\\odot$. \\cite{wis93} observed with IRAS a larger sample of 56 clusters, from which only two may show evidence for extended diffuse emission. Further ISO observations revealed  FIR excess emission towards Coma \\citep{sti98}, possibly due to a dust mass of $\\sim 10^7 - 10^9$ M$_\\odot$ in the inner 0.4 Mpc region, but none towards five other Abell clusters \\citep{sti02}. More recently, \\cite{mon05} adopted a statistical approach and co-added IRAS maps of a large number ($> 11,000$) of clusters and detected a significant emission in all four IRAS band (12, 25, 60 and 100 $\\mu$m). The origin of the emission, tentatively attributed to ICM dust, is, however, questionable. Indeed, in addition to the challenging detection of low surface brightness and extended emission, IR observations of ICM dust might be affected by contamination from Galactic cirrus (e.g., \\citealt{wis93,bai07}), as well as from cluster members \\citep{qui99}, dusty star forming galaxies in clusters (\\citealt{gea06}), dust associated with the central dominant galaxies or the presence of a cooling flow \\citep{bre90,gra90,cox95,edg99}, or even from background (possibly magnified) galaxies.  The first claim for ICM dust was, nevertheless, historically made by \\cite{zwi62}, based on the extinction of light from distant clusters by nearby ones. He estimated A$_V \\sim 0.4$ for Coma. \\cite{kar69} derived a similar value of A$_V \\sim 0.3$ mag for Coma, and found a mean cluster extinction of A$_V \\sim 0.20 \\pm 0.05$ mag by averaging over 15 clusters. \\cite{bog73} measured that distant Abell clusters appear to be inversely correlated on the sky with nearby ones. To account for this effect, they argued for an extinction corresponding to A$_{V} \\sim 0.4$ mag and extending to $\\sim 2.5$ times the optical radii of the nearby clusters. However, these first results were probably severly affected by the difficulties of identifying background clusters.  Obvious background sources to search for dust extinction and reddening are quasars. \\cite{boy88} measured a value of A$_B = 0.2$ mag (A$_ V \\sim 0.15$ mag) from a deficiency of ultraviolet-excess-selected objects behind clusters of galaxies. Similarly, \\cite{rom92} concluded that an area-averaged extinction of A$_B \\sim 0.5$ mag was required to explain the anti-correlation between quasars and Abell clusters. On the other hand, \\cite{mao95} could not measure a significant difference [E($B-V$)$\\leq 0.05$] between the color distribution of radio-selected quasars located within 1$\\degr$ of a foreground Abell cluster, and quasars with no cluster in their line of sight. He reached the conclusion that the apparent deficiency of optically discovered quasars in the background of clusters could be attributed to selection effects, reflecting the difficulty to identify quasars in crowded fields.  More recently, \\cite{nol03} proposed, to our knowledge, the first and only search so far of ICM dust based on extinction and reddening of background galaxies. They used a catalog of $\\sim 140$ nearby ($z < 0.08$) clusters with $B$ and $R$ photometric data, complete to $B \\sim 20.5$ mag and $R \\sim 19.5$ mag. For each cluster, they constructed two samples. The ``cluster'' sample includes all galaxies lying in a circular region centered on the cluster, and with an angular radius equivalent to the physical radius of 1.3 Mpc at the distance of the cluster. The second sample, namely ``control sample'', contains galaxies projected within a ring of inner radius 1.3 Mpc and of outer radius such that the area covered by the cluster and control samples are identical. The cluster and control samples for all clusters were then respectively merged to increase the number of galaxies and improve the statistics. The cluster sample inevitably suffers from contamination of cluster members and foreground galaxies, although \\cite{nol03} estimate that these latest amount to less than 10\\% contamination. The cluster and control group galaxies were then distributed in two color-magnitude diagrams (CMD), and any extinction and reddening of the cluster group was searched by shifting the cluster CMD both in magnitude and color with respect to the control CMD, until the two CMD matched. The pixel resolution (in magnitude and color) of the CMD is constrained by the number of galaxies in each pixel, which should be high enough. Given the number of galaxies in each of their samples ($<10^5$), \\cite{nol03} fixed the smallest reasonable pixel size for $\\Delta$($B-R$) and $\\Delta R$ to 0.025. They could not detect a significant difference between the cluster and control samples, thus yielding upper limits of A$_R < 0.025$ and E($B-R$)$< 0.025$ mag (A$_V < 0.044$) for ICM dust.  We note briefly that dust extinction and reddening of background galaxies have been successfully detected in galactic disks (see, e.g., \\citealt{gon98}, \\citealt{hol05}), galactic halos \\citep{zar94}, and, recently, in the intergalactic medium towards the M81 group \\citep{xil06}. This simple method should become more and more popular with the advent of wide field cameras on board of large telescopes.  Finally, in a very recent paper, \\cite{che07} claim the detection of reddening of background quasars behind $0.1 < z < 0.3$ galaxy clusters using the {\\em Sloan Digital Sky Survey} (SDSS) photometric and spectroscopic data. They report E($g-i$) of $0.003 \\pm 0.002$ from photometry analysis, on one hand, and $0.008 \\pm 0.003$ from spectroscopic extinction curve, on the other hand, towards the central $\\sim 1$ Mpc of clusters. Assuming a Galactic extinction curve (\\citealt{sch98}), this converts to A$_V \\sim 0.005 \\pm 0.003$ and $0.013 \\pm 0.004$, respectively. The reference quasars were taken at large projected distances ($> 7$ Mpc) from the cluster centers. Their measurement of average reddening at different radial annuli seems to indicate both a large covering factor and a large scale distribution of dust, up to $5-6$ Mpc.  In the present paper, we propose to address the question of the presence of dust in the ICM by studying the statistical reddening of galaxies located in the background of clusters, as compared to a reference sample of galaxies located away from the line of sight to clusters. We have constructed the samples based on the cluster and galaxy catalogs from the Red-sequence Cluster Survey (RCS, \\citealt{gla05}), and using galaxy photometric redshifts. In \\S2, we briefly describe the RCS data and present our galaxy sample selection and analysis. Our results are discussed in \\S3, and a summary is given in \\S4. ",
        "conclusions": "In this paper, we report a search for the presence of dust in the intracluster medium. Our method relies on the statistical measurement of the color differences between galaxies located in the background of a cluster and reference galaxies. We take advantage of data acquired with the CFHT in the frame of the Red Sequence Cluster Survey, providing us with deep and homogeneous catalogs of galaxies and galaxy clusters. For our purpose, a major asset of these RCS data is the availability of photometric redshifts for more than a million galaxies. We extracted 458 clusters up to a redshift of 0.5, and a total of $\\sim 90,000$ galaxies, selected from their photometric redshift $0.5 < z_{ph} < 0.8$ and with photometric redshift uncertainty $\\delta z_{ph}/(1+z_{ph}) < 0.06$. In particular, our selection criteria ensure low contamination by foreground galaxies and cluster members. Special attention was given to the propagation of photometry and photometric redshift uncertainties, which were taken into account by a combination of Monte-Carlo and bootstrap methods.  A sample of about 8,000 galaxies projected within $1\\times R_{200}$ of a cluster is compared to different reference samples composed of galaxies projected at distances [1-2], [2-3], [3-4] and [6-7]$\\times R_{200}$ from a cluster center. We do not detect significant color differences between the first sample and the different references. Unless intracluster dust extends over scales larger than $\\sim 10$ Mpc in radius, we therefore conclude that the dust extinction is statistically low, with visual extinction $<$A$_V$$>$ = $0.004 \\pm 0.010$ for the 458 RCS clusters with $z < 0.5$, corresponding to an upper dust mass limit of $8 \\times 10^9$ M$_\\odot$ within $1 \\times R_{200}$. Given the uncertainties, our results are consistent with the recent measurements of reddening of background quasars by galaxy clusters from \\cite{che07}.  The on going RCS2 survey, extending the RCS coverage to $\\sim 1,000$ deg$^2$ in four photometric bands with the CFHT MegaCam will soon be available. The observed sky surface will be multiplied by a factor of 30, providing a much larger number of background galaxies. These new data could allow us to improve our measurements and compare the dust reddening towards clusters with different properties, such as richness, X-ray brightness or merging activity."
    },
    "0801/0801.4173_arXiv.txt": {
        "abstract": "SWIFT~J1756.9-2508 is one of the few accreting millisecond pulsars (AMPs) discovered to date. We report here the results of our analysis of its aperiodic X-ray variability, as measured with the Rossi X-ray Timing Explorer during the 2007 outburst of the source. We detect strong ($\\sim$35\\%) flat-topped broadband noise throughout the outburst with low characteristic frequencies ($\\sim$0.1~Hz). This makes SWIFT~J1756.9-2508 similar to the rest of AMPs and to other low luminosity accreting neutron stars when they are in their hard states, and enables us to classify this AMP as an atoll source in the extreme island state. We also find a hard tail in its energy spectrum extending up to 100~keV, fully consistent with such source and state classification. ",
        "introduction": "\\label{sec:intro}  Nine years after the discovery of the first accreting millisecond pulsar (AMP), the number and variety of this type of neutron star low-mass X-ray binary (NS-LMXB) continues to grow. The idea that a neutron star could be spun up by accretion torques in the bosom of a LMXB \\citep{Alpar82, BhatHeuv91} was confirmed in 1998, when the first millisecond pulses were seen from a LMXB \\citep{Wijnands98}. Besides coherent pulsations, two types of quasi-periodic millisecond X-ray variability have been seen in AMPs: twin kilohertz quasi-periodic oscillations (kHz QPOs) in SAX~J1808.4--3658 and XTE~J1807--294 \\citep[][]{Wijnands03,Linares05} and burst oscillations also in SAX~J1808.4--3658 and in XTE~J1814--338 \\citep[][]{Chakrabarty03,Strohmayer03}. Together, these phenomena place important constraints on kHz QPO and type I X-ray burst models, offering a new insight onto accretion physics \\citep[][for a recent review]{vanderKlis06}. As far as low frequency variability is concerned, the fastest spinning AMP, IGR~J00291+5934, showed the strongest broadband noise with the lowest characteristic frequencies observed to date in a NS-LMXB \\citep{Linares07}.  Even though AMPs have provided new input to better understand accretion onto low magnetic field NSs, important questions about their physics remain open. For instance, if the magnetic field of the neutron star in AMPs disrupts and channels the accretion flow, thereby producing the observed pulsations, it is not clear yet why the strength or configuration of such magnetic field should differ drastically from that of non-pulsating NS-LMXBs \\citep[see][for a possible explanation]{Cumming01}. In this context it is interesting to note that in three systems millisecond X-ray pulsations have been seen to appear in and disappear from the persistent emission, producing predominant \\citep[HETE~J1900.1-2455;][]{Galloway07} intermittent \\citep[SAX~J1748.9-2021;][]{Altamirano08} or very rare \\citep[Aql~X-1;][]{Casella08} episodes of pulsations. This implies that the AMP within them is only active or visible during a relatively small fraction of the time, which may provide a link with non-pulsating NS-LMXBs.  On June $7^{th}$, 2007, a new X-ray transient was discovered \\citep{Krimm07a} with the burst alert telescope (BAT) onboard {\\it Swift}. Subsequent {\\it RXTE} observations revealed that this was the eighth discovered AMP, with a pulse frequency of $\\sim$182~Hz and an orbital period of $\\sim$54~minutes \\citep{Markwardt07a,Markwardt07b, Krimm07c}. The outburst lasted about two weeks and a possible infrared counterpart was identified \\citep{Burderi07}. No radio pulsations were found in quiescence at 8.7 and 1.4~GHz \\citep[][]{Possenti07,Hessels07}.  In the present work we analyze the aperiodic variability of the source and compare it with other AMPs and with atoll sources \\citep[a low-luminosity class of NS-LMXB, see][]{HK89}. Using both PCA and HEXTE data we also measure its 2-200~keV energy spectrum. We are thereby able to classify SWIFT~J1756.9-2508 as an atoll source and to characterize its accretion state.  \\begin{table}[h] \\center \\tiny \\caption{Log of the observations.} \\begin{minipage}{0.5\\textwidth} \\begin{tabular}{l c c c c c c c} \\hline\\hline Set & ObsID & Date\\footnote{Observation start date in  MJD} & Detectors\\footnote{Average number of active detectors} &  Count Rate \\footnote{Average and standard deviation of the $\\sim$2-35 keV PCA count rate, including all active detectors and not corrected for background (which we estimate from set E to be $\\sim$22.3 c/s/PCU; see Section~\\ref{sec:data})} & 512-s PDS\\footnote{Number of power density spectra extracted from the observation, each of them 512 s long} \\\\ & & & & (c/s) & \\\\ \\hline A & 93065-01-01-02 & 54265.2 & 2.0 & 111.0$\\pm$1.4 & 6  \\\\ & 93065-01-02-00 & 54266.0 & 2.0 & 102.3$\\pm$3.8 & 16  \\\\ \\hline B & 93065-01-02-01 & 54267.0 & 2.0 & 95.3$\\pm$0.9 & 4 \\\\ & 92050-01-01-01 & 54267.1 & 3.0 & 136.7$\\pm$1.0 & 6  \\\\ & 92050-01-01-00 & 54267.4 & 1.7 & 75.3$\\pm$21.1 & 24  \\\\ & 92050-01-01-02 & 54268.1 & 2.5 & 109.3$\\pm$20.9 & 12  \\\\ \\hline C & 92050-01-01-03 & 54269.0 & 2.0 & 71.3$\\pm$1.3 & 6  \\\\ & 92050-01-01-04 & 54269.2 & 2.0 & 68.8$\\pm$4.2 & 5  \\\\ \\hline D & 92050-01-01-06 & 54270.2 & 2.0 & 46.0$\\pm$2.2 & 10 \\\\ & 92050-01-01-07 & 54270.6 & 2.4 & 62.4$\\pm$12.6 & 20  \\\\ \\hline E & 92050-01-01-10 & 54271.6 & 2.0 & 49.4$\\pm$0.9 & 3  \\\\ & 92050-01-01-11 & 54271.8 & 2.0 & 50.0$\\pm$2.2 & 4  \\\\ & 92050-01-01-12 & 54271.9 & 3.0 & 67.0$\\pm$1.6 & 5  \\\\ & 92050-01-01-09 & 54271.9 & 1.0 & 20.9$\\pm$0.3 & 5  \\\\ & 92050-01-01-08 & 54272.4 & 1.3 & 28.6$\\pm$11.1 & 19  \\\\ & 92050-01-01-13 & 54273.0 & 2.0 & 47.4$\\pm$1.2 & 5  \\\\ & 92050-01-02-00 & 54273.5 & 1.0 & 21.8$\\pm$0.6 & 11  \\\\ \\hline\\hline \\end{tabular} \\end{minipage} \\label{table:obs} \\end{table} ",
        "conclusions": "\\label{sec:results}  Figure \\ref{fig:ps} shows the broadband power spectrum of SWIFT~J1756.9-2508. Its overall shape is typical of island and extreme island states (IS/EIS) of atoll sources \\citep[see e.g.][]{Barret00,Belloni02}, as well as similar to low (luminosity) and hard (spectrum) states of black hole X-ray binaries (e.g. XTE~J1118+480; \\citealt{Revni00}, XTE~J1550--564; \\citealt{Cui99}). Furthermore, the integrated (0.01-100~Hz) fractional rms variability was very high ($\\sim$35\\%), also characteristic of such low-hard (Comptonized) states. This allows us to classify this AMP as an atoll source, something common to all the other AMPs known to date \\citep{Wijnands05}. The low characteristic frequencies of the variability (see below) indicate that SWIFT~J1756.9-2508 was in the EIS during its 2007 outburst \\citep[see][for an overview of states]{vanderKlis06}.  \\begin{table}[h] \\center \\tiny \\caption{Parameters of the 4-Lorentzian fits to sets A and B.} \\begin{minipage}{0.5\\textwidth} \\begin{tabular}{l c c c c r} \\hline\\hline Parameter & $L_b$ & $L_h$ & $L_{\\ell ow}$ & $L_u$ & $\\chi^2$/d.o.f. \\\\ \\hline Set A & & & & & \\\\ \\hline ${\\nu}_{max}$ (Hz) & 0.12$\\pm$0.01 & 0.45$\\pm$0.03 & 3.1$\\pm$0.9 & 53$\\pm$30 & 167/150 \\\\ rms (\\%) & 17.7$\\pm$0.8 & 13.1$\\pm$1.6 & 17.9$\\pm$1.3 & 21$\\pm$3 & \\\\ Q & 0 (fixed) & 0.9$\\pm$0.3 & 0 (fixed) & 0 (fixed) & \\\\ \\hline Set B & & & & & \\\\ \\hline ${\\nu}_{max}$ (Hz) & 0.09$\\pm$0.01 & 0.39$\\pm$0.02 & 3.1$\\pm$0.6 & 141$\\pm$41 & 142/150 \\\\ rms (\\%) & 18.3$\\pm$0.8 & 15.1$\\pm$1.5 & 21.0$\\pm$0.9 & 34$\\pm$3 & \\\\ Q & 0 (fixed) & 0.7$\\pm$0.2 & 0 (fixed) & 0 (fixed) & \\\\ \\hline\\hline \\footnotetext{See Section~\\ref{sec:data} for details of the fit function and notation. All errors are 1$\\sigma$.} \\end{tabular} \\end{minipage} \\label{table:psfits} \\end{table}  The most prominent variability feature is a strong, band-limited and flat-topped noise component ($L_b$, for ``break''), with a break or cutoff at a frequency ${\\nu}_b$$\\sim$0.1~Hz. This break frequency lies in the low frequency end of the range usually observed in NS-LMXBs. We also detect a narrower (Q$\\sim$0.8) component at higher frequencies ($\\sim$0.4~Hz). The correlations between characteristic frequencies of power spectral features are a useful tool to identify such features \\citep{WK99,PBK99,Straaten05}. The narrow component at $\\sim$0.4~Hz follows the correlation found by \\citet{WK99} between $\\nu_b$ and the frequency of the ``hump'' ($L_h$) present above $L_b$ in both black hole and neutron star systems (see Figure~\\ref{fig:bhwk}), and we therefore identify it as $L_h$. The next component in order of increasing frequency has a characteristic frequency of 3.1~Hz, close to (but a bit lower than) what we would expect from extrapolating the relation between $L_{\\ell ow}$ and $L_b$ found in atoll sources \\citep{Straaten02,Straaten03}. This leads us to tentatively identify it as $L_{\\ell ow}$. The highest frequency component that we find in the power spectrum of SWIFT~J1756.9-2508 is most likely $L_u$, also present also present in the same frequency range in IS/EIS of atoll sources. A less likely candidate is the so-called ``hecto-Hertz'' Lorentzian, as that component seems to be present only when the variability is at higher frequencies, with $\\nu_b$ more than one order of magnitude higher \\citep{Straaten02}. All parameters from the fits to sets A and B are shown in Table~\\ref{table:psfits}.  In disk accreting X-ray binaries as the flat-topped noise shifts down in frequency its flat top becomes higher. In other words, the rms level at the break of the power spectrum is anticorrelated with the frequency of that break. This was found by \\citet{BH90} in Cyg~X-1 and extended to other BH and NS systems by \\citet{Belloni02}. SWIFT~J1756.9-2508 also follows this anticorrelation, as can be clearly seen from Figure~\\ref{fig:bhwk}.  The broadband 2-200~keV spectra from all observations within sets A, B and C were reasonably well fitted (reduced $\\chi^2 <$ 1.5 for 61 degrees of freedom) with the adopted (Sec.~\\ref{sec:data}) model (see Figure~\\ref{fig:spec}). The resulting parameters are shown in Table~\\ref{table:spfits}. A hard tail extending up to $\\sim$100~keV was present in the energy spectra during most of the outburst, similar to island states of other atoll sources \\citep{DiSalvo04}. As can be seen from Figure~\\ref{fig:tev}, the resulting power law index increases when the outburst decays. Several models have been proposed for the emission process at work in the transition from outburst to quiescence, including pulsar shock emission and thermal emission from the neutron star \\citep[][for a review]{Campana98}. Our data are not sensitive enough to yield precise tests of such models, and in general sensitivity limits and time constrains make this task a difficult one.  \\begin{table}[h] \\tiny \\caption{Parameters of the spectral fits.} \\begin{minipage}{0.5\\textwidth} \\begin{tabular}{l c c c r r} \\hline\\hline ObsID & Ph.I. & P.L. & kT  & $\\chi^2$/d.o.f. & 2-200~keV unabs.\\\\ &  & norm &  (keV) &  & Flux ($erg/s/cm^2$)\\\\ \\hline Set A & & & & &\\\\ \\hline 93065-01-01-02  &  1.85$\\pm$0.02 &  (13.0$\\pm$0.6)$\\times 10^{-2}$ & 0.5$\\pm$0.2  &  43/61 & (13.0$\\pm$0.4)$\\times 10^{-10}$ \\\\ 93065-01-02-00  &  1.84$\\pm$0.02 &  (11.0$\\pm$0.6)$\\times 10^{-2}$ & 0.7$\\pm$0.1  &  53/61 & (11.7$\\pm$0.2)$\\times 10^{-10}$ \\\\ \\hline Set B & & & & &\\\\ \\hline 93065-01-02-01  &  1.92$\\pm$0.04 &  (11.0$\\pm$0.9)$\\times 10^{-2}$ & 0.5$\\pm$0.1  &  59/61  & (9.7$\\pm$0.5)$\\times 10^{-10}$\\\\ 92050-01-01-01  &  1.94$\\pm$0.04 &  (11.0$\\pm$1.1)$\\times 10^{-2}$ & 0.6$\\pm$0.4  &  61/61  & (8.7$\\pm$0.4)$\\times 10^{-10}$\\\\ 92050-01-01-00  &  1.88$\\pm$0.02 &  (9.6$\\pm$0.5)$\\times 10^{-2}$ & 0.6$\\pm$0.1  &  70/61  & (8.0$\\pm$0.2)$\\times 10^{-10}$\\\\ 92050-01-01-02  &  1.90$\\pm$0.03 &  (9.1$\\pm$0.6)$\\times 10^{-2}$ & 0.7$\\pm$0.1  &  83/61  & (6.8$\\pm$0.3)$\\times 10^{-10}$\\\\ \\hline Set C & & & & &\\\\ \\hline 92050-01-01-03  &  2.03$\\pm$0.05 &  (6.4$\\pm$0.6)$\\times 10^{-2}$ & 0.5$\\pm$0.1  &  73/61  & (2.8$\\pm$0.3)$\\times 10^{-10}$\\\\ 92050-01-01-04  &  1.98$\\pm$0.06 &  (4.7$\\pm$0.7)$\\times 10^{-2}$ & 0.4$\\pm$0.2  &  41/61  & (2.8$\\pm$0.3)$\\times 10^{-10}$\\\\ \\hline\\hline \\footnotetext{NOTE.-- In all fits the column density was fixed to 5.4$\\times 10^{22}$ $cm^{-2}$ \\citep{Krimm07c}. The power law (P.L.) is characterized by the photon index (Ph.I.) and the normalization, which corresponds to the photon flux density at 1~keV. In five cases the B.B. normalization was not significant but in all 8 observations it gave a $2-4~\\sigma$ improvement in the fit, according to the f-test for additional terms. All errors are 1$\\sigma$.} \\end{tabular} \\end{minipage} \\label{table:spfits} \\end{table}   We do not detect a significant change in the fractional rms of $L_b$ with energy in the studied energy range, suggesting that the spectrum of this component is flat. Previous work has shown that in LMXBs the X-ray luminosity does not uniquely determine the characteristic frequencies of the variability \\citep{Ford00,Straaten05,vanderKlis06} so that another parameter must act to decouple luminosity and variability frequencies. Figure~\\ref{fig:tev} shows that as the outburst decays and the inferred mass accretion rate drops the fractional rms variability seems to increase, whereas the frequency of the break decreases. The same behavior has been observed e.g. in IGR~J00291+5934 \\citep{Linares07}. A successful model for broadband noise in LMXBs should explain both this short-term correlation between luminosity and break (or ``cutoff'') frequency and the long-term luminosity-frequency decoupling discussed above.   \\textbf{Acknowledgments:}  {\\it Swift}-BAT transient monitor results provided by the {\\it Swift}-BAT team \\newline ({\\it http://swift.gsfc.nasa.gov/docs/swift/results/transients/}).  \\clearpage"
    },
    "0801/0801.3753_arXiv.txt": {
        "abstract": "We present the results of an optical and X-ray  monitoring campaign on the short-period massive SB2 binary HD~152218.  Combining our HiRes spectroscopic data with previous observations, we unveil the contradictions between the published orbital solutions. In particular, we solve the aliasing on the period and derive a value close to 5.604~d. Our eccentricity $e=0.259\\pm0.006$ is slightly lower than previously admitted. We show that HD~152218 is probably undergoing a relatively rapid apsidal motion of about 3\\degr~yr$^{-1}$ and we confirm the O9IV+O9.7V classification. We derive minimal masses of $15.82 \\pm 0.26$~\\msol\\ and   $12.00 \\pm 0.19$~\\msol\\ and constrain the radius of the components to  $R_1=10.3\\pm1.3$~\\rsol\\ and $R_2=7.8\\pm1.7$~\\rsol. We also report the results of a \\xmm\\  monitoring of the  HD~152218 X-ray emission throughout its orbital motion. The averaged X-ray spectrum is relatively soft and it is well reproduced by a 2-T optically thin thermal plasma model with component temperatures about 0.3 and 0.7~keV. The system presents an increase of its  X-ray flux by about 30\\% near apastron compared to periastron, which is interpreted as the signature of an ongoing wind-wind interaction process occurring within the wind acceleration region. ",
        "introduction": "\\label{sect: intro}  Early-type stars of spectral type O are characterized by strong stellar winds. Though not as extreme as those of  their evolved descendants, the Wolf-Rayet stars, these winds combine  terminal velocities of a few thousand \\kms\\ and important mass-loss rates (about $10^{-7}-10^{-5}$~\\msol~yr$^{-1}$) that significantly affect both the surroundings of the star and its evolution. In a binary system, it is expected that the winds from the two stars collide, producing a density enhanced region known as the wind interaction zone. The shock-heated plasma within this zone is expected to produce an additional contribution to the X-ray emission from the early-type system and, indeed, the early-type binaries are known to be statistically X-ray overluminous compared to single stars of the same spectral type \\citep{ChG91}. This extra X-ray emission might display phase-locked modulations, either due to a change of the absorption properties along the line of sight, or to a modulation of the shock strength due e.g. to a varying separation between the components in an eccentric binary. The properties of the shock thus strongly depend both on the geometry of the binary system and on the characteristics of the individual winds. However these properties often cruelly lack accurate constraints. In this framework, our team is involved in a long-standing effort to study early-type stars in a number of open clusters.  The present work focuses on the  short-period binary \\astrobj{HD152218}, located in \\astrobj{NGC 6231} at the core of the \\astrobj{Sco OB 1} association. This star is a long known SB2 system \\citep{Str44}. Previous papers \\citep{HCB74, LM83, SLP97, GM01} however reported conflicting results and we thus decided to re-appraise the orbital properties of this apparently well-studied system. The aim is to bring  its orbital and physical parameters on firm ground. As mentioned above, accurate ephemeris and a detailed knowledge of the orbital and physical properties of a system are indeed crucial ingredients for interpreting the X-ray data and for uncovering the possible signature of a wind interaction.   The remaining of this paper is organised as follows. The next section describes the observational material and the data reduction processes. Sect. \\ref{sect: orbit} aims at unveiling the discrepancies of the previously published orbital solutions. Sect. \\ref{sect: physic} discusses the  physical properties and evolutionary status of HD~152218 and, in Sect. \\ref{sect: xmm}, we analyse the X-ray observations in the light of the upgraded orbital solution. Finally Sect. \\ref{sect: ccl} provides a summary of the present work. ",
        "conclusions": "\\label{sect: ccl}  We have presented the results of a long-term spectroscopic monitoring campaign on the massive binary \\astrobj{HD 152218}. The last two investigations on the object were presenting conflicting results. The present data  complete  the existing observational sets and allow to definitely overcome the ambiguity on the orbital period of the system, bringing to firm ground the period of  5.604~d. From our analysis, we reject the shorter period and higher eccentricity proposed by \\citet{GM01} and rather confirm the earlier work of \\citet{SLP97}, based on IUE observations.  Our newly derived orbital solution is characterized by a slightly lower eccentricity than previously accepted: $e=0.259\\pm0.006$. The system is most probably formed by an O9 sub-giant and an O9.7 main sequence star. We derived minimal masses of $15.82 \\pm 0.26$~\\msol\\ and   $12.00 \\pm 0.19$~\\msol\\ and we constrained the component radii to values of $R_1=10.1\\pm1.0$~\\rsol\\ and $R_2=8.9\\pm2.0$~\\rsol.  These values are in good agreement with previous findings and further indicate that \\astrobj{HD 152218} should have an orbital inclination of about 60--70\\degr. This corresponds to a limiting case in which eclipses or ellipsoidal variations might be observed.  We also report the results of the monitoring of the system in the X-rays thanks to \\xmm\\ observations of the cluster. The averaged X-ray spectrum is relatively soft. It is well reproduced by a 2-T \\mek\\ model with component temperatures about 0.3 and 0.7~keV. We showed that the system presents an increase of its  X-ray flux  of about 30\\% near apastron compared to periastron. We note that this could be the signature of an ongoing wind-wind interaction process occurring within the wind acceleration region. Such a scenario indeed predicts a stronger shock, and thus a larger X-ray emission from the colliding zone, when the separation between the components is larger, allowing the winds to reach larger pre-shock velocities. This is also supported by the apparent modulation of the hardness ratio, which seems higher when the emission level is stronger. Preliminary computations however indicate that no simple modeling, reproducing the main properties of the interaction, can be obtained. By contrast, the same computations rather point towards second order effects, such as radiative inhibition, to play a major role in this system. These effects could be strong enough to govern the structure and shape of the interaction. HD~152218 might thus be an excellent case to test and/or refine wind collision models, including detailed physics on the wind acceleration and inhibition."
    },
    "0801/0801.4729_arXiv.txt": {
        "abstract": "The development of turbulent gas flows in the intra-cluster medium and in the core of a galaxy cluster is studied by means of adaptive mesh refinement (AMR) cosmological simulations. A series of six runs was performed, employing identical simulation parameters but different criteria for triggering the mesh refinement. In particular, two different AMR strategies were followed, based on the regional variability of control variables of the flow and on the overdensity of subclumps, respectively. We show that both approaches, albeit with different results, are useful to get an improved resolution of the turbulent flow in the ICM. The vorticity is used as a diagnostic for turbulence, showing that the turbulent flow is not highly volume-filling but has a large area-covering factor, in agreement with previous theoretical expectations. The measured turbulent velocity in the cluster core is larger than $200\\ \\rmn{km\\ s^{-1}}$, and the level of turbulent pressure contribution to the cluster hydrostatic equilibrium is increased by using the improved AMR criteria. ",
        "introduction": "\\label{intro}  The potential importance of turbulence for the physics of galaxy clusters has been widely recognised in recent years (see the discussion below for details and references). The precise nature of turbulence in the intra-cluster medium (ICM), however, remains controversial. The relevant dimensionless parameter for the onset of turbulence is the Reynolds number $Re$: \\begin{equation} \\label{reynolds} Re = \\frac{LV}{\\nu}\\,\\,, \\end{equation} where $L$ is the integral scale of the flow instability, $V$ is the characteristic velocity at scale $L$, and $\\nu$ is the kinematic viscosity of the fluid. For a fully turbulent flow, $Re \\ga 1000$. In the framework of galaxy cluster physics, the uncertainty in $Re$ results mostly from the determination of $\\nu$. In the unmagnetised case, the Braginskii formulation \\citep{b65} for the viscosity is generally used, leading to the frequently reported estimate of $Re \\sim 100$ \\citep[e.g.][]{snb03} in the ICM.  This assumption does not hold in a magnetised plasma. As known from theory \\citep{s62}, the presence of magnetic fields suppresses the transport coefficients below the unmagnetised value by a factor $f$ which depends on the tangled field structure. While highly uncertain, it is estimated in the range $10^{-2}$ -- $1$ (\\citealt{rmf05}; see also \\citealt{nm01}). An effectively suppressed ($f \\sim 10^{-2}$) viscosity would lead to $Re \\gg 1$ both in the cluster cores and in the hotter and less dense ICM \\citep{j07}. On the other hand, it has been claimed that a factor $f \\ga 0.1$ is consistent with the observation of filaments in the Perseus cluster \\citep{fsc03}\\footnote{A weakly suppressed, or unsuppressed viscosity is also required for preserving the morphological stability of AGN-driven bubbles \\citep[e.g.][]{rmf05,ss06} and for the cluster heating by viscous dissipation \\citep{fsa03,fsc03}, though the numerical simulations of \\citet{ss06} question the importance of this contribution to the cluster energy budget.}. Irrespective of a more precise knowledge of $\\nu$, it appears that the flow in the ICM is, even with a conservative estimate, mildly turbulent.  The turbulent nature of the flow in the intra-cluster medium (ICM) could be directly confirmed with the help of high-resolution X-ray spectroscopy of emission line broadening \\citep{snb03,dvb05,bhr05,rebu07}. The unfortunate flaw in the main instrument of the {\\it Suzaku} satellite postponed this test to the near future. Nevertheless, other observational clues have been interpreted as evidence for the turbulent state of the ICM, such as the analysis of pressure maps of the Coma cluster \\citep{sfm04}, the lack of resonant scattering in the $6.7\\ \\rmn{keV}$ He-like iron $\\rmn{K_\\alpha}$ line in the Perseus cluster \\citep{cfj04}, the broadening of the iron abundance profile in the core of Perseus \\citep{rcb05} and other galaxy clusters \\citep{rcb06}, and the Faraday rotation maps of the Hydra cluster \\citep{ve05,ev06}. In addition to the astrophysical problems mentioned above, an improved knowledge of turbulence in the ICM can shed light on the amplification of magnetic fields \\citep{dbl02,brs05,ssh06}, non-thermal emission in clusters \\citep{b04}, and acceleration of cosmic rays \\citep{mjk01,mrk01,bl07}. The role of turbulence has also been investigated as a source of heating in cooling cores, as described in the analytical model by \\citet{dc05}, where both the dissipation of turbulent energy and the turbulent diffusion are taken into account (cf.~also \\citealt{kn03,fmw04,fsw04}). In this context, the turbulence driven by galaxy motions in the ICM has been studied e.g.~by \\citet{bd89} and \\citet{k07}.  From a theoretical point of view, the production of turbulence in galaxy clusters has been ascribed to two main mechanisms. This work does not address the outflow of active galactic nuclei (AGN), which inflate buoyant bubbles and eventually stir the ICM, but will be focused on turbulence produced during the hierarchical formation of cosmic structures.  Several numerical simulations of galaxy cluster formation show that episodes of active merging can produce turbulence in the ICM \\citep{r98,nb99,t00,rs01,dvb05}, with typical velocities of $300$ -- $600\\ \\rmn{km\\ s^{-1}}$ and injection scales of $300$ -- $500\\ \\rmn{kpc}$. In their simulations, \\citet{dvb05} use a low-viscosity version of the {\\sc gadget-2} SPH implementation \\citep{spri05}, which helps to better resolve turbulent flows. \\citet{ams06} showed that grid methods are more suitable than SPH in modelling dynamical instabilities.  Since the properties of grid-based methods with regard to modelling turbulence are substantially better understood than SPH, we intend to explore their capability in the context of cosmological simulations. These, in turn, rely strongly on adaptive mesh refinement (AMR) in order to follow the evolution of strongly clumped media. Using AMR for turbulent flows is a new and rapidly evolving field \\citep{knp06,knp07,sfh08} which we extend to cluster simulations in this work.  The importance of a proper definition of the criteria for triggering grid refinement in turbulent flows has been studied by \\citet{ias07} (hereafter paper I) by testing novel AMR criteria designed for resolving turbulent flows, and applying them in a simplified subcluster merger scenario. In this paper, we extend this work to cosmological AMR simulations of galaxy clusters. In this framework, we investigate how the AMR criteria tested in paper I can be profitably (from a physical and computational point of view) used for resolving the flow in the ICM. In AMR simulations of structure formation, both the refinement of turbulent flows and the the identification of small subclumps with a relatively low overdensity \\citep{ons05,hlf07} may be equally important; the relative emphasis has to be determined on a case-by-case basis. Both aspects and the related refinement strategies will be addressed in this work.  We focus our analysis on one single galaxy cluster out of the several objects forming in the simulated computational volume. An extended study of turbulence should be based on average cluster properties, sampling several objects of different masses (cf.~\\citealt{dvb05,vtc06}). Such a study is beyond the scope of this work and is left for future analysis.  The paper is organised as follows: in Sec.~\\ref{simulations}, we present the setup of the cosmological simulations and the set of employed AMR criteria. We discuss some general properties of the simulated cluster and present a performance comparison of the simulations in Sec.~\\ref{profiles}. In Sec.~\\ref{properties}, the results are shown together with comparisons between the different runs. The results are summarised and discussed in Sec.~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  The study of turbulent flows in the ICM is important for a thorough understanding of galaxy clusters and for the physics of the plasma in these objects. The turbulent state of the ICM is, from an observational point of view, still a debated issue, and the theoretical determination of the kinematic viscosity contains large uncertainties. Besides these open problems, a specific question has been addressed in this work: are the existing numerical techniques of grid-based, AMR cosmological simulations suitable for studying the main properties of turbulence in the ICM? In particular, our attention was focused on the AMR criteria employed in simulations of strongly clumped objects. A galaxy cluster was therefore simulated with a standard, reference setup, and then the simulation was repeated with five different choices of refinement criteria.  The runs were subdivided into two groups, with an underlying difference in the refinement strategy. In the first group, we used AMR criteria based on the regional variability of control variables of the flow, designed for refining the grid at the locations of significant velocity fluctuations. In the second group, a more accurate tracking of the overdense subclumps was enforced.  The distinction between the two groups of runs is important for the interpretation of the results of Sects.~\\ref{filling} and \\ref{core}. In the runs $B$, $C$ and $D$, the values of $\\max{(|\\bomega|)}$ (Table \\ref{vorticity}) and  $v_{\\rmn{rms}}$ (Table \\ref{core-table}) are larger with respect to run $A$, and to the runs of the second group. The AMR criteria used in these runs are therefore very suitable for refining the grid where the velocity fluctuations are underresolved, and result in an increased magnitude of the turbulence. In the second group of runs, the AMR does not refine explicitly on turbulence, but the lower threshold on overdense region allows a better tracking of subclumps. In other words, the stirring mechanism for turbulence generation in the ICM is better resolved. Therefore comparatively larger $\\langle |\\bomega| \\rangle$ and a larger volume-filling factor are obtained.  In both groups of simulations, the results in the innermost $0.1\\ R_{\\rmn{vir}}$ are consistent with a better modelling of the turbulent flow. In addition to an increased rms velocity, the change in $P_{\\rmn{turb}}/P_{\\rmn{tot}}$ indicates that the turbulent component plays a non-negligible role in the pressure support. Recently \\citet{cfv07}, based on X-ray and optical observations of M87 and NGC 1399, put a constraint of the order of $10\\ \\%$ on the non-thermal contribution to the thermal pressure, due to turbulent motions, cosmic rays and magnetic fields. This is not in contradiction to the findings of our work, because we find that the turbulent pressure is only a few percent of the thermal pressure at most.  The presented comparison shows clearly that run $A$ (which implements AMR criteria widely used in grid-based cosmological simulations) fails to reproduce both the magnitude of the rms velocity and the spatial extent of the turbulent flow. This central issue should be carefully taken into account in future investigations of the ICM turbulent features. These results suggest that computationally more demanding simulations will be able to resolve more flow and substructures, which in turn can further contribute to the turbulent properties.  Apparently the presented AMR approach to the modelling of turbulent flows in the ICM has limited utility outside of the cluster core. As discussed in Sec.~\\ref{core}, this is partly due to the volume-filling properties of turbulence in clusters, and partly to the constraint on the number of produced AMR grids which, if too many, would make the simulation computationally unaffordable.  The presented results are similar to previous simulations \\citep{nb99,dvb05}, as far as the magnitude of the turbulent velocity is concerned. Interestingly, \\citet{k07} shows that the level of ICM turbulence driven by galaxy motions saturates around 220 km/s, which is not very different from our results. As stated in the Introduction, a detailed comparison of the features of turbulence in SPH and grid-based simulation is out of the scope of this paper, but we observe that the trends that we inferred from Table \\ref{core-table} are also present in \\citet{dvb05}. In that case, when the low-viscosity SPH scheme is used, the variation of some of the variables is more marked that in our study. However it should be stressed that in this work we do not make use of a improved numerical scheme to better resolve turbulence, but only of more suitable choices of AMR criteria.  In many astrophysical problems (including galaxy clusters) the range of length scales needed to follow the turbulent cascade down to the Kolmogorov dissipation scale extends well beyond the grid spatial resolution limit, even using AMR. A strong improvement in the modelling of turbulence would be the application of Large Eddy Simulations, where only the largest scales are resolved and the dynamics at length scales smaller than the spatial resolution is handled by a subgrid scale model \\citep{sag01,snh06,snhr06}. Such a tool will allow to consistently evaluate the level of subgrid turbulent energy and will make the numerical information about the level of turbulence in the flow more directly available. We expect that the magnitude of the turbulent motions, also outside $0.1\\  R_{\\rmn{vir}}$, could be substantially increased. The use of this tool in the framework of cosmological simulations will be explored in future work."
    },
    "0801/0801.2538_arXiv.txt": {
        "abstract": "We solve equations, describing in a simplified way the newtonian dynamics of a selfgravitating nonrotating spheroidal body after loss of stability. We find that contraction to a singularity happens only in a pure spherical collapse, and deviations from the spherical symmetry stop the contraction by the stabilising action of nonlinear nonspherical oscillations. A real collapse happens after damping of the oscillations due to energy losses, shock wave formation or viscosity. Detailed analysis of the nonlinear oscillations is performed using a Poincar\\'{e} map construction. Regions of regular and chaotic oscillations are localized on this map. ",
        "introduction": "Dynamic stability of spherical stars is determined by an average adiabatic power $\\gamma=\\frac{\\partial\\log P}{\\partial\\log\\rho}|_S$. For a density distribution $\\rho=\\rho_0 \\varphi(m/M)$, the star in the newtonian gravity is stable against dynamical collapse when $\\int_0^R{(\\gamma-\\frac{4}{3})P\\frac{dm}{\\varphi(m/M)}}>0$ \\citep{ZN, BK1989}. This approximate criterium becomes exact for adiabatic stars with constant $\\gamma$. Here $\\rho_0$ is a central density, $M$ is a stellar mass, $m$ is the mass inside a Lagrangian radius $r$, so that $m=4\\pi\\int_0^r{\\rho r^2 dr}$, $M=m(R)$, $R$ is a stellar radius. Collapse of a spherical star may be stopped only by a stiffening of the equation of state, like neutron star formation at late stages of evolution, or formation of fully ionized stellar core with $\\gamma=\\frac{5}{3}$ at collapse of clouds during star formation. Without a stiffening a spherical star in the newtonian theory would collapse into a point with an infinite density (singularity).  In the presence of a rotation a star is becoming more dynamically stable against collapse. Due to the more rapid increase of a centrifugal force during contraction, in comparison with the newtonian gravitational force, collapse of a rotating star will be always stopped at finite density by centrifugal forces. Here we show, that deviation from the spherical symmetry in a non-rotating star with zero angular momentum leads to a similar stabilization, and non-spherical star without dissipative processes  never will reach a singularity. Therefore collapse to a singularity is connected with a secular type of instability, even without rotation.  When a uniform sphere of cold gas is set into free-fall collapse under its self-gravity its shape is unstable. The linear instability of large-scale modes (second-order harmonics) during the collapse was discovered by \\citet{LB64} and for the case of non-rotating sphere by \\citet{LB79} (see also \\citet{LB96}). Numerical investigations of collapsing pressureless spheroids have been done in the papers of \\citet{LB64} and \\citet{LMS}. In \\citet{LB79} it was also shown that the pressure prevents the development of large-scale shape instability if initially the gravity is more than three-fifths pressure resisted. \\citet{Fuji} derived the equations of motion for a rotating ellipsoid from a system of hydrodynamical equations. However, the accounting for pressure effects in his approach was inconsistent with thermal processes, leading to the wrong results for the dynamical behaviour of a system with radiative losses (see \\citet{BKTs2005}). Using the variational principle and deriving the equation for the entropy function from energy balance give the correct relations for the pressure and the total energy. In \\citet{RT91} the virial equations for rotating Riemann ellipsoids of incompressible fluid are demonstrated to form a Hamiltonian dynamical system.  We calculate a dynamical behavior of a non-spherical, non-rotating star after its loss of a linear stability, and investigate nonlinear stages of contraction. We use approximate system of dynamic equations, describing 3 degrees of freedom of a uniform self-gravitating compressible ellipsoidal body (\\citet{BK2004}, \\citet {BKTs2005}). We obtain that the development of instability leads to the formation of a regularly or chaotically oscillating body, in which dynamical motion prevents the formation of the singularity. We find regions of chaotic and regular pulsations by constructing a Poincar\\'e diagram for different values of the initial eccentricity and initial entropy. For simplicity we restrict ourself by calculating only spheroidal figures with $\\gamma=\\frac{4}{3}$, and only briefly represent results for $\\gamma=\\frac{6}{5}$. At the end we discuss qualitatively effects of general relativity in a non-spherical collapse of a non-rotating body. ",
        "conclusions": "The main result following from our calculations is the indication to a degenerate nature of formation of a singularity  in unstable newtonian self-gravitating gaseous bodies. Only pure spherical models can collapse to singularity, but any kind of nonsphericity leads to nonlinear stabilization of the collapse by a dynamic motion, and formation of regularly or chaotically oscillating body. This conclusion is valid for all unstable equations of state, namely, for adiabatic with $\\gamma <4/3$. In addition to the case with $\\gamma =4/3$, we have calculated the dynamics of the model with $\\gamma=6/5$, and have obtained similar results. The Poincar\\'{e} map for this case is represented in Fig.9. For $\\gamma=6/5$ we have the entropy function $\\varepsilon = E_{th} (abc)^{1/5}$, nondimensional entropy function $\\tilde{\\varepsilon} = \\varepsilon / U_0 a_0^{3/5}$, see (\\ref{scal-par}). For the spherical star with $\\gamma=6/5$ the nondimensional Hamiltonian instead of (\\ref{sphere-H}) is \\begin{equation} \\label{gamma6} \\tilde{H} = \\frac{3m}{10} \\dot{a}^2 - \\frac{3}{5a} m^2 + \\frac{\\tilde{\\varepsilon}}{a^{3/5}} . \\end{equation} Note, that region of chaotic behavior on the  Poincar\\'{e} map is gradually increasing for $\\gamma=4/3$ with decreasing of the entropy $\\varepsilon$ and the total energy $H$. At $\\varepsilon=9/10$ and $H=-3/50$ we have found only regular oscillations (Fig.8), at $\\varepsilon=2/3$ and $H=-1/5$ both kind of oscillations are present (Figs. 4-6), and only chaotic behavior is found at $\\varepsilon=1/6$ and $H=-1/2$ (Fig.7). We connect \\citep{BKTs2005} this chaotic behavior with development of anisotropic instability, when radial velocities strongly exceed the transversal ones \\citep{ant, fp85}.  In reality a presence of dissipation leads to damping of these oscillations, and to final collapse of nonrotating model, when total energy of the body is negative. In the case of core-collapse supernova the main dissipation is due to emission of neutrino. The time of the neutrino losses is much larger than the characteristic time of the collapse, so we may expect that the collapse leads to formation of a neutron star where nonspherical modes are excited and exist during several seconds after the collapse. In addition to the damping due to neutrino emission the shock waves will be generated, determining highly variable energy losses during the oscillations. Besides viscosity and radiation which can damp the ellipsoid-like motions and allow collapse, there is the possibility that inertial interactions with higher order modes that must be present in the real stratified bodies may cause an inertial cascade and drain the energy from the second order modes an a non-secular way that does not depend on dissipative coefficients. It is very seductive to connect chaotic oscillations with highly variable emission observed in the prompt gamma ray emission of cosmic gamma ray bursts \\citep{mg81}. A presence of rotation and magnetic field strongly complicate the picture of the core-collapse supernova explosion \\citep{abkm05}.  In this paper we consider in details only spheroidal bodies. In reality the spheroid will become a triaxial ellipsoid during the motion. In addition to spheroids we calculate many variants with triaxial figures (see also \\citet{BKTs2005}). Qualitatively we obtain the same results for ellipsoids: no singularity was reached for any $\\varepsilon > 0$ and establishing oscillatory (regular or chaotic) regime under negative total energy prevents the collapse. However, in the case of the ellipsoid with semi-axes $a \\neq b \\neq c$ we have a system with three degrees of freedom and six-dimensional phase space. Therefore we could not carry out rigorous investigation of regular and chaotic types of motion by the constructing Poincar\\'{e} map as it was done for spheroid with two degrees of freedom, and restricted ourselfs by a description of the spheroidal case.  In the frame of a general relativity dynamic stabilization against collapse by nonlinear nonspherical oscillations cannot be universal. When the size of the body approaches gravitational radius no stabilization is possible at any $\\gamma$. Nevertheless, the nonlinear stabilization may happen at larger radii, so after damping of the oscillations the star would collapse to the black hole. Due to development of nonspherical oscillations there is a possibility for emission of gravitational waves during the collapse of nonrotating stars with the intensity similar to rotating bodies, or even larger.  Account of general relativity will introduce a new non-dimensional parameter, which can be written as $p_g = \\frac{2 G m_0}{c^2 a_0}$. The fate of the gravitating body will depend on the value of this parameter, and we may expect a direct relativistic collapse to a black hole at increasing $p_g$, approaching unity. It is known, that a nonrotating black hole is characterized only by its mass \\citep{ZN}. In absence of other dissipative processes, the excess of energy, connected with a nonspherical motion should be emitted by gravitational waves during a formation of the black hole.  \\begin{figure} \\centerline{\\hbox{\\includegraphics[width=0.5\\textwidth]{Gamma6-5.eps}}} \\caption{The Poincar\\'{e} map for one chaotic and four regular trajectories in case of $\\gamma=6/5$, $m=1$, $\\tilde H = - 3/50$, $\\tilde\\varepsilon=27/50$, see (\\ref{gamma6}). The $(a,\\dot a)$ values are taken in the minimum of c. Full black line is the bounding curve. The point inside the regular region corresponds to coherent oscillations with the same period for $a$ and $c$ values, similar to those represented in Fig.2.} \\end{figure}"
    },
    "0801/0801.0017_arXiv.txt": {
        "abstract": "The collisionally pumped, ground-state 1720\\,MHz maser line of OH is widely recognized as a tracer for shocked regions and observed in star forming regions and supernova remnants. Whereas some lines of excited states of OH have been detected and studied in star forming regions, the subject of excited-state OH in supernova remnants -- where high collision rates are to be expected -- is only recently being addressed. Modeling of collisional excitation of OH demonstrates that 1720, 4765 and 6049\\,MHz masers can occur under similar conditions in regions of shocked gas. In particular, the 6049 and 4765\\,MHz masers become more significant at increased OH column densities where the 1720\\,MHz masers begin to be quenched. In supernova remnants, the detection of excited-state OH line maser emission could therefore serve as a probe of regions of higher column densities. Using the Very Large Array, we searched for excited-state OH in the 4.7, 7.8, 8.2 and 23.8\\,GHz lines in four well studied supernova remnants with strong 1720\\,MHz maser emission (Sgr\\,A\\,East, W\\,28, W\\,44 and IC\\,443). No detections were made, at typical detection limits of around 10 mJy\\,beam$^{-1}$. The search for the 6\\,GHz lines were done using Effelsberg since the VLA receivers did not cover those frequencies, and are reported on in an accompanying letter \\citep{fish07}. We also cross-correlated the positions of known supernova remnants with the positions of 1612\\,MHz maser emission obtained from blind surveys. No probable associations were found, perhaps except in the Sgr\\,A\\,East region. The lack of detections of excited-state OH indicates that the OH column densities suffice for 1720\\,MHz inversion but not for inversion of excited-state transitions, consistent with the expected results for C-type shocks. ",
        "introduction": "Masers observed in the 1720\\,MHz satellite line of OH are often associated with supernova remnants (SNRs). They originate in the shocked region where the expanding SNR collides with a molecular cloud. During the collision, a non-dissociative C-type shock can produce the temperature and density conditions required for 1720\\,MHz maser emission to occur \\citep{wardle99, lockett99}. The C-type shock model requires shock speeds of the order of 25 km\\,s$^{-1}$, and shock chemistry predicts the conditions to be right for masers behind the shock wave. Such velocities and spatial positions are in good agreement with previous observations of SNRs \\citep[e.g.,][]{claussen97, yusef-zadeh03b, frail98}. A natural, next step would be to perform very long baseline interferometric (VLBI) observations to measure the physical sizes of the masing regions, and to track the SNR expansion with time. However, a major hurdle at 1720\\,MHz is the angular broadening due to interstellar scattering. For example, in the direction of the Galactic center, pronounced scattering ($\\Theta_{obs}\\sim 500$ mas at 1.6\\,GHz) has been measured \\citep{vanlangevelde92}. To overcome this problem, \\citet{hoffman03} conducted MERLIN and VLBA observations of the 1720\\,MHz masers in IC\\,443, positioned at a Galactic longitude believed to be little affected by scattering. In another project, \\citet{claussen02} used a nearby pulsar to estimate the interstellar scattering effect close to W\\,28.  A more general method that could work at any position on the sky would be using higher frequency transitions tracing the same type of gas (shocked by the SNR/cloud collisions). Higher frequency masers are less susceptible to scattering effects, which scale as $\\lambda^{2}$. Finding higher frequency lines would be particularly interesting for the Sgr\\,A\\,East region, since some 1720\\,MHz masers found here have velocities very offset from what is expected in the SNR model \\citep{yusef-zadeh96, karlsson03}. These offset velocity masers coincide on the sky with the circumnuclear disk and might therefore arise under different conditions. Proper motion studies of masers in the Sgr\\,A complex would thus indicate whether the kinematics really differ for the offset velocity masers and the Sgr\\,A\\,East masers. Higher frequency excited-state OH masers thus could be used for astrometric and proper motion VLBI studies along the heavily scattered Galactic plane and in the Galactic center.  Modeling of the three lower rotational states of OH has demonstrated that in star forming regions (SFRs) satellite line 1720 and 4765\\,MHz, and mainline 6035\\,MHz masers can occur under similar conditions in regions of shocked gas \\citep{gray91, gray92}. The 4765 and 6035\\,MHz masers become more significant at increased column densities where the 1720\\,MHz masers begin to switch off. Detection of excited-state OH maser emission could therefore serve as a probe of regions of higher OH column densities in SNRs ($N_{\\rm OH}\\sim 3\\times 10^{17}$\\,cm$^{-2}$ instead of $N_{\\rm OH}\\sim 3\\times 10^{16}$\\,cm$^{-2}$ for the 1720\\,MHz masers). However, calculations of the $N_{\\rm OH}$ resulting from C-type shocks are only able to produce relatively low values, $N_{\\rm OH} \\simeq 10^{16}$\\,cm$^{-2}$ \\citep{lockett99, wardle99} which may be a limiting factor in the formation of excited-state OH masers in SNRs. It should be noted that detailed observations of OH and H$_2$O in IC\\,443 require an OH production scenario involving both J- and C-type shocks \\citep{snell05, hewitt06}, indicating that estimating the OH column density could be more complicated than previously thought.  In this work we present the details and results of our search for higher frequency OH transitions in SNRs with the Very Large Array (VLA). We also report on the results of a literature search for collisionally excited 1612\\,MHz masers associated with SNRs.   \\begin{deluxetable*}{llllrccr}[t] \\tabletypesize{\\footnotesize} \\tablecolumns{8} \\tablewidth{0pt} \\tablecaption{Observational parameters} \\tablehead{\\colhead{Pointing}                     & \\multicolumn{2}{c}{RA   (J2000) Dec (J2000)}   & \\colhead{Transition}                          & \\colhead{V$_{\\rm sys}$\\tablenotemark{a}}                        & \\colhead{$\\Delta$V$_{\\rm ch}$\\,\\tablenotemark{b}}        & \\colhead{Beam size\\tablenotemark{c}}           & \\colhead{Mean rms\\tablenotemark{c}}           \\\\ \\colhead{}                                    & \\colhead{{h}\\phn{m}\\phn{s}}                   & \\colhead{\\phn{\\arcdeg}~\\phn{\\arcmin}~\\phn{\\arcsec}} & \\colhead{MHz}                                & \\colhead{km~s$^{-1}$}                         & \\colhead{km~s$^{-1}$}                         & \\colhead{\\phn{\\arcsec}$\\times$\\phn{\\arcsec}}  & \\colhead{mJy beam$^{-1}$}                      }  \\startdata IC443 & 06 16 43.61 & $+$22 32 39.78 & 4750.7, 4660.2, 4765.6&$-$4.6&0.9 & 5.1$\\times$4.3 & 13.2 \\\\ &    &  & 7761.7, 7820.1, 7832.0, 7749.9  & $-$4.6 & 0.6 & 3.0$\\times$2.8 & 8.4 \\\\ &    &  & 8135.9, 8189,6, 8207.4, 8118.1  & $-$4.6 & 0.5 & 3.0$\\times$2.7 & 6.5 \\\\ Sgr\\,A\\,East & 17 45 44.31 & $-$29 01 18.34 & 4750.7 & 64.0 & 0.8 & 7.7$\\times$3.4 & 13.8 \\\\ &              &                 & 4750.7 & 132.0 & 0.8 & 7.7$\\times$3.4 & 14.5 \\\\ &    &                   & 4660.2, 4765.6 & 64.0 & 0.8 & 7.7$\\times$3.4 & 14.1 \\\\ &    &  & 7761.7, 7820.1, 7832.0, 7749.9 &  64.0 & 0.5 & 7.0$\\times$2.5 & 11.4 \\\\ &    &  & 8135.9, 8189,6, 8207.4, 8118.1 &  64.0 & 0.5 & 5.7$\\times$2.6 & 9.5 \\\\ & & & 23817.9, 23826.6, 23838.9, 23805.3 &  64.0 & 0.7 & 1.4$\\times$0.8 & 9.4 \\\\ Sgr\\,A\\,East & 17 45 40.62 & $-$28 59 43.98 & 4750.7 & 64.0 & 0.8 & 5.8$\\times$3.9 & 15.1 \\\\ &              &                 & 4750.7 & 132.0 & 0.8 & 5.8$\\times$3.9 & 14.1 \\\\ &    &                   & 4660.2, 4765.6 & 132.0 & 0.8 & 5.8$\\times$3.9 & 14.6 \\\\ &    &  & 7761.7, 7820.1, 7832.0, 7749.9  & 132.0 & 0.5 & 5.4$\\times$2.6 & 11.9 \\\\ &    &  & 8135.9, 8189,6, 8207.4, 8118.1  & 132.0 & 0.5 & 5.3$\\times$2.5 & 9.5 \\\\ & & & 23817.9, 23826.6, 23838.9, 23805.3 & 132.0 & 0.7 & 1.2$\\times$0.9 & 9.5 \\\\ W28A &18 00 45.55&$-$23 17 43.33& 4750.7 & 6.3  & 0.9 & 5.9$\\times$4.0 & 8.6 \\\\ &              &                & 4750.7 & 76.3 & 0.9 & 5.9$\\times$4.0 & 8.8 \\\\ &              &        & 4660.2, 4765.6 &  6.3 & 0.9 & 6.2$\\times$4.2 & 8.7 \\\\ &    &  & 7761.7, 7820.1, 7832.0, 7749.9 &  6.3 & 0.6 & 4.2$\\times$2.6 & 9.1 \\\\ &    &  & 8135.9, 8189,6, 8207.4, 8118.1 &  6.3 & 0.6 & 4.5$\\times$2.6 & 7.4 \\\\ W28CD&18 01 39.35&$-$23 25 01.97 & 4750.7 & 14.1 & 0.9 & 5.7$\\times$4.0 & 9.2 \\\\ &              &                 & 4750.7 & 84.1 & 0.9 & 5.9$\\times$4.1 & 8.0 \\\\ &              &         & 4660.2, 4765.6 & 14.1 & 0.9 & 6.2$\\times$3.9 & 8.6 \\\\ &     &  & 7761.7, 7820.1, 7832.0, 7749.9 & 14.1 & 0.6 & 4.8$\\times$2.6 & 9.1 \\\\ &     &  & 8135.9, 8189,6, 8207.4, 8118.1 & 14.1 & 0.5 & 4.1$\\times$2.8 & 7.2 \\\\ W28EF&18 01 51.64&$-$23 18 21.02 & 4750.7 & 11.7 & 0.9 & 6.7$\\times$3.6 & 8.3 \\\\ &              &                 & 4750.7 & 81.7 & 0.9 & 6.9$\\times$3.5 & 9.2 \\\\ &     &                  & 4660.2, 4765.6 & 11.7 & 0.9 & 7.2$\\times$4.3 & 8.9 \\\\ &     &  & 7761.7, 7820.1, 7832.0, 7749.9 & 11.7 & 0.6 & 4.6$\\times$2.8 & 9.1 \\\\ &     &  & 8135.9, 8189,6, 8207.4, 8118.1 & 11.7 & 0.5 & 5.0$\\times$2.6 & 7.1 \\\\ W44EF\\tablenotemark{d} & 18 56 33.02 & $+$01 27 54.40 & 4750.7 & 6.3  & 0.9 & 5.3$\\times$3.9 & 8.9 \\\\ &              &                 & 4750.7 & 76.3 & 0.9 & 5.3$\\times$3.9 & 10.3 \\\\ &    &                   & 4660.2, 4765.6 &  6.3 & 0.9 & 6.2$\\times$3.7 & 8.9 \\\\ W44E\\tablenotemark{d}  & 18 56 29.14 & +01 29 14.15 & 7761.7, 7820.1, 7832.0, 7749.9 & 6.3 & 0.6 &3.9$\\times$2.5&8.1 \\\\ &    &  & 8135.9, 8189,6, 8207.4, 8118.1 &  6.3 & 0.6 & 3.9$\\times$2.4 & 6.3 \\\\ W44F\\tablenotemark{d} & 18 56 36.90 & $+$01 26 34.65 & 7761.7, 7820.1, 7832.0, 7749.9 & 6.3 & 0.6 &4.4$\\times$2.4&8.0 \\\\ &    &  & 8135.9, 8189,6, 8207.4, 8118.1 &  6.3 & 0.6 & 4.1$\\times$2.4 & 6.5 \\\\ \\enddata  \\tablenotetext{a}{The observed bandwidth was centered on the systemic LSR velocity, and covered $\\sim$96, 58, 55 and 36 km\\,s$^{-1}$ for the 4.7, 7.8, 8.2 and 23.8\\,GHz lines respectively. } \\tablenotetext{b}{Velocity resolution, with no on-line Hanning smoothing applied during the observations} \\tablenotetext{c}{The transitions were imaged separately, but lines close in frequency have similar beam sizes and channel rms. The listed value is the mean of the transitions.}  \\tablenotetext{d}{At lower frequencies W44E and W44F were observed in one beam centered at their mean position, W44EF. At higher frequencies, they were observed separately as W44E and W44F.} \\end{deluxetable*} ",
        "conclusions": "For OH column densities $N_{\\rm OH}\\sim 10^{16}-10^{18}$ cm$^{-2}$, models of collisionally-pumped excited-state OH in SNRs predict inversions in the 1720, 4765 and 6049\\,MHz lines, while no significant maser optical depths will be produced in other transitions of the 4.7, 7.8, 8.2 and 23.8\\,GHz lines.  We have presented the results from a search of all these excited-state OH transitions in four, well-known SNRs, with no detections. These VLA observations could not tune to the 6049\\,MHz line, which is also predicted to have large optical depths under conditions similar to those of 1720\\,MHz masers. However, a few recent searches for the 6049\\,MHz line in SNRs have yielded no detections \\citep{fish07, mcdonnell07}. In the future, the EVLA upgrade will allow for a targeted, deeper search for 6049\\,MHz in these SNRs.  Our non-detections are consistent with regions of lower column densities ($N_{\\rm OH}\\leq5\\times10^{16}$cm$^{-2}$), where 1720\\,MHz masers are strong, and little inversion occurs in the higher transitions. Based on VLA and single-dish observations by \\citet{yusef-zadeh95} \\& \\citet{hewitt07}, such a post-shock medium may have a large filling factor. \\citet{hewitt07} found that a large part of the 1720\\,MHz maser flux is undetected using VLA baselines, indicating a widespread distribution of weaker 1720\\,MHz emission. This likely reflects large post-shock regions of low column density, or alternatively, regions of different temperature and density from what is normally expected for 1720\\,MHz maser production. In turn, such a gas component will provide little chance of excited-state OH maser emission. If this is the case, excited-state OH will not be detected either coincident with or offset from detected 1720\\,MHz masers in SNRs. The non-detections imply a low OH column density in SNRs, in agreement with the rather low estimates of the $N_{\\rm OH}\\simeq 10^{16}$\\,cm$^{-2}$ expected to be produced in C-type shocks \\citep{wardle99, lockett99}.  From these results we draw the conclusion that in the regions where 1720\\,MHz masers are found, the OH column densities are insufficient for excited-state OH masers to exist. This is supported by the absence of 1612\\,MHz masers in SNRs, which require 1--2 orders of magnitude higher column densities than the 1720\\,MHz masers. However, the upper limit to the column density must be even stricter to avoid producing 6049 and 4765\\,MHz masers."
    },
    "0801/0801.0221_arXiv.txt": {
        "abstract": "Double-peaked oxygen lines in the nebular spectra of two peculiar Type Ib/c Supernovae (SN Ib/c) have been interpreted as off-axis views of a GRB-jet or unipolar blob ejections. Here we present late-time spectra of \\nsn\\ SN IIb, Ib and Ic and show that this phenomenon is common and should not be so firmly linked to extraordinary explosion physics. The line profiles are most likely caused by ejecta expanding with a torus- or disk-like geometry. Double-peaked oxygen profiles are not necessarily the indicator of a mis-directed GRB jet. ",
        "introduction": "To understand the connection between long-duration Gamma Ray Bursts (GRBs) and the supernovae (SNe) associated with them, it would be significant to detect the effects of jets that are produced in the explosion even when the observer is not in the relativistic beam \\citep{woosley06_rev}). In the broad-lined Type Ic, SN 2003jd, an off-axis GRB jet was proposed as the cause for the double-peaked oxygen profile observed in its nebular spectrum \\citep{mazzali05_03jd,valenti08}. A similar double-peaked oxygen profile, with a large blueshift, was detected in the peculiar Type Ib, SN 2005bf \\citep{maeda07,modjaz07_thesis}, where it was interpreted as coming from a unipolar blob or jet \\citep{maeda07}.  If the double-peaked oxygen shape were exclusively caused by highly relativistic jets, then we might expect only broad-lined SN Ic, whose line widths approach 30,000 \\kms , and which are the only subtype of stripped-envelope SN seen at the sites of GRBs to exhibit them. However, if we see double-peaked profiles in a range of supernova types, we might conclude that asphericities are present in more typical core-collapse events.  We set out to obtain a set of nebular spectra for supernovae of various subtypes to measure the emission line profiles with adequate resolution and signal to detect signs of these double-peaked profiles in nebular spectra of supernovae. As the supernova turns optically thin a few months after maximum light, the emission line shapes can provide information on the velocity distribution of the ejecta \\citep{matheson01,foley03,maeda06}. Structure in emission line profiles have been reported in SN II \\citep{spyromilio91}, SN IIb \\citep{spyromilio94_93j,matheson00_93jdetail} and SN Ib \\citep{sollerman98,elmhamdi04}, and interpreted as clumping of oxygen in the ejecta.   In \\S~\\ref{obs_sec} we describe the observational sample. In \\S~\\ref{doubleOxy_sec} we discuss the line profiles and show that the double-peaked line profiles are unlikely to be caused by optical depth effects. We speculate on the implications of our results in \\S~\\ref{interpretation_sec} and summarize in \\S~\\ref{conclusion_sec}. ",
        "conclusions": "In summary, we detect clear double-peaked line profiles of \\OxyOne\\ in three SN IIb and SN Ib/c out of our observed sample of \\nsn\\ and additionally, in one out of four published SN Ib. Our sample does not include broad-lined SN Ic or peculiar objects.  These line profiles are probably not caused by optical depth effects, and suggest global anisotropies in the ejecta. Prior to this work, double-peaked oxygen lines had only been reported in two peculiar SN Ib/c and interpreted as off-axis GRB-jet or unipolar blob ejections. It seems more likely that asphericities are present in a wide variety of core-collapse events and they are not strictly confined to the supernovae associated with GRB jets. Although special models have been proposed to account for the line profiles in peculiar supernovae, our investigation suggests that double-peaked profiles, and underlying disks, are not unusual. These ordinary SN from our sample have a variety of subtypes, reflecting the diversity of mass loss prior to the explosion.  We recommend a meta-analysis of available nebular spectra of all core-collapse SN, spanning the full range from SN II to IIb, Ib, Ic and finally, to broad-lined SN Ic, in order to quantify the kinematics and geometry of the ejecta and to seek trends as a function of SN type. Further high-SN/N multi-epoch observations of a sample of SN II/IIb/Ib/Ic coupled with radiative transfer models should help to elucidate the observed blue- and redshifts of the line profiles."
    },
    "0801/0801.4185.txt": {
        "abstract": "The Stokes $V$ parameter characterizes asymmetry of amplitudes between right- and left-handed waves, and non-vanishing value of the $V$ parameter yields a circularly polarized signal. Cosmologically, $V$ parameter may be a direct probe for parity violation in the  universe. In this paper, we theoretically investigate a measurement of this parameter, particularly focusing on the gravitational-wave backgrounds observed via ground-based interferometers. In contrast to the traditional analysis that only considers the total amplitude (or equivalently $\\Omega_{\\rm GW}$), the signal analysis including a circular-polarized mode has a rich structure due to the multi-dimensionality of target parameters. We show that, by using the network of next-generation detectors, separation between polarized and unpolarized modes can be performed with small statistical loss induced by their correlation. ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   Because of the extremely weak signal, a direct detection of gravitational waves is a technically challenging issue, and we have not yet succeeded the direct detection despite extensive efforts.  Nevertheless, the weakness of the gravitational interaction may be a great advantage for astronomy and cosmology, because gravitational waves can propagate to us from very early universe almost without scattering and absorption \\cite{Thorne_K:1987,Cutler:2002me}. In this respect, stochastic background of gravitational waves is one of the most important targets for gravitational wave astronomy \\cite{Allen:1996vm}. If detected, the stochastic background will serve as an invaluable fossil to study the physics at extremely high-energy scale for which other methods cannot be attainable.    Over the last decade, sensitivity of gravitational wave detectors to the stochastic background has drastically improved. We will soon reach at the sensitivity level $\\Omega_{\\rm GW}\\lsim 10^{-5}$ around 100Hz \\cite{Abbott:2006zx}, where $\\Omega_{\\rm GW}$ is the energy density of the gravitational waves normalized by critical density of the universe. This level is below the indirect cosmological constraints, such as derived from the observed abundance of light elements \\cite{Allen:1996vm} (see \\cite{Smith:2006nka} for the constraints from cosmic microwave background), and in this sense, gravitational wave detectors will provide a unique opportunity to directly constrain the early universe. In order to further get a stringent constraint and/or valuable information from the next-generation detectors, one important approach is to improve the statistical analysis of gravitational wave backgrounds. So far,  most of theoretical studies on the gravitational-wave backgrounds have been directed  to its energy  spectrum (for its anisotropies, see, {\\it e.g.,} \\cite{Giampieri:1997ie}). The authors recently provided a brief sketch for measurement of the Stokes $V$ parameter of the gravitational-wave background via correlation analysis of ground-based detectors \\cite{Seto:2007tn} (see \\cite{Lue:1998mq} for cosmic microwave background and \\cite{Seto:2006hf} for space gravitational wave detectors such as  LISA \\cite{lisa} or BBO/DECIGO \\cite{bbo,Seto:2001qf}).  The Stokes $V$ parameter may be basic observable to quantify the parity violation process. One of such parity violation process is through the Chern-Simons term that might be originated from string theory \\cite{Alexander:2004us}. This paper is a follow-up study to the preceding short report. In addition to detailed explanations and supplementary materials to the previous paper, we developed a new statistical framework to deal with multiple parameters of gravitational wave backgrounds, and we specifically applied it to simultaneous estimation of amplitudes of both the unpolarized and polarized modes of the gravitational-wave background.   This paper is organized as follows. In section \\ref{sec:circular_pol}, we describe polarization decomposition of a gravitational-wave  background,  and define its Stokes $V$ parameter. The basic framework to treat polarized gravitational waves is essentially the same one as in the case of electromagnetic waves \\cite{radipro}. In section \\ref{sec:overlap_func}, we explain the correlation analysis for the gravitational-wave background and introduce the overlap functions that characterize sensitivities to the polarized and unpolarized modes.  Then, we discuss basic properties of the overlap functions, and calculate them for the planed next-generation detectors, such as advanced LIGO. In section \\ref{sec:broadband_SNR}, broadband analysis of the gravitational-wave background is studied, taking into account the measurement noises for each detector. In section \\ref{sec:separation}, we discuss how well we can separately measure the polarized and unpolarized modes.  In contrast to the traditional arguments only for the unpolarized mode, the situation considered here is more complicated. We provide a statistical framework to analyze multiple parameters of the stochastic background with correlation analysis. Finally, section \\ref{sec:summary} is a brief summary of this paper. Appendix \\ref{sec:tensor_analysis} presents the derivation for the expressions of the overlap functions.  This geometrical derivation is similar to that given in Ref.\\cite{Flanagan:1993ix}. In appendix \\ref{sec:PDF_for_corr}, we discuss the probability distribution functions (PDFs) of basic observational quantities with correlation analysis. In appendix \\ref{sec:derivation_optimal_SNR}, we derive the formal expressions for optimal signal-to-noise ratios for detectors more than four. In appendix \\ref{sec:moon}, we comment on the surface of the Moon as potential sites for laser interferometers.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec:summary} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   In this paper we present prospects for measuring the Stokes $V$ parameter of  stochastic gravitational wave backgrounds via the correlation analysis.  This parameter characterizes the asymmetry of amplitudes of right- and left-handed waves.  As the parity transformation interchanges the two polarization modes,   it can be regarded as the basic observational measure to probe parity violation.  We  made detailed analyses for the basic properties of the overlap functions $\\gamma_I$ and $\\gamma_V$, especially their dependencies on geometry of detector configurations.  In contrast to studies only for the unpolarized $I$ mode (equivalently,  the energy spectrum $\\Omega_{\\rm GW}$),  we need to develop a new statistical framework to deal with rich  structures caused by  multi-dimensionality of target parameters.   We provide an optimal method that will be applicable to  various problems of gravitational-wave backgrounds.  Based on our new method, we estimated sensitivities of the planned and proposed next-generation interferometers to the  $V$ modes.   We  found that it is important to have a large number of detectors in order to reduce possible effects due to correlation between target parameters.   \\bigskip We would like to thank M. Ando, N. Kanda and M. Ohashi for useful conversations. This work was in part supported by a Grant-in-Aid for Scientific Research from the Japan Society for the Promotion of Science (No.~18740132).  \\appendix %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0801/0801.2965_arXiv.txt": {
        "abstract": "In this paper we discuss the properties of the quasi-steady state cosmological model (QSSC) developed in 1993 in its role as a cyclic model of the universe driven by a negative energy scalar field.  We discuss the origin of such a scalar field in the primary creation process first described by F. Hoyle and J. V. Narlikar forty years ago.  It is shown that the creation processes which takes place in the nuclei of galaxies are closely linked to the high energy and explosive phenomena, which are commonly observed in galaxies at all redshifts.  The cyclic nature of the universe provides a natural link between the places of origin of the microwave background radiation (arising in hydrogen burning in stars), and the origin of the lightest nuclei (H, D, He$^3$ and He$^4$).  It also allows us to relate the large scale cyclic properties of the universe to events taking place in the nuclei of galaxies.  Observational evidence shows that ejection of matter and energy from these centers in the form of compact objects, gas and relativistic particles is responsible for the population of quasi-stellar objects (QSOs) and gamma-ray burst sources in the universe.  In the later parts of the paper we briefly discuss the major unsolved problems of this integrated cosmological and cosmogonical scheme.  These are the understanding of the origin of the intrinsic redshifts, and the periodicities in the redshift distribution of the QSOs. ",
        "introduction": "\\subsection{Cosmological models}  The standard cosmological model accepted by the majority at present is centered about the big bang which involves the creation of matter and energy in an initial explosion.  Since we have overwhelming evidence that the universe is expanding, the only alternative to this picture appears to be the classical steady-state cosmology, of Bondi, Gold and Hoyle, (Bondi and Gold, 1948, Hoyle, 1948) or a model in which the universe is cyclic with an oscillation period which can be estimated from observation.  In this latter class of model the bounce at a finite minimum of the scale factor is produced by a negative energy scalar field.  Long ago Hoyle and Narlikar (1964) emphasized the fact that such a scalar field will produce models which oscillate between finite ranges of scale.  In the 1960s theoretical physicists shied away from scalar fields, and more so those involving negative energy.  Later Narlikar and Padmanabhan (1985) discussed how the scalar creation field helps resolve the problems of singularity, flatness and horizon in cosmology.  It now appears that the popularity of inflation and the so-called new physics of the 1980s have changed the 1960s' mind-set.  Thus Steinhardt and Turok (2002) introduced a negative potential energy field and used it to cause a bounce from a non-singular high density state.  It is unfortunate that they did not cite the earlier work of Hoyle and Narlikar which had pioneered the concept of non-singular bounce through the agency of a negative energy field, at a time when the physics community was hostile to these ideas. Such a field is required to ensure that matter creation does not violate the conservation of matter and energy.  Following the discovery of the expansion of the universe by Hubble in 1929, practically all of the theoretical models considered were of the Friedmann type, until the proposal by Bondi, Gold and Hoyle in 1948 of the classical steady state model which first invoked the creation of matter.  A classical test of this model lay in the fact that, as distinct from all of the big bang models, it predicted that the universe must be accelerating (cf Hoyle and Sandage, 1956). For many years it was claimed that the observations indicated that the universe is decelerating, and that this finding disproved the steady state model.  Not until much later was it conceded that it was really not possible to determine the deceleration parameter by the classical methods then being used.  Gunn and Oke (1975) were the first to highlight the observational uncertainties associated with this test. Of course many other arguments were used against the classical steady state model (for a discussion of the history see Hoyle, Burbidge and Narlikar 2000 Chapters 7 and 8). But starting in 1998 studies of the redshift-apparent magnitude relation for supernovae of Type 1A showed that the universe $is$ apparently accelerating (Riess et al. 1998, Perlmutter et al. 1999).  The normal and indeed the proper way to proceed after this result was obtained should have been at least to acknowledge that, despite the difficulties associated with the steady state model, this model had all along been advocating an accelerating universe.  It is worth mentioning that McCrea (1951) was the first to introduce vacuum related stresses with equation of state $p = -\\rho$ in the context of the steady state theory.  Later Gliner (1970) discussed how vacuum-like state of the medium can serve as original (non singular) state of a Friedmann model.  The introduction of dark energy is typical of the way the standard cosmology has developed; viz,  a new assumption is introduced specifically to sustain the model against some new observation.  Thus, when the amount of dark matter proved to be too high to sustain the primordial origin of deuterium, the assumption was introduced that most of the dark matter has to be non-baryonic.  Further assumptions about this dark matter became necessary, e.g., cold, hot, warm, to sustain the structure formation scenarios.  The assumption of inflation was introduced to get rid of the horizon and flatness problems and to do away with an embarrassingly high density of relic magnetic monopoles. As far as the dark energy is concerned, until 1998 the general attitude towards the cosmological constant was typically as summarized by Longair in the Beijing cosmology symposium: ``None of the observations to date require the cosmological constant\" (Longair 1987).  Yet, when the supernovae observations could not be fitted without this constant, it came back with a vengeance as dark energy.  Although the popularity of the cosmological constant and dark energy picked up in the late 1990s, there had been earlier attempts at extending the Friedmann models to include effects of vacuum energy.  A review of these models, vis-a-vis observations may be found in the article by Carroll and Press (1992).  We concede that with the assumptions of dark energy, non-baryonic dark matter, inflation etc. an overall self consistent picture has been provided within the framework of the standard model.  One demonstration of this convergence to self consistency is seen from a comparison of a review of the values of cosmological parameters of the standard model by Bagla, et al. (1996), with the present values. Except for the evidence from high redshift supernovae, in favour of an accelerating universe which came 2-3 years later than the above review, there is an overall consistency of the picture within the last decade or so, including a firmer belief in the flat ($\\Omega = 1$) model with narrower error bars.  Nevertheless we also like to emphasize that the inputs required in fundamental physics through these assumptions have so far no experimental checks from laboratory physics.  Moreover an epoch dependent scenario providing self-consistency checks, e.g. CMB anisotropies, cluster baryon fraction as a function of redshift does not meet the criterion of `repeatability of scientific experiment'.  We contrast this situation with that in stellar evolution where stars of different masses constitute repeated experimental checks on the theoretical stellar models thus improving their credibility.  Given the speculative nature of our understanding of the universe, a sceptic of the standard model is justified in exploring an alternative avenue wherein the observed features of the universe are explained with fewer speculative assumptions.  We review here the progress of such an alternative model.  In this model creation of matter is brought in as a physical phenomenon and a negative kinetic energy scalar field is required to ensure that it does  not violate the law of conservation of matter and energy. A simple approach  based on Mach's principle leads naturally to such a field within the curved spacetime of general relativity described briefly in $\\S$ 2. The resulting field equations have the two simplest types of solutions for a homogeneous and isotropic universe:  (i) those in  which the  universe oscillates but there is no creation of matter,  and (ii) those in which the universe steadily expands with a constant value  of $H_o$ being driven by continuous creation of  matter.  The simplest model including features of both these solutions is  the {\\it Quasi-Steady  State Cosmology} (QSSC), first proposed  by Hoyle, Burbidge and Narlikar (1993). It has the scale factor in the form:   \\begin{equation} S (t) = \\exp \\bigg({t\\over P}\\bigg) \\{1 + \\eta \\cos \\theta (t)\\}, ~\\theta (t) \\approx \\frac{2 \\pi t}{Q}, \\end{equation}  \\noindent where $P$ is the long term `steady state' time scale of expansion while $Q$ is the period of a single oscillation.  Note that it is essential for the universe to have a long term expansion; for a universe that has only oscillations without long term expansion would run into problems like the Olbers paradox.  It is also a challenge in such a model to avoid running into `heat death' through a steady increase of entropy from one cycle to next. These difficulties are avoided if there is creation of new matter at the start of each oscillation as happens in the QSSC, and also, if the universe has a steady long term expansion in addition to the oscillations.  New matter in such a case is of low entropy and the event horizon ensures a constant entropy within as the universe expands.  The  QSSC  has an additional attractive feature if one  uses  the criterion   of  the  Wheeler  and  Feynman  absorber  theory   of electromagnetic   radiation  (Wheeler  and  Feynman, 1945, 1949). This theory provided a very natural explanation of why in actuality the electromagnetic signals propagate into the future, i.e., via retarded solutions, despite the time-symmetry of the basic equations.  By writing the theory in a relativistically invariant action-at-a-distance form, Wheeler and Feynman showed that suitable absorptive properties of the universe can lead to the breaking of time-symmetry.  As was discussed by Hogarth (1962) and later by Hoyle and Narlikar (1963, 1969, 1971) who also extended the argument to quantum  electrodynamics,  the Wheeler-Feynman theory gives results consistent with observations only if the past absorber is imperfect and the future absorber is perfect.   This requirement is $not$ satisfied by a  simply cyclic universe or by an ever-expanding big bang universe but is satisfied by the QSSC because of expansion  being coupled with cyclicity.  One may question as to why one needs to have the Wheeler-Feynman approach to electrodynamics in preference to field theory.  The advantages are many, including (i) a satisfactory explanation of the Dirac formula of radiative reaction, (ii) the unambiguous deduction of why one uses retarded solutions in preference to advanced ones and (iii) a resolution of the ultraviolet divergences in quantum electrodynamics.  Rather than go into these aspects in detail we refer the reader to a recent review by Hoyle and Narlikar (1995).  Since cosmology seeks to deal with the large-scale properties of the universe, it inevitably requires a strong connection with fundamental physics.  In the big bang cosmology particle physics at very high energy is considered very relevant towards understanding cosmology.  In the same spirit we believe that the action at a distance approach to fundamental physics brings about an intimate link of microphysics with cosmology.  The Wheeler-Feynman approach is an excellent demonstration of such a connection.  \\subsection{Cosmogony}  In this paper we shall discuss this cosmological model, but first we want to indicate the importance of the observed behavior of the galaxies (the observed cosmogony) in this approach.  Now that theoretical cosmologists have begun to look with favor on the concepts of scalar negative energy fields, and the creation process, they have taken the position that this subject can only be investigated by working out models based on  classical approaches of high energy physics and their effects on the global scale.  In all of the discussions of what is called precision cosmology there is no discussion of the remarkable phenomena which have been found in the comparatively nearby universe showing that galaxies themselves can eject what may become, new galaxies. We believe that only when we really understand how individual galaxies and clusters etc. have formed, evolve, and die (if they ever do) shall we really understand the overall cosmology of the universe.  As was mentioned earlier, the method currently used in the standard model is to suppose that initial quantum fluctuations were present at an unobservable epoch in the early universe, and then try to mimic the building of galaxies using numerical methods, invoking the dominance of non-baryonic matter and dark energy for which there is no independent evidence.  In one sense we believe that the deficiency of the current standard approach is already obvious.  The model is based on only some parts of the observational data.  These are:  all of the details of the microwave background, the abundances of the light elements, the observed dimming of distant supernovae, and the large scale distribution of the observed galaxies.  This has led to the conclusion that most of the mass-energy making up the universe has properties which are completely unknown to physics.  This is hardly a rational position, since it depends heavily on the belief that all of the laws of physics known to us today can be extrapolated back to scales and epochs where nothing is really testable; and that there is nothing new to be learned.  In spite of this, a very persuasive case has been made that all of the observational parameters can be fitted together to develop what is now becoming widely accepted as a new standard model, the so-called $\\Lambda$CDM model (Spergel et al., 2003).   There have been some publications casting doubt on this model, particularly as far as the reality of dark energy and cold, dark matter are concerned (Meyers et al. 2004; Blanchard et al. 2003).  It is usual to dismiss them as controversial and to argue that a few dissenting ideas on the periphery of a generally accepted paradigm are but  natural.  However,  it is unfortunately the case that a large fraction of our understanding of the extragalactic universe is being based on the belief that there \\textit{was} a beginning and an inflationary phase, and that the seeds of galaxies all originate from that very early phase.  We believe that an alternative approach should be considered and tested by observers and theorists alike. In this scheme the major themes are (1) that the universe is cyclic and there was no initial big bang, and (2) \\textit{all} of the observational evidence should be used to test the model.  As we shall show, this not only includes the observations which are used in the current standard model, but also the properties and interactions of galaxies and QSOs which are present in the local ($z < 0.1$) universe.  Possibly the most perceptive astronomer in recent history was Viktor Ambartsumian the famous Armenian theorist.  Starting in the 1950s and 1960s (Ambartsumian, 1965) he stressed the role of explosions in the universe arguing that the associations of galaxies (groups, clusters, etc.) showed a tendency to expand with far larger kinetic energy than is expected by assuming that the gravitational virial condition holds.  We shall discuss the implications of the cluster dynamics in Section 6. Here we take up the issue emphasized by Ambartsumian that there apparently exist phenomena in nuclei of galaxies where matter seems to appear with large kinetic energy of motion directed outwards. In Section 6 we will also include other phenomena that share the same property, namely explosive creation of matter and energy.  We shall refer to such events as mini-creation events.  Since these phenomena appear on the extragalactic scale and involve quasi-stellar objects, active galaxies, powerful radio sources and clusters and groups of galaxies at all redshifts, we believe they must have an intimate connection with cosmology. Indeed, if one looks at standard cosmology, there too the paradigm centers around the `big bang' which is itself an explosive creation of matter and energy. In the big bang scenario the origin of all of the phenomena is ultimately attributed to a single origin in the very early universe.  No connection has been considered by the standard cosmologists between this primordial event and the {\\it mini-creation events} (MCEs, hereafter) that Ambartsumian talked about.  In fact, the QSOs and AGN are commonly ascribed to supermassive black holes as `prime movers'. In this interpretation the only connection with cosmology is that it must be argued that the central black holes are a result of the processes of galaxy formation in the early universe.  In  the QSSC we have been trying to relate such mini-creation events (MCEs) directly  to the  large  scale  dynamics of the universe.  We  show in Sections 2 - 4 that  the dynamics  of the universe is governed by the frequency and  power of  the MCEs, and there is a two-way feedback between  the  two. That is, the universe expands when there is a large MCE  activity and contracts when the activity is switched off.  Likewise,  the MCE activity  is  large  when the density  of  the  universe  is relatively  large and negligible when the density  is  relatively small.   In  short,  the universe oscillates  between  states  of finite maximum and minimum densities as do the creation phases in the MCEs.  This was the model proposed by Hoyle, Burbidge and Narlikar (1993) and called the {\\it quasi-steady state cosmology} (QSSC in brief). The model was motivated partly by Ambartsumian's ideas and partly by the growing number of explosive phenomena that are being discovered in extragalactic astronomy.  In the following sections we discuss the cosmological model and then turn to the various phenomena which are beginning to help us understand the basic cosmogony.  Then we discuss and look at the phenomena themselves in the framework of this cosmology.  Finally, we discuss some of the basic problems that have been uncovered by the new observations for which no theoretical explanation has so far been proposed. ",
        "conclusions": "The oscillating universe in the QSSC, together with a long-term expansion, driven by a population of mini-creation events provides the missing dynamical connection between cosmology and the `local' explosive phenomena. The QSSC additionally fulfills the roles normally expected of a cosmological theory, namely (i) it provides an explanation of the cosmic microwave background with temperature, spectrum and inhomogeneities related to astrophysical processes (Narlikar et al. 2003), (ii) it offers a purely stellar-based interpretation of all observed nuclei ({\\it including} light ones)(Burbidge et al. 1957; Burbidge and Hoyle 1998), (iii) it generates baryonic dark matter as part of stellar evolution (Hoyle et al. 1994), (iv) it accounts for the extra dimming of distant supernovae {\\it without} having recourse to dark energy (Narlikar, Vishwakarma and Burbidge 2002; Vishwakarma and Narlikar 2005), and it also suggests a possible role of MCEs in the overall scenario of structure formation (Nayeri et al. 1999).  The last mentioned work shows that preferential creation of new matter near existing concentrations of mass can lead to growth of clustering.  A toy model based on million-body simulations demonstrates this effect and leads to clustering with a 2-point correlation function with index close to $-1.8$. Because of repulsive effect of the $C$-field, it is felt that this process may be more important than gravitational clustering.  However, we need to demonstrate this through simulations like those in our toy model, {\\it together with} gravitational clustering.  There are two challenges that still remain, namely understanding the $origin$ of anomalous redshifts and the observed $periodicities$ in the redshifts.  Given the QSSC framework, one needs to find a scenario in which the hitherto classical interpretation of redshifts is enriched further with inputs of quantum theory.  These are huge problems which we continue to wrestle with.  \\bigskip  \\noindent {\\it Acknowledgements}  One of us, (JVN) thanks College de France, Paris for hospitality when this work was in process. RGV is grateful to IUCAA for hospitality which facilitated this collaboration.   \\clearpage"
    },
    "0801/0801.0960_arXiv.txt": {
        "abstract": "The success of the SWIFT/BAT and INTEGRAL missions has definitely opened a new window for follow-up and deep study of the transient gamma-ray sky. This now appears as the access key to important progresses in the area of cosmological research and deep understanding of the physics of compact objects. To detect in near real-time explosive events like Gamma-Ray bursts, thermonuclear flashes from Neutron Stars and other types of X-ray outbursts we have developed a concept for a wide-field gamma-ray coded mask instrument working in the range 8-200 keV, having a sensitivity of 0.4~ph~cm$^{-2}$~s$^{-1}$ in 1~s (15-150 keV) and arcmin location accuracy over a sky region as wide as 3~sr. This scientific requirement can be achieved by means of two large area, high spatial resolution CZT detection planes made of arrays of relatively large ($\\sim1$~cm$^2$) crystals, which are in turn read out as matrices of smaller pixels. To achieve such a wide Field-Of-View the two units can be placed at the sides of a S/C platform serving a payload with a complex of powerful X-ray instruments, as designed for the EDGE mission. The two units will be equipped with powerful signal read out system and data handling electronics, providing accurate on-board reconstruction of the source positions for fast, autonomous target acquisition by the X-ray telescopes. ",
        "introduction": "\\label{sect:sections}  The relevance of X- and $\\gamma$-ray broadband instrumentation has been largely established during the last decade, for the study of thermal and non-thermal physical processes in astrophysical sites, as well as their tight interaction spanning a very broad region of the electromagnetic spectrum (from the keV to the TeV range). Still very difficult remain the investigations of transient phenomena mainly due to the need of fast follow up. The sensitivity of wide field telescopes is intrinsically worse than what is needed for dedicated, sophisticated spectroscopy measurements and the time scale of the events to be studied is often in the range of seconds or minutes. Especially after the results of BeppoSAX (Boella et al, 1987) and SWIFT (Gehrels et al. 2004), it is recognized that important achievements can be obtained by short-term reaction to transient events like Gamma-Ray Bursts (GRBs), outbursts from compact objects and thermonuclear flashes from accreting Neutron Stars (NS).  Despite these phenomena have been known for long, new types of transient sources are appearing in most recent years. Among them, the Supergiant Fast X-ray Transients are a new class showing hour-long, hard X-ray outbursts (Sguera et al. 2006). These are accreting objects associated to High Mass X-ray Binaries with a Supergiant companion. Their growing number (despite the difficulty of detecting them) is suggesting they are a dominant class of binaries in our Galaxy, very interesting as they could be the progenitors of double NS binaries or NS/Black Hole binaries. Other recent, interesting class is that of Anomalous X-ray Pulsars (AXP), which have been discovered to undergo short, extremely hard outbursts (Kaspi 2007, Sguera et al. 2007).  Flaring behaviour of Galactic compact sources at time scales of hours to min is, therefore, a rather established but poorly studied phenomenon. Other fast and strong events have been found in Black Hole (BH) systems, like rapid flux variability of V4641 SGR (e.g. Wijnands \\& Van der klis 2000) and possibly GRB~070610, likely to be associated to a stellar black hole (Kasliwal et al. 2007). We also list giant outbursts of Cyg X-1 (Golenetskii et al. 2003) and bright flares from XTE J1650-500 (Tomsick et al 2003). These are probably not related to disk instability but can be interpreted in many ways, as violent activity of the inner accretion disk (e.g. feeding of material into jets), fluctuations of direct wind accretion, presence of thick, possibly clumpy absorbers or outflows. High resolution spectroscopy of these bright transient events is therefore a tool for important diagnostics, as well as for persistent X-ray emission.  In accreting stellar Black Holes, the detection of Fe fluorescence lines with \u201cskewed\u201d profile can be explained by  relativistic effects at the inner edge of accretion disk. On the other hand, broad and narrow Fe lines have been observed from Cyg X-1 while the source was in its \u201cintermediate\u201d state, which may be interpreted as originating from reflection on both inner region and disk (Miller et al. 2004).  Fine X-ray spectroscopy is also of key importance in the study of NS.  The measurement of the gravitational redshift in observations of X-ray bursts, through detection of redshifted lines can lead to determination of the NS mass-radius relationship, hence constraining the equation of state of  the dense nuclear matter.  An important class of thermonuclear bursts are the \u201csuperbursts\u201d, i.e. very long ($\\sim$~hours) events due to unstable burning of carbon and/or other products of the hydrogen/helium burning occurring on the NS surface. Despite these events are $\\sim$~1000 times more energetic than normal type-I bursts, only a few of them are known, while it is very important to study their frequency and characteristics as a function of accretion rate (see e.g. Sthromayer \\& Bildsten 2003). Discrete components by Fe have been found in the spectrum of a superburst by 4U1820-30 (Sthromayer \\& Brown 2002).  We have developed a concept for an instrument exploiting a very wide region of the sky, to serve as an affordable, sensitive and moderate resource consuming payload complement for re-pointing capability in a fast-slewing satellite. The former success of the BeppoSAX/WFC (Jager et al. 1997) in the study of GRB and transient sources has motivated us to develop similar concept where source positioning is obtained by the use of coded masks, proven to be very successful also in the case of SWIFT and INTEGRAL (Winkler et al.~2003). This study, as reported hereafter has been targeted to the EDGE satellite (den Herder et al. 2007), to be flown on Low Earth Orbit (LEO). EDGE is a proposal for a class-M mission in the framework of the ESA Cosmic Vision Programme and carries a payload complex consisting of powerful X-ray spectroscopy and imaging telescopes, a Wide Field Monitor (WFM, described in this work) and a GRB detector.  EDGE looks as the ideal mission to exploit the above type of transient phenomena. Rapid response to X-ray transient events and capability of high resolution spectroscopy are the key to probe the extreme relativistic regime of the inner regions in BH accretion disks and physical conditions at the neutron star surface, e.g. by detection of line features in superbursts. In such a way the phenomena currently observed sporadically will be upgraded to real classes allowing to more detailed and well constrained models.  For an overview of the EDGE scientific objectives and payload see den~Herder et al. (2007). ",
        "conclusions": ""
    },
    "0801/0801.4772_arXiv.txt": {
        "abstract": "We present the results of an experiment to image the interacting binary star $\\beta$~Lyrae with data from the Navy Prototype Optical Interferometer (NPOI), using a differential phase technique to correct for the effects of the instrument and atmosphere on the interferometer phases.  We take advantage of the fact that the visual primary of $\\beta$~Lyrae and the visibility calibrator we used are both nearly unresolved and nearly centrally symmetric, and consequently have interferometric phases near zero.  We used this property to detect and correct for the effects of the instrument and atmosphere on the phases of $\\beta$~Lyrae and to obtain differential phases in the channel containing the H$\\alpha$ emission line.  Combining the H$\\alpha$-channel phases with information about the line strength, we recovered complex visibilities and imaged the H$\\alpha$ emission using standard radio interferometry methods.  We find that the results from our differential phase technique are consistent with those obtained from a more-standard analysis using squared visibilities ($V^2$'s).  Our images show the position of the H$\\alpha$ emitting regions relative to the continuum photocenter as a function of orbital phase and indicate that the major axis of the orbit is oriented along p.a.$=248.8\\pm1.7^{\\circ}$.  The orbit is smaller than previously predicted, a discrepancy that can be alleviated if we assume that the system is at a larger distance from us, or that the contribution of the stellar continuum to the H$\\alpha$ channel is larger than estimated.  Finally, we also detected a differential phase signal in the channels containing He~I emission lines at 587.6 and 706.5~nm, with orbital behavior different from that of the H$\\alpha$, indicating that it originates from a different part of this interacting system. ",
        "introduction": "%  One of the major challenges in optical intererometry is recovering images, since fringe visibility phases are badly corrupted by the atmosphere on short timescales.  A few attempts have been made to reconstruct images from the squared visibility ($V^2$), an unbiased estimator of the visibility amplitude (e.g., \\citealt{Baldwin96}, \\citealt{Quirrenbach94}), but in most cases, images in optical interferometry are produced from $V^2$ and closure phases (e.g., \\citealt{Mon07}).  A closure phase is the sum of visibility phases around a triangle of interferometer array elements; the atmospheric turbulence effects cancel in the sum, but only a fraction $1-2/N$ of the phase information from an array of $N$ elements is recovered.  The challenge in optical interferometric imaging is to recover as much of the phase information as possible.  Here we describe the technique we developed to obtain H$\\alpha$ differential phases from NPOI observations of the interacting binary star $\\beta$ Lyrae. We were able to recover complex visibilities of this source, and, for the first time, obtain H$\\alpha$ images using standard radio interferometric reconstruction techniques.  Differential phases, also known as phase referencing (see \\citealt{Quirrenbach99}, \\citealt{Mon03} for a discussion on the subject) is an interesting and powerful technique. Here one needs multiwavelength observations, obtained simultaneously, and a priori knowledge about the structure of the source in a region of the observed spectrum.  From that knowledge, one calculates the expected phases in that spectral region, corrects for the atmospheric and instrumental effects, and interpolates or extrapolates the corrections to the wavelengths of interest.  One of the first attempts to use this technique was presented by \\cite{Vakili97,Vakili98}, who used H$\\alpha$ and He~I~667.8~nm GI2T observations of P~Cygni and $\\zeta$~Tauri to detect structures in the line-emitting regions of these stars. However, they did not reconstruct images based on their measurements; their conclusions were based only on the analysis of the phases and visibilities of the stars. Another application of the differential phase technique was presented by \\cite{akeson00}, who used the Keck Interferometer to detect variations of phase vs.\\ wavelength in data from the binary star $\\iota$~Peg.  An ideal class of objects to which the differential phase technique can be applied is Be stars, which are composed of a B star usually unresolved by interferometric observations ($<1$~mas), surrounded by a disk of gas with a diameter of a few mas, which is detected in H$\\alpha$ line emission (\\citealt{Tycner05}).  The eclipsing binary star $\\beta$ Lyrae (HD~174638, HR~7106, FK5~705) has been extensively studied since it was discovered to be an eclipsing system over 220 years ago \\citep{Goodricke1785}. Nevertheless, it remains one of the most baffling binary systems known (see \\citealt{Harmanec02} for a comprehensive review). The current model consists of a $\\sim$3 M$_{\\odot}$ B6-8II star (star 2 in Harmanec's notation) which has filled its Roche lobe and is losing mass at a rate of 2$\\times10^{-5}$ M$_{\\odot}$ year$^{-1}$ to a $\\sim$13 M$_{\\odot}$ early B star (star 1), which is completely embedded in and hidden by a thick accretion disk. This accretion disk simulates a pseudophotosphere, which appears to an observer to have a spectrum of an A5III star. The orbit is very nearly circular, the orbital period is a little over 12.9 days, and the period is increasing by 19 seconds per year. The system lies at a distance of $270 \\pm 39$~pc, as measured by Hipparcos \\citep{Perryman97}. The spectrum of $\\beta$ Lyrae contains at least six spectroscopic line components \\citep{Harmanec02}. The three major components are: (a) strong absorption lines from star 2 used to derive the radial velocity curve having an amplitude of 370 km~s$^{-1}$, (b) ``shell'' absorption components seen in the hydrogen, helium, and many metallic lines, believed to arise from accretion disk, and (c) strong lines of hydrogen and helium seen in emission which vary with orbital phase and with time. One of the features that makes $\\beta$ Lyrae unique among contact binaries is the likelihood that the strong emission lines come from a bipolar jet that has formed near the point where material streaming from star 2 interacts with the accretion disk around star 1 \\citep{Harmanec96}.   This paper is organized in the following way. In Section 2 we present the observations and reductions. In Section 3 we describe the corrections applied to the data in order to obtain differential phases and describe the results obtained for $\\beta$ Lyrae. In Section~4 we describe the imaging of the H$\\alpha$ emission of this source. In Section~5 we obtain information about the orbit of this system from the displacement of the H$\\alpha$ emission relative to the continuum photocenter. In Section~6 we compare the results obtained from the analysis of the H$\\alpha$ results with those obtained using a traditional V$^2$ analysis, and in Section~7 we give a summary of our results. ",
        "conclusions": "In this paper we presented the development of a differential phase technique using NPOI data. We described the methods used to correct the instrumental and atmospheric effects on the phases of $\\beta$ Lyrae, which allowed us to recover the complex visibility of the H$\\alpha$ channel. We also described the procedure used to correct the continuum contribution to the H$\\alpha$ channel and successfully imaged the line emission of this system using standard interferometric techniques. We found that the major axis is oriented along p.a. 248.8$^{\\circ}$, consistent with previous spectropolarimetric measurements, radio and optical interferometry results. From the comparison between the observed position of H$\\alpha$ relative to the continuum photocenter and values obtained from models, we find that our measurements indicate a semi major axis smaller than that predicted by the models. We suggest that the most likely causes for this discrepancy are a wrong distance, brightness ratio between the 2 stars and the uncertainty in the correction of the continuum contribution to the H$\\alpha$ channel. We also found that the results obtained using the technique developed here are consistent with those obtained from a more traditional analysis that uses only V$^2$ measurements.  The technique presented here represents a major step in the process of obtaining images from optical interferometric observations. Due to several limitation, the most important being the rapid time scales in which the atmosphere changes, the phases of a source wrap around too fast. Consequently, this information is lost and one cannot apply the usual image reconstruction techniques employed in radio interferometry. By making some assumptions about the structure of the source, our technique allows us to recover this information and obtain images for the emission line channels. In the case of $\\beta$ Lyrae this allowed us to measure the orientation and size of the orbit using the relative location of the H$\\alpha$ and continuum photocenters. We also found that the HeI emission has a different structure from that of H$\\alpha$.  Besides $\\beta$ Lyrae we are also applying this technique to other Be systems. This will allow us to image their circumstellar disk, map structures and study in better detail the accretion flow and effects such as the ionization of the disk by a binary companion."
    },
    "0801/0801.4258_arXiv.txt": {
        "abstract": " ",
        "introduction": "The binary stars are crucial for our knowledge about the universe. Especially eclipsing binaries provide us an unique insight to the basic physical parameters of the stars, stellar clusters, interstellar medium and galaxies. They are excellent distance indicators. We are able to learn more about the matter composition of the stars, about their evolution status, or the presence of planets or other components in these systems.  The very first task is the data acquisition, because only with precise input data is one able to get precise results. In last few decades mainly due to excellent satellite observatories (and not only in the visible part of the spectrum) our knowledge of them has rapidly grown.  Regarding the astrometry, there is still decreasing the number of observations of the wide pairs. On the other hand, due to the new interferometers, which could resolve the milli- and micro- arcsecond angular distances, the observable semimajor axes of the astrometric binaries are still decreasing. Unfortunately, most of the systems analyzed below have the angular size of the astrometric orbit from 1 arcsec down to 100 mas, which is beyond the limits for the modern multi-aperture interferometers. And the lack of recent observations lead to the low accuracy of the results.  Another approach is photometry and the classical observation of minimum light. Due to the large \"baseline\" of observers in our country and the interest of amateur astronomers, the number of these observations is growing very rapidly and the cooperation between professional and amateur astronomers is very intensive. Many of the observations of minimum timings used in this study came from amateur astronomers and these measurements are as accurate as from the professional observatories. Thanks to the large minimum times data set we are able to analyze many of the eclipsing binary systems for their long-term period variations.  The whole thesis is divided into several parts. In the first one is presented the theory needed for the analysis of multiple stellar systems by photometric and astrometric techniques, description of such systems and some limitations which have to be considered. In the second part are introduced several systems which show apparent period changes in their $O-C$ diagrams. And in the third one are the systems analyzed by photometry and astrometry simultaneously. Also the catalogue of other suggested systems for simultaneous analysis is included in this chapter. This is the crucial part of this thesis. The method itself is introduced in the chapter \\ref{Ch1} and the results are in the chapter \\ref{Ch3}.   \\chapter{Theory} \\label{Ch1} ",
        "conclusions": ""
    },
    "0801/0801.1227_arXiv.txt": {
        "abstract": "Recently the Milagro experiment observed diffuse multi-TeV gamma ray emission in the Cygnus region, which is significantly stronger than what predicted by the Galactic cosmic ray model. However, the sub-GeV observation by EGRET shows no excess to the prediction based on the same model. This TeV excess implies possible high energy cosmic rays populated in the region with harder spectrum than that observed on the Earth. % In the work we studied this theoretical speculation in detail. We find that, a diffuse proton source with power index $\\alpha_p\\lesssim 2.3$, or a diffuse electron source with power index $\\alpha_e\\lesssim2.6$ can reproduce the Milagro's observation without conflicting with the EGRET data. Further detections on neutrinos, a diagnostic of  the hadronic model, and hard X-ray synchrontron radiation, a diagnostic of the lepton model, help to break this degeneracy. In combination with the gamma ray observations to several hundred GeV by Fermi, we will be able to understand the diffuse emission mechanisms in the Cygnus region better. ",
        "introduction": "The Galactic diffuse gamma-ray emission provides important information on the origin and propagation of the Galactic cosmic rays (GCRs) \\citep{2000ApJ...537..763S,2004ApJ...613..962S}. In general, there exist three possible mechanisms for the diffuse gamma-ray emission: decay of neutral pion produced by hadronic interaction between cosmic ray (CR) nuclei and interstellar gas; inverse Compton (IC) scattering between CR electron and interstellar radiation field (ISRF); and the bremsstrahlung radiation by interaction between CR electron and interstellar gas. In addition, dark matter annihilation is emerging as an alternative possibility to the diffuse $\\gamma$-rays \\citep{2005NewAR..49..213D,2005A&A...444...51D,2007arXiv0711.1912D, 2008PhRvD..78d3001B,2008PhLB..668...87B}.  The Cygnus region, defined as $65^{\\circ}<l<85^{\\circ}$ and $-3^{\\circ}<b<3^{\\circ}$ following the Milagro measurements, is rich in molecular clouds and is one of the richest star formation regions in the Galaxy \\citep{1996ApJ...466..282D,2006A&A...458..855S}. Observations in radio \\citep{1952AuSRA...5...17P}, infrared \\citep{2000A&A...360..539K,2002A&A...389..874C,2003ApJ...597..957H}, optical \\citep{1969A&A.....1..270D}, X-ray \\citep{1967ApJ...148L.119G} and $\\gamma$-ray \\citep{1996A&AS..120C.423C} bands found many interesting sources in this region. These sources have provided us valuable information about the astrophysical processes in this region. Ground-based very high energy (VHE) $\\gamma$-ray observatories also detected TeV $\\gamma$-ray emissions in the Cygnus region \\citep{2002A&A...393L..37A, 2007ApJ...658L..33A}. Such detection is of great importance for CR physics, since it indicates the existence of high energy CR accelerators. Especially the Milagro experiment\\footnote{Milagro homepage, http://umdgrb.umd.edu/cosmic/milagro.html} has made remarkable progress recently.  The Milagro experiment discovered two sources together with a diffuse TeV $\\gamma$-ray emission in the Cygnus region \\citep{2007ApJ...658L..33A,2008arXiv0805.0417A,2007ApJ...664L..91A}. For the two sources, MGRO J2031+41 is possibly the counterpart of the unidentified TeV source TeV J2032+4130, first discovered by HEGRA \\citep{2002A&A...393L..37A}, and MGRO J2019+37 is observed for the first time without an obvious counterpart. There are extensive discussions about the nature of the two TeV sources \\citep{2002A&A...393L..37A,2003ApJ...597..494B,2003ApJ...589..487M, 2007PhRvD..75h3001B}. However, no consensus has been reached yet. As for the diffuse $\\gamma$-ray emission, the Milagro measurement shows an evident excess (hereafter ``TeV excess'') compared with the diffuse emission predicted by the Galactic cosmic ray model GALPROP \\citep[$\\sim7$ times of the conventional model prediction and $\\sim2$ times of the optimized model prediction,][]{2008arXiv0805.0417A}. On the other hand, the sub-GeV band measurements by EGRET are well consistent with the GALPROP prediction \\citep{1997ApJ...481..205H,2007ApJ...658L..33A} \\footnote{It should be pointed out that, there are ``GeV excesses'' of EGRET observations compared with the conventional model of diffuse $\\gamma$-ray emission in $\\gtrsim$GeV energy band \\citep{1997ApJ...481..205H}. However, since it is not the main purpose of the present work, we just briefly discuss the ``GeV excess'' problem for the sake of completeness and do not intend to go deep in this topic.}.  The ``TeV excess'' may be due to some unidentified TeV sources or diffuse high energy proton or electron population  in the Cygnus region\\citep{2007ApJ...658L..33A}. If the latter is true we can draw two general conclusions according to the Milagro and EGRET data: 1), the spectrum of the proton or electron should be harder than the one observed at the Earth; 2), the CRs of the source population need to reside in the Cygnus region in order not to exceed the locally measured CR fluxes. In this work we investigate this possible explanation in detail by building realistic models and deriving their implications for future experiments.  This paper is organized as follows: in Sec.2 we give an introduction to the diffuse $\\gamma$-ray emission predicted by conventional GCR model. In Sec.3 we present a brief introduction to the ``GeV excess'' problem. Our models to solve the ``TeV excess'' problem and possible observational effects for future experiments are discussed in Sec.4. Finally the conclusion is drawn in Sec.5. ",
        "conclusions": "The TeV $\\gamma$-ray emission of the Cygnus region observed by Milagro shows significant excess compared with the conventional CR model. In this work we introduce a high energy proton or electron source to explain this high energy $\\gamma$-ray emission. It is shown that both the hadronic and leptonic models can give good explanation to the current data. The total energy needed for the source population is consistent with the typical energy output of a supernova.  The associated neutrino emission for the hadronic model and synchrotron radiation for the leptonic model are discussed as possible probes to discriminate these two scenarios. For the hadronic model, we estimate a $3\\sigma$ significance for 15-year exposure of IceCube in case the protons cutoff energy reaches $\\sim10$ PeV. The non-null result of the neutrino detection will support the hadronic model. For the leptonic scenario, the synchrotron radiation by the electrons is more luminous at the X-ray wavelength. We show that if the energy cut-off of the electron source is as high as $\\sim1$ PeV, the synchrotron radiation will peak in the hard X-ray bands and it might be observable on the satellite HXMT.  In the calculation, the ISM density, the ISRF intensity and the magnetic field are adopted as typical Galactic values. Since there are a large number of massive stars and molecular clouds in the Cygnus region, the ISRF and ISM density may be higher than the Galactic average values. We should point out that this does not change the basic conclusion in this work. This is because a higher ISM density or ISRF intensity can be effectively compensated by a lower CR flux. For the hadronic model, the neutrino flux will keep the same since neutrino flux is proportional to the $\\gamma$-ray flux. For the leptonic model a higher ISRF intensity will correspond to a lower electron luminosity. However, considering the magnetic field in the Cygnus region may be also stronger than the Galactic value we may expect similar synchrotron radiation as given in the work.  Finally, since the present data are lack in the energy interval from $\\sim10$GeV to $\\sim1$TeV due to the transition of detection technology from space to ground there are large uncertainties for model construction. The space telescope Fermi will extend the detection energy up to $\\sim300$GeV \\citep{2004HEAD....8.2101R}. Together with some ground-based experiments, such as ARGO, which may lower the energy threshold down to $\\sim100$GeV \\citep{1999NuPhS..78...38B}, we will have better understanding in the emission mechanism of this region in the future."
    },
    "0801/0801.1011_arXiv.txt": {
        "abstract": "We discuss the different physical processes that are important to understand the thermal X-ray emission and absorption spectra of the diffuse gas in clusters of galaxies and the warm-hot intergalactic medium. The ionisation balance, line and continuum emission and absorption properties are reviewed and several practical examples are given that illustrate the most important diagnostic features in the X-ray spectra. ",
        "introduction": "\\label{Introduction}  Thermal X-ray radiation is an important diagnostic tool for studying cosmic sources where high-energy processes are important. Examples are the hot corona of the Sun and of stars, Solar and stellar flares, supernova remnants, cataclysmic variables, accretion disks in binary stars and around black holes (Galactic and extragalactic), the diffuse interstellar medium of our Galaxy or external galaxies, the outer parts of Active Galactic Nuclei (AGN), the hot intracluster medium, the diffuse intercluster medium. In all these cases there is thermal X-ray emission or absorption.  In the present paper we focus upon the properties of the X-ray spectrum in the diffuse gas within and between galaxies, clusters and the large scale structures of the Universe. As this gas has very low densities, the level populations in the atoms are not governed by Saha-like equations, but instead most of the atoms will be in or near the ground state. This simplifies the problem considerably. Furthermore, in these tenuous media we need to consider only gas with small or moderate optical depth; the radiative transport is therefore simple. Stellar coronae have similar physical conditions, but there the densities are higher so that in several cases the emergent spectrum has density-dependent features. Due to the low densities in our sources, these density effects can be ignored in most cases; however, for the lowest density gas photoionisation effects must be taken into account.  We show in this paper that it is possible to derive many different physical parameters from an X-ray spectrum: temperature, density, chemical abundances, plasma age, degree of ionisation, irradiating continuum, geometry etc.  The outline of this paper is as follows. First, we give a brief overview of atomic structure (Sect.~\\ref{sect:atomic}). We then discuss a few basic processes that play an important role for the thermal plasmas considered here (Sect.~\\ref{sect:basic}). For the proper calculation of an X-ray spectrum, three different steps should be considered: \\begin{enumerate} \\item the determination of the ionisation balance (Sect.~\\ref{sect:balance}), \\item the determination of the emitted spectrum (Sect.~\\ref{sect:emission}), \\item possible absorption of the photons on their way towards Earth (Sect.~\\ref{sect:absorption}). \\end{enumerate} We then briefly discuss issues like the Galactic foreground emission (Sect.~\\ref{sect:galem}), plasma cooling (Sect.~\\ref{sect:cooling}), the role of non-thermal electrons (Sect.~\\ref{sect:nonthermal}), and conclude with a section on plasma modelling (Sect.~\\ref{sect:plasmamodelling}). ",
        "conclusions": ""
    },
    "0801/0801.3825_arXiv.txt": {
        "abstract": "This paper addresses the nature of the near infrared background.  We investigate whether there is an excess background at 1.4 microns, what is the source of the near infrared background and whether that background after the subtraction of all known sources contains the signature of high redshift objects ($Z > 10$).  Based on NICMOS observations in the Hubble Ultra Deep Field and the Northern Hubble Deep Field we find that there is no excess in the background at 1.4 microns and that the claimed excess is due to inaccurate models of the zodiacal background. We find that the near infrared background is now spatially resolved and is dominated by galaxies in the redshift range between 0.5 and 1.5. We find no signature than can be attributed to high redshift sources after subtraction of all known sources either in the residual background or in the fluctuations of the residual background.  We show that the color of the fluctuations from both NICMOS and \\emph{Spitzer} observations are consistent with low redshift objects and inconsistent with objects at redshifts greater than 10. It is most likely that the residual fluctuation power after source subtraction is due to the outer regions of low redshift galaxies that are below the source detection limit and therefore not removed during the source subtraction. ",
        "introduction": "The nature of the near infrared background is the subject of intense current investigation.  Much of this interest centers on whether the near infrared background contains the signature of very high redshift ($Z > 10$) sources. Claims for such a signature have been spurred by the claim of an excess in the background with a sharp cutoff to the blue of 1.4$\\mu$ \\cite{ref:mat05} and from fluctuation analyses of deep \\emph{Spitzer} data \\cite{ref:kas05,ref:kas07b}. These findings are in contrast to earlier analyses of NICMOS images from the Northern Hubble Deep Field (NHDF) \\cite{ref:thm03} and the Hubble Ultra Deep Field (HUDF) \\cite{ref:thm06} where no excess was found. ",
        "conclusions": "We conclude that the Near Infrared Extragalactic Background has been resolved and is due primarily to normal galaxies at redshifts near 1.  The claimed excess at 1.4 microns is false and arose from the inadequacy of zodiacal models to predict the background level to the accuracy need to determine the true source flux.  We further conclude that none of the properties of the fluctuation spectrum after source subtraction in any of the 2MASS, NICMOS or \\emph{Spitzer} images require very high redshift ($z>10$) objects to account for them."
    },
    "0801/0801.2318_arXiv.txt": {
        "abstract": "Globular Cluster Systems (GCSs) of most early-type galaxies feature two peaks in their optical colour distributions. Blue-peak GCs are believed to be old and metal-poor, whereas the ages, metallicities, and the origin of the red-peak GCs are still being debated. We obtained deep K-band photometry and combined it with HST observations in g and z to yield a full SED from optical to near-infrared. This now allows us to break the age-metallicity degeneracy.  We used our evolutionary synthesis models GALEV for star clusters to compute a large grid of models with different metallicities and a wide range of ages. Comparing these models to our observations revealed a large population of intermediate-age (1--3 Gyr) and metal-rich ($\\approx$ solar metallicity) globular clusters, that will give us further insights into the formation history of this galaxy. ",
        "introduction": "\\label{s:intro} Globular Cluster Systems (GCSs) are now recognized as powerful tracers of their parent galaxy's formation history \\citep{west04,fritze04,brodie06}. From their age and metallicity distributions one can reconstruct the parent galaxy's (violent) star formation and chemical enrichment history all the way from the very onset of star formation in the early universe to the present.  Most early-type galaxies show a bimodal color distribution for their GCSs \\citep[e.g.][]{gebhardt99,kundu01a,kundu01b,peng06}: A universal blue peak and a red peak with colors and height relative to the blue peak varying from galaxy to galaxy. The blue peak GCs are believed to be old and metal-poor, the origin of the red peak is still unclear.  \\subsection{Our approach to lift the age-metallicity degeneracy}  Optical data alone do \\textbf{not} allow to disentangle ages and metallicities: Colour-to-metallicity transformations have to assume an age, while colour-to-age transformations are only valid for one metallicity. This degeneracy, however, can be broken by including near-infrared data that are more sensitive to changes of metallicity rather than age.  \\cite{anders04b} used extensive artificial star cluster tests and showed that observations in \\begin{itemize} \\item 3 passbands for GCSs in dust-free E/S0s or \\item 4 passbands for (young) star clusters in dusty environments, \\item spanning as wide as possible a wavelength-basis (U through K) and \\item including at least one NIR-band (e.g. K) with \\item accuracies $\\leq 0.05$ mag in the optical and $\\leq 0.1$ mag in the NIR \\end{itemize} allow to disentangle ages and metallicities and determine \\textbf{individual GC metallicities to $< \\pm 0.2$ dex, and ages to $< \\pm 0.3$ dex}, i.e. they allow to distinguish $\\leq 7\\,\\rm Gyr$ old GCs from those $\\geq 13\\,\\rm Gyr$ old.  Similar studies also using NIR-data to determine ages and metallicities of globular clusters have only been done for a few galaxies until now \\citep{puzia02,kisslerpatig02,hempel03,larsen05,hempel07}. More than half of these are found to host a population of GCs that is younger and/or more metal-rich than the old and metal-poor GC population in the Milky Way.  We note that many of these previous studies relied on cumulative distribution. The most important change compared to our analysis is that we derive the physical parameters for each individual cluster. That also enables us to study correlations of these parameters, e.g. with their spatial distributions. ",
        "conclusions": "We obtained deep K-band photometry of the Virgo lenticular NGC4570 and combined it with archival optical data from HST. We selected globular cluster candidates based on the HST data and their intrinsic optical sizes.  GALEV evolutionary synthesis models were used in combination with AnalySED to automatically derive physical parameters age, metallicity and mass for each individual cluster.  We detect a significant population of intermediate-age (1--3 Gyr) and metal-rich ($\\rm [Fe/H] > -0.4$) cluster population that has not been reported for this galaxy before."
    },
    "0801/0801.0547_arXiv.txt": {
        "abstract": "{ Galactic foreground emission fluctuations are a limiting factor for precise cosmic microwave background (CMB) anisotropy measurements.} { We intend to improve current estimates of the influence of Galactic synchrotron emission on the analysis of CMB anisotropies.  } { We perform an angular power spectrum analysis (APS) of all-sky total intensity maps at 408~MHz and 1420~MHz, which are dominated by synchrotron emission out of the Galactic plane. We subtract the brighter sources from the maps, which turns out to be essential for the results obtained. We study the APS as a function of Galactic latitude by considering various cuts and as a function of sky position by dividing the sky into patches of $\\sim 15^{\\circ}\\times 15^{\\circ}$ in size. } { The APS of the Galactic radio diffuse synchrotron emission is best fitted by a power law, $C_{\\ell} \\sim k \\ell^{\\alpha}$, with $\\alpha \\in [-3.0,-2.6]$, where the lower values of $\\alpha$ typically correspond to the higher latitudes. Nevertheless, the analysis of the patches reveals that strong local variations exist. A good correlation is found between the APS normalized amplitude, $k_{100}=k \\times 100^\\alpha$, at 408~MHz and 1420~MHz. The mean APS for $\\ell \\in [20,40]$ is used to determine the mean spectral index between 408 MHz and 1420 MHz, $\\beta_{(0.408-1.4){\\rm GHz}} \\in [-3.2,-2.9]$ ($C_{\\ell}(\\nu) \\propto \\nu^{-2\\beta}$), which is then adopted to extrapolate the synchrotron APS results to the microwave range. } { We use the 408 MHz and 1420 MHz APS results to predict the Galactic synchrotron emission fluctuations at frequencies above 20~GHz. A simple extrapolation to 23~GHz of the synchrotron emission APS found at these radio frequencies does not explain all the power in the WMAP synchrotron component even at middle/high Galactic latitudes. This suggests a significant microwave contribution (of about $50\\%$ of the signal) by other components such as free-free or spinning dust emission. The comparison between the extrapolated synchrotron APS and the CMB APS shows that a mask excluding the regions with $|b_{gal}| \\lesssim 5^{\\circ}$ would reduce the foreground fluctuations to about half of the cosmological ones at 70~GHz even at the lowest multipoles. The main implications of our analysis for the cosmological exploitation of microwave temperature anisotropy maps are discussed. } ",
        "introduction": "The possibility to study the primordial phases of our Universe and its properties and evolution through CMB anisotropies relies on our capability to precisely extract the cosmological signal from observations \\citep{teg00_fore}. Maps of the microwave sky include many Galactic and extragalactic astrophysical contributions. A correct recovery of the CMB anisotropy field requires an accurate removal of those foreground signals from the observed maps. Current knowledge of the foreground components permits one to retrieve the bulk of the cosmological information encoded in the CMB anisotropy, and, in particular, its angular power spectrum (APS). Nevertheless, a deeper understanding of the foregrounds is crucial to settle important cosmological issues, which arose from the {\\sc WMAP}~\\footnote{http://lambda.gsfc.nasa.gov/} results \\citep{naselski06_wmap1_et_fore,naselski06_wmap_ell4,chiang07_lowLanom} and could potentially be addressed by the forthcoming {\\sc Planck}~\\footnote{http://www.rssd.esa.int/Planck} mission \\citep{tauber04_cospar}. Moreover, it would enable a precise reconstruction of the individual foreground components and therefore a complete astrophysical exploitation of the satellite data.   Galactic synchrotron emission is the major source of contamination at frequencies below 50~GHz for intermediate and large angular scales, as recently confirmed by the impressive {\\sc WMAP} results \\citep{ben03_wmap_1yr_fore,hins06_wmap_3yr_temp}. Synchrotron radiation \\citep{rybicki79} arises from cosmic ray electrons gyrating in the magnetic field of our Galaxy. The energy spectrum and the density of the cosmic ray electrons as well as the magnetic field strength vary across the Galaxy, therefore the observed synchrotron emission will depend on the frequency and on the region of the sky. Radio observations at $\\nu \\lesssim 5$~GHz provide the clearest picture of the Galactic synchrotron morphology, since at these frequencies the diffuse non-thermal radiation clearly dominates over all other emission components outside the Galactic plane. In the past, the 408~MHz all-sky map \\citep{haslam82_408mhz} has extensively been used as a template for the Galactic synchrotron emission in foreground separation attempts (e.g. Bouchet \\& Gispert~1999; Bennett et al.~2003). In addition, that map as well as other less suited surveys have been exploited to find an appropriate parametrization of the synchrotron emission APS to be $i)$ extrapolated to the microwave range for estimating the contamination of CMB anisotropies at different angular scales \\citep{giard01_hardtrao,bacci01_synch} and $ii)$ used to set priors in foreground separation applications \\citep{teg96_lf,bouchet99_wf,bouchet99_foresim}. The outcome of these analyses is that the APS of the synchrotron emission computed over large portions of the sky can be modelled by a power law, i.e. $C_{\\ell} \\sim k\\; \\ell^{\\alpha}$ ($\\ell \\sim 180^{\\circ}/{\\theta}$), with a spectral index $\\alpha \\sim [-3,-2.5]$ for $\\ell \\lesssim 200$, corresponding to angular scales $\\theta \\gtrsim 1^{\\circ}$. Clearly, such a general result is just a first step as it does not describe the complexity of the synchrotron emission APS, whose parameters are expected to change with frequency and with sky direction.   We have carried out a detailed analysis of all-sky radio maps and improved on previous attempts by providing an accurate characterization of the synchrotron emission APS. The results obtained for the new 1.4~GHz polarization all-sky survey (Reich et al., in prep.) will be reported in a forthcoming companion paper. The analysis presented in this paper focuses on the synchrotron emission APS in total intensity. A substantial improvement was possible by using a new all-sky map at 1.42~GHz (Reich et al., in prep.), which has a higher angular resolution and better sensitivity, in addition to the all-sky map at 408~MHz. These maps are currently the best suited data for studying the Galactic synchrotron emission at largest angular scales. A more detailed description of some technical aspects related to this work can be found in \\citet{phdthesis}. We extensively used the {\\tt HEALPix}\\footnote{http://healpix.jpl.nasa.gov/} software package \\citep{gorski05_healpix}. {\\tt HEALPix} (Hierarchical Equal Area isoLatitude Pixelization) is a curvilinear partition of the sphere optimized for fast spherical harmonics transforms and angular power spectrum estimation. The latter task is performed by the facility {\\tt Anafast}. We also made use of the data reduction package based on the NOD2-software \\citep{haslam74_nod}. \\\\   The paper is organized as follows. Sect.~2 describes the characteristics of the 408~MHz and 1.42~GHz total intensity surveys, their projection onto {\\tt HEALPix} maps and noise considerations. In Sect.~3 the Galactic radio emission APS over large areas is examined, which reveals the necessity of a discrete source subtraction for a correct evaluation of the diffuse synchrotron APS. The two all-sky maps are decomposed into a map of the diffuse component and a map of discrete sources. Their angular power spectra are derived and discussed. The results obtained by fitting the angular power spectra of the diffuse component maps are presented. In Sect. 4 the radio survey angular power spectra are extrapolated to the microwave range for a comparison with the WMAP 3-yr results. Sect. 5 is dedicated to a local analysis of the radio map APS. We summarize our results and conclusions in Sect.~6. \\\\ ",
        "conclusions": "The aim of our analysis is to improve our understanding of the Galactic synchrotron emission as a foreground for CMB dedicated experiments. For this purpose, we carried out an unprecedented detailed study of the Galactic radio emission, in terms of its angular power spectrum (APS), using total intensity all-sky maps at 408 MHz and 1420 MHz. \\\\ An accurate modeling of the synchrotron APS is missing in the literature so far, but is urgently required for a more precise and complete exploitation of the information awaited from the {\\sc Planck} satellite. It constitutes a precious input for component separation activities, both for the realization of spatial templates of the foreground and for the definition of priors on its spatial and frequency dependence. \\begin{itemize} \\item[{\\bf 1.}] Being interested in the diffuse component of the synchrotron emission, the brighter discrete sources (DS) have been eliminated from the radio maps by 2-dimensional Gaussian fitting. This approach is very flexible and also permits the removal of extended structures. \\item[{\\bf 2.}] The APS was computed for both large areas and small patches and several consistency tests were used to check the reliability of the recovered APS in the case of limited sky coverage (see Appendix~B of La Porta 2007). \\\\ The study of the APS for various cuts, i.e. of regions with Galactic latitudes above or below a certain value $b_{cut}$ ($|b_{cut}| \\in [5^{\\circ},60^{\\circ}]$), allowed us to explore a possible dependence of the mean properties of the Galactic synchrotron emission on latitude, preserving at the same time the largest possible coverage, which is important when estimating the CMB APS because of the sampling variance. Such cuts provided information for $\\ell \\in [\\ell_{min},\\ell_{max}]$, where $10 \\lesssim \\ell_{min} \\lesssim 30$ for increasing $b_{cut}$ and $\\ell_{max} \\sim 200 -300$ at 408 MHz and 1420 MHz, respectively. \\\\ The patches correspond to the pixels of an {\\tt HEALPix} map at $n_{side}=4$, which have an angular dimension of $\\sim 15^{\\circ}$ and permit us to investigate the local variations of the synchrotron APS for multipoles larger than $\\sim 60$. \\\\ The derived angular power spectra were modelled in both cases according to Eq.~\\ref{aps_best_model} and a specific method was set up to find the best least square fit on adaptive grids of the parameter space and to evaluate the uncertanties on the retrieved parameters (see Appendix~C of La Porta~2007 for details). \\item[{\\bf 3.}] An indirect cross-check of the fit result reliability was provided in the case of the Galactic cuts by the estimated source terms, which are consistent with the expectations from extragalactic source counts at both frequencies. Nowadays, source counts at 1.4 GHz are well established down to very low flux limits. It is remarkable that the source angular power spectra obtained independently from source counts and from the fit of the survey APS are in good agreement. \\item[{\\bf 4.}] The slope of the synchrotron APS, $C_{\\ell} \\sim k \\ell^{\\alpha}$, changes with $b_{cut}$ without showing a well-defined regular trend, although it is found to be typically steeper for $b_{gal} \\gtrsim 20^{\\circ}$. For the cuts, $\\alpha$ varies in the range $\\sim [-3.0,-2.6]$ at both frequencies. However, the analysis of the small patches gives evidence that locally the synchrotron APS can be much flatter for $\\ell \\gtrsim 60$, reaching in some cases values of $\\alpha \\sim -0.8$. The normalized amplitude, $k_{100}=k \\times 100^{\\alpha}$, gradually increases toward the Galactic plane, following, as expected, the background radio emission gradient. A good correlation exists between the results obtained for $k_{100}$ at 408 MHz and 1420 MHz, for both cuts and patches. This is expected, given that the spectral properties of the electron density distribution responsible for the Galactic diffuse non-thermal emission should be the same in that frequency range and further supports the reliability of the obtained estimates of the synchrotron APS. \\item[{\\bf 5.}] The maps of $k_{100}$ and $\\alpha$ resulting from the local analysis represent the starting point for the simulation of small-scale fluctuation fields to be added to the DS-subtracted maps to build phenomenological templates of the Galactic synchrotron emission. At present, an empirical approach is the most reliable way to proceed in the realization of realistic templates of the foreground, given the poor knowledge of the Galactic magnetic field and of the cosmic ray electron density distribution needed for a 3-dimensional physical modelling. \\end{itemize} The issues (discussed in Sect.~\\ref{extrap_kmem_comp}) related to the K-band synchrotron component retrieved by \\citet{hins06_wmap_3yr_temp} are a clear example of the difficulties encountered in the foreground separation due to the lack of adequate priors and guess. We performed a source subtraction on that map and produced a map of the Galactic diffuse non-thermal emission at 23 GHz to be compared with those at lower frequencies. \\begin{itemize} \\item[{\\bf 6.}] The extrapolation to 23 GHz of the APS obtained at 408 MHz and 1420 MHz for higher latitude regions ($|b_{gal}|\\gtrsim 40^{\\circ}$) reveals that the mean spectral index ($C_{\\ell}(\\nu) \\propto \\nu^{2\\beta}$) $\\beta_{(0.408-23){\\rm GHz}} \\lesssim \\beta_{(1.4-23){\\rm GHz}}$, which is the opposite of what is expected for synchrotron emission. We estimate the excess of the signal in the 23 GHz map to be $\\gtrsim 50\\%$ by using the mean value of the APS at lower multipoles. This result can be interpreted in terms of additional contributions to the 23 GHz Galactic non-thermal emission, which could be mainly due to anomalous dust. \\end{itemize}  A direct application of the presented analysis is to determine the level of contamination of the CMB anisotropies due to the Galactic diffuse synchrotron emission at different angular scales. The conclusions reported below refer to the cuts, which are more relevant for CMB measurements because of the large coverage. However, similar considerations could be repeated using the results of the local analysis, thus allowing the identification of the clearest sky areas, which is essential for the success of ground-based CMB experiments. \\begin{itemize} \\item[{\\bf 7.}] An important - although not unexpected - outcome of the cut analysis is that the amplitude of the APS for the northern hemisphere cuts above $\\sim 20^{\\circ}$ is raised by the presence of the NPS. Consequently, the results obtained separately for the two Galactic hemispheres were used to bracket the APS of the synchrotron emission. \\item[{\\bf 8.}] We used the APS results at 408 MHz and 1420 MHz to determine the frequency spectral index to be adopted in the extrapolation to the microwave range for each coverage case. In particular, we found that $\\beta_{(0.408-1.4){\\rm GHz}} \\in [-3.2,-2.9]$ with an uncertainty of a few percent. \\item[{\\bf 9.}] The extrapolation to 30~GHz of the synchrotron APS obtained at 1420 MHz led to a signal in good agreement with the upper limit fixed by COBE-DMR, thus supporting the reliability of our results. At this frequency the synchrotron emission constitutes a severe contamination of CMB anisotropies at the largest angular scales ($\\ell \\lesssim 40$). Nevertheless, a cut at $\\sim 20^{\\circ}$ would reduce the synchrotron fluctuations to about half the cosmological ones. The same holds at $\\nu \\sim 70$~GHz for a cut at $\\sim 5^{\\circ}$ and the situation further improves when excluding a larger portion of the sky around the Galactic plane. This implies that even though the current treatment of the foregrounds does not permit an accurate removal of the synchrotron emission, the latter does not prevent the recovery of the bulk of the cosmological information encoded in the CMB temperature APS.\\\\ \\end{itemize} A deeper understanding of the foreground is indispensable to settle other important issues in the perspective of the forthcoming {\\sc Planck} mission. {\\sc Planck} should achieve a sensitivity comparable to the cosmic variance and its performance should be limited mainly by foregrounds, thanks to the extremely accurate control of all instrumental systematic effects \\citep{buri04_PafterW,menn04_lfi}. \\\\ A first issue is the estimate of the CMB temperature-polarization APS, which carries information about the reionization history and the tensor-to-scalar ratio (see, e.g. Kogut~2003, Kogut et al.~2003, and references therein). Possible foreground residual contamination in the total intensity CMB anisotropy map would affect fine analysis based on the estimate of the cross-correlation APS, also because the polarized component of the cosmological signal is orders of magnitude lower. \\\\ Another issue is the evaluation of the Gaussianity of the primordial fluctuations, which in the standard inflationary paradigm\\footnote{ Simple (standard) inflationary scenarios (see the reviews by Lyth \\& Riotto~1999 and Linde~2005) predict the existence of Gaussian density fluctuations and of gravitational waves with a nearly scale-invariant spectrum. } generate the structures observed in the Universe today. Gaussianity tests are a powerful tool, complementary to the tests exploiting the APS, which allow us to probe the ``concordance'' model \\citep{sperg06_wmap_3yr_param} and also to distinguish among inflationary models \\citep{bart04_nonG}. The level of non-Gaussianity predicted by the ``concordance'' model cannot be detected by {\\sc WMAP}, but in principle it should be observable with {\\sc Planck}. Galactic foregrounds are non-Gaussian and anisotropic and even low-level contamination in the maps can produce detectable non-Gaussianities (see, e.g. Naselsky et al.~2005), although they have minimal effects on the estimated APS \\citep{hin03_wmap_1yr_aps}. Consequently, the foreground removal has to be extremely accurate, so as not to limit {\\sc Planck} in verifying this crucial\\footnote{Standard cosmologies predict a minimum level of non-Gaussianity for the the primordial perturbations. In alternative cosmologies \\citep{lyth03_curvaton,ark04_ghostinfl,alish04_altercosmol} such a lower limit is even  higher. Consequently, detection or non-detection of non-Gaussianities sheds light on the physics of the early Universe.} prediction."
    },
    "0801/0801.2774_arXiv.txt": {
        "abstract": "We have analysed the distribution of inclination-corrected galaxy concentrations in the Sloan Digital Sky Survey. We find that unlike most galaxy properties, which are distributed bimodally, the distribution of concentrations is trimodal: it exhibits three distinct peaks. The newly-discovered intermediate peak, which consists of early-type spirals and lenticulars, may contain $\\sim 60\\%$ of the number density and $\\sim 50\\%$ of the luminosity density of $\\absmag < -17$ galaxies in the local universe. These galaxies are generally red and quiescent, although the distribution contains a tail of blue star-forming galaxies and also shows evidence of dust. The intermediate-type galaxies have higher apparent ellipticities than either disc or elliptical galaxies, most likely because some of the face-on intermediate types are misidentified as ellipticals. Their physical half-light radii are smaller than the radii of either the disc or elliptical galaxies, which may be evidence that they form from disc fading. The existence of a distinct peak in parameter space associated with early-type spiral galaxies and lenticulars implies that they have a distinct formation mechanism and are not simply the smooth transition between disc-dominated and spheroid-dominated galaxies. ",
        "introduction": "Morphological classification of galaxies dates back to the \\citet{hubble26} ``Tuning Fork'' that classified galaxies into ellipticals and spirals (and further into barred and unbarred spirals). More recent and quantitative classifications of large numbers of galaxies from current galaxy redshift surveys have upheld this bimodality, finding that morphological properties such as the global concentration of the light (e.g. as measured by the Petrosian concentration parameter \\cpetro), the functional form of the surface brightness profile (e.g. exponential versus $R^{1/4}$ law, or the value of the \\citet{sersic-profile} $n$ index), and the asymmetry are best decomposed into two populations: concentrated, smooth, de~Vaucouleurs (or high $n$) ellipticals and diffuse, clumpy, exponential spirals. Properties of galaxies sensitive to their star formation histories, such as the global colour, the presence of nebular emission lines, the presence of spectral features symptomatic of young stars, and the detailed stellar populations when they can be resolved, are also decomposed into two populations: red, quiescent ellipticals full of old stars and blue, star-forming spirals with young stellar populations \\citep[see][and references therein]{shimasaku-etal01,strateva-etal01, blanton-etal03-properties,baldry-etal04,driver-etal06}.  At the joining point of the tuning fork lies the S0 or lenticular class. These are galaxies that are too flattened not to have a disc morphology, but that appear dominated by their spheroidal component and appear more similar to ellipticals than discs in their colour and environment. These intermediate galaxies may be a unique class of object with a particular formation scenario, such as disc fading, or they may simply represent the smooth transition from galaxies with low bulge-to-disc ratios, $B/D$, to those with high $B/D$.  If lenticulars have a unique formation scenario, they are likely to congregate in a particular region of galactic parameter space. In this paper, we test this possibility by studying the morphological distribution of galaxies in the Sloan Digital Sky Survey. Section~\\ref{sdss-section} details the galaxy sample. In Section~\\ref{trimodality}, we investigate the distribution of galaxy concentrations and find that they are trimodal. In Section~\\ref{intermediate properties}, we investigate the properties of the newly-discovered intermediate class of galaxies. In Section~\\ref{discussion} we discuss the implications of our results, and Section~\\ref{conclusions} contains our conclusions. ",
        "conclusions": "\\label{conclusions} We have analysed the locations of SDSS galaxies in the Petrosian concentration \\cpetro\\ vs. isophotal axis ratio $b/a$ plane. Disc galaxies occupy a slanted curve in this plane due to inclination effects: more edge-on galaxies have higher \\cpetro. We have corrected for this and defined a new inclination-independent concentration index \\cnorm. Unlike most galaxy properties (including \\cpetro), which are distributed bimodally, the distribution of \\cnorm\\ is \\textit{tri}modal, with a third peak intermediate between the low-concentration disc galaxies and the high-concentration elliptical galaxies.  We have studied the properties of this newly-identified intermediate class of galaxy, which represents $\\sim 60\\%$ of the total number density and $\\sim 50\\%$ of the luminosity density of $\\absmag < -17$ galaxies detectable in SDSS. Most intermediate-type galaxies have colours and spectra typical of early-type galaxies, although a non-negligible number are blue and star forming. The location of these galaxies on the \\cpetro-$b/a$ plane indicates that they have discs. The concentrations are indicative of typical bulge-to-disc ratios of $3$, although this may be an underestimate. The colours and spectra of intermediate-type galaxies vary with inclination (edge-on galaxies are redder and appear to be more passive), indicative of dust. The fraction of blue star-forming galaxies among the intermediate class may be underestimated due to dust obscuration of star-forming regions. The typical absolute magnitudes and half-light radii of the SDSS intermediate-type galaxies are $-21$ and $2.4~\\mathrm{kpc}$ respectively. The physical sizes of the intermediate-type galaxies are significantly smaller than those of both ellipticals and discs.  Intermediate-type galaxies have much higher apparent ellipticities than either disc or elliptical galaxies. Our preferred explanation is that these galaxies contain residual central star formation that is obscured by dust when viewed edge-on; therefore, the measured concentration is higher when viewed face-on and these galaxies merge into the region of parameter space occupied by high-concentration ellipticals.  Associating our intermediate class of galaxies with lenticulars and early-type spirals, we propose that the existence of a distinct peak in parameter space implies that there is a particular formation mechanism responsible for the creation of the earliest-type spiral galaxies."
    },
    "0801/0801.3258_arXiv.txt": {
        "abstract": "Boxy/peanut bulges in disc galaxies have been associated to stellar bars. We analyse their properties in a large sample of $N$-body simulations, using different methods to measure their strength, shape and possible asymmetry, and then inter-compare the results. Some of these methods can be applied to both simulations and observations. In particular, we seek correlations between bar and peanut properties, which, when applied to real galaxies, will give information on bars in edge-on galaxies, and on peanuts in face-on galaxies. ",
        "introduction": "Both simulations and theoretical studies have shown that bars are not vertically thin morphological features, but have a considerable vertical extent and a vertical structure, known as the Boxy/Peanut bulges (hereafter B/P; Combes \\& Sanders 1981, Combes et al. 1990). Comparisons between observations and $N$-body simulations have established this direct connection firmer (Athanassoula 2005 and references therein). Observations have shown that both bars and B/P bulges are quite predominant in disc galaxies and that the corresponding frequencies are in good agreement with the link between the two structures (L\\\"utticke, Dettmar \\& Pohlen 2000).  We measure the peanut properties in a large sample of several hundred $N$-body simulations ran by one of us (EA) for different purposes. More information on these simulations and on their properties can be found in Athanassoula \\& Misiriotis (2002) and Athanassoula (2003, 2007). In particular, we seek correlations between the properties of the bar and the properties of the B/P bulge. ",
        "conclusions": "We presented several methods to calculate the strength of the bar and the strength, shape and asymmetry of the B/P bulge and found strong correlations between their results. The most important correlation relates the strength of the bar with the strength of the B/P bulge, the strongest bars having the strongest peanuts. We also find that the strength of the peanut depends on the number of buckling episodes it underwent, the strongest bars having undergone more buckling episodes (Fig. 3). Finally, we find a very interesting result about $C_{2,z}(R)$, i.e. about the shape of the radial density profiles along cuts perpendicular to the equatorial plane. For strong bars, having a strong peanut or X-shaped bulge, this profile is more flat-topped, while for weaker bars, with more boxy-like bulges, it is more peaked. All the results summarized here are discussed in length by Athanassoula \\& Martinez-Valpuesta (2008, in preparation).  \\begin{figure} \\begin{center} \\includegraphics[scale=0.8]{Martinez-Valpuesta_fig4.eps} \\caption{Correlations between bar and peanut properties. Each symbol corresponds to one simulation. The type of symbol is related to the number of buckling events suffered by the bar during its evolution. {\\it Left panel}: Strength of the B/P bulge measured with our Fourier based method vs. the strength of the bar. {\\it Right panel}: Shape of the B/P bulge (i.e. shape of the radial density profiles along cuts perpendicular to the equatorial plane, measured by the minimum of the kurtosis) plotted as a function of bar strength. } \\end{center} \\end{figure}"
    },
    "0801/0801.3772_arXiv.txt": {
        "abstract": "When calculating stellar initial mass functions (IMFs) for young clusters, one has to take into account that most massive stars are born in multiple systems and that most IMFs are derived from data that cannot resolve such systems. It is also common to measure IMFs for clusters that are located at distances where multiple chance superpositions between members are expected to happen. In this article I model the consequences of both of those phenomena, real multiple systems and chance superpositions, on the observed color-magnitude diagrams and the IMFs derived from them. Using numerical experiments I quantify their influence on the IMF slope for massive stars and on the generation of systems with apparent masses above the stellar upper mass limit. The results in this paper can be used to correct for the biases induced by real and chance-alignment multiple systems when the effects are small and to identify when they are so large that most information about the IMF in the observed color-magnitude diagram is lost. Real multiple systems affect the observed or apparent massive-star MF slope little but can create a significant population of apparently ultramassive stars. Chance superpositions produce only small biases when the number of superimposed stars is low but, once a certain number threshold is reached, they can affect both the observed slope and the apparent stellar upper mass limit. In the second part of the paper, I apply the experiments to two well known massive young clusters in the Local Group, NGC 3603 and R136. In both cases I show that the observed population of stars with masses above 120 \\Ms/ can be explained by the effects of unresolved objects, mostly real multiple systems for NGC 3603 and a combination of real and chance-alignment multiple systems for R136. Therefore, the case for the reality of a stellar upper mass limit at solar or near-solar metallicities is strengthened, with a possible value even lower than 150 \\Ms/. An IMF slope somewhat flatter than Salpeter or Kroupa with $\\gamma$ between $-1.6$ and $-2.0$ is derived for the central region of NGC 3603, with a significant contribution to the uncertainty arising from the imprecise knowledge of the distance to the cluster. The IMF at the very center of R136 cannot be measured with the currently available data but the situation could change with new Hubble Space Telescope (HST) observations. ",
        "introduction": "This paper is the second one of a series where we explore the effects of different biases on the determination of the stellar initial mass function (IMF). In paper I \\citep{MaizUbed05} we analyzed the numerical biases induced by using bins of equal width when fitting power-laws to binned data (an effect that is more general than its application to the calculation of mass functions). Those biases can be eliminated in several ways, of which a simple one is by grouping the data in equal-number bins (as opposed to equal-width bins). In this second paper I explore the effect of unresolved multiple systems, either physical or chance alignments, especially for the high-mass end of the IMF. In a future paper we will analyze the effect of random uncertainties in the mass determinations on the computed initial mass function\\footnote{Most of the analysis in these papers is also applicable to the calculation of the present-day mass function (PDMF).}.  \\subsection{Binaries and the IMF}  Why worry about binaries and higher-order multiple systems when calculating the IMF of a young stellar cluster such as NGC 3603 or R136? The quick answer is because one expects most of the observed point sources in clusters beyond a few kpc to be multiple. Therefore, deriving their masses by assuming they are single stars should introduce biases in the determination of the stellar mass functions of the clusters.  \\citet{Masoetal98} studied nearby O stars from a sample that included field objects and stars in low- and intermediate-mass clusters and associations. They analyzed the observed multiple fractions as a function of environment, including both spectroscopic and visual/speckle binaries and discovered that at least 72\\% of the O stars in clusters and associations had detected companions. Furthermore, the period distribution had a clear bimodal structure, with one peak around 10 days and another one around $10^5$ years. They argued that such an effect was quite likely to be instrumental because binaries with periods $P$ near 1-1000 years are hard to detect and because the period distribution for B stars is quite flat\\footnote{Such a distribution, expressed in terms of separation $r$ instead of $P$, is close to $f(\\log r) \\propto (\\log r)^0$ and is known as \\\"Opik's law, after \\citet{Opik24}.}. Therefore, \\citet{Masoetal98} predicted that when better spatial resolution were achieved and more intensive radial-velocity variation surveys were undertaken in the future, the gap between the two peaks might be filled and the true multiple fraction for massive stars in clusters might be found to be close to unity. Another interesting result of that paper was that the mass-ratio distribution for O binaries was flat or nearly so, i.e. massive stars tend to have companions that are more massive on average than what would be expected from a random sampling of the IMF. The shape of the secondary mass distribution function is relevant because different slopes should affect the observed or apparent mass function derived from unresolved data in different ways. \\citet{SagaRich91} calculated that the apparent MF slope for intermediate-mass stars in the random pairing case becomes significantly flatter when adding unresolved binaries but the effect is likely to be different in the flat mass-ratio case.  In the ten years after \\citet{Masoetal98} other works have validated their conclusions. \\citet{GarcMerm01} measured that 11 out of 14 O-systems in the cluster NGC 6231 are spectroscopic binaries. This means that at least 79\\% of its O/WR systems are multiple and at least 88\\% of their O/WR (and some early B) stars are in multiple systems. Their numbers are also similar for early-B stars. Also, new studies are indeed filling the Mason period gap: \\citet{Gameetal07} reported 11 new massive spectroscopic binaries and \\citet{Maizetal08b} cite two new visual systems in the solar neighborhood (see also \\citealt{Nelaetal04}). With respect to the secondary mass distribution function, \\citet{Kouwetal05} studied intermediate-mass stars in the Sco OB2 association and found that they could exclude random pairing in binaries, since the more massive stars clearly tend to associate with other massive objects. Another recent study of Cyg OB2 by \\citet{KobuFrye07} also finds a large ($>$80\\%) binary fraction for massive stars. Those authors favor a secondary mass distribution function that is either (a) nearly flat or (b) a combination of a 60\\% Kroupa-like distribution and a 40\\% ``twin'' ($M_2 > 0.95 M_1$) component.  But what about more massive clusters? Is it possible that the multiple fraction there is significantly lower? Studying distant objects is more complicated but a recent survey by \\citet{Evanetal06} found that NGC 346 (in the SMC), NGC 2004, and N11 (in the LMC) have a minimum spectroscopic binary fraction for O- and early B-type stars of 23 to 36\\%. Considering that their program did not allow them to detect binaries with periods of more than several tens of days and that at the distance of the Magellanic Clouds it is very hard to resolve visual binaries, the real binary fraction should be considerably higher. Therefore, it is safe to conclude that the effect of unresolved massive binaries in the observed IMF is important and needs to be analyzed. Such an analysis is the first goal of this paper.  \\subsection{The stellar upper mass limit}  From the point of view of the origin of the IMF, it has become apparent in the last decade that massive stars must form in a different way than low-mass objects because of the short time scales involved and the limiting effect of radiation pressure on the mass accretion rate \\citep{BallZinn05,ZinnYork07}. There are currently two viable alternatives to form a massive star: by accretion through a disk, either in a relatively isolated environment or in a more crowded one where competition for mass among different protostars may be important, and by stellar collisions followed by mergers. The second alternative can only be relevant when the stellar densities are very high, such as at the core of a dense and massive cluster, but cannot be significant in relatively low-density environments such as OB associations.  If stellar mergers are indeed important at the core of dense clusters, there should be large consequences on their stellar IMFs. \\citet{Portetal99} predicted that runaway collisions in such cores will lead to stars with masses above 100 \\Ms/ in a time scale of one or a few million years, which is comparable to the expected life times of those stars. The process may keep going on to the point of producing objects above 1000~\\Ms/ that would end up as intermediate-mass black holes \\citep{Portetal04a}. Therefore, massive dense clusters should show an stellar upper mass limit $m_{\\rm max}$ significantly higher than those of lower density or mass. Can such stellar behemoths exist? Recently, \\citet{Belketal07} and \\citet{Yungetal08} have attempted modeling of solar-metallicity stars up to 1000~\\Ms/, which would be formed by mergers. However, such modeling encounters problems with the determination of the Eddington limit, its associated instabilities and interaction with rotation, and with the large extrapolation required for mass-loss rates (which are already uncertain at the present time even for run-of-the-mill main-sequence O stars). Hence, theory has not given its final word on the possible existence of such ultramassive stars.  What do observations tell us about those hypothetical ultramassive stars? Several papers on R136 and other Local Group massive young clusters \\citep{WeidKrou04,Fige05,OeyClar05,Koen06} indicate that $m_{\\rm max} \\approx$ 120-200 \\Ms/ when evolutionary tracks are used to obtain stellar masses\\footnote{Those values are called evolutionary masses and refer to the initial mass of the star.}. Some luminous blue variables (LBVs) such as $\\eta$~Car could have slightly higher evolutionary masses but LBVs are notoriously difficult to observe in detail due to the existence of circumstellar material. With respect to the more accurate masses measured from the orbital motion of the stars (keplerian masses), the highest published values correspond to WR 20a, a binary with 83.0~$\\pm$~5.0~\\Ms/ and 82.0~$\\pm$~5.0~\\Ms/ (\\citealt{Bonaetal04}, see also \\citealt{Rauwetal04}). Therefore, on a first impression it appears that observations do not favor the existence of ultramassive stars at near-solar metallicity.  Unresolved multiple systems can also play a role in the identification of ultramassive stars by making e.g. a binary made out of two stars close to the stellar upper mass limit appear as a single object with an evolutionary mass well above that limit. For example, \\citet{Maizetal07} found that Pismis 24-1, formerly an ultramassive star candidate, is in reality composed of three stars with masses below 120 \\Ms/. It is also possible that the observed $m_{\\rm max}$ measured from evolutionary masses is affected not only by real (bound) multiple systems but also by chance alignments in the crowded environments where some of these objects are found. Analyzing the possibility that some of the objects observed close to $m_{\\rm max}$ may be blended sources composed of stars of significantly lower mass is the second goal of this paper. ",
        "conclusions": ""
    },
    "0801/0801.3544_arXiv.txt": {
        "abstract": "We aim to understand the correlation between cloud formation and alkali line formation in substellar atmospheres.We perform line profile calculations for Na~I and K~I based on the coupling of our kinetic model for the formation and composition of dust grains with 1D radiative transfer calculations in atmosphere models for brown dwarfs and giant gas planets.  The Na~I and K~I line profiles sensibly depend on the way clouds are treated in substellar atmosphere simulations. The kinetic dust formation model results in the highest pseudo-continuum compared to the limiting cases. ",
        "introduction": "Dust influences the environment from which it forms by consuming elements from the gas phase. Hence, dust selectively alters the local metallicity in the atmosphere of a brown dwarf as well as in planetary atmospheres (Woitke \\& Helling 2004, Helling, Woitke \\& Thi 2008).  Dust clouds furthermore efficiently absorb photons which can result in $\\mu$m-broad absorption features (Helling et al. 2006, 2007; Cushing et al. 2006, Burgasser et al. 2007). The absorbed energy is redistributed inside the lattice structure of the solid grain causing an isotropic irradiation of IR photons into the atmosphere, a process, which considerably alters the atmospheric temperature profile (Tsuji et al. 1996, Tsuji 2005, Ackerman \\& Marley 2001, Allard et al. 2001). Both processes, the selective alteration of local element abundances and the changing atmospheric temperature, have a distinct influence on the gas-phase chemistry which determines the atomic and molecular concentrations. These intricate effects are important even for elements not or only marginally involved into the dust formation process as we will show in this letter.  Alkali lines are dominant features in the optical spectrum of brown dwarfs and the far line wings of the most abundant alkali metals, sodium and potassium, determine the pseudo-continuum in this spectral range (Johnas et al. 2007a, Allard et al. 2007).  Another alkali, Li, is used to determine fundamental parameter of substellar objects (e.g. Martin, Rebolo \\& Maguzza 1994, Johnas et al. 2007c).  They all can be a powerful tool to study the cloud formation in substellar objects if combined with an atmosphere simulation which allows to correlate the line profiles with an analysis of the atmosphere regarding the $(T, p_{\\rm gas}, v_{\\rm conv})$ structure ( $T$ - gas temperature, $p_{\\rm gas}$ - gas pressure, $v_{\\rm conv}$ - convective velocity) and the chemistry.  In this letter we show that the modelling of dust formation has a strong impact on the line shapes and the depth of the Na~I and K~I alkali lines in dust-enshrouded L-dwarfs. We utilise our kinetic model for dust cloud formation in atmosphere simulations of substellar objects (brown dwarfs and extrasolar planets) in comparison to two limiting cases in order to demonstrate the sensitive dependence of alkali line profiles on the local $(T, p_{\\rm gas}, \\epsilon_i)$ profile ($\\epsilon_i$ - element abundance) of the atmosphere. We show that the treatment of dust cloud formation has a much stronger impact on the alkali line profiles than the details of the line profile calculation.   \\section[]{Atmosphere Models}  We present our studies based on simulations carried out with the general purpose model atmosphere code {\\texttt{PHOENIX}}\\footnote{This is the same code as used by Allard et al. (2001) except for the dust chemistry and line-profile calculations.}, version 15, (Hauschildt \\& Baron 1999). We combine the study of line profiles of Johnas (2007) with the kinetic dust cloud modelling by (Helling \\& Woitke 2007) which has recently been incorporated into {\\texttt{PHOENIX}} by Dehn (2007) (also Dehn et al. 2008).  We concentrate here on the following two different line profile approaches for which the model atmospheres are calculated: \\begin{itemize} \\item{\\sc impact -- line profiles:}\\\\ van der Waals profiles applying the impact approximation for near line wings (Schweitzer et al. 1996) \\item{\\sc modern1 -- line profiles: }\\\\ Detailed non-analytical, semi-classical line profiles (Allard et al. 2005, Allard \\& Spiegelman 2006, Johnas et al. 2006) including H$_2$ in two symmetries, $C_{2v}$ and $C_{\\infty v}$, and He as perturber. The alkali line profiles are represented by two terms, one describing the near line wing with the impact approximation (Baranger  1958b, Royer 1971) and the other describing the far line wing with the one--perturber approximation in the density expansion (Royer 1971).  \\footnote{ Johnas et al. (2007a,b)  % show that alternative alkali near line wing profiles (Mullamphy et al. 2007) % result in negligible effects on the atmospheric structure and synthetic spectrum.} \\end{itemize}  \\noindent For our study of the influence of different dust model approaches on alkali line profiles, we compare simulations for the two limiting cases {\\sc Cond} and {\\sc Dusty} (Allard et al. 2001), and our new {\\sc Drift} module (Dehn et al. 2008) which includes a detailed modelling of dust formation: \\begin{itemize} \\item{{\\sc Cond} approach:}\\\\ The {\\sc Cond} approximation assumes the dust to be in chemical and phase equilibrium with the surrounding gas phase.  The dust is ignored as opacity source because it is {\\sl assumed} that after the dust grains have formed they sink down below the photosphere, thus resulting in a depleted gas phase with no dust absorption or emission. This scenario was used to model T-type brown dwarfs (e.g. Lodieu et al. 2007) or planetary objects (e.g. Seifahrt et al. 2007) with T$_\\mathrm{eff}$ lower than $\\sim$1400K but does not allow to consider the cloud {\\sl formation} process.  These models are not realistic for the parameter range considered here and are included for comparison only.  \\item{{\\sc Dusty} approach:}\\\\ The dust is assumed to have formed in chemical and phase equilibrium, as in the {\\sc Cond} models but is {\\sl assumed} to remain at the place of formation. In both {\\sc Cond} and {\\sc Dusty}, the amount of elements bound in grains is subtracted from the gas phase. No dust settling is taken into account. The interstellar grain size distribution is used for the dust opacity calculation.  {\\sc Dusty} models were used at higher T$_{\\rm eff}$ representing the case of a thick cloud layers inside the atmosphere of L\\,-\\,dwarfs (e.g. Leggett et al. 2001) and do not consider the actual formation process.  \\item{{\\sc Drift} approach:}\\\\ The {\\sc Drift} package (Dehn et al. 2008) is based on a kinetic treatment of dust formation (Woitke \\& Helling 2003, 2004; Helling \\& Woitke 2006; Helling, Woitke \\& Thi 2008).  A stationary dust formation process is assumed for application in static model atmospheres. In these model atmospheres, seed particles (here: TiO$_2$) nucleate from the gas phase if an appropriate super-saturation is achieved, subsequently grow a mantle made of various compounds (SiO$_2$[s], Al$_2$O$_3$[s], Fe[s], MgO[s], MgSiO$_3$[s], Mg$_2$SiO$_4$[s], and TiO$_2$[s]) by 32 chemical surface reactions (Dehn 2007), settle gravitationally, and evaporate.  The dust formation cycle is completed by convective overshooting of uncondensed material which is modelled by an exponential decrease of the mass exchange frequency into the radiative zone (Ludwig et al. 2006). The material composition of the grains, the amount and the size of dust particles formed -- which are needed to evaluate the cloud's opacity in the radiative transfer calculation -- are results of our model. \\end{itemize} ",
        "conclusions": "\\label{chapt:con}   We have demonstrated that the dust treatment in an atmosphere simulation for substellar objects has a large influence on the resulting line-profile of the alkalies Na~I and K~I.  We conclude that the main cause for these large differences is the strongly changing local temperature which is directly linked to the dust treatment by the resulting dust opacity. The different treatments of the alkali line profiles result in negligible changes compared to the effect of dust treatment.  The pseudo-continuum of alkali line-profiles are strongest if based on a microphysical dust model taking into account the kinetic nature of dust formation ({\\sc Drift-Phoenix}). This findings may foster studies on Li~I and Rb~I lines since Johnas (2007) has shown that in particular the far K~I line wings can mask the Rb~I doublet, but both the Na~I and the K~I overlap the Li~I doublet cores."
    },
    "0801/0801.1407_arXiv.txt": {
        "abstract": "We study perturbation of holographic dark energy and find it be stable. We study the fate of the universe when interacting holographic dark energy is present, and discuss a simple phenomenological classification of the interacting holographic dark energy models. We also discuss the cosmic coincidence problem in the context of holographic dark energy. We find that the coincidence problem can not be completely solved by adding an interacting term. Inflation may provide a better solution of the coincidence problem. ",
        "introduction": "The cosmological constant problem\\cite{ccref} is a longstanding problem in theoretical physics, and has been taken more and more seriously since the discovery of accelerated expansion of our universe\\cite{supernova}\\cite{WMAP}\\cite{SDSS}. In addition to the problem why the cosmological constant is nonvanishing, there is also the ``cosmic coincidence problem''\\cite{ccref}.  Based on the validity of effective quantum field theory, Cohen {\\it et al} \\cite{Cohen:1998zx} pointed out that the quantum zero-point energy of a system should not exceed the mass of a black hole of the same size. This observation relates the UV cutoff of a system to its IR cutoff. As a cosmological application, Li \\cite{mli04} suggested to choose the future event horizon as the IR cut-off, the energy density of vacuum is given by \\begin{equation}\\label{rhod} \\rho_D=3c^2M_p^2R_h^{-2}~, \\end{equation} where $R_h\\equiv a \\int_{t}^{\\infty}dt'/a(t')$ is the size of future event horizon. This is called holographic dark energy (HDE) and has been studied extensively \\cite{refHDE}.  As a phenomenological model, the stability of holographic dark energy is an important issue and has been investigated by Myung \\cite{Myung:2007pn} first. Myung \\cite{Myung:2007pn} assumed that holographic dark energy is a usual fluid component. He calculated the square of sound speed of holographic dark energy and found it be negative, this leads to an instability of perturbation of holographic dark energy. However, holographic dark energy is given by the holographic vacuum energy, whose perturbation should be treated globally. A calculation of perturbation of holographic dark energy will be presented in sect.2. It is shown that perturbation of holographic dark energy is stable, and we do not need to face the negative sound speed square problem.  The interacting holographic dark energy, assuming that dark energy interacts with matter, has become a popular topic recently \\cite{interactingref}. With the interacting term, the story of holographic dark energy becomes more interesting. There is often an attractor solution to the evolution equation, in which the effective equations of state of dark energy and matter become identical in the far future. In sect.3, we will give a simple and phenomenological classification of the interacting terms, we will show that we can tune the interacting parameter to avoid the phantom-like universe.  In sect.4, we shall discuss the coincidence problem of holographic dark energy in a more natural way. We end this paper with conclusion and discussion in sect.5. ",
        "conclusions": "In this paper, we study perturbation of holographic dark energy. Since holographic dark energy is just the holographic vacuum energy, its perturbation is global. We calculate perturbation of the holographic dark energy, and find it stable.  Many numerical and  analytic works have been done on the interacting holographic dark energy. We made a simple and phenomenological classification of interacting holographic dark energy in this paper, and derived the sufficient and necessary condition of no phantom (the big rip). Needless to say, this classification has not been done previously. It is worth to note that we write the interacting term just by hand for lack of knowledge of the first principle of holographic dark energy. We hope we will return to this issue in future projects.  We also discussed the coincidence problem. It is shown that the interacting holographic dark energy approach only solves the coincidence problem partially.  The original solution to the coincidence problem proposed in \\cite{mli04} stands a better resolution."
    },
    "0801/0801.4825_arXiv.txt": {
        "abstract": "If a microlensing event is caused by a star, the event can exhibit change in color due to the light from the lens.  In the previous and current lensing surveys, the color shift could not be used to constrain the lens population because the blended light responsible for the color shift is mostly attributed to nearby background stars rather than the lens. However, events to be observed in future space-based surveys do not suffer from blending and thus the color information can be used to constrain lenses.  In this paper, we demonstrate the usefulness of future surveys in measuring color shifts. By conducting simulation of galactic lensing events based on the specification of a proposed space-based lensing  survey, we estimate that the shift in the color of $R-H$ will be measured at 5$\\sigma$ level for $\\sim 12\\%$ of events that occur on source stars with apparent magnitudes brighter than $J=22.5$.  Color-shifted events tend to have high magnifications and the lenses will have brightnesses equivalent to those of source stars.  The time scales of the color-shifted events tend to be longer than those without color shifts. From the mass distribution of lenses, we find that most of the color-shifted events will be produced by stellar lenses with spectral types down to mid M-type main sequence stars. ",
        "introduction": "One important characteristic of a microlensing event is that the color of the lensed star does not change during the lensing magnification.  However, events observed in actual lensing surveys can exhibit color shift.  This color shift occurs when the flux from the lensed source star is blended with the flux from other stars located in the seeing disk of the star.  The lensed and blended stars generally have different colors. Then, color shift occurs as the source/blend relative brightness changes during lensing magnification.   \\citet{kamionkowski95} and \\citet{buchalter96} pointed out that if the object responsible for a lensing event is a star, it is possible to obtain information about the lens from the measured color shift.  However, this idea could not be implemented for the actual analysis of lens population.  The most important reason for this is that under the previous and current lensing surveys conducted by using ground-based telescopes, the blended flux is mostly attributed to nearby background stars rather than the lens.   However, the situation will be different in future space-based lensing surveys, such as the {\\it Microlensing Planet Finder} ({\\it MPF}).  The {\\it MPF} is a proposed space mission to NASA's Discovery Program exclusive for microlensing and it will be equipped with a 1.1 m aperture telescope \\citep{bennett02, bennett04}. Such a survey is ideal for the measurement of the color shift in lensing events due to the following reasons.  First, thanks to the high resolution provided from space observation, the survey does not suffer from blending caused by background stars.  Then, one can safely attribute the observed color shift to the lens. Second, the fraction of color-shifted events would be higher in the space-based survey because it will monitor much fainter stars than current surveys.  As the monitored star becomes fainter, the relative flux of the lens to the source star increases and thus the event is more likely to produce large color shift. Third, with high photometric precision, the survey can detect much smaller color shift than ground-based surveys.  With the increased fraction of large color-shifted events and higher sensitivity to small color shift combined with nearly an order higher event rate, the future space-based lensing survey would be able to detect color shifts for a large number of events.   In this paper, we demonstrate the usefulness of future space-based lensing surveys in measuring color shifts of lensing events. For this, we estimate the fraction of events with measurable color shifts among the events to be detected from the proposed {\\it MPF} survey.  We also investigate the properties of the color-shifted events by conducting detailed simulation of galactic lensing events under realistic observational conditions. The fraction of color-shifted events was previously estimated by \\citet{buchalter96}.  However, their work was intended as a template for general comparison with various observational assumptions and thus is not focused on a specific survey. In addition, they did not provide information about the properties of color-shifted events.   The paper is organized as follows.  In \\S\\ 2, we introduce formalism used for the description of the color shift in general cases and under various limiting cases.  In \\S\\ 3, we describe the simulation conducted to estimate the fraction of color-shift events and their properties.  In \\S\\ 4, we present the resulting distribution of color shift.  We probe the properties of the color-shifted events by investigating the distributions of various quantities related to the events including the source and lens brightness, lens light fraction, magnification, lens mass, and event time scale.  In \\S\\ 5, we summarize and conclude. ",
        "conclusions": "We investigated the usefulness of future space-based lensing surveys in measuring color shifts of microlening events caused by the flux from the lens.  By conducting simulation of galactic microlensing events based on the specification of the proposed {\\it MPF} survey, we estimated that the shift in the color of $R-H$ would be measured at 5$\\sigma$ level for $\\sim 12\\%$ of events that would occur on source stars with apparent magnitudes brighter than $J=22.5$.  Color-shifted events tend to have high magnifications and the lenses would have brightnesses equivalent to those of source stars. We estimated that $\\sim 31\\%$ of the color-shifted events would have magnifications $A > 10$ and the lenses of $\\sim 45\\%$ of the color-shifted events would be brighter than source stars.  The time scales of the color-shifted events would be longer than those without color shifts because these events tend to be brighter and thus heavier than those without color shifts.  We estimated that the most frequent time scale of the color-shifted events would be $\\sim 18$ days, while the mode value of all events would be $\\sim 6$ days.  The spectral type of the lens stars of the color-shifted events would extend down to mid M type because the {\\it MPF} survey will observe in near infrared through which the lens/source flux ratio is not negligible.  Being able to extract information about the lens from the measured color shift, future space-based lensing surveys would be able to better constrain the lens population of galactic microlensing events."
    },
    "0801/0801.3128_arXiv.txt": {
        "abstract": "Simultaneous measurement of line-of-sight (LOS) magnetic and velocity fields at the photosphere and chromosphere are presented. Fe I line at $\\lambda6569$ and $H_{\\alpha}$ at $\\lambda6563$ are used respectively for deriving the physical parameters at photospheric and chromospheric heights. The LOS magnetic field obtained through the center-of-gravity method show a linear relation between photospheric and chromospheric field for field strengths less than $700$ G. But in strong field regions, the LOS magnetic field values derived from $H_{\\alpha}$ are much weaker than what one gets from the linear relationship and also from those expected from the extrapolation of the photospheric magnetic field. We discuss in detail the properties of magnetic field observed in $H_{\\alpha}$ from the point of view of observed velocity gradients. The bisector analysis of $H_{\\alpha}$ Stokes $I$ profiles show larger velocity gradients in those places where strong photospheric magnetic fields are observed. These observations may support the view that the stronger fields diverge faster with height compared to weaker fields. ",
        "introduction": "It is inferred through various observations that the magnetic field is playing a central role in the solar energetics that take place in the higher layers of the atmosphere (see for eg. \\citet{regnier06}). But, it is somewhat difficult to obtain a reliable vector magnetic field measurements at the chromosphere and corona \\citep{hector05, judge07}. The measurements of the photospheric magnetic fields are relatively well established. However, simultaneous vector magnetic field observations at the photospheric and chromospheric heights will give a better handle on the understanding of the magnetic structuring of the solar atmosphere. More direct and reliable measurements of the magnetic field at the chromosphere will also serve as a better boundary condition for extrapolating the magnetic fields to the coronal heights.  Measurements of the vector magnetic field in $H_{\\alpha}$ is particularly important in understanding the connection between photospheric, chromospheric and coronal magnetic field as its height of formation ranges almost from the upper photosphere to upper chromosphere \\citep{vernaza81}. Also, this is one of the most widely used spectral lines for the study of solar chromosphere (see a recent review by \\citet{rutten07}).  Comparison of the active region magnetic fields measured in $H_{\\alpha}$ and in the photosphere show a one-to-one correspondence in the weak field regions. But in the strong field regions, like umbra, there is a considerable deviation from the linearity \\citep{bala04, hanaoka05}. Infact the LOS magnetic field measured in $H_{\\alpha}$ weakens much faster for the corresponding strong field regions of the photosphere. At this point it may not be appropriate to demarcate the field strength above which this deviation happens because different observations show different deviation points. \\citet{gosain03} have also reported the systematic weakening of the magnetic field derived from Mg I $\\lambda5173/5184$ lines in comparison with that of Fe I $\\lambda6301.5/6302.5$ lines. In their case, the magnetic field measured by Mg I lines agree with the potential field extrapolation of the photospheric LOS magnetic field in weak field regime. In strong field regions there is a systematic shift towards lower values but still linear. Simultaneous observations of He I at $\\lambda10830$ and Si I lines at $\\lambda10827.1$ by \\citet{choudhary02} suggest that the field diverges faster in the upper layers of the chromosphere.  However there is no conclusive explanation for the observed weaker chromospheric magnetic fields for the corresponding strong photospheric fields. Various possibilities have been discussed by \\citet{bala04} \\& \\citet{hanaoka05} about weaker chromospheric fields observed in umbra. \\citet{hanaoka05} have discussed the possibility of scattered light and/or peculiarity of the atmosphere (radiative transfer effects) for the decrease in the polarization signal which in turn cause the underestimation of the magnetic field. \\citet{bala04} have suggested that the strongest fields measured at the photosphere diverge spatially and more quickly than the weak fields when propagating upward. In this paper, we address these issues from the point of view of the observed velocity gradients.  It is well known that, velocities and magnetic fields are coupled to each other in the solar atmosphere (see for eg. \\citet{rajaguru06}). Any change in the magnetic field is expected to alter the plasma motion and hence the observed velocity. If the field strength decreases with height then the velocity is expected to increase and vice versa, for the simple reason that the inhibition of the plasma depends on the magnetic field strength \\citep{spruit79}. It is also well known that the magnetic field strength in the umbral region at the photosphere is large. If the magnetic field measured with $H_{\\alpha}$ in the umbral region is weaker, then it shows that the magnetic field gradient is larger than expected which  implies larger velocity gradients in the umbral region compared to the penumbral region at chromospheric heights. In order to observationally verify this we have analyzed the velocity and magnetic fields estimated from the simultaneous spectropolarimetric observations of an active region using $H_{\\alpha}$ and Fe I spectral lines.  The spectropolarimetric observations which were carried out at the Kodaikanal solar telescope using a dual beam polarimeter are discussed in section(\\ref{sect:instr&obs}). Sections (3) and (4) discuss the data reduction and analysis procedures respectively. A comparison of LOS magnetic field at the chromosphere and photosphere is presented in section (\\ref{sect:losmag}). Velocity gradients obtained through bisector technique are discussed in section (\\ref{sect:velgrad}). ",
        "conclusions": "\\label{sect:concl}  In this paper we have used the multiwavelength spectropolarimetric tool to understand the stratification of the magnetic and velocity fields in the solar atmosphere. Out of the two lines considered for spectropolarimetry, one forms at photospheric height (Fe I $\\lambda6569$) and the other spans almost from the upper photosphere to upper chromosphere ($H_\\alpha$). Hence, these lines were useful in studying the connection between physical parameters at the photospheric and chromospheric heights. The main physical parameters studied in this paper are the magnetic and velocity fields in an active region.  As discussed in section \\ref{sect:losmag}, the LOS magnetic field measured in chromosphere is in general weak compared to its photospheric counterpart. The weakening of the chromospheric field is much faster for the corresponding strong photospheric field. The magnetic field strengths observed in the umbral chromosphere are much weaker than those expected from the extrapolation of the photospheric magnetic field.  For instance, the field strength inferred through  $H_{\\alpha}$ observations is about 400 G where as the field strength obtained through the extrapolation of the observed photosperic field under potential field approximation to an height of 2000 km is about 1000 G (assuming that the height of formation of $H_{\\alpha}$ is about 2000 km). Various possibilities have been discussed which can cause the weaker fields observed in the umbral chromosphere. The most probable ones are the fast divergence of the stronger fields when they propagate upward in the atmosphere \\citep{bala04} and  ambiguity in the height resolution of $H_{\\alpha}$ magnetic sensitivity  which may be photospheric and/or chromospheric depending on the region of observation \\citep{hector04}. If former is the reason, then it can explain the observed properties of both velocity and the magnetic fields presented in this paper. This is because, as shown in Fig. \\ref{fig3}, observed chromospheric fields are systematically weaker at the locations where strong photospheric fields are observed. If the weaker field strengths observed at the umbral chromosphere are truly of solar origin then this implies that the umbral fields decrease more rapidly with height compared to penumbral fields. Rapid decrease in field strengths cause rapid increase in velocity with height as there is a small upflow in the umbral photosphere. Observations also show that the velocity increases more rapidly in the umbral region compared to penumbral region (section \\ref{sect:velgrad}). Earlier observations by \\citet{gosain03} have also indicated the weakening of the magnetic field which is larger for stronger fields. In their observations quicker weakening of the stronger fields is not apparent, probably, because of the lines (Mg I  b1 and b2 at $\\lambda5173$ and $\\lambda5184$) used to infer chromospheric magnetic field that originate at the lower chromosphere. While the observations based on $H_{\\alpha}$ presented in this paper as well as earlier by \\citet{bala04,hanaoka05} show clearly that the strong fields weaken quickly with height because in these works, to infer chromospheric magnetic field only the spectral region close to its line core is considered which samples mostly the higher layers of the chromosphere \\citep{rutten07}. Hence there is a possibility that the weaker LOS field strengths observed in the umbral chromosphere are caused due to the faster divergence of the stronger fields. However, we would like to caution that the reliability of the COG method in estimating the LOS magnetic field strengths discussed in this paper (section \\ref{sect:losb}) is for simple solar atmospheric model. But, in reality the solar atmosphere may be complicated. More studies on the magnetic and velocity response functions for $H_{\\alpha}$ in different regions will help in better interpretation of the observations. Simultaneous multiline spectropolarimetry, which includes $H_{\\alpha}$ and at least one more line which is formed at the chromosphere (preferably Infrared $Ca$ Triplet line at $\\lambda$8542) and a photospheric line, is needed to get further insight into the physical processes that take place in the chromosphere."
    },
    "0801/0801.1625_arXiv.txt": {
        "abstract": "We use 2MASS and \\emph{MSX} infrared observations, along with new molecular line (CO) observations, to examine the distribution of young stellar objects (YSOs) in the molecular cloud surrounding the halo \\mbox{H\\,{\\sc ii}} region KR~140 in order to determine if the ongoing star-formation activity in this region is dominated by sequential star formation within the photodissociation region (PDR) surrounding the \\mbox{H\\,{\\sc ii}} region. We find that KR~140 has an extensive population of YSOs that have ``spontaneously'' formed due to processes not related to the expansion of the \\mbox{H\\,{\\sc ii}} region. Much of the YSO population in the molecular cloud is concentrated along a dense filamentary molecular structure, traced by C$^{18}$O, that has not been erased by the formation of the exciting O star. Some of the previously observed submillimetre clumps surrounding the \\mbox{H\\,{\\sc ii}} region are shown to be sites of recent intermediate and low-mass star formation while other massive starless clumps clearly associated with the PDR may be the next sites of sequential star formation. ",
        "introduction": "\\label{sec:intro}  This study is the fourth in a series of papers exploring the structure of the \\mbox{H\\,{\\sc ii}} region KR~140 and its associated star-forming activity. In \\citet*{ker99} VES~735, the O-star powering the \\mbox{H\\,{\\sc ii}} region was examined. A second paper, \\citet*{bal00}, was a multiwavelength study of the structure, energetics, and kinematics of the \\mbox{H\\,{\\sc ii}} region. Finally, \\citet{ker01} presented an analysis of submillimetre (submm hereafter) observations of the region at 450 and 850~$\\mu$m. The main result of \\citet{ker01} was the discovery of numerous submm dust cores located within the molecular gas surrounding the \\mbox{H\\,{\\sc ii}} region, including a large number of cores that were clearly located at the interface between ionized and molecular gas, a likely location for star-formation induced or ``triggered'' by the expansion of the \\mbox{H\\,{\\sc ii}} region \\citep{elm98}. Two of the more isolated cores were clearly associated with \\emph{IRAS} sources, but the low resolution of \\emph{IRAS} combined with the extensive diffuse emission of dust associated with the \\mbox{H\\,{\\sc ii}} region made it impossible to determine if any of the other cores were associated with star-formation activity.  The goal of this study is to determine the distribution of young stellar objects (YSOs) throughout the KR~140 molecular cloud. We focus especially on the embedded stellar content of the submm cores as these are potentially the youngest star-forming regions. We use the spatial distribution of the YSOs relative to the photodissociation region (PDR) surrounding the \\mbox{H\\,{\\sc ii}} region as a means of gauging the relative importance of spontaneous and sequential, or triggered, star formation in this region. To achieve this we have analyzed a combination of submm data from \\citet{ker01} and \\citet{moo07}, 2MASS (2 Micron All Sky Survey) and Midcourse Space Experiment (\\emph{MSX}) infrared data, and newly acquired $^{12}$CO, $^{13}$CO and C$^{18}$O molecular line data from the Five College Radio Astronomy Observatory (FCRAO). In the next section we review the pertinent parameters of KR~140 and provide information on the new observations of the region. In \\S~\\ref{sec:2mass} and \\S~\\ref{sec:dist} we first show how 2MASS data can be utilized to identify the YSO population and then examine its spatial distribution. Our results are discussed in \\S~\\ref{sec:discuss} and conclusions are presented in \\S~\\ref{sec:conclude}. ",
        "conclusions": "\\label{sec:conclude}  While much of the focus of studies of star formation in and around \\mbox{H\\,{\\sc ii}} regions is on the interface regions, our study of KR~140 illustrates that spontaneous star formation occurring away from the interface regions in these molecular clouds can be as important as sequential or triggered star formation related to the expansion of the \\mbox{H\\,{\\sc ii}} region. Our analysis shows that the molecular cloud associated with KR~140 contains an extensive population of YSOs that is not associated with the PDR.  We have also shown that this YSO population is not uniformly distributed though the molecular cloud but is concentrated in four distinct regions. Many of the YSOs are also clearly associated with a filamentary structure within the molecular cloud which is traced by C$^{18}$O emission. Finally, we see that both low-mass and intermediate-mass YSOs are forming throughout the molecular cloud. Particularly interesting is the region around KMJB~1 (\\emph{IRAS}~02174+6052) which contains a tight grouping of a number of embedded intermediate-mass and solar-mass YSOs including the very highly embedded intermediate-mass YSO 2MASS~02210610+6106043."
    },
    "0801/0801.1555_arXiv.txt": {
        "abstract": "We report the discovery of two concentric partial Einstein rings around the gravitational lens \\lname, as part of the Sloan Lens ACS Survey. The main lens is at redshift $\\zl=0.222$, while the inner ring (1) is at redshift $\\zsi=0.609$ and Einstein radius ${\\rein}_1=1.43\\pm0.01 \\arcsec$. The wider image separation (${\\rein}_2=2.07\\pm 0.02\\arcsec$) of the outer ring (2) implies that it is at higher redshift than Ring 1. Although no spectroscopic feature was detected in $\\sim9$ hours of spectroscopy at the Keck I Telescope, the detection of Ring 2 in the F814W ACS filter implies an upper limit on the redshift of $\\zsii\\lesssim6.9$. The lens configuration can be well described by a power law total mass density profile for the main lens $\\rho_{\\rm tot}\\propto r^{-\\gamma'}$ with logarithmic slope $\\gamma'=2.00\\pm0.03$ (i.e. close to isothermal), velocity dispersion $\\sigma_{\\rm SIE}=287\\pm5\\kms$ (in good agreement with the stellar velocity dispersion $\\sigma_{v,*}=284\\pm24\\kms$) with little dependence upon cosmological parameters or the redshift of Ring 2. Using strong lensing constraints only we show that the enclosed mass to light ratio increases as a function of radius, inconsistent with mass following light. Adopting a prior on the stellar mass to light ratio from previous SLACS work we infer that $73\\pm9\\%$ of the mass is in form of dark matter within the cylinder of radius equal to the effective radius of the lens. We consider whether the double source plane configuration can be used to constrain cosmological parameters exploiting the ratios of angular distance ratios entering the set of lens equations. We find that constraints for \\lname\\, are uninteresting due to the sub-optimal lens and source redshifts for this application. We then consider the perturbing effect of the mass associated with Ring 1 (modeled as a singular isothermal sphere) building a double lens plane compound lens model. This introduces minor changes to the mass of the main lens, allows to estimate the redshift of the Ring 2 $(\\zsii=3.1\\mypm{2.0}{1.0})$, and  the mass of the source responsible for Ring 1 $(\\sigma_{\\rm SIE,s1}=94\\mypm{27}{47}\\kms)$. We conclude by examining the prospects of doing cosmography with a sample of 50 double source plane gravitational lenses, expected from future space based surveys such as DUNE or JDEM. Taking full account of the uncertainty in the mass density profile of the main lens, and of the effect of the perturber, and assuming known redshifts for both sources, we find that such a sample could be used to measure $\\Omega_{\\rm m}$ and $w$ with 10\\% accuracy, assuming a flat cosmological model. ",
        "introduction": "\\label{sec:intro}  Measuring the mass distribution of galaxies is essential for understanding a variety of astrophysical processes. Extended mass profiles of galaxies provide evidence for dark matter either using rotation curves \\citep[\\eg][]{rubin80,albada85,swaters03}, weak lensing \\citep[\\eg][]{brainerd96,hoekstra04,sheldon04,mandelbaum06}, or dynamics of satellite galaxies \\citep[\\eg][]{prada03,conroy07} which is one of the main ingredients of the standard $\\Lambda$ cold dark matter ($\\Lambda$CDM) cosmological model. At galactic and subgalactic scales, numerical cosmological simulations make quantitative predictions regarding, e.g., the inner slope of mass density profiles and the existence of dark matter substructure. Precise mass measurements are key to test the predictions and provide empirical input to further improve the models.  Gravitational lensing has emerged in the last two decades as one of the most powerful ways to measure the mass distributions of galaxies, by itself or in combination with other diagnostics. Although strong gravitational lenses are relatively rare in the sky \\citep[$\\lesssim20$ per square degree at space-based depth and resolution;][]{marshall05,moustakas07}, the number of known galaxy-scale gravitational lens systems has increased well beyond a hundred as a result of a number of dedicated efforts exploiting a variety of techniques \\citep[\\eg][]{warren96,ratnatunga99,kochanek99,myers03b,bolton04,cabanac07}. The increased number of systems, together with the improvement of modeling techniques \\citep[\\eg][]{kochanek92,warren03,treu04,brewer06,suyu06,wayth06,barnabe07}, has not only enabled considerable progress in the use of this diagnostic for the study of the mass distribution of early and most recently late-type galaxies, but also for cosmography, i.e. the determination of cosmological parameters \\citep[\\eg][]{golse02a,soucail04,dalal05}.  Given the already small optical depth for strong lensing, the lensing of multiple background sources by a single foreground galaxy is an extremely rare event. At Hubble Space Telescope (HST) resolution (FWHM $\\sim0\\farcs12$) and depth (I$_{\\rm AB}\\sim27$) it is expected that one massive early-type galaxy (which dominate the lensing cross-section) in about 200 is a strong lens \\citep{marshall05}. Taking into account the strong dependence of the lensing cross-section on lens galaxy velocity dispersion ($\\propto \\sigma^4$), and the population of lens galaxies, we estimate that about one lens galaxy in $\\sim 40-80$ could be a double source plane strong gravitational lens (see appendix \\ref{append:prob}). For these reasons, at most a handful of double lenses are to be found in the largest spectroscopic surveys of early-type galaxies such as the luminous red galaxies of the Sloan Digital Sky Survey. However, future high resolution imaging surveys such as those planned for JDEM and DUNE \\citep{aldering04,refregier06} will increase the number of known lenses by 2-3 orders of magnitude \\citep{marshall05}, and hence should be able to provide large statistical samples of double source plane gravitational lenses, opening up the possibility of qualitatively new applications of gravitational lensing for the study of galaxy formation and cosmography.  We report here the discovery of the first double source plane partial Einstein Ring. The gravitational lens system \\lname~, was discovered as part of the Sloan Lens ACS (SLACS) Survey \\citep{bolton05a,bolton06,treu06,koopmans06,bolton07a,gavazzi07b}. The object was first selected by the presence of multiple emission lines at higher redshift in the residuals of an absorption line spectrum from the SDSS database as described by \\citet{bolton04} and then confirmed as a strong lens by high resolution imaging with the Advanced Camera for Surveys aboard HST. In addition to an Einstein ring due to the source (hereafter source 1) responsible for the emission lines detected in the SDSS spectrum, the Hubble image also shows a second multiply imaged system forming a broken Einstein Ring with a larger diameter then the inner ring (hereafter source 2). This configuration can only arise if the two lensed systems are at different redshifts and well aligned with the center of the lensing galaxy. It is a great opportunity that a double source plane lens has been found among the approximately 90 lenses discovered by the SLACS collaboration to date \\citep[][in prep.]{bolton08}.  The goal of this paper is to study and model this peculiar system in detail, as an illustration of some astrophysical applications of double source plane compound lenses, including i) the determination of the mass density profile of the lens galaxy independent of dynamical constraints; ii) placing limits on the mass of source 1 based on multiple lens plane modeling; iii) estimating the redshift of source 2 and the cosmological parameters from the angular distance size ratios. The paper is therefore organized as follows. Section \\ref{sec:data} summarizes the observations, photometric and spectroscopic measurements, and discusses the morphology of the lens system. Section \\ref{sec:model} describes our gravitational lens modeling methodology. Section \\ref{ssec:resodel} gives the main results in terms of constraints on the mass distribution of the lens galaxy and of source 1. Section \\ref{sec:beyond} discusses the use of double source plane lenses as a tool for cosmography using the example of \\lname~and also addresses the potential of large samples of such double source plane lenses for the same purpose. In section \\ref{sec:conclu} we summarize our results and briefly conclude.  Unless otherwise stated we assume a concordance cosmology with $H_0 = 70\\h\\,\\kms\\Mpc^{-1}$, $\\Omega_{\\rm m}=0.3$ and $\\Omega_\\Lambda=0.7$. All magnitudes are expressed in the AB system.   \\begin{figure*}[htb] \\centering \\includegraphics[width=\\hsize]{ima_field.eps} \\caption{\\small HST/F814W overview of the lens system \\lname. The right panel is a zoom onto the lens showing two concentric partial ring-like structures after subtracting the lens surface brightness.} \\label{fig:overview} \\end{figure*} ",
        "conclusions": "\\label{sec:conclu}  In this paper we report the discovery of the first galaxy-scale double lensing event made of a foreground lens galaxy at redshift $\\zl=0.222$, a first source at redshift $\\zsi=0.609$ (Ring 1) and a more distant source (Ring 2) with unknown redshift, despite an attempt to measure its redshift with deep optical spectroscopy using LRIS on the Keck I Telescope. The detection of Ring 2 in a single orbit HST-ACS F814W filter image, sets an upper limit to its redshift $\\zsii < 6.9$.  Modeling the geometry of the lensed features at different source planes we determine the mass density profile of the lens galaxy which is found to be close to isothermal. The best fit lens model predicts a stellar velocity dispersion in very good agreement with that measured from SDSS spectroscopy.  The model requires a relatively large amount of dark matter inside the effective radius $f_{\\rm DM,2D}(<\\reff)\\simeq 73\\pm9\\%$ (corresponding to a projected total mass-to-light ratio $M/L_V=11.54\\pm0.13 \\hmmolsun$), assuming the stellar mass to light ratio measured in paper IV. Along with the complex isophotes of the lens galaxy and the presence of several other (less luminous) galaxies at similar photometric redshifts, the high dark matter fraction suggests that the lens may be the central galaxy of a group scale halo. The high precision of this measurement -- far superior to that attainable from a single multiply imaged systems -- demonstrates that double source plane lenses are extremely valuable tools to study the mass profile of galaxies and groups.  In order to constrain the redshift of Ring 2 and assess the feasibility of determining cosmological parameters using double source plane lenses, we constructed multiple lens plane mass models. In that case the lensing effect of Ring 1 on Ring 2 is taken into account and modeled as a singular isothermal sphere. Although the extra mass component adds additional uncertainty to the derived $\\zsii$ and cosmological parameters, it provides a unique way to determine the total mass of the intermediate galaxy (the inner Ring). Discarding the family {\\it ii} range of solutions (disfavored by kinematics of stars on the main lens galaxy), the two-lens-plane mass model results can be summarized as follows:  \\begin{itemize}  \\item The redshift of Ring 2 is found to be $\\zsii=2.6\\mypm{1.0}{0.7}$. This is a genuine prediction that can be tested with the help of deep Hubble images at shorter wavelengths, using the dropout technique.  \\item No interesting constraints on cosmological parameters can be obtained from the lensing analysis of the system \\lname, due to the unknown redshift of Ring 2, the overall degeneracy of cosmography with the slope of the mass density profile between the rings, the degeneracy with the mass of the inner ring, and the suboptimal combination of lens and source redshifts.  \\item The velocity dispersion of Ring 1 is found to be $\\sigma_{\\rm SIE}=56\\pm30 \\kms$, in good agreement with the value expected based on the extrapolation of Tully-Fisher relation at this redshift \\citep{moran07b} and on weak-lensing measurements \\citep{hoekstra05a}. Given that lensing cross section increases with the fourth power of the velocity dispersion, individual lenses in this mass range are expected to be very rare. In addition, it is observationally more difficult to identify lensed background structures embedded in foreground late-type galaxies with small Einstein radius (both for imaging or spectroscopic lens searches). Thus by exploiting the boost of the primary lens, double source plane, also seen as double lens plane systems, may be an effective way to determine the lensing mass of small distant galaxies, complementing detailed photometric studies \\citep{marshall07}, and kinematic studies with integral field spectrographs on large ground based telescopes with adaptive optics.  \\end{itemize}  Future planned space missions like JDEM or DUNE are expected to deliver several tens of thousands of single source plane lenses \\citep{aldering04,marshall05,refregier06} and several tens of double source plane lens galaxies. Given the great utility of multiple source plane lenses as tools to study distant galaxies, and the relatively small number of expected systems, we argue that the necessary effort of spectroscopic follow-up would be easily affordable and well motivated.  In addition, a relatively large sample of double source plane galaxy-scale gravitational lenses will be a practical tool for cosmography. As an example, we calculated the constraints on $\\Omega_{\\rm m}$ and the equation of state of Dark Energy $w=p_{\\rm DE}/\\rho_{\\rm DE}$ that can be obtained from a sample of 50 double source plane lenses, assuming both source redshifts are known and are realistically distributed. Spectroscopic follow-up of such systems rings is also required to control systematic effects such as the change in the mean density profile slope as a function of the lens galaxy redshift. A careful analysis taking into account the uncertainty on the mass profile of the main lens and of the perturber shows that cosmological parameters can be measured with an accuracy of 10\\% comparable to that obtained from the Hubble diagram of Type Ia supernovae."
    },
    "0801/0801.3285_arXiv.txt": {
        "abstract": "Feedback processes by active galactic nuclei (AGN) appear to be a key for understanding the nature of the very X--ray luminous cool cores found in many clusters of galaxies. We investigate a numerical model for AGN feedback where for the first time a relativistic particle population in AGN-inflated bubbles is followed within a full cosmological context. In our high-resolution simulations of galaxy cluster formation, we assume that BH accretion is accompanied by energy feedback that occurs in two different modes, depending on the accretion rate itself. At high accretion rates, a small fraction of the radiated energy is coupled thermally to the gas surrounding the quasar, while in a low accretion state, mechanically efficient feedback in the form of hot, buoyant bubbles that are inflated by radio activity is considered. Unlike in previous work, we inject a non-thermal particle population of relativistic protons into the AGN bubbles, instead of adopting a purely thermal heating. We then follow the subsequent evolution of the cosmic ray (CR) plasma inside the bubbles, considering both its hydrodynamical interactions and dissipation processes relevant for the CR population. This permits us to analyze how CR bubbles impact the surrounding intracluster medium, and in particular, how this contrasts with the purely thermal case. Due to the different buoyancy of relativistic plasma and the comparatively long CR dissipation timescale we find substantial changes in the evolution of clusters as a result of CR feedback. In particular, the non-thermal population can provide significant pressure support in central cluster regions at low thermal temperatures, providing a natural explanation for the decreasing temperature profiles found in cool core clusters. At the same time, the morphologies of the bubbles and of the induced X-ray cavities show a striking similarity to observational findings. AGN feedback with CRs also proves efficient in regulating cluster cooling flows so that the total baryon fraction in stars becomes limited to realistic values of the order of $\\sim 10\\%$, more than a factor of 3 reduction compared with cosmological simulations that only consider radiative cooling and supernova feedback. We find that the partial CR support of the intracluster gas also affects the expected signal of the thermal Sunyaev-Zel'dovich effect, with typical modifications of the integrated Compton-$y$ parameter within the virial radius of the order of $\\sim 10\\%$. ",
        "introduction": "\\renewcommand{\\thefootnote}{\\fnsymbol{footnote}} \\footnotetext[1]{E-mail: deboras@mpa-garching.mpg.de }  There is a growing body of observational evidence indicating that many galaxy clusters harbor in their centre supermassive black holes (BHs), which interact in a complex way with their surroundings. Recently, spectacular images from the XMM-Newton and Chandra X-ray telescopes \\citep[e.g.][]{McNamara2005, Fabian2006, Forman2006} have shed some light on this intricate interaction, showing clear imprints of central active galactic nuclei (AGN) in the intracluster medium of their host clusters. In particular, by comparing observations obtained at radio wavelengths with X--ray images, it is often seen that there are X--ray depressions in the innermost cluster regions which spatially correspond to significant radio emission \\citep[e.g.][]{Owen2000, Blanton2001, Clarke2005}. The most plausible candidate for inducing these features is the central BH, which is generating jet-inflated radio lobes, often simply dubbed `bubbles'.  There has been considerable theoretical effort to understand the interplay of AGN-driven bubbles with the intracluster medium (ICM) \\citep[e.g.][]{Churazov2001, Quilis2001, Brueggen2002, Ruszkowski2002, DallaVecchia2004, Omma2004a, Sijacki2006a, Sijacki2006b, Sijacki2007}. Most of these studies have, for simplicity, considered bubbles that are filled with hot thermal gas, and by following their evolution with time, they tried to constrain how much heat AGN bubbles can deliver to the surrounding medium. However, the synchrotron and inverse Compton emission detected from the radio lobes suggests that they contain a significant amount of relativistic electrons and that they are permeated with magnetic fields. Besides relativistic electrons and magnetic fields, it is also very plausible that an important part of the pressure support in the bubbles stems from relativistic protons, especially if the AGN jets are heavy. To date it is however not clear, both from an observational and a theoretical point of view, what the precise composition of the bubbles, the relative mixture of hot gas, and non-thermal particle components really is. Also, we have no detailed knowledge about the strength and configuration of the magnetic fields within the bubbles.   Regarding the contribution from hot thermal gas, the work of \\citet{Mazzotta2002} suggested that the bubble in the MKW 3s galaxy cluster may have a predominantly thermal origin, while \\citet{Schmidt2002} found that in the Perseus cluster a thermal gas component cooler than $\\sim 11\\,{\\rm keV}$ filling the bubbles seems ruled out. Moreover, it is not clear whether entrainment of thermal gas by the radio jet is effective and how this should depend on the BH properties. A study of a larger sample of galaxy clusters with central X-ray cavities by \\citet{Dunn2005} hints towards a rather complex picture, leaving several possibilities still open, both for significant pressure support by non-thermal protons or by hot thermal plasma, and for scenarios where the magnetic fields are filamentary at least in some of the cases.  Bearing in mind the outlined uncertainties it is still highly interesting to investigate the possible effects of cosmic rays (CRs) in AGN-inflated bubbles on cool core clusters, as considered already by several authors \\citep[e.g.][]{Boehringer1988, Loewenstein1991, Mathews2007, Guo2007, Sanders2007, Ruszkowski2007}. These studies have highlighted different channels of CR interaction with the thermal intracluster gas and have demonstrated that CRs can be dynamically important in galaxy clusters, contributing up to $50\\%$ of the central cluster pressure (but see \\citet{Churazov2007} for somewhat lower estimates of CR pressure support of order of $10-20\\%$). Nevertheless, due to the complicated nature of CR physics, most of these works were based on analytical or 1D numerical calculations, and neglected any cosmological evolution of the host clusters, and hence any possible dependence on its dynamical state. On the other hand, cosmological simulations of CRs produced at structure formation shocks by, for example, \\citet{Miniati2001, Miniati2002, Ryu2003, Ryu2004} have highlighted the importance of a realistic cosmological setting for a more fateful generation, distribution and following evolution of CRs.  Recently, \\citet{Ensslin2007} have proposed a simplified formalism for the treatment of CR protons that it suitable for implementation and use in self-consistent cosmological codes, as subsequently demonstrated by \\citet{Jubelgas2007} and \\citet{Pfrommer2006}. In these numerical studies, the CR source processes considered were restricted to supernovae (SNe) and structure formation shock waves. While it was found that these CR sources can affect the interstellar medium of low mass galaxies significantly, they did not prevent the overcooling in the centres of galaxy groups and clusters. However, the possible impact of CR protons generated by central AGN has not been explored in these works, and this forms the primary objective of our current study.  This paper is organized as follows. In Section~\\ref{Methodology} we outline the methodology we have adopted to simulate CR bubbles and to follow their cosmological evolution. In Section~\\ref{Isolated}, we present test runs performed for isolated halo simulations in order to analyze our numerical model in detail, while the bulk of our results from cosmological simulations of galaxy cluster formation is described in Section~\\ref{Cosmological}. Finally, we discuss our findings and draw our conclusions in Section~\\ref{Discussion}. ",
        "conclusions": "\\label{Discussion}  In this study we have investigated a model for self-regulated BH feedback with non-thermal relativistic particles in AGN-inflated bubbles. To describe the cosmic ray particles, we have adopted the formalism developed by \\citet{Ensslin2007} and \\citet{Jubelgas2007} for the treatment of CR protons, and combined it with prescriptions for BH growth and feedback, outlined in \\citet{Springel2005b} and \\citet{Sijacki2007}. This allowed us to study the influence of CR bubbles on their host clusters in self-consistent cosmological simulations of galaxy cluster formation, and in particular, to compare with the case where the bubbles are filled with purely thermal gas. Our methods also allowed us to gain some insight on how strongly the bubble morphologies and their propagation through the cluster atmosphere are affected by a cluster's dynamical state and the gaseous bulk motions in the ICM.  We have found that simulations with CR bubbles much more faithfully resemble observational findings compared with the results for thermal bubbles. This is mostly driven by the softer equation of state of gas that is supported by non-thermal pressure, and by the longer dissipation timescale of the CR component relative to thermal cooling. As a result, CR bubbles can rise to the cluster outskirts and even leak into the surrounding intergalactic medium. Still, their pressure support becomes gradually less important as they move away from cluster central regions. We have also found that the bubble morphologies in forming clusters are rather complex, especially during major merger events, where bubbles can become significantly perturbed from their initially spherical shape and are displaced from their initial injection axis. This implies that based on the spatial position of bubbles alone it is not readily possible to infer whether the AGN jet has changed its orientation between two successive outbursts or whether bulk motions in the ICM have displaced buoyantly rising bubbles.   Interestingly, we have found in our cosmological simulations that neither the BHARs nor the SFRs are very sensitive to the nature of the bubble feedback. However, the temperature distribution in the ICM is affected noticeably, given that CR bubbles provide non-thermal pressure support in central cluster regions. Above all, this changes the galaxy cluster temperature profiles which decline towards the innermost regions, being in much better agreement with the observed temperature profiles of cool core clusters. Correspondingly, the thermal pressure distribution is altered as well by CR bubbles, with a reduction of the thermal pressure in the centre and an increase in the outskirts. This redistribution of the thermal pressure can modify the integrated Compton-$y$ parameter at a level of $\\sim 10\\%$, an effect that we expect to contribute to the scatter in the relationship between cluster mass and $Y$.   We have evaluated how many significant bubble outburst the central BH in our cluster is undergoing during the time interval from $z=2$ to $z=0$, when the BH is in the low accretion, `radio-mode' regime. Based on this we obtain an average duty cycle of $\\sim 10^8\\,{\\rm yrs}$. We note however that sudden inflows of gas towards the central BH triggered by merging activity -- especially at $z=0.6$, which corresponds to the last major merger of the host cluster -- can cause much more frequent bubble events than the average value during a limited period. There is hence considerable variability in the duty cycle related to the merger history of the host halo, an issue we plan to investigate in more detail in a forthcoming study.   In our simulations, the sizes of the bubbles are not determined by calculating the jet physics from first principles, given that we cannot resolve the initial stage of jet formation and bubble inflation by a jet even in the most advanced state-of-art cosmological simulations. Instead, we rely on simple physical parameterizations and observational guidance to inject the bubbles with plausible sizes. Certainly this is an important limitation of our work and may in part be responsible for the rather large bubble sizes of up to a few hundred kpc that we find in some cases. On the other hand, there are two additional factors that may favor rather large bubble dimensions in our current modelling. One is that we have kept the injection axis constant with time. In some cases, the bubble injection time matches the buoyant rise time in such a way that fresh CR protons are injected into an already existing bubbles, making them reach a larger size. If this is also realized in nature, it is potentially ambiguous whether large bubbles are created by particularly powerful outbursts or result from several smaller outbursts which combine together.  However, we note that if CR diffusion is efficient (which we have not considered in this study) it would be less probable that several bubble outbursts mix together into a large X--ray cavity. In this case an old bubble would already disperse before a new one could have time to reach it and combine with it. CR diffusion would also reduce the probability of finding a bubble at large enough distance from the cluster centre (and would also possibly modify the heating rates), thus more sophisticated studies and more stringent observational constrains are needed to pin-down the relevance of CR diffusion out of the bubbles.  The second factor which may bias the sizes of our simulated CR bubbles high is the fact that observationally only relativistic electrons have been detected so far in radio lobes, while we have conjectured in our simulations that there should be a spatially extended distribution of CR protons as well. Many of our CR bubbles, especially the ones that are further away from the central AGN, are expected to have a rather aged population of CR electrons, which would not have detectable radio emission. In our picture these bubbles correspond to the so-called ghost cavities which are identified as depressions in X--ray emission without any notable radio emission, or only radio emission at lower frequencies. Because of the very low contrast in projected surface brightness maps, many of these ghost cavities are missed both in observations and simulations. It will be an important task for future work to better quantify this bias, and to find powerful observational constraints for the relative content of relativistic electrons and protons in AGN-inflated bubbles. Our simulations certainly suggest that CRs may play a crucial role in shaping the central ICM properties and in regulating AGN activity in clusters of galaxies."
    },
    "0801/0801.2884_arXiv.txt": {
        "abstract": "The solar magnetic field is key to understanding the physical processes in the solar atmosphere. Nonlinear force-free codes  have been shown to be useful in extrapolating the coronal field from underlying vector boundary data \\citep[see][for an overview]{Schrijver:etal06}. However, we can only measure the magnetic field vector routinely with high accuracy in the photosphere with, e.g., Hinode/SOT, and unfortunately these data do not fulfill the force-free consistency condition as defined by \\cite{aly89}. We must therefore apply some transformations to these data before nonlinear force-free extrapolation codes can be legitimately applied.  To this end, we have developed a minimization procedure  that uses the measured photospheric field vectors as input to approximate a more chromospheric like field \\citep[The method was dubbed preprocessing. See][for details]{wiegelmann:etal06}. The procedure includes force-free consistency integrals and spatial smoothing. The method has been intensively tested with model active regions \\citep[see][]{metcalf:etal07} and been applied to several ground based vector magnetogram data before. Here we apply the preprocessing program to photospheric magnetic field measurements with the Hinode/SOT instrument. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0801/0801.2410_arXiv.txt": {
        "abstract": "We present a structure finding algorithm designed to identify galaxy groups in photometric redshift data sets: the probability friends-of-friends (pFoF) algorithm. This algorithm is derived by combining the friends-of-friends algorithm in the transverse direction and the photometric redshift probability densities in the radial dimension. The innovative characteristic of our group-finding algorithm is the improvement of redshift estimation via the constraints given by the transversely connected galaxies in a group, based on the assumption that all galaxies in a group have the same redshift. Tests using the Virgo Consortium Millennium Simulation mock catalogs allow us to show that the recovery rate of the pFoF algorithm is larger than 80\\% for mock groups of at least $2\\times10^{13}M_{\\sun}$, while the false detection rate is about 10\\% for pFoF groups containing at least $\\sim8$ net members. Applying the algorithm to the CNOC2 group catalogs gives results which are consistent with the mock catalog tests. From all these results, we conclude that our group-finding algorithm offers an effective yet simple way to identify galaxy groups in photometric redshift catalogs. ",
        "introduction": "Galaxy groups are sites where local galaxy number density is relatively higher than the field. The majority ($\\sim60\\%$) of galaxies lies in groups \\citep[e.g.,][]{2004MNRAS.348..866E,2006ApJS..167....1B,2006AN....327..365T}, so that galaxy groups provide an excellent location to study the effect of local environment on galaxy formation and evolution. Unlike galaxy clusters, galaxy groups, especially those at high redshift, are not easy to detect because of their smaller size and the significantly lower hot gas density. The current published galaxy group catalogs are constructed based on large-scale galaxy redshift surveys using automated group finding schemes. The techniques include the popular friends-of-friends algorithm \\citep[e.g.,][]{1983ApJS...52...61G, 2005ApJ...630..759M, 2004MNRAS.348..866E} and the Voronoi partition technique \\citep[e.g.,][]{2005ApJ...625....6G}. Most of these catalogs list galaxy groups either in the nearby Universe ($ z < 0.1$) or over a small sky area. Galaxy groups of large sample sizes in intermediate and higher redshift space still remain largely unexplored.  Up to now, most structure finding techniques use spectroscopic redshift or simulated catalogs, both containing accurate three-dimensional position information. With the development of the photometric redshift method, the approximate redshifts of all galaxies in a photometric multi-band survey can be obtained without the time-consuming spectroscopic measurements. The photometric redshift method involves either SED (spectral energy distribution) fitting \\citep[e.g.,][]{2000A&A...363..476B,2003ApJ...586..745C,2004ApJ...600L.167M, 2004ApJS..155..243W,2005ApJ...619L..31B,2006ApJS..162...20B} or the use of a spectral `training set' to compute the photometric redshift via an empirical polynomial of galaxy colors and magnitudes \\citep[e.g.,][]{1995AJ....110.2655C, 2005ApJS..158..161H}. Since the redshifts are derived from broadband galaxy colors rather than from spectra, the photometric redshift method can estimate the redshift of objects which are too faint for spectroscopy. On the other hand, photometric redshifts have larger uncertainties by a factor of $50-100$ than those measured from spectroscopy. Due to the less accurate distance information in photometric redshift catalogs, the main problem of structure finding is the blurring of configurations in redshift space, producing unrealistic or elongated structures caused by the large photometric redshift uncertainties \\citep{2004MNRAS.349..425B}. Even with excellent photometric redshift estimation ($\\sigma_z \\sim 0.03$), the structures on the small scale will still be largely smeared out. Furthermore, projection effects make the subtraction of foreground and background galaxy contamination important in analyzing structures found using photometric redshift.  In order to overcome some of these problems, we propose here a method of finding galaxy groups in photometric redshift catalogs. The knowledge of galaxy photometric redshift uncertainty or probability density is required for this method. This group finding methodology is based on the idea of the standard friends-of-friends algorithm in the transverse direction, but takes into account the photometric redshift probability density to determine the friendship in the radial direction. We describe the photometric redshift technique and the error estimation for individual galaxies in $\\S$\\ref{sec_photoz}, and present our photometric sample selection criteria in $\\S$\\ref{sec_sample}. The group-finding parameters and the algorithm are detailed in $\\S$\\ref{sec_pFoFpar} and $\\S$\\ref{sec_pFoF}. The basic properties of galaxy groups are quantified in $\\S$\\ref{sec_grp}. This algorithm is tested in $\\S$\\ref{sec_test} using mock catalogs constructed from the \\it Virgo Consortium Millennium Simulation \\rm \\citep{2005Natur.435..629S}, and applied to the real observed groups in the \\it Canadian Network for Observational Cosmology Survey \\rm \\citep[CNOC2;][]{2000ApJS..129..475Y} in $\\S$\\ref{cnoc2grp}. Finally, we present a summary in $\\S$\\ref{sec_conclusion}. The analyses of galaxy group samples from a number of surveys will be presented in future papers. We adopt the standard cosmological parameters of $H_0$=70 km/s/Mpc, $\\Omega_m=0.3$, and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "} We have presented a new group-finding algorithm, pFoF, for identifying galaxy groups using photometric redshift catalogs. We have tested our pFoF algorithm on both mock catalogs and the CNOC2 groups. We summarize the most relevant results below.  Using the sample in which the simulated photometric redshifts resemble the real data, the comparisons between the pFoF and mock groups show that our algorithm produces reasonable results: (1) more than $80\\%$ of the mock groups with $1.2 \\times 10^{13} M_{\\sun}$ halo mass are recovered, (2) the fraction of false groups is $10\\%$ for the groups of $N_{gal}\\geq 7.91$, and (3) $\\sim 5\\%$ of pFoF groups are flagged as `serious projection' for which the pFoF group members are contributed by multiple mock groups. We find that the pFoF results strongly depend on the sample depth. The samples should be sufficiently deep ($\\sim M^*_{R_c}+1.5$) into the luminosity function for reliable group finding results. The use of samples with accurate redshift measurements reveals that the false detection rate depends strongly on the photometric redshift measurement accuracy. Application of the pFoF algorithm to the RCS-CNOC2 patches shows good agreement for the CNOC2 groups with $0.19 \\leq z < 0.41$.  The basic working principle of our pFoF algorithm is to improve the group redshift by joining new members. The average uncertainty in the estimated pFoF group redshift in our mock group tests is $\\sim 0.017$, compared with the average uncertainty of 0.070 for the photometric redshifts of individual galaxies. While such group redshift uncertainty is still very large compared with groups spectroscopically identified, our results show that our pFoF algorithm reduces the photometric redshift uncertainties significantly.  With our test results, we have demonstrated that our group-finding algorithm is able to identify galaxy groups with the capability of dealing with photometric redshift uncertainties. The purpose of this paper is to provide a method for searching galaxy groups (and clusters) in photometric redshift data sets as the first in a series of papers. We will apply this pFoF algorithm to the CNOC1 and RCS data sets. These data sets will provide us with a large sample of galaxy groups at $0.2 \\leq z < 0.6$, and enable us to study environmental dependence of galaxy properties and their evolution."
    },
    "0801/0801.1143_arXiv.txt": {
        "abstract": "We describe a sample of thermally emitting neutron stars discovered in the \\rosat\\ All-Sky Survey.  We discuss the basic observational properties of these objects and conclude that they are nearby, middle-aged pulsars with moderate magnetic fields that we see through their cooling radiation. While these objects are potentially very useful as probes of matter at very high densities and magnetic fields, our lack of understanding of their surface emission limits their current utility.  We discuss this and other outstanding problems: the spectral evolution of one sources and the relation of this population to the overall pulsar population. ",
        "introduction": "While almost 2000 isolated neutron stars have now been discovered as radio pulsars, the total number in the Galaxy is much larger.  Radio pulsars emit pulsations for $\\sim10^7$~yr and are visible due to radio beams that subtend 1--10\\% of the sky, so the total number of neutron stars of all ages just in the local region of the Galaxy (where radio pulsars are detectable) should be $\\gsim 10^6$ \\citep[e.g.,][]{vml+04,fgk06}.  Are those neutron stars not detectable as radio pulsars objects invisible, or is there some chance of observing them (there is always the exception of Geminga, where we see no radio emission but is otherwise a standard pulsar)?  For years astronomers have proposed that a large fraction of these objects would be visible through one of two mechanisms: accretion \\citep{ors70,tc91,bm93} or cooling \\citep{ttzc00}.  The first mechanism could revive old, dead pulsars, while the second would primarily work for younger sources.  Both mechanisms, however, make the neutron stars visible in the soft X-ray regime, not in the radio regime that had dominated the study of neutron stars.      These  neutron stars should be identifiable by \\citep{tc91}: \\begin{enumerate} \\item Largely thermal emission peaking in the soft X-ray or far-UV band, requiring small hydrogen column densities to remain visible \\item The absence of bright optical counterparts \\item Significant ($\\gsim 0.1\\mbox{ arcsec yr}^{-1}$) proper motions \\item Preferred locations in the Galactic plane \\end{enumerate} The first two criteria relate to the spectra of the neutron stars, and serve to rule out the active galaxies and stars that dominate X-ray surveys \\citep{hg88}.  The third criterion reflects the proximity of the sources (with maximum distances of $\\sim 1$~kpc) and the large space velocities of known neutron stars \\citep{hllk05,fgk06}.  The final criterion comes from the Galactic nature of the sources, and is similar to the distribution of radio pulsars.   \\subsection{The Legacy of \\rosat} While some had anticipated up to 5,000 objects discovered through soft X-ray surveys \\citep{tc91,bm93}, the \\rosat\\ All-Sky Survey (RASS) discovered just over half a dozen\\footnote{The difference is largely from poor assumptions about the pulsar velocity distribution \\citep{nt99,ttzc00} and the effects of magnetic fields \\citep{is75,pnr+03}.} (as of 2007) nearby cooling neutron stars (see Tab.~\\ref{tab:sum}).  These neutron stars are then all the more valuable because of their rarity.  The first such source to be discovered was \\Rxjw\\ (hereafter \\rxjw). It was originally identified serendipitously as a soft, bright X-ray source with no obvious optical counterpart \\citep{wwn96}.  Its location in front of the R~CrA molecular cloud meant that it had to be nearby ($\\lsim 200$~pc \\citep{kh98}) and hence small (not a white dwarf or anything larger)---otherwise the X-ray emission would have been absorbed\\footnote{As it turned out, this argument was false, as observations of more distant stars did not have high extinction \\citep{vkk01}.  Nonetheless, \\rxjw\\ is quite close \\citep{kvka02,wl02}.}.  Confirmation of its nature came with the discovery of a very faint ($m_B\\approx 25.8$~mag), blue optical counterpart \\citep{wm97}.   \\begin{table} \\setlength{\\tabcolsep}{4pt} {% \\begin{tabular}{cccccccccccccc} \\hline \\tablehead{1}{c}{c}{RX J} & \\mc{3}{c}{\\textbf{Spin}\\tablenote{We give the spin-period, period derivative, and rms pulsed fraction.} } & & \\mc{5}{c}{\\textbf{Spectrum}\\tablenote{We give the hydrogen column density, blackbody temperature, \\xmm\\ EPIC-pn count-rate, central energy of any absorption features in the X-ray spectrum, and $B$-band magnitude.  We also give energies of any secondary lines, if known.  Note that all spectral estimates are covariant, and that the details of the X-ray absorption depend on the specific model.} }  && \\mc{2}{c}{\\textbf{Astrometry}\\tablenote{We give the proper motion and distance, using the parallax if known, else that inferred from \\citet{pph+07}.}} & \\tablehead{1}{c}{c}{References}  \\\\ \\cline{2-4} \\cline{6-10} \\cline{12-13} & \\tablehead{1}{c}{b}{$\\mathbf{P}$} & \\tablehead{1}{c}{b}{$\\mathbf{\\dot P}$} & \\tablehead{1}{c}{b}{PF} && \\tablehead{1}{c}{b}{$\\mathbf{N_{H,20}}$} & \\tablehead{1}{c}{b}{$\\mathbf{kT}$} & \\tablehead{1}{c}{b}{PN}&\\tablehead{1}{c}{b}{$\\mathbf{E_{\\rm abs}}$}& \\tablehead{1}{c}{b}{$\\mathbf{m_B}$} && \\tablehead{1}{c}{b}{$\\mathbf{\\mu}$} & \\tablehead{1}{c}{b}{d} \\\\ & \\tablehead{1}{c}{b}{(s)} & \\tablehead{1}{c}{b}{($\\mathbf{10^{-14}}$)} & \\tablehead{1}{c}{b}{(\\%)}& & \\tablehead{1}{c}{b}{($\\mathbf{\\mbox{cm}^{-2}}$)} & \\tablehead{1}{c}{b}{(keV)} & \\tablehead{1}{c}{b}{($\\mathbf{s^{-1}}$)} &\\tablehead{1}{c}{b}{(keV)}&\\tablehead{1}{c}{b}{(mag)}&& \\tablehead{1}{c}{b}{($\\mathbf{\\mbox{mas yr}^{-1}}$)} & \\tablehead{1}{c}{b}{(pc)} \\\\ \\hline 1856.5$-$3754 & \\phn7.06 & \\nodata & \\phn1  && 0.8 & \\phn62 &8.3 & \\nodata &25.2&& 333 & 160 & \\citenum{kvka02,bhn+03,tm07,vkk01,vkk07}\\\\ 0720.4$-$3125\\tablenote{The spectral parameters are average quantities.} & \\phn8.39 & \\phn7 & 11 && 1.0 & \\phn87 &7.6 &  0.3 &26.6 &&\\phn97& 360& \\citenum{kvkm+03,mzh03,kvk05,kvka07,vkkpm07,hztb04}\\\\ 1605.3+3249   & \\nodata & \\nodata &  $<3$  && 0.8 & \\phn93 &5.6 & 0.5(0.6,0.8) &27.2 && 155 & 390 & \\citenum{kkvk03,vkkd+04,haberl07,msh+05,zdlmt06}\\\\ 1308.6+2127   & 10.31 & 11 & 18 && 1.8 & 102 &2.5 & 0.2(0.4) &28.4\\tablenote{Inferred from STIS 50CCD.} && 200\\tablenote{Based on Motch~et~al., these proceedings.} & \\nodata& \\citenum{kkvk02,kvk05b,hsh+03,shhm05,shhm07}\\\\ 2143.0+0654 & \\phn9.44 & \\nodata & \\phn4 && 3.6 & 102 &2.0 & 0.7 &$>26$\\tablenote{Inferred from $V$ and $r^{\\prime}$ bands.} && \\nodata & 430 & \\citenum{zct+01,zct+05,rtj+07}\\\\ 0806.4$-$4123 & 11.37 & \\nodata & \\phn6 && 1.1 & \\phn92 &1.8 &  0.3(0.6) &$>24$&&\\nodata & 250 & \\citenum{hmz+04,haberl07}\\\\ 0420.0$-$5022 & \\phn3.45 & \\nodata & 17 && 2.1 & \\phn45 &0.2 & 0.3 &26.6 &&\\nodata & 345 & \\citenum{hmz+04,haberl07}\\\\ \\hline \\end{tabular}} \\caption{Observed Properties of the Seven Isolated Neutron Stars\\label{tab:sum}} \\end{table}   Since then, six other similar sources have been identified with considerable effort \\citep{hmb+97,hmp98,shs+99,mhz+99,hpm99,zct+01}. While a variety of names and acronyms exist for these objects, we call them simply ``Isolated Neutron Stars'' (INS).  Identification of additional sources that may still be present \\citep{rfbm03} in the \\rosat\\ Bright Sources Catalog (containing $\\approx 18000$ sources with $>0.05\\mbox{ counts s}^{-1}$ with the Position-Sensitive Proportional Counter, or PSPC; \\citep{rbs}) is extremely difficult given the poor positional accuracy of the PSPC.  We will return to this later. ",
        "conclusions": ""
    },
    "0801/0801.1375_arXiv.txt": {
        "abstract": "The classical tests of general relativity (perihelion precession, deflection of light, and the radar echo delay) are considered for the Dadhich, Maartens, Papadopoulos and Rezania (DMPR) solution of the spherically symmetric static vacuum field equations in brane world models. For this solution the metric in the vacuum exterior to a brane world star is similar to the Reissner-Nordstrom form of classical general relativity, with the role of the charge played by the tidal effects arising from projections of the fifth dimension. The existing observational solar system data on the perihelion shift of Mercury, on the light bending around the Sun (obtained using long-baseline radio interferometry), and ranging to Mars using the Viking lander, constrain the numerical values of the bulk tidal parameter and of the brane tension. ",
        "introduction": "The idea that our four-dimensional Universe might be a three-brane \\cite{RS99a}, embedded in a higher dimensional space-time and inspired by superstring theory, has recently attracted much attention. In this context, the ten-dimensional $E_{8}\\times E_{8}$ heterotic string theory, which contains the standard model of elementary particles, could be a promising candidate for the description of the real Universe. This theory is connected with an eleven-dimensional theory, the $M$-theory, compactified on the orbifold $R^{10}\\times S^{1}/Z_{2}$ \\cite{HW96}. According to the brane-world scenario, the physical fields in our four-dimensional space-time, which are assumed to arise as fluctuations of branes in string theories, are confined to the three brane. Only gravity can freely propagate in the bulk space-time, with the gravitational self-couplings not significantly modified. The model originated from the study of a single $3$-brane embedded in five dimensions, with the $5D$ metric given by $ds^{2}=e^{-f(y)}\\eta _{\\mu \\nu }dx^{\\mu }dx^{\\nu }+dy^{2}$, which due to the appearance of the warp factor, could produce a large hierarchy between the scale of particle physics and gravity. Even if the fifth dimension is uncompactified, standard $4D$ gravity is reproduced on the brane. Hence this model allows the presence of large, or even infinite non-compact extra dimensions. Our brane is identified as a domain wall in a $5$-dimensional anti-de Sitter space-time. For a review of dynamics and geometry of brane Universes see \\cite{Ma01}.  Due to the correction terms coming from the extra dimensions, significant deviations from the Einstein theory occur in brane world models at very high energies \\cite{SMS00,SSM00}. In particular, gravity is largely modified at the electro-weak scale, 1TeV. The cosmological and astrophysical implications of the brane world theories have been extensively investigated in the literature \\cite{all}. Gravitational collapse can also produce high energies, with the five dimensional effects playing an important role in the formation of black holes \\cite{all1}.  Note that for standard general relativistic spherical compact objects the exterior space-time is usually described by the Schwarzschild metric. However, in five dimensional brane world models, the high energy corrections to the energy density, together with Weyl stresses from bulk gravitons, imply that on the brane the exterior metric of a static star is no longer the Schwarzschild metric \\cite{Da00}. The presence of the Weyl stresses also mean that the matching conditions do not have a unique solution on the brane; the knowledge of the five-dimensional Weyl tensor is needed as a minimum condition for uniqueness.  Static and spherically symmetric exterior vacuum solutions of the brane world models were initially proposed by Dadhich et al. \\cite{Da00} and Germani and Maartens \\cite{GeMa01}. The former solutions has the mathematical form of the Reissner-Nordstrom solution, in which a tidal Weyl parameter plays the role of the electric charge of the general relativistic solution  \\cite{Da00}. The solution was obtained by imposing the null energy condition on the 3-brane for a bulk having non-zero Weyl curvature, and a specific case was obtained by matching to an interior solution, corresponding to a constant density brane world star. An exact interior uniform-density stellar solution on the brane was also found in \\cite{GeMa01}. In the latter model, it was found that the general relativistic upper bound for the mass-radius ratio, $M/R<4/9$, was reduced by 5-dimensional high-energy effects. It was also found that the existence of brane world neutron stars leads to a constraint on the brane tension, which is stronger than the big-bang nucleosynthesis constraint, but weaker than the Newton-law experimental constraints \\cite{GeMa01}. We refer the reader to \\cite{branesolutions1,branesolutions2} and references therein for further static and spherically symmetric brane world solutions.   There are several possibilities of observationally testing the brane world models at an astrophysical/cosmological scale, such as using the time delay of gamma ray bursts \\cite{HaE} or by using the luminosity distance--redshift relation  for supernovae at higher redshifts \\cite{GeE}. The classical tests of general relativity, namely, light deflection, time delay and perihelion shift, have been analyzed, for gravitational theories with large non-compactified extra-dimensions, in the framework of the five-dimensional extension of the Kaluza-Klein theory, using an analogue of the four-dimensional Schwarzschild metric in \\cite{Li00}. Solar system data also imposes some strong constraints on Kaluza-Klein type theories. The existence of extra-dimensions and of the brane-world models can also be tested via the gravitational radiation coming from primordial black holes, with masses of the order of the lunar mass, $M \\sim 10^{-7}M_{\\odot }$, which might have been produced when the temperature of the universe was around 1TeV. If a significant fraction of the dark halo of our galaxy consists of these lunar mass black holes, a huge number of black hole binaries could exist. The detection of the gravitational waves from these binaries could confirm the existence of extra-dimensions \\cite{In03}.  It is the purpose of the present paper to consider the classical tests (perihelion precession, light bending and radar echo delay) of general relativity for static gravitational fields in the framework of brane world models. To do this we shall adopt for the geometry of the brane outside a compact, stellar type object, the spherically symmetric, Reissner-Nordstrom type, static brane metric obtained by Dadhich, Maartens, Papadopoulos and Rezania (DMPR for short) \\cite{Da00}. For this metric, we first consider the motion of a particle (planet), and the contributions of the five-dimensional effects to the perihelion precession are calculated.  By considering the motion of a photon in the static brane gravitational field we obtain the corrections, due to the projected bulk Weyl tensor, to the bending of light by massive astrophysical objects and to the radar echo delay, respectively. Existing data on light-bending around the Sun, using long-baseline radio interferometry, ranging to Mars using the Viking lander, and the perihelion precession of Mercury, can all give significant and detectable solar system constraints associated with the extra-dimensional part of the metric. More exactly, the study of the classical general relativistic tests, by taking into account the corrections coming from the extra dimensions, constrain the tidal bulk parameter and, via the junction conditions, the brane tension. The advance of the perihelion for the (charged) standard general relativistic Reissner-Nordstrom metric has been considered in \\cite{Ch01}, where the formula for the shift in the perihelion of a charged particle has been derived. The gravitational lensing by a Reissner-Nordstrom black hole in the weak field limit has also been analysed in \\cite{Sereno:2003nd}.  The present paper is organized as follows. The static and spherically symmetric vacuum solution on the brane is presented in Section \\ref{Sect2}. In Section \\ref{Sect3} we consider the classical solar system tests, namely, the perihelion shift, the light deflection and the radar echo delay, for the brane world model stars. We conclude our results in Section \\ref{Sect4}. ",
        "conclusions": "\\label{Sect4}  In the present paper, we considered the observational and experimental possibilities for testing at the level of the solar system the DMPR solution of the vacuum field equations in brane world models. The classical tests of general relativity in the solar system give strong constraints on the numerical values of the brane tension and of the bulk tidal parameter, respectively. Perihelion precession, light deflection and radar echo delay all give definite constraints on the numerical value of $Q$. While the two estimates obtained by using electromagnetic waves propagation data (light deflection and radar delay) are relatively consistent with each other, giving $\\left| Q\\right| \\leq 10^{18}-10^{19}\\,{\\rm cm}^{2}$, the perihelion precession gives a much stronger constraint $\\left| Q\\right| \\leq 6\\times 10^{7}-5\\times 10^{8}\\,{\\rm cm}^{2}$. An improvement of one order of magnitude in the observational data on Mercury's perihelion shift could provide a very precise estimate of the bulk tidal parameter.  Relative to the brane tension, the numerical value of $\\lambda$ has been constrained by using big bang nucleosynthesis data, which gives $\\lambda \\geq 1$ MeV$^{4}$ \\cite{MaWa00}. A much stronger constraint for the brane tension has been obtained by null results of Newton's law at sub-millimeter scales. Bulk effects lead to the modification of Newton's law on the brane.  The computation of the Newton potential on the brane shows that the correction terms to the Newton potential involve a logarithmic factor \\cite{Ga00}. When the distance between two point masses is very small with respect to the AdS radius, the contribution of the Kaluza-Klein spectrum becomes dominant as compared to the usual inverse square law. This type of behavior of the Newtonian gravitational potential may be used to prove the existence of an extra dimension experimentally \\cite{Ga00}. The null results of deviations from Newton's law give the constraint $\\lambda \\geq 10^{8}$ GeV$^{4}$ \\cite{MaWa00}.  Junction conditions relate the tidal parameter to the brane tension and mass, and radius of a compact astrophysical object. This can be used to calculate the numerical value of the brane tension $\\lambda$. However, despite the fact that these relations have been derived for high constant density stars, for which general relativistic effects are extremely strong (which is definitely not the case for the Sun), their application to the case of the solar system can give some approximate estimates of $\\lambda$. Stronger estimate, however, are implied from the perihelion precession, giving $\\lambda \\geq 10^{5}\\,{\\rm MeV}^{4}$. All estimates could be very much improved by the use of observational data obtained for high density compact astrophysical objects, like neutron stars.  On the other hand, there are several other vacuum solutions of the spherically symmetric static gravitational field equations on the brane \\cite{branesolutions1}. Indeed, the effects due to the projections of the Weyl tensor specify the deviations of brane world models from general relativity. Since the generic form of the Weyl tensor in the full five-dimensional theory is yet unknown, the effects of known solutions must be studied on a case by case basis. While one can in principle constrain the projections, this only yields very mild constraints on the five-dimensional Weyl tensor.  The study of the classical tests of the general relativity could provide a very powerful method for constraining the allowed parameter space of solutions, and to provide a deeper insight into the physical nature and properties of the corresponding space-time metrics. Therefore, this opens the possibility of testing brane world models by using astronomical and astrophysical observations at the solar system scale. Of course, this analysis must be extended to the study of all vacuum solutions on the brane, which requires developing general methods for the high precision study of the classical tests in arbitrary spherically symmetric space-times. In the present paper we have provided some basic theoretical tools necessary for the in depth comparison of the predictions of the brane world model and of the observational/experimental results.  \\ack We thank Roy Maartens for helpful discussions. The work of TH was supported by the RGC grant No.~7027/06P of the government of the Hong Kong SAR. FSNL was funded by Funda\\c{c}\\~{a}o para a Ci\\^{e}ncia e a Tecnologia (FCT)--Portugal through the grant SFRH/BPD/26269/2006."
    },
    "0801/0801.2140_arXiv.txt": {
        "abstract": "We study and develop constraint preserving boundary conditions for the Newtonian magnetohydrodynamic equations and analyze the behavior of the numerical solution upon considering different possible options. ",
        "introduction": "Magnetic fields play an important role in the behavior of plasmas and are thought mediate important effects like dynamos in the core of planets and the formation of jets in active galactic nuclei and gamma ray bursts; induce a variety of magnetic instabilities; realize solar flares, etc. (see e.g. \\cite{book1,book2}). Understanding the role of magnetic fields in these and other phenomena have spurred through the years many efforts to obtain solutions of the magnetohydrodynamic equations. The non-linear nature of these equations limits the understanding that can be gained in a particular problem via analytical techniques. This implies that solutions for complex systems must be obtained by numerical means and a suitable numerical implementation must be constructed for this purpose. Such implementation must be able to evolve the solution to the future of some initial configuration and guarantee its quality. A delicate, subsidiary quantity,  can be monitored in part to estimate this. This quantity is the ``monopole constraint'' $\\partial_i B^i$ which must be zero at the analytical level for a consistent solution. This quantity is not a part of the main variables, rather it is a derived quantity which should be satisfied by a true physical solution. In practice, unless a numerical implementation  of the MHD equations is carefully designed this constraint can be severely violated. This, in turn, signals (and is sometimes the cause) of a degrading numerical solution. For these reasons, several approaches have been investigated and developed for guaranteeing a controlled behavior of this quantity. One such approach is known as the {\\it constraint transport technique}~\\cite{constrainttransport,toth} which adopts a particular algorithm that staggers the variables appropriately to ensure the satisfaction of the constraint at round-off level within Finite Difference and Finite Elements techniques. This approach has been quite successful in a number of applications across different disciplines and particularly relevant in astrophysics applications~\\cite{Gammie:2003rj,MHD_stone,MHD_hawley,MHD_shibata,MHD_shapiro,MHD_rezzolla,Li:2006jj,Obergaulinger:2005xd}. However, by design it imposes limits on the algorithmic options available to an implementation. This fact can be at odds, or introduce complications, with applications where adaptive mesh refinement is required and/or advanced numerical techniques (that exploit useful properties of the equations) are adopted. An alternative approach, which controls the constraint at truncation-error level maintains complete freedom in the numerical techniques to be adopted. This approach, referred to as {\\it divergence cleaning} puts the burden to control the constraint not on the algorithm to be employed but rather on the system of equations to be solved itself \\cite{kimclean,dedner_2002}. This is achieved by considering an additional variable suitably coupled to the system through another equation so that, through the evolution, the constraint behavior is kept under control. While this method has, to date, received less attention than the constraint transport one, successful applications in diverse scenarios have already illustrated its usefulness (e.g. \\cite{dedner_2002,balsarakim,toth,neilsen1,neilsen2}).  Regardless of the technique employed, boundary conditions can play a significant role and may spoil all efforts towards a correct implementation in the absence of boundaries. Clearly, even when a stable option is adopted, it need not be consistent with the goal of preserving the constraint (be it at round-off or truncation levels) and so carefully designed boundary conditions must be formulated. Commonly employed options include straightforward outflow-type conditions --which do not enforce the constraint-- or ``absorbing'' boundary conditions which aim to reduce the influence of spurious effects induced at the boundary in the numerical solution~\\cite{dednerbdry}.  In the present work we concentrate on formulating constraint preserving boundary conditions for the Newtonian ideal MHD equations (for systems with and without divergence cleaning). Such a task is intimately related to the hyperbolic properties of the equations and so we re-analyze the system of equations and discuss alternatives for controlling the constraints through the divergence cleaning technique.  We therefore organize our presentation along the following lines. In section II we analyze the system of equations and the formulation of boundary conditions beginning with a simplified model from which we draw the strategy to apply in the complete MHD system. Section III presents a series of tests that highlight the benefits gained by our construction. We conclude in section IV with some final comments. ",
        "conclusions": "We have investigated the numerical solution to the (ideal) Newtonian magnetohydrodynamics equations and developed boundary conditions which at the continuum level are constraint preserving and so would not introduce further violations in the computational domain. At the numerical level these conditions do introduce some truncation-error level violations which converge away with resolution. These boundary conditions developed are consistent with maximally dissipative ones and so the discrete energy of the system remains bounded. We also examined the constraint's behavior when enlarging the system so as to couple an extra field. The addition of such field, together with a suitable modification of the equations induces a non-zero speed in the propagation of the constraint and allows for driving its value to zero. We have studied, in particular, two of the many equivalent systems. One which is fully conservative --which is only weakly hyperbolic without the addition of the extra field which renders it symmetric hyperbolic-- and a Galilean invariant one which is symmetric hyperbolic --even without adding the extra field-- but is not expressible in conservative form.  In order to examine the solution's behavior we implemented the equations in Cartesian coordinates and with a spatial discretization that preserves the initial constraint error. This allows us to separate bulk from boundary effects and observe that the main violations do indeed occur at boundaries. To examine the boundary condition effects on the evolution we adopt three types of boundary conditions: \\begin{itemize} \\item The most direct one, and crudest, sets to zero all right hand sides in a buffer boundary region. This might lead to inconsistencies for one might be prescribing conditions to outgoing modes. This condition is essentially equivalent the most commonly employed one of flux-copying near boundary points. \\item The second one is the ``no-incoming'' boundary condition where one projects out all incoming modes in the evolution equations leaving the rest (tangential and outgoing modes) intact. The implementation of this condition is slightly delicate as modes can change directions over time, and so may turn from being incoming to being outgoing and vice-versa. \\item The third condition preserves the constraint in the sense that is deduced by setting to zero the incoming mode related to the constraint propagation. This condition is a modification to the previous one where now instead of projecting out all incoming modes, one incoming mode is fixed so that the incoming constraint mode is zero as a result. \\end{itemize}  Not surprisingly the first option displays the worst behavior for the constraint as not only nothing is done to minimize the constraint violations introduced through the boundaries but further inconsistencies in the solution are induced there. As a result, in the best case the constraint mode bounces back from the boundary into the integration region. The second option performed substantially better displaying a gain of an order of magnitude in constraint preservation. This is a result of the boundary condition allowing for one of the $\\phi$ modes --carrying constraint violations-- to leave  the integration region. With the third alternative an additional order of magnitude (at least) is gained with respect to the non-incoming condition. This condition not only allows for constraint violations to leave the computational domain but does not introduce significant violations at the boundary.  Thus we see that appropriately handling the boundary conditions the problem of  constraint behavior can be significantly controlled. It should be mentioned that we still do not have a mathematically rigorous proof that the constraint preserving boundary conditions is well posed. However, in simpler systems --like the toy model we discussed at the beginning of the paper-- a proof could be devised at least for the case where the eigenvalues do not change sign. Further work in this direction is needed to put the analysis and results obtained in stronger grounds."
    },
    "0801/0801.2971_arXiv.txt": {
        "abstract": "This is the second part of an \\ha\\ kinematics follow--up survey of the {\\it Spitzer\\/} Infrared Nearby Galaxies Survey (SINGS) sample. The aim of this program is to shed new light on the role of baryons and their kinematics and on the dark/luminous matter relation in the star forming regions of galaxies, in relation with studies at other wavelengths. The data for 37 galaxies are presented. The observations were made using Fabry--Perot interferometry with the photon--counting camera \\FM\\ on 4 different telescopes, namely the Canada--France--Hawaii 3.6m, the ESO La Silla 3.6m, the William Herschel 4.2m, and the Observatoire du mont M\\'egantic 1.6m telescopes. The velocity fields are computed using custom IDL routines designed for an optimal use of the data. The kinematical parameters and rotation curves are derived using the {\\it GIPSY} software.  It is shown that non--circular motions associated with galactic bars affect the kinematical parameters fitting and the velocity gradient of the rotation curves. This leads to incorrect determinations of the baryonic and dark matter distributions in the mass models derived from those rotation curves.  \\\\ ",
        "introduction": "Understanding star formation mechanisms and their physical connection to the interstellar medium (ISM) properties of galaxies is crucial for resolving astrophysical issues such as the physical nature of the Hubble sequence, the nature and triggering of starbursts, and the interpretation of observations of the high--redshift universe. These scientific objectives constitute the core science program of the {\\it Spitzer} Infrared Nearby Galaxies Survey (SINGS) \\citep{2003PASP..115..928K}.  Since the science drivers of the project are dependent of many variables, the SINGS project requires a comprehensive set of data including infrared imaging and spectroscopic data, broadband imaging in the visible and near--infrared as well as UV imaging and spectrophotometry.  The {\\it Spitzer} data and ancillary observations of the SINGS galaxies will also provide valuable tools for understanding the physics of galaxy formation and evolution.  The formation of individual stars from the collapse of dense molecular clouds is relatively well known, e.g. the intensity of star formation is strongly correlated with the column density of gas and stars \\citep{1998ApJ...498..541K}.  However, the large--scale processes driving star formation are still poorly understood.  For instance, the effect of gas dynamics on the regulation of star formation is somewhat unknown.  Since young stars are often associated with spiral arms, it is thought that protostars are formed from compressed gas along large--scale shock fronts. Moreover, the star formation history varies greatly along the Hubble sequence. Elliptical galaxies, being gravitationally supported by velocity dispersion, have exhausted their gas reservoir and hence star--forming processes have ceased \\citep{1998ARA&A..36..189K}. Nevertheless, some elliptical galaxies having a rotating disc are still forming stars (e.g. NGC 2974, \\citealt{2007MNRAS.376.1021J}). Spiral galaxies, on the other hand, continue to form stars and are supported by rotation.  Understanding the effect of gas dynamics on star formation would certainly help improving our understanding of galaxy formation. Indeed, a Schmidt law modulated by rotation seems to better fit the data than a simple (gas column density) Schmidt law \\citep{2003MNRAS.346.1215B, 2001ApJ...555..301M, 1989ApJ...344..685K}.  This paper presents the second part of an \\ha\\ kinematics survey of the SINGS galaxies. The 37 galaxies showing \\ha\\ emission were observed by means of Fabry--Perot (FP) interferometry. The paper is organized as follows. Section \\ref{observations} describes the composition of the SINGS sample and the hardware used for the observations.  The data reduction process, using a custom IDL pipeline designed for an optimal use of the data, is introduced in section \\ref{reduction}. In section \\ref{gipsy}, the details of the kinematical analysis that has been done on the SINGS galaxies are described.  Section \\ref{results} presents the observational results in the form of velocity fields, monochromatic maps, position--velocity diagrams, and rotation curves. In section \\ref{discussion}, the effect of the bar on the observed kinematics is discussed. Finally, the scientific applications of the kinematical results presented in this paper are reviewed in section \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  We have presented in this paper the second and last part of the \\ha\\ kinematics follow--up survey of the Spitzer Infrared Nearby Galaxies Survey (SINGS) sample. The goal of this kinematical follow up is to better understand the role of baryons and of the dark/luminous matter relation in star forming regions of galaxies.  The shape of the velocity field in the central galactic regions, drawn through its H$\\alpha$ component, is indeed directly related to the baryonic luminous disk and the star formation processes. The SINGS sample will provide a unique opportunity to link the kinematics with numerous observations and studies at other wavelengths.  The data have been obtained from high resolution Fabry--Perot observations using the \\FM\\ camera and a L3CCD detector. The SINGS sample of galaxies has been observed at the OmM 1.6m telescope, ESO La Silla 3.6m telescope, CFHT 3.6m telescope, and WHT 4.2m telescope.  The velocity fields were obtained by using a data reduction pipeline written in IDL and the rotation curves were computed by the \\textit{rotcur} task from the {\\it GIPSY} software. When fitting the kinematical parameters, care was taken to avoid the zones obviously affected by non--circular motions. However, we have demonstrated that for barred systems, different bar characteristics considerably modify the central velocity gradient of the computed rotation curves in the pure circular rotation hypothesis. Therefore, numerical modeling of barred galaxies is crucial in order to extract rotation curves that are truly representative of the gravitational potential and hence of the mass distribution in those galaxies. In the meantime, the dark matter distribution of the SINGS galaxies, using the rotation curves derived here, will be presented in a forthcoming paper.  Not only will these observations provide the high spatial resolution data needed for constraining the dark matter density profiles of galaxies, but they will be helpful for studying these profiles as a function of morphological type.  Furthermore, they will help to delineate the role of gas kinematics in regulating the star formation rate.  For instance, \\cite{2002Ap&SS.281..101P} have argued that the probability of collapse of molecular clouds leading to star formation is greatly enhanced in slowly rotating gas discs compared to high velocity rotation.  Moreover, \\cite{2003A&A...405...89C} and \\cite{2007A&A...466..905F} have suggested that the star forming inner and nuclear rings in the nearby galaxies NGC 3627 and NGC 628 (respectively) are driven by a rotating asymmetry because their location in the host disk are in agreement with Lindblad resonances caused by a bar pattern speed. Gas dynamical processes therefore are important in regulating the star formation history of galaxies and the \\ha\\ kinematics presented in this paper will help understanding the star formation processes."
    },
    "0801/0801.0832_arXiv.txt": {
        "abstract": "We present results from a pilot {\\em HST} ACS deep imaging study in broad-band $V$ of five low-redshift QSO host galaxies classified in the literature as ellipticals. The aim of our study is to determine whether these early-type hosts formed at high redshift and have since evolved passively, or whether they have undergone relatively recent mergers that may be related to the triggering of the nuclear activity. We perform two-dimensional modeling of the light distributions to analyze the host galaxies' morphology. We find that, while each host galaxy is reasonably well fitted by a de Vaucouleurs profile, the majority of them (4/5) reveal significant fine structure such as shells and tidal tails.  These structures contribute between $\\sim$5\\% and 10\\% to the total $V$-band luminosity of each host galaxy within a region of $r \\sim$ 3\\,$r_{\\rm eff}$ and are indicative of merger events that occurred between a few hundred Myr and a Gyr ago. These timescales are comparable to starburst ages in the QSO hosts previously inferred from Keck spectroscopy. Our results thus support a consistent scenario in which most of the QSO host galaxies suffered mergers with accompanying starbursts that likely also triggered the QSO activity in some way, but we are also left with considerable uncertainty on physical mechanisms that might have delayed this triggering for several hundred Myr after the merger. ",
        "introduction": "Quasi-stellar objects (QSOs) are the most luminous active galactic nuclei (AGNs), believed to be powered by accreting supermassive black holes (BH). The growing observational evidence that massive BHs exist not only in the centers of active galaxies, but also in inactive early-type galaxies indicates that more than the mere presence of a massive BH is needed to trigger the activity. And while the ``bright quasar phase'' is now often cited as an essential phase in the evolution of galaxies \\citep[e.g.,][]{spr05a}, we still do not have a clear picture of what causes the triggering, or even what the necessary conditions are for the triggering to occur.  For more than two decades now, mergers have been suspected to be responsible for triggering QSOs \\citep{sto82,san88}.  There is an abundance of circumstantial evidence in the literature connecting QSO activity at low redshifts with mergers and interactions \\citep[e.g.,][]{hec84,hut88,can01,urr07}. Not only do a number of low-redshift QSOs show signs of interactions and mergers, but there is also a tendency for QSO host galaxies to exist in only moderately rich environments \\citep[e.g.,][see, however, \\citealt{mar07} for a counterexample]{dre85,ell91,wol01} consistent with the results of quasar clustering studies, which suggests that quasars exist in haloes of mass $\\sim$10$^{12}$ - 10$^{13}$ $M_{\\odot}$ \\citep{hop07a, daa07}. In such environments, velocity dispersions are moderate, and interactions  have their greatest effect.   At least in the case of QSOs found in ultraluminous infrared galaxies (ULIRGs), there is a close connection between mergers and QSO activity. ULIRGs are essentially always found in merging systems \\citep{san96}, and far too many of them harbor QSOs to be explained by the joint probability of finding a given $L^*$ or brighter galaxy to be both a ULIRG and a QSO host galaxy by chance \\citep{can01}.  So, at least for this subclass of QSO, some of the gas that the merger feeds into the central starburst also finds its way much deeper to the central black hole.  Nevertheless, recent high-resolution imaging studies show that a large fraction of QSOs reside in hosts with relaxed morphologies like elliptical galaxies \\citep[e.g.,][]{dis95,bah97,mcl99,flo04}.  These observations have elicited claims that tidal interactions in QSO host galaxies are the exception rather than the rule, and that even in those cases in which there are clear signs of interaction, the interactions are not necessarily related to the nuclear activity \\citep[e.g.,][hereafter D03]{dun03}. From two-dimensional fitting of a well matched sample of 33 AGNs (radio-loud and radio-quiet QSOs as well as radio galaxies), using images obtained with the Wide Field \\& Planetary Camera 2 (WFPC2) on board {\\em HST}, D03 concluded that ``for nuclear luminosities $M_V < -23.5$, the hosts of both radio-loud {\\em and} radio-quiet AGN are virtually all massive elliptical galaxies with basic properties that are indistinguishable from those of quiescent, evolved, low-redshift ellipticals of comparable mass''.  Such conclusions conflict with recent findings of significant amounts of dust and cold molecular gas, as well as young stellar populations, in many radio-loud and radio-quiet QSOs \\citep[e.g.][]{mar99, can01, eva01, sco03, tad05, bar06}. Although the hosts of most luminous AGNs are bulge-dominated, they seem to be significantly bluer than inactive elliptical galaxies and indeed show evidence for starbursts in the relatively recent past (1--2 Gyr) \\citep[e.g.,][]{kau03, san04, can06, sch06, jah07}.  In addition, both numerical simulations and observations show that merger remnants that have undergone violent relaxation have surface-brightness profiles that follow $r^{1/4}$ laws \\citep[e.g.,][]{too72,bar88,hib96}, even while still showing clear signs of the tidal interactions \\citep{rot04}. Thus, if QSO host galaxies are indeed relatively recent merger remnants, they would most likely have elliptical profiles, but they would also show fine structure indicative of past interactions.  Detecting fine structure in the already elusive QSO host galaxies is a challenging task.   Studies of nearby merger remnants show that their fine structure is often faint compared to the galaxies themselves.   For example, the total luminosity of the shells in shell galaxies accounts only for 5--15\\% of the total luminosity of the galaxy \\citep[e.g.,][]{for86,pri88,wil00}, and their detection often requires a resolution of at least $\\sim$0.5 kpc \\citep[e.g.,][]{sch92}. Therefore, detecting fine structure in QSO hosts requires high angular resolution and higher signal-to-noise (S/N) observations than those obtained in previous studies.  Thus, we are conducting a study with deep (5 orbit) HST ACS and WFPC2 observations of a sample of QSO host galaxies that are classified as ellipticals. We presented results for the first object in \\citet[][hereafter Paper I]{can07}. In the present paper, we describe results for the four remaining objects from our pilot ACS study of 5 objects. We will present results for the rest of the sample (14 additional objects) in subsequent papers. Throughout this paper, we adopt $\\Omega_{\\lambda} = 0.7$, $\\Omega_{m} = 0.3$ and $H_0$ = 71 km\\,s$^{-1}$\\,Mpc$^{-1}$. ",
        "conclusions": "\\label{mergerage} We find significant fine structure in the deep {\\em HST} ACS images of at least four of the five QSO host galaxies (PHL\\,909, PKS\\,0736+01, MC2\\,1635+119, and OX\\,169) in our pilot sample. While the prominent linear structure in OX\\,169 has been previously reported by several authors \\citep{sto78,hut84,geh84,smi86,hec86,mcl99}, the structure in the other three objects is revealed here for the first time. Only one of the objects studied, PG\\,0923+201, does not show any obvious fine structure.  Our findings are supported by the WFPC2 images of D03 and \\citet{mcl99} since some of the structure that we find is also visible in their images: Part of the north-western tidal tail in PHL\\,909 is visible (Fig.\\ A11, \\citealt{mcl99}); some of the spiral-like structure in PKS\\,0736+01 can be seen (Fig.\\ A6 in \\citealt{mcl99}); and the shell in MC\\,1635+119 labeled ``d'' in Fig.~3 of Paper I is visible to the north-east of the QSO nucleus at a distance of $\\sim$4\\arcsec~(Fig.\\ A18 in \\citealt{mcl99}). However, the depth reached in our ACS images is needed to positively identify the ``fuzz'' seen in these earlier images as real. Our results clearly show the need of high S/N imaging for an accurate interpretation of QSO host galaxy morphologies.  The host galaxy of PG\\,0923+201 does not show obvious signs of fine structure. However, the surface-brightness profile reveals a bump in the inner 2\\arcsec~that could reflect the presence of an inner disk. Interestingly, PG\\,0923+201 resides in a galaxy group of at least 6 galaxies. It is the only QSO in our sample that has been observed with the {\\em Spitzer} Space Telescope, both with the Infrared Spectrograph (IRS) and the Multiband Imaging Photometer for Spitzer (MIPS; 24$\\mu$m). Neither FIR emission nor Polycyclic-Aromatic Hydrocarbon (PAH) features have been detected \\citep{shi07}. Interestingly, this lack of signs of star-formation activity in the infrared goes hand-in-hand with the lack of obvious signs for recent merger activity.  Much of the fine structure observed in our deep {\\em HST} ACS images of the other four QSO host galaxies is indicative of merger events in the relatively recent past. The possible type and timescales for the merger in MC2\\,1635+119 were discussed in detail in Paper I. Briefly, we found that the fine structure seen in the host galaxy of this object was most likely produced by a strong tidal interaction within the last $\\sim$1\\,Gyr.   On the one hand,  the regularity of the inner shell structure is suggestive  of a near radial collision of a dwarf galaxy with a giant elliptical galaxy in a minor merger. On the other hand, the arcs observed at much larger distances than the inner shells, and other tidal debris off-axis from the direction of the encounter implied by the inner shells, are more indicative of a major merger.   Using simple $N$-body simulations, we estimated that the time scale needed to form the observed structure is less than $\\sim$1.7 Gyr (see Paper I for details). In comparison, the spectrum of the host galaxy is dominated by a population of intermediate-age ($\\sim$1.4 Gyr) stars, indicative of a strong starburst that may have occurred during the merger event. Thus, the observed QSO activity may have been triggered in the recent past either by a minor merger or by debris from an older ($\\sim$Gyr) major merger that is currently ``raining'' back into the central regions of the merger remnant.  The prominent extended structure observed in OX\\,169 is most likely a tidal tail seen edge on. Prominent tidal tails are also observed in the host of PHL\\,909. The tails in both QSO hosts are strongly suggestive of relatively recent major merger events, where at least some of the parent galaxies were disk galaxies.  The structure surrounding PKS\\,0736+201 resembles that of the Seyfert 1.5 galaxy NGC\\,5548, an Sa galaxy with ripples, tidal arms and a faint long tidal tail that has been described in detail by \\citet{sch88}. The observed structure in NGC\\,5548 has been interpreted in terms of the spatial wrapping of material from an obliquely infalling companion. When seen in projection, such material can appear in the form of relatively sharp-edged ripples \\citep{sch88}. Compared with NGC\\,5548, the faint spiral-like structure seen in PKS\\,0736+01 is more extended ($\\sim$ 50\\,kpc vs. 10\\,kpc) and more regular.  Structure similar to that observed in PKS 0736+201 can be reproduced in numerical simulations of minor mergers. For example, the merger of a dwarf elliptical galaxy with a large spiral and a total mass ratio of 1:8 results in very similar structure $\\sim$1\\,Gyr after the first passage of the two galaxies, but also again at later stages in the merger event (\\citealt{you07}; T. J. Cox, private communication).  By comparing the structure observed in these hosts to results from numerous published $N$-body simulations of mergers (e.g., \\citealt{hib95}), we know that the ages of these tidal features are likely of the order of a few orbital times. Given the masses of giant elliptical galaxies and the distances at which we observe these features, this time span ranges from a few hundred Myr to about 1 Gyr. In the cases in which we can clearly see tidal tails, we can infer that the dynamical timescales cannot be much less than $\\sim 250$ Myr, since that is roughly the amount of time that it would take for stars with typical orbital velocities ($\\sim 200$\\,km\\,s$^{-1}$) to travel in a straight line to the projected maximum distances at which we observe the features (e.g., $\\sim$ 50 kpc for PHL\\,909 or $\\sim$ 42 kpc for OX\\,169). On timescales longer than $\\sim$1 Gyr, tidal tails are expected to rapidly disperse and lose contrast.  Further evidence that these are relatively evolved mergers is the fact that enough time has elapsed since the initial tidal encounter for each host to acquire an essentially elliptical morphology. While it is possible that some of the objects, such as MC\\,1635+119 and PG\\,0923+201, may have started off as elliptical galaxies, at least OX\\,169 and PHL\\,909 must have resulted from the merger of disk galaxies.   The fact that they are at the present time reasonably well fitted by de Vaucouleurs profiles implies merger timescales on the order of a Gyr.  To address the question of whether the AGN activity may indeed have been triggered by the merger events, we need to obtain estimates of the different timescales involved. How do the ages of the stellar populations compare to the time elapsed since the merger events, and what does that imply for the QSO duty cycle? What is the connection between merger, star formation, and BH accretion?  As mentioned above, the timescales for the merger events observed in our sample is of the order of a Gyr or less.   G. Canalizo \\& A. Stockton (2008, in preparation) find similar timescales for strong starburst events in the host galaxies of these QSOs.  By modeling deep Keck spectra of these objects, they find traces of major starburst episodes in intermediate-age starburst populations that involve a substantial fraction of the stellar mass. These starburst ages range from $\\sim0.7$ Gyr (OX\\,169) to 2.2 Gyr (PKS\\,0736+017).  Hence, it seems possible that the starbursts were triggered by the merger events.  The most difficult timescale to estimate is that of the duration of the AGN activity. Current estimates of QSO lifetimes range between 10$^6$ and 10$^8$ years and are derived primarily from demographics (see review by \\citealt{mar04}). For example, from theoretical calculations \\citet{yu02} estimate a mean lifetime for a luminous QSO ($L_{\\rm bol}$ $\\ge$ 10$^{46}$\\,erg\\,s$^{-1}$) of (3--13) $\\times$ 10$^7$ yr which is comparable to the Salpeter time scale ($\\sim$ 4.5 $\\times$ 10$^7$ yr for $\\epsilon$ = 0.1 and $L/L_{\\rm edd}$ = 1; \\citealt{sal64,mar04}). Apart from demographics, there have been several attempts to deduce the AGN duty cycle for individual galaxies. For example, the length of radio jets can yield a lower estimate of QSO lifetimes, depending on the expansion speed of the jet \\citep[e.g.,][]{wil78,sch95,blu99}. However, this method cannot be applied to our sample as the two radio-loud QSOs in the sample each feature only a compact flat-spectrum radio source. Also, the size of the NLR has been proposed as a straightforward geometric measure of AGN lifetimes \\citep{mar04}. However, the NLR sizes for the five QSOs of our sample have not yet been determined.  Thus, as we do not have a direct way to estimate the time elapsed since the triggering of the nuclear activity for the objects in our sample, it is difficult to work out the physical relationship between the mergers we detect and the fueling of the QSOs.  An overly simplistic scenario, in which each QSO was triggered by the merger at the time of the starburst episode would imply QSO lifetimes on the order of a Gyr, in stark contrast to the theoretical estimates cited above.  Instead, our observations may provide evidence in support of significant time delays between tidal interactions and the fueling of the central black holes, as predicted by hydrodynamic simulations (see, e.g., \\citealt{bar98,spr05b,hop07b}).   However, while QSOs often reside in bulge-dominated galaxies like the ones presented in this study, they are also frequently found in young mergers (e.g., \\citealt{urr07,guy06,can01,hut94}). If, as the accumulating evidence suggests, mergers are indeed essential for the triggering of QSOs, then our results imply at least one of two scenarios: Either QSO activity is episodic \\citep[e.g.,][]{nor88,hop06} and occurs over longer timescales than previously speculated, or there is a large range of values for the time delays between the merger and the onset of the nuclear activity."
    }
}